{
    "0710/0710.5082.txt": {
        "abstract": "% context heading (optional) % {} leave it empty if necessary {} % aims heading (mandatory) {Photoionization models so far are unable to account for the high electron temperature $T_e$([\\ion{O}{iii}]) implied by the line intensity ratio [\\ion{O}{iii}]$\\lambda$4363\\AA/[\\ion{O}{iii}]$\\lambda$5007\\AA\\ in low-metallicity blue compact dwarf galaxies, casting doubts on the assumption of photoionization by hot stars as the dominant source of heating of the gas in these objects of large cosmological significance.} % methods heading (mandatory) {Combinations of runs of the 1-D photoionization code NEBU are used to explore alternative models for the prototype giant \\ion{H}{ii} region shell I\\,Zw\\,18\\,NW, with no reference to the filling factor concept and with due consideration for geometrical and stellar evolution constraints.} % results heading (mandatory) {Acceptable models for I\\,Zw\\,18\\,NW are obtained, which represent schematically an incomplete shell comprising radiation-bounded condensations embedded in a low-density matter-bounded diffuse medium. The thermal pressure contrast between gas components is about a factor 7. The diffuse phase can be in pressure balance with the hot superbubble fed by mechanical energy from the inner massive star cluster. The failure of previous modellings is ascribed to (1) the adoption of an inadequate small-scale gas density distribution, which proves critical when the collisional excitation of hydrogen contributes significantly to the cooling of the gas, and possibly (2) a too restrictive implementation of Wolf-Rayet stars in synthetic stellar cluster spectral energy distributions. A neutral gas component heated by soft X-rays, whose power is less than 1\\% of the star cluster luminosity and consistent with CHANDRA data, can explain the low-ionization fine-structure lines detected by SPITZER. [O/Fe] is slightly smaller in I\\,Zw\\,18\\,NW than in Galactic Halo stars of similar metallicity and [C/O] is correlatively large.} % conclusions heading (optional), leave it empty if necessary {Extra heating by, \\eg, dissipation of mechanical energy is not required to explain $T_e$([\\ion{O}{iii}]) in I\\,Zw\\,18. Important astrophysical developments are at stakes in the 5\\% uncertainty attached to \\oiii\\ collision strengths.} ",
        "introduction": "\\label{intro}  The optical properties of Blue Compact Dwarf (BCD) galaxies are similar to those of Giant Extragalactic \\hii\\ Regions (GEHIIR). Their blue continuum arises from one or several young Massive Star Clusters (MSC), which harbour extremely large numbers of massive stars.  BCDs are relatively isolated, small-sized, metal-poor galaxies (Kunth \\& \\\"Ostlin 2000) and may be the rare `living fossils' of a formerly common population. BCDs can provide invaluable pieces of information about the primordial abundance of helium (\\eg, Davidson \\& Kinman 1985), the chemical composition of the InterStellar Medium (ISM, \\eg, Izotov et al. 2006), the formation and evolution of massive stars, and the early evolution of galaxies at large redshift. Among them, \\IZ\\ stands out as one of the most oxygen-poor BCDs known (\\eg, Izotov et al. 1999) and a young galaxy candidate in the Local Universe (\\eg, Izotov \\& Thuan 2004).  The line emission of \\hii\\ regions is believed to be governed by radiation from massive stars, but spectroscopic diagnostics most often indicate spatial fluctuations of the electron temperature \\Te\\ (see the dimensionless parameter $t^2$, Peimbert, 1967), that appear larger than those computed in {\\sl usual} photoionization models, suggesting an {\\sl extra heating} of the emitting gas (\\eg, Peimbert, 1995; Luridiana et al. 1999). Until the cause(s) of this failure of photoionization models can be identified, a sword of Damocles is hung over a basic tool of astrophysics.  Tsamis \\& P\\'equignot (2005) showed that, in the GEHIIR 30\\,Dor of the LMC, the various \\Te\\ diagnostics could be made compatible with one another if the ionized gas were {\\sl chemically inhomogeneous} over small spatial scales. A pure photoionization model could then account for the spectrum of a bright filament of this nebula. Although this new model needs confirmation, it is in suggestive agreement with a scenario by Tenorio-Tagle (1996) of a recycling of supernova ejecta through a rain of metal-rich droplets cooling and condensing in the Galaxy halo, then falling back on to the Galactic disc and incorporating into the ISM without significant mixing until a new \\hii\\ region eventually forms. If this class of photoionization models is finally accepted, extra heating will not be required for objects like 30\\,Dor, with near Galactic metallicity.  Another problem is encountered in low-metallicity (`low-Z') BCDs (Appendix~A). In BCDs, available spectroscopic data do not provide signatures for $t^2$'s, but a major concern of photoionization models is explaining the high temperature \\Te(\\oiii) infered from the observed intensity ratio $r$(\\oiii) = \\oiii\\la4363/(\\oiii\\la5007+4959). Thus, Stasi\\'nska \\& Schaerer (1999, SS99) conclude that photoionization by stars fails to explain $r$(\\oiii) in the GEHIIR \\IZ\\,NW and that photoionization must be supplemented by other heating mechanisms. A requirement for extra heating is indirectly stated by Luridiana et al. (1999) for NGC\\,2363.  A possible heating mechanism is conversion of mechanical energy provided by stellar winds and supernovae, although a conclusion of Luridiana et al. (2001) does not invite to optimism. A limitation of this mechanism is that most of this mechanical energy is likely to dissipate in hot, steadily expanding superbubbles (Martin, 1996; Tenorio-Tagle et al. 2006). It is doubtful that heat conduction from this coronal gas could induce enough localized enhancement of \\Te\\ in the photoionized gas (\\eg, Maciejewski et al. 1996), even though Slavin et al. (1993) suggest that turbulent mixing may favour an energy transfer. Martin (1997) suggests that shocks could help to explain the trend of ionization throughout the diffuse interstellar gas of BCDs, but concedes that ``shocks are only being invoked as a secondary signal in gas with very low surface brightness''. Finally, photoelectric heating from dust is inefficient in metal-poor hot gas conditions (Bakes \\& Tielens 1994).  Nevertheless, the conclusion of SS99 is now accepted in many studies of GEHIIRs. It entails so far-reaching consequences concerning the physics of galaxies at large redshifts as to deserve close scrutiny. If, for exemple, the difference between observed and computed \\Te(\\oiii) in the model by SS99 were to be accounted for by artificially raising the heat input proportionally to the photoionization heating, then the total heat input in the emitting gas should be doubled. {\\sl This problem therefore deals with the global energetics of the early universe}.  After reviewing previous models for \\IZ\\,NW (Sect.~\\ref{prev_IZ}), observations and new photoionization models are described in Sects.~\\ref{obs} \\& \\ref{newmod}. Results presented in Sect.~\\ref{res} are discussed in Sect.~\\ref{disc}. Concluding remarks appear in Sect.~\\ref{concl}. Models for other GEHIIRs are reviewed in Appendix~A. Concepts undelying the new photoionization models are stated in Appendices~B and C. ",
        "conclusions": "\\label{concl}  Owing to its small heavy element content, \\IZ\\ stands at the high-\\Te\\ boundary of photoionized nebulae. Where ionization and temperature are sufficiently high, the cooling is little dependent on conditions, except through the concentration of H$^0$, controlled by density. Therefore, in the photoionization model logics, {\\sl \\Te\\ is then a density indicator}, in the same way it is an O/H indicator in usual \\hii\\ regions. It is for not having recognized implications of this new logics, that low-metallicity BCD models failed.  In a photoionization model study of \\IZ\\,NW, SS99 employed a  filling-factor description and concluded that \\Te(\\oiii) was fundamentally unaccountable. The {\\sl vogue} for this simple description of the ionized gas distribution resides in its apparent success for usual \\hii\\ regions, a success falling in fact to the strong dependence of cooling on abundances. Universally adopted along past decades, this concept led all authors to conclude that photoionization by hot stars did not provide enough energy to low-Z GEHIIRs. This conclusion is in line with a movement of calling into question photoionization by stars as the overwhelmingly dominant source of heat and ionization in gaseous nebulae, a movement cristalizing on the `$t^2$ problem' (Esteban et al. 2002; Peimbert et al. 2004), since the presence of \\Te\\ fluctuations {\\sl supposedly larger} than those reachable assuming photoionization by stars implies additional heating.  A conclusion of the present study is that the gas distribution is no less critical than the radiation source in determining the line spectrum of \\hii\\ regions. Assuming pure photoionization by stars, the remarkable piece of information carried by the large \\Te(\\oiii) of \\IZ\\,NW is that the mean density of the \\oiii\\ emitting region is much less than \\Ne(\\sii), a low \\Ne\\ confirmed by line ratios \\oiv\\la25.9$\\mu$/\\heii\\la4686 and \\feiii\\la4986/\\feiii\\la4658. \\IZ\\,NW models comprising a plausible SED and respecting geometrical constraints can closely match almost all observed lines from UV to IR, including the crucial \\oiii\\la4363 (\\siv\\la10.5$\\mu$ is a factor 2 off, however). Thus, extra heating by, \\eg, dissipation of mechanical energy in the photoionized gas of low-metallicity BCD galaxies like \\IZ\\ is {\\sl not} required to solve the `\\Te(\\oiii) problem'. Moreover, since low-ionization fine-structure lines can be explained by soft X-rays, (hydrodynamical) heating is {\\sl not even} required in warm \\hi\\ regions protected from ionization and heating by star radiation.  As a final note, on close scrutiny, the solutions found here {\\sl tend to be just marginally consistent} with observed $r$(\\oiii). Given the claimed accuracy in the different fields of physics and astrophysics involved, postulating a mechanical source of heating is premature, whereas a {\\sl 2--3\\% upward correction} to the collision strength for transition O$^{2+}$($^3$P\\,--\\,$^1$S) at \\Te~$\\sim$~2$\\times$10$^4$\\,K is an alternative worth exploring by atomic physics. Yet, another possibility is a substantial increase of the distance to \\IZ. Owing to accurate spectroscopy and peculiar conditions in \\IZ, important astrophysical developments are at stakes in the 5\\% uncertainty attached to \\oiii\\ collision strengths.  If photoionized nebulae are shaped by shocks and other hydrodynamic effects, this does not imply that the emission-line intensities are detectably influenced by the thermal energy deposited by these processes. Unravelling this extra thermal energy by means of spectroscopic diagnostics and models is an exciting prospect, whose success depends on a recognition of all resources of the photoionization paradigm. Adopting the view that photoionization by radiation from young hot stars, including WR stars, is the only excitation source of nebular spectra in BCD galaxies, yet without undue simplifications, may well be a way to help progresses in the mysteries of stellar evolution, stellar atmosphere structure, stellar supercluster properties, giant \\hii\\ region structure and, {\\sl at last}, possible extra sources of thermal energy in BCDs."
    },
    "0710/0710.4400_arXiv.txt": {
        "abstract": "{} {Bright HNC 1--0 emission, rivalling that of HCN 1--0, has been found towards several Seyfert galaxies. This is unexpected since traditionally HNC is a tracer of cold (10 K) gas, and the molecular gas of luminous galaxies like Seyferts is thought to have bulk kinetic temperatures surpassing 50 K. There are four possible explanations for the bright HNC: (a) Large masses of hidden cold gas; (b) chemistry dominated by ion-neutral reactions; (c) chemistry dominated by X-ray radiation; and (d) HNC enhanced through mid-IR pumping. In this work we aim to distinguish the cause of the bright HNC and to model the physical conditions of the HNC and HCN emitting gas.} {We have used SEST, JCMT and IRAM 30m telescopes to observe HNC 3--2 and HCN 3--2 line emission in a selection of 5 HNC-luminous Seyfert galaxies. We estimate and discuss the excitation conditions of HCN and HNC in NGC~1068, NGC~3079, NGC~2623 and NGC~7469, based on the observed 3--2/1--0 line intensity ratios. We also observed CN 1--0 and 2--1 emission and discuss its role in photon and X-ray dominated regions.} {HNC 3--2 was detected in 3 galaxies (NGC~3079, NGC~1068 and NGC~2623). Not detected in NGC~7469. HCN 3--2 was detected in NGC~3079, NGC~1068 and NGC~1365, it was not detected in NGC~2623. The HCN 3--2/1--0 ratio is lower than 0.3 only in NGC~3079, whereas the HNC 3--2/1--0 ratio is larger than 0.3 only in NGC~2623. The HCN/HNC 1--0 and 3--2 line ratios are larger than unity in all the galaxies. The HCN/HNC 3--2 line ratio is lower than unity only in NGC~2623, which makes it comparable to galaxies like Arp~220, Mrk~231 and NGC~4418. } {We conclude that in three of the galaxies the HNC emissions emerge from gas of densities $n\\lesssim10^5~\\3cm$, where the chemistry is dominated by ion-neutral reactions. The line shapes observed in NGC~1365 and NGC~3079 show that these galaxies have no circumnuclear disk. In NGC~1068 the emission of HNC emerges from lower ($<10^5~\\3cm$) density gas than HCN ($>10^5~\\3cm$). Instead, we conclude that the emissions of HNC and HCN emerge from the same gas in NGC~3079. The observed HCN/HNC and CN/HCN line ratios favor a PDR scenario, rather than an XDR one, which is consistent with previous indications of a starburst component in the central regions of these galaxies. However, the $N({\\rm HNC})/N({\\rm HCN})$ column density ratios obtained for NGC~3079 can be found only in XDR environments.} ",
        "introduction": "The hydrogene cyanide, HCN molecule, is commonly used as an extragalactic tracer of molecular gas with densities $n$(H$_2$) larger than $10^4~\\3cm$ (e.g. Solomon et al. \\cite{solomon92}; Curran et al. \\cite{curran00}; Kohno \\cite{kohno05}). The HCN to CO intensity ratio varies significantly, from 1/3 to 1/40 in starburst galaxies, and it has not been determined whether this variation depends on the dense molecular gas content or on the abundance and/or excitation conditions. In addition, recent results seems to indicate that HCN may not be an unbiased tracer of the dense molecular gas content in LIRGs and ULIRGs (Graci\\'a-Carpio et al. \\cite{gracia06}). It is essential, therefore, to use other molecular tracers than HCN, in order to understand the physical conditions in the dense gas.  A molecule of particular interest, for comparison with HCN, is its isomer HNC. The detection of interstellar HNC supports the theory of dominant ion-molecule chemistry in dark molecular clouds. Both species are thought to be created by the same dissociative recombination of HCNH$^+$. This ion can produce HCN and HNC, with approximately equal abundances. Models based only on this scheme would predict then an HNC/HCN ratio $\\approx 1$. However, the observed HNC/HCN abundance ratios vary significantly between different kinds of molecular clouds - the ratio ranges from 0.03 to 0.4 in warm cores ($T_k > 15$ K), and can be as high as 4.4 in cold cores ($T_k < 15$ K).  The CN (cyanogen radical) molecule is another tracer of dense gas, with a lower (by a factor of 5) critical density than HCN. CN is also chemically linked to HCN and HNC by photodissociations (e.g. Hirota \\& Yamamoto \\cite{hirota99}). Surveys of the 1--0 transition of CN and HNC have been done in order to trace a cold, dense phase of the gas in luminous galaxies (Aalto et al. \\cite{aalto02}). It was found that the HNC 1--0 luminosities often rivalled those of HCN 1--0. These results seem to contradict the idea of warm ($T_k \\gtrsim 50$ K) gas in the centers of luminous galaxies (e.g. Wild et al. \\cite{wild92}; Wall et al. \\cite{wall93}) whose IR luminosities were suggested to originate from star formation rather than AGN activity (Solomon et al. \\cite{solomon92}).  \\begin{table*}[!t]  \\begin{minipage}{18cm} \\caption[]{Sample of galaxies$^{~\\rm a}$.} \\label{tab:galaxies} \\centering \\begin{tabular}{lccccccc} \\hline \\noalign{\\smallskip} Galaxy & Seyfert  & RA & DEC & v$_{sys}$ & Distance$^{~\\rm b}$ & $\\Omega_S$(CO)$^{~\\rm c}$ & $\\Omega_S$(HCN)$^{~\\rm d}$\\\\ &     & [hh mm ss] & [$^{\\circ}~{}'~{}''$] & [\\kms] & [Mpc] & [$''~^2$] & [$''~^2$]\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} NGC~3079 & 2   & 10 01 57.805 & +55 40 47.20 & 1116$\\pm$1 & 15.0$\\pm$1.1 & $15\\times 7.5$ & $5\\times 5$ \\\\ NGC~1068 & 2   & 02 42 40.711 & -00 00 47.81 & 1137$\\pm$3 & 15.3$\\pm$1.1 & $30\\times 30$ & $10\\times 10$ \\\\ NGC~2623 & 2   & 08 38 24.090 & +25 45 16.80 & 5549$\\pm$1 & 74.6$\\pm$5.4 & $8\\times 8$ & $2.6\\times 2.6$ \\\\ NGC~1365 & 1.8 & 03 33 36.371 & -36 08 25.45 & 1636$\\pm$1 & 22.0$\\pm$1.6 & $50\\times 50$ & $16.5\\times 16.5$ \\\\ NGC~7469 & 1.2 & 23 03 15.623 & +08 52 26.39 & 4892$\\pm$2 & 65.8$\\pm$4.8 & $8\\times 8$ & $4\\times 6$ \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\begin{list}{}{} \\item[${\\mathrm{a}}$)] The Seyfert classification, positions (in equatorial J2000 coordinates) and heliocentric radial velocities were taken from NED.  \\item[${\\mathrm{b}}$)] The distances were calculated using the Hubble constant (H$_0 \\approx 74.37$ \\kms~Mpc$^{-1}$) estimated by Ngeow and Kanbur (\\cite{ngeow06}).  \\item[${\\mathrm{c}}$)] The source sizes of the CO 1--0 transition line were estimated from the maps presented in (Koda et al. \\cite{koda02}) for NGC~3079, (Schinnerer et al. \\cite{schinnerer00}) for NGC~1068, (Bryant et al. \\cite{bryant99}) for NGC~2623, (Sandqvist \\cite{sandqvist99}) for NGC~1365, and (Papadopoulos \\& Allen \\cite{papadopoulos00}) for NGC~7469. The source sizes for the $J = 2-1$ transition line were assumed equal to that of the $J = 1-0$ line. For NGC~7469, the source size of the CO 2--1 emission estimated from the corresponding map presented by Davies et al. (\\cite{davies04}) agrees well with the source size estimated for the CO 1--0 line.  \\item[${\\mathrm{d}}$)] Source sizes of HCN 1--0 were estimated from the corresponding maps published in (Kohno et al. \\cite{kohno00}) for NGC~3079, (Kohno et al. \\cite{kohno01} and Helfer \\& Blitz \\cite{helfer95}) for NGC~1068, (Davies et al. \\cite{davies04}) for NGC~7469. The source sizes of NGC~2623 and NGC~1365 were estimated using proportions found in NGC~1068 (read text in \\S3.5). Because of their chemical link, the source sizes of the CN and HNC molecules were assumed the same as that of HCN. Due to the lack of high resolution maps, the source sizes corresponding to the emission of the higher transition lines were assumed equal to that of the $J = 1-0$ line. \\end{list}  \\end{minipage}  \\end{table*}  According to observations in the vicinity of the hot core of Orion KL, experimental data, chemical steady state and shock models, the HNC/HCN ratio decreases as the temperature and density increase (e.g. Schilke et al. \\cite{schilke92}; Talbi et al. \\cite{talbi96}; Tachikawa et al. \\cite{tachikawa03}). If a bright HNC 1--0 transition line is nevertheless detected under these conditions, it could be due to the following possible explanations: (a) the presence of large masses of hidden cold gas and dust at high densities ($n > 10^5~\\3cm$); (b) chemistry dominated by ion-molecule reactions with HCNH$^+$ at low density ($n \\approx 10^4~\\3cm$) in regions where the temperature dependence of the HNC abundance becomes weaker; (c) enhancement by mid-IR pumping, also in low density regions where the lines would not be collisionally excited; and (d) the influence of UV-rays in Photon Dominated Regions (PDRs) and/or X-rays in X-ray Dominated Regions (XDRs) at densities $n \\gtrsim 10^4~\\3cm$ and at total column densities $N_{\\rm H} > 3\\times 10^{21}~\\2cm$ (Meijerink \\& Spaans \\cite{meijerink05}).  In the case of CN, observations of its emission towards the Orion A molecular complex (Rodr\\'iguez-Franco et al. \\cite{rodriguez98}) suggest that this molecule is also enhanced in PDRs, but particularly in XDRs, where a CN/HCN abundance ratio larger than unity is expected (e.g., Lepp \\& Dalgarno \\cite{lepp96} and Meijerink, Spaans, Israel \\cite{meijerink07}).  We have observed low and high transition lines of the HCN, HNC and CN molecules in a group of Seyfert galaxies, which are supposed to host both sources of power, AGN and starburst activity, in their central region. Our interest is to assess the excitation conditions of HCN and HNC, distinguish between the above possible causes of the bright HNC, and to explore the relation between the CN emission, XDRs and dense PDRs in these sources.  In \\S2 we describe the observations. The results (spectral lines, line intensities and line ratios) are presented in \\S3. The interpretation of line shapes and gas distribution in the most relevant cases, as well as the possible explanations for the bright HNC and the modelling of the excitation conditions of HCN and HNC are discussed in \\S4. The conclusions and final remarks of this work are presented in \\S5. ",
        "conclusions": "We have used the SEST and JCMT telescopes to carry out a survey of CN 2--1, HCN 3--2 and HNC 3--2 line emission in a sample of 4 Seyfert galaxies, plus NGC~3079 which was observed with the IRAM 30m telescope. The conclusions we draw are as follows: \\begin{list}{}{} \\item[1)] We detected HNC 3--2 emission in 3 of the 5 galaxies, while we obtain an upper limit for one of them (NGC~7469). HCN 3--2 was also detected in 3 galaxies (NGC~3079, NGC~1068 and NGC~1365), while it was not detected in NGC~2623. CN 2--1, along with the spingroups (\\textit{J} = 5/2 -- 3/2, \\textit{F} = 7/2 -- 5/2) and (\\textit{J} = 3/2 -- 1/2, \\textit{F} = 5/2 -- 3/2) was also detected in NGC~3079, NGC~1068 and NGC~1365. \\item[2)] {The line shapes} observed in NGC~1365 and NGC~3079 suggests that there is no circumnuclear disk in these galaxies. \\item[3)] We find that in 3 of the galaxies the HNC 3--2/1--0 line ratios suggest that the HNC emissions emerge from gas of densities $n\\lesssim10^5~\\3cm$, where the chemistry is dominated by ion-neutral reactions. In NGC~2623 a model of large masses of hidden cold (10 K) gas and dust, as well as a chemistry dominated by ion-neutral reactions, are yet to be distinguished as the correct interpretation for the bright HNC observed in this galaxy. \\item[4)] The 3--2/1--0 line ratios and the modelled excitation conditions imply that the HNC emission emerges from a more diffuse ($n<10^5~\\3cm$) gas region than the HCN emission ($n>10^5~\\3cm$) in NGC~1068, whereas they emerge from the same lower density ($n\\lesssim10^5~\\3cm$) gas in NGC~3079. \\item[5)] The HCN/HNC and CN/HCN line ratios tentatively favor a PDR scenario, rather than an XDR one, in the 3 Seyfert galaxies where we have CN, HNC and HCN data. The $N({\\rm HNC})/N({\\rm HCN})$ column density ratios obtained for NGC~3079 can be found only in XDR environments. \\end{list}  In order to complete the sample, we plan to observe HCN 3--2 and CN 2--1 in NGC~7469, CN 2--1 in NGC~2623 and HNC 3--2 in NGC~1365. We plan to perform high resolution observations to further study the distribution and source sizes of CN and HNC.  Modeling of the collision data for the CN molecule would be useful to estimate the $N({\\rm CN})/N({\\rm HCN})$ column density ratio, which would complement the $N({\\rm HNC})/N({\\rm HCN})$ ratio in order to have a more sophisticated tool to estimate and distinguish the prevalent environment conditions of the high density gas in the nuclear region of Seyfert galaxies.  The AGN contribution (through XDR effects) is typically of a small angular scale and can be seriously affected by beam dilution at the transition lines studied in this work. On the other hand, the starburst contribution is of a larger angular scale than the AGN, and it effects can be contaminating our observations, and hence leading to the favored PDR scenario found with our models. Hence, our suggested interpretations could change if we zoom in on these sources. Therefore, high resolution maps of HNC and CN molecules are necessary to complement those of HCN, and to do a more accurate estimate of molecular abundances and line intensity ratios, which take source size into account. Observations of the higher transition lines (e.g. $J$=4--3) can also aid to disentangle the effects of the AGN and the starburst ring, due to the smaller beam size obtained at higher frequencies."
    },
    "0710/0710.3685_arXiv.txt": {
        "abstract": "{ Relic neutralinos produced after the Big Bang are favoured candidates for Dark Matter. They can accumulate at the centre of massive celestial objects like our Sun. Their annihilation can result in a high-energy neutrino flux that could be detectable as a localised emission with earth-based neutrino telescopes like ANTARES. In this paper a brief overview of the prospects of the indirect search for Dark Matter particles with the ANTARES detector will be given. The analysis method and expected performance for the detection of the expected neutrinos will be discussed. \\PACS{ {95.35.+d}{Dark matter (stellar, interstellar, galactic, and cosmological)}   \\and {95.55.Vj}{Neutrino, muon, pion, and other elementary particle detectors; cosmic ray detectors} } % } % ",
        "introduction": "\\label{intro} It is now a well-established fact that according to our present theory of gravity, 85\\%~of the matter content of our universe is missing. Observational evidence for this discrepancy ranges from macroscopic to microscopic scales, e.g. gravitational lensing in galaxy clusters, galactic rotation curves and fluctuations measured in the Cosmic Microwave Background. This has resulted in the hypothesised existence of a new type of matter called Dark Matter. Particle physics provides a well-motivated explanation for this hypothesis: The existence of (until now unobserved) massive weakly interacting particles (WIMPs). A favorite amongst the several WIMP candidates is the neutralino, the lightest particle predicted by Supersymmetry, itself a well-motivated extension of the Standard Model.  If Supersymmetry is indeed realised in Nature, Supersymmetric particles would have been copiously produced at the start of our Universe in the Big Bang. Initially these particles would have been in thermal equilibrium. After the temperature of the Universe dropped below the neutralino mass, the neutralino number density would have decreased exponentially. Eventually the expansion rate of the Universe would have overcome the neutralino annihilation rate, resulting in a neutralino density in our Universe today similar to the cosmic microwave background.  These relic neutralinos could then accumulate in massive celestial bodies in our Universe like our Sun through weak interactions with normal matter and gravity. Over time the neutralino density in the core of the object would increase considerably, thereby increasing the local neutralino annihilation probability. In the annihilation process new particles would be created, amongst which neutrinos. This neutrino flux could be detectable as a localised emission with earth-based neutrino telescopes like ANTARES.  This paper gives a brief overview of the prospects for the detection of neutrinos originating from neutralino annihilation in the Sun with the ANTARES neutrino telescope.  \\begin{figure}[b] \\center{ \\includegraphics[width=0.45\\textwidth,angle=0]{NEA_60kHz0XOFF_off.png} \\caption{The ANTARES Neutrino Effective Area vs. $E_\\nu$.} \\label{fig:1}       % } \\end{figure}  \\begin{figure*}[t] \\center{ \\includegraphics[width=0.8\\textwidth,angle=0]{psflux.png} \\caption{Predicted $\\nu_\\mu+\\bar{\\nu}_\\mu$ flux from the Sun in mSUGRA parameter space.} \\label{fig:2}       % } \\end{figure*} ",
        "conclusions": "Nearly half of the ANTARES detector has been operational since January this year. The detector is foreseen to be completed in early 2008. The data show that the detector is working within the design specifications.  As can be seen from Fig.~\\ref{fig:4}, mSUGRA models that are excludable by ANTARES at 90\\%~CL are found in the Focus Point region. The same models should also be excludable by future direct detection experiments, as is shown in Fig.~\\ref{fig:7}.  To improve the ANTARES sensitivity, a directional trigger algorithm has recently been implemented in the data acquisition system. In this algorithm, the known position of the potential neutrino source is used to lower the trigger condition. This increases the trigger efficiency, resulting in a larger $A_{\\rm eff}^{\\nu}$. In Fig.~\\ref{fig:8}, the $A_{\\rm eff}^{\\nu}$ at the trigger level for the standard- and the directional trigger algorithm are shown in black (``{\\em trigger3D}'') and red (``{\\em triggerMX}'') respectively."
    },
    "0710/0710.2360_arXiv.txt": {
        "abstract": "{ The similarity of the observed mass densities of baryons and cold dark matter may be a sign they have a related origin. The baryon-to-dark matter ratio can be understood in the MSSM with right-handed (RH) neutrinos if CDM is due to a d = 4 flat direction condensate of very weakly coupled RH sneutrino LSPs and the baryon asymmetry is generated by Affleck-Dine leptogenesis along a d = 4 $\\left(H_{u}L\\right)^2$ flat direction. Observable signatures of the model include CDM and baryon isocurvature perturbations and a distinctive long-lived NLSP phenomenology. \\PACS{ {98.80.Cq}{Cosmology} } % } % ",
        "introduction": "\\label{intro}    A striking feature of the observed Universe is the similar mass density in baryons and cold dark matter. (The 'Baryon-to-Dark Matter' (BDM) ratio.) From the WMAP three-year results for the $\\Lambda$CDM model, $\\Omega_{DM}/\\Omega_{B} = 5.65 \\pm 0.58$ \\cite{wmap}. However, in most models the physics of baryogenesis and dark matter production are unrelated. So why is the mass density in baryons within an order of magnitude of that of dark matter? There are three possibilities: \\newline (i)  A remarkable coincidence. \\newline (ii) Some anthropic selection mechanism, usually assumed but undefined (e.g. in the case of thermal relic neutralino dark matter). \\newline (iii) The mechanisms for the origin of the baryon asymmetry and dark matter are related.  The latter possibility seems the simplest interpretation of the BDM ratio. Indeed, we may be ignoring a {\\it big clue} to the nature of the correct BSM particle theory. It is highly non-trivial for a particle physics theory to have within its structure (without contrivance) a mechanism that can account for the BDM ratio. Therefore if the BDM ratio is due to such a mechanism it would provide us with a powerful principle by which to select the best canadiate particle physics models.   BDM models broadly divide into two classes: \\newline {\\bf 1).Charge conservation based:} The dark matter particle and baryon number are related by a conserved charge, $Q_{B} + Q_{cdm} = 0 \\Rightarrow n_{cdm} \\sim n_{B}$. The CDM particle mass satisfies $ m_{cdm} = m_{n}n_{B}/n_{cdm}$ and so $m_{cdm} \\sim 1 GeV$ is necessary. However, this does fit well with SUSY if the LSP mass is $O(m_{W})$ or larger. \\newline {\\bf 2). Dynamics based:} In this case the dark matter and baryon densities are related by {\\it similar} physical mechanisms for their origin. This implies a less rigid relation between $n_{B}$ and $n_{CDM}$, which may allow us to understand why it is the mass rather than number densities that are observed to be similar. ",
        "conclusions": "RH sneutrino condensate CDM combined with Affleck-Dine baryogenesis can plausibly account for the observed similarity of the baryon and dark matter mass densities in the Universe. Seen as a selection principle, the requirement that a particle physics model can {\\it without contrivance} account for the BDM ratio favours the MSSM with neutrino masses and  RH sneutrino condensate CDM. It is quite remarkable that the MSSM with neutrino masses has the ability to account for the BDM ratio as a natural consequence its structure.  CDM and baryon isocurvature perturbations are possible, the ratio of which gives information on the nature of SUSY inflation. Long-lived NLSP phenomenology is also expected, which can be distinguished from gravitino and axino LSP phenomenology via trapped stau final states, and from thermal relic RH sneutrino LSP phenomenology once the parameters of the MSSM are known.  There are a number of issues which remain to be addressed. One is the possibility that $M_{N} \\neq 0$. In this case it is possible that the heavier generation condensate RH sneutrinos, which will have a lifetime longer than the age of the Universe so long as $\\lambda_{\\nu}$ is still small,  would decay into the LSP RH sneutrino plus $e^{+}e^{-}$. This could produce a potentially observable diffuse $\\gamma$-ray background. In addition, the phenomenology and cosmology of MSSM-LSPs in this model should be studied in detail."
    },
    "0710/0710.2683_arXiv.txt": {
        "abstract": "We study the dynamics of a twisted tilted disc under the influence of an external radiation field. Assuming the effect of absorption and reemission/scattering  is that a pressure is applied to the disc surface where the local optical depth is of order unity, we determine the response of the vertical structure and the influence it has on the possibility of instability to warping.  We derive a pair of equations describing the evolution of a  small tilt as a function of radius in the small amplitude regime that applies to both the diffusive and bending wave regimes. We also study the non linear vertical response of the disc numerically using an analogous  one dimensional slab model. For global warps, we find that in order for the disc vertical structure to respond as a quasi uniform shift or tilt, as has been assumed in previous work,  the product of the ratio of the  external radiation momentum flux to the local disc mid plane pressure, where it is absorbed,  with the disc aspect ratio should be significantly less than unity. Namely, this quantity should be of the order of or smaller than the ratio of the disc gas density corresponding to the layer intercepting radiation to the mid plane density, $\\lambda \\ll 1$.  When this condition is not satisfied the disc surface tends to adjust so that the local normal becomes perpendicular to the radiation propagation direction. In this case dynamical quantities determined by the disc twist and warp tend to oscillate with a large characteristic period $T_{*}\\sim \\lambda^{-1}T_{K}$, where $T_{K}$ is some 'typical' orbital period of a gas element in the disc. The possibility of warping instability then becomes significantly reduced.  In addition, when the vertical response is non uniform, the  possible  production of shocks may lead to an important dissipation mechanism. ",
        "introduction": "\\noindent  A significant number of X-ray binaries  exhibit long-term  periodicities on time-scales of ~10-100 d. Examples are Her X-1, SS 433 and LMC X-4  see e.g. Clarkson et al. 2003. Precession and warping  of a tilted accretion disc has been proposed as an explanation (Katz 1973; Petterson 1975) which has subsequently  been  found to have observational support (Clarkson et al. 2003). The effect of radiation pressure  on a twisted tilted accretion disc was first considered by Petterson 1977a,b who noted that when radiation from the central source is absorbed at the disc surface and  re-emitted, a potentially important torque will result. Iping \\& Petterson 1990 later suggested that such torques determined the shape of the disc and its precession rate. Pringle 1996 subsequently showed that an initially axisymmetric thin disc could be unstable to  warping as a result of interaction with an external radiation field (see e.g.  Maloney et a. 1998; Ogilvie \\& Dubus 2001 and Foulkes et al. 2006 for later developments).  The analysis considered the disc to behave as a collection of rings, which could interact by  transferring  angular momentum through the action of viscosity, but otherwise behaved as if they were rigid. In particular, the vertical displacement or tilt was assumed to be essentially uniform and independent of the vertical coordinate. Thus the pressure applied at the surface at optical depth unity is assumed to be effectively communicated through the disc vertical structure so as to give a near uniform response. In this paper we extend the theoretical treatment of the interaction of a disc with an external radiation field to take account of possible significant departures of the tilt or displacement from uniformity in the vertical direction. One of our objectives is to determine the conditions under which the assumption of a uniform response is valid, and then to estimate some of the consequences when they are not satisfied, including potential additional dissipation resulting from non linear effects such as the production of shock waves. The latter is done using a one dimensional slab analogue model with the required high resolution in the vertical direction.  Although  we focus on the effects of a surface pressure induced by an external radiation field, very similar considerations are likely to apply when warps are induced by a surface pressure resulting from interaction with an external  wind (e.g. Quillen, 2001) or surface forces produced by the interaction of the disc with an external magnetic field originating in the central star ( e.g. Pfeiffer \\& Lai, 2004).  We begin by considering the relevant issues using simple physical arguments. Following  Pringle 1996 we consider a thin disc immersed in an external radiation field with mid plane initially coinciding with a Cartesian $(x,y)$ plane. It is supposed that the upper surface is parallel to the external rays so that there is initially no interaction with the radiation.  The upper surface is then given a vertical elevation $h(r,\\phi),$ where we now use polar coordinates $(r,\\phi).$ As a result, a disc element with surface area $d{\\cal A}$ absorbs  momentum from the radiation field at a rate \\be  {\\dot {\\cal F}} = F_0  d{\\cal A} \\left(r{\\partial (h/r)\\over \\partial r}\\right).\\ee Here $ F_0$ is the momentum flux at radius $r$ associated with the external radiation field. The factor in brackets gives the angle through which the local normal is rotated as a result of the perturbation (see Appendix \\ref{A} for more details). Assuming the absorbed momentum is reradiated isotropically above the disc by the surface layers at an  optical depth of unity, there will be an applied pressure there of magnitude \\be  {\\cal P} = {2F_0\\over 3}   \\left(r{\\partial (h/r)\\over \\partial r}\\right).\\ee This external pressure, when applied to a complete elementary ring, produces a net torque per unit  length of magnitude \\be { d {\\cal T} \\over dr} = {2\\pi F_0 \\over 3} \\left(r^2{\\partial (h/r)\\over \\partial r}\\right){\\bf l},\\ee where the complex  vector ${\\bf l}$ has  components equal to $(i,1,0)$ in the Cartesian coordinate system. Here, in performing   the azimuthal integration,  we take into account that the azimuthal dependence of $h$ is through a factor of the form $ \\exp(-i\\phi)$ and work with the radial amplitude from now on.  To find the consequent evolution of the disc, one requires the component of the  angular momentum per unit length perpendicular to its unperturbed direction. This is given by  \\be {d{\\cal J}/dr} = 2\\pi\\Sigma r^2 \\Omega {\\langle h \\rangle}i{\\bf l},\\ee with $\\Omega$ and $\\Sigma $  being the near Keplerian local disc angular velocity and the disc surface density, respectively. $\\langle h \\rangle =\\int  d\\zeta \\rho h /\\Sigma $ where $\\zeta $ is a vertical coordinate and $\\rho $ is the gas density.  For an elementary ring the condition that the rate of change of angular momentum equal the applied torque gives a tilt evolution equation of the form \\be {\\partial {\\langle h \\rangle}\\over \\partial t}= -i{ F_0 \\over 3\\Omega\\Sigma} \\left({\\partial (h/r)\\over \\partial r}\\right). \\ee Assuming the vertical response is uniform, so that we can set $h= {\\langle h \\rangle} \\equiv r{\\bf W}$, we obtain a description of warp evolution equivalent to Pringle 1996 when effects due to viscosity and bending wave propagation are neglected (see sections \\ref{simpa} and \\ref{Dyneq} below). This description indicates the possibility of instability to radiation pressure warping.  However, here we stress the fact that the averaged  elevation $\\langle h \\rangle ,$ enclosed in angled brackets, applies to the bulk of the inertia of the disc and can therefore be shown to be close to the elevation of the mid plane. On the other hand $h$ as used in the torque formula applies to the disc surface elevation. An important aspect of this paper is to distinguish these two elevations and investigate under what conditions they can be taken to be equal. This would be possible if the vertical structure responds as a rigid body to the external pressure forcing and we find the conditions for this to occur. The general requirement is found to be that the density at the surface of the disc where the pressure is applied should not be too small.   It is possible to estimate in a simple way when the vertical displacement response of the disc to the external pressure forcing becomes non uniform. Let us suppose that the external pressure $ {\\cal P}$ is applied at the surface layer where the density $\\rho = \\rho_*.$ The induced vertical displacement $h,$  now should be considered to  be also  a function of  the vertical coordinate $\\zeta.$   Assuming a linear response, which should be appropriate for sufficiently small elevations and, for simplicity an isothermal equation of state with sound speed $c_s,$ it can be easily shown that in the upper layers of the disc where vertical motions dominate the Lagrangian pressure perturbation has the form determined by presence of $h$, $\\Delta P = -\\rho_{*}c_{s}^{2} {\\partial h\\over \\partial \\zeta}$, see equation (\\ref{e24}) below. Equating this to the external pressure we have \\be {\\partial h\\over \\partial \\zeta}= -{{\\cal P}\\over \\rho_* c_s^2}.\\ee If the vertical extent of the disc is $\\zeta_0,$ and $h$ varies on a radial scale comparable to $r,$  the characteristic change in $h$ induced over the vertical thickness is easily estimated to be  \\be \\Delta h \\sim  \\zeta_0{{\\cal P}\\over \\rho_* c_s^2} \\sim {2F_0r \\over 3\\rho_* \\Omega^2\\zeta_0} \\left({\\partial (h/r)\\over \\partial r}\\right),\\label{i1}\\ee where we use the approximate relation $c_s=\\Omega\\zeta_0.$  From the condition $\\Delta h \\sim h $, using (\\ref{i1}) and estimating ${\\partial h\\over \\partial r}\\sim h/r$, it is clear that whether $h$ is uniform or not is governed by the  the parameter $ F_0 r /(\\rho_*  \\Omega^2\\zeta_0^3)(\\zeta_0/r)^2 = \\lambda^{-1} \\epsilon \\delta^2,$ with $\\lambda =\\rho_*/\\rho_c$ and $\\delta = \\zeta_0/r$  so defining $\\epsilon= F_0 r /(\\rho_c  \\Omega^2\\zeta_0^3).$  For a uniform response, one requires that $\\epsilon \\delta ^2 \\ll \\lambda.$ This is equivalent to the requirement that the product of the ratio of the  external radiation momentum flux to the local disc pressure with the disc aspect ratio should be significantly less than unity. We recall here that because of the geometrical configuration,  the  momentum flux locally reradiated by  the disc is in general much less than the external momentum flux at that location.  In  this paper we extend the treatment of disc warping induced by the action of an external pressure to include the effects of the response of the vertical structure particularly when this is non uniform as is the case when the above condition is not satisfied. In that case we find that the disc dynamics enters a different regime. This is such that the exposed surface tends to align so as to reduce the momentum absorption and hence the applied pressure. In the extreme limit of this regime the upper surface acts as if it is in contact with a rigid wall with the tendency to radiation warping instability tending to vanish. In this limit quantities characterising the disc twist and warp tend to oscillate at typical frequency $\\sim \\lambda \\Omega $.  To illustrate these effects we derive a description of the one dimensional evolution of the disc inclination in radius and time that incorporates the effects of the vertical structure response, radiation torques and  which applies both to the high viscosity regime, when the evolution is diffusive, and to the low viscosity regime when the evolution is wavelike. In all cases, the efficacy of surface  radiation pressure driven instabilities is found to be reduced once the model parameters are such that the vertical response is significantly non uniform.  We go on to estimate conditions for the response to be nonlinear and investigate the development of shock waves in the response using a one dimensional slab analogue model which has the same characteristic behaviour of  linear perturbations as the full disc model. We find that such shocks potentially  provide an important dissipation mechanism.   The plan of this paper is as follows. In section \\ref{sec1} we describe some aspects of the thin disc model used. In section \\ref{Coords} we go on to introduce the twisted coordinate system used  together with the notation convention, giving the basic equations in section \\ref {Beq}. By integrating over the vertical direction we use these to  derive  a single equation governing   the dynamics of a twisted disc  in section \\ref{Derveq}. When the disc behaves like a set of rigid rings for which the vertical response is uniform, this equation can be used to give a complete specification of the warp evolution. We note that this provides an  extension of  previous formulations to be able to consider the case when  warps propagate as waves rather than diffuse radially (e.g. Nelson \\& Papaloizou 1999).  In this Paper we use the  twisted coordinate system  formalism first introduced by Petterson 1977a, 1978 where the dynamical equations take the most simple form. When this formalism is adopted and an accurate description of all components of the equations of motion is needed as in the problem on hand we show that, in general, there is an ambiguity in choice of twisted coordinates with a set of these describing the same physical situation. As discussed in section \\ref{Derveq} a choice of a most appropriate twisted coordinate system can be motivated by the condition that perturbations of all dynamical quantities determined by the disc twist and warp are small. The transformation law between different twisted coordinates corresponding to the same physical situation  is  derived in appendix C for an inviscid disc.   In order to obtain a complete description of the warp evolution when the vertical response is non uniform, we begin by  obtaining a complete solution of the vertical problem for a polytropic model in section \\ref{poly}. This is used to obtain a pair of equations governing the radial warp evolution. We also indicate how the results can be extended to apply to  more general models. We confirm the condition for the disc response to be like that of a series of rigid rings as $\\lambda \\gg \\epsilon (\\zeta_0/r)^2.$    We go on to perform a linear stability analysis of the radial evolution equations adopting a WKB approach in section \\ref {qualan} obtaining a maximum potential growth/decay rate in section \\ref{7.3}. In section \\ref{8} we discuss and confirm the analogy between the response of the disc and the  linear and non linear dynamics of a vertically stratified one dimensional slab. We consider the development of shock waves and the formation of a rarefied hot atmosphere in section \\ref{8.3} giving a crude estimate of the warp dissipation rate in section \\ref{8.4}. Finally in section \\ref {Conc} we discuss and summarize our results. ",
        "conclusions": "\\label{Conc}  In this paper we have presented a self-consistent calculation of the response of a twisted disc to the action of a radiation pressure force acting in upper layers of the disc. This is assumed to be due to the interception and subsequent reradiation  of radiation from  a central source. The  radiation pressure force is assumed to be  applied at a single density level $\\rho_{*}$ corresponding to optical thickness unity. The degree of twist and warping is assumed to be small enough that linear theory can be used to calculate the response.  The  analysis of  Pringle 1996  modelled the disc as a set of concentric rigid rings in a state of Keplerian rotation. These were assumed to communicate with each other through the exchange angular momentum because of the action of viscous forces. Up to now there has been no consideration of effects arising from the response of the disc vertical structure to the externally applied pressure.  In this paper we have considered the warping and twisting of an accretion disc taking account the response of the disc vertical structure to an external radiation field assuming that the effect of self-shielding of radiation by the global twisted disc is not significant. We developed a description of the evolution of the disc in terms of a pair of equations governing   the one dimensional evolution of the inclination as a function of  radius and time (see section \\ref{sec6} and equations (\\ref{e26}), (\\ref{e48}) and (\\ref{e49})). This description extends earlier ones, so that  the influence of radiation pressure on  discs for which warps  occur in the low  viscosity bending wave regime, as well as the higher viscosity diffusive regime,  may be considered.  We found several qualitatively different dynamical regimes that may be realised in astrophysical discs. These are related to whether the character of the response of the disc vertical structure causes significant departures from what would be obtained for a rigid body.  \\subsection{ Conditions for  a quasi-rigid response} We found that to avoid such departures, the momentum flux  due to radiation from the central source $F_0,$  should be  smaller than a critical value, $F_{*},$  given by $F_{*}=(\\lambda /\\delta )P_{c},$ where $\\delta = \\zeta_0/r $ is the disc aspect ratio, $P_{c}$ is the disc mid plane pressure and it has been assumed that the the ratio $(\\lambda /\\delta )$ is  small (see sections \\ref{Dyneq} and \\ref{7.1.2}). Then, when $F_0 > F_{crit}\\approx 0.1(\\delta /\\alpha )P_{c},$ warping instability of the disc becomes a possibility ( eg. Pringle, 1996).  Thus for radiation driven warping of a disc that behaves like an ensemble of rigid rings, we find two requirements, namely  that  $F_{*} > F_{crit},$  together with $F_0 < F _{*}.$ These conditions together imply that the ration of the surface to mid plane density should be sufficiently large and such that $\\lambda > \\lambda_{crit}\\approx 0.1\\delta^{2}/\\alpha.$   \\subsection{ Large external radiation momentum flux} In the opposite limit of large  $F_0 > max(F_{*}, F_{crit}),$ consideration of the disc structure in the vertical direction is essential as the vertical displacement ceases to be uniform. The perturbed gas motion in the disc is mainly determined by the vertical component of velocity and the perturbed quantities tend to oscillate at a characteristic frequency $\\omega = (\\rho_c \\zeta_0 / 2\\Sigma)\\lambda \\Omega $, with $\\Omega$ being the local Keplerian angular velocity. In this situation, vertical motion tends to be suppressed by the external pressure field on the irradiated side of the disc while  warping motions persist in the bulk of the disc and on the shielded side of the disc (see sections \\ref{Dyneq} and \\ref{smalllam}).  The disc surface intercepting the radiation tends to align parallel to the radiation rays, thus decreasing the amount of intercepted radiation, while the inclination angles associated with the disc mid plane and the opposite free boundary of the disc oscillate, being approximately equal to each other. In this limit the upper surface of optical thickness one plays the role of an impenetrable wall. Thus the development of instability of the inclination angle due to radiation pressure effects (e.g. Pringle 1996) would be inhibited in this limit.  In principle, the presence of warping motions in this regime would be associated with variability on  a large characteristic time scale $T_{ch}\\sim \\lambda^{-1}T_{K}$, where $T_{K}$ and $\\lambda $ are some 'typical' values of the orbital period and the density ratio.  \\subsection{Limits on the WKB growth rate} Following Pringle 1996 we have considered possible instabilities using a WKB approach. As this neglects potentially important global effects and boundary conditions results are not definitive, nonetheless they should give a good indication of when instability could be possible.  We find that when $F_{0} < F_{*}$ the growth rate increases with $F_{0}$ while in the formal limit $F_{0} \\rightarrow \\infty$ it approaches zero. Thus  in the linear theory there is an upper limit for the growth rate of the instability of the inclination angle which is always $\\le (\\rho_c \\zeta_0/ \\Sigma)\\lambda\\Omega $ in the absence of viscosity (see section \\ref{7.3}). When present, viscosity acts to damp any instability at a rate $0.1\\delta^2\\Omega /\\alpha,$ leading again to the condition $\\lambda > \\lambda_{crit}\\approx 0.1\\delta^{2}/\\alpha$ for radiation warping instability to be feasible.  \\subsection{Conditions for non-linearity} Our results described above rely upon the applicability of linear perturbation theory. As discussed in section \\ref{7.5}, the breakdown of linear theory is expected when the  Lagrangian change of pressure induced in the surface layers becomes  of the  order of the unperturbed pressure even though the change of inclination angle  could be very small. The corresponding constraints on the inclination angle are especially strong for disc models with a significant drop of temperature toward the boundaries of the disc.  \\subsection{Non linear calculations for the one dimensional slab analogue}  A direct numerical approach to the issues discussed above is not yet feasible due to three-dimensional nature of the problem and  the many physical processes involved. However, when vertical motions dominate, the problem becomes analogous to the one dimensional problem of vertical motions induced  in a stratified gas column orbiting around a central source with Keplerian angular velocity.  As discussed in Section \\ref{8} the one dimensional slab gives the same fundamental period of oscillation as obtained from  the full disc theory when an appropriate boundary condition  on the upper surface of the column is specified. The dependence of the eigenfrequency on theoretical parameters as well as the effective presence, in the limit of small $\\lambda,$ of an impenetrable wall at the upper surface of the column have been checked numerically.  Where they can be compared, agreement between our analytical and numerical results is very good.   We also used the one dimensional model in order to investigate possible consequences of non-linear behaviour in the system. To do this we studied the motion of the slab ensuing from the imposition of a vertical velocity profile with varying amplitude. We found that when the ratio of the Lagrangian pressure perturbation to the local value of the pressure at the upper boundary of the disc, $\\Delta P/P_{*}$,  becomes larger than $1-10,$ strong shocks propagating downward into the column are observed (see section \\ref{8.3}). In principle, these shocks may lead to non-linear dissipation of energy stored in the vertical motion and in section \\ref{8.4} we made a very crude estimate of a possible time scale of $200$ orbits  for $\\Delta P/P_{*}\\sim 10$. However, this result must be checked in framework of a more sophisticated numerical approach which takes into account physical processes neglected in this study, most notably the effects of radiation transfer in the upper layers of the column.  \\subsection{Issues for future consideration}  In a fully  self-consistent study the dynamical effects of radiation pressure must be studied together with effects determined by the  radiation heating of the disc atmosphere. This can be done in  the most convenient way within the  framework of the one dimensional model discussed above.  In a realistic disc model, where effects due to the  flaring of the disc lead to the interception  and scattering of radiation in the disc photosphere,  when axisymmetric and unperturbed, radiation heating can significantly influence on or even determine the value of the density ratio $\\lambda.$  This parameter, being the ratio of the density at the absorbing surface to the mid plane density defines the boundary between different dynamical regimes for a twisted disc. In this connection it is interesting to note that in certain vertical models of X-ray heated accretion discs, the ratio $\\lambda $ can be quite small, being of order of $10^{-4}-10^{-5}$, e.g. Jimenez-Garate et al 2002. In such models the new dynamical effects discussed in this Paper would certainly play an important role."
    },
    "0710/0710.0686_arXiv.txt": {
        "abstract": "A survey of currently known extrasolar planets indicates that close to 20\\% of their hosting stars are members of binary systems. While the majority of these binaries are wide (i.e., with separations between 250 and 6500 AU), the detection of Jovian-type planets in the three binaries of $\\gamma$ Cephei (separation of 18.5 AU), GL 86 (separation of 21 AU), and HD 41004 (separation of 23 AU) have brought to the forefront questions on the formation of giant planets and the possibility of the existence of smaller bodies in moderately close binary star systems. This paper discusses the late stage of the formation of habitable planets in binary systems that host Jovian-type bodies, and reviews the effects of the binary companion on the formation of Earth-like planets in the system's habitable zone. The results of a large survey of the parameter-space of binary-planetary systems in search of regions where habitable planets can form and have long-term stable orbits are also presented. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.2226_arXiv.txt": {
        "abstract": "Gamma Ray Bursts (GRBs) show evidence of different spectral shapes, light curves, duration, host galaxies and they explode within a wide redshift range. However, the most of them seems to follow very tight correlations among some observed quantities relating to their energetic. If true, these correlations have significant implications on burst physics, giving constraints on theoretical models. Moreover, several suggestions have been made to use these correlations in order to calibrate GRBs as standard candles and to constrain the cosmological parameters. \\\\ We investigate the cosmological relation between low energy $\\alpha$ index in GRBs prompt spectra and the redshift $z$. We present a statistical analysis of the relation between the total isotropic energy $E_{iso}$ and the peak energy $E_p$ (also known as Amati relation) in GRBs spectra searching for possible functional biases. \\\\ Possible implications on the $E_{iso}$ vs $E_p$ relation of the $\\alpha$ vs $(1+z)$ correlation are evaluated. We used MonteCarlo simulations and the boostrap method to evaluate how large are the effects of functional biases on the $E_{iso}$ vs $E_p$. We show that high values of the linear correlation coefficent, up to about 0.8, in the $E_{iso}$ vs $E_p$ relation are  obtained for random generated samples of GRBs, confirming the relevance of functional biases. \\\\ Astrophysical consequences from $E_{iso}$ vs $E_p$ relation are then to be revised after a more accurate and possibly bias free analysis. ",
        "introduction": "Gamma-ray Bursts (GRBs) are brief and intense flashes of high energy radiation emitted mostly in the $\\gamma$-ray band. They are detected from wholly random directions in the sky at the rate of about once a day and typically last from a few milliseconds to several minutes. Within a few years, the BATSE experiment on board the NASA's Compton Gamma Ray Observatory satellite (Fishman et al. 1989) has recorded over 2700 GRB events with an isotropic distribution in the sky (Meegan et al. 1996). However, although BATSE was very sensitive to high-energy photons, it could not discern the direction of a burst to better than a few degrees uncertainty, too large to pinpoint the location of individual explosions.  The real revolutionary step forward occurred in the 1997, thanks to the Italian-Dutch BeppoSAX satellite (Boella et al. 1997). This satellite was not as sensitive as BATSE to $\\gamma$ rays, but its relative quick response of pointing system, coupled with good accuracy position information, permitted the first detection of an X-ray {\\it afterglow}, the radiation emitted after the initial burst of $\\gamma$-ray (Costa et al. 1997).  This discovery of afterglows made redshift measurements possible and confirmed that GRBs lie at cosmological distance ($0.0085$ (Galama et al. 1999) $<z<$ $6.29$ (Kawai et al. 2006)).  It is well known that in the past years several correlations have been discovered linking various energies characterizing GRBs. All of them involve $E_p$, the energy peak of time integrated spectral energy distribution (SED). The first relation found links the rest frame isotropic energy $E_{iso}$ with $E_p$ (Amati et al. 2002, AM02 hereinafter, also known as 'Amati relation', see also Amati, 2006, AM06 hereinafter, for an updated version). This correlation is seen by several authors as an useful method to standardizing GRB energetics.  Subsequently, $E_p$ was found to be tightly correlated also with the collimation corrected energy $E_{\\gamma}$ (Ghirlanda et al. 2004, GH04 hereinafter, also known as 'Ghirlanda relation'), and this relation is used to constrain cosmological parameters using GRBs as 'known' candles. The same relations can be expressed in terms of luminosity using spectra not time integrated: $L_{iso}\\propto E_{p}^k$ (Yonetoku et al. 2004), $L_{\\gamma}\\propto E_{p}^k$ (Ghirlanda, Ghisellini \\& Firmani 2006).  However, this kind of relations seems to contradict some observational evidences, as the large variety of light curves, spectra, redshifts, durations and host galaxies, leading us to suppose that the nature of GRB explosions is not unique.  Band et al. (2005) and  Nakar et al. (2005) tested the consistency of a large sample of BATSE GRBs (with unknown redshift) with the Amati and Ghirlanda relations. Their results suggest that they may be artifacts of selection effects, inferring that about half (Nakar et al. 2005) or even $\\sim90\\%$ (Band et al. 2005) of the whole GRB population cannot satisfy the correlation for any value of the redshift. However, these conclusions have been questioned by several other authors (Ghirlanda et al. 2005, Bosnjak et al. 2005, Pizzichini et al. 2005), that found instead that the peak energy and the fluences of BATSE GRBs with unknown redshift are fully consistent with the $E_{iso}~vs~E_{p}$ and the $E_{\\gamma}~vs~E_{p}$ relations.  In AM02, and in AM06, $E_{iso}$ and $E_{p}$ are not evaluated directly by a fitting procedure but they are calculated using  analytic relations that include the same parameters. These procedures could introduce functional biases in the spectral correlation quoted above. The influence of these functional biases has not been never considered in literature.  In this work we present a statistical analysis of the $E_{iso}$ vs $E_p$ relation, based on Monte Carlo and bootstrap simulations, searching for possible intrinsic correlation terms. We previously study the cosmological relation between the redshift $z$ and the low energy spectral index $\\alpha$ (see Sec.3), because a correlation between these spectral parameters can make more tight the $E_{iso}$ vs $E_p$ relation. In Sect. 4 we evaluate functional and intrinsic correlation terms in the $E_{iso}$ vs $E_p$ relation, in order to estimate if the method used to derive it is statistically biased. ",
        "conclusions": "In this paper a statistical analysis of the GRBs parameters correlations $\\alpha$ vs $(1+z)$ and $E_{iso}$ vs $E_p$ is presented. We studied the $\\alpha$ vs $(1+z)$ correlation observing that it is present not only in the AM02 GRB sample but also in those analysed by GH04. As stated by different authors (Lloyd et al. 2000, Band et al. 1993), there is no clear explanation for it, but we suggested that it could be a direct observational consequence of a selection effects: faint high redshift GRBs are under the detectability threshold.  As a consequence of this, the decrease of high redshift GRBs introduce a correlation between $\\alpha$ and $z$. This correlation can have important consequences when we attempt to use GRBs as standard candles and, moreover, it introduces a bias in searching relations between quantities builded using $\\alpha$ and $z$.  We investigated how the presence of functional correlation terms affects the $E_{iso}$ vs $E_p$ relation. As builded in AM02, $E_{iso}$ and $E_p$ depend on a common set of spectral parameters. We performed several Monte Carlo simulations , initially considering only the intrinsic correlation term at a fixed redshift. We found that the $E_{iso}$ vs $E_p$ relation is biased, with a not negligible probability to reproduce correlation coefficents higher than 0.5. When we introduced in the generation of simulated sample also the $\\alpha$ vs $(1+z)$ correlation, the mean value of the correlation coefficent increase up to $0.64$. Higher values were obtained using Eq.8 and Eq.9 instead of Eq.4 and Eq.5. A boostrap method was also applied to the AM02 sample. The good agreement with the previous analysis proves that our results are independent on the simulation code.  An already discussed (e.g. Band et al. 2005) relevant subject is the contribution from selection effects in making more tight the $E_{iso}$ vs $E_p$ relation with respect to our simulations.  It is possible that instrumental sensivity limits would produce forbidden regions in the plane $E_{iso}, E_p$. As an example, a combination of the spectral parameters in the GRB rest frame could make a GRB with a very low luminosity and therefore below the detectability threshold. We also stress that the SED energy peak is an averaged quantity, because during a single GRB that shows more bumps in its lightcurve we expect that $E_p$ changes, making difficult the link of this variable to physical quantities.\\\\  In the literature appear several other tight correlations concerning the energetics and the spectral distributions of GRBs (e.g $E_{\\gamma}~vs~E_p$ (GH04), $L~vs~E_p$ (Yonetoku et al. 2004)), some of them are used to constrain cosmological parameters, using GRBs as 'known' candles. All of them are based on the $E_{iso}$ vs $E_p$ relation and this implies that they are affected by functional biases. In the GH04 relation, for example, the collimation corrected energy is computed considering the jet opening angle, that is in turn derived from $E_{iso}$ and $z$. \\\\  In any case, we do not exclude that a physical relation between the $E_{iso}$ and $E_p$ can really exist, but it can be safely established only after removing all the biases.  A good solution could be to use spectral laws explicitly written in terms of the quantities for which a correlation is searched (as made, for example, in Tramecere, Massaro \\& Cavaliere (2007) for the spectrum of the near HBL object Mkn 421).  A further step could also be to work with homogeneous samples of GRBs in terms of some relevant characteristics (e.g. time evolution, spectral shape, ...) to verify that the correlation is actually followed by them."
    },
    "0710/0710.0015_arXiv.txt": {
        "abstract": "We present observations of $^{13}$CO, C$^{18}$O, HCO$^{+}$, H$^{13}$CO$^{+}$, DCO$^{+}$ and N$_{2}$H$^{+}$ line emission towards the Barnard 68 starless core. The line profiles are interpreted using a chemical network coupled with a radiative transfer code in order to reconstruct the radial velocity profile of the core. Our observations and modeling indicate the presence of complex radial motions, with the inward motions in the outer layers of the core but outward motions in the inner part, suggesting radial oscillations. The presence of such oscillation would imply that B68 is relatively old, typically one order of magnitude older than the age inferred from its chemical evolution and statistical core lifetimes. Our study demonstrates that chemistry can be used as a tool to constrain the radial velocity profiles of starless cores. ",
        "introduction": "\\label{sec:introduction}  Starless cores likely represent the earliest stage of the formation of a star. In the standard view of star formation, cores form out an initially magnetically sub-critical cloud, and evolve in a quasi-static fashion through ambipolar diffusion \\citep{Shu87,Mouschiovas91}. In the opposite picture, cores are dynamic objects that form by shocks in a turbulent flow, fragment and collapse (or disappear) in a few crossing times \\citep{Padoan01,Ballesteros-Paredes2003}. Velocity measurements can provide important constraints on these two scenarios. For example, \\cite{Tafalla98} examined the kinematics of L1544 and inferred infall velocities up to 0.1 km s$^{-1}$ \\emph{a priori} incompatible with sub-critical ambipolar diffusion models. However, \\cite{Ciolek00} argued that these velocities can be understood if the core is super-critical.  Line of sight velocities in cores can be established from observations of self-absorbed line profiles. In an infalling core, if the excitation temperature increases towards the center (as one can expect in a centrally condensed core with a roughly constant temperature), self-absorbed lines are expected to be asymmetric, with a blue peak brighter than the red peak \\cite[see][]{Evans99}. This technique has been used in the past to detect collapse motions in starless cores and protostars \\cite[see][for a review]{Myers00}. One of the major difficulties of this technique is the large degree of chemical complexity within these objects. For example, numerous molecules are observed to deplete in the densest regions of starless cores \\cite[e.g.][]{Bergin02a,Tafalla02}, thus hampering our ability to measure the velocity in the dense central regions of these objects.  Nevertheless, with a detailed knowledge of the chemistry of these objects, one can choose appropriate molecular transitions to trace different part of the cores, and unveil their radial kinematic structure \\citep{Bergin03b,vanderTak05}. In this paper, we use line observations of $^{13}$CO, C$^{18}$O, HCO$^{+}$, H$^{13}$CO$^{+}$, DCO$^{+}$ and N$_{2}$H$^{+}$ to reconstruct the velocity profile of the Barnard 68 core (hereinafter B68). With a well defined physical \\citep{Alves01,Bergin06b} and chemical \\citep{Bergin02a,Maret06,Maret07a} structure, this core is an ideal target for such a study. For this we compare the observed line profiles with the predictions of a chemical network coupled with a radiative transfer code. The paper is organized as follow. Observations are presented in \\S \\ref{sec:observations}. The analysis is detailed in \\S \\ref{sec:analysis-results}, and the results are discussed in \\S \\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  Our observations and modelling indicate that the outer part of the core is collapsing while the inner part is expanding. The velocities in both regions are relatively small (a few tens of meters per second) and are largely sub-sonic. The transition between the collapsing and expanding region occurs at a radius of about 8,000 AU. It is important to note that the fit to the observations is probably not unique. However, large changes in the transition and velocities between the regions are likely ruled out. As already mentioned, our model does not reproduce the asymmetries of the N$_{2}$H$^{+}$ (1-0) and DCO$^{+}$ line profiles, due to an underestimating of the opacity. There are several possible explanations for this discrepancy. First, our chemical model might not predict the correct abundance profile for these two species. For example, increasing the DCO$^{+}$ abundance at intermediate radii, with a subsequent reduction at small radii to keep the column density constant would increase the opacity of line without changing the line flux significantly (which is correctly predicted). Indeed, our model predicts a DCO$^{+}$ (2-1) line intensity at $10 < A_{v} < 20$ that is lower than observed \\citep[see Fig. 3 in][]{Maret07a}. Second, the physical structure of the core is also somewhat uncertain. For example, the temperature profile may be slightly different than the one predicted by \\cite{Bergin06b} \\citep[see][]{Crapsi07}, although small changes ($\\pm 3$ K) in the temperature profile were found to have little effect on the line intensities.  Turbulence at the center of the core might also be different that what assumed here.  The turbulent profile was determined by \\cite{Bergin06b} using C$^{18}$O observations, which are not a good probe of the innermost regions of the core because of heavy depletion.  Decreasing the turbulent velocity profile in the innermost region of the core would also increase the opacity of these two lines. Regardless of the discrepancy with our model predictions, we emphasize these two line profiles, together with the HCO$^{+}$ (4-3) line profile, unambiguously indicate the presence of outwards motions at the center of the core. We found that no constant infalling velocity profile could reproduce the observations.  The step-like velocity profile we obtain is suggestive of the presence of radial oscillations in the core. Small non-radial oscillations of the outer layers of the core around a stable equilibrium have been proposed by \\citet{Lada03} and \\cite{Redman06} to interpret the asymmetry of the CS (2-1) and HCO$^{+}$ (3-2) lines observed in Barnard 68. \\cite{Lada03} also suggested that, in addition to the surface oscillations seen in the CS (2-1) line, some radial oscillation might also be present. Our observations and modeling are in agreement with this hypothesis.  \\cite{Keto05} proposed a hydrodynamic model for the evolution of cores in quasi-equilibrium. In their modeling, a small perturbation of a stable core (according to the Bonnor-Ebert criterion) is found to result in damped oscillations around the equilibrium position. \\cite{Keto06} showed that the surface pattern observed in the CS (2-1) line towards B68 is consistent with the (n = 1, l = m = 2) quadripolar oscillation mode for a specific orientation ($\\lambda = \\theta = 30^\\circ$). For the same oscillation mode, the radial velocity is qualitatively similar to the one obtained in this study, but is out of phase: the inner part of the core has a negative radial velocity, while the outer part of the core has a positive velocity. At the other half of the oscillation cycle, the velocity would be reversed and in qualitative agreement with the profile we obtain (E. Keto, \\emph{priv. comm.})  As pointed out by \\cite{Keto06} and \\cite{Redman06}, the presence of oscillations in B68 suggests that the core is relatively old, typically older than a few sound speed crossing times ($\\sim 10^{6}$ yr). This is comparable to the ambipolar diffusion timescale at the center of the core and more than one order of magnitude higher than the free-fall time scale \\citep[$10^{6}$ and $7 \\times 10^{4}$ yr respectively;][]{Maret07a}. This is also about an order of magnitude higher than the age of the core as determined from the degree of CO depletion \\citep[$10^{5}$yr;][]{Bergin02a,Bergin06b}\\footnote{Strictly speaking, the age determined by this method is a lower limit of the real age of the cloud, because CO is assumed to be pre-existing, and the density is supposed to be constant \\citep{Bergin06b}.}. For $t = 3 \\times 10^{5}$ yr, \\cite{Bergin06b} model predicts a C$^{18}$O (1-0) line intensity that is three time smaller than the observations, while at $t = 10^{6}$ yr, the greater CO depletion leads to an even greater mismatch between observations and model predictions. The N$_{2}$H$^{+}$, HCO$^{+}$ H$^{13}$CO$^{+}$, and DCO$^{+}$ would also be underestimated by the model because of the freeze-out of their parent of their precursors, N$_{2}$, CO and $^{13}$CO. Thus the ``chemical age'' of B68 inferred from the CO depletion appears to be hard to reconcile with the ``dynamical age'' implied by the presence of surface oscillations. It is also interesting to compare the sound crossing time in B68 to typical statistical starless core lifetimes \\cite[see][for a review]{WardThompson07}. Using sub-millimeter continuum emission maps, \\cite{Kirk05} obtained lifetimes of $1-3 \\times 10^{5}$ years which are in good agreement with the age of the B68 determined from CO depletion, but again about an order of magnitude higher than the sound crossing time.  The present study emphasizes the importance of understanding the chemistry to constrain the kinematics of starless using spectral line profiles. Because of the strong chemical gradients that exist in these cores -- as well as excitation effects -- different lines can be used to selectively trace different parts of the core. We have demonstrated that, using lines that are appropriately chosen, it is possible to reconstruct the core radial velocity profile. In B68, our observations and modelling suggest the presence of complex line-of-sight motions that are consistent radial oscillations. The presence of such oscillation indicates that some core are long-lived, with lifetimes about an order of magnitude higher than those derived from their chemical evolution or sub-millimeter surveys. So far, oscillations have been inferred in only one other core, FeSt 1-457 \\citep{Aguti07}. Clearly, similar studies on a larger number of cores are needed to establish if such cores are common."
    },
    "0710/0710.3400_arXiv.txt": {
        "abstract": "{} { We investigate whether the globular clusters (GCs) in the recently published sample of GCs in the Fornax cluster by Bergond and coworkers are indeed intra--cluster objects.} {We combine the catalogue of radial velocity measurements by Bergond et al.~with our CTIO MOSAIC photometry in the  Washington system  and analyse the relation of metal--poor and metal--rich GCs with their  host galaxies.  } {The metal--rich GCs appear to be kinematically  associated with their respective host galaxies. The vast majority of the metal--poor GCs found in between the galaxies of the Fornax cluster have velocities which are  consistent with them being members of the very extended NGC\\,1399 GC system. { We find  that when the sample is restricted to the most accurate velocity measurements, the GC velocity dispersion profile can be described with a mass model derived for the NGC\\,1399 GC system within 80\\,kpc. We identify one ``vagrant'' GC whose radial velocity suggests that it is not bound to any galaxy unless its orbit has a very large apogalactic distance.}} {} {} ",
        "introduction": "The existence of intra--cluster globular clusters (ICGCs), i.e. globular clusters (GCs) which are not bound to individual galaxies but, rather, move freely through the potential well of a galaxy cluster as a whole, was proposed by  White~(\\cite{white87}) and West et al.~(\\cite{west95}). In the nearby ($D=19\\,\\rm{Mpc}$) Fornax cluster of galaxies, ICGC candidates were identified as an excess population of GC candidates in the vicinity of dwarf galaxies (Bassino et al.~\\cite{bassino03}), and as a region of enhanced number density of GC candidates in between the central dominant galaxy NGC\\,1399 and its second--nearest (giant) neighbour, NGC\\,1387 (Bassino et al.~\\cite{b06a}, cf.~their Fig.~9).  Tamura et al.~(\\cite{tamura06}) performed a wide--field survey of the central Virgo cluster and detected an excess population of GC candidates far away from any major galaxy. The spectroscopic confirmation of these candidate ICGCs, however, is still pending. \\par Recently, Bergond et al.~(\\cite{bergond07}) presented the velocities of a sample of GCs in the Fornax cluster, with some objects having projected distances of more than 230\\,kpc from NGC\\,1399. These authors  labelled a significant fraction of their objects as ICGCs. Yet, one has to bear in mind that, with the photometric study by Bassino et al.~(\\cite{b06a}), it became clear that the very populous globular cluster system (GCS) of the central galaxy, NGC\\,1399, has an extent of at least 250\\,kpc, which is comparable to the core radius of the cluster ($R_{King}=0.7^{\\circ}\\simeq 230\\,\\rm{kpc}$, Ferguson \\cite{ferguson89}). \\par In this \\emph{Letter}, we present Washington photometry for 116 of the 149 GCs of the Bergond et al.~kinematic sample and demonstrate that the photometric and dynamical properties of the ``ICGCs'' are rather consistent with those of GCs belonging to NGC\\,1399. ",
        "conclusions": "The red (metal--rich) GCs seem to be kinematically associated with their respective host galaxy or with NGC\\,1399.  The population of ICGCs is predominantly blue, as expected from the shallower number density profile of the metal--poor NGC\\,1399 GCs.  For the blue GCs, it is harder to determine whether a given GC ``belongs'' to a minor Fornax member or if it is part of the ICGC/NGC\\,1399 population.  We conclude that the vast majority of the blue ICGCs in fact belongs to the very extended NGC\\,1399 GCS.  Out to at least 165\\,kpc, the velocity dispersion profile of the GCs is consistent with the mass profile derived from the dynamics of the NGC\\,1399 GCS within 80\\,kpc.\\par We propose to distinguish between intra--cluster GCs i.e.~GCs which are found at large distances from the galaxies in a cluster (i.e.~a geometrical classification) and \\emph{vagrant} GCs, which are characterised by radial velocities suggesting that they are not bound to any galaxy in particular, but rather belong to the galaxy cluster as a whole.  The example of Fornax illustrates the difficulties one faces when trying to identify vagrant clusters. The detection of stripped GCs is made difficult by the richness and extent of the NGC\\,1399 GCS and its high velocity dispersion.  For NGC\\,1404, which has a velocity of about $500\\,\\rm{km\\,s}^{-1}$ w.r.t.~NGC\\,1399, but a (projected) distance of only $\\sim\\!50\\,\\rm{kpc}$, it will be extremely hard to distinguish between a superposition along the line--of--sight and stripping. NGC\\,1387, at about 100\\,kpc has a systemic velocity which is just $\\sim150\\,\\rm{km\\,s}^{-1}$ lower than that of NGC\\,1399, making it hard to single out stripped GCs, although the photometry of Bassino et al.~(\\cite{b06a}) suggests their presence.  We suggest that a dynamical study of the NGC\\,1399/NGC\\,1404 and NGC\\,1387 region might yield evidence for tidal structures."
    },
    "0710/0710.1405_arXiv.txt": {
        "abstract": "We analyze \\ion{Fe}{1} 630~nm observations of the quiet Sun at disk center taken with the spectropolarimeter of the Solar Optical Telescope aboard the {\\em Hinode} satellite. A significant fraction of the scanned area, including granules, turns out to be covered by magnetic fields. We derive field strength and inclination probability density functions from a Milne-Eddington inversion of the observed Stokes profiles. They show that the internetwork consists of very inclined, hG fields.  As expected, network areas exhibit a predominance of kG field concentrations. The high spatial resolution of {\\em Hinode}'s spectropolarimetric measurements brings to an agreement the results obtained from the analysis of visible and near-infrared lines. ",
        "introduction": "\\label{sec:intro}  Most of the studied aimed at determining the distribution of field strengths in the internetwork (IN) quiet Sun have used polarimetric measurements in the spectral regions around 630~nm and 1565~nm, but their results do not agree. The visible \\ion{Fe}{1} lines at 630.2~nm indicate a predominance of kG fields (S\\'anchez Almeida \\& Lites 2000; Dom\\'{\\i}nguez Cerde\\~na et al.\\ 2003; Socas-Navarro \\& Lites 2004), whereas the infrared lines at 1565~nm suggest hG fields (Lin 1995; Lin \\& Rimele 1999; Khomenko et al.\\ 2003; Mart\\'{\\i}nez Gonz\\'alez et al.\\ 2006a; Dom\\'{\\i}nguez Cerde\\~na et al.\\ 2006). The distribution of IN field inclinations has only been studied by Lites et al.\\ (1996) and Khomenko et al.\\ (2003).  Here we analyze \\ion{Fe}{1} 630~nm measurements of the quiet Sun taken by the spectropolarimeter aboard {\\em Hinode} at the unprecedented spatial resolution of 0\\farcs32. The observed Stokes spectra are inverted to determine the distribution of field strengths and inclinations in the observed region.  Our results show that most of the IN fields are weak, opposite to what has been found from ground-based measurements of the same lines at 1\\arcsec. ",
        "conclusions": "The high spatial resolution spectropolarimetric measurements of {\\em Hinode} indicate that most IN fields are weak. This is in agreement with the picture derived from the more magnetically sensitive \\ion{Fe}{1} lines at 1565~nm (Lin 1995; Collados 2001; Khomenko et al.\\ 2003; Mart\\'{\\i}nez Gonz\\'alez et al.\\ 2006b) and from lines showing hyperfine structure such as \\ion{Mn}{1} 553~nm (L\\'opez Ariste et al.\\ 2006) and \\ion{Mn}{1} 1526.2~nm (Asensio Ramos et al.\\ 2007). Keller et al.\\ (1994) also found weak fields in the internetwork using the \\ion{Fe}{1} 525.0~nm lines, although at a lower spatial resolution and without inclination information.  Our results seem to confirm the mean IN field strength of $\\sim$100~G derived by Trujillo Bueno et al.\\ (2004) from a Hanle-effect interpretation of \\ion{Sr}{1} 460.7~nm measurements.  Interestingly, the slope of the field strength distribution in the IN is similar to that obtained from magneto-convection simulations of comparable mean flux density. The observed field inclinations, however, turn out to be significantly larger than those predicted by the simulations. The scenario of an IN filled by nearly horizontal hG fields is compatible with the large trasverse magnetic fluxes found in the IN by Lites et al.\\ (2007a,b). We still do not know the origin of such ubiquitous horizontal IN fields, but Lites et al.\\ (2007b) have suggested a number of plausible mechanisms.    In summary, Milne-Eddington inversions of the \\ion{Fe}{1} 630~nm lines observed by {\\em Hinode} at 0\\farcs32 reveal a predominance of hG fields in quiet Sun internetwork regions, contrary to what is obtained from the same lines at 1\\arcsec\\/. This is the first time that \\ion{Fe}{1} 630~nm observations confirm the weak IN fields indicated by near-infrared measurements, which may definitely close the discrepancy between the results derived from both spectral regions."
    },
    "0710/0710.3636_arXiv.txt": {
        "abstract": "$UBVRI$ photometry and medium resolution optical spectroscopy of peculiar Type Ia supernova SN 2005hk are presented and analysed, covering the pre-maximum phase to around 400 days after explosion. The supernova is found to be underluminous compared to \"normal\" Type Ia supernovae. The photometric and spectroscopic evolution of SN 2005hk is remarkably similar to the peculiar Type Ia event SN 2002cx. The expansion velocity of the supernova ejecta is found to be lower than  normal Type Ia events. The  spectra obtained $\\gsim  200$ days since explosion do not show the presence of forbidden [\\ion{Fe}{ii}], [\\ion{Fe}{iii}] and [\\ion{Co}{iii}] lines, but are dominated by narrow, permitted \\ion{Fe}{ii}, NIR \\ion{Ca}{ii} and \\ion{Na}{i} lines with P-Cygni profiles. Thermonuclear explosion model with Chandrasekhar mass ejecta and a kinetic energy smaller ($\\KE = 0.3 \\times 10^{51} {\\rm ergs}$) than that of canonical Type Ia supernovae is found to well explain the observed bolometric light curve. The mass of \\Nifs\\ synthesized in this explosion is $0.18 \\Msun$. The early spectra are successfully modeled with this less energetic model with some modifications of the abundance distribution. The late spectrum is explained as a combination of a photospheric component and a nebular component. ",
        "introduction": "\\label{sec:introduction}  An impressive homogeneity in the light curves and peak luminosities make Type Ia supernovae (SNe Ia) good candidates in the determination of the extragalactic distance scale. Though a majority of the observed SNe Ia belong to the \"normal\"   type (Branch et al.\\ 1993), a number of studies indicate significant photometric as well as spectroscopic differences. For example, studies of nearby supernovae by Li et al.\\ (2001) indicate that 64\\% SNe Ia are \"normal\", while 20\\% are of the overluminous SN 1991T type and 16\\% of the underluminous SN 1991bg type. The peak absolute luminosities of SNe Ia are well correlated with their immediate post-maximum decline rate, forming a photometric sequence from the luminous blue events with a relatively slow decline rate to the faster, red, subluminous events (Hamuy et al.\\ 1996a,b; Phillips et al.\\ 1999).  However, there are a few SNe Ia that are known to deviate from this relation, the most notable amongst them being SN 2002cx. This supernova was found to be underluminous, but had a light curve decline rate $\\Delta$m$_{15}(B)=1.29$, comparable to normal SNe (Li et al.\\ 2003). The early phase ($\\lsim 100$ days after explosion) spectra indicate line velocities lower by a factor of 2 compared to those of normal SNe Ia (Li et al.\\ 2003; Branch et al.\\ 2004). Furthermore, the late phase ($\\sim250$ days after explosion) spectra are also quite dissimilar compared to normal SNe Ia and possibly consist of P-Cygni profiles (Jha et al.\\ 2006).  Interestingly, SN 2002cx is not a unique event. Jha et al.\\ (2006) and Phillips et al.\\ (2007) have shown that SN 2005hk shows photometric and spectroscopic behaviour almost identical to that of SN 2002cx. Further, Jha et al.\\ (2006)  have shown that SNe 2003gq and 2005P also show spectra similar to SN 2002cx. Spectropolarimetric observations of SN 2005hk at the early phases by Chornock et al.\\ (2006) indicate low polarization levels, indicating that the peculiarities of SN 2002cx-like SNe do not result from an extreme asphericity. Based on the analyses of the early phase spectra of SN 2002cx, Branch et al.\\ (2004) suggest that the observed lower line velocities are consistent with the deflagration models of explosion. Jha et al.\\ (2006) report a possible detection of \\ion{O}{i} lines in the late phase spectra and suggest large scale mixing in the central region, which is also consistent with three-dimensional (3D) deflagration models.  Phillips et al.\\ (2007) find that, qualitatively the observed light curves of SN 2005hk are in reasonable agreement with  model calculations of a 3D deflagaration model that produces $\\sim 0.2~\\Msun$ of $^{56}$Ni.  Understanding the nature of this class of SNe Ia is thus quite important for the study of  homogeneity and heterogeneity of SNe Ia. It may provide a caution for the cosmological use of SNe Ia and a strong constraint to the explosion models.  We present in this paper the photometric and spectroscopic development of SN 2005hk over $\\sim400$ days since explosion. ",
        "conclusions": "Photometric and spectrosopic data on the peculiar SN 2005hk are presented. The $B$ band light curve of SN 2005hk shows comparable pre-maximum brightening, faster decline in the initial $\\sim$20 days past maximum and slower decline beyond $\\sim$50 days past maximum, as compared to normal Type Ia supernovae. The fainter peak bolometric luminosity indicates synthesis of small amount of \\Nifs \\ in the explosion. Low expansion velocity of the ejecta together with fainter peak luminosity is explained by an explosion with lower kinetic energy. A reasonable fit to the bolometric light curve of SN 2005hk is achieved with a less energetic ($0.30 \\times 10^{51}$erg) model which synthesized  $0.18 \\Msun$ of \\Nifs. The light curve evolution is similar to SN 2002cx.  The pre-maximum spectrum of SN 2005hk is similar to that of SN 1991T, with much lower expansion velocity. The spectral evolution of SN 2005hk is very similar to SN 2002cx. The spectrum of SN 2005hk at late phases ($> 200$ days) is similar to SN 2002cx, except for higher expansion velocities and higher velocity dispersions. The presence of P-Cygni profiles in the late phase spectrum of SN 2005hk indicates that the ejecta have  not become optically thin till our last observation. Modeling of the pre-maximum spectra of SN 2005hk indicates a relatively higher temperature which makes \\ion{Fe}{iii} lines strong. The presence of weak \\ion{O}{i} line at $\\lambda$7774 at almost all epochs is modeled as a consequence of high abundance of completely mixed unburned oxygen in the ejecta. The late phase spectra of SN 2005hk are modeled as a combination of the photospheric and nebular components, with the nebular line emitting region being more centrally concentrated than is expected in the case of a lower energetic model."
    },
    "0710/0710.1069_arXiv.txt": {
        "abstract": "We utilize existing imaging and spectroscopic data for the galaxy clusters MS2137-23 and Abell 383 to present improved measures of the distribution of dark and baryonic material in the clusters' central regions.  Our method, based on the combination of gravitational lensing and dynamical data, is uniquely capable of separating the distribution of dark and baryonic components at scales below 100 kpc. Our mass models include pseudo-elliptical generalized NFW profiles for constraining the inner dark matter slope, and our lens modeling takes into account both the ellipticity and substructure associated with cluster galaxies as necessary in order to account for the detailed properties of multiply-imaged sources revealed in Hubble Space Telescope images.  We find a variety of strong lensing models fit the available data, including some with dark matter profiles as steep as expected from recent simulations. However, when combined with stellar velocity dispersion data for the brightest member, shallower inner slopes than predicted by numerical simulations are preferred, in general agreement with our earlier work in these clusters.  For Abell 383, the preferred shallow inner slopes are statistically a good fit only when the multiple image position uncertainties associated with our lens model are assumed to be 0\\farcs5, to account for unknown substructure.  No statistically satisfactory fit was obtained matching both the multiple image lensing data and the velocity dispersion profile of the brightest cluster galaxy in MS2137-23.  This suggests that the mass model we are using, which comprises a pseudo-elliptical generalized NFW profile and a brightest cluster galaxy component may inadequately represent the inner cluster regions.  This may plausibly arise due to halo triaxiality or by the gravitational interaction of baryons and dark matter in cluster cores.  The intriguing results for Abell 383 and MS2137-23 emphasize the need for a larger sample of clusters with radial arcs. However, the progress made via this detailed study highlights the key role that complementary observations of lensed features and stellar dynamics offer in understanding the interaction between dark and baryonic matter on non-linear scales in the central regions of clusters. ",
        "introduction": "Cold dark matter (CDM) simulations (both with and without the inclusion of baryonic physics) are a crucial tool and proving ground for understanding the physics of the universe on nonlinear scales. One of the most active aspects of research in this area concerns the form of the dark matter density profile. Key questions raised in recent years include: Is there a universal dark matter density profile that spans a wide range of halo masses? What is the form of this profile and how uniform is it from one halo to another? To what extent do baryons modify the dark matter distribution?  Drawing on a suite of N-body simulations, \\citet{NFW97} originally proposed that the dark matter density profiles in halos ranging in size from those hosting dwarf galaxies to those with galaxy clusters have a universal form. This 3-D density distribution, termed the NFW profile, follows $\\rho_{DM}\\propto r^{-1}$ within some scale radius, $r_{sc}$, and falls off as $\\rho_{DM}\\propto r^{-3}$ beyond.  Subsequent simulations indicated that the inner density profile could be yet steeper - $\\rho_{DM}\\propto r^{-1.5}$ \\citep{M98,Ghigna00}.  As computing power increases and numerical techniques improve, it is now unclear whether the inner dark matter distribution converges to a power law form rather than becoming progressively shallower in slope at smaller radii \\citep{P03,Navarro04,Diemand04,Diemand05}.  For comparisons with data, such simulations need to account for the presence of baryons. This is particularly the case in the cores of rich clusters. Although baryons represent only a small fraction of the overall cluster mass, they may be crucially important on scales comparable to the extent of typical brightest cluster galaxies. Much work is being done to understand the likely interactions between baryons and DM \\citep{Gnedin04,Nagai05,Faltenbacher05}.  These simulations will provide refined predictions of the relative distributions of baryons and DM.  This paper is a further step in a series which aims to present an observational analog to progress described above in the numerical simulations. At each stage it is desirable to confront numerical predictions with observations. Whereas some workers have made good progress in constraining the {\\em total} density profile (e.g.~\\citet{Broadhurst05b}), in order to address the relevance of the numerical simulations we consider it important to develop and test techniques capable of  separating the distributions of dark and baryonic components (e.g. ~\\citet{Sand02,Zappacosta06,Biviano06,Mahdavi07}).  This paper presents a refined version of the method first proposed by \\citet{Sand02}, exploited more fully in \\citet{Sand04} (hereafter S04). S04 sought to combine constraints from the velocity dispersion profile of a central brightest cluster galaxy (BCG) with a strong gravitational lensing analysis in six carefully selected galaxy clusters in order to separate the baryonic and dark matter distributions.  S04 carefully selected clusters to have simple, apparently 'relaxed' gravitational potentials in order to match broadly the 'equilibrium' status of the simulated dark matter halos originally analyzed by \\citet{NFW97} and subsequent simulators.  For example, Abell 383 and MS2137-23 have almost circular BCGs ($b/a$=0.90 and 0.83 respectively), require a single cluster dark matter halo to fit the strong lensing constraints (in contrast to the more typical clusters that require a multi-modal dark matter morphology -- Smith et al. 2005), have previously published lens models with a relatively round dark matter halo ($b/a$=0.88 and 0.78 respectively - Smith et al. 2001; Gavazzi 2005), and display no evidence for dynamical disturbance in the X-ray morphology of the clusters (Smith et al. 2005; Schmidt \\& Allen 2006).  The merit of the approach resides in combining two techniques that collectively probe scales from the inner $\\sim$10 kpc (using the BCG kinematics) to the $\\sim$100 kpc scales typical of strong lensing. Whereas three of the clusters contained tangential arcs, constraining the total enclosed mass within the Einstein radius, three contained both radial and tangential gravitational arcs, the former providing additional constraints on the derivative of the total mass profile. In their analysis, S04 found the gradient of the inner dark matter density distribution varied considerably from cluster to cluster, with a mean value substantially flatter than that predicted in the numerical simulations.  S04 adopted a number of assumptions in their analysis whose effect on the derived mass profiles were discussed at the time. The most important of these included ignoring cluster substructure and adopting spherically-symmetric mass distributions centered on the BCG. The simplifying assumptions were considered sources of systematic uncertainties, of order 0.2 on the inner slope. Although the six clusters studied by S04 were carefully chosen to be smooth and round, several workers attributed the discrepancy between the final results and those of the simulations as likely to arise from these simplifying assumptions \\citep{Bartelmann04,Dalal04b,Meneghetti05}.  The goal of this paper is to refine the data analysis for two of the clusters (MS2137-23 and Abell 383) originally introduced by S04 using fully 2-D strong gravitational lensing models, thus avoiding any assumptions about substructure or spherical symmetry. The lensing models are based on an improved version of the LENSTOOL program (\\citealt{Kneibphd,Kneib96}; see Appendix; http://www.oamp.fr/cosmology/lenstool/). A major development is the implementation, in the code, of a pseudo-elliptical parameterization for the NFW mass profile, i.e. a generalization of those seen in CDM simulations, viz: \\begin{equation}\\label{eqn:gnfw} \\label{eq:gnfw} \\rho_d(r)=\\frac{\\rho_{c} \\delta_{c}}{(r/r_{sc})^{\\beta}(1+(r/r_{sc}))^{3-\\beta}} \\end{equation} where the asymptotic DM inner slope is $\\beta$. This formalism allows us to overcome an important limitation of previous work and takes into account the ellipticity of the DM halo and the presence of galaxy-scale subhalos. Furthermore the 2-D lensing model fully exploits the numerous multiply-imaged lensing constraints available for MS2137-23 and Abell 383.  The combination of gravitational lensing and stellar dynamics is the most powerful way to separate baryons and dark matter in the inner regions of clusters.  However, it is important to keep a few caveats in mind.  Galaxy clusters are structurally heterogeneous objects that are possibly not well-represented by simple parameterized mass models. To gain a full picture of their mass distribution and the relative contribution of their major mass components will ultimately require a variety of measurements applied simultaneously across a range of radii.  Steps in this direction are already being made with the combined use of strong and weak gravitational lensing (e.g.~\\citet{Limousin06,Bradac06}), which may be able to benefit further from information provided from X-ray analyses (e.g.~\\citet{Schmidt06}) and kinematic studies (e.g.~\\citet{Lokas03}). A recent analysis has synthesized weak-lensing, X-ray and Sunyaev-Zeldovich observations in the cluster Abell 478 -- similar cross-disciplinary work will lend further insights into the mass distribution of clusters \\citep{Mahdavi07}.  Of equal importance are mass models with an appropriate amount of flexibility and sophistication.  For instance, incorporating models that take into account the interaction of baryons and dark matter may shed light into the halo formation process and provide more accurate representations of dark matter structure.  Halo triaxiality, multiple structures along the line of sight and other geometric effects will also be important to characterize. At the moment, incorporating these complexities and securing good parameter estimates is computationally expensive even with sophisticated techniques such as the Markov-Chain Monte Carlo method.  Numerical simulation results are often presented as the average profile found in the suite of calculations performed. Instead, the distribution of inner slopes would be a more useful quantity for comparison with individual cluster observations.  Also, comparisons between simulations and observations would be simplified if {\\it projected} density profiles of simulated halos along multiple lines of sight were to be made available.  These issues should be resolvable as large samples of observed mass profiles are obtained.  For the reasons above, comparing observational results with numerical simulations is nontrivial.  The observational task should be regarded as one of developing mass modeling techniques of increasing sophistication that separate dark and baryonic matter, so as to provide the most stringent constraints to high resolution simulations which include baryons as they also increase in sophistication.  The combination of stellar dynamics and strong lensing is the first crucial step in this process.  Its diagnostic power will be further enhanced by including other major mass components (i.e.~the hot gas of the intracluster medium or the stellar contribution from galaxies) out to larger radii.  A plan of the paper follows.  In \\S~\\ref{sec:methods} we explain the methodology used to model the cluster density profile by combining strong lensing with the BCG velocity dispersion profile.  In \\S~\\ref{sec:obsresults} we focus on translating observational measurements into strong lensing input parameters.  This section includes the final strong lensing interpretation of MS2137-23 and Abell 383.  In \\S~\\ref{sec:stronglens} we present the results of our combined lensing and dynamical analysis.  In \\S~\\ref{sec:systematics} we discuss further systematic effects, limitations and degeneracies that our technique is susceptible to -- with an eye towards future refinements.  Finally, in \\S~\\ref{sec:finale} we summarize and discuss our conclusions. Throughout this paper, we adopt $r$ as the radial coordinate in 3-D space and $R$ as the radial coordinate in 2-D projected space.  When necessary, we assume $H_{0}$=65 \\kms Mpc$^{-1}$, $\\Omega_{m}$=0.3, and $\\Omega_{\\Lambda}$=0.7. ",
        "conclusions": "\\label{sec:systematics}  In the previous section we have presented the results of our analysis, which showed that a mass model comprising a stellar component for the BCG following a Jaffe profile together with a generalized NFW DM cluster halo is able to adequately reproduce the observations for Abell 383 (albeit {\\it only} for the coarse lensing positional accuracy scenario) but is unable to simultaneously reproduce the observed multiple image configuration and BCG velocity dispersion profile for MS2137-23.  In the case of Abell 383, the inner DM profile is flatter than $\\beta$=1, supporting the earlier work of S04.  This indicates that at least some galaxy clusters have inner DM slopes which are shallower than those seen in numerical simulations -- but only if the mass parameterization used in the current work is reflective of reality.  Further work in this interesting cluster using other observational probes will further refine the mass model, and determine if the generalized NFW DM form is a good fit to the cluster profile.  In this section we discuss systematic uncertainties in our method and possible refinements that could be made to reconcile the mass model with the observations for MS2137-23. We hope that many of these suggestions will become important as cluster mass models improve and thus will present fruitful avenues of research.  \\subsection{Systematic Errors}  We focus first on systematic errors associated particularly with the troublesome stellar velocity dispersion profile for MS2137-23.  Errors could conceivably arise from (i) significant non-Gaussianity in the absorption lines (which are fit by Gaussians), (ii) uncertain measurement of the instrumental resolution used to calibrate the velocity dispersion scale, and (iii) template mismatch.  Non-gaussianity introduces an error that we consider too small to significantly alter the goodness of fit \\citep{Gavazzi05}. The instrumental resolution of ESI (the Keck II instrument used to measure the velocity dispersion profile; \\citet{Sheinis02}) is $\\sim$30 \\kms; this is much smaller than the measured dispersion. Even if the instrumental resolution was in error by a factor of two, the systematic shift in $\\sigma$ would only be 3 \\kms (using Eq~3 in Treu et al.\\ 1999). This would affect all measurements and not reverse the trend with radius.  Concerning template mismatch, S04 estimated a possible systematic shift of up to 15-20 \\kms . This could play a role especially as the signal to noise diminishes at large radii, where the discrepancy with the model profile is greatest. To test this hypothesis, we added 20 \\kms in quadrature to only those velocity dispersion data points in MS2137-23 at $R > 4$ kpc and recalculated the best-fitting $\\chi^{2}$ values. $\\chi^{2}$ is reduced from 31 to 28.8, a modest reduction which fails to explain the poor fit.  Although selectively increasing the error bars on those data points most discrepant with the model is somewhat contrived, our result does highlight the need for high S/N velocity dispersion measures out to large radii.  A high quality velocity dispersion profile has been measured locally for Abell 2199 to $\\sim$20 kpc \\citep{Kelsonetal02}. Interestingly, these high S/N measures display similar trends to those for MS2137-23 in the overlap regime, i.e. a slightly decreasing profile at $R\\lesssim 10$kpc.  The dip witnessed in MS2137-23 is thus not a unique feature, although with deeper measurements we might expect to see a rise at larger radii as a result of the shallow DM profile.  A final potential limitation in the dynamical analysis is the assumption of orbital isotropy. Both S04 and \\citet{Gavazzi05} explored the consequences of mild orbital anisotropy, concluding a possible offset of $\\Delta \\beta \\sim0.15$ might result.  Even including orbital anisotropy into his analysis, \\citet{Gavazzi05} was unable to fit the observed velocity dispersion profile.  Since we determine our lensing $\\chi^{2}$ values in the source plane, we checked to make sure that no extra images were seen after remapping our best-fit lensing + velocity dispersion models back to the image plane. No unexpected images were found, although several images that were explicitly not used as constraints were found, such as the mirror image of radial arc image 2a in MS2137 and the complex of multiple images associated with 3abc, 5ab, and 6ab in Abell 383 (see Figures~\\ref{fig:mulplot} and \\ref{fig:mulplota383}).  As discussed in \\S~\\ref{sec:lensinterpms2137} and \\ref{sec:lensinterpa383}, some of these multiple images were not used as constraints because we could not confidently identify their position either due to galaxy subtraction residuals or blending with other possible multiple image systems.  \\begin{figure*} \\begin{center} \\mbox{ \\mbox{\\epsfysize=4.5cm \\epsfbox{f5a.eps}} \\mbox{\\epsfysize=4.5cm \\epsfbox{f5b.eps}} } \\mbox{ \\mbox{\\epsfysize=4.5cm \\epsfbox{f5c.eps}} \\mbox{\\epsfysize=4.5cm \\epsfbox{f5d.eps}} } \\caption{Confidence contours (68\\%,95\\%, and 99\\%) when we allow the dark matter scale radius to be fixed at values a factor of two beyond our observationally motivated prior.  Top Row -- Contours when we fix the dark matter scale radius to $r_{sc}$=50 and 400 kpc in MS2137.  Although the $r_{sc}$=400 kpc scenario provides a relatively good fit to the data ($\\chi^{2}\\sim$26), this value for the scale radius is much larger than that observed from weak lensing data.  The $r_{sc}$=50 kpc scenario is a significantly worse fit to the data, with $\\chi^{2}\\sim$39.  Note that the DM inner slope is $\\beta < 1$ in both scenarios.  Bottom Row -- Contours when we fix the dark matter scale radius to $r_{sc}$=50 and 400 kpc in A383. The large discrepancy in inner slope values obtained emphasize the need for a mass probe at larger radii.  The best-fitting model for either fixed scale radius is significantly worse than the best-fitting $r_{sc}$=100 kpc result ($\\chi^{2}\\sim$26.5 and 31.3 for $r_{sc}$=50 and 400 kpc respectively). \\label{fig:diffrsc}} \\end{center} \\end{figure*}  We finally comment on the uncertainties assigned to the multiple image systems for our lens models.  We have presented two sets of results in this work; with assigned image positional accuracies of $\\sigma_{I}$=0\\farcs2 and 0\\farcs5.  We find a variety of lens models are compatible with the $\\sigma_{I}$=0\\farcs2 case and only when the velocity dispersion data is included into the analysis does the data fail to be reproduced by the model.  Certainly if we were to further increase the positional errors, at some point a good velocity dispersion fit could conceivably be obtained, but we will refrain from doing so in the present work.  Increasing the positional uncertainties is only justified if there is evidence that there are significant missing components in the mass models.  Further observations that can probe the mass distribution on fine scales to larger radii and higher quality models which can account for phenomena such as adiabatic contraction in the inner regions of galaxy clusters and triaxiality represent the best way to obtain a more precise picture of the cluster mass distribution.   \\subsection{Improving the Mass Model}  We now turn our attention to possible inadequacies in the mass model. It is important to stress that the two diagnostics (lensing and dynamics) adopted in this study probe different scales.  The lensing data tightly constrains the mass profile at and outside the radial arc ($\\sim$20 kpc), while the velocity dispersion constrains the mass profile inside $R \\lesssim 10$ kpc.  Since multiple images are numerous and their positions can be more precisely measured than velocity dispersion \\footnote{The error on the astrometry with respect to the relevant scale, the Einstein Radius $\\theta_{\\rm E}$ is much smaller than the relative error on velocity dispersion, i.e. $\\delta \\theta / \\theta_{\\rm E} << \\delta \\sigma/\\sigma$}, they carry more weight in the $\\chi^2$ statistic than the kinematic points, producing a best overall fitting model (which is lensing dominated) that is a relatively poor fit to the kinematic data.  To improve the model, one must admit that either one of the two components of the modeling is incorrect, or that the functional form of the mass profile chosen to extrapolate the lensing information at the scales relevant for dynamics is insufficient.  In this section we discuss several areas where the mass model presented in this paper could be improved.  \\subsubsection{The Contribution of the Brightest Cluster Galaxy}  We might query the assumption of a Jaffe density profile for the BCG. This seems an unlikely avenue for improvement given the Jaffe profile fits the observed BCG surface brightness profile remarkably well (see Figure 2 of S04).  Moreover, \\citet{Gavazzi05} utilized a Hernquist mass profile in his analysis of MS2137-23, which also matches the observations, and \\citet{Gavazzi05} was likewise unable to reproduce the observed S04 velocity dispersion profile.  We have additionally checked our assumptions by altering the PIEMD fit to the BCG surface brightness data so that it is matched not to the derived Jaffe profile fit to the BCG but directly to the HST surface brightness profile.  With this setup, we found a $r_{cut}$ value of 23.70 kpc for MS2137 and 28.65 kpc for Abell 383 (compare this with the numbers in Table~\\ref{tab:lensfixed}).  Redoing our analysis for the best-fitting $r_{sc}$ scenario only, our constraints on $\\beta$ for both Abell 383 and MS2137 did not change by more than 0.05, and so it is not likely that our method for constraining the BCG mass contribution is the root cause of our inability to fit the data to a BCG + gNFW cluster DM halo mass model.  Conceivably the BCG may not be coincident with the center of the cluster DM halo, as has been assumed throughout this work.  It is often the case that small subarcsecond off sets between BCGs and cluster DM halos are necessary to fit lensing constraints (e.g.~\\citet{gps05}).  There is strong evidence that the BCG is nearly coincident with the general cluster DM halo {\\it in projection} from the strong lensing work presented here and by others \\citep{gavazzi03,Gavazzi05}.  However, an offset could be responsible for the flat to falling observed velocity dispersion profile if the BCG were actually in a less dark matter dominated portion of the cluster.  Another possibility is that there are multiple massive structures along the line of sight, which would be probed by the strong lensing analysis, but not with the velocity dispersion profile of the BCG.  A comprehensive redshift survey of MS2137-23 could provide further information on structures along the line of sight.   \\subsubsection{The Advantage of a Mass Probe at Larger Radii}\\label{sec:highrad}  With our presented data set, we have seen that it is difficult to constrain the DM scale radius, $r_{sc}$ because both of our mass probes are only effective within the central $\\sim$100 kpc of the clusters -- within the typical DM scale radius observed and seen in CDM simulations.  For this reason, the inferred DM scale radius for both Abell 383 and MS2137-23 lay at the boundary of our assumed prior range.  Future work will benefit from weak lensing data, along with galaxy kinematics and X-ray data of the hot ICM which can each probe out to large clustercentric radii.  Although not the focus of the current work, pinning down the correct DM scale radius will be crucial for constraining other DM mass parameters.  For instance, there is a well-known degeneracy between $r_{sc}$ and the inner slope $\\beta$ (e.g.~\\citet{gavazzi03,Gavazzi05}). To briefly explore this, we have reran our analysis (for the coarse positioning lensing case) for both clusters with a $r_{sc}$ of 50 and 400 kpc -- factors of two beyond our chosen $r_{sc}$ prior.  We show our confidence contours in Figure~\\ref{fig:diffrsc}, which are noteworthy.  For example in the case of MS2137-23, if we fix $r_{sc}$=50 kpc, then the best-fitting $\\beta = 0.05$.  However, if $r_{sc}$=400 kpc then $\\beta=0.7$, more in accordance with simulations.  Interestingly, the $r_{sc}$=400kpc scenario returns a better overall $\\chi^{2}\\sim26$ than any model with $r_{sc}$=100-200 kpc -- even though a $r_{sc}$ of 400 kpc is clearly ruled our by extant weak lensing observations.  None of the other $r_{sc}$=50,400 kpc scenarios produced $\\chi^{2}$ values that were comparable to those seen with $r_{sc}$=100-200 kpc.  Any further knowledge of the DM scale radius would aid greatly in constraining $\\beta$ and determining the overall goodness of fit of the generalized NFW DM profile to the cluster data.  X-ray studies assuming hydrostatic equilibrium \\citep{Allen01vir,Schmidt06} and a combined strong and weak lensing analysis \\citep{Gavazzi05} have presented data on MS2137-23 to radii much larger than that probed in this study. To check that the mass model derived from data within $\\sim$ 100 kpc do not lead to results at variance with published data at larger radii, we have taken the \\citet{Gavazzi05} results and compared their derived mass at large radii with an extrapolation of our mass models.   Examining Figure~3 from \\citet{Gavazzi05} we estimate that from his weak lensing analysis a 2D projected mass enclosed between $1.6 \\times 10^{14}$ and $1.1 \\times 10^{15} M_{\\odot}$ at $\\sim 1.08$ Mpc using the cosmology adopted in this paper.  Correspondingly, if we take all of the $\\Delta \\chi^{2}<1.0$ models using our analysis method (the coarse positional accuracy case was use) and calculate the expected 2D projected mass enclosed at 1.08 Mpc we find values between $6.9 \\times 10^{14}$ and $8.4 \\times 10^{14} M_{\\odot}$, well within the expected range.   Note that no attempt was made to extrapolate the mass {\\it profiles} derived in our analysis to larger radii than the data in this paper allow, although we are acquring weak lensing data for a large sample of galaxy clusters to perform a more extensive analysis. The purpose of this consistency check is to only ensure that the masses we derive for such large radii are not too discrepant with existing analyses.  The consistency check is satisfied and lends some credence to the models.  \\subsubsection{Dark matter baryons interactions and Triaxiality}\\label{sec:ac}  The central regions of DM halos can be strongly affected by the gravitational interaction with baryons during halo formation.  If stars form and condense much earlier than the DM, it is expected that the baryons will adiabatically compress the DM resulting in a halo that is {\\it steeper} than that of the original \\citep{Blumenthal86,Gnedin04}.  Alternatively, dark matter heating through dynamical friction with cluster galaxies can counteract adiabatic contraction, leading to a shallower DM profile \\citep{Elzant04,Nipoti03}.  The present work takes into account neither of the above scenarios, and if any baryon-DM interaction greatly changes the cluster density profile, our assumed parameterized gNFW profile may be inappropriate.  Recently, \\citet{Zappacosta06} have used X-ray mass measurements in the cluster Abell 2589 to conclude that processes in galaxy cluster formation serve to counteract adiabatic contraction in the cluster environment. Certainly, more observational work is needed to understand the interplay between baryons and DM in clusters, and extended velocity dispersion profiles of BCGs in conjunction with other mass tracers at larger radii could serve as the best testing ground for the interplay of dark and luminous matter.  Not only is there likely significant interplay between baryons and DM in the central regions of clusters, but real galaxy clusters are certainly triaxial and, if ignored, this may lead to biased parameter estimations and discrepancies when combining mass measurement techniques that are a combination of two- and three-dimensional. Several recent studies have considered the effects of halo triaxiality on observations.  Using an N-body hydrodynamical simulation of a disk galaxy and performing a 'long slit' rotation curve observation, \\citet{Hayashi04} found that orientation and triaxial effects can mistake a cuspy DM profile for one that has a constant density core. At the galaxy cluster scale, \\citet{Clowe04} performed mock weak lensing observations of simulated galaxy clusters and found that the NFW concentration parameter recovered was correlated with the 3D galaxy cluster orientation.  In order to investigate the recent rash of galaxy clusters with observed high concentration parameters in seeming contradiction to the CDM paradigm \\citep{Kneib03,Gavazzi05,Broadhurst05b}, \\citet{Oguri05} used strong and weak lensing data in Abell 1689 along with a set of models that included halo triaxiality and projection effects.  Again, it was seen that halo shape causes a bias in mass (and mass profile) determination, although it should be kept in mind that measurements of concentration are extremely difficult (e.g. Halkola et al.\\ 2006), and the recent study of \\citet{Limousin06} has seemed to clear up the concentration parameter controversy for at least Abell 1689.  In terms of the current work, \\citet{Gavazzi05} has pointed out that the inability of his lensing model to fit the MS2137-23 BCG velocity dispersion profile may be due to halo triaxiality or projected mass along the line of sight (which would increase the mass measured in the lensing analysis but would not be seen in the stellar velocity dispersion).  \\citet{Gavazzi05} showed that an idealized prolate halo with an axis ratio of $\\sim$ 0.4 could explain the velocity dispersion profile in MS2137-23.  Halo triaxiality could also explain the high concentration previously seen in this cluster.  Again, the gap between simulations and observations may be bridged with respect to triaxiality if further steps were taken to compare the two directly. One step in this direction would be the publication of detailed density profiles for the simulations (in 3-D or along numerous projected sight-lines).  The most recent DM only simulations have indicated that the standard NFW profile representation of a DM profile (and its \\citet{M99} counterpart with an inner slope $\\beta \\sim 1.5$) can be significantly improved by slightly altering the model to a profile with a slope that becomes systematically shallower at small radii (e.g.~\\citet{Navarro04}, but see \\citet{Diemand05}).  While we have adopted the traditional generalized NFW profile in this study, future work with parameterized models should move towards the latest fitting functions along with an implementation of adiabatic contraction as has already been attempted by \\citet{Zappacosta06}.  Note, however, that both \\citet{Navarro04} and \\citet{Diemand04} have stated that all fitted functions have their weaknesses when describing complicated N-body simulations and that when possible simulations and observations should be compared directly."
    },
    "0710/0710.1891_arXiv.txt": {
        "abstract": "Our high time resolution observations of individual pulses from the Crab pulsar show that the main pulse and interpulse differ in temporal behavior, spectral behavior, polarization and dispersion. The main pulse properties are consistent with one current model of pulsar radio emission, namely, soliton collapse in strong plasma turbulence.  The high-frequency interpulse is quite another story. Its dynamic spectrum cannot easily be explained by any current emission model; its excess dispersion must come from propagation through the star's magnetosphere. We suspect the high-frequency interpulse does not follow the ``standard model'', but rather comes from some unexpected region within the star's magnetosphere. Similar observations of other pulsars will reveal whether the radio emission mechanisms operating in the Crab pulsar are unique to that star, or can  be identified in the general population. ",
        "introduction": "Between 1994 and 2002 our group observed strong pulses from the Crab  pulsar between 1 and 5 GHz at the VLA and Arecibo.  Our Arecibo observations were designed with 2-ns time resolution, in order to test competing models of the radio emission mechanism.  We initially concentrated on the MP, because it is stronger than the IP below $\\sim$5 GHz in the star's mean profile, and also because giant pulses are more common at the phase of the MP at these frequencies \\cite{Cordes2004}.  Our results, and our data acquisition system, are described in \\cite{HKWE}.    \\begin{figure}[ht] \\vspace{-1.5in} \\rotatebox{-90}{ \\includegraphics[trim = 40 0 60 0, width=0.6\\textwidth,clip]{EilekFig1.ps}} \\vspace{-2.5in} \\caption{An example of a Main Pulse, observed with 2.2-GHz bandwidth at 9 GHz, and coherently dedispersed \\cite{HKWE}. The pulse seen in total intensity (upper panel, plotted with total intensity time resolution 6.4 ns) contains several short-lived microbursts. The dynamic spectrum (lower panel, plotted with resolution 102 ns and 19.5 MHz) shows that the microburst emission spans the full receiver bandwidth.  In a few MPs individual, short-lived nanoshots are sparse enough in time to be separately identified (shown in \\cite{HKWE, HE2007}); these examples reveal the dynamic spectrum of the nanoshots is relatively narrow. } \\label{MPfig} \\end{figure}  \\begin{figure}[ht] \\rotatebox{-90}{ \\includegraphics[trim = 40 0 60 0, width=0.6\\textwidth,clip]{EilekFig2.ps}} \\caption{An example of an Interpulse, observed with 2.2 GHz bandwidth at 9 GHz, and coherently dedispersed.  The IP seen in total intensity (upper panel) typically contains 1 or 2 sub-bursts; thus it has a simpler time signature than the MP (as in the example of Figure \\ref{MPfig}).  The regular {\\it emission bands} in the dynamic spectrum (lower panel) are not due to instrumental or interstellar effects, but are characteristic to the emission physics of the IP. Note that the secondary burst, seen in total intensity, coincides with the appearance of new band sets in the dynamic spectrum. Plotted with total intensity time resolution 51.2 ns, and dynamic spectrum resolution 104 ns and 19.5 MHz. } \\label{IPfig} \\end{figure}  To follow up on these results, we  extended our data acquisition system to higher time resolution.  We went to higher frequencies (5 to 10 GHz) in order to take advantage of the 2.2-GHz bandwidths (and corresponding sub-ns time resolution) available at Arecibo. We recorded individual IPs as well as MPs, because at these frequencies strong pulses are much more common at the phase of the IP \\cite{Cordes2004}.  These new observations, carried out between 2003 and 2006, are reported in \\cite{HE2007}.  We were astonished to find that IPs are very different from MPs at these frequencies. The IP differs from the MP in polarization, time signature, spectrum and dispersion, as discussed below.    All of our IP observations were taken betweeen 5 and 10 GHz. Technical limitations, as well as the scarcity of strong IPs at lower frequencies, kept us from observing the IP below 4 GHz.  We are therefore describing the ``high-frequency IP'', which occurs at a slighly earlier rotation phase from the ``regular'' IP (as seen at lower radio frequencies as well as in  high-energy bands; {\\em e.g.}, \\cite{MH1996}).  This phase offset suggests that the high-frequency IP may not be related at all to the regular IP; it may come from a very different part of the magnetosphere. ",
        "conclusions": "We have identified two different types of coherent radio emission from the Crab pulsar, one associated with the MP, the other associated with the high-frequency IP.  Two different radiation mechanisms seem to be be operating within the star's magnetosphere.  In addition, the higher dispersion of the high-frequency IP suggests the signal has passed through an unusually large plasma column before leaving the pulsar.   Conventional wisdom ascribes the MP to emission from the open field line region above one of the star's magnetic axes.  If this is the case, the high-frequency IP is probably not simply radiation from the other magnetic pole. It is more likely to come from some unexpected part of the magnetosphere; its phase offset, relative to the regular IP \\cite{MH1996}, corroborates this idea. The suggestion from \\cite{Maxim} that the high-frequency IP is emitted within the closed field line region may be on the right track."
    },
    "0710/0710.2776_arXiv.txt": {
        "abstract": "{} {The nature of \\sori~(S\\,Ori~J053810.1$-$023626), a faint mid-T type object found towards the direction of the young \\so~cluster, is still under debate. We intend to disentangle whether it is a field brown dwarf or a 3-Myr old planetary-mass member of the cluster.} {We report on near-infrared $JHK_s$ and mid-infrared [3.6] and [4.5] IRAC/Spitzer photometry recently obtained for \\sori. The new near-infrared images (taken 3.82\\,yr after the discovery data) have allowed us to derive the first proper motion measurement for this object. } {The colors $(H-K_s)$, $(J-K_s)$ and $K_s$\\,$-$\\,[3.6] appear discrepant when compared to T4--T7 dwarfs in the field. This behavior could be ascribed to a low-gravity atmosphere or alternatively to an atmosphere with a metallicity significantly different than solar. The small proper motion of \\sori~(11.0\\,$\\pm$\\,5.9~mas\\,yr$^{-1}$) indicates that this object is further away than expected if it were a single field T dwarf lying in the foreground of the \\so~cluster. Our measurement is consistent with the proper motion of the cluster within 1.5\\,$\\sigma$ the astrometric uncertainty. } {Taking into account \\sori's proper motion and the new near- and mid-infrared colors, a low-gravity atmosphere remains as the most likely explanation to account for our observations. This supports \\sori's membership in \\so, with an estimated mass in the interval 2--7~\\mj, in agreement with our previous derivation.} ",
        "introduction": "Knowledge of the initial mass function is crucial for understanding the formation processes of stars, brown dwarfs and free-floating planetary-mass objects. Whether and where there is a limit for the creation of objects by direct collapse and fragmentation of molecular clouds has become one of the major goals in the study of very young populations. Planetary-mass candidates with masses in the interval 3--13 Jovian masses (\\mj) have been found in various star-forming regions (e.g., Lucas \\& Roche \\cite{lucas00}; Zapatero Osorio et al. \\cite{osorio00}; Chauvin et al. \\cite{chauvin04}; Lucas et al. \\cite{lucas05}; Luhman et al. \\cite{luhman05}; Jayawardhana \\& Ivanov \\cite{ray06}; Allers et al. \\cite{allers06}; Gonz\\'alez-Garc\\'\\i a et al. \\cite{gonzalez06}; Caballero et al. \\cite{caballero07}). These objects are mostly free-floating but in in a few cases appear as wide companions to young brown dwarfs or low-mass stars.  \\sori~(S\\,Ori~J053810.1$-$023626) is the coolest free-floating, planetary-mass candidate so far reported in the literature. It was discovered by Zapatero Osorio et al$.$ (\\cite{osorio02a}) and lies in the direction of the \\so~cluster (352~pc and 1--8~Myr, with a best estimate at 3~Myr; Perryman et al. \\cite{perryman97}; Oliveira et al. \\cite{oliveira02}; Zapatero Osorio et al. \\cite{osorio02b}; Sherry et al. \\cite{sherry04}). The spectral type of \\sori~was determined at T5.5\\,$\\pm$\\,1.0 from molecular indices measured over near-infrared $H$- and $K$-band low-resolution spectra. Mart\\'\\i n \\& Zapatero Osorio (\\cite{martin03}) obtained an intermediate-resolution spectrum from 1.17 to 1.37~$\\mu$m ($J$-band), in which the K\\,{\\sc i} doublet at 1.25\\,$\\mu$m was detected. After comparison with theoretical spectra from Allard et al. (\\cite{allard01}), the authors inferred an effective temperature and surface gravity of $T_{\\rm eff}$\\,=\\,1100\\,$^{+200}_{-100}$~K and log\\,$g$\\,=\\,3.5\\,$\\pm$\\,0.5~cm~s$^{-2}$, in agreement with the expectations for a few megayears-old T dwarf. State-of-the-art evolutionary models (Chabrier \\& Baraffe \\cite{chabrier00}; Burrows et al. \\cite{burrows97}; Baraffe et al. \\cite{baraffe98}) yield a mass of 3\\,$^{+5}_{-1}$~\\mj~if \\sori's very young age is finally confirmed.  Burgasser et al. (\\cite{burgasser04}), in contrast, have raised doubts about the low-gravity atmosphere and true cluster membership of \\sori. Based on the supposed similarity of the observed spectra to field T6--T7 dwarfs, these authors argued that the S\\,Ori object is ``an old, massive field brown dwarf lying in the foreground of the \\so~cluster''. However, this work relied on low signal-to-noise ratio data. Better quality photometry and spectra are needed to assess the true nature of this candidate.  Here we present astrometric measurements, IRAC/Spitzer data and $JHK_s$ photometry for \\sori. We find that this object has colors unexpected for its spectral classification, which is measured in the range T4.5--T7 with a best estimate at T6. We ascribe this to a low gravity atmosphere, with a different metallicity being an alternative, but less likely, explanation.    \\begin{table*} \\caption[]{Log of near-infrared observations of \\sori.} \\label{obslog} \\centering \\begin{tabular}{llcccc} \\hline\\hline Telescope   & Instrument & Field of view & Pixel    & Observing dates & Exposure time \\\\ &            & (arcmin$^2$)  & (arcsec) &                 & (s)           \\\\ \\hline 3.5\\,m CAHA & Omega-2000 & 225           & 0.45     & 2005 Oct 22 ($JH$), 2005 Oct 25 ($K_s$) & 3600 ($J$), 6000 ($H$), 7200 ($K_s$) \\\\ 10\\,m KeckII& NIRSPEC    & 0.59          & 0.18     & 2005 Oct 26 ($JK'$)                     & 405 ($J$), 270 ($K'$)                \\\\ \\hline \\end{tabular} \\end{table*}  \\begin{table*} \\caption[]{Near-infrared (2MASS photometric system) and IRAC/Spitzer photometry of \\sori.} \\label{phot} \\centering \\begin{tabular}{lccccccc} \\hline\\hline Telescope  &         $J$        &         $H$        &        $K_s$       &         [3.6]       &        [4.5]       &         $H-K_s$      &         $J-K_s$      \\\\ &        (mag)       &        (mag)       &        (mag)       &         (mag)       &        (mag)       &          (mag)       &          (mag)       \\\\ \\hline 3.5 m CAHA & 19.98\\,$\\pm$\\,0.06 & 20.07\\,$\\pm$\\,0.07 & 19.60\\,$\\pm$\\,0.08 &      ...            &       ...          & $+$0.48\\,$\\pm$\\,0.11 & $+$0.38\\,$\\pm$\\,0.10 \\\\ Keck\\,II   & 19.96\\,$\\pm$\\,0.07 &       ...          & 19.58\\,$\\pm$\\,0.07 &      ...            &       ...          &        ...           & $+$0.38\\,$\\pm$\\,0.10 \\\\ Spitzer    &        ...         &       ...          &        ...         &  18.62\\,$\\pm$\\,0.30 & 17.19\\,$\\pm$\\,0.15 &        ...           &         ...          \\\\ \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "The low proper motion of \\sori~($\\mu$\\,=\\,11.0\\,$\\pm$\\,5.9~mas~yr$^{-1}$) makes it unlikely that it is a nearby ($\\le$30~pc) T dwarf. On the one hand, we compared our measurement with the motion of 192 Hipparcos stars (Perryman et al. \\cite{perryman97}) within a radius of $15^\\circ$ around \\so~and at a distance between 80 and 130~pc, which is the distance interval expected for \\sori~if it were a field, single T6 dwarf. About 70\\%~of the Hipparcos stars show larger motion than \\sori, suggesting that the S\\,Ori object is located farther away. On the other hand, our measurement is consistent (within 1.5~$\\sigma$) with the proper motion of the O9.5V--B0.5V-type star \\so~AB, which is the most massive member of the cluster of the same name. We note, however, that the relative motion of the Orion OB association is directed away from the Sun (de Zeeuw et al. \\cite{zeeuw99}). This makes it very hard to detect cluster members via proper motion analysis. On the contrary, radial velocity studies can be more discriminant (Jeffries et al. \\cite{jeffries06}) but the extreme faintness of \\sori~prevents any accurate radial velocity measurement with current instrumentation.  \\subsection{Color-color diagrams}  Color-color diagrams are depicted in Fig.~\\ref{colcol}. To put \\sori~into context, we included $JHK_s$ data of more than 100 field T-dwarfs and IRAC/Spitzer photometry of 36 field T-dwarfs compiled from the literature (Knapp et al. \\cite{knapp04}; Tinney et al. \\cite{tinney05}; Patten et al. \\cite{patten06}; Burgasser et al. \\cite{burgasser06a}; Artigau et al. \\cite{artigau06}; Mugrauer et al. \\cite{mugrauer06}; Leggett et al. \\cite{leggett99}, \\cite{leggett02}, \\cite{leggett07}; Liu et al. \\cite{liu07}; Luhman et al. \\cite{luhman07}; Looper et al. \\cite{looper07}). All near-infrared colors were conveniently transformed into the 2MASS photometric system using equations quoted in Stephens \\& Leggett (\\cite{stephens04}), which are appropriate for ultracool dwarfs.  The location of \\sori~in Fig.~\\ref{colcol} is challenging since this object appears as an outlier, particularly when the $K$-band magnitude is involved. Two field dwarfs lie near it in the near-infrared color-color panels: 2MASS\\,J00501994$-$3322402 (Tinney et al. \\cite{tinney05}), whose photometric errors are quite large, and 2MASS\\,J13243559$+$6358284 (Looper et al. \\cite{looper07}). The latter object is widely discussed by Looper et al. (\\cite{looper07}) in terms of binarity and low-gravity atmosphere. We applied the criterion defined by Covey et al. (\\cite{covey07}, eq.~2) to distinguish objects with photometric properties deviating from the properties typical of field dwarfs and found that \\sori~lies more than 2\\,$\\sigma$ away from the near-infrared sequence defined by the field T-type brown dwarfs. In the mid-infrared wavelengths, the photometry of \\sori~deviates from the field on a 1--2\\,$\\sigma$ level. Covey et al.'s equation takes into account the color uncertainties of \\sori~and the width of the field distribution.  As compared to T4--T7 field dwarfs, \\sori~presents redder $(H-K_s)$, $(J-K_s)$, and [3.6]\\,$-$\\,[4.5] colors and a bluer $K_s$\\,$-$\\,[3.6] index than expected for its spectral type ($\\sim$T6). However, the $(J-H)$ index is similar to that of T5--T6 field dwarfs. The $K$-band reddish nature of \\sori~is also apparent in its low-resolution $HK$ spectrum. Fig.~4 of Burgasser et al. (\\cite{burgasser04}) shows these data along with the spectrum of the field T6.5 2MASS\\,J10475385$+$2124234. Both spectra were obtained with similar instrumentation and are normalized to unity at 1.57~$\\mu$m. While the field dwarf matches reasonably well the $H$-band region of \\sori, it underestimates the flux at $K$-band, supporting the redder $(H-K_s)$ index of \\sori. Burgasser et al. (\\cite{burgasser04}) argued that this may be indicative of a ``lower surface gravity for \\sori~relative to 2MASS\\,J10475385$+$2124234''.  Multiplicity cannot explain the observed photometric properties of \\sori. We artificially produced near-infrared colors of L--T and T--T pairs using the absolute magnitudes provided by Liu et al. (\\cite{liu06}). None of the combinations were able to reproduce our observations.  The comparison of our data to theory is shown in the right panel of Fig.~\\ref{colcol}. The models depicted are those of Tsuji et al. (\\cite{tsuji04}), but in our analysis we also employed cloudless models by Marley et al. (\\cite{marley02}) and Burrows et al. (\\cite{burrows06}) obtaining similar results. The agreement between the models and the field dwarf observations is reasonably good. The great majority of the mid- and late-T dwarfs lie within the log\\,$g$\\,=\\,4.0 and 5.5~dex tracks, as expected for ``old'' dwarfs in the solar neighborhood. This is also consistent with the recent results of the spectral fitting work by Burgasser et al. (\\cite{burgasser06b}). These authors employed models by Burrows et al. (\\cite{burrows06}).  The near-infrared photometry of \\sori~and current state-of-the-art theory of ultracool dwarfs indicate that this object may have a lower-gravity atmosphere than similarly classified T dwarfs in the solar vicinity. Because of the different pressure and density conditions at which H$_2$, CH$_4$, and CO absorptions are produced, low-gravity objects tend to be brighter at $K$ and redder in all near-infrared and [3.6]\\,$-$\\,[4.5] colors than comparable high-gravity objects (see discussions and Figures in Knapp et al. \\cite{knapp04}; Patten et al. \\cite{patten06}; Leggett et al. \\cite{leggett07}; Liebert \\& Burgasser \\cite{liebert07}). This is what we qualitatively observe in \\sori.  From the right panel of Fig.~\\ref{colcol} and using the solar-metallicity models by Tsuji et al. (\\cite{tsuji04}), we derive log\\,$g$\\,$\\sim$\\,3.0 dex and $T_{\\rm eff}$\\,$\\sim$\\,1000--1100~K. This is consistent with previous results obtained from the spectral fitting analysis of low- and intermediate-resolution near-infrared spectra: $T_{\\rm eff}$\\,$\\sim$\\,800$^{+200}_{-100}$ K and log\\,$g$\\,$\\sim$\\,4.0\\,$\\pm$\\,1.0 dex (Zapatero Osorio et al. \\cite{osorio02a}), $T_{\\rm eff}$\\,$\\sim$\\,1100$^{+200}_{-100}$ K and log\\,$g$\\,$\\sim$\\,3.5\\,$\\pm$0.5 dex (Mart\\'\\i n \\& Zapatero Osorio \\cite{martin03}), respectively. These authors compared observations to theoretical data computed by Allard et al. (\\cite{allard01}). Recently, Liu et al. (\\cite{liu07}) have quantified the sensitivity of near-infrared spectra with $T_{\\rm eff}$\\,$\\sim$\\,700--900~K to changes in metallicity and surface gravity using a different grid of synthetic spectra by Burrows et al. (\\cite{burrows06}). \\sori~is brighter at $K$ relative to $J$ or $H$ by a factor of $\\sim$1.4. Table~5 and Fig.~4 of Liu et al. (\\cite{liu07}) suggest that log\\,$g$ is thus lower than the field by about 1.0\\,dex, in agreement with previous determinations. We caution that current state-of-the-art synthetic spectra do not provide detailed fits to the observed data (e.g., Burrows et al. \\cite{burrows06}). Therefore, any quantitative result derived from the direct comparison of observations to models awaits further confirmation. On the contrary, for a given temperature qualitative predictions on the atmospheric relative behavior as a function of gravity and metal content can be more reliable.  Metallicity might also be an issue. The \\so~cluster has solar abundance ([Fe/H]\\,=\\,0.0\\,$\\pm$\\,0.1~dex; Caballero \\cite{caballero06}); we do not expect any metallicity effect when comparing cluster members to the field. From theoretical considerations, the effects of increasing abundance and decreasing gravity on the near-infrared spectra of cool T dwarfs are similar. From Fig.~20 of Burrows et al. (\\cite{burrows06}), which shows $(J-K)$ against $T_{\\rm eff}$ for various gravities and metal abundances, we infer a metallicity of [Fe/H]\\,$\\sim$\\,$+$0.5 dex for \\sori~if its reddish effect was all due to metallicity. A similar rich metal content is obtained from Liu et al. (\\cite{liu07}). Nevertheless, the super-solar metallicity explanation, although possible, seems unlikely. On the one hand, the metallicity distribution of F, G, and K dwarf stars in the solar neighborhood peaks at around [Fe/H]\\,=\\,0.0 dex and roughly extends up to $+$0.5~dex; less than $\\sim$10\\%~of the stars are more metal-rich than $+$0.3 dex (Valenti et al. \\cite{valenti05}; Santos et al. \\cite{santos05}; Boone et al. \\cite{boone06}). On the other hand, the IRAC/Spitzer data of the S\\,Ori object do not support the high metallicity case. Burgasser et al. (\\cite{burgasser06b}) and Liebert \\& Burgasser (\\cite{liebert07}) have demonstrated 2MASS\\,J12373919$+$6526148 (T6.5) and 2MASS J09373487$+$2931409 (T6p) to be old, high surface gravity brown dwarfs with sub-solar abundance. These field dwarfs display [3.6]\\,$-$\\,[4.5] colors slightly redder (by 0.15\\,mag) than expected for their assigned spectral types. In contrast to what could be inferred from the near-infrared colors, this would indicate that \\sori~is a low metallicity T dwarf. Thus, a low-gravity atmosphere remains as the most likely explanation to account for the observed photometry of \\sori.  \\begin{figure*} \\centering \\includegraphics[width=9cm]{4aanda.ps}~~~ \\includegraphics[width=9cm]{5aanda.ps} \\caption{Color-magnitude diagram of \\so~low mass members. \\sori~is labeled. The 3-Myr isochrone by Chabrier \\& Baraffe (\\cite{chabrier00}) is overplotted onto the data as a solid line in the left panel (see text for $T_{\\rm eff}$, log\\,$L/L_\\odot$ conversion into observables). The right panel shows the sequence of M5.5--T8 field dwarfs at the distance of the cluster (solid line, spectral types are indicated). The dotted line stands for the field sequence shifted by $-$1.8 mag in the $J$-band to match the photometric trend delineated by \\so~members. Objects with $J$\\,$-$\\,[3.6] colors significantly redder than the cluster sequence show infrared flux excesses likely due to circum(sub)stellar disks (Caballero et al. \\cite{caballero07}).} \\label{colmag} \\end{figure*}  No obvious infrared flux excesses are detectable in the IRAC/Spitzer [3.6] and [4.5] bands, suggesting that there is no envelope or disk around \\sori~emitting intensively at these wavelengths. A possitive detection would have provided strong evidence for its youth. However, we remark that disks around young, low-mass brown dwarfs (close to the deuterium burning mass limit) are seen at wavelengths longer than 5\\,$\\mu$m (Luhman et al. \\cite{luhman05}; Caballero et al. \\cite{caballero07}; Zapatero Osorio et al. \\cite{osorio07}), while the observed fluxes in the near-infrared up to 5\\,$\\mu$m are photospheric in origin. The public [5.8]- and [8.0]-band IRAC/Spitzer images are not conclusive for \\sori.  \\subsection{Color-magnitude diagram}  The photometric sequence of \\so~substellar members, including \\sori, is shown in the $J$ vs $J$\\,$-$\\,[3.6] color-magnitude diagram of Fig.~\\ref{colmag} (Caballero et al. \\cite{caballero07}; Zapatero Osorio et al. \\cite{osorio07}). This sequence follows a relatively smooth progression with increasing color down to $J$\\,$\\sim$20 and $J$\\,$-$\\,[3.6]\\,=\\,2.8~mag. The location of \\sori~suggests that the $J$\\,$-$\\,[3.6] index suddenly turns toward bluer values at a nearly unchanged $J$ magnitude. A similar turnover (ocurring at spectral types L7--L8, 1400--1300~K) is also observed in field ultracool dwarfs. The field sequence of objects with spectral types M5.5--T8 moved to the distance of the \\so~cluster is displayed in the right panel of Fig.~\\ref{colmag} (absolute magnitudes and colors are adopted from Patten et al. \\cite{patten06}, and references therein). Because of their very young age, \\so~low-mass stars and substellar objects are in the phase of gravitational contraction (e.g., Chabrier \\& Baraffe \\cite{chabrier00}). Cluster members thus show larger size and luminosity than their older counterparts of related colors in the field. As seen in Fig.~\\ref{colmag}, the average cluster photometric sequence appears brighter than the field by about 1.8~mag in the $J$-band. The dotted line in Fig.~\\ref{colmag} (right panel) represents the field sequence normalized to the \\so~locus of late-M and early-L cluster members. \\sori~nicely sits on the location expected for \\so~T-type members.  We have also compared our data to the {\\sc cond} and {\\sc dusty} solar metallicity evolutionary models by Chabrier \\& Baraffe (\\cite{chabrier00}). For the age range 1--8\\,Myr, substellar objects with $T_{\\rm eff}$ between 700 and 1300~K, corresponding to the mass interval $\\sim$1--6\\,\\mj, show log\\,$g$\\,=\\,3.0--4.0 dex, which coincides within the large error bar with the surface gravity estimation for \\sori.  The 3-Myr isochrone (Chabrier \\& Baraffe \\cite{chabrier00}) is displayed along the photometric sequence of \\so~in the left panel of Fig.~\\ref{colmag}. Theoretical surface temperatures and luminosities were converted into observed magnitudes and colors using the color-temperature-spectral type and spectral type-bolometric correction relationships given in the literature (Dahn et al. \\cite{dahn02}; Vrba et al. \\cite{vrba04}; Knapp et al. \\cite{knapp04}; Patten et al. \\cite{patten06}). The model convincingly reproduces the cluster low mass sequence except for the fact that \\sori~appears overluminous by about 1~mag, which might suggest binarity. This was also discussed in Zapatero Osorio et al. (\\cite{osorio02a}). However, there are issues that prevent us from concluding whether this object is double or whether models make wrong predictions for the smallest masses and young ages: {\\sl (i)} as seen from the field sequence (Fig.~\\ref{colmag}, right panel), there is a $J$-band brightening across the color turnover that theory fails to reproduce (Vrba et al. \\cite{vrba04}; Knapp et al. \\cite{knapp04}). {\\sl (ii)} The blue color turnover takes place at a roughly constant temperature in the field ($\\sim$1300--1400~K), and all color-temperature-bolometric correction transformations show a sharp change at this point; on the contrary, the evolutionary models available to us do not have a complete temperature sampling (this may explain the abrupt color reversal at the bottom of the isochrone in Fig.~\\ref{colmag}). {\\sl (iii)} The relations used to transform theoretical predictions into observables are obtained for high-gravity field objects. It is now known that gravity impacts significantly the near- and mid-infrared colors of T dwarfs (Leggett et al. \\cite{leggett07}; Burrows et al. \\cite{burrows06}), whereas the colors of the warmer M and L types are not so sensitive to the gravity parameter. Transformations are thus expected to be gravity-dependent for the coolest temperatures. It becomes necessary to find a physical explanation for the $J$ brightening feature and to discover more \\so~T-type, planetary-mass members for a proper comparison with evolutionary tracks."
    },
    "0710/0710.2295_arXiv.txt": {
        "abstract": "I discuss the development and resolution of the solar neutrino problem, as well as opportunities now open to us to extend our knowledge of main-sequence stellar evolution and neutrino astrophysics. ",
        "introduction": "} This paper is based on a talk given at the Caltech conference \\cite{caltech}  ``Nuclear Astrophysics 1957-2007: Beyond the First 50 Years,\" July 23-27, 2007, which focused on the state of nuclear astrophysics fifty years after the seminal paper of Burbidge, Burbidge, Fowler, and Hoyle \\cite{bbfh}. The quest to measure solar neutrinos, and later to resolve the solar neutrino problem, began in the early days of nuclear astrophysics, with the first efforts to understand proton burning in main sequence stars. I would like to review that history, our current understanding of solar neutrinos, and open questions in neutrino physics, and discuss some opportunities for further solar neutrino measurements.  Solar neutrino physics brings together stellar modeling, nuclear reactions, and observation.  A key early development was the 1959 Holmgren and Johnston \\cite{holmgren} measurement of the S-factor for the pp-chain reaction $^3$He($\\alpha,\\gamma)^7$Be, which proved to be surprisingly large.   This implied that the sun could produce some of its energy through the more temperature dependent ppII and ppIII cycles of the pp chain, elevating the neutrino fluxes expected from  $^7$Be electron capture and $^8$B $\\beta$ decay.  (See Fig. 1.)   These neutrinos contribute to ground- and excited-state transitions in $^{37}$Cl($\\nu$,e)$^{37}$Ar, a reaction for detecting neutrinos that had been proposed by Pontecorvo \\cite{pontecorvo} in 1946 and considered by Alvarez \\cite{alvarez} in 1949 (who studied backgrounds that might inhibit, for example, a reactor neutrino experiment).  Alvarez's envisioned experiment -- proposed as a test of the distinguishability of the neutrino and antineutrino, prior to the discovery of parity violation  -- was later conducted by Davis at Savannah River \\cite{jandr,davis55}.  \\begin{figure} \\begin{center} \\includegraphics[width=12cm]{fig1.pdf} \\end{center} \\caption{The three cycles of the pp chain and associated neutrinos.} \\label{fig1} \\end{figure}  The Caltech effort in nuclear astrophysics brought three young researchers, Bahcall, Iben, and Sears, together in the summer of 1962.   Stimulated in part by the Holmgren and Johnston measurement, they began construction of a solar model to predict the solar core temperature, the most important parameter governing the competition between the ppI, ppII, and ppIII cycles, and to provide the first quantitative estimate of the resulting neutrino fluxes.  The Bahcall, Fowler, Iben, and Sears model, published in 1963, predicted a counting rate for Davis's proposed 100,000-gallon chlorine solar neutrino detector of about one event per day.  A key development occurred in 1963 when, during a seminar by Bahcall at Copenhagen, Mottelson inquired about the importance of neutrino excitation of excited states in $^{37}$Ar \\cite{jandr}.   In 1964 Bahcall and Barnes \\cite{barnes} pointed out that a calibration of the excited state contribution to the $^{37}$Cl cross section could be made by measuring the delayed protons from the analog $\\beta$ decay of $^{37}$Ca, the isospin mirror of $^{37}$Cl.  Effectively the lifetime of $^{37}$Ca would be a test of the elevated cross section predicted on the basis of the excited-state contribution to $^8$B neutrino absorption.  Subsequent measurements by Hardy and Verrall \\cite{hardy} and Reeder, Poskanzer, and Esterlund \\cite{poskanzer} established the importance of the excited-state contribution. [These experiments were later repeated in a manner that was kinematically complete: see Adelberger et al. \\cite{adelberger} for a discussion.]  Bahcall \\cite{bahcall64} and Davis \\cite{davis64} published companion letters in March 1964 arguing the adequacy and feasibility of a 100,000-gallon Cl experiment to measure solar neutrinos.  Excavation of the detector cavity in the Homestake Mine began in summer 1965.  The tank was installed and filled by the next year.  The first results were announced by Davis, Harmer, and Hoffman \\cite{davis68} in 1968, an upper bound of 3 SNU (1 SNU = 10$^{-36}$ capures/Cl atom/sec), below the standard solar model (SSM) prediction of Bahcall, Bahcall, and Shaviv of 7.5 $\\pm$ 3 SNU \\cite{bbs}.  This result and associated theoretical work on suggested solutions led to a series of experiments -- Gallex/GNO/SAGE \\cite{gallex,gno,sage}, Kamiokande \\cite{kamiokande}, Super-Kamiokande \\cite{sk} , and the Sudbury Neutrino Observatory \\cite{sno}.  Efforts by Borexino \\cite{borexino} and KamLAND \\cite{kamland} to measure the $^7$Be neutrinos are currently underway or in preparation. These experiments are important as tests of the SSM and of the new neutrino physics that proved to be the source of the solar neutrino problem. ",
        "conclusions": "Neutrino astrophysics and the theories of the origin of the elements, the main theme of this conference, share a common history.  Laboratory astrophysics has made solar neutrino physics into a quantitative field, and allowed experimenters to anticipate the kinds of major discoveries that justified experiments like SNO and Super-Kamiokande.  The results -- discovery of neutrino mass and flavor mixing characterized by large angles -- are of great importance, providing our first constraints on physics beyond the SM of particle physics.  But as summarized here, the list of remaining laboratory neutrino physics questions is long.  The answers to the open questions will be important in helping us characterize extreme astrophysical and cosmological neutrino environments. The needed 20-year program of laboratory and astrophysical neutrino studies is not unlike the laboratory/astrophysics interface that Willie Fowler cultivated to help us understand the origin of the elements.  Despite the current focus on particle physics properties of neutrinos, solar neutrino spectroscopy remains an important probe of the SSM and stellar evolution.  The arguments for measuring the CNO flux, using our sun as a calibrated laboratory, seem particularly strong.  Such a program would effectively test our understanding of the hydrogen burning mechanism for massive main-sequence stars.  It would also address the primary discrepancy in the SSM, the tension between helioseismology and neutrino flux predictions that follows from new analyses of surface metallicity."
    },
    "0710/0710.2892_arXiv.txt": {
        "abstract": "The recently recognized class of ``transitional disk\" systems consists of young stars with optically-thick outer disks but inner disks which are mostly devoid of small dust. Here we introduce a further class of ``pre-transitional disks\" with significant near-infrared excesses which indicate the presence of an optically thick inner disk separated from an optically thick outer disk; thus, the spectral energy distributions of pre-transitional disks suggest the incipient development of disk gaps rather than inner holes. In UX Tau A, our analysis of the {\\it Spitzer} IRS spectrum finds that the near-infrared excess is produced by an inner optically thick disk and a gap of $\\sim$56 AU is present. The {\\it Spitzer} IRS spectrum of LkCa 15 is suggestive of a gap of $\\sim$46 AU, confirming previous millimeter imaging. In addition, UX Tau A contains crystalline silicates in its disk at radii $\\gtrsim$ 56 AU which poses a challenge to our understanding of the production of this crystalline material. In contrast, LkCa 15's silicates are amorphous and pristine. UX Tau A and LkCa 15 increase our knowledge of the diversity of dust clearing in low-mass star formation. ",
        "introduction": "Previous studies have revealed stars with inner disks that are mostly devoid of small dust, and these ``transitional disks\" have been proposed as the bridge between Class II objects, young stars surrounded by full disks accreting material onto the central star, and Class III objects, stars where the protoplanetary disk is mostly dissipated and accretion has stopped (e.g. Strom et al. 1989, Skrutskie et al. 1990; Stassun et al. 2001).   New spectra from the {\\it Spitzer Space Telescope} which greatly improve our resolution in the infrared have been used to define the class of ``transitional disks\" as those with spectral energy distributions (SEDs) characterized by a significant deficit of flux in the near-infrared relative to optically thick full disks, and a substantial infrared excess in the mid- and far-infrared. Extensive modeling studies of several transitional disks around T Tauri stars \\citep{dalessio05, uchida04, calvet05, espaillat07} and F-G stars \\citep{brown07} have been presented. In particular, the SEDs of the transitional disks of the T Tauri stars (TTS) CoKu Tau$/$4 \\citep{dalessio05}, TW Hya \\citep{calvet02, uchida04}, GM Aur, DM Tau \\citep{calvet05}, and CS Cha \\citep{espaillat07} have been explained by modeling the transitional disks with truncated optically thick disks with most of the mid-infrared emission originating in the inner edge or ``wall\" of the truncated disk. In all these cases, except in CoKu Tau$/$4, material is accreting onto the star, so gas remains inside the truncated disk, but with a small or negligible amount of small dust, making these regions optically thin.  Here we present models of UX Tau A and LkCa 15, low-mass pre-main sequence stars in the young, $\\sim$1 Myr old Taurus star-forming region which have been previously reported as transitional disks \\citep{furlan06, bergin04}. We present evidence for gaps in optically thick disks, as opposed to ``inner holes\", that is, large reductions of small dust from the star out to an outer optically thick wall. ",
        "conclusions": "Here we introduce the ``pre-transitional disk\" class where we see the incipient development of disk gaps in optically thick protoplanetary disks as evidenced by significant near-infrared excesses when compared to the Taurus median SED and previously studied transitional disks \\citep{dalessio05, calvet02, uchida04, calvet05, espaillat07}. The pre-transitional disk of UX Tau A has a $\\sim$56 AU gap as opposed to an inner hole.  It is also possible to fit LkCa 15's SED with a $\\sim$46 AU gap that contains some optically thin dust; a model that has a hole rather than a gap also fits its SED and future near-infrared interferometry may be able to discriminate between these models.  However, the truncation of LkCa 15's outer disk at $\\sim$46 AU is consistent with resolved millimeter interferometric observations \\citep{pietu06} which makes it one of three inner disk holes imaged in the millimeter (TW Hya: Hughes et al. 2007; GM Aur: Wilner et al. in preparation).  In addition to our sample, the disks around F-G stars studied by \\citet{brown07} also belong to the pre-transitional disk category.  The large gaps that are being detected in pre-transitional disks are most likely due to observational bias since larger gaps will create larger mid-infrared deficits in the SED.  Smaller gaps will most likely have less apparent dips in their SEDS and be more difficult to identify, however, if their gaps contain some optically thin material the silicate emission in these objects should be much stronger than can be explained by a full disk model.  The existence of an inner optically thick disk may be an indicator of the first stages of disk clearing that will eventually lead to the the inner holes that have been seen in previously reported transitional disks; this has important implications on disk evolution theories since only planet-formation can account for this structure. Hydrodynamical simulations have shown that a newly formed planet could accrete and sweep out the material around it through tidal disturbances and this is sufficient in producing the hole size in CoKu Tau$/$4 \\citep{quillen04}, even maintaining substantial accretion rates \\citep{varniere06}. Moreover, \\citet{najita07} have found that the intrinsic properties of transitional disks may favor planet formation. Another proposed formation mechanism for the holes in transitional disks is photoevaporation, in which a photoevaporative wind halts mass accretion towards the inner disk and material in this inner disk is rapidly evacuated creating an inner hole \\citep{clarke01}; the hole then increases in size as the edge continues photoevaporating \\citep{alexanderarmitage}. Neither this model nor the inside-out evacuation induced by the MRI \\citep{chiang07} would explain how an optically thick disk accreting at a sizable accretion rate (see Table 1) would remain inside the hole. Rapid dust growth and settling has also been proposed to explain the holes in disks \\citep{lin04}.  Again, this does not account for the presence of optically thick inner disk material given that theory suggests grain growth should be fastest in the inner disk, not at some intermediate radius \\citep{weiden97, chiang99}.  Our sample also has interesting dust compositions (Watson et al. 2007, Sargent et al. in preparation).  LkCa 15 has an amorphous silicate feature indicating little if any processing leading to the crystallization seen in other young stars.  Amorphous silicates are also seen in CoKu Tau$/4$, DM Tau, and GM Aur. In contrast, UX Tau A is different from all the other transitional disks because it has crystalline silicate emission features in addition to amorphous silicate emission features (Fig. 1 inset). The wall at $\\sim$56 AU is the main contributor to the crystalline silicate emission since it dominates the flux in the mid- and far-infrared.  This raises the question of whether crystalline silicates are created close to the star or if they can be created in situ at $\\sim$56 AU.  If the former, it challenges current radial-mixing theories, none of which can get significant amounts of crystalline silicates out to this distance \\citep{gail01, kellergail04}.  One possibility for {\\it in situ} processing may be collisions of larger bodies, which might produce small grains heated sufficiently to create crystals (S. Kenyon, personal communication).  Pre-transitional disks offer further insight into the diversity of the ``transitional disk\" class and future studies of these disks will greatly advance our understanding of disk evolution and planet formation.     \\vskip -0.1in"
    },
    "0710/0710.0545_arXiv.txt": {
        "abstract": "We present a model to estimate the synchrotron radio emission generated in microquasar (MQ) jets due to secondary pairs created via decay of charged pions produced in proton-proton collisions between stellar wind ions and jet relativistic protons. Signatures of electrons/positrons are obtained from consistent particle energy distributions that take into account energy losses due to synchrotron and inverse Compton (IC) processes, as well as adiabatic expansion. The space parameter for the model is explored and the corresponding spectral energy distributions (SEDs) are presented. We conclude that secondary leptonic emission represents a significant though hardly dominant contribution to the total radio emission in MQs, with observational consequences that can be used to test some still unknown processes occurring in these objects as well as the nature of the matter outflowing in their jets. ",
        "introduction": "\\label{intro}  X-ray binary systems (XRBs) are composed by either a stellar mass black hole or a neutron star, and a normal (non degenerated) star which supplies matter to the compact object through the formation of an accretion disk. Some ~260 XRBs are known up to now \\cite{liu06} probably corresponding to an underlying population of some tens of thousands of compact objects in our Galaxy. A few of these sources present also non-thermal radio emission, hence evidencing the existence of mechanism(s) capable of injecting and/or accelerating large numbers of relativistic particles. Some radio emitting X-ray binary systems (REXBs) have been observed showing ejection of material at relativistic velocities and to display jets like those seen in quasars and active galactic nuclei but at $\\sim 10^{-6}$ times shorter scales. This analogy is the reason for calling them microquasars (MQs) \\cite{mirabel99} and making them some of the most interesting objects for astrophysics. Furthermore, attention on these objects has grown since the proposal of Paredes et al. (2000) \\cite{paredes00} of MQs as counterparts of some of the unidentified gamma-ray sources of the EGRET catalog \\cite{hartman99} and hence pointing them as plausible high energy emitters. A strong confirmation of this association has come from the detections of the MQs LS 5039 and LS I +61 303 at Tev energies using respectively the ground-based Cherenkov telescopes HESS \\cite{aharonian05} and MAGIC \\cite{albert06}, giving support and empowering at the same time a number of previous detailed studies centered on the mechanisms operating in these sources in order to explain the gamma ray domain (see, e.g., \\cite{bosch05} and \\cite{romero05}). Moreover, a jet origin of the emission from MQs has been suggested from the observation of syncrothron emission of relativistic electrons/positrons extending from the radio all the way into the X-ray regime. In this sense jet-like models have focused on different approaches regarding the particle origin that could generate the required emission properties in a consistent way. Some of them consider leptons directly injected at the base of the jet and, in extending outwards, Compton-interaction with external/self-created photon fields produces high energy radiation. Other models deal with an hadronic origin of the high energy emission, through proton-proton interactions and pion decay producing  gamma rays and leaving the resulting co-generated leptons as low energy emitters. The present work refers to the later procedure, focusing on the modelisation of the secondary leptonic synchrotron emission in order to constrain the characterization of MQ jets. An outline of the model is given in the next, followed by the results showing the SEDs and lightcurves under different parameter assumptions and the conclusions that can be extracted from them. ",
        "conclusions": "SEDs are obtained for different magnetic field values, electron/positron spectral indices and spatially distributed disks. We have estimated also the expected emission along the jet at 1 and 5 GHz. Leptons are injected in the context of hadronic secondaries generation within a detailed model that takes into account in a consistent way particle injection mechanisms and cooling due to radiation processes and adiabatic expansion. The luminosities obtained are slightly lower than in the models based on primary leptons injection, and must be considered complementary to them. However, we note that within our model there is no requirements of acceleration processes along the jet to obtain the final emission results. Such acceleration processes are still not well understood, although. They could come from diffusive shock acceleration along the jet when fresh ejecta interact with previous blobs of plasma already outflowing at lower velocities. Other scenarios assume a continuous energy transfer mechanism from the magnetic field to the matter content of the jet in such a way that the resulting parsec-scale radio emission can be explained. The fact of studying alternative models were particles are directly injected until a certain height along the jet can constrain the amount of acceleration required and contribute to the understanding of the physical mechanisms that can lead to such processes.   Signatures at different distances along the jet and specific spectral features detectable for reasonable parameter values treated in our numerical simulations have the potential to be an important clue for determining the matter content of jets. In particluar, highly resolved observations at 1 and 5 GHz could determine if leptons are present at heights $10^{12-13}$ cm at the edge of the binary system typical region where wind matter from the companion is still significant. If electrons/positrons still show high energies due to a recent injection from hadronic interactions at these parts of the jet, it could be a signature of secondary generation without the necessity of invoking additional acceleration processes."
    },
    "0710/0710.5010_arXiv.txt": {
        "abstract": "Recent observations of XTE J1739-285 suggest that it contains a neutron star rotating at 1122 Hz\\cite{Kaaret2007}. Such rotational frequency would be the first for which the effects of rotation are significant. We study the consequences of very fast rotating neutron stars for the potentially observable quantities as stellar mass and pulsar period. ",
        "introduction": "\\label{sect:introd}  Neutron stars with their very strong gravity can be very fast rotators. Theoretical studies show that they could rotate at sub-millisecond periods, i.e., at frequency $f=1/{\\rm period}>$1000 Hz\\cite{CST1994,Salgado1994}. The first millisecond pulsar B1937+21, rotating at $f=641$ Hz\\cite{Backer1982}, remained the most rapid one during 24 years after its detection. In January 2006, discovery of a more rapid pulsar J1748-2446ad rotating at $f=716$ Hz was announced \\citep{Hessels2006}. However, such  sub-kHz frequencies are still too low to significantly affect the  structure of neutron stars with $M>1M_\\odot$ \\cite{STW1983}. Actually, they belong to a {\\it slow rotation} regime, because their $f$ is significantly smaller than the mass shedding (Keplerian) frequency $f_{\\rm K}$. Effects of rotation on neutron star structure are then $\\propto (f/f_{\\rm K})^2\\ll 1$. Rapid rotation regime for $M>1M_\\odot$ requires submillisecond pulsars with supra-kHz frequencies $f>1000$ Hz.  Very recently Kaaret et al.\\cite{Kaaret2007} reported a discovery of oscillation frequency $f=1122$ Hz in an X-ray burst from the X-ray transient,  XTE J1739-285. According to Kaaret et al.\\cite{Kaaret2007} \"this oscillation frequency suggests that XTE J1739-285 contains the fastest rotating neutron star yet found\". If confirmed, this would be the first detection of a sub-millisecond pulsar (discovery of a 0.5 ms pulsar in SN1987A remnant announced in January 1989 was withdrawn one year later).  Rotation at $f>1000$ Hz is sensitive to the stellar mass and to the equation of state (EOS).  Hydrostatic, stationary configurations of neutron stars rotating at given rotation frequency $f$ form a one-parameter family, labeled by the central density. This family - a curve in the mass - equatorial radius plane - is limited by two instabilities. On the high central density side, it is instability with respect to axi-symmetric perturbations, making the star collapse into a Kerr black hole. The low central density boundary results from the mass shedding from the equator.  In the present paper we show how rotation at $f>1000$ Hz is sensitive to the EOS, and what constraints on the EOS of neutron stars result from future  observations of stably rotating sub-millisecond pulsars. ",
        "conclusions": "The $M(R_{\\rm eq})$ curve for $f\\gtrsim 1400$ Hz is flat. Therefore, for given EOS the mass of NS is quite well defined. Conversely, measured mass of a NS rotating at $f\\gtrsim 1400$ Hz will allow us to unveil the actual EOS. The \"Newtonian\" formula for the Keplerian frequency works surprisingly well for precise 2-D simulations and sets a firm upper limit on $R_{\\rm eq}$ for a given $f$. Finally, observation of $f\\gtrsim 1200$ Hz  sets stringent limits on the initial mass of the nonrotating star which was spun up to this frequency by accretion. \\label{sect:discuss}"
    },
    "0710/0710.0621_arXiv.txt": {
        "abstract": "X-ray spectra from stellar coronae are reprocessed by the underlying photosphere through scattering and photoionization events.  While reprocessed X-ray spectra reaching a distant observer are at a flux level of only a few percent of that of the corona itself, characteristic lines formed by inner shell photoionization of some abundant elements can be significantly stronger.  The emergent photospheric spectra are sensitive to the distance and location of the fluorescing radiation and can provide diagnostics of coronal geometry and abundance.  Here we present Monte Carlo simulations of the photospheric K$\\alpha_1,\\alpha_2$ doublet arising from quasi-neutral Fe irradiated by a coronal X-ray source.  Fluorescent line strengths have been computed as a function of the height of the radiation source, the temperature of the ionising X-ray spectrum, and the viewing angle.  We also illustrate how the fluorescence efficiencies scale with the photospheric metallicity and the Fe abundance.  Based on the results we make three comments: (1) fluorescent Fe lines seen from pre-main sequence stars mostly suggest flared disk geometries and/or super-solar disk Fe abundances; (2) the extreme $\\approx 1400$~m\\AA\\ line observed from a flare on V~1486~Ori can be explained entirely by X-ray fluorescence if the flare itself were partially eclipsed by the limb of the star; and (3) the fluorescent Fe line detected by {\\it Swift} during a large flare on II~Peg is consistent with X-ray excitation and does not require a collisional ionisation contribution.  There is no convincing evidence supporting the energetically challenging explanation of electron impact excitation for observed stellar Fe~K$\\alpha$ lines. ",
        "introduction": "\\label{s:intro}  It is well-established from surveys of the sky at EUV and X-ray wavelengths that all stars with spectral types later than mid-F, except for giants later than mid-K, possess hot outer atmospheres akin to that of the Sun \\citep[e.g.][]{Vaiana.etal:81, Schmitt:97}.  While much observational and theoretical effort has been devoted to understanding solar coronal spectra and, in more recent years, toward understanding stellar coronal emission and spectra, comparatively little attention has been devoted to the reprocessing and line fluorescence resulting from this coronal emission by the underlying solar and stellar photospheres.  In contrast, considerable effort has been spent on the study of X-ray reprocessing by ``cold'' gas in much more complex systems with more prominent fluorescent features but more uncertain geometries and physical conditions, such as the accretion disks around black holes and non-degenerate objects in X-ray binaries \\citep[e.g.][]{Felsteiner.Opher:76,Hatchett.Weaver:77,Fabian.etal:89, George.Fabian:91,Laor:91, Matt.etal:97, Ballantyne.etal:02, Beckwith.Done:04, Cadez.Calvani:05, Dovciak.etal:04, Laming.Titarchuk:04, Brenneman.Reynolds:06}. The processes involved in photospheric fluorescence by coronal irradiation are the same as those discussed in these works; the main difference here is in the specific geometry of the X-ray source above a quasi-neutral sphere and of the extended, shell-like nature of the coronal source above the photosphere during quiescent conditions.  X-rays emitted from a hot ($T\\ga 10^6$~K) corona incident on the underlying photosphere can undergo either Compton scattering or photoabsorption events through the ionization of atoms or weakly ionized species. Through scattering events, photons can be reflected back in a direction towards the stellar surface where they have a finite chance of escape.  Compton scattering redistributes the spectrum to lower energies by $\\sim E^2/m_ec^2$ per collision, where $E$ is the photon energy and $m_e$ the electron rest mass.  The spectrum reflected from a stellar surface by scattering events is then shifted and broadened towards lower energies.  Photoionization events involving X-ray photons directed toward the photosphere are predominantly inner shell interactions with astrophysically abundant elements, the outer and valance cross-sections being very small at these energies.  Observable fluorescent lines can then arise as a result of the finite escape probabilities of photons emitted in outward directions by hole transitions in these atoms photoionized in their inner shells.  These processes have been described in the solar context by, e.g., \\citet{Tomblin:72} and \\citet{Bai:79} (B79), and more recently for arbitrarily photoionized slabs by \\citet{Kallman.etal:04}.  The strongest of the fluorescent lines for a plasma of approximately cosmic composition is the $2s$-$1p$ 6.4~keV Fe K$\\alpha$ doublet occurring following ejection of a $1s$ electron.  It has been observed in solar spectra on numerous occasions \\citep[e.g.][]{Neupert.etal:67, Doschek.etal:71, Feldman.etal:80, Tanaka.etal:84, Parmar.etal:84, Zarro.etal:92}.  The mechanism of fluorescence by the thermal X-ray coronal continuum was suggested by \\citet{Neupert.etal:67}, and was firmly established on more theoretical grounds by \\citet{Basko:78,Basko:79} and B79. \\citet{Parmar.etal:84} provided convincing observational confirmation based on flare spectra obtained by the {\\it Solar Maximum Mission}, though it has also been noted that contributions from non-thermal electron impact might also be present during hard X-ray bursts \\citep*[e.g.][]{Emslie.etal:86, Zarro.etal:92}.  B79 pointed out that, for a given source spectrum, the observed flux of Fe~K$\\alpha$ photons from the photosphere depends on essentially three parameters: the photospheric iron abundance; the height of the emitting source; and the heliocentric angle between the emitting source and observer.  \\citet{Phillips.etal:94} used Fe K$\\beta$ observations to probe the difference between the photospheric and coronal iron abundance for flares observed by the {\\it Yohkoh} satellite.  More importantly for the stellar case, the spatial aspects of photospheric fluorescent line formation suggest its application to understanding the spatial distribution of coronal structures and flares on stars of different spectral type and activity level to the Sun \\citep{Drake.etal:99}.  Indeed, Fe~K fluorescence has recently been detected during flares on the active binary II~Peg \\citep{Osten.etal:07} and on the single giant HR~9024 \\citep{Testa.etal:07}.  The Fe~K line has also been seen in a growing sample of pre-main sequence (PMS) stars \\citep{Imanishi.etal:01,Favata.etal:05,Tsujimoto.etal:05, Giardino.etal:07}, in which the line is thought to originate predominantly from the irradiated protoplanetary disk rather than the photosphere.    Since the work of \\cite{Bai:79}, there have been no concerted efforts to extend models of photospheric fluorescence for coronal excitation sources with other characteristics.  Fluorescent lines other than Fe~K$\\alpha$ have also not yet, to our knowledge, been studied by other workers in any detail in this context.  A reasonably strong feature observed in solar spectra near 17.62~\\AA\\ had been identified with Fe L$\\alpha$ photospheric fluorescence \\citep{McKenzie.etal:80, Phillips.etal:82}, but calculations of the expected line strength was shown by \\citet{Drake.etal:99} to be much too weak to explain the feature, and these authors instead identified the line with a transition in Fe~XVIII arising from configurational mixing and both seen in {\\it Electron Beam Ion Trap} spectra and predicted by theory \\citep{Cornille.etal:92}.  However, given the potential diagnostic value of photospheric fluorescence, other lines, such as O~K$\\alpha$, are possibly observable with very high quality observations and warrant further study.  The capabilities of current X-ray missions such as {\\it Chandra}, {\\it XMM-Newton}, {\\it Swift} and {\\it Suzaku} to detect fluorescent lines further motivates a re-examination of the photospheric fluorescence problem in the context of stellar coronae, photospheres and protoplanetary disks.  We restrict the study in hand to Fe~K fluorescence from stellar photospheres and defer detailed discussions of protoplanetary disks and fluorescence from other elements to future work. ",
        "conclusions": "The main scientific motivation for this work is to provide the foundation to use Fe fluorescence as a quantitative diagnostic of coronal and flare geometry.  There now exists a handful of detections of fluorescent emission from stars.  Sensitivity is currently limited to a large extent by the low spectral resolution of available instruments and progress is expected to accelerate dramatically with the future availability of X-ray calorimeters.  Since modern X-ray spectral analyses based on low-resolution CCD pulse-height spectra tend to express line strengths in terms of the line equivalent width, we have computed this quantity for the case of $\\theta=0$ and the ranges of heights and X-ray temperatures investigated in \\S\\ref{s:newcalcs}.  The equivalent width in this context refers to the fluorescent, processed line seen on top of the continuum of the ionising coronal spectrum. While the $\\theta=0$ case gives the most optimistic line strength, we note that $f(\\theta)$ is quite slowly varying for angles $\\theta \\la 45^\\circ$ for the flare heights for which significant Fe\\,K${\\alpha}$ might be observed.  The equivalent widths are illustrated in Figure~\\ref{f:ew} and listed in Table~\\ref{t:fekaEW1}.  \\subsection{Fluorescence from Pre-Main Sequence Stars}  The observability of the cold Fe~K$\\alpha$ line is of course strongly dependent on the quality of the X-ray spectrum obtained.  The most extensive study of PMS Fe fluorescence to date is that based on Chandra observations of the Orion Nebula Cluster by \\citet{Tsujimoto.etal:05}.  This study detected significant 6.4~keV excesses attributable to Fe fluorescence for 7 out of 127 sources found to have significant counts in the 6-9~keV band.  Equivalent widths were in the range 110-270~eV at plasma temperatures of $\\sim 3$-10~keV.  There is clearly a strong selection effect here and these fluorescent line strengths likely represent the upper end of the distribution.  Our calculations for a flare at scale height $h=0 R_\\star$ are also appropriate for a flare occurring above an infinite plane, such as might approximate a disk-encircled PMS star. As in the photospheric case, the fluorescence problem can be treated orthogonally from the ionisation structure of the disk, which is not greatly altered from its very largely neutral overall state by X-rays from a typical flare.  Any small degree of X-ray photoionisation will also not affect Fe~K$\\alpha$ line strengths because fluorescence yields are essentially invariant for lower Fe ions. Our calculations indicate that attaining equivalent widths much in excess of 100~eV is not straightforward for such a simple geometry for the plasma temperatures observed in the fluorescing X-ray spectra.  This finding is in agreement with earlier calculations by \\citet{Matt.etal:91} and \\citet{George.Fabian:91}, who find equivalent widths of $\\sim 150$~eV for a flat disk illuminated by X-rays with power-law photon spectral energy distributions.  There are at least four ways in which equivalent widths might be elevated above the values we find: (1) super-solar Fe abundance in the disk material, possibly arising as a result of an elevated dust-to-gas ratio; (2) disk flaring, resulting in a solid angle coverage $> 2\\pi$; (3) line-of-sight obscuration of the central flaring source (but not fluorescent line photons) by optically thick structures such as the star itself (ie the flare occurring on the far hemisphere); and (4) fluorescence contributions from ionisation by non-thermal electrons.  By analogy with the solar case, in which excitation by non-thermal electrons is usually negligible \\citep{Parmar.etal:84,Emslie.etal:86} and is much more difficult on energetic grounds, we consider (4) the {\\em least} plausible of these.  \\citet{Ballantyne.Fabian:03} have also shown in the accretion disk context that Fe~K production by non-thermal electron bombardment requires 2-4 orders of magnitude greater energy dissipation in the electron beam than is required for an X-ray photoionization source.  Disk flaring can give rise to increased line strengths by factors $< 2$ simply from consideration of the increased solid angle coverage possible compared with an infinite flat disk.  It is also difficult to envisage enhanced Fe abundances in the disk being able to elevate line strengths by more than a factor of a few.  Of some interest, then, is the observation of an Fe~K$\\alpha$ line equivalent width of $\\approx1400$~m\\AA\\ during the rise phase of a flare on the PMS Orion nebula star forming region object V~1486~Ori by \\citet{Czesla.Schmitt:07}.  Such an enhancement over an infinite disk value of $\\sim 150$~m\\AA , even with a large degree of disk flaring, would still require extreme enhancements of the disk Fe abundance by an order of magnitude or more \\citep[e.g. \\S\\ref{s:fesens} and][]{Matt.etal:97} were the line due to photoionisation by the {\\em directly observed} continuum.  We point out, however, that fluorescent line photons from a PMS disk can still be observed when the X-ray flaring source is located behind the star and obscured from the line-of-sight.  In the case of the V~1486~Ori flare, the large observed Fe~K$\\alpha$ equivalent width is simply and plausibly associated with a partially obscured flare whose rise phase was not fully observed directly owing to line-of-sight obscuration by the star itself.  Such obscured flares will inevitably be the cause of some fraction of observed Fe~K$\\alpha$ lines from PMS stellar disks.  This explanation is also more consistent than one relying on preferential disk geometry and Fe abundance with the non-detection of Fe~K$\\alpha$ from a second less extreme flare whose impulsive phase instead appears to have been quite visible.    \\subsection{Fluorescence from stellar photospheres}    The strong flare in II~Peg observed by {\\it Swift} and analysed by \\citet{Osten.etal:07} presents another interesting case.  Photospheric Fe~K$\\alpha$ was clearly detected throughout the event.  Equivalent widths for different times in the flare ranged from 18 to 61~eV, with uncertainties in the 20-45\\%\\ range.  The authors favoured a collisional excitation mechanism for the line, arguing that fluorescence would be unlikely to produce an observable feature.  This assessment employed a simple analytical formula applicable to optically-thin cases in which only a minor fraction of the incident X-ray flux is subject to photoabsorption or scattering \\citep[see e.g.][]{Liedahl:99,Krolik.Kallman:87}.  \\citet{Osten.etal:07} correctly noted that the path length required to obtain the observed equivalent widths under such an approximation was similar to the $\\tau=1$ Compton scattering depth, but discounted fluorescence as a possibility on these grounds.  Other than the inapplicability of the optically-thin formula for the photospheric fluorescence case, one reason such an argument is overly pessimistic is that incidence angles on the photosphere range from $\\sim 0$--$90^\\circ$ for small scale heights and path lengths for escape are less than penetration depths by the factor of the inverse cosine of these angles.  We defer a more detailed treatment of the event to future work, but note here that equivalent widths of 50~eV are achieved for flare heights up to $\\sim 0.2 R_\\star$ or so for the $10^8$~K model in our grid, a temperature similar to the average of the values found for the flare by \\citet{Osten.etal:07}.  While collisional ionisation cannot be ruled out observationally as the source of the observed Fe fluorescence, it is not a requirement."
    },
    "0710/0710.5374_arXiv.txt": {
        "abstract": "While  isolated   neutron  stars   (INSs)  are  among   the  brightest $\\gamma$-ray sources, they are among the faintest ones in the optical, and their study is a  challenging task which require the most powerful telescopes.  \\hst\\  has lead  neutron star optical  astronomy yielding nearly all  the identifications achieved since the  early 1990s. Here, the major \\hst\\ contributions in the optical studies of INSs and their relevance for neutron stars' astronomy are reviewed. ",
        "introduction": "\\label{sec:1}  Before  the  launch  of  \\hst,   optical  studies  of  INSs  were  the exception. In the first 20 years since the pulsars discovery, only the Crab and Vela pulsars were identified (Cocke 1969; Lasker 1976), while optical pulsations  were detected from  an unidentified source  at the center of SNR  B0540-69 in the LMC (Middleditch  \\& Pennypacker 1985), and  only  a  candidate  counterpart  was  found  for  the  misterious $\\gamma$-ray source  Geminga (Bignami et  al.  1987).  This  score was expected  to be  considerably improved  by  \\hst, thanks  to its  much larger sensitivity with respect to ground based telescopes, and to the sharp spatial resolution of the \\wfpc\\  as well as to the near-UV view of  the ESA's \\foc.   Unfortunately, the  spherical aberration  of the \\hst\\ optics affected the execution of most approved proposals, except for those aimed  at the brightest targets. So, in  the early 1990s the leadership in  the INSs  optical astronomy was  still in the  hands of ground-based  observatories, mainly in  those of  the ESO  \\ntt\\ which secured the identification of Geminga through the proper motion of its counterpart, a technique soon become  the standard one, and the likely identifications  of the  optical pulsar  in  SNR B0540-69  and of  PSR B0656+14 (see  Mignani et al.   2000).  However, the  refurbishment of \\hst\\ in SM-1 (Dec. 2003) brought its performance back to the original expectations and  gave it a  leading role in INSs'  optical astronomy, mantained even after  the advent of the 10-m  class telescopes.  Since \\hst\\ has provided 8 new  INSs identifications, against the 2 of the \\vlt\\  and  the  \\keck\\  (see  Mignani et  al.   2004),  boosting  the identification rate by  a factor 4.  This could have been higher if not for the \\foc\\ removal in SM-3B (March 2002) and for the \\stis\\ failure  (Aug. 2004), which  alone have yielded nearly  all the \\hst\\  INSs identifications,  depriving the  telescope of  its near-UV view.   Thus, \\hst\\  observations have  opened wide  a  new, important observing window  on INSs and triggered  the interest of  a larger and larger fraction of the neutron star community. ",
        "conclusions": ""
    },
    "0710/0710.2248_arXiv.txt": {
        "abstract": "{} {We intend to establish the X-ray properties of Swift J0732.5--1331 and therefore confirm its status as an intermediate polar.} {We analysed 36\\,240~s of X-ray data from {\\em RXTE}. Frequency analysis was used to constrain temporal variations and spectral analysis used to characterise the emission and absorption properties.} {The X-ray spin period is confirmed to be 512.4(3)~s with a strong first harmonic. No modulation is detected at the candidate orbital period of 5.6~h, but a coherent modulation is present at the candidate 11.3~h period. The spectrum is consistent with a 37~keV bremsstrahlung continuum with an iron line at 6.4~keV absorbed by an equivalent hydrogen column density of around $10^{22}$ atoms~cm$^{-2}$.} {Swift J0732--1331 is confirmed to be an intermediate polar.} ",
        "introduction": "Intermediate polars (IPs) are a sub-class of cataclysmic variables (CVs). They fill the phase space, in terms of magnetic field strength, and spin and orbital periods, between non-magnetic CVs and the strongly magnetic synchronously rotating polars. The magnetic field strength is believed to be in the range of a few MG to tens of MG at the white dwarf surface. This is large enough to dramatically alter the accretion flow, yet not large enough to synchronize the spin and orbital periods. This magnetic field gives rise to the defining characteristic of the sub-class, that of X-ray variation pulsed at the spin period of the white dwarf. For an exhaustive review of CVs see e.g. \\cite{warner95}.  There are between twenty six\\footnote{http://asd.gsfc.nasa.gov/Koji.Mukai/iphome/iphome.html as of 23/8/7.} and fifty IPs currently known (depending on the selection criteria used). The hard X-ray selected object, Swift~J0732.5--1331 (hereafter J0732), is a suspected IP in need of confirmation. The circumstance of its discovery makes J0732 similar to the host of candidate IPs that have been discovered to be powerful emitters of hard X-rays/soft gamma-rays in the 20--100~keV range in the {\\sl INTEGRAL\\/}/IBIS survey \\citep{barlow06}. We have embarked on a campaign of pointed {\\sl RXTE\\/} observations of these hard X-ray discovered candidate IPs. Here we present the first results of our campaign on J0732. ",
        "conclusions": "The unambiguous X-ray spin period detection at 512.4(3)~s, along with the spectral fit to an absorbed 37~keV bremsstrahlung model with an iron line, confirm the intermediate polar status of Swift J0732.5--1331. We are unable to determine the orbital period from these {\\em RXTE} data although there is some indication of modulation at the previously suggested photometric period of 11.3~h and none at the spectroscopic period of 5.6~h. To conclude we note that this system is similar, in terms of its small $P_{\\rm spin}/P_{\\rm orb}$ value, to the IPs RX~J2133.7+5107 and NY~Lup (IGR~J15479--4529). Both of these are {\\em INTEGRAL} hard X-ray sources and both also have soft X-ray components. We might therefore expect that Swift~J0732.5--1331 would also display such characteristics upon further study."
    },
    "0710/0710.0767_arXiv.txt": {
        "abstract": "We investigate the possibility that near future observations of ultra-high-energy cosmic rays (UHECRs) can unveil their local source distribution, which reflects the observed local structures if their origins are astrophysical objects. In order to discuss this possibility, we calculate the arrival distribution of UHE protons taking into account their propagation process in intergalactic space i.e. energy losses and deflections by extragalactic magnetic field (EGMF). For a realistic simulation, we construct and adopt a model of a structured EGMF and UHECR source distribution, which reproduce the local structures actually observed around the Milky Way. The arrival distribution is compared statistically to their source distribution using correlation coefficient. We specially find that UHECRs above $10^{19.8}$eV are best indicators to decipher their source distribution within 100 Mpc, and detection of about 500 events on all the sky allows us to unveil the local structure of UHE universe for plausible EGMF strength and the source number density. This number of events can be detected by five years observation by Pierre Auger Observatory. ",
        "introduction": "\\label{intro}  The origin of ultra-high-energy cosmic rays (UHECRs) above $10^{19}$eV is one of the most intriguing problems in astroparticle physics. Akeno Giant Air Shower Array (AGASA) found statistically significant small-scale clusterings of observed UHECR events with large-scale isotropy \\citep*{takeda99}. The AGASA data set of 57 events above $4 \\times 10^{19}$eV contains four doublets and one triplet within separation angle of $2^{\\circ}.5$, consistent with the experimental angular resolution. The chance probability of observing such multiplets under an isotropic distribution is only about 1\\% \\citep*{hayashida00}. A combination of the results of many UHECR experiments (including AGASA) also revealed eight doublets and two triplets within $4^{\\circ}$ on a totally 92 events above $4 \\times 10^{19}$eV \\citep*{uchihori00}. These multiplets suggest that the origins of UHECRs are point-like sources. For identification of UHECR sources, arrival directions of UHECRs have been observed in detail by High Resolution Fly's Eye (HiRes) and Pierre Auger Observatory (Auger). However, so far, these experiments have reported no significant clustering on the arrival distribution above $4 \\times 10^{19}$eV \\citep*{abbasi05,mollerach07}.  Recently, several classes of astrophysical objects in many literature has tested for positional correlations with observed arrival directions of UHECRs. The correlations with BL Lac objects were discussed on the assumptions of smaller deflection angles of UHECRs than the experimental angular resolution and/or neutral primaries \\citep*{tinyakov01}, and in consideration with the deflection by Galactic magnetic field (GMF) \\citep*{tinyakov02}. \\cite{gorbunov05} considered various classes of powerful extragalactic sources for the correlation study including small corrections of UHECR arrival directions by GMF. \\cite{hague07} discussed the correlation with nearby active galactic nuclei (AGNs) from RXTE catalog of AGNs. However, these studies have not taken into account UHECR propagation in extragalactic space. UHECRs above $8 \\times 10^{19}$eV lose a significant fraction of their energies by photopion production in collision with the cosmic microwave background (CMB) photons during their propagation \\citep*{berezinsky88,yoshida93}. Thus, UHECRs have {\\it horizons}, which are the maximum distances of their sources that UHECRs can reach the Earth, even if their energies are below $8 \\times 10^{19}$eV at the Earth. The positional correlations between arrival directions of UHECRs and their source candidates outside the horizons are not significant. (In \\cite{hague07}, only nearby AGNs within the horizons are considered.) In addition to the UHECR horizons, deflections due to extragalactic magnetic field (EGMF) are also important since extragalactic cosmic rays are propagated for a much greater distance than in Galactic space. Propagation process of UHECRs should be considered in such correlation studies.  \\cite{yoshiguchi03} investigated the correlation between the arrival distribution of UHECRs and their source distribution taking into account UHECR propagation in intergalactic space with a uniform turbulent EGMF whose strength is 1 nG and coherent length is 1 Mpc. The authors adopted a source distribution with $10^{-6}~{\\rm Mpc}^{-3}$ that reproduced the local structures and the AGASA results. They concluded that detection of a few thousand events above $4 \\times 10^{19}$eV reveal observable correlation with the sources within 100 Mpc.  However, a uniform turbulent field is not realistic EGMF model. Faraday rotation measurements indicate magnetic field strengths at the $\\mu$G level within inner region ($\\sim$ central Mpc) of galaxy clusters \\citep*{kronberg94}. The evidence for synchrotron emission in numerous galaxy clusters \\citep*{giovannini00} and in a few cases of filaments \\citep*{kim89,bagchi02} also seems to suggest the presence of magnetic fields with $0.1-1.0\\mu$G at cosmological structures. Several numerical simulations of large-scale structure formation have confirmed these magnetic structures \\citep*{sigl03,dolag05}.  Based on these studies, in recent years, we calculated propagation of UHE protons in a structured EGMF which well reproduces the local structures actually observed and simulated their arrival distributions with several normalizations of EGMF strength and several number density of UHECR sources \\citep*{takami07}. We constrained the source number density to best reproduce the AGASA results. As a result, $10^{-5}~{\\rm Mpc}^{-3}$ is the most appropriate number density, which is weakly dependent on EGMF strength. (In rectilinear propagation, similar number density is also obtained in \\cite{blasi04,kachelriess05}) However, this has large uncertainty due to the small number of observed events at present. $10^{-4}~{\\rm Mpc}^{-3}$ and $10^{-6}~{\\rm Mpc}^{-3}$ are also statistically allowed. Therefore, it is useful to deliberate the correlation between the arrival distribution and the source distribution in the case of these number densities. Note that we revealed in the paper that this uncertainty will be solved by future increase of detected events.  In this study, we calculate the arrival distribution of UHECRs, taking their propagation process into account, and investigate the correlation the arrival distribution and their source distribution in the future. A structured EGMF model and source distribution which can reproduce the local universe actually observed are adopted. The source number density and the EGMF strength are treated as parameters since these have some uncertainty. A goal of this study is that we understand the number of observed events to start to observe the UHECR source distribution by UHECRs and how much the correlation is expected in the future.  Auger has already detected more events above $10^{19}$eV than those observed by AGASA \\citep*{roth07}. Nevertheless, the event clustering has not observed, as mentioned above. It might be due to EGMF and/or GMF strong enough not to generate the multiplets or statistical fluctuation for small number of observed events at highest energies. In any case, we should predict and discuss how the arrival distribution reflects UHECR source distribution.  Chemical composition of UHECRs is very important for the correlation. If UHECRs are heavier components, magnetic deflections are larger and the correlation is worse. One of observables for study of UHECR composition is the depth of shower maximum, $X_{\\rm max}$, which can be measured by fluorescence detectors. Its average value $\\left< X_{\\rm max} \\right>$ is dependent on UHECR composition and energy. HiRes reported that composition of cosmic rays above $10^{19}$eV is dominated by protons as a result of $X_{\\rm max}$ measurement \\citep*{abbasi05b}. Recent result by Auger is compatible to the HiRes result within systematic uncertainties \\citep{unger07}. However, they concluded that the interpretation of $\\left< X_{\\rm max} \\right>$ distribution is ambiguous because of the uncertainties of hadronic interaction at highest energies. Thus, UHECR composition at highest energies is controversial at present. Despite that, in this study, we assume that all UHECRs are protons since composition above $10^{19}$eV has proton-like feature.   This paper is organized as follows: in section \\ref{model} we provide our models of UHECR source distribution and a structured EGMF. In section \\ref{method} we explain our calculation method for the arrival distribution with UHECR propagation and statistical method. In section \\ref{results}, The results of the correlation between the arrival distribution of UHE protons and their source distribution. We summarizes this study in section \\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  In this paper, we calculated the arrival distribution of UHE protons taking into account energy losses and deflections by EGMF during propagation in intergalactic space in order to investigate the possibility that future observations of UHECRs can unveil the local structure of UHE universe. In order to reproduce a realistic situation, we adopted a structured EGMF model and source distributions which reproduce the observed local structures. The arrival distribution of UHE protons was compared statistically to their source distribution using the correlation coefficients. As the number of observed events increases, the correlation coefficient increases and converges to some value which represents the ability to unveil the source distribution by UHE protons, i.e. charged particles. Thus, the number of events that the correlation coefficient starts to converge is an important number for UHECR observations. In other words, detection of such number of events allows us to unravel UHECR source distribution. We found that UHECRs above $10^{19.8}$eV are best indicators to decipher their source distribution within 100 Mpc from discussion based on the final values of the correlation coefficients and GZK mechanism, and 5000, 500, and 200 event detections above $10^{19.8}$eV on all the sky can unveil their source distribution for the source number densities of $10^{-4}$, $10^{-5}$, and $10^{-6}~{\\rm Mpc}^{-3}$ respectively. Note that ground based detectors observe only about half hemisphere, so only half of such numbers are requested.  In this study, we took only EGMF into account as magnetic field, i.e., neglected GMF GMF deflects trajectories of UHE protons efficiently by its regular components, which consist in spiral and dipole components \\citep*{alvarez02,yoshiguchi03b}. A turbulent component of GMF very weakly change the arrival directions of UHE protons \\citep*{yoshiguchi04}. The deflection angles of UHECR protons are a few degree at around $10^{20}$eV except for the direction of Galactic Center. Such deflection disturbs the spatial pattern of UHECR arrival distribution at a few degree scale. The effect of GMF is one of our future investigations.  A lot of inquiries on UHECR source number density result in around $10^{-5}~{\\rm Mpc}^{-3}$ on ground of the AGASA results as introduced in section \\ref{intro}. If this number density is true, five year observation by Auger and future observation by TA and Extreme Universe Space Observatory \\citep*{euso} will reveal the distribution of nearby UHECR sources. The dawn of the UHE particle astronomy is just around the corner."
    },
    "0710/0710.1557_arXiv.txt": {
        "abstract": "\\noindent The field of astroparticle physics is currently developing rapidly, since new experiments challenge our understanding of the investigated processes. Three messengers can be used to extract information on the properties of astrophysical sources: photons, charged Cosmic Rays and neutrinos. This review focuses on high-energy neutrinos ($E_{\\nu}>100$~GeV) with the main topics as follows. \\begin{itemize} \\item The production mechanism of high-energy neutrinos in astrophysical shocks. The connection between the observed photon spectra and charged Cosmic Rays is described and the source properties as they are known from photon observations and from charged Cosmic Rays are presented. \\item High-energy neutrino detection. Current detection methods are described and the status of the next generation neutrino telescopes are reviewed. In particular, water and ice Cherenkov detectors as well as radio measurements in ice and with balloon experiments are presented. In addition, future perspectives for optical, radio and acoustic detection of neutrinos are reviewed. \\item Sources of neutrino emission. The main source classes are reviewed, i.e.~galactic sources, Active Galactic Nuclei, starburst galaxies and Gamma Ray Bursts. The interaction of high energy protons with the cosmic microwave background implies the production of neutrinos, referred to as GZK neutrinos. \\item Implications of neutrino flux limits. Recent limits given by the AMANDA experiment and their implications regarding the physics of the sources are presented. \\end{itemize} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.2232_arXiv.txt": {
        "abstract": "We have detected 523 sources in a survey of the Small Magellanic Cloud (SMC) Wing with {\\it Chandra}.  By cross-correlating the X-ray data with optical and near-infrared catalogues we have found 300 matches.  Using a technique that combines X-ray colours and X-ray to optical flux ratios we have been able to assign preliminary classifications to 265 of the objects.  Our identifications include four pulsars, one high-mass X-ray binary (HMXB) candidate, 34 stars and 185 active galactic nuclei (AGNs).  In addition, we have classified 32 sources as 'hard' AGNs which are likely absorbed by local gas and dust, and nine 'soft' AGNs whose nature is still unclear.  Considering the abundance of HMXBs discovered so far in the Bar of the SMC the number that we have detected in the Wing is low. ",
        "introduction": "Multi-wavelength studies of the Small Magellanic Cloud (SMC) have shown that it contains a large number of X-ray binary pulsars.  From analysis of H$\\alpha$ measurements \\citep{ken91} and supernova birth rates \\citep{fil98} the star formation rate (SFR) for the SMC is estimated to lie in the range 0.04--0.4 $M_{\\odot}$ yr$^{-1}$.  \\citet{sht05} used these upper and lower SFR estimates and the linear relation between the number of high-mass X-ray binaries (HMXBs) and the SFR of the host galaxy from \\citet{gri03} to predict the number of HMXBs expected in the SMC with luminosities $\\ge 10^{35}$ erg s$^{-1}$.  They found that between 6 and 49 of these systems should be present.  Currently $\\sim60$ known or probable HMXBs have been detected in the SMC \\citep[see e.g.][]{hab04,coe05,mcg07}.  It is believed that the considerable number of pulsars can be explained in terms of a dramatic phase of star formation, probably related to the most recent closest approach of the SMC and the Large Magellanic Cloud \\citep[LMC;][]{gar96}.  To date most of the X-ray studies of the SMC have concentrated on the Bar which has proved to be a significant source of HMXBs. These systems not only provide an homogeneous sample for study, but also give direct insights into the history of our neighbouring galaxy as they are tracers of star formation rates.  Part of the puzzle of the X-ray population of the SMC is the missing or under represented components.  In particular, there are no known low-mass X-ray binaries (LMXBs) or black hole binaries and only one confirmed supergiant X-ray binary detected to date \\citep[see also][]{mcb07a}.  A survey of the X-ray binary population of the LMC by \\citet{neg02} revealed a similar distribution (within small number statistics) to that in our galaxy - all types were present.  It is therefore important to try and identify the ``missing'' X-ray binary types in the SMC.  We recently completed the first X-ray survey in the SMC Wing with \\chan\\ (see Section \\ref{sect:obs} for more details).  A study of the brightest ($>50$ counts) X-ray sources uncovered two new pulsars, and detected two previously known pulsars \\citep{mcg07}.  In addition to the four pulsars, the sample included two foreground stars, 12 probable AGNs and five unclassified sources. We found that the pulsars had harder spectra than the other bright X-ray sources.  In this paper we report on the analysis of the whole survey and present preliminary classifications for a large fraction of the sources detected.  \\begin{figure} \\includegraphics[width=84mm]{fig1.eps} \\caption{The location of the 20 fields studied by Chandra in this work, overlaid on a neutral hydrogen density image of the SMC \\citep{sta99}.  The Wing and Bar of the SMC are marked.} \\label{fig:smc_fields} \\end{figure} ",
        "conclusions": "\\label{sect:disc}  For the 523 sources detected in the SMC Wing survey we have been able to find optical matches for 300 of them, and assign preliminary classifications to 265 objects.  Our classification method has the advantage that it does not require optical spectra, however, it still requires optical counterparts to be identified.  We also note that to classify the remaining 49\\% of the survey deeper optical surveys are needed, and in some cases better coverage of the Wing.  The majority of the Wing sources are found to be AGNs.  In the whole survey we only identify four pulsars \\citep[see][]{mcg07} and one HMXB candidate, which compared to the Bar is a small sample. The relatively few pulsars detected in the Wing is perhaps not surprising given the accepted link between regions of H$\\alpha$ and star formation, with the main regions of star formation coinciding with the high density H$\\alpha$ region in the Bar \\citep{ken95}.  However, in general, the pulsars we detected in the Wing have harder spectra than those in the Bar.  It is also remarkable that the only supergiant system so far detected in the SMC, SMC X-1, lies in the Wing.  We note that, despite appearances, the SMC is a very three dimensional object.  Studies of the Cepheid population by \\citet{lan86} have revealed that the depth of the SMC is up to 10 times its observed width.  The two main structures, the Bar and the Wing, could be separated by 10--20 kpc. Could different populations be represented in the two regions?  In the case of the HMXBs if we based our response on the X-ray results alone we could perhaps draw the conclusion that the sources in the Wing and Bar are in fact different.  However, taking into account the optical spectral analysis in which the optical counterparts for the pulsars were found to be typical of other HMXBs in the SMC \\citep{sch07,mcb07b}, different populations seem less likely.  This could imply that there is absorption local to the sources which effects the X-ray spectral results.  There is also the possibility that a greater population of HMXBs does exist in the Wing of the SMC, but we were not fortunate enough to catch more than a handful of them when they were switched on.  From our studies of 10 years of {\\it RXTE} data we find that the probability of a Be X-ray transient being in an active phase is only, on average, $\\sim 10$\\% \\citep[Figure 4.62,][]{gal06}. Quiescent X-ray transients have been detected previously in the Milky Way with luminosities $<10^{34}$ erg s$^{-1}$ \\citep[e.g.][]{neg00,cam02}.  The origin of the quiescent luminosity in Be X-ray transients is still under debate, with a number of processes suggested to account for the detected emission \\citep[see e.g.][]{cam02,kre04}.  The two mechanisms detectable from sources located in the SMC are: accretion onto the magnetospheric boundary, the propeller regime \\citep{ill75,cam00}, and very low rate accretion onto the surface of the neutron star, i.e. residual/leaking accretion \\citep[e.g.][]{ste94}.  The one HMXB candidate that we have identified has a luminosity (at the distance to the SMC) of $3.2\\times10^{33}$ erg s$^{-1}$ so it could be a quiescent source.  The lack of HMXBs in the Wing indicates that we are looking at an older population which is confirmed by optical studies of the star formation history of the SMC \\citep[e.g.][]{har04}.  In theory this should increase our chances of detecting LMXBs.  Arguably, LMXBs should be well distributed within the SMC, i.e. they should lie in the Bar and the Wing, however, deep looks of the SMC Bar \\citep{naz03} have been unsuccessful in detecting any.  The number of LMXBs expected in the SMC is proportional to the total stellar mass of the galaxy, resulting in a prediction of only one system with an X-ray luminosity of $\\geq 10^{35}$ erg s$^{-1}$ \\citep[see][]{sht05}.  However, \\citet{gar01} have shown that quiescent LMXBs can be as faint as $2\\times 10^{30}$ erg s$^{-1}$.  To go as deep as that is beyond the capability of current X-ray telescopes, but in 100 ks it would be possible to reach a limit of $\\sim 10^{32}$ erg s$^{-1}$, sufficient to detect a sample of fainter sources and study their characteristics.  If an observation like this were performed in the Wing it could be compared directly with the deep exposures of the Bar \\citep{naz03,zez05} and help quantify the LMXB population in the SMC."
    },
    "0710/0710.0371_arXiv.txt": {
        "abstract": "One well-known way to constrain the hydrogen neutral fraction, $\\bxhi$, of the high-redshift intergalactic medium (IGM) is through the shape of the red damping wing of the \\lya absorption line.  We examine this method's effectiveness in light of recent models showing that the IGM neutral fraction is highly inhomogeneous on large scales during reionization.  Using both analytic models and ``semi-numeric\" simulations, we show that the ``picket-fence\" absorption typical in reionization models introduces both scatter and a systematic bias to the measurement of $\\bxhi$.  In particular, we show that simple fits to the damping wing tend to \\emph{overestimate} the true neutral fraction in a partially ionized universe, with a fractional error of $\\sim 30\\%$ near the middle of reionization.  This bias is generic to any inhomogeneous model.  However, the bias is reduced and can even underestimate $\\bxhi$ if the observational sample only probes a subset of the entire halo population, such as quasars with large HII regions.  We also find that the damping wing absorption profile is generally steeper than one would naively expect in a homogeneously ionized universe.  The profile steepens and the sightline-to-sightline scatter increases as reionization progresses.  Of course, the bias and scatter also depend on $\\bxhi$ and so can, at least in principle, be used to constrain it.  Damping wing constraints \\emph{must} therefore be interpreted by comparison to theoretical models of inhomogeneous reionization. ",
        "introduction": "\\label{intro}  The reionization of hydrogen in the intergalactic medium (IGM) is a landmark event in the early history of structure formation, because it defines the moment at which galaxies (and black holes) affected every baryon in the Universe.  As such, it has received a great deal of attention -- both observationally and theoretically -- in the past several years.  Unfortunately, the existing observational evidence is enigmatic (see \\citealt{fan06-review} for a recent review).  Electron scattering of cosmic microwave background photons implies that reionization occurred at $z \\sim 10$, albeit with a large uncertainty \\citep{page06}.  On the other hand, \\lya forest spectra of  quasars at $z \\sim 6$ show some evidence for a rapid transition in the globally-averaged neutral fraction, $\\bxhi$ (e.g., \\citealt{fan06}).  However the \\lya absorption is so saturated in the \\citet{gunn65} trough (with optical depth $\\tau_{\\rm GP} \\ga 10^5 \\bxhi$) that constraints derived from that spectral region \\citep{fan06, maselli07} are difficult to interpret (e.g, \\citealt{lidz06, becker07}).  Another probe is the red damping wing of the IGM \\lya absorption:  the line is so saturated at these redshifts that even photons that are emitted redward of the \\lya resonance can suffer significant absorption from the strong damping wings of that transition.  This has a number of consequences for high-redshift observations.  For example, surveys that search for high-$z$ galaxies through their \\lya emission lines will find fewer and fewer galaxies as the IGM becomes more and more neutral \\citep{haiman02-lya, santos04}, although galaxy clustering strongly moderates this decline \\citep{furl04-lya, furl06-lya, mcquinn07, mesinger07-lya}.  Such surveys have now detected objects at $z \\sim 6.5$--$9$ (e.g., \\citealt{kashikawa06, iye06, stark07}), but their implications for reionization are unclear \\citep{malhotra04, haiman05-lya, malhotra06, kashikawa06, dawson07, dijkstra07, mcquinn07-lya, mesinger07-lya}.  The evolution of galaxy abundances and clustering measures the damping wing absorption in a statistical sense, but even more information can potentially be gleaned from the damping wing absorption profiles in individual objects \\citep{miralda98}.  For the galaxies described above, this information is difficult to extract because of their faintness and the complicated origins of their \\lya emission lines \\citep{mcquinn07-lya}.  However, high signal-to-noise spectra of bright objects could be extremely helpful. If the damping wing profile from IGM absorption can be isolated from these spectra, this would provide detailed information on the neutral gas along each particular line of sight (LOS) -- rather than the statistical information available from most other probes.  This is very useful, as reionization is expected to be highly inhomogeneous.  There are two candidates for such high signal-to-noise spectra at high-redshifts:  quasars and gamma-ray bursts (GRBs).  Quasars present several challenges: complicated intrinsic spectra, biased IGM environments \\citep{barkana04-grb, lidz07}, and large HII regions (which significantly weaken the damping wing absorption redward of the quasar \\lya line, and can necessitate detailed spectral analysis of the blue side of the line; \\citealt{madau00, mesinger04-mockprox}).  Nevertheless, there have already been two claims of damping wing detections in high-redshift spectra, both using quasars from the Sloan Digital Sky Survey (SDSS).  \\citet{mesinger04} detected a $\\bxhi \\gsim 0.2$ damping wing through the decreased fluctuations in the total Ly$\\alpha$ optical depth near the edge of the HII region surrounding \\qnametwoeight\\ ($z_S=6.28$).  Similarly, by simulating the optical depth distributions blueward of the Ly$\\alpha$ line center and comparing them with deep observations, \\citet{mesinger07-prox} detected the presence of a $\\bxhi \\gsim 0.033$ damping wing in the spectra of \\qnametwoeight\\ and \\qnametwotwo\\  ($z_S=6.22$).  The maximum likelihood was at $\\bxhi=1$ for both quasars.  The second set of candidates, GRBs, have fewer obstacles to overcome.   Long-duration GRBs are believed to be remnants of massive stars (and so trace the bulk of the star formation, which probably occurs in lower-mass halos with more ``typical\" IGM environments), and their afterglows have extremely simple power-law intrinsic spectra (see, e.g., \\citealt{piran05} for a review).  The event rates at high redshifts may be quite high, and cosmological time-dilation helps to identify the sources when they are still bright \\citep{bromm02-grb, ciardi00, lamb00, mesinger05}.  As a result, there is a great deal of optimism in the literature regarding their potential for damping-wing measurements (e.g. \\citealt{miralda98, barkana04-grb}).  The highest-redshift GRB afterglow observed so far (at $z \\approx 6.3$), has already been used to constrain the global neutral fraction at that time \\citep{kawai06, totani06}.  Unfortunately, this object illustrates the major difficulty with the red damping wing test for GRBs:  intrinsic absorption in the host galaxies \\citep{miralda98}.  Most GRBs are now known to have large columns of associated neutral hydrogen \\citep{vreeswijk04, chen04-grb}. Roughly $20\\%$ of well-studied objects have $N_{\\rm HI} \\la 10^{20} \\colden$ \\citep{chen07}, although nearly all of the objects in this sample are at $z \\la 6$.  The $z \\approx 6.3$ GRB does appear to have intrinsic absorption with $N_{\\rm HI} \\sim 10^{21.6} \\colden$ \\citep{totani06}, which makes it difficult to constrain the IGM absorption.  In principle, it is still possible because isolated HI absorbers have different spectral profiles than the IGM (with the optical depth inversely proportional to the wavelength offset squared for isolated absorbers, and to the wavelength offset itself for the IGM).  The two sources can then be separated by looking at the shape of the absorption.  \\citet{totani06} found a best fit with $\\bxhi=0$ and estimated that $\\bxhi \\la 0.17$ ($0.60$) at 68\\% (95\\%) confidence.  Better constraints will require faster followup (when the afterglow is brighter) and systems with less intrinsic absorption.  To date, the red damping wing test has generally been assumed to be simple and straightforward.  It is usually argued that the absorption is sensitive to a large path length in the IGM, so that small-scale clumpiness can be ignored and that the ionized fraction can be taken to be uniform (for an exception, see \\citealt{barkana02}).  However, most models of reionization have much more inhomogeneous distributions of neutral and ionized gas, with discrete HII regions surrounding clusters of galaxies, and a sea of nearly neutral gas separating them (e.g., \\citealt{arons72, shapiro87}).  Such a picture is inevitable when hot stars ionize the gas.  Moreover, the most recent models show that the ionized bubbles can become quite large even relatively early in reionization, with sizes $\\ga 10 \\Mpc$ when $\\bxhi \\sim 0.5$ \\citep{furl04-bub, furl05-charsize, iliev05-sim, zahn07-comp, mcquinn07, mesinger07}.  Because the damping wing is sensitive to fluctuations on Mpc scales, it is actually not a good approximation to take the IGM ionized fraction to be constant.  In this paper, we will examine whether (and how) the damping wing can actually be used to constrain the reionization process.  We summarize the basic physics of the line in \\S \\ref{lya}.  We then examine a series of toy models of the ``picket-fence\" absorption typical of the IGM during reionization in \\S \\ref{toy}.  In particular, we show that interpreting measurements with the naive view of a uniform IGM is not only subject to significant scatter (from the different networks of ionized bubbles intersected along different lines of sight) but also a substantial systematic bias.  In \\S \\ref{sim}, we describe the ``semi-numeric\" simulations used to generate our main results, which we present in \\S \\ref{results}.  This more detailed picture confirms that scatter between different lines of sight and bias relative to the naive view will be critical in interpreting any observed sources.  Finally, we conclude in \\S \\ref{disc}.  When this project was nearing completion, we learned of a similar effort by \\citet{mcquinn07-damp} and refer the reader there for a complementary discussion.  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}.  Unless otherwise specified, we use comoving units for all distances. ",
        "conclusions": "\\label{disc}  In this paper, we have examined how the shape of the \\lya red damping wing can be used to constrain the IGM before reionization is complete.  In the past, it has usually been assumed that the absorbing gas can be well-approximated by a uniform density medium with constant ionized fraction.  However, recent reionization models have shown that ionized bubbles can be quite large, so the latter is not a good approximation.  We have therefore critically examined how well the damping wing constrains the neutral fraction during inhomogeneous reionization.  We have identified two major issues with its interpretation.  First, there is substantial scatter in the optical depth along different lines of sight.  Most of this is due to the scatter in the distance between the source and the nearest patch of neutral gas; however, there is still non-negligible scatter even if this distance can be measured from the shape of the damping wing.  In our semi-numeric simulations, the fractional r.m.s. fluctuation in $\\bxhi$ thus estimated increases from 0.1 to 1 over the range $0.9\\gsim\\bxhi\\gsim0.2$.   Fortunately, this statistical uncertainty can be reduced simply by finding more lines of sight.  The other problem is more severe:  we have shown that the ``picket-fence\" absorption from inhomogeneous reionization adds a systematic, and often large, \\emph{bias} to measurements of the neutral fraction.  Although the damping wing is indeed sensitive to a large path length through the IGM, it is most sensitive to the closest gas.  As a result, simple fits to the damping wing will always \\emph{overestimate} the true neutral fraction in a partially ionized universe, with an error of $\\sim 30\\%$ near the middle of reionization.  This bias is generic to any inhomogeneous model.  The bias is reduced and can even become negative if observations only probe a subset of the entire halo population, such as quasars with large HII regions.  Both the systematic and statistical uncertainty can be reduced by a careful fit to the damping wing spectral profile, which is typically steeper than the naively expected $(\\Delta \\lobs)^{-1}$ profile.  However, because the absorption typically comes from many neutral patches, a large number of parameters are required for a detailed fit, and given the relatively modest difference from the expected behavior, these will be difficult to measure, probably only possible in systems with intrinsically large optical depths. Moreover, the scatter in the profiles, even at fixed $\\tau_D$, is sufficient that large samples will be required to put strong constraints on reionization from the spectral shape.  Of course, the bias and scatter also depend on $\\bxhi$ and so can, at least in principle, be used to constrain it. For example, large dispersion in the inferred neutral fractions could be an indicator of $\\bxhi \\la 0.2$. If an independent estimate of $\\bxhi$ exists, one could reverse the direction of analysis, and use the bias and scatter to constrain the reionization model and topology.  Fortunately, for a given model of reionization, the dispersion and bias can be calibrated by theoretical models.  We therefore argue that the most efficient way to constrain reionization with the damping wing is through comparison with detailed models.  Of course, any such constraints will be model-dependent, but we believe that the morphology of reionization is now sufficiently well-understood (see, e.g., \\citealt{furl05-charsize, mcquinn07}) that these uncertainties will likely not dominate the statistical uncertainties from the small number of accessible sources, at least in the relatively near future.  For example, the reionization morphology is nearly independent of redshift \\citep{furl04-bub, mcquinn07}. Also, we have found only a modest dependence of the $x_D$ distribution on halo mass (mostly due to the variation in bubble size with mass).  However, toward the end of reionization, when the absorption is dominated by rare, narrow sheets of neutral hydrogen, the details of the radiative transfer algorithm (or an approximation to it, as in our models) and of the sample selection will be extremely important.  Nevertheless, the task is challenging, as the damping wing profile must be separated from the rapidly varying resonance absorption for quasars (as in \\citealt{mesinger04, mesinger07-prox}) or from intrinsic absorbers for GRBs.  Fortunately, in the latter case $\\sim 20\\%$ of moderate-redshift GRBs have only modest absorbers and will still be useful for these purposes \\citep{chen07}.  So far, the damping wing analysis has been performed on three high-redshift quasars: \\qnamefourtwo\\ ($z_S=6.42$), \\qnametwoeight\\ ($z_S=6.28$), \\qnametwotwo\\ ($z_S=6.22$) \\citep{mesinger04, mesinger07-prox}, as well as GRB 050904 ($z_S \\approx 6.3$) \\citep{kawai06, totani06}.  This paper highlights the need to calibrate these and future damping wing analysis with simulations of the reionization morphology.  Obviously we cannot set firm constraints without detailed simulations of the observations.  Nevertheless, the mean bias we find from our simulations seems to work in the direction of strengthening the upper limit (on $\\bxhi$) from the \\citet{totani06} measurements, and weakening the lower limit from the \\citet{mesinger04, mesinger07-prox} constraints at $\\bxhi\\lsim0.6$ (although, interestingly, it would strengthen them if $\\bxhi\\gsim0.6$; see the sign change for the bias in Fig.~\\ref{fig:high_mass_mean_sig}).  Conversely, the steeper-than-expected absorption profile seems to work in the direction of weakening the \\citet{totani06} constraints (especially because it must be distinguished from strong internal absorption) while strengthening the \\citet{mesinger04, mesinger07-prox} constraints. The absorption profile might be more relevant than the bias for these studies, as an overall bias can be partially degenerate with other free parameters in the fit:  that is, when the absorption profile \\emph{can} be detected to high precision, its shape will certainly be useful in constraining $\\bxhi$.  The scatter in both effects would probably somewhat erode the confidence contours for all of these studies. On the other hand, our model predicts large scatter between different LOSs at the end of reionization, which is consistent with the measurements at $z \\sim 6.3$. More precise limits will require simultaneous fits to the intrinsic absorption and the range of possible IGM absorber profiles, and we defer them to future work.  Another intriguing possibility is to try to measure damping wing characteristics from stacked spectra of many \\lyans-emitting galaxies.  \\citet{mcquinn07} have shown that the wing shape is difficult to separate from uncertainties in the line for individual objects, and the scatter we have described will also make the interpretation of individual faint emitters problematic.  But, if the characteristics of the population are relatively constant, stacking may increase the signal to noise sufficiently to allow a detection of a ``mean\" damping wing at each redshift, even far redward of line center.   SRF thanks Crystal Martin and Josh Bloom for conversations that stimulated this work.  We thank Z. Haiman for helpful comments on this manuscript.  This research was partially supported by grant NSF-AST-0607470."
    },
    "0710/0710.2697_arXiv.txt": {
        "abstract": "We studied the clustering properties and multiwavelength spectral energy distributions of a complete sample of 162 Ly$\\alpha$-emitting (LAE) galaxies at $z\\simeq 3.1$ discovered in deep narrow-band MUSYC imaging of the Extended Chandra Deep Field South. LAEs were selected to have observed frame equivalent widths $>$80\\AA\\ and emission line fluxes $>1.5\\times10^{-17}$ergs cm$^{-2}$ s$^{-1}$. Only 1\\% of our LAE sample appears to host AGN. The LAEs exhibit a moderate spatial correlation length of $r_0=3.6^{+0.8}_{-1.0}$Mpc, corresponding to a bias factor $b=1.7^{+0.3}_{-0.4}$, which implies median dark matter halo masses of $\\log_{10}\\mathrm{M_{med}} = 10.9^{+0.5}_{-0.9}$M$_\\odot$. Comparing the number density of LAEs, $1.5\\pm0.3\\times10^{-3}$Mpc$^{-3}$, with the number density of these halos finds a mean halo occupation $\\sim$1--10\\%. The evolution of galaxy bias with redshift implies that most $z=3.1$ LAEs evolve into present-day galaxies with $L<2.5L^*$, whereas other $z>3$ galaxy populations typically evolve into more massive galaxies. Halo merger trees show that $z=0$ descendants occupy halos with a wide range of masses, with a median descendant mass close to that of $L^*$. Only 30\\% of LAEs have sufficient stellar mass ($>\\sim3\\times10^9$M$_\\odot$) to yield detections in deep Spitzer-IRAC imaging. A two-population SED fit to the stacked $UBVRIzJK$+[3.6,4.5,5.6,8.0]$\\mu$m fluxes of the IRAC-undetected objects finds that the typical LAE has low stellar mass ($1.0^{+0.6}_{-0.4}\\times10^9$M$_\\odot$), moderate star formation rate ($2\\pm1 $M$_\\odot$yr$^{-1}$), a young component age of $20^{+30}_{-10}$Myr, and little dust ($A_V<0.2$). The best fit model has 20\\% of the mass in the young stellar component, but models without evolved stars are also allowed. ",
        "introduction": "\\label{sec:intro}  The discovery of high-redshift Ly$\\alpha$-emitting galaxies (LAEs) opened a new frontier in astronomy \\citep{cowieh98,huetal98}. Because the Ly$\\alpha$ line is easily quenched, a galaxy with detectable Ly$\\alpha$ emission is likely dust-free, i.e., in the initial phases of a burst of star formation. The Ly$\\alpha$ lines have large equivalent widths ($20${\\AA}$<$EW$_{\\mathrm{rest}}<\\sim 100${\\AA}) and broad velocity widths (100 km~s$^{-1}$$<$FWHM$<$800 km~s$^{-1}$) and are often asymmetric, indicative of high-redshift galaxies undergoing active star-formation \\citep[e.g.,][]{manningetal00, kudritzkietal00,arnaboldietal02,rhoadsetal03, dawsonetal04,huetal04,venemansetal05,matsudaetal06,gronwalletal07}.  Other high-redshift galaxy populations (including Lyman break galaxies, Distant Red Galaxies, Sub-Millimeter Galaxies) exhibit strong clustering and should evolve into elliptical and giant elliptical galaxies today \\citep{adelbergeretal05a,quadrietal07a,webbetal03}. These objects were selected by unusually strong rest-frame continuum emission in the ultraviolet, optical, and far-IR respectively, resulting in $10^{10}L_\\odot \\leq L_{bol} \\leq 10^{12}L_\\odot$ \\citep{reddyetal06b}. Such strong continua appear to occur primarily in deep potential wells that are strongly biased versus the general distribution of dark matter halos. LAEs instead offer the chance to probe the faint end of the (bolometric) luminosity function of high-redshift galaxies, which contains the majority of galaxies. The strong Ly$\\alpha$ emission line allows detection and spectroscopic confirmation of LAEs with typical bolometric luminosities $\\simeq 10^{10} L_\\odot$. A detailed calculation of the LAE luminosity function at $z\\simeq 3.1$ is given in \\citet{gronwalletal07}.  Spectral energy distribution (SED) modelling of the stacked $UBVRIzJK$ photometry of 18 LAEs in the Extended Chandra Deep Field-South (ECDF-S) \\citep{gawiseretal06b} showed the average galaxy to have low stellar mass ($<10^9$ M$_\\odot$) and minimal dust \\citep[see also][]{nilssonetal07}. LAEs have the highest {\\it specific} star formation rates (defined as SFR divided by stellar mass) of any type of galaxy, implying the youngest ages \\citep{castroceronetal06}. Because Ly$\\alpha$ emission is easily quenched by dust, LAEs have often been characterized as protogalaxies experiencing their first burst of star formation \\citep[e.g.][]{hum96}. However, the differing behavior of Ly$\\alpha$ and continuum photons encountering dust and neutral gas makes it possible for older galaxies to exhibit strong Ly$\\alpha$ emission when morphology and kinematics favor the escape of these photons towards Earth \\citep{neufeld91,haimans99,hanseno06}. This could allow older, dusty galaxies with actively star-forming regions to exhibit Ly$\\alpha$ emission with high equivalent width.  SED modelling of LAEs using Spitzer-IRAC \\citep{fazioetal04} to probe rest-frame near-infrared wavelengths, where old stars dominate the emission, has yielded mixed results. \\citet{pirzkaletal07} report extremely young ages of a few Myr and low stellar masses ($10^6$M$_\\odot<$M$_*<10^8$M$_\\odot$) from SED modelling of 9 LAEs at $4.0<z<5.7$ in the Hubble Ultra Deep Field. However, \\citet{laietal07a} performed SED fitting to 3 LAEs with IRAC detections out of a sample of 12 $z=5.7$ LAEs in GOODS-N and found ages as high as 700 Myr and significant stellar masses ($10^9<$ M$_\\odot$ $<10^{10}$), making it appear that these LAEs were not undergoing their first burst of star formation. The 9/12 LAEs lacking IRAC detections are presumably younger and less massive and might lack an evolved population. Investigating the nature of the LAEs without IRAC detections requires a stacking approach to see if the typical LAE stellar mass is low enough to have been generated in a single ongoing starburst.  Stacking will yield the best results when applied to a large statistical sample of LAEs in a region with deep IRAC imaging.  \\citet{gronwalletal07} present the largest available sample of LAEs in an unbiased field, 162 LAEs at $z=3.1$ in the ECDF-S discovered as part of the MUSYC survey (\\citealp{gawiseretal06a}, \\url{http://www.astro.yale.edu/MUSYC}). The ECDF-S has been targeted with deep narrow-band imaging and multi-object spectroscopy, complemented by public broad-band $UBVRIzJK$, Spitzer+IRAC and Chandra+ACIS-I imaging. We improve the constraints of \\citet{gawiseretal06b} using the MUSYC $UBVRIzJK$ photometry of our larger sample of LAEs and adding IRAC [3.6$\\mu$m,4.5$\\mu$m,5.8$\\mu$m,8.0$\\mu$m] Cycle 2 legacy images from SIMPLE (Spitzer IRAC/MUSYC Public Legacy of the ECDF-S, \\url{http://www.astro.yale.edu/dokkum/simple}).  This Letter summarizes our imaging and spectroscopic observations of LAEs, presents our results from clustering analysis and SED modelling, and discusses the implications for the formation process of typical present-day galaxies. We assume a $\\Lambda$CDM cosmology consistent with WMAP results \\citep{spergeletal07} with $\\Omega_m=0.3, \\Omega_\\Lambda=0.7$, $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, and rms dark matter fluctuations on $8h^{-1}$ Mpc scales given by $\\sigma_8 = 0.9$ .  All correlation lengths and number densities are comoving.  We have suppressed factors of $h_{70}$ in reporting correlation lengths, number densities, dark matter masses, stellar masses and star formation rates. ",
        "conclusions": "\\label{sec:discussion}   In CDM cosmology, galaxy formation is an ongoing process caused by merging of lower-mass dark matter halos, which may already possess stars. Finding stellar population ages of $<100$ Myr is interesting. Our analysis of halo merger trees from the Milli-Millenium simulation \\citep{springeletal05} found the median age of dark matter halos with M$>10^{10.6}$M$_\\odot$ at $z=3.1$ (defined as the age since half of the dark matter mass was accumulated) to be $\\sim$600 Myr, with only $<$10\\% of halos younger than 100 Myr. If repeated LAE phases occur, the mean halo occupation of $\\sim1$--$10$\\% can be interpreted as a ``duty cycle'' telling us what fraction of each halo's lifetime is spent in the early phases of starbursts before significant dust is generated, and the population-averaged young age of $\\sim$20 Myr would imply that this phase typically lasts $\\sim$40 Myr. Alternatively, if all dark matter halos experience a single LAE phase shortly after their ``formation'' in a major merger, the mean halo occupation implies that LAEs will only be found in the youngest $1$--$10$\\% of halos, which is barely consistent with their single-population best-fit age of 100 Myr. If LAEs represent a subset of dark matter halos selected to have ages less than 100 Myr, their observed clustering may underestimate their dark matter halo masses by up to a factor of two \\citep[see][]{gaow07}.    \\begin{figure}[h!] \\epsscale{1.25} \\plotone{f6.eps} \\caption[] { Tracks show the evolution of bias with redshift calculated using the no-merging model.  The filled circle shows our result for the bias of LAEs at $z=3.1$. Previous results at high-redshift are shown for LAEs at $z=4.5$ and $z=4.86$ (stars, \\citealp{kovacetal07} and \\citealp{ouchietal03}, respectively), LBGs at $z\\sim4$ (stars, \\citealp{ouchietal04b}), K-selected galaxies (diamonds, \\citealp{quadrietal07a}), bright LBGs at $z\\sim 3$ (triangle, \\citealp{leeetal06}), and BM, BX, and LBG galaxies (asterisks, \\citealp{adelbergeretal05b}, \\citealp{adelbergeretal05a}). Local galaxy clustering is shown for SDSS galaxies (open circles, \\citealp{zehavietal05}) and for rich clusters (cross, \\citealp{bahcalletal03}). K-band limits are in Vega magnitudes. } \\label{fig:bias} \\end{figure}    Figure \\ref{fig:bias} shows the reported bias values for LAEs to be lower than those of other $z>3$ galaxy populations \\citep[bias values determined as in][]{quadrietal07a}. The expected evolution of bias is shown for the ``no-merging'' model \\citep{fry96,whiteetal07}. A realistic amount of merging will cause the bias to drop somewhat faster, so the plotted trajectories provide an upper limit on the bias factor of a given point at lower redshifts.  This shows that typical $z=3.1$ LAEs will evolve into galaxies of at most a few times $L^*$ at $z=0$. The bias values imply that LAEs at $z=3.1$ might evolve into the subset of BX galaxies at $z\\simeq 2.2$ dimmer than $K=21.5$, which also show relatively weak clustering \\citep{adelbergeretal05a}. The $K>21.5$ BX galaxies have average M$_*=1.5\\times10^{10}$M$_\\odot$, so the $z=3.1$ LAEs would need to form stars at an average rate of 14 M$_\\odot$yr$^{-1}$ over the intervening Gyr.  This could be achieved with a constant specific SFR and no merging or with a constant SFR and $\\sim$2 major mergers.  The only previous measurements of LAE clustering in unbiased fields are at $z=4.5$ \\citep{kovacetal07} and $z=4.86$ \\citep{ouchietal03}, and this earlier LAE population appears to have significantly stronger clustering, consistent with possibly evolving into typical Lyman break galaxies at $z\\simeq 3$.  The models of \\citet{ledelliouetal06} predict stellar and dark matter masses and star formation rates for $z=3$ LAEs within a factor of two of our results, despite assuming a top-heavy IMF and a very low escape fraction $f_{esc}=0.02$ that appears inconsistent with the observed lack of dust \\citep[see][for an alternative approach]{kobayashietal07}. \\citet{maoetal07} used the stellar mass of $5\\times10^8$M$_\\odot$ observed by \\citet{gawiseretal06b} to predict LAE dark matter masses of $10^{10}<$M$<10^{11}$M$_\\odot$, in the lower end of our allowed range. The stellar ages of $\\sim 20$ Myr preferred by the two-population fit are noticeably lower than the maximum values of 100 to 500 Myr predicted by these authors, \\citet{moriu06}, and \\citet{haimans99}, but the ages of 60 to 350 Myr preferred for the case of no evolved stars would be compatible.    None of the current models and numerical simulations of LAEs \\citep[see also][]{thommesm05,razoumovs06,tasitsiomi06} predict their present-day descendants. Nonetheless, the evolution of a significant fraction of $z=3.1$ LAEs into $z=0$ $L^*$ galaxies with dark matter mass  $M_{DM}\\simeq2\\times10^{12}$M$_\\odot$ and stellar mass $M_*\\simeq4\\times10^{10}$M$_\\odot$ \\citep{ichikawaetal07} appears reasonable. Fig. \\ref{fig:mass} shows the histogram of present-day masses of dark matter halos in the Milli-Millenium merger trees that have progenitors with M$>5\\times10^{10}$M$_\\odot$ at $z=3.1$.  The median present-day halo mass is $1.2\\times10^{12}$M$_\\odot$, and this would increase if LAEs found in sub-halos of massive $z=3.1$ halos were included. \\citet{lietal07} predict that the main progenitor of a present-day $L^*$ galaxy had a dark matter mass of $\\sim10^{11}$M$_\\odot$ at $z=3$ and that these galaxies experienced several major mergers at $1.5<z<7$. To form an $L^*$ galaxy at $z=0$, several LAEs could merge while experiencing a mild reduction in average SFR, with accretion of lower-mass dark matter halos through minor mergers providing most of the final dark matter mass. However, the halo mass distribution of $z=0$ descendants in Fig.\\ref{fig:mass} is very broad, with 25th and 75th percentile values of $2.9\\times10^{11}$M$_\\odot$ and $7.6\\times10^{12}$M$_\\odot$.  While $z=0$ $L^*$ galaxies like the Milky Way are roughly the median descendants of $z=3.1$ LAEs, the descendant halos include a wide range from dwarf galaxies to rich galaxy groups.   \\begin{figure}[h!] \\epsscale{1.25} \\plotone{f7.eps} \\caption[] { Histogram of dark matter halo masses of present-day descendants of halos with M$>5\\times10^{10}$M$_\\odot$ at $z=3.1$.  The dashed lines show the median halo mass of $1.2\\times10^{12}$M$_\\odot$ and the 25th and 75th percentile values of  $2.9\\times10^{11}$M$_\\odot$  and $7.6\\times10^{12}$M$_\\odot$. } \\label{fig:mass} \\end{figure}   The typical LAE stellar mass at $z=3.1$ is lower than that of any other studied high-redshift population \\citep[see][]{reddyetal06b} but is close to that of dim ($i<26.3$) Lyman break galaxies (LBGs) at $z\\sim5$ \\citep{vermaetal07}. LAEs at $z=3.1$ have much lower star formation rate, stellar age, stellar mass, dark matter halo mass, and dust extinction than the $\\sim30$M$_\\odot$yr$^{-1}$, $\\sim500$ Myr, $\\sim2\\times10^{10}$M$_\\odot$, $\\sim3\\times10^{11}$M$_\\odot$, $A_V\\simeq1$ LBG population at $z\\sim3$ \\citep[$R<25.5$,][]{shapleyetal01,adelbergeretal05a} or the $\\sim100$M$_\\odot$yr$^{-1}$, $\\sim2$ Gyr, $\\sim10^{11}$M$_\\odot$, $\\sim10^{13}$M$_\\odot$, $A_V\\simeq2.5$ Distant Red Galaxy (DRG) population \\citep{webbetal06,forsterschreiberetal04,quadrietal07a}. The high-redshift  Sub-Millimeter Galaxies \\citep{chapmanetal03b} appear to be the most massive and dusty, with the highest SFR. LAEs may represent the beginning of an evolutionary sequence where galaxies gradually become more massive and dusty due to mergers and star formation, but most LAEs at $z=3.1$ will never reach the DRG stage since DRG stellar and dark matter masses are already greater than those of present-day $L^*$ galaxies.  The Damped Ly$\\alpha$ Absorption systems (DLAs, \\citealp{wolfeetal05}) are another high-redshift population that probes the faint end of the luminosity function. The dark matter halo masses of DLAs at $z\\sim3$  were determined by \\citet{cookeetal06b} to lie in the range $10^9<$M$<10^{12}$M$_\\odot$ i.e., $1.3<b<4$, which overlaps with the range of both $L^*$ and super-$L^*$ progenitors in Fig. \\ref{fig:bias}. At least half of the DLAs appear to have ongoing star formation \\citep{wolfeetal04} and two of the three DLAs detected in emission were seen in Ly$\\alpha$.  Further study is needed to determine the relationship between DLAs and LAEs.   The observed properties of LAEs at $z=3.1$ make them the most promising candidates to be high-redshift progenitors of present-day $L^*$ galaxies like the Milky Way. Our results suggest that LAEs are observed during the early phases of a burst of star formation, perhaps caused by a major merger of smaller dark matter halos.  The input halos appear to have already contained stars, accounting for the evolved stellar population that appears to contribute most of the LAE stellar mass, although starburst-only models are also allowed. It is clear that not all progenitors of L$^*$ galaxies were LAEs at $z=3.1$. The comoving number density of our sample of LAEs is a factor of 15 less than $\\phi^*$ for local galaxies \\citep{linetal96}, plus we expect several high-redshift halos to merge into a single galaxy today. It remains possible that all progenitors of present-day galaxies experienced an LAE phase at {\\it some} redshift. Clustering and SED studies of LAEs at various redshifts are needed to assess the validity of this hypothesis."
    },
    "0710/0710.1694_arXiv.txt": {
        "abstract": "An eight stage balanced modulation scheme for dual beam polarimetry is presented in this paper. The four Stokes parameters are weighted equally in all the eight stages of modulation resulting in total polarimetric efficiency of unity. The gain table error inherent in dual beam system is reduced by using the well known beam swapping technique. The wavelength dependent polarimetric efficiencies of Stokes parameters due to the chromatic nature of the waveplates are presented. The proposed modulation scheme produces better Stokes $Q$ and $V$ efficiencies for wavelengths larger than the design wavelength whereas Stokes $U$ has better efficiency in the shorter wavelength region. Calibration of the polarimeter installed as a backend instrument of the Kodaikanal Tower Telescope is presented. It is found through computer simulation that a $14\\%$ sky transparency variation during calibration of the polarimeter can introduce $\\approx 1.8\\%$ uncertainty in the determination of its response matrix. ",
        "introduction": "\\label{sec:intro} Polarimetric accuracy is one of the most important goals in modern astronomy. It is limited since most optical elements encountered by the light on its path from the source to the detector, can alter its state of polarization (for eg. telescope optics, imaging system, grating, etc). Apart from these, variation in sky transparency, image motion and blurring due to the atmosphere  are a major concern in high precision ground based solar polarimetry. The effect of atmosphere, which is commonly known as seeing induced effect, can be reduced by fast modulation schemes(Stenflo and Povel, 1985). The modulation frequencies in these schemes are generally higher than seeing fluctuations, which is $\\approx 1kHz$(Stenflo and Povel, 1985 and Lites, 1987). Large format CCDs, which are required to cover reasonable spectral and spatial range, will pose difficulty in reading out the data at $kHz$ speed. Stenflo and Povel(1985) proposed a scheme whereby rapidly modulated signal is demodulated by optical means, thereby avoiding the need to read the detectors at a rapid rate. Lites(1987) has proposed a rotating waveplate modulation scheme as an alternative to minimize the seeing induced cross-talk among Stokes parameters. There he has shown that the faster the rotation rate of the modulator, the lower the cross-talk among Stokes parameters. And the seeing induced cross-talk levels of a dual beam system are factors 3-5 smaller than those of a single beam system. However, in dual beam system, the error introduced due to flat field residual is a matter of concern in high precision polarimetry. A possible solution to the above mentioned problems can be found by using a mixed scheme in which spatial and temporal modulations are performed(Elmore et al. 1992, Martinez Pillett et al. 1999 and Sankarasubramanian et al. 2003). The gain table uncertainties are avoided using the beam swapping technique(Donati et al. 1990, Semel et al. 1993 and Bianda et al. 1998)  A low cost dual beam polarimeter has been installed as a backend instrument for the Kodaikanal Tower Telescope (KTT). Different modulation schemes were studied and an optimum scheme is identified. The proposed scheme requires eight stages of modulation of input light in order to obtain the maximum polarimetric efficiency. Laboratory experiments have been performed to verify the theoretical understanding of the proposed scheme. The studies are extended to other wavelengths apart from the design wavelength of $\\lambda6300$.   The outline of this paper is as follows. The proposed eight stage modulation scheme for the measurement of general state of polarization is discussed in section(2). Wavelength dependence of the efficiency of the polarimeter in measuring Stokes parameters is presented in section(3). In section(4), the performance of the polarimeter at KTT is presented. ",
        "conclusions": "An eight stage modulation scheme to measure the general state of polarization is presented here. Beam swapping technique is incorporated in this scheme, which helps in alleviating the gain correction errors. The total polarimetric efficiency is close to unity as the Stokes parameters are weighted equally in all the stages of modulation. The final Stokes parameters are demodulated using all the stages of intensity measurements. Hence, the derived input Stokes parameters are equally weighted time averaged quantities over the time of measurement.  Since the retarders used in the polarimeter are chromatic, the efficiency of the polarimeter in measuring Stokes $QUV$ is wavelength dependent. The laboratory experiments performed to study the wavelength dependence of efficiency of the polarimeter confirms the theoretical expectations.  It is found through computer simulation that a $14\\%$ sky transparency variation can cause $\\approx 1.8\\%$ uncertainty in the elements of the polarimetric response matrix during its calibration, for the modulation/demodulation scheme presented here. The non-zero values of the off-diagonal elements are not a serious concern if those values do not change drastically in short time scales. During any solar polarimetric measurements, data for the calibration of the polarimeter are taken at least once a day. Calibration of the polarimeter are carried out on a few days over a period of 10-day and the response matrix derived over this period did not show any appreciable variations. The variations in the off-diagonal elements are less than the fit errors ($<$ 0.5\\%). The polarization signals observed on the Sun is always less than 40\\% and hence an uncertainty of 0.5\\% in the calibration will produce an inaccuracy of 0.2\\% in the polarization signals.  The measured total polarimetric efficiency of the polarimeter installed at KTT is $\\approx 0.986$ at $\\lambda6563$ wavelength region which is better than some of the polarimeters such as ZIMPOL(0.72), ASP(0.88), TIP(0.92) and POLIS(0.84)."
    },
    "0710/0710.3691_arXiv.txt": {
        "abstract": "We developed a source detection algorithm based on the Minimal Spanning Tree (MST), that is a graph-theoretical method useful for finding clusters in a given set of points. This algorithm is applied to $\\gamma$-ray bidimensional images where the points correspond to the arrival direction of  photons, and the possible sources are associated with the regions where they clusterize. Some filters to select these clusters and to reduce the spurious detections are introduced. An empirical study of the statistical properties of MST on random fields is carried in order to derive some criteria to estimate the best filter values. We introduce also two parameters useful to verify the goodness of candidate sources. To show how the MST algorithm works in the practice, we present an application to an EGRET observation of the Virgo field, at high galactic latitude and with a low and rather uniform background, in which several sources are detected. ",
        "introduction": "Telescopes for satellite-based high energy  $\\gamma$-ray astronomy detect individual photons by means of the electron-positron pair that they generate through the detector. From the pair trajectories it is possible to reconstruct the original direction of the photon with an uncertainty that decreases with the energy, from a few degrees below 100 MeV to less than a degree above 1 GeV. This technique was applied to the past $\\gamma$-ray observatories SAS-2 (Fichtel et al. 1975), COS-B (Bennett 1990) and EGRET-CGRO (Kanbach et~al. 1988; Thompson et~al. 1993), all equipped with spark chambers. Pair tracking is also used in the current AGILE mission (Tavani et al. 2006) and in the LAT telescope on board the next GLAST mission, both employing silicon microstrip detectors (Gehrels et al. 1999). The resulting product is an image where each photon is associated with a direction in the sky: discrete sources thus correspond to regions in which a number of photons higher than those found in the surroundings are observed. When the size of this region is consistent with the instrumental Point Spread Function the source is considered as point-like, otherwise it can be extended or a group of near sources.  Various algorithms are applied to the detection of point-like or extended sources in $\\gamma$-ray astronomy: the most extensively used one is based on the Maximum Likelihood (Mattox et al. 1996), whereas others based on Wavelet Transform analysis (Damiani et al. 1997), Optimal Filter (Sanz et al. 2001), Scale-Adaptive Filter (Herranz et al. 2002), etc., were variously applied to real and simulated data to study their performances. Some of them are based on deconvolution techniques of the instrumental Point Spread Function (PSF). Many methods work directly on the pixellated images, i.e. count or intensity maps. Other methods search for clusters in the arrival directions of photon that, if statistically significant, are considered an indication of a source.  The approach considered by us is essentially a cluster search based on a \\emph{minimal spanning tree} (MST) algorithm. This technique has its root in graph theory, and highlights the \\emph{topological} pattern of connectedness of the detected photons. Given a graph $G(V, E)$, where V is the set of vertices (or \\emph{nodes}) and E is the set of weighted \\emph{edges} connecting them, a MST (Kruskal 1956; Prim 1957; Zahn 1971) is the tree (a subgraph of $G$ without closed circuits) that connects all the points with the minimum total weight, defined as the sum of the weight of each tree's edge. In a data set consisting of points in a Cartesian frame of reference, we can consider them as the nodes of a graph, the edges being the lines joining the nodes, weighted by their length.   The MST method was originally proposed for $\\gamma$-ray source detection by Di~Ges\\`u and Sacco (1983), who investigated also the statistical properties in uniform fields. This work was developed by Di~Ges\\`u and Maccarone (1986), and De~Biase et al. (1986) applied MST for detecting extended sources in EXOSAT X-ray images. Other authors applied MST methods to the goal of finding galaxy clusters, both in 2 and 3-dimensional surveys and simulations (Barrow et al. 1985; Bhavsar \\& Ling 1988a,b; Plionis et al. 1992; Krzevina \\& Saslaw 1996, Doroshkevich et al. 2001, 2004) and showed the capabilites of the method as a filament-finding algorithm.  In this paper we investigate the MST approach in $\\gamma$-ray source detection, and present a new study of its statistical properties and the definition of selection criteria. We also introduce some parameters useful to classify the reliabilty of detected clusters to be associated with source candidates. We would like to emphasize here that this method is not \\emph{alternative} to other source detection algorithms, but it is \\emph{complementary}, in the sense that it can give a list of possible candidate sources (identified via their photons' clusterization properties) that could be further investigated by other means.   This paper is structured as follows. In Sect.~2 we describe our MST algorithm, and in Sect.~3 and 4 we investigate by means of numerical simulations the statistical distributions of edge length and node number, and we introduce some criteria useful for the source detection with our method. An example of application to an EGRET field is shown in Sect.~5, while in Sect.~6 we summarize and discuss our results. ",
        "conclusions": "We presented an application of a Minimal Spanning Tree algorithm to the problem of source detection in $\\gamma$-ray images. This method does not involves in the computation the instrumental response functions and works recognizing the regions of the sky where arrival directions of photons clusterize. It has the advantages of a fast calculation but did not provide directly estimates of the source flux. We have shown that a MST based algorithm is a viable method to detect $\\gamma$-ray sources both in simulated images and in real $\\gamma$-ray observations of the EGRET experiment on board Compton-GRO. We proposed some tools to optimize the filtering parameters and to assess the reliability of source detections, like the clustering degree and the bootstrap detection stability. These tools are based on a study, although empirical, of the statistical properties of the Minimal Spanning Tree on random fields.  The MST application to an EGRET field around the two famous $\\gamma$-ray loud quasars 3C~273 and 3C~279 found almost all the 3EG sources already detected in the same pointing and confirmed the presence of another source, detected in a different pointing. We consider this result a good indication that MST method is particularly efficient. We found also evidence of a new possible source with a significance comparable to that of other well established sources. We expect that future experiments with a better sensivity, like the LAT instrument on board GLAST, will confirm or disprove this finding.  There are, however, several possible effects that make difficult the source detection and require even more attention when the MST method is used. These problems can be divided into four main categories: $i$) problems due to the presence of strong sources, $ii$) problems arising from energy spectra of the sources different from that of the background; moreover, different spectral indices between the sources will result in different probabilities to be detected, due to the energy dependence of the PSF, $iii$) problems originated by images with a non-homogeneous background, $iv$) problems due to the geometrical distortions from the arriving celestial photons in projection onto the $\\gamma$-ray telescope, that will result not necessarily in a circular shape to characterize proper cluster selections. At present we have not developed a well established strategy to solve these problems and in the following we will briefly discuss some aspects useful for the understanding of results.  One or more strong sources in the field have various possible consequences. A first relevant effect is that they are characterized by a high clustering degree and consequently reduce the value of $\\Lambda_{\\mathrm{m}}$ with respect to the one expected in the field if they were absent. A value of $\\Lambda_{\\mathrm{c}}$ very close to $\\Lambda_{\\mathrm{m}}$ would here be good to detect strong sources but this selection criterion could miss other possible sources of lower flux. Another effect is the presence of possible ``satellites'' in the surroundings of a strong source, even closer than expected from the PSF, originated by cutting an edge whose length is just smaller than $\\Lambda_{\\mathrm{c}}$. For example, the cluster detected in the EGRET field (see Sect.~4) with no obvious counterpart, is at a distance of about 3.3$^{\\circ}$ from the strong radio quasar 3C~279, and therefore we cannot exclude that it could be a satellite of the latter. Usually, the satellites do not have a high frequency in the bootstrap fields.  The energy distribution of the photons also affects the source detection, because the PSF of $\\gamma$-ray telescopes changes with the energy becoming much narrower at high energies. This implies that sources with spectra harder than the background are better detected in high energy images because their clustering degree increases. At variance, sources with soft spectra give more disperse clusters and cannot be easily found.  Another class of problems is present when the background is markedly non-homogeneous, as in the case where the field contain a portion of the galactic disc. In this case, using an unique $\\Lambda_{\\mathrm{c}}$ in all the image would correspond to a long cutting in the dense region and to a short cutting in the region of low density with the consequence of missing real sources and producing more spurious clusters.  A general approach to be used for $\\gamma$-ray source detection is that of using several methods, possibly based on different techniques, and to compare their results. In this way it will be possible to reduce the number of spurious detections, because of the different criteria and \\emph{a priori} assumptions applied in the source recognition. Accordingly, MST method can be used to obtain a quick list of photon clusterization regions, that could correspond to possible sources, to be studied indipendently with other methods.  There are other clustering algorithms that can be applied to $\\gamma$-ray source detection, like the Voronoi tessellation (Icke \\& van de Weygaert 1987, Aurenhammer 1991). In particular, this method is based on the construction of its dual graph, the Delaunay triangulation, of which MST is a subset. We think, therefore, that at least in principle, they would provide similar result and that a combined figure of merit for source detection should be defined.  Here we discussed gamma-ray astronomy as a prime candidate for the application of MST method, but it could be even better applicable to the study of data clusterization in ultra-high energy cosmic rays (UHECR) and hemispherical neutrino experiments, that are characterized to the absence of structured background. We think also that it will be possible to extend MST to higher dimensional spaces introducing time and energy as additional dimensions. Basically there are two approaches: $i$) to search for clusters in separate, dimensionally homogeneous subspaces, and then to search for the intersection of the detected clusters and $ii$) to define a new metric for the tree edges that combine together the various dimensions in a suitable way for the MST computation. Preliminary numerical attempts based on the second approach, with energy as third coordinate, seem to be very promising to identify sources having spectra different from that of the background. Another possible 3-dimensional generalization is to take into account also the time, thus searching for variable or stable sources. We will discuss a possible application of such a generalized MST in a subsequent work."
    },
    "0710/0710.4552_arXiv.txt": {
        "abstract": "We have analyzed the redshift-dependent fraction of galactic bars over 0.2$<$z$<$0.84 in 2,157 luminous face-on spiral galaxies from the COSMOS 2-square degree field.  Our sample is an order of magnitude larger than that used in any previous investigation, and is based on substantially deeper imaging data than that available from earlier wide-area studies of high-redshift galaxy morphology. We find that the fraction of barred spirals declines rapidly with redshift.  Whereas in the local Universe about 65\\% of luminous spiral galaxies contain bars (SB$+$SAB), at $z\\sim0.84$ this fraction drops to about 20\\%. Over this redshift range the fraction of {\\em strong} (SB) bars drops from about 30\\% to under 10\\%.  It is clear that when the Universe was half its present age, the census of galaxies on the Hubble sequence was fundamentally different from that of the present day.  A major clue to understanding this phenomenon has also emerged from our analysis, which shows that the bar fraction in spiral galaxies is a strong function of stellar mass, integrated color and bulge prominence. The bar fraction in very massive, luminous spirals is about constant out to z$\\sim$0.84 whereas for the low mass, blue spirals it declines significantly with redshift beyond z=0.3.  There is also a slight preference for bars in bulge dominated systems at high redshifts which may be an important clue towards the co-evolution of bars, bulges and black holes.  Our results thus have important ramifications for the processes responsible for galactic downsizing, suggesting that massive galaxies matured early in a dynamical sense, and not just as a result of the regulation of their star formation rate. ",
        "introduction": "How, when and at what rate did the Hubble sequence form?  This question is central to the field of galaxy formation and evolution. We examine it by measuring the evolution of the bar fraction with redshift using the 2-square degree Cosmic Evolution Survey (COSMOS). In nearly all simulations, the formation timescale for a bar is rapid {\\sl once the necessary conditions (a massive, dynamically cold and rotationally-supported disk)} are met.  Therefore the redshift evolution of the bar fraction is a fundamental probe of the evolutionary history of disk galaxies.  The bar fraction is defined simply as: \\begin{equation} f_{\\rm bar} = {{\\rm number \\ of\\  barred\\  spirals}\\over {\\rm number\\  of\\  all\\ spirals}}. \\end{equation} In the local Universe the value of $f_{bar}$ is quite well established.  When only strongly barred\\footnote{Bars that are highly elliptical and have rectangular isophotes are classified as strongly barred (SB) galaxies whereas those with more oval shapes are classified as SAB or ovally distorted galaxies} galaxies (SB) are counted, the RSA, RC3 and UGC \\citep{sandage87,devau91,nilson73} all give values of $f_{bar}=0.25-0.3$. When ovally distorted (SAB) are also counted the situation becomes a little less clear-cut, because, unlike the RC3, the UGC and RSA do not attempt to carefully compile an inventory of such galaxies. If ovally distorted systems in the RC3 are included in the computation of $f_{bar}$ then the local bar fraction rises to $f_{bar}\\sim0.6$. This result is in good agreement with recent infrared studies which have measured the local bar fraction to be $\\sim$ 0.65 \\citep{eskridge00, whyte02, menendez07, marinova07}. In the infrared, a majority of the SAB galaxies are classified as strongly barred SB systems \\citep{eskridge00}.  As noted by \\citet{eskridge00} and \\citet{menendez07}, the overall bar fraction is the same in the infrared and the optical (although there is a small number of cases where bars are unveiled at infrared wavelengths). This is not surprising, because bars are primarily stellar structures whose visibility only declines sharply at ultraviolet wavelengths, short wards of the Balmer break (see also \\S \\ref{bandshifting}).  We conclude that the consensus value of the local barred fraction is $f_{bar}\\sim0.3$ for strongly barred systems, and $f_{bar}\\sim0.65$ for all barred galaxies, and that these values are so well-known that they have not changed significantly in over four decades.  In sharp contrast with the rapid and stable consensus reached on the local bar fraction, attempts to measure the bar fraction at high redshift have proven difficult.  The earliest analyses of the bar fraction in the Hubble Deep Fields (HDFs) found a dramatic paucity of barred spirals at z$>$0.5 \\citep{abraham96, vanden96, abraham99}. These authors concluded that at lookback times greater than 5 Gyr disks were either dark matter dominated or dynamically too hot (perhaps due to the increased merging activity) to host bars. However, the small volume probed by the HDFs (only thirty bright, face-on spiral galaxies between $0<z<1$) led to concerns that the bar fraction at high redshift may not be adequately measured. \\citet{sheth03} investigated whether a significant number of bars could have been missed, as suggested by \\citet{bunker99} using the H-band NICMOS HDF.  \\citet{sheth03} found four bars and two candidate bars out of 95 galaxies at z$>$0.7.  Overall, the fraction of barred spirals in the NICMOS HDF remained extremely low, as in the optical HDF studies.  But \\citet{sheth03} noted that their study was limited by the coarse NICMOS resolution (0$\\farcs$15) such that only the largest (and rarest) bars could be identified (bars with semi-major axis $>$5 kpc).  When the fraction of these large bars at z$>$0.7 was compared to local samples, there was no compelling evidence for a decline in barred spirals, but likewise the NICMOS data did not unveil any new bars at low redshifts; all except one of the four bars in the \\citet{sheth03} study are at z$>$0.9, where k-correction effects are important (\\S \\ref{bandshifting}).  A major advance in spatial resolution was possible with the Advanced Camera for Surveys (ACS) whose 0$\\farcs$05 pixels are able to resolve all but the smallest (nuclear, $<$2 kpc diameter) bars at all redshifts.  Using ACS data, two studies \\citep{elm04,jogee04} found that contrary to the previous HDF results, the bar fraction is constant at 30\\% over the last 8 Gyr (since $z=1.2$).  The sample sizes, however remained modest in these studies (186 in \\citealt{elm04}, and 258 in \\citealt{jogee04}).  In this paper we examine in detail the redshift evolution of the bar fraction using the unparalleled wide and deep 2-square degree COSMOS data set.  The plan for the paper is as follows: in Section 2 we describe our sample selection procedure. The classification methodology we have adopted is described in Section 3. Our main results are presented in Section 4, before being discussed in Section 5. Our conclusions are summarized in Section 6. An Appendix to this paper provides a detailed analysis of possible selection effects at high redshift and a discussion of our local calibration sample of 139 galaxies from the Sloan Digital Sky Survey (SDSS) Data Release 4 \\citep{adelman06}. Throughout this paper we adopt a flat $\\Lambda$-dominated cosmology with $H_0$=70 km\\ s$^{-1}$\\ Mpc$^{-1}$, $\\Omega_M=0.3$, and $\\Omega_\\Lambda=0.7$. ",
        "conclusions": "Bars are an important signpost of galaxy evolution because once a galaxy disk is sufficiently massive, dynamically cold and rotationally supported it forms a bar.  Therefore the evolution of the bar fraction over time is an important indicator of the evolutionary history of disk galaxies and the assembly of the Hubble sequence.  Using a detailed analysis of 2,157 L$^*$ face-on, spiral galaxies from 0.0$<$z$<$0.84 in the COSMOS 2-square degree survey we have investigated the evolution of the bar fraction over the last 7 Gyr. We have undertaken an extensive and careful analysis of selection effects (k-correction, surface brightness dimming, inclination, spatial resolution, etc.)  which is detailed in the Appendix.  Our main results are as follows:  $\\bullet$ The bar fraction for L$^*$ galaxies drops from about 65\\% in the local Universe to about 20\\% at z=0.84. Over this redshift range the fraction of {\\em strong} bars (SB) drops from about 30\\% to under 10\\%.  Thus at a lookback time of 7 Gyr, when the Universe was half its present age, fundamental aspects of Hubble's `tuning fork' classification sequence had not yet fallen into place.  Only about one fifth of all spiral galaxies were ``mature'' enough (dynamically cold, massive and rotationally supported) to host galactic structures of the type we see today.  $\\bullet$ For the total f$_{bar}$ (SB+SAB), the change is far less dramatic between z=0.3 and z=0.0 indicating slow evolution in galactic structures in L$^*$ galaxies over the last 4 Gyr.  It is likely that there is significant evolution in the formation of bars in the sub-L$^*$ galaxies over this period.  $\\bullet$ One of the most significant findings in this study is the correlation between f$_{bar}$ and the galaxy mass, luminosity and color. We find that in the highest redshift bins f$_{bar}$ is higher in the more massive, luminous and redder systems. In fact, in the most massive systems, f$_{bar}$ is already as high at z=0.8 as the local value.  These systems thus had already arrived with their present Hubble types at a lookback time of 7 Gyr.  In the subsequent 3 Gyr, from z=0.84 to z=0.3, the lower mass, bluer systems evolved more slowly toward their present Hubble types.  Thus the signature of downsizing is intimately connected with dynamical maturity of disks and is present in the formation of galactic structure.  $\\bullet$ Finally, we find a slight preference for barred galaxies to be more bulge-dominated in the high redshift bin.  This correlation is consistent with the dynamical downsizing found for bars in general if bars and bulges both form earlier and more prominently in the most massive galaxies. The lack of a stronger correlation may be related to the variety of bulges: bars are also likely to be involved with the inflow that builds pseudo-bulges.  Given the strong correlation between bulge properties and black hole mass seen today, there may be a co-evolution of bars, bulges and black holes in some galaxies.  The exact details of these processes remain to be investigated."
    },
    "0710/0710.3528_arXiv.txt": {
        "abstract": "{The High Energy Stereoscopic System (H.E.S.S.), located in the Khomas Highlands of Namibia, is an array of four imaging atmospheric-Cherenkov telescopes designed to detect $\\gamma$-rays in the very high energy (VHE; $>$ 100 GeV) domain. Its high sensitivity and large field-of-view (5$^{\\mbox{\\tiny o}}$) make it an ideal instrument to perform a survey within the Galactic plane for new VHE sources. Previous observations in 2004/2005 resulted in numerous detections of VHE gamma-ray emitters in the region l = 330$^{\\mbox{\\tiny o}}$ - 30$^{\\mbox{\\tiny o}}$ Galactic longitude. Recently the survey was extended, covering the regions l = 280$^{\\mbox{\\tiny o}}$ - 330$^{\\mbox{\\tiny o}}$ and l = 30$^{\\mbox{\\tiny o}}$ - 60$^{\\mbox{\\tiny o}}$, leading to the discovery of several previously unknown sources with high statistical significance. The current status of the survey will be presented.}  \\begin{document} ",
        "introduction": "The majority of the newly discovered sources of very high energy (VHE; $>$ 100 GeV) $\\gamma$-rays are related to late phases of stellar evolution, either directly to massive stars or to the compact objects they form after their collapse. The possible associations include pulsar wind nebulae (PWN) of high spin-down luminosity pulsars such as G\\,18.0$-$0.7 \\cite{hess_j1825}, supernova remnants like RX\\,J1713.7$-$3946 \\cite{RXJ1713}, and open star clusters like Westerlund\\,2 \\cite{westerlund2}. As these objects cluster closely along the Galactic plane, a survey of this region is an effective approach to discover new sources and source classes of VHE $\\gamma$-ray emission. ",
        "conclusions": "The H.E.S.S. Galactic plane survey, which started in the year 2004, now reaches from $-$85$^{\\circ}$ longitude to 60$^{\\circ}$ longitude, and covers an approximately 6$^{\\circ}$ broad band around latitude b = 0$^{\\circ}$.  In total, more than 950\\,hours of data were taken in this region, including survey mode observations, re-observations of source candidates and dedicated observations of known or suspected $\\gamma$-ray sources. The first stage of the survey, covering the inner 60$^{\\circ}$ of the Galactic plane, has increased the number of known VHE $\\gamma$-sources within this region from three at the beginning of 2004 to seventeen. Further follow-up observations within this region and the extension of the survey along the Galactic plane resulted in the discovery of even more additional VHE $\\gamma$-ray emitters. Most of them were presented during this conference. Multi-wavelength follow-up observations and archival searches have already begun, and will be crucial for understanding the underlying processes at work in these astrophysical objects."
    },
    "0710/0710.1836_arXiv.txt": {
        "abstract": "We report on the first detection of maser emission in the $J$=11-10, $J$=14-13 and $J$=15-14 transitions of the $v$=0 vibrational state of SiS toward the C-rich star IRC+10216. These masers seem to be produced in the very inhomogeneous region between the star and the inner dust formation zone, placed at $\\simeq$5-7~R$_*$, with expansion velocities below 10~km~s$^{-1}$. We interpret the pumping mechanism as due to overlaps between $v$=1-0 ro-vibrational lines of SiS and mid-IR lines of C$_2$H$_2$, HCN and their $^{13}$C isotopologues. The large number of overlaps found suggests the existence of strong masers for high-$J$ $v$=0 and $v$=1 SiS transitions, located in the submillimeter range. In addition, it could be possible to find several rotational lines of the SiS isotopologues displaying maser emission. ",
        "introduction": "The detection of strong maser emission at the frequencies of pure rotational transitions of some molecules is a common phenomenon in circumstellar envelopes (CSE's) of evolved stars \\citep{elitzur_1992,gray_1999}. The maser is usually produced in a small region of the envelope and sometimes provides valuable information on the physical conditions of the emitting region.  Due to the different chemistry, masers are produced by different molecules in O- and C-rich stars. In O-rich stars, SiO exhibits strong maser emission in different rotational transitions within several vibrational states, from $v$=1 to 4 \\citep{cernicharo_1993,pardo_1998}. These masers are formed in a region of the CSE very close to the stellar surface and seem to be driven by NIR radiation \\citep{pardo_2004}. In C-rich stars, although SiO is present with similar abundances than in O-rich stars \\citep{schoier_2006}, no SiO maser has been detected. The explanation could be that SiO is formed at $\\simeq$3-5~R$_*$, where the angular dilution of the star is high and the density and temperature lower than in the regions where SiO masers are produced in O-rich stars \\citep{agundez_2006}. In C-rich stars only HCN shows strong maser emission in several pure rotational lines within vibrational states from $\\nu_2$=1 to 4 \\citep{lucas_1989,schilke_2003}. These masers must be formed in the innermost regions of the CSE. SiS has been previously found to show weak maser emission in the $J$=1-0 $v$=0 transition in IRC+10216 \\citep{henkel_1983}.  In this letter, we report on the first detection of maser emission from the $J$=11-10, 14-13 and 15-14 transitions in the $v$=0 vibrational state of SiS (hereafter M$_1$, M$_2$, and M$_3$) observed toward the C-rich star IRC+10216. We have also obtained observations of $v$=1 rotational lines which exhibit thermal emission. We propose that overlaps of $v$=1-0 ro-vibrational transitions of SiS with mid-IR lines of C$_2$H$_2$ and HCN could provide the pumping mechanism for these masers as well as higher-$J$ $v$=0 SiS masers in the submillimeter range. This discovery is interesting because this species could play in C-rich stars a role similar to that of SiO in O-rich stars: the energy level pattern of both molecules is similar and it is also formed close to the star, as chemical equilibrium and interferometric observations imply \\citep{bieging_1989}. ",
        "conclusions": "The $v$=1 rotational lines show single cusped profiles and their relative intensities indicate that the rotational levels are thermally populated. The linewidths correspond to velocities of 9-11~km~s$^{-1}$, lower than the terminal velocity $\\simeq$14.5~km~s$^{-1}$ \\citep{cernicharo_2000}, thus the emission arises from the innermost region of the CSE, between the photosphere and the inner dust formation zone, placed at $\\simeq$5~R$_*$ \\citep{keady_1993}. We can derive the SiS abundance in that region from the $v$=1 lines assuming an uniform sphere at a distance of 180 pc with a radius of 10 R$_*$, illuminated by the central star (R$_*$=5$\\times$10$^{13}$ cm, T$_*$=2300~K), with T$_\\textnormal{\\scriptsize{k}}$=1000~K and n$_{\\textnormal{\\scriptsize{H}}_2}$=1.6$\\times$10$^{9}$~cm$^{-3}$ (mean values for this region derived by \\citealt{fonfria_2006}) and an expansion velocity of 11~km~s$^{-1}$. We solve the statistical equilibrium for SiS considering 100 rotational levels and 3 vibrational states and apply the LVG radiative transfer formalism using the code developed by J. Cernicharo. With the assumed n$_{\\textnormal{\\scriptsize{H}}_2}$, the column density for H$_2$ is $\\simeq$7$\\times$10$^{23}$~cm$^{-2}$ and the derived one for SiS is $\\simeq$5.0$\\times$10$^{18}$~cm$^{-2}$. Hence, the SiS abundance is $\\simeq$7$\\times$10$^{-6}$. This result is compatible with a higher SiS abundance in the innermost CSE (from LTE chemistry models, 3$\\times$10$^{-5}$, \\citealt{agundez_2006}) and an abundance further away of 6.5$\\times$$10^{-7}$ according to observations of $v$=0 rotational lines over the outer CSE by \\citet{bieging_1989}.  Most $v$=0 rotational lines show a rounded or slightly double peaked profile with the blue part absorbed by cold SiS through the envelope. However, M$_1$, M$_2$ and M$_3$ show extra emission in the form of narrow peaks (FWHM=1-3~km s$^{-1}$). The velocities of these features, within the $-10$ to 10~km s$^{-1}$ range, indicate that SiS maser emission arises from different regions located between the star and the inner dust formation zone ($r$$\\simeq$5~R$_*$). The lines M$_2$ and M$_3$ have been previously observed by \\citet{sahai_1984} but no maser emission was noticed. This could be due either to the limited sensitivity of their observations or to a time variability of the SiS maser phenomenon. The bottom-center panel of Fig.~\\ref{fig:figure} shows in detail the line profiles of the observed masers. Up to ten maser features labelled ($a$,\\ldots,$j$) are identified. The most complex line profile is that of M$_3$. It is formed by 5 main features: $a$, $d$, $f$, $i$, $j$ with $v$=$-8.8$, $-3.7$, $0.88$, $6.3$ and $8.4$~km s$^{-1}$, respectively. The line profiles show that the strongest peaks are at negative velocities, having their red counterparts rather weak. This behavior was previously found by \\citet{henkel_1983} for the $v$=0 $J$=1-0 line, which mostly consists of a narrow peak (FWHM=0.3~km~s$^{-1}$) centered at $-13.5$~km~s$^{-1}$. The strong asymmetry of the lineshapes can be either due to blanking by the star of the redshifted maser or by amplification of the blueshifted emission by the foreground stellar environment.  The strongest observed maser, M$_3$, with T$_\\textnormal{\\scriptsize{MB}}$$\\simeq$60~K (F$_\\textnormal{\\scriptsize{obs}}$$\\simeq$300~Jy; with thermal and non-thermal emission), is weak compared to some SiO masers detected in O-rich stars \\citep*[e.g.][]{cernicharo_1993} or to HCN masers observed in IRC+10216 \\citep{lucas_1989,schilke_2003}. However, M$_1$, M$_2$, and M$_3$ are stronger than the SiS $J$=1-0 maser observed towards IRC+10216 by \\citet{henkel_1983}. The similarity between maser features in M$_2$ and M$_3$ indicates that they may arise from the same regions and produced by the same pumping mechanism; the maser in M$_1$ is probably formed in other regions. Hence, we suggest two possible geometries of the innermost CSE to explain the observed features:  ($i$) An onion-like innermost region, where each maser is produced in a shell. This hypothesis is supported by the symmetry of features $a$--$j$ and $d$--$i$. The peaks at extreme velocities, $a$ and $j$, would be produced just in front of and behind the star near the inner dust formation shell ($r$$\\simeq$5~R$_*$) with expansion velocities of $\\simeq$5-11~km~s$^{-1}$. The features $d$ and $i$ would be produced in a similar way but in an inner shell with a lower expansion velocity. Finally, the central peak, $f$, would be formed in a shell very close to the star, with the whole shell contributing to the maser emission. M$_1$ would be produced in a cap-shaped region in front of the star.  ($ii$) All the masers are formed in different positions of a clumpy shell. The different features in M$_2$ and M$_3$ would be produced in different regions of the shell: peaks $a$ and $j$ in front of and behind the star and the other peaks ($d$, $f$, and $i$) in different clumps, as occur with the only feature of M$_1$.  The classic pumping mechanism for the SiO $v$$>$0 masers observed in O-rich stars resides in the increase of the trapping lifetime (A/$\\tau$)$_{v \\rightarrow v-1}$ with $J$ for $v$$\\rightarrow$$v$-1 transitions, when they become optically thick \\citep{kwan_1974}. Such mechanism produces masers in adjacent rotational lines of the $v$ state, and explains the $v$=1 and 2 SiO masers \\citep{bujarrabal_1981,lockett_1992}. However, the masers observed in rotational transitions of $^{29}$SiO, $^{30}$SiO, and in $v$=3 and 4 of SiO do not show the latter behavior and have been interpreted as due to IR overlaps between ro-vibrational lines of SiO isotopologues \\citep{cernicharo_1991,gonzalezalfonso_1997}. For SiS, the absence of maser emission in $v$=1 rotational lines and the odd $v$=0 pattern also exclude the Kwan \\& Scoville pumping mechanism. This suggests that overlaps of $v$=1-0 ro-vibrational transitions of SiS with those of mode $\\nu_5$ of C$_2$H$_2$ and mode $\\nu_2$ of HCN, could provide the pumping mechanism. C$_2$H$_2$ and HCN are abundant in the inner CSE of IRC+10216 and dominate the 11-14 $\\mu$m spectrum \\citep{fonfria_2006}. Overlaps with these two species have been already proposed by \\citet{sahai_1984} to explain the different profiles of adjacent $J$ lines of SiS. However, the SiS frequencies used in that work were not as accurate ($\\sigma$$\\sim$10$^{-1}$~cm$^{-1}$) as those available today. We have calculated those frequencies from the Dunham coefficients determined by \\citet{sanz_2003}, for which the error of the band center is $<$$10^{-4}$~cm$^{-1}$ ($\\simeq$0.04~km~s$^{-1}$; the relative accuracy of P and R lines is much better). The frequencies of C$_2$H$_2$, H$^{13}$CCH, HCN, and H$^{13}$CN lines have been taken from the HITRAN Database 2004 \\citep{rothman_2005}, with an accuracy better than $10^{-3}$ cm$^{-1}$ ($\\simeq$0.4~km~s$^{-1}$) for C$_2$H$_2$ and H$^{13}$CCH, and $10^{-4}$~cm$^{-1}$ for HCN and H$^{13}$CN.  Table~\\ref{tab:frequencies_sis} shows the mid-IR line overlaps of SiS with C$_2$H$_2$, HCN, and their most abundant isotopologues. For the overlap search we selected coincidences within $|\\Delta v|$$<$10~km~s$^{-1}$. However, since the CSE is expanding, every region of the envelope is receding from the others. Hence, if the population of the SiS levels is affected by an overlap with a strong line of other species, the frequency of this overlapping transition must be higher than the SiS one. This would restrict the condition to positive $\\Delta v$. Nevertheless, due to the linewidth, lines at $\\Delta v$$<$0 can overlap the SiS lines. Hence, we have set the negative cutoff to one half of the typical linewidth in the innermost CSE ($\\simeq$5~km~s$^{-1}$; \\citealt{fonfria_2006}). Therefore, our search is restricted to $-2.5\\le\\Delta v\\le 10$~km~s$^{-1}$. All the ro-vibrational SiS lines commented hereafter refer to $v$=1$\\rightarrow$0 transitions and will be labelled with the usual spectroscopic nomenclature R, Q, P (see footnote of Table~\\ref{tab:frequencies_sis}).  \\begin{deluxetable}{c@{}c|c@{}c@{}c} \\tabletypesize{\\scriptsize} \\tablewidth{0pc} \\tablecolumns{5} \\tablecaption{Mid-infrared Line Overlaps of SiS with C$_2$H$_2$, HCN, and Their Most Abundant Isotopologues} \\tablehead{\\colhead{Line} & \\colhead{$\\nu$ (cm$^{-1}$)} & \\colhead{Mol.} & \\colhead{Transition} & \\colhead{$\\Delta v$(km/s)}} \\startdata \\multicolumn{5}{c}{\\textbf{Overlaps Involving SiS Levels of Observed Lines}}\\\\[2pt] R 9 & 750.3695 & HCN         &  $01^{1}0-00^{0}0$ R$_e\\left(12\\right)$        & $-6.6$  \\\\ P10 & 738.2889 & C$_2$H$_2$  &  $1^{-1}1^{1}-1^{-1}0^{0}$ R$_f\\left(3\\right)$ & \\phs{}1.2  \\\\ R13 & 752.6431 & C$_2$H$_2$  &  $0^{0}1^{1}-0^{0}0^{0}$ R$_e\\left(9\\right)$   & \\phs{}5.9  \\\\[2pt] P14 & 735.7324 & C$_2$H$_2$  &  $0^{0}1^{-1}-0^{0}0^{0}$ Q$_e\\left(38\\right)$ & $-6.1$  \\\\ P15 & 735.0861 & C$_2$H$_2$  &  $0^{0}1^{-1}-0^{0}0^{0}$ Q$_e\\left(36\\right)$ & $-9.9$  \\\\[2pt] \\multicolumn{5}{c}{\\textbf{Overlaps Involving SiS Levels of Unobserved Lines}}\\\\[2pt] P 2 & 743.2624 & C$_2$H$_2$  &  $0^{0}1^{1}-0^{0}0^{0}$ R$_e\\left(5\\right)$    & \\phs{}0.6 \\\\ P16 & 734.4368 & C$_2$H$_2$  &  $0^{0}1^{-1}-0^{0}0^{0}$ Q$_e\\left(34\\right)$  & \\phs{}1.1 \\\\ P20 & 731.8114 & C$_2$H$_2$  &  $0^{0}1^{-1}-0^{0}0^{0}$ Q$_e\\left(24\\right)$  & \\phs{}9.5 \\\\ P22 & 730.4815 & C$_2$H$_2$  &  $1^{-1}1^{1}-1^{1}0^{0}$ Q$_e\\left(18\\right)$  & $-2.1$ \\\\ R22 & 757.5819 & C$_2$H$_2$  &  $1^{1}1^{1}-1^{1}0^{0}$ R$_e\\left(10\\right)$   & \\phs{}0.1 \\\\ R22 & 757.5819 & H$^{13}$CN  &  $01^{1}0-00^{0}0$ R$_e\\left(17\\right)$         & \\phs{}5.2 \\\\ P23 & 729.8123 & H$^{13}$CCH &  $0^{0}1^{-1}-0^{0}0^{0}$ Q$_e\\left(19\\right)$  & \\phs{}8.7 \\\\ P24 & 729.1402 & C$_2$H$_2$  &  $0^{0}1^{-1}-0^{0}0^{0}$ Q$_e\\left(1\\right)$   & \\phs{}9.4 \\\\ P25 & 728.4654 & H$^{13}$CCH &  $0^{0}1^{-1}-0^{0}0^{0}$ Q$_e\\left(7\\right)$   & $-0.0$ \\\\ R25 & 759.1733 & HCN         &  $01^{1}0-00^{0}0$ R$_e\\left(15\\right)$         & \\phs{}3.6 \\\\ R26 & 759.6977 & C$_2$H$_2$  &  $0^{0}1^{1}-0^{0}0^{0}$ R$_e\\left(12\\right)$   & $-1.0$ \\\\ R33 & 763.2816 & H$^{13}$CN  &  $01^{1}0-00^{0}0$ R$_e\\left(19\\right)$         & \\phs{}6.1 \\\\ P38 & 719.4363 & C$_2$H$_2$  &  $1^{1}1^{1}-1^{1}0^{0}$ P$_e\\left(5\\right)$    & \\phs{}4.5 \\\\ R40 & 766.7130 & C$_2$H$_2$  &  $0^{0}1^{1}-0^{0}0^{0}$ R$_e\\left(15\\right)$   & \\phs{}4.3 \\\\ R42 & 767.6652 & C$_2$H$_2$  &  $1^{1}1^{-1}-1^{1}0^{0}$ R$_e\\left(20\\right)$  & \\phs{}8.7 \\\\ R45 & 769.0697 & C$_2$H$_2$  &  $0^{0}1^{1}-0^{0}0^{0}$ R$_e\\left(16\\right)$   & $-1.8$ \\\\ P48 & 712.1720 & C$_2$H$_2$  &  $1^{-1}1^{-1}-1^{-1}0^{0}$ P$_f\\left(8\\right)$ & \\phs{}4.7 \\enddata \\tablecomments{Overlaps of SiS with C$_2$H$_2$, H$^{13}$CCH, HCN and H$^{13}$CN found in the mid-IR range $\\left[690,780\\right]$~cm$^{-1}$ with $-2.5\\le\\Delta v\\le 10$ km~s$^{-1}$, where $\\Delta v/c=[\\nu (\\textnormal{X})-\\nu (\\textnormal{SiS})]/ \\nu (\\textnormal{SiS})$ and $J\\le 50$, (corresponding SiS $v$=0 rotational frequencies below 900~GHz). The lines and frequencies at the left correspond to the vibrational transitions $v$=1$\\rightarrow$0 of SiS. The errors on the velocities are $< 0.5$~km~s$^{-1}$. The notation for the C$_2$H$_2$ and H$^{13}$CCH vibrational states involved in the ro-vibrational transitions is $\\nu_4^{\\ell_4}\\nu_5^{\\ell_5}$, whereas for HCN and H$^{13}$CN is $\\nu_1\\nu_2^\\ell\\nu_3$. The parity of the lower level is even ($e$) and odd ($f$). The transitions are labelled as R, P, and Q for $J_{up}-J_{low}=+1$, $-1$, 0.} \\label{tab:frequencies_sis} \\end{deluxetable}  In order to qualitatively interpret the effects of IR overlaps on maser emission, we have used the same LVG radiative transfer code modified to account for overlaps, changing the intensity at the overlapping frequency and the escape probability for photons from the overlapped lines. Thus, for some SiS lines, the excitation temperature becomes negative (maser activation) and the brightness temperature is considerably enhanced. M$_2$ and M$_3$ are naturally explained by the overlap of the R$\\left(13\\right)$ line of SiS with the strong C$_2$H$_2$ line $0^01^1-0^00^0$R$_e\\left(9\\right)$ (Table~\\ref{tab:frequencies_sis}). SiS molecules are easily excited from $v$=0 $J$=13 to $v$=1 $J$=14 through R$\\left(13\\right)$ and decay to the $v$=0 $J$=15 via P$\\left(15\\right)$, creating a population inversion between the $v$=0 $J$=13 and 15, and producing maser emission in M$_2$ and M$_3$. M$_1$ may be produced by the overlap of the P$\\left(10\\right)$ SiS line with the strong C$_2$H$_2$ line $1^{-1}1^1-1^{-1}0^0$R$_f\\left(3\\right)$. This overlap can pump SiS molecules from $v$=0 $J$=10 to $v$=1 $J$=9 through P$\\left(10\\right)$, depopulates the $v$=0 $J$=10 level and produces the inversion between $v$=0 $J$=10 and 11.  We have also looked for overlaps of $v$=1-0 higher-$J$ SiS lines with C$_2$H$_2$ and HCN transitions to try to predict SiS masers at submillimeter wavelengths. Some of them are shown in the second block of Table~\\ref{tab:frequencies_sis}. They suggest, for example, that a maser could be found in rotational transitions involving the $v$=0 $J$=23 level (likely $J$=24-23 and maybe 23-22), or the $v$=0 $J$=25 and 26 states (perhaps in the $v$=0 $J$=27-26 and maybe 26-25). These overlaps could also produce masers in $v$=1 rotational transitions.  There are many overlaps of other SiS isotopologues with lines of C$_2$H$_2$, HCN and $^{28}$Si$^{32}$S. With the adopted criteria for the overlap search (see footnote of Table~\\ref{tab:frequencies_sis}), we have found 91, 93, 76, and 94 coincidences for $^{29}$SiS, $^{30}$SiS, Si$^{34}$S, and Si$^{33}$S, respectively. Consequently, although these species are less abundant than $^{28}$Si$^{32}$S, the population of some levels could be inverted producing maser emission.  This study represents the discovery of three new SiS masers and should be complemented with future observations of higher-$J$ $v$=0 and 1 rotational transitions. Furthermore, a detailed multi-molecule non-local radiative transfer model would help to understand the dust formation region and the role of SiS in C-rich evolved stars."
    },
    "0710/0710.0600_arXiv.txt": {
        "abstract": "% We describe the status of a project whose main goal is to detect variability along the extreme horizontal branch of the globular cluster NGC~6752. Based on Magellan 6.5m data, preliminary light curves are presented for some candidate variables. By combining our time-series data, we also produce a deep CMD of unprecedented quality for the cluster which reveals a remarkable lack of main sequence binaries, possibly pointing to a low primordial binary fraction. ",
        "introduction": "Among field B-type subdwarf (sdB) stars, three types of non-radial pulsators have so far been detected, namely:  \\begin{itemize} \\item EC 14026 (sdBV) stars: these are $p$-mode pulsators whose temperatures fall in the range between 29,000 and 36,000~K. Their periods are typically found in the range 100-200~sec, and their amplitudes cover the range from 0.4 to 25\\%.  \\item PG1716+426 (``Betsy'') stars: these are $g$-mode pulsators, with temperatures in the range between 25,000 and 30,000~K. Their periods are much longer, typically falling in the range between 2000 and 9000~sec, and their amplitudes are smaller than 0.5\\%.  \\item Hybrid stars: these are stars that present simultaneous $p$- and $g$- mode oscillations. At present, only two examples have been reported in the literature \\citep{abea05,roea05,ssea06}. \\end{itemize}   Performing asteroseismology on these stars holds the promise to unveil their innermost secrets, thus providing an exciting new route toward the solution of the so-called ``second-parameter problem'' \\citep[][and references therein]{mc05}. The great promise of the technique notwithstanding, such variables have never been detected in previous searches in globular clusters \\citep[e.g.,][]{mrea06}.  Accordingly, the main purpose of the present study is to perform a new search for this type of variables in the relatively nearby southern globular cluster NGC~6752, which contains a very long blue HB ``tail,'' and thus many potential candidates for the three aforementioned variability types. ",
        "conclusions": "Our analysis of time-series observations collected with the Magellan 6.5m MagIC camera already reveals a few intriguing variable candidates, including at least one on the hot extension of the HB, in two small fields located away from the cluster center. We have recently acquired extensive time-series observations using IMACS over a much larger field. These data will help confirm the nature of the suspected variables, and will presumably reveal a host of other faint/low-amplitude variables in NGC~6752. In addition, the new data will allow us to map the MS binary fraction in the cluster as a function of radius, which will reveal, for the first time, how in detail the binary fraction decreases as a function of radius in a globular star cluster."
    },
    "0710/0710.2433_arXiv.txt": {
        "abstract": "We investigate the wave effect in the gravitational lensing by a black hole with very tiny mass less than $10^{-19} M_{\\rm sun}$ (solar mass), which is called attolensing, motivated by a recent report that the lensing signature might be a possible probe of a modified gravity theory in the braneworld scenario. We focus on the finite source size effect and the effect of the relative motion of the source to the lens, which are influential to the wave effect in the attolensing. Astrophysical condition that the lensed interference signature can be a probe of the modified gravity theory is demonstrated. The interference signature in the microlensing system is also discussed. ",
        "introduction": "Many physicists have drawn attention to extra dimensional physics for several years due to recent development in testing Randall-Sundrum type II (RS-II) scenario \\cite{RSII}. In this scenario they considered a four dimensional positive-tension brane embedded in five dimensional AdS bulk which allows us to reconsider our understanding about the history of our universe in the early stages. Some investigations have been carried out to modify the existence of primordial black holes (PBH)s in this RS-II scenario that the life time of five dimensional PBHs against the Hawking radiation becomes longer compared with the standard PBH in four dimensions \\cite{GCL}. This is because the five dimensional feature becomes significant in the black hole with very tiny mass. The ratio of the life time against the Hawking radiation of such five dimensional PBH to that of the four dimensional PBH may be estimated, $lM_4/l_4M$, where $l$ is the AdS radius of the braneworld model, $l_4$ and $M_4$ are the four dimensional Planck length and mass, respectively, and $M$ is the black hole mass \\cite{GCL}.  Since the braneworld PBHs can live longer, then it is possible that such black holes still exist and spread out in our universe. Therefore, it is natural to consider the possibility of the gravitational lensing phenomenon by such the black hole. More recently, the authors \\cite{KPIII} investigated the wave effect in gravitational lensing by the black hole with the very tiny mass smaller than $10^{-19} M_{\\rm sun}$, where $M_{\\rm sun}$ is the solar mass, which is called attolensing. They showed that the interference signature in the energy spectrum in gamma ray burst due to the attolensing might be a possible probe of the modified gravity theory. In the standard general relativity, PBH with the mass smaller $10^{-18} M_{\\rm sun}$ will be evaporated through the Hawking radiation within the cosmic age. Then, the detection of such the interference signature would be a probe of extra dimension of our universe.  In the present paper, we consider the two effects in the lensing phenomenon that are influential for measurement of the interference signature in the energy spectrum in the attolensing: One is the finite source size effect. The other is the effect of the relative motion of the black hole (lens object) to the source. We demonstrate the condition that these effects become influential. This is one of the astrophysical conditions that the attolensing can be a probe of the modified gravity theory. Throughout this paper, $ H_0 = 72$ km s$^{-1}$ Mpc$^{-1}$ is the Hubble parameter, and we use the unit in which the light velocity equals 1. ",
        "conclusions": "We investigated the condition that the finite source size effect and the relative motion of the source becomes substantial in the wave effect of the gravitational lensing. The condition is expressed by the formula (\\ref{condtwo}), whose physical meaning is that the difference of the phase between the two light paths from $y_{\\rm max}$ and $y_{\\rm min}$ is larger than $2\\pi$. We have shown that the finite source size effect is important in the attolensing by the black hole at the cosmological distance.  Also the relative motion of the source to the lens can be influential. For the attolensing by the black hole at the galactic distance, the constraint is relaxed by the factor $\\sqrt{D_S/D_L}\\sim 10^{3}$. However, we should also note that the detection of the attolensing is limited in practice \\cite{KPIII}, even when the finite source size is not taken into account.   In general, the signature of the interference becomes remarkable when $w\\sim 1$, i.e., \\begin{eqnarray} w \\sim 0.3\\times (1+z_L)\\left({h\\nu\\over 100{\\rm MeV}}\\right) \\left({M\\over 10^{-19}M_{\\rm sun}}\\right)\\sim 1. \\label{condfirst} \\end{eqnarray} In the case $w\\simlt 1$, the amplification due to the lensing becomes negligible, because the wavelength of the light becomes larger than the size of deflector. Combining this and the condition (\\ref{condtwo}), we may conclude that the condition for the possible observation of the interference signature in the gravitational lensing needs \\begin{eqnarray} 1\\simlt w \\simlt {\\pi \\over |y_{\\rm min}-y_{\\rm max}|}. \\label{finalcond} \\end{eqnarray} Therefore, this means $|y_{\\rm min}-y_{\\rm max}|\\simlt \\pi$, the angular size of the source must be less than the Einstein radius, is always necessary for a possible observation of the interference signature. It will be useful to demonstrate the region satisfying (\\ref{finalcond}) clearly. In Figure 8, the dashed line is $w=1$, and the solid line is $w|y_{\\rm min}-y_{\\rm max}| =\\pi$, where we fixed $\\hat R=10^3$ km and $D_{LS}D_S/D_L=H_0^{-1}$. The shaded region satisfies the condition (\\ref{finalcond}).  It might also be interesting to consider whether other physical system satisfy the condition (\\ref{finalcond}) or not. Figure 9 shows a region satisfying the condition with the parameter associated with the microlensing. Similar to Figure 8, the shaded region satisfies the condition. Here the dashed line in Figure 9 is $w=1$, and the solid line is $w|y_{\\rm min}-y_{\\rm max}| =\\pi$, where we fixed $\\hat R=7\\times 10^5$ km (solar radius) and $D_{LS}D_S/D_L=30$ kpc. The point of the intersection of these two lines is \\begin{eqnarray} && \\left({\\nu\\over {\\rm GHz}}\\right)=8.9\\times 10^2 \\left({\\hat R\\over 7\\times 10^5{\\rm km}} \\right)^{-2} \\left({30{\\rm kpc}\\over D_{LS}D_S/D_L}\\right)^{-1} \\\\ && \\left({M\\over M_{\\rm sun}}\\right)=0.9\\times 10^{-8} \\left({\\hat R\\over 7\\times 10^5{\\rm km}} \\right)^{2} \\left({30{\\rm kpc}\\over D_{LS}D_S/D_L}\\right). \\end{eqnarray} Thus the microlensing can be a possible system that satisfies the condition (\\ref{finalcond}). Especially, for the microlensing by an Earth like planet $M\\sim 10^{-5} M_{\\rm sun}$, the relevant range of the frequency is $1$ GHz $\\sim 100$ GHz. Then, the measurement of the microlensing event through the frequency might be relevant to the interference signature. However, it will be very difficult to detect the signal because the stars at $1\\sim 100$ GHz frequency band at the galactic distance is very dark in general.   \\vspace{0.3cm}"
    },
    "0710/0710.5355_arXiv.txt": {
        "abstract": "{ The ANTARES Collaboration is building a high-energy neutrino telescope at 2500 m depth in the Mediterranean Sea. The experiment aims to search for high-energy cosmic neutrinos through the detection of Cerenkov light induced by muons and showers resulting from neutrino interactions with the surrounding medium. The detector will consist of a three-dimensional array of 900 optical modules housing photomultipliers. It will be composed of 12 strings, 5 of them being already in operation since January 2007. The muon track is reconstructed from the arrival time and the charge of the signals obtained from the photomultipliers, whose positions are known by means of an acoustic positioning system. The reconstruction strategies include several steps among which there are: optical background filtering, algorithms for first estimations of the track parameters, and a final fit aiming to reach an angular resolution better than 0.3 degree above 10 TeV in the full detector. Different reconstruction strategies will be presented and their application to the present real data analysis will be reviewed.}  \\begin{document} ",
        "introduction": "When an ultra-relativistic particle ($\\beta \\simeq 1$) moves in a medium, Cerenkov light is emitted at an angle depending on the refraction index $n$ of the medium. In case of sea water $n$ is about $1.34$, and thus the Cerenkov emission angle becomes \\begin{equation} cos (\\theta_C) = \\frac{1}{n} \\simeq cos(42^\\circ). \\end{equation} The emitted Cerenkov photons are detected by an array of photomultipliers installed at the bottom of the sea. Since January 2007 the ANTARES detector is a full three-dimensional array of photomultipliers consisting of 5 strings detecting muons at a depth of 2475 meters below the sea level. ANTARES 10-inch photomultipliers are housed in pressure resistant glass spheres called optical modules (OM). The 5-line detector data taking has allowed the Collaboration to tune the various processes needed to collect data, and muon events have been detected on top of the optical background (60-100 kHz most of the time). The calibration system has been proven to be efficient, and the knowledge of the charge, of the arrival time of the signals and the positioning of the OMs has allowed the first reconstruction codes to be tested in realistic conditions. The status of the experiment is discussed in \\cite{antoine}. Preliminary data studies show that a flux of atmospheric downward-going muons is triggering the detector at a rate of about 1 Hz and that upgoing atmospheric neutrino candidates have been identified. Among the various muon reconstruction algorithms under study, two different alternative methods are presented in this paper. Other well developed reconstruction methods are described in \\cite{aart}, \\cite{carmona}, and the implementation of the {\\sl{simulated annealing}} algorithm \\cite{simann} is under study. A discussion on event reconstruction techniques for Cerenkov neutrino telescopes can also be found in \\cite{amanda}. ",
        "conclusions": ""
    },
    "0710/0710.5163_arXiv.txt": {
        "abstract": "We report the discovery of 11 new cataclysmic variable (CV) candidates by the Isaac Newton Telescope (INT) Photometric H$\\alpha$ Survey of the northern Galactic plane (IPHAS). Three of the systems have been the subject of further follow-up observations.  For the CV candidates IPHAS J013031.90+622132.4 and IPHAS J051814.34+294113.2, time-resolved optical spectroscopy has been obtained and radial-velocity measurements of the H$\\alpha$ emission-line have been used to estimate their orbital periods.  A third CV candidate (IPHAS J062746.41+ 014811.3) was observed photometrically and found to be eclipsing. All three systems have orbital periods above the CV period-gap of 2--3\\,h. We also highlight one other system, IPHAS J025827.88+635234.9, whose spectrum distinguishes it as a likely high luminosity object with unusual C and N abundances. ",
        "introduction": "\\label{intro} Cataclysmic variables (CVs) are semi-detached interacting binary systems containing a white dwarf (WD) primary and a late-type main-sequence secondary. The secondary fills its Roche lobe, and matter is transferred to the primary through the inner Lagrangian point. The mass-transfer process in CVs means that line emission is a common observational feature of the majority of CVs, especially in the Balmer series. Observed lines may originate from optically thin or irradiated parts of the accretion disc (if present, \\citealt{1980ApJ...235..939W}). Line emission may also be produced in the accretion stream and on the irradiated face of the secondary star. \\citet{1995cvs..book.....W} provides a comprehensive review of CVs.  Population synthesis models suggest that the intrinsically faintest low accretion rate systems should dominate the Galactic CV population and should be found predominantly at short orbital periods, that is, below the well-known CV ``period gap'' between 2 hrs and 3 hrs \\citep{1993A&A...271..149K, 1997MNRAS.287..929H}. This dominant population of faint, short-period CVs has proven quite elusive. This is partly because most known CVs have been discovered with techniques that are biased against detecting CVs with low mass accretion rates. For example, all flux-limited surveys with relatively bright limiting magnitudes and also variability searches (e.g. for dwarf nova outbursts) are intrinsically biased against the detection of faint, short-period CVs with potentially long inter-outburst recurrence times (such as WZ Sge). For a closer look at the period distributions of CVs found with different discovery methods, see \\citet{2005ASPC..330....3G}. Recently \\citet{2007MNRAS.374.1495P} have shown that the relative dearth of short-period CVs is not just due to selection effects, implying a serious flaw in our understanding of CV evolution (see also \\citealt{2006A&A...455..659A}).  Nevertheless, selection effects must be at least partly responsible for the lack of known faint short-period CVs, and so many such systems still remain to be found.       Given the ubiquity of line emission amongst CVs, searches for objects displaying H$\\alpha$ emission offer a powerful way to find new CVs. Examples of such CV searches in the Galactic field include those of \\citet{2002ASPC..261..190G}, \\citet{2005A&A...443..995A}, \\citet{2002A&A...395L..47H}, and the Chandra Multiwavelength Plane (ChaMPlane) survey (\\citealt{2005ApJ...635..920G}, \\citealt{2005ApJS..161..429Z} and \\citealt{2006ApJS..163..160R}). Similar searches have also been performed in globular clusters have been performed by \\citet{1995ApJ...439..695C}, \\citet{1995ApJ...455L..47G}, \\citet{1996ApJ...473L..31B} and \\citet{2000ApJ...532..461C}. The Isaac Newton Telescope (INT) Photometric H$\\alpha$ Survey of the northern Galactic plane (IPHAS) is currently surveying the Milky Way in broad-band Sloan $r^\\prime$ and $i^\\prime$ and narrow-band H$\\alpha$ and provides an excellent data base for a detailed CV search at low Galactic latitudes.  The survey goes to a depth of $r^\\prime\\simeq20$\\,mag and covers the latitude range $-5^o < b < +5^o$. A detailed introduction to the survey, including the transmission profiles of the filters it uses, is given by \\citet{2005MNRAS.362..753D}. What makes IPHAS particularly promising for finding CVs is that, empirically, the intrinsically  faintest, low mass transfer rate ($\\dot{M}$) systems tend to have the largest Balmer line  equivalent widths (EWs; \\citealt{1984ApJS...54..443P}; \\citealt{2006MNRAS.369..581W}). Unless the inverse relationship between Balmer EW and orbital period breaks down for the faintest, short-period systems, emission-line surveys should be very good at finding low-$\\dot{M}$ CVs. IPHAS should then be an excellent way of detecting a large population of faint, short-period CVs.     In this paper we announce the discovery of eleven new CV candidates from the IPHAS survey data. We present spectroscopic follow-up observations of two of these CVs candidates and determine their orbital periods. We also present the results of time-series photometry of another CV candidate which establishes it as a long period, eclipsing, system. ",
        "conclusions": "\\label{conc} Spectroscopy of several objects from the IPHAS H$\\alpha$ excess catalogue has led to the discovery of 11 new CV candidates by virtue of their strong H$\\alpha$ emission. The identification spectra of the CVs candidates include cases where the secondary star is particularly prominent (IPHAS J1856), and examples of double-peaked emission lines (for example IPHAS J0518). One particular candidate (IPHAS J0258) has similarities to the extreme objects V Sge and QU Car.  Time-resolved optical spectroscopy of two of the new CVs (IPHAS J0130 and IPHAS J0518) has been obtained using CAFOS on the 2.2\\,m telescope at the Calar Alto Observatory. Radial-velocity measurements of the H$\\alpha$ emission-line were used to determine their orbital periods.  Periodograms and a Monte Carlo analysis were used to estimate the orbital period of $P_{orb}=0.130061 \\pm 0.000001$\\,d for IPHAS J0130.  The periodogram obtained from the radial-velocity data of IPHAS J0518 reveals four peaks which could be the orbital period of the system: 0.2383, 0.2203, 0.2595 and 0.2049\\,d, each with an uncertainty of 0.0007\\,d. Time-series photometry of IPHAS J0627 using the SAAO 1.9\\,m telescope found the system to be a long-period eclipser with four possible orbital periods: $1.020 \\pm 0.002$\\,d, $0.5101 \\pm 0.0008$\\,d, $0.3401 \\pm 0.0006$\\,d and $0.2551 \\pm 0.0004$\\,d. Further observations are necessary to obtain accurate values of the orbital period of all three systems, and to find accurate binary parameters."
    },
    "0710/0710.4954_arXiv.txt": {
        "abstract": "We study the signal for the detection of quasi-stable supersymmetric particle produced in interactions of cosmogenic neutrinos.  We consider energy loss of high energy staus due to photonuclear and weak interactions.  We show that there are optimal nadir angles for which the stau signal is a factor of several hundred larger than muons.  We discuss how one could potentially eliminate muon background by considering the energy loss of muons in the detector. We also show results for the showers produced by weak interactions of staus that reach the detector. ",
        "introduction": "Ultrahigh energy cosmic neutrinos could potentially probe physics beyond the Standard Model \\cite{ringwald}.  Interactions of UHE neutrinos ($E_\\nu \\geq 10^{17}$eV) with nucleons probe center of mass energies above 14 TeV.  Some fraction of these neutrinos may produce supersymmetric particles or some other exotic particles. These processes are suppressed relative to standard model processes, however, in some models interesting signals may arise from supersymmetric particles with long lifetimes.  In most SUSY scenarios, particles produced in high energy collisions decay immediately into the lightest one and are thus hard to detect. However in some low scale supersymmetric models in which gravitino is the lightest supersymmetric particle (LSP) and R-parity is conserved, the next-to-lightest particle (NLSP) is the charged superpartner of the right-handed tau, the stau \\cite{susy_models}. Due to its weak coupling to the gravitino, the stau is a long-lived particle in these models. For the supersymmetric breaking scale, $\\sqrt F >5\\times10^6$ GeV, the long-lived stau could travel distances of the order of $10^4$ km before decaying into the gravitino.  The distance that staus travel before decaying depends on the gravitino mass (or equivalently on the supersymmetry breaking scale) and the stau mass. Limits on the stau mass of about 100 GeV come from its non-observation in accelerator experiments \\cite{stmass1,stmass2,stmass3,stmass4}. Recently it was proposed that staus produced in high energy neutrino interactions, where neutrinos originate in astrophysical sources, might be detectable in neutrino telescopes \\cite{ABC,Ahlers}.  The cross section for the production of staus in neutrino-nucleon scattering \\cite{ABC} is several orders of magnitude smaller than the neutrino charged-current or neutral-current cross section \\cite{gqrs}.   However, once produced, the long-lived staus have the potential to travel through the earth  without decaying and thus open up a possibility to be detected in neutrino telescopes.  The long range of staus could potentially compensate for the suppression in the production cross section by increasing the effective detector volume and therefore enhancing the signal.  The detection of staus depends on the stau lifetime and range, so it is important to determine the energy loss as it traverses the earth. The details of the range depend in part on the supersymmetry breaking and how the quasi-stable stau particle is comprised of the SUSY partners of the right-handed and left-handed taus. The electromagnetic energy loss has been shown to have the largest contribution from photonuclear interactions for stau energies between $10^6$-$10^{12}$ GeV, resulting in a range of $10^4$ km.w.e. for masses of the order of a few hundred GeV \\cite{RSS}. Weak interactions may come into play as well. The stau range has been shown to be sensitive to the mixing angle of right-handed and left-handed staus.  When the mixing is maximal, weak interactions act to suppress the range at energies above $\\sim10^9$ GeV \\cite{RSU1}, however, their weak interactions have the potential to produce signals in neutrino detectors such as the Antarctic Impulse Transient Array (ANITA) \\cite{anita} and the Antarctic Ross Iceshelf Antenna Neutrino Array (ARIANNA) \\cite{arianna}.  The high energies required for stau production lead us to focus on the production of staus in interactions of cosmogenic neutrinos as they traverse the Earth and/or in the detector. These neutrinos originate from cosmic ray protons interacting with the cosmic microwave background, $$ p\\gamma(3{\\rm K}) \\rightarrow \\Delta \\rightarrow N\\pi$$ followed by charged pion, muon and neutron decays. This flux is guaranteed as cosmic ray fluxes are measured as well as the 3K microwave background.  We use a conservative cosmogenic neutrino flux evaluated by Engel, Seckel and Stanev (ESS) in Ref. \\cite{ess}. They evaluate the neutrino flux associated with the measured cosmic ray flux by tracing back cosmic ray propagation through the background radiation. Depending on the cosmological evolution assumed, the overall normalization of this flux has an uncertainty of about a factor of four. In addition, neutrinos could be produced at the sources of the high energy cosmic rays and those are not included in the evaluation of ESS neutrino flux. Thus, the ESS neutrino flux is a conservative estimate of the cosmogenic flux.  The cosmogenic neutrino flux, when neutrino flavor oscillations are not included, peaks at high energies, around $10^8$ GeV, and thus it is in the energy range where ANITA and ARIANNA have very good sensitivity \\cite{anita,arianna}. The neutrino flavor ratio for cosmogenic neutrinos deviates from the common 2:1 ratio, due to the neutron decay contribution to electron neutrinos \\cite{ess}. This implies that one needs to consider three flavor oscillations when considering the cosomogenic neutrino flux that arrives at the Earth \\cite{us_JMRS}.  We consider cosmogenic neutrinos, their propagation, stau production, and subsequent energy loss as it traverses the earth, for a region of parameter space where the staus do not decay over the distances required. We compare the resulting stau flux when there is no mixing between right-handed and left-handed stau and when there is maximal mixing. We consider muon-like signals (charged tracks) produced by staus and its associated background. We discuss the potential for eliminating the background by measuring the energy loss, which requires large volume detectors. Finally we discuss the showers produced in the ice due to stau interactions and its background from neutrino-induced showers. ",
        "conclusions": "We have studied signals of staus produced in interactions of cosmogenic neutrinos.  We have considered two types of signals, muon-like charged tracks and showers.   We have focused on low scale supersymmetric models that have stau as NLSP, which decays into the lightest SUSY particle, the gravitino. For a sufficiently large scale of supersymmetry breaking, the stau has a very long lifetime.  Our focus has been on cosmogenic neutrino fluxes and their associated stau production in the Earth. Energy losses, both through electromagnetic and weak interactions, are important in evaluating stau signals. The energy loss of staus, however, is relatively small in comparison with muons.  Thus, for some nadir angles, the stau flux is much larger than the muon flux produced in neutrino charged-current interactions.   The enhancement of the stau flux is larger from an input cosmogenic neutrino flux than for the Waxman-Bahcall neutrino flux \\cite{WB}. This is because the cosmogenic neutrino flux is peaked at energies of about $10^8$ GeV, while the WB flux is characterized by a steep power law with index of two. The large ratio of staus to muons from cosmogenic neutrinos is encouraging for experimental detection, but in order to see this signal one needs to be able to distinguish between staus and muons. Using the average energy loss per unit distance is not a good way to distinguish staus and muons, since the scaling of the energy loss parameter $\\beta$ has the effect of making a high-energy stau look like a lower energy muon.  We have proposed a way to distinguish between stau and muon tracks by measuring the energy loss of muons via their interactions in the ice, and to use this method to reduce the background.  We also considered showers produced by staus interacting in the ice via charged-current interactions.  The backgrounds for this signal are showers induced directly by neutrinos that reach the detector and interact inside the detector via charged-current or neutral-current interactions.  The only way that staus would produce showers in the ice is if there is a weak mixing, but this process also contributes to reducing stau range.  These effects combine with the small stau production probability to give fluxes of attenuated staus that are several orders of magnitude less than attenuated neutrino fluxes. This small stau to neutrino ratio translates directly to shower rates.  In addition to weak interactions, another possibility for shower production would be stau decays in the detector. For the parameter space considered here, with long-lived staus, the signal from decays is suppressed relative to their weak interactions.  In conclusion, stau signals at high energies are best identified by muon-like tracks. The most important detection issue is distinguishing between staus and muons, which may be possible by looking at the incremental electromagnetic energy loss as the charged particle moves through the detection volume. With very large volumes, there is a potential for detection of staus with future large neutrino telescopes."
    },
    "0710/0710.3868_arXiv.txt": {
        "abstract": "Two dimensional hydrodynamical disks are nonlinearly unstable to the formation of vortices. Once formed, these vortices essentially survive forever. What happens in three dimensions? We show with incompressible shearing box simulations that  in 3D a vortex in a short box forms and survives just as in 2D.  But a vortex in a tall box is unstable and is destroyed. In our simulation, the unstable vortex decays into a transient turbulent-like state that transports angular momentum outward at a nearly constant rate for hundreds of orbital times. The 3D  instability that destroys vortices is a generalization of  the 2D instability that forms them. We derive  the conditions for these nonlinear instabilities to act by calculating the coupling between linear modes, and thereby derive the criterion for a vortex to survive in 3D as it does in 2D: {\\it the azimuthal extent of the vortex must be larger than the scale height  of the accretion disk}. When this criterion is violated, the vortex is unstable and decays. Because vortices are longer in azimuthal than in radial extent by a factor that is inversely proportional to their excess vorticity, a vortex with given radial extent will only survive in a 3D disk if it is sufficiently weak.  This counterintuitive result explains why previous 3D simulations always yielded decaying vortices: their vortices were too strong. Weak vortices behave two-dimensionally even if their width is much less than their height because they are stabilized by rotation, and behave as Taylor-Proudman columns. We conclude that in protoplanetary disks  weak vortices can trap dust and serve as the nurseries of planet formation.  Decaying strong vortices might be responsible for the outwards transport of angular momentum that is required to make accretion disks accrete. ",
        "introduction": "Matter accretes onto a wide variety of objects, such as young stars, black holes, and white dwarfs, through accretion disks.  In highly ionized disks  magnetic fields are important, and they trigger turbulence via the magnetorotational instability \\citep{BH98}.  However, many disks, such as those around young stars or dwarf novae, are nearly neutral  \\citep[e.g.,][]{SMUN00,GM98}. In these disks, the fluid motions are well described by hydrodynamics.  Numerical simulations of hydrodynamical disks in two-dimensions---in the plane of the disk---often produce long-lived vortices \\citep{GL99,UR04,JG05}. If vortices really exist in accretion disks, they can have important consequences.  First and foremost, they might generate turbulence. Since turbulence naturally transports angular momentum outwards\\footnote{ Energy conservation implies that turbulence transports angular momentum outwards; see \\S \\ref{sec:pseudo}. Nonetheless, if an external energy source  (e.g., the radiative energy from the central star) drives the turbulence, then angular momentum could in principle be transported inwards. }, as is required for mass to fall inwards, it might be vortices that cause accretion disks to accrete. Second, in disks around young stars, long-lived vortices can trap solid particles and initiate the formation of planets \\citep{BS95}.   Why do vortices naturally form in 2D simulations? Hydrodynamical disks are stable to linear perturbations. However, they are nonlinearly unstable, despite some claims to the contrary in the astrophysical literature. In two dimensions, the incompressible hydrodynamical equations of a disk are equivalent to those of a non-rotating linear shear flow \\citep[e.g.,][hereafter L07]{L07}. And it has long been known that such flows are nonlinearly unstable (\\mycite{Gill65}; \\mycite{LK88}; L07).  This nonlinear instability is just a special case of the Kelvin-Helmholtz instability. Consider a linear shear flow extending throughout the $x$-$y$ plane with velocity profile $\\bld{v}=-qx\\bld{\\hat{y}}$, where $q>0$ is the constant shear rate, so that $-q$ is the flow's vorticity. (In the equivalent accretion disk, the local angular speed is $\\Omega=2q/3$.)  This shear flow is linearly stable to infinitesimal perturbations.  But if the shear profile is altered by a small amount, the alteration can itself be unstable to infinitesimal perturbations.  To be specific, let the alteration be confined within a band of width $\\Delta x$, and let it have vorticity $\\omega=\\omega(x)$ (with $|\\omega|\\lesssim q$), so that it induces a velocity field in excess of the linear shear with components $u_y\\sim \\omega\\Delta x$ and $u_x=0$. Then this band is unstable to infinitesimal nonaxisymmetric (i.e. non-stream-aligned) perturbations provided roughly that \\be \\left|k_y \\right| \\lesssim {1\\over q}{|\\omega|\\over \\Delta x} \\ \\ \\Rightarrow {\\rm 2D\\ instability} \\label{eq:2dinst} \\ee where $k_y$ is the wavenumber of the nonaxisymmetric perturbation.\\footnote{ More precisely, the necessary and sufficient condition for instability in the limit $|\\omega|\\ll q$ is that $|k_y|<{1\\over 2 q}\\int_{-\\infty}^\\infty {d\\omega/dx\\over x-x_0}dx$, where $x_0$ is any value of $x$ at which $d\\omega/dx=0$ \\citep[][L07]{Gill65,LK88}. For arbitrarily large  $\\omega$, Rayleigh's inflection point theorem and Fj\\o rtoft's theorem give necessary (though insufficient) criteria for instability \\citep{DR04}.  The former states that for instability, it is required that $d\\omega/dx=0$ somewhere in the flow, i.e. that the velocity field must have an inflection point.  \\cite{Love99} generalize Rayleigh's inflection point theorem to compressible and nonhomentropic disks. \\label{foot:inst} } For any value of $|\\omega|$ and $\\Delta x$, the band is always unstable to perturbations with long enough wavelength. Remarkably, instability even occurs when $|\\omega|$ is infinitesimal. Hence we may regard this  as a true nonlinear instability. \\cite{BH06} assert that detailed numerical simulations have not shown evidence for nonlinear instability. The reason many simulations fail to see it is that their boxes are not long enough in the $y$-direction to encompass a small enough non-zero $|k_y|$.   In two dimensions, the outcome of this instability is a long-lived vortex (e.g., L07). A vortex that has been studied in detail is the Moore-Saffman vortex, which is a localized patch of spatially constant vorticity  superimposed on a linear shear flow  \\citep{Saffman95}. When $|\\omega|\\lesssim q$, where $\\omega$ here refers to the spatially constant excess vorticity within the patch, and when the vorticity within the patch ($\\omega-q$) is stronger than that of the background shear, then the patch forms a stable vortex that is elongated in $y$ relative to $x$ by the factor \\be {\\Delta y\\over \\Delta x} \\sim {q\\over |\\omega|} \\ . \\label{eq:ms} \\ee This relation applies not only to Moore-Saffman vortices, but also to vortices whose $\\omega$ is not spatially constant. It may be understood as follows. A patch with characteristic excess vorticity $\\sim\\omega$ and with $\\Delta y\\gg \\Delta x$ induces a velocity field in the $x$-direction with amplitude $u_x\\sim|\\omega|\\Delta x$, independent of the value of $\\Delta y$ (e.g., \\S 6 in L07). As long as $|\\omega|\\lesssim q$, the $y$-velocity within the vortex is predominantly due to the background shear, and is $\\sim q\\Delta x$. Therefore the time to cross the width of the vortex is $t_x\\sim \\Delta x/u_x \\sim 1/|\\omega|$, and the time to cross its length is $t_y\\sim \\Delta y/(q\\Delta x)$.  Since these times must be comparable in a vortex, equation (\\ref{eq:ms}) follows. Equation (\\ref{eq:ms}) is very similar to equation (\\ref{eq:2dinst}). The 2D instability naturally forms into a 2D vortex. Futhermore, the exponential growth rate of the instability is $\\sim |\\omega|$, which is comparable to the rate at which fluid circulates around the vortex.  More generally, an arbitrary axisymmetric profile of $\\omega(x)$ tends to evolve into a distinctive banded structure.  Roughly speaking, bands where $\\omega<0$ contain vortices, and these are interspersed with bands where $\\omega>0$, which  contain no vortices.   (Recall that we take the background vorticity to be negative; otherwise, the converse would hold.) The reason for this  is that only regions that have $\\omega<0$ can be unstable, as may be inferred either from the integral criterion for instability given in footnote \\ref{foot:inst}, or from Fj\\o rtoft's theorem. For more detail on vortex dynamics in shear flows, see the review by \\cite{Marcus93}.   What happens in three dimensions? To date, numerical simulations of vortices in 3D disks have been reported in two papers. \\cite{BM05b} initialized their simulation with a Moore-Saffman vortex, and solved the anelastic equations in a stratified disk. They found that this vortex decayed.  As it decayed, new vortices were formed in the disk's atmosphere, two scale heights above the midplane. The new vortices survived for the duration of the simulation. \\cite{SSG06} performed both 2D and 3D simulations of the compressible hydrodynamical equations in an unstratified disk, initialized with large random fluctuations.  They found that whereas the 2D simulations produced long-lived vortices, in three dimensions  vortices rapidly decayed.  Intuitively, it seems clear that a vortex in a very thin disk will behave as it does in  2D. And from the  3D simulations described above it may be inferred that placing this vortex in a very thick disk will induce its decay. Our main goal in this paper is to understand  these two behaviors, and the transition between them. A crude explanation of our final result is that vortices decay when the 2D vortex motion couples resonantly to 3D modes, i.e., to modes that have vertical wavenumber $k_z\\ne 0$.  As described above, a vortex with excess vorticity $|\\omega|$ has circulation frequency $\\sim|\\omega|$, and $k_y/k_x\\sim |\\omega|/q$, where $k_x$ and $k_y$ are its ``typical'' wavenumbers. Furthermore, it is well-known that the frequency of axisymmetric ($k_y=0$) inertial waves is $\\Omega k_z/\\sqrt{k_x^2+k_z^2}$  (see eq.  [\\ref{eq:axi3d}]). Equating the two frequencies, and taking the $k_x$ of the 3D mode to be comparable to the $k_x$ of the vortex, as well as setting $q=3\\Omega/2$ for a Keplerian disk, we find \\be k_z\\sim k_y \\label{eq:res} \\ee as the condition for resonance. Therefore a vortex with length $\\Delta y$  will survive in a box with height $\\Delta z\\lesssim \\Delta y$, because in such a box all 3D modes have too high a frequency to couple with the vortex, i.e., all nonzero $k_z$ exceed the characteristic $k_y\\sim 1/\\Delta y$. But when $\\Delta z\\gtrsim \\Delta y$, there exist $k_z$ in the box that satisfy the resonance condition (\\ref{eq:res}), leading to the vortex's destruction. This conclusion suggests that vortices live indefinitely in disks with scale height less than their length ($h\\lesssim \\Delta y$) because in such disks all 3D modes have too high a frequency for resonant coupling. This conclusion is also consistent  with the simulations of \\cite{BM05b} and \\cite{SSG06}. Both of these works initialized their simulations with strong excess vorticity $|\\omega|\\sim q$, corresponding to nearly circular vortices. Both had vertical domains that were comparable to the vortices' width. Therefore both saw that their vortices decayed.  Had they initialized their simulations with smaller $|\\omega|$, and increased the box length $L_y$ to encompass the resulting elongated vortices, both would have found long-lived 3D vortices. \\citeauthor{BM05b}'s discovery of long-lived vortices in the disk's atmosphere is simple to understand because  the local scale height is reduced in inverse proportion to the height above the midplane.   Therefore higher up in the atmosphere the dynamics becomes more two-dimensional, and a given vortex is better able to survive the higher it is.\\footnote{ However, \\cite{BM05b} also include buoyancy forces in their simulations, which we ignore here.  How buoyancy affects the stability of vortices is a topic for future work. }   \\subsection{Organization of the Paper}  In  \\S \\ref{sec:eom} we introduce the equations of motion, and in \\S \\ref{sec:pseudo} we present two pseudospectral simulations. One illustrates the formation and survival of a vortex in a short box, and the other illustrates the destruction of a vortex in a tall box.  In \\S\\S \\ref{sec:lin}-\\ref{sec:nonlin} we develop a theory explaining this behavior. The reader who is satisfied by the qualitative description leading to equation (\\ref{eq:res}) may skip those two sections. The theory that we develop is indirectly related to the transient amplification scenario for the generation of turbulence. Even though hydrodynamical disks are linearly stable, linear perturbations can be transiently amplified before they decay, often by a large factor.  It has been proposed that sufficiently amplified modes might couple nonlinearly, leading to turbulence \\citep[e.g.,][]{CZTL03,Yecko04,AMN05}. However, to make this proposal more concrete, one must work out how modes couple nonlinearly.  In  L07, we did that in two dimensions. We showed that the 2D nonlinear instability of equation (\\ref{eq:2dinst}) is a consequence of the coupling of an axisymmetric mode with a transiently amplified mode, which may be called a ``swinging mode'' because its phasefronts are swung around by the background shear. In \\S \\ref{sec:nonlin} we show that the 3D instability responsible for the destruction of vortices is a generalization of this 2D instability. It may be understood by examining the coupling of a 3D swinging mode with an axisymmetric one.    3D modes become increasingly unstable as $|k_z|$ decreases, and in the limit that $k_z\\rightarrow 0$, the 3D instability matches smoothly onto the 2D one. Thicker disks are more prone to 3D instability because they encompass smaller $|k_z|$. ",
        "conclusions": "\\label{sec:conc}   Our main result follows from Figure \\ref{fig:marg}, which maps out the stability of a ``mother mode'' (i.e., a mode with wavevector $\\bar{k}\\bld{\\hat{x}}$ and amplitude $\\bar{\\omega}$) to nonaxisymmetric 3D perturbations. A mother mode is unstable provided that the $k_y$ and $k_z$ of the nonaxisymmetric perturbations  satisfy both $|k_y|\\lesssim \\bar{k}\\bar{\\omega}/q$ and $|k_z|\\lesssim |k_y|$, dropping order-unity constants. Based on this result, we may understand the formation, survival, and destruction of vortices. Vortices form out of mother modes that are unstable to 2D ($k_z=0$) perturbations. Mother modes that are unstable to 2D modes but stable to 3D ($k_z\\ne 0$) ones, form into long-lived vortices.  Mother modes that are unstable to both 2D and 3D modes are destroyed. Therefore a mother mode with given $\\bar{k}$ and $\\bar{\\omega}$ will form into a vortex if the disk has a sufficiently large circumferential extent and a sufficiently small scale height, i.e., if $r\\gtrsim \\bar{k}^{-1}q/\\bar{\\omega}$ and $h\\lesssim \\bar{k}^{-1}q/\\bar{\\omega}$, where $r$ is the distance to the center of the disk, and $h$ is the scale height. Alternatively, the mother mode will be destroyed in a turbulent-like state if  both $r$ and $h$ are sufficiently large ($r\\gtrsim \\bar{k}^{-1}q/\\bar{\\omega}$ and $h\\gtrsim \\bar{k}^{-1}q/\\bar{\\omega}$).   Our result has a number of astrophysical consequences. In protoplanetary disks that do not contain any vortices, solid particles drift inward. Gas disks orbit at sub-Keplerian speeds, $v_{\\rm gas}\\sim \\Omega r (1-\\eta)$, where $\\Omega r$ is the Keplerian speed and $\\eta\\sim  (c_s/\\Omega r)^2$, with $c_s$  the sound speed.  Since solid particles  would orbit at the Keplerian speed in the absence of gas, the mismatch of speeds between solids and gas produces a drag on the solid particles, removing their angular momentum and causing them to fall into the star. For example, in the minimum mass solar nebula, meter-sized particles fall in from 1 AU in around a hundred years. This rapid infall presents a serious problem for theories of planet formation, since it is difficult to produce planets out of dust in under a hundred years. Vortices can solve this problem \\citep{BS95}. A vortex that has excess vorticity $-\\bar{\\omega}$ and radial width $1/\\bar{k}$ can halt the infall of particles provided that $\\bar{\\omega}/\\bar{k}\\gtrsim (\\Omega r)\\eta$, because the gas speed induced by such a vortex more than compensates for the sub-Keplerian speed induced by gas pressure.\\footnote{ We implicitly assume here that the stopping time of the particle due to gas drag is comparable to the orbital time, which is true for meter-sized particles at 1 AU in the minimum mass solar nebula.  A more careful treatment shows that a vortex can stop a particle with stopping time $t_s$ provided that  $\\bar{\\omega}/\\bar{k}\\gtrsim (\\Omega r)(\\Omega t_s)\\eta$ \\citep{Youdin08}. } Previous simulations implied that 3D vortices rapidly decay, and so cannot  prevent the rapid infall of solid particles \\citep{BM05b,SSG06}.  Our result shows that vortices can survive within disks, and so restores the viability of vortices as a solution to the infall problem.   A more important---and more speculative---application of our result is to the transport of angular momentum within neutral accretion disks. In our simulation of a vortex in a tall box, we found that as the vortex decayed it transported angular momentum outward at a nearly constant rate for hundreds of orbital times. If decaying vortices transport a significant amount of angular momentum in disks, they would resolve one of the most important outstanding questions in astrophysics today: what causes hydrodynamical accretion disks to accrete? To make this speculation more concrete, one must understand the amplitude and duration of the ``turbulence'' that results from decaying vortices. This is a topic for future research.  In this paper, we considered only the effects of rotation and shear on the stability of vortices, while we neglected the effect of vertical gravity. There has been a lot of research in the geophysical community on the dynamics of fluids in the presence of vertical gravity, since stably stratified fluids are very common on Earth---in the atmosphere, oceans, and lakes.  In numerical and laboratory experiments of strongly stratified flows, thin horizontal ``pancake vortices'' often form, and fully developed turbulence is characterized by thin horizontal layers. \\citep[e.g., ][]{BBLC07}.  Pancake vortices are stabilized by vertical gravity, in contrast to the vortices studied in this paper which are stabilized by rotation. Gravity inhibits vertical motions because of buoyancy: it costs gravitational energy for fluid to move vertically. The resulting quasi-two-dimensional flow can form into a vortex.\\footnote{\\cite{BC00} show that a vertically uniform vortex column in a stratified (and non-rotating and non-shearing) fluid suffers an instability (the ``zigzag instability'') that is characterized by a typical vertical lengthscale $\\lambda_z\\sim U/N$, where $U$ is the horizontal speed induced by the vortex, $N$ is the Brunt-V\\\"ais\\\"al\\\"a frequency, and the horizontal lengthscale of the vortex $L_h$ is assumed to be much greater that  $\\lambda_z$ (hence the pancake structure). We may understand Billant \\& Chomaz's result in a crude fashion with an argument similar to that employed in the introduction to explain the destruction of rotation-stabilized vortices: since the frequency of buoyancy waves is $Nk_x/k_z$ (when $|k_x|\\ll |k_z|$), and since the frequency at which fluid circulates around a vortex is $U/L_h\\sim k_xU$, there is a resonance between these two frequencies for vertical lengthscale $1/k_z\\sim U/N$.} We may speculate that in an astrophysical disk vertical gravity provides an additional means to stabilize vortices, in addition to rotation. But to make this speculation concrete, the theory presented in this paper should be extended to include vertical gravity.    We have not addressed in this paper the origin of the axisymmetric structure (the mother modes) that give rise to surviving or decaying vortices. One possibility is that decaying vortices can produce more axisymmetric structure, and therefore they can lead to self-sustaining turbulence. This seems to us unlikely. We have not seen evidence for it in our simulations, but this could be  because of the modest resolution of our simulations. Other possibilities for the generation of axisymmetric structure include thermal instabilities, such as the baroclinic instability, or convection, or stirring by planets.  This, too, is a topic for future research.      \\appendix"
    },
    "0710/0710.3059_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "\\label{sec:discussion}  We considered the relation between the HMXB population and the SFH of the host galaxy. The number of HMXBs can be represented as a convolution (Eq. (4)) of the star formation history SFR(t) with the function $\\eta_{HMXB}(t)$ describing the dependence of the HMXB number on the time elapsed since the star formation event. Thus, the evolution of the HMXB population after the star formation event can be reconstructed by analyzing the distribution of HMXBs in stellar complexes with different SFHs.  Using archival optical observations, we reconstructed the spatially resolved SFH in the SMC over the past $\\sim$100~Myr (Fig. 10). For this purpose, the observed color-magnitude diagrams of the stellar population were approximated by linear combinations of model isochrones. We analyzed the stability and errors of this method for reconstructing the recent SFH and showed that its accuracy is limited by the uncertainties in the currently available models for the evolution of massive stars. However, the systematic error introduced by this factor may be ignored, since the main source of uncertainty in the solution is the Poisson noise due to the relatively small number of HMXBs in the part of the SMC investigated by XMM-Newton.  Using the derived SFHs and the spatial distribution of HMXBs in the SMC from Shtykovskiy and Gilfanov (2005b), we reconstructed the function $\\eta_{HMXB}(t)$ that describes the dependence of the HMXB number on the time elapsed since the star formation event (Fig. 12). We compared the derived dependence with the behavior of the SN II rate. The HMXB number reaches its maximum $\\sim$20--50~Myr after the star formation event, which is comparable to or exceeds the lifetime of a  $8M_\\odot$ star. This is much later than the maximum of the SN II rate. In addition, note the shortage of the youngest systems. Observationally, this manifests itself in the absence (or an extremely small number) of HMXBs with black holes in the SMC. This behavior is related to the evolution of the companion star and the neutron star spin period and is consistent with the population synthesis model calculations (Popov et al. 1998). When these results are interpreted, it should be kept in mind that the function $\\eta_{HMXB}(t)$ depends on the luminosity threshold used to select the X-ray sources. In our analysis, we used a sample with a low luminosity threshold, L$_{min}\\sim 10^{34}$~erg/s. In such a sample, low-luminosity sources, mostly Be/X systems, mainly contribute to the number of sources, while the relative contribution from systems with black holes and/or O/B supergiants, which must constitute the majority of sources in the lower time bin in Fig. 12, is small. Therefore, the time dependence of the number of bright sources (e.g.,  $>10^{37}$~erg/s) will differ from that shown in Fig. 12.  The HMXB formation efficiency in the SMC does not exceed the prediction of the N$_{HMXB}$--SFR calibration (Grimm et al. 2003). Their abnormal abundance compared to the predictions based on the emission in standard SFR indicators, such as the H$_{\\alpha}$ line, can result from a peculiarity of the SFH in the SMC."
    },
    "0710/0710.4509_arXiv.txt": {
        "abstract": "We present a scheme to solve the nonlinear multigroup radiation diffusion (MGD) equations.  The method is incorporated into a massively parallel, multidimensional, Eulerian radiation-hydrodynamic code with adaptive mesh refinement (AMR).  The patch-based AMR algorithm refines in both space and time creating a hierarchy of levels, coarsest to finest. The physics modules are time-advanced using operator splitting. On each level, separate ``level-solve'' packages advance the modules.  Our multigroup level-solve adapts an implicit procedure which leads to a two-step iterative scheme that alternates between elliptic solves for each group with intra-cell group coupling. For robustness, we introduce pseudo transient continuation ($\\ptc$).  We analyze the magnitude of the $\\ptc$ parameter to ensure positivity of the resulting linear system, diagonal dominance and convergence of the two-step scheme.  For AMR, a level defines a subdomain for refinement.   For diffusive processes such as MGD, the refined level uses Dirichet boundary data at the coarse-fine interface and the data is derived from the coarse level solution. After advancing on the fine level, an additional procedure, the sync-solve (SS), is required in order to enforce conservation.  The MGD SS reduces to an elliptic solve on a combined grid for a system of $G$ equations, where $G$ is the number of groups. We adapt the ``partial temperature'' scheme for the SS; hence, we reuse the infrastructure developed for scalar equations.  Results are presented. We consider a multigroup test problem with a known analytic solution.  We demonstrate utility of $\\ptc$ by running with increasingly larger timesteps.  Lastly, we simulate the sudden release of energy $Y$ inside an Al sphere ($r = 15$ cm) suspended in air at STP\\@.  For $Y = 11$ kT, we find that gray radiation diffusion and MGD produce similar results. However, if $Y = 1$ MT, the two packages yield different results.  Our large $Y$ simulation contradicts a long-standing theory and demonstrates the inadequacy of gray diffusion. ",
        "introduction": "\\label{intro}  This paper describes a numerical method to solve the radiation multigroup diffusion (MGD) equations.  Two themes are presented. One is the scheme itself.  We add Pseudo Transient Continuation $(\\ptc)$ to the familiar ``fully implicit'' method of Axelrod et al \\cite{AxDuRh}.  The second theme is code-specific.  Our MGD solver is embedded in a multidimensional, massively parallel, Eulerian radiation-hydrodynamic code, which has patch-based, time-and-space Adaptive Mesh Refinement (AMR) capability.   Our code's AMR framework stems from the Berger and Oliger idea \\cite{BeOl} developed for hyperbolic, compressible hydrodynamic schemes. The idea was expanded by Almgren et al \\cite{Alm} and applied to the type of elliptic solvers required for the incompressible equations of Navier-Stokes.  Howell and Greenough \\cite{HowGre} applied the Almgren et al framework to the scalar, parabolic ``gray'' radiation diffusion equation, thereby creating the start of our radiation-hydrodynamic code.  The AMR framework works as follows. A domain, referred to as the ``coarse'' or L0 level, is discretized using a uniform, coarse spatial mesh size $h_c$.\\footnote{In multiple dimensions, coordinates have their own mesh spacing.}  After advancing with a timestep $\\dtc$, the result is scanned for possible improvement.  One may refine subregions containing a chosen material, at material interface(s), or at shocks, etc. Whatever refinement criteria are used, after the subdomains are identified, specific routines define a collection of ``patches,'' which cover the subdomains.  In two dimensions, the patches are unions of rectangles; in 3D, they are unions of hexahedra.  The patches need not be connected, but they must be contained within the coarse level.  The patches denote the ``fine'' or L1 level and are discretized with a uniform, spatial mesh size $h_f$.  A typical refinement ratio $h_c/h_f$ equals two, but higher multiples of two are also allowed.  Because the original framework was designed for temporally explicit hyperbolic schemes, $\\dtc$ is restricted by a CFL condition. This implies a similar restriction for the L1 level timestep $\\dtf$. For the case, $h_c/h_f = 2$, level L1 time-advances twice using $\\dtf = \\dtc / 2$. Boundary conditions for level L1 are supplied as follows. Wherever level L1 extends to the physical boundary, the level uses the conditions prescribed by the problem. Portions of level L1's boundary which lie inside the physical domain have conditions prescribed by time and space interpolated data obtained from the L0 solution. For diffusion equations, these conditions are of Dirichlet type. The numerical solution consists of both coarse and fine grid results.  Unfortunately, as it stands, the composite solution does not guarantee conservative fluxes across the level boundaries. To maintain conservation, a separate procedure, dubbed a sync-solve (SS) is required.  The SS reduces to an elliptic unstructured grid solve on the composite grid of L0 and L1 levels.  The AMR procedure may be recursive.  That is, a level L1 grid may generate its own subdomain for refinement, i.e., a level L2.  In that case, one SS couples results from levels L1 and L2.  Once the levels advance to the L0 level time, a SS coupling all three levels ensues.  For the multigroup equations, the SS requires an unstructured grid solve for a coupled system of reaction-diffusion equations.  Our scheme for a multigroup SS is an important theme of this paper.  The MGD equations stem from a discretization of the multifrequency radiation diffusion equations. The latter is an approximation to the equations of radiation transfer, obtained by assuming the matter to be optically thick, which suppresses the directional dependence of the radiation intensity. Details of the derivation may be found in various sources: Mihalas and Mihalas \\cite{MM2}, Zel'dovich and Raizer \\cite{ZelRai}, Pomraning \\cite{Pom}.  The gray radiation diffusion equation is a simplification of the MGD equations.  It is essentially a one-group equation and is derived by integrating over all frequencies.  Surprisingly, it gives very good results in many cases.  However, it clearly cannot display frequency-dependent effects.  When those are important, it gives incorrect results.  Unfortunately, unless one solves a problem with both gray and MGD, one never knows when the former is adequate.  We now summarize the paper. Our MGD scheme consists of two parts. Sections~\\ref{levelsolve} and \\ref{MGanlsys} develop the level-solve algorithm, which is applied on each level. Section~\\ref{levelsolve} develops the equations, the discretization, and our $\\ptc$ scheme. Section~\\ref{MGanlsys} proves three lemmas which determine the initial magnitude of the $\\ptc$ parameter $\\gs$.  % Our philosophy for $\\gs$ is as follows.  The result of the level solve is the time-advanced radiation group energy density, which physics dictates to be nonnegative.  Zeroing anomalously negative values is not an option since they are the correct conservative solution to the linear system that stems from the discretization of the system.  Thus, the unphysical result nonetheless conserves energy.  The difficulty is avoided if in the original formulation of the linear system $ A x = b$, $A$ is an M-matrix and the right-hand-side (RS) is nonnegative. Since we solve $ A x = b$ using an iterative scheme, the magnitude of $\\gs$ is determined to ensure $b \\ge 0$, a diagonally dominant $A$, and that the iterations converge. To a large extent, we are guided by Pert~\\cite{Pert}, who discusses how and why the solution to a discretization of an equation may be unacceptable from a physical standpoint. For a first reading, section~\\ref{MGanlsys} may be skipped; the analysis of the required magnitude of $\\gs$ is not needed for the subsequent sections.  We note that $\\ptc$ is widely used to solve nonlinear systems of equations.  It is closely related to the Inexact Newton Backtracking Method by Shahid et al \\cite{ShTuWa}.  When applying $\\ptc$ to a Newton solver, the basic idea is to limit the change to the iterates when one is far from the root but not restrict the change as one approaches the root.  With $\\ptc$, limiting is done by the magnitude of the pseudo-timestep.  Kelley and Keyes \\cite{KelKey} put $\\ptc$ on a solid analytic framework by examining the three regimes of $\\ptc$: small, medium, and large pseudo-timesteps.  In the last regime, $\\ptc$ recovers Newton's second order of convergence.  Our $\\ptc$ implementation differs from the norm. Standard applications typically detect when a problem is ``hard'' and then reduce the timestep or some other parameter by an arbitrary amount. However, this method will not work for us because our solver is embedded in a time-dependent multiphysics code with separate modules for compressible gasdynamics, heat conduction, radiation transport.  Our MGD solver is called numerous times during the course of a simulation.  (If running, with AMR, multiple times per physical time advance.) Although the physical $\\dt$ is controlled by various means, and depending on the problem can vary many orders of magnitude, we require a MGD solver that works under all conditions. Our $\\ptc$ approach is similar to the one of Shestakov et al \\cite{ShHoGr}. We set the initial magnitude of the $\\ptc$ parameter to ensure that for the first step, our iteration scheme converges and that the result is physical. We note that our usage of $\\ptc$ is nearly equivalent to having the MGD module time-advance not in a single (physical) step $\\dt$, but in smaller time increments until the desired time $t^0 + \\dt$ is reached. Some colleagues refer to the process as ``sub-cycling'' the radiation module. It is easy to show that the lemmas of Sec.~\\ref{MGanlsys} still apply for sub-cycling.  Section~\\ref{mgamr} describes the second part of our solver, viz., the sync-solve. Section~\\ref{rapAMR} contains results. Three problems are presented.  The first, in Sec.~\\ref{linwin}, displays the accuracy of the method and its convergence properties: first order in time and second order in space.  Section~\\ref{PTCrobust} demonstrates the utility afforded by $\\ptc$.  For hard problems, it accelerates convergence; for very hard problems, $\\ptc$ is indispensable.  Section~\\ref{hotball} models the explosive expansion of a hot metal sphere suspended in cold air. The simulation couples all of the code's physics modules. The problem is an ideal candidate for AMR since effects propagate a large distance away from the source, yet in early times, resolution is needed only near the sphere.  The problem also demonstrates the necessity of multigroup diffusion. We find that if the sphere's energy is very high, gray diffusion gives the wrong answer.  For a 1 MT energy source, our MGD simulation contradicts results of Brode \\cite{Bro}, who used gray diffusion. Section~\\ref{conclusion} contains concluding remarks.  There are three appendices.  Appendix~\\ref{table} gives a table of exact values for the test problem described in section~\\ref{linwin}. Appendix~\\ref{apb} discusses situations that may complicate attaining a diagonally dominant matrix when discretizing the multigroup system. Appendix~\\ref{apc} presents a spatial convergence analysis of the multigroup system when running in ``production'' mode, that is, with a dominant flux limiter and with AMR. ",
        "conclusions": "\\label{conclusion}  We have described a numerical scheme to solve the radiation multigroup diffusion equations.  The scheme is implemented in a radiation-hydrodynamic code with the patch-based AMR methodology, originally proposed by Berger and Oliger \\cite{BeOl} for hyperbolic partial differential equations.  Our scheme consists of two parts. The first, described in Sections~\\ref{levelsolve} and \\ref{MGanlsys}, is applied on a {\\em level\\/} of the AMR grid layout and may be adapted to any code. This part consists of adding $\\ptc$ to the ``fully-implicit'' iterative scheme of Axelrod et al \\cite{AxDuRh}. $\\ptc$ brings an extra degree of robustness and enhances convergence of the Axelrod scheme.  We have developed lemmas that determine the minimum magnitude for the $\\ptc$ parameter $\\tau$ to ensure that the iterations converge and the result is physically meaningful. The appropriate magnitude depends on the problem.  Our implementation of $\\ptc$ is not optimal---at least for our AMR code architecture.  In our code, for each AMR level, we compute a {\\em single\\/} scalar parameter $\\tau$.  However, the levels consist of a collection of grids (rectangles in 2D) that need not be connected. If the grids are not connected, they form independent problems.  Hence, it would be more efficient to use different $\\tau$ for disconnected grids.  The second part of our scheme,  the sync-solve (SS), addresses a specific need of our code, viz., the requirement of having an energy-conserving result on the composite grid of multiple AMR levels.  For the multigroup equations, this part reduces to a coupled system of elliptic equations on the unstructured grid combining all levels. Since the SS is intended to be a small correction to the result of the level solves, we adapted the key element of the ``partial temperature'' scheme of Lund and Wilson \\cite{LunWil}.  This allowed reducing the multigroup SS to a collection of scalar SS's.  We were then able to reuse existing software.%  This paper included simulations of three problems. The first two are idealized tests of only the multigroup module. The third is a ``real'' problem, which uses the full capability of the code: AMR, multiple materials, etc. The first problem was chosen because of its non-triviality and the availability of analytic results with which to compare.  We obtained excellent agreement and verified the convergence properties of the scheme.  The second problem illustrated the benefits brought by $\\ptc$.  We compared the conventional scheme of Axelrod et al \\cite{AxDuRh} with our $\\ptc$-modified version. For hard problems, $\\ptc$ either decreased run times or ensured convergence in regimes where the conventional scheme diverged.  The third problem showed that our multigroup module has been fully integrated into the code and has already extended the scientific frontier. For a high yield air burst at STP, we found that gray diffusion gives an incorrect result during the radiation-dominated regime because gray fails to capture the frequency-dependent effects of the air opacity."
    },
    "0710/0710.2213_arXiv.txt": {
        "abstract": "We investigate gravitino dark matter scenarios in which the primordial $\\Lisix$ production is catalyzed by bound-state formation of long-lived negatively charged particles $\\champ$ with $\\Hefour$. In the constrained minimal supersymmetric Standard Model (CMSSM) with the stau $\\stau^-$ as the $\\champ$, the observationally inferred bound on the primordial $\\Lisix$ abundance allows us to derive a rigid lower limit on the gaugino mass parameter for a standard cosmological history. This limit can have severe implications for supersymmetry searches at the Large Hadron Collider and for the reheating temperature after inflation. ",
        "introduction": "\\label{sec:introduction}  Big Bang Nucleosynthesis (BBN) is a powerful tool to test physics beyond the Standard Mo\\-del. Recently, it has been realized that the presence of heavy long-lived negatively charged particles $\\champ$ can have a substantial impact on the primordial light element abundances via bound-state formation% ~\\cite{Pospelov:2006sc,Kohri:2006cn,Kaplinghat:2006qr,Cyburt:2006uv,Hamaguchi:2007mp,Bird:2007ge,Kawasaki:2007xb,Jittoh:2007fr,Jedamzik:2007cp}. In particular, when $\\champ$ and $\\Hefour$ form Coulomb bound states, $(\\Hefour\\champ)$, too much $\\Lisix$ can be produced via the catalyzed BBN (CBBN) reaction~\\cite{Pospelov:2006sc} \\begin{align} \\label{eq:CBBN-reaction} (\\Hefour\\champ)+\\mathrm{D} \\rightarrow \\Lisix + \\champ . \\end{align} The formation of $(\\Hefour\\champ)$ and hence the CBBN production of $\\Lisix$ becomes efficient at temperatures $T \\sim 10\\ \\keV$, i.e., at cosmic times $t> 10^3\\ \\seconds$ at which standard BBN (SBBN) processes are already frozen out.  The observationally inferred bound on the primordial $\\Lisix$ abundance then restricts severely the $\\champ$ abundance at such times.  A long-lived $\\champ$ may be realized if the gravitino is the lightest supersymmetric particle (LSP).  In particular, it is reasonable to consider gravitino LSP scenarios within the constrained minimal supersymmetric Standard Model (CMSSM)~\\cite{Ellis:2003dn,Cerdeno:2005eu,Jedamzik:2005dh,Cyburt:2006uv,Pradler:2006hh} in which the gaugino masses, the scalar masses, and the trilinear scalar couplings are parameterized by their respective universal values $\\monetwo$, $\\mzero$, and $A_0$ at the scale of grand unification $\\mgut \\simeq 2\\times 10^{16}\\ \\GeV$.  Within this framework, the lighter stau $\\stau$ is the lightest Standard Model superpartner in a large region of the parameter space and thus a well-motivated candidate for the next-to-lightest supersymmetric particle (NLSP).  Since its couplings to the gravitino LSP are suppressed by the (reduced) Planck scale, $\\MP=2.4\\times 10^{18}\\,\\GeV$, the stau will typically be long-lived for conserved $\\mathrm{R}$-parity% \\footnote{For the case of broken R-parity, see, e.g.~\\cite{Buchmuller:2007ui}} and thus $\\stau^-$ can play the role of $\\champ$.  In scenarios with conserved $\\mathrm{R}$-parity, the gravitino LSP is stable and a promisig dark matter candidate. Gravitinos can be produced efficiently in thermal scattering of particles in the primordial plasma. If the Universe, after inflation, enters the radiation dominated epoch with a high reheating temperature $\\TR$, the resulting gravitino density $\\Omegatp$ will contribute substantially to the dark matter density $\\Omega_{\\mathrm{dm}}$~\\cite{Bolz:2000fu,Pradler:2006qh,Rychkov:2007uq}.  In this work we calculate the amount of $\\Lisix$ produced in~(\\ref{eq:CBBN-reaction}) by following the treatment of Ref.~\\cite{Takayama:2007du}. In particluar, we employ a recent state-of-the-art result for the CBBN reaction cross reaction~\\cite{Hamaguchi:2007mp}. The obtained upper limit on the $\\champ$ abundance from possible $\\Lisix$ overproduction vanishes for sufficiently short $\\tau_{\\champ}$. This allows us to extract a lower limit on the universal gaugino mass parameter $\\monetwo$ within minimal supergravity scenarios where the gravitino is the LSP and the $\\champ$ is the $\\stau^-$ NLSP.\\footnote{In this work we assume a standard cosmological history with a reheating temperature $\\TR$ that exceeds the freeze-out temperature $\\Tf$ of the $\\stau$ NLSP.} This limit leads directly to an upper bound on $\\TR$ since $\\Omegatp$ cannot exceed the observed dark matter density. The bounds on $\\monetwo$ and $\\TR$ derived below depend on the gravitino mass but are independent of the CMSSM parameters.  Before proceeding, let us comment on the present status of BBN constraints on gravitino dark matter scenarios with a long-lived charged slepton NLSP. In a recent ambitious study~\\cite{Jedamzik:2007cp} it is argued that bound-state formation of $\\champ$ with protons at $T \\sim 1\\ \\keV$ might well reprocess large fractions of the previously synthesized $\\Lisix$.  This seems to relax the bound on the $\\champ$ abundance for $\\tau_{\\champ}> 10^6~\\seconds$.  However, at present, the uncertainties in the relevant nuclear reaction rates in~\\cite{Jedamzik:2007cp} make it difficult to decide whether a new cosmologically allowed region will open up.  In this work we assume that this is not the case, in particular, since the $^3$He/D constraint on electromagnetic energy release~\\cite{Sigl:1995kk} becomes severe in this region and excludes stau lifetimes $\\tau_{\\stau}\\gtrsim10^6~\\seconds$~\\cite{Cerdeno:2005eu,Cyburt:2006uv,Kawasaki:2007xb,Jedamzik:2007cp}. Then only the constraint from hadronic energy release on D~\\cite{Kawasaki:2004qu,Feng:2004mt,Cerdeno:2005eu,Jedamzik:2006xz,Steffen:2006hw} can be slightly more severe than the one from catalyzed $^6$Li production~\\cite{Cyburt:2006uv,Steffen:2006wx,Pradler:2006hh,Kawasaki:2007xb}. We neglect the D constraint in this work since it can only tighten the bounds on $\\monetwo$ and $\\TR$ as can be seen, e.g., in Figs.~4~(b--d) and~5 of Ref.~\\cite{Pradler:2006hh}. For deriving conservative bounds on $\\monetwo$ and $\\TR$, it is thus sufficient to consider the CBBN reaction~(\\ref{eq:CBBN-reaction}) exclusively. ",
        "conclusions": "\\label{sec:conclusion} We have considered the catalysis of $\\Lisix$ production in CMSSM scenarios with the gravitino LSP and the stau NLSP. Within a standard cosmological history, the calculated $\\Lisix$ abundance drops below the observational limit on primordial $\\Lisix$ for $\\tau_{\\stau} \\lesssim 5\\times 10^3\\,\\seconds$. Taken at face value, we find that this constraint translates into a lower limit $\\monetwo \\ge 0.9\\, \\TeV ( \\mgr / 10\\, \\GeV )^{2/5}$ in the entire natural region of the CMSSM parameter space. This implies a conservative upper bound $\\TR\\lesssim 4.9\\times 10^7\\GeV( \\mgr / 10\\, \\GeV )^{1/5}$. The bounds on $\\monetwo$ and $\\TR$ not only confirm our previous findings~\\cite{Pradler:2006hh} but are also independent of the particular values of the CMSSM parameters for the considered $\\stau$ NLSP abundances."
    },
    "0710/0710.2814_arXiv.txt": {
        "abstract": "CO isotopes are able to probe the different components in protostellar clouds. These components, core, envelope and outflow have distinct physical conditions and sometimes more than one component contributes to the observed line profile. In this study we determine how CO isotope abundances are  altered by the physical conditions in the different components. We use a 3D molecular line transport code to simulate the emission of four CO isotopomers, $^{12}$CO $J=2\\rightarrow1$, $^{13}$CO $J=2\\rightarrow1$, C$^{18}$O $J=2\\rightarrow1$ and C$^{17}$O $J=2\\rightarrow1$ from the Class 0/1 object L483, which contains a cold quiescent core, an infalling envelope and a clear outflow. Our models replicate JCMT (James Clerk Maxwell Telescope) line observations with the inclusion of freeze-out, a density profile and infall. Our model profiles of $^{12}$CO and $^{13}$CO have a large linewidth due to a high velocity jet. These profiles replicate the process of more abundant material being susceptible to a jet. C$^{18}$O and C$^{17}$O do not display such a large linewidth as they trace denser quiescent material deep in the cloud. ",
        "introduction": "Molecules, particularly CO, are used as tracers of H$_2$ and thus of gas density in cold dark clouds. CO is highly abundant with a low critical density and typically exhibits strong optical depth effects in cold dark molecular clouds. The four most common CO isotopes differ in abundance by as much as three orders of magnitude. Thus they become optically thick at different column densities. Taken together, observations of CO isotopes can trace the gas density in all the main components of cold, dark clouds: the intermediate optical depth envelope, the high optical depth core, and the optically thin bipolar outflows. However, we know from observational and theoretical studies that the abundance of CO depends on conditions in the clouds such as, shock heating \\citep{van_dishoeck_95,nisini_07}, U.V. excitation \\citep*{goldsmith_07}, freeze-out \\citep*{lee_04} and varies from place to place. Therefore we cannot use CO as an H$_2$ tracer without understanding its chemical variation. To better understand CO variations in cold dark clouds, we observed and modeled one particular cloud, Lynds 483. We chose L483 as a prototype for a study of CO abundances because it is a well studied nearby ($\\sim 200~{\\rm pc}$) molecular cloud. It contains an IRAS source 18148-0440 that is in transition between a Class 0 and Class 1 object \\citep{tafalla.et.al00}. It exhibits an infalling envelope \\citep{park.et.al99,park.et.al00,tafalla.et.al00} and a slow bipolar molecular outflow \\citep{fuller.et.al95,buckle.et.al99} yet the core and envelope are still cold and dense \\citep{ladd.et.91,fuller.et.92,fuller&wootten00}. Thus many of the physical properties and kinematic features that are present in either less or more evolved clouds are all present in L483.  Our aim is to combine these components into a single model for L483 to strongly constrain the structure and dynamics of the system and hence then to infer the CO abundance throughout the cloud.  \\begin{figure*} \\centering \\includegraphics[width=180mm]{fig1} \\caption{An example line profile of each of the four isotopes used in this study where the offsets with respect to IRAS 18140-0440 are indicated in the top right corner of each panel. The $^{12}$CO y-axis is in units of $T_{r}$ and $^{13}$CO, C$^{18}$O and C$^{17}$O are in units of $T_{mb}$ } \\label{overview} \\end{figure*}  We obtained emission line profile data from L483 for transitions of the four most common CO isotopes (section ~\\ref{obs}). In addition, we used an archival dust continuum emission map of L483 as an unbiased mass tracer to be compared with the CO data. We used a radiative transfer model to analyse the data and to produce synthetic spectra to be compared in detail with individual observations (section ~\\ref{models}).  Estimates for density, temperature and a physical model of L483 were constrained with help from the observational literature. The abundances of the CO isotopes were then varied in the model to give a good match with the observed line profiles. This yields the abundance variation throughout L483 as well as self-consistent temperatures, densities and velocities and abundance ratios. ",
        "conclusions": "We modelled L483 in four isotopomers of CO in order of decreasing abundance $^{12}$CO, $^{13}$CO, C$^{18}$O and C$^{17}$O. Each species delineates a different region and therefore we get a clearer picture of cloud dynamics.  C$^{18}$O and C$^{17}$O are optically thin and are important tracers of the denser regions of star forming clouds. Optically thick species such as $^{12}$CO and $^{13}$CO are more abundant and more susceptible to the jet motion. They trace regions farther out from the centre of the cloud and have line profiles showing a large line width consistent with material exposed to a high velocity jet. Using a radiative transfer code and a dynamical model for L483 we were able to self-consistently calculate for the first time the abundance of the CO isotopes in the different regions of such a cloud.  Our principal finding  is that the CO line profiles in L483 are well fitted with a self-consistent envelope plus boundary layer model and that the CO abundances increase substantially in this boundary layer. The most likely reason for this is that molecular ices on dust grains are heated and released back into the gas phase in the boundary layer. A constant abundance model was found to overestimate the abundance towards the centre of the cloud and only freeze-out of material towards the centre was able to produce modeled profiles consistent with observations. Our tanh geometry is chosen because it matches the observed morphology seen in other protostellar outflows \\citep{tafalla.et.al00}.  Other geometries are possible for the outflow e.g. conical, cylindrical outflow but it is unlikely that it would have a substantial effect on our results because such a detailed treatment of the boundary layer components is difficult to achieve until there is sub-arcsecond resolution resolution.  We emphasize that our results provide an abundance enhancement measurement rather than proving an exact mechanism by which the CO is enhanced, e.g. chemical reactions or dissociation. A more detailed treatment would involve a full dark cloud and gas-grain chemistry whilst accounting for localized shock heating. The most enhanced species in our study, by a factor of $\\sim 30$, is the $^{13}$CO material. The exothermic reaction leading to the creation of $^{13}$CO \\citep{duley_williams} is shown below \\begin{equation} \\; \\; \\; \\; \\; ^{13}{\\rm C}^{+} + ^{12}{\\rm CO} \\rightleftharpoons ^{12}{\\rm C}^{+} + ^{13}{\\rm CO} + \\Delta{\\rm E} \\label{13co_creation} \\end{equation} where the zero-point energy difference $\\Delta{\\rm E}$ is equivalent to a temperature $\\Delta{\\rm E}/{\\rm k}$ of 35~K. This mechanism may be the source for the enhanced abundance observed from our modeling.  The enhanced abundance seen in the boundary layer effect may also be detectable in other molecules. \\citet{park.et.al00} used interferometric observations of HCO$^{+}$ and observed anti-infall profiles close to the centre of the cloud. They concluded the HCO$^{+}$ was tracing the outlying regions of the outflow, i.e. a region between the envelope and the jet. The reason the HCO$^{+}$ emission is predominately seen here rather than in the more extensive envelope is also likely due to an enhancement of HCO$^{+}$ caused by the shock-heated release of icy grain mantles followed by chemical reaction. CO and H$_2$O are liberated into the gas phase and the shock-induced radiation field then can photodissociate CO to C$^+$. This then reacts with the H$_2$O to form HCO$^{+}$. Such a model was successfully used to explain the enhancement of HCO$^{+}$ commonly seen at the bases of molecular outflows  \\citep{rawlings.et.al00,rawlings.et.al04}. The results in this paper demonstrate that a combination of datasets with several lines and transitions coupled with a 3D molecular line transport code is a powerful way to determine the properties of dense star forming cores."
    },
    "0710/0710.5119_arXiv.txt": {
        "abstract": "{The excess above 1 GeV in the energy spectrum of the diffuse Galactic gamma radiation, measured with the EGRET experiment, can be interpreted as the annihilation of Dark Matter (DM) particles. The DM is distributed in a halo around the Milky Way. Considering the directionality of the gamma ray flux it is possible to determine the halo profile. The DM within the halo has a smooth and a clumpy component. These components can have different profiles as suggested by N-body simulations and the data is indeed compatible with a NFW profile for the diffuse component and a cored profile for the clumpy component. These DM clumps can be partly destroyed by tidal forces from interactions with stars and the gravitational potential of the Galactic disc. This effect mainly decreases the annihilation signal from the Galactic centre (GC). In this paper constraints on the different profiles and the survival probability of the clumps are discussed. \\PACS{ {95.35.+d}{Dark Matter}   \\and {98.35.Gi}{Galactic halo} } % } % ",
        "introduction": "\\label{seq:intro} From WMAP measurements of the temperature aniso- tropies in the Cosmic Microwave Background (CMB) in combination with data on the Hubble expansion and the density fluctuation in the Universe \\cite{RefSpergel} we gather that Cold Dark Matter (CDM) makes up 23\\% of the energy of the Universe. The nature of the Dark Matter (DM) is unknown, but one of the most promising particles is the \"weakly interacting massive particle\" (WIMP). Assuming that WIMPS are Majorana particles they can annihilate each other and produce a large amount of secondary particles. For the determination of the density distribution, the so-called DM halo profile of the WIMP particles, the gamma radia\\-tion is most important because it is not influenced by the magnetic field of the galaxy and points back directly to its source.\\\\ The observation of the diffuse Galactic gamma radiation of the Milky Way with EGRET showed an excess above 1 GeV in the photon energy spectrum. This excess is different for various sky directions and can be interpreted as a WIMP annihilation signal \\cite{RefSander}. Therefore it can be used to determine the DM halo profile.\\\\ In section \\ref{seq:halo}, we will describe the mechanism of the determination of the DM halo profile from the EGRET excess and explain how to calculate the DM annihilation flux of gamma rays. Then, after differentiating between diffuse and clumpy DM, we will introduce a survival probability for DM clumps as well as a ringlike substructure of DM within the Galactic disc. ",
        "conclusions": "CDM probably consists of WIMPS which are heavy and slow particles. If these particles are Majorana particles they can annihilate each other and produce Galactic gamma radiation which can be used to determine the density profile of the DM. In this analysis the directionality of the DM annihilation flux measured with EGRET was used to find a possible DM halo profile. After dividing the Galactic gamma ray flux into 4 latitude and 45 longitude regions the background and the DM annihilation signal were fitted to the data for each of the 180 bins. The DM annihilation flux is dominated by the annihilation flux of the clumpy DM which is proportional to $\\rho_{\\chi, clump}$, not $\\rho^2_{\\chi, clump}$. While the diffuse DM component has a cuspy NFW profile a shallower cored distribution was obtained for the clumpy component. The DM annihilation flux is dominated by the clumpy DM component, but the clumpy component yields a mass below the required mass $> 10^{12}$ solar masses \\cite{RefBattaglia}. However, if combined with the diffuse cuspy NFW profile, both the EGRET data and the mass constraint can be fulfilled. In order to take the tidal disruption of DM clumps in the vicinity of stars into account a survival probability for clumps was introduced. Most of the clumps are expected to be destroyed near the Galactic centre, although a steep cusp may survive. The strong signal observed from the Galactic centre yielded a survival probability at the centre of $P(0)=0.7$. This means that the DM clumps are not completely destroyed, which is in good agreement with more detailed calculations in Ref. \\cite{RefDokuchaev2}. In addition to the DM halo profiles two ringlike substructure were required at radii of 4 and 13 kpc. The halo and ring parameters were obtained by minimizing a $\\chi^2$ function comparing the flux of the excess from the various sky directions with the line-of-sight integral in the halo. Figure \\ref{fig:longitudes} shows that the halo model fits the measured data very well.\\\\ In summary, the EGRET excess of diffuse Galactic gamma rays is in good agreement with the expectations of a cored clumpy halo component plus a cuspy diffuse one. The ringlike substructure, expected from the tidal disruption of the nearby Canis Major dwarf galaxy, is clearly seen and its heavy mass above $10^{10}$ solar masses as obtained from the EGRET data, has been recently confirmed by the reduced gas flaring in this region."
    },
    "0710/0710.5269_arXiv.txt": {
        "abstract": "We investigate the full $5D$ dynamics of general braneworld models. Without making any further assumptions we show that cyclic behavior can arise naturally in a fraction of physically accepted solutions. The model does not require brane collisions, which in the stationary case remain fixed, and cyclicity takes place on the branes. We indicate that the cosmological constants play the central role for the realization of cyclic solutions and we show that its extremely small value on the observable universe makes the period of the cycles and the maximum scale factor astronomically large. ",
        "introduction": "The last decade proves to be really exciting for cosmology. Observational data indicated, among other very interesting results, that the expansion of the universe is accelerated \\cite{observ}. At the same time the braneworld scenario appeared in the literature \\cite{Horawa,RS99}. Though the exciting idea that we live in a fundamentally higher-dimensional spacetime which is greatly curved by vacuum energy was older \\cite{Rubakov83}, the new class of ``warped\" geometries offered a simple way of localizing the low energy gravitons on the brane.  In this novel background the old idea of a cyclic Universe was reheated. Started as ekpyrotic \\cite{Khoury.best,Khoury.rebound}, enriched to ekpyrotic/cyclic \\cite{ekpyrotic1,Khoury.rebound,Turok.simplified,Turok.sing,ekpyrotic.pert,cyclic.clifton,chargeBH} and recently to new ekpyrotic \\cite{ekpyrotic3,ekpyrotic4,ekpyrotic2,ekpyrotic.fields}, the new paradigm tries to be established as an alternative to standard cosmology. According to its basic contents, our universe experiences an infinite or extremely large number of cycles, each one consisting of a hot bang phase, a phase of accelerated expansion, a phase of slow-ekpyrotic contraction and a bounce-bang that triggers the next cycle. Starting with a simplified notional framework (infinite and not ``created\" time) cyclic cosmology have many advantages. It successfully faces the homogeneity, isotropy, topological and flatness problems, it handles the issue of initial conditions, it incorporates the dark energy and transforms it to an important factor, and it provides the mechanism of the generation of cosmic perturbations and of structure formation. However, there are some key issues that do not have a consistent and efficient approach so far, despite the great progress. These are the settlement of the singularity, although temperature and density remain finite, the entropy evolution, and the fate of the perturbations through the bounce. Through this research, cyclic scenarios have become more complicated, by the insertion of more complex potentials, of more branes \\cite{Khoury.best}, of the mechanism of ghost condensation \\cite{ghcond,ekpyrotic3}, of more scalar fields \\cite{ekpyrotic.fields} and of procedures which cancel the tachyonic instabilities \\cite{ekpyrotic4}.  Most of the works on cyclic cosmology involve, initially or at some stage, the transition to effective $4D$ equations. However, as it was mentioned in \\cite{Linde01,4Dbreakdown}, such a procedure does not lead to reliable results since one cannot return to the $5D$ description self-consistently. Furthermore, the old $4D$-singularity problem (of both Big Bang and traditional cyclic universes), has been replaced by a new one (equally annoying) concerning the singularity of extra dimension(s). This later case is accompanied by the brane collision phenomenon, which seems to be a basic constituent of the ekpyrotic scenario.  In this work we desire to investigate the full $5D$ dynamics of general braneworld models and examine if a cyclic behavior is possible. This is an essential procedure in order to consistently confront the arguments of the authors of \\cite{Linde01}, which claim that cyclic behavior cannot arise from a complete $5D$ description, and our study must not include any additional assumptions or fine tunings in order to remain general and therefore convincing. Secondly, we are interested to explore if a cyclic behavior of $5D$ dynamics is necessarily related to brane collisions. This work is organized as follows: In section \\ref{model} we present the $5D$ braneworld model and we derive the equations of motion. In section \\ref{analyt} we provide analytical solutions for two simplified stationary solution subclasses, while in \\ref{numer} we investigate numerically the full stationary dynamics. Finally, in section \\ref{discussion} we discuss the physical implications of our analysis and we summarize the obtained results. ",
        "conclusions": "\\label{discussion}  In the aforementioned analysis we considered general braneworld models characterized by the action (\\ref{action}), the conformal metric (\\ref{metric}), and the general potentials (\\ref{bulkpot}) and (\\ref{branepot}). Performing both analytical and numerical calculations we showed that the full $5D$ dynamics allows for stationary solutions corresponding to oscillatory scale factor of the physical brane and therefore to cyclic universes. In statistical terms cyclicity corresponds to $\\approx 4\\%$ of the physical solutions. Our investigation is completely $5D$, cyclic behavior arises naturally and is induced on the brane by the full dynamics, and it is not a result of a modified $4D$ dynamics, with fine-tuned parameters or specific assumptions in the Friedmann equation. Furthermore, we do not use an explicit brane state equation, considering just the bulk scalar field (the decays and interactions of which will eventually fill the physical brane with the conventional content \\cite{apostol}). As we mentioned in the introduction this full $5D$ approach is necessary in order to confront the arguments of the authors of \\cite{Linde01}. Indeed, their allegations that one cannot transit to an effective $4D$ theory (integrating the action over $y$), solve the equations there and then return naively to the $5D$ description (adding time-dependence by hand), are correct. Doing so, the results are not self-consistent (especially the boundary conditions are not satisfied) and the authors of \\cite{Linde01} use this fact as a central argument against the cyclic scenario. However, our consistent $5D$ analysis reveals that cyclic behavior is possible.  Another important feature of the present study is that cyclic universes do not require brane collisions. Thus, we avoid the known problems concerning such a description, which force ekpyrotic model to successively more complicated versions. On the contrary, the branes do not move at all and the system is stable (stationary solutions are a stable fixed point \\cite{brcod,Tetradis01}). Furthermore, in our model, expansion and contraction take place in the 3+1 branes, and in all 3+1 slices in general, while the fifth dimension remains unaffected. The $4$ spacial dimensions shrink periodically to an $1D$ string and re-expand. This is in a radical contrast with the cyclic models with extra dimensions, where the extra dimension is the one that gets contracted (the fifth in \\cite{ekpyrotic1} or the eleventh in \\cite{Khoury.rebound}). Cyclicity seems to re-obtain its ``physical\" meaning.  Our $5D$ investigation is general and does not involve extra assumptions, fine-tunings or specific potential forms. We result to periodic, cyclic, homogenous and isotropic universes, where the scale factor changes smoothly from expanding to contracting. An observer on the physical brane feels successively accelerated expansion, decelerated expansion, turnabout, accelerated contraction, decelerated contraction, bounce e.t.c, and a promising signature of the cyclic behavior would be the measure of the varying rate of the Hubble constant. The cycles period, given in (\\ref{period}), can be arbitrary, depending on $B(0)$, i.e. on the value of the warp factor on the physical brane ($\\theta$ is bounded from above and therefore cannot act as a period-decreasing factor). A very interesting conclusion comes from the insertion of observational results in our model, which was not made above in order to remain as general as possible. Explicitly, if we use the fact that the cosmological constant of our Universe is extremely small ($\\approx{\\cal O}(10^{-47})\\ \\text{GeV}^4$), and assuming a reasonable $M_5$ value of ${\\cal O}(10^{19})$ GeV, the first two boundary conditions in relation (\\ref{junctions}) provide in general a huge value for $e^{B(0)}$ ($\\approx{\\cal O}(10^{45})$). This is in consistency with the scaling transformation of \\cite{brcod}, which allows us, in a solution, to scale the parameters by $e^{-S}$ and add to the warp factor the constant $S$, and acquire another solution. Therefore, the extremely small cosmological constant of the observable universe leads the cyclicity period to be around $T\\approx{\\cal O}(10^{13})$ years and the maximum scale factor value, given by (\\ref{minmax}), to be $a_{max}\\approx {\\cal O}(10^{28})$ m (where the decimal exponents in these rough estimations can vary by 1 or 2, depending on $B'(0)$ and $\\theta$ values). Luckily enough, the smallness of the cosmological constant excludes oscillatory models with small periods in astronomical terms. In more foundational words, the reason that made the cosmological constant that small, is the same that makes the cycle period and the size of the Universe that large.  In this work we have been restricted to stationary solutions, where the subclass of them that possesses $H^2<0$ corresponds to eternal cyclic behavior with constant period. Numerical investigation of the full dynamics seem to consist of such stationary solutions and the transitions between them \\cite{brcod,Tetradis01}. In such transitions $H_0^2$ on the physical brane can chance sign, leading to a form of ``chaotic cyclicity\", where large intervals of (non-periodic in general) oscillatory behavior could be followed by large intervals of conventional evolution and vice versa. In this case, an initial Big Bang and/or a final Big Rip or Big Crunch (in conventional terms) could be possible. Another interesting possibility would be the exploration of our model with cosmological constants being piecewise constant functions of time, reflecting cosmological phase transitions, which could also lead to chaotic cyclicity. Note however that numerical confirmation of such behaviors is very hard due to the small probability of cyclic stationary solutions ($\\approx 10^{-2}\\%$ as we have already mentioned). These subjects are under investigation.  In order for a model to serve as a description of nature, it has to explain the basic physical key issues. Especially for cyclic cosmology, amongst others these are the entropy evolution and, probably the most pressing issue, that of a fuller understanding of the bounce and the handling of the singularity. Our model provides a consistent background for cyclicity and it reveals how such a behavior arises from the full $5D$ dynamics. However, since braneworlds and brane cosmology in general arise as limits of a multi-dimensional theory unknown up to now, the $5D$ results have a phenomenological character and must be considered from this point of view. Definitely, a complete explanation and apprehension, and a successful confrontation of the aforementioned subjects, can only come through a higher-dimensional, fundamental theory of nature. For the moment we have to rely on the relevant research on cyclic cosmology, linearized gravity, M-theory and strings, which has improved our knowledge on these issues. These results can be embodied in our analysis. The most hopeful effort is the use of quantum fluctuations in order to tame the singularity, which effectively is translated into a modification of gravity by the scalar field \\cite{Bojowald01,ekpyrotic3,ekpyrotic4}. Alternatively, using loop quantum gravity we could modify non-perturbatively the dynamical equations leading to a singularity resolution as in \\cite{Maartens}. Concerning the entropy, we could include the relevant discussion in our investigation. The argument of the authors of \\cite{Turok.sing,Turok.simplified} about maximum amount of entropy possible in de Sitter spacetime, may lead our model to have a maximum cycle number between $10^{20}$ and $10^{30}$. However, the idea of the causal patch \\cite{Turok.sing} is probably the best way of handling the entropy problem so forth, and there are some interesting recent works on the subject which give a boost on cyclic cosmology \\cite{entropy}.  Let us close this discussion section with some comments on the role of the brane tensions and of the bulk cosmological constant in our model. As can be numerically confirmed, setting them to zero makes it almost impossible to satisfy the boundary conditions obtaining $H^2<0$ and singularity absence in [0,1] (this can be achieved only through a careful fine-tuning since our random choice procedure gives an one-digit number of such solutions in $10^6$ parameter multiplets). On the other hand, as we showed in \\ref{B}, in the case where $\\Lambda$, $\\lambda_0$ and $\\lambda_1$ are the only non-zero parameters, an $\\approx10^{-1}\\%$ of the random parameter choices, or $\\approx6\\%$ of the solutions, correspond to $H^2<0$. In mathematical terms, $\\Lambda$, $\\lambda_0$ and $\\lambda_1$ are requisite in order to acquire a solution with $H^2<0$ in the full dynamics, in a natural and not in a fine-tuning way. In terms of physics, it is the dark energy that lies in the background of the oscillatory mechanism and allows for cyclicity to realize. Adding the fact that it determines the cycles period and the maximum scale factor value, we conclude that dark energy is crucial in the described model. This brings it closer to the ekpyrotic paradigm of the literature.  In this work we examine general braneworld models and we show that cyclic behavior can naturally arise from the full $5D$ dynamics. One important feature is that brane collisions are not required, on the contrary the branes remain stable, and the cyclicity takes place on the $4D$ geometry not on the extra dimension. Another significant result is that the smallness of the cosmological constant of the observable universe pushes the cyclic period and the scale factor to astronomical large values, an essential requirement for the establishment of cyclic cosmology as a realistic alternative paradigm. Furthermore, we indicate the possibility of a ``chaotic cyclicity\", that is extremely large, non-periodic, cyclic intervals followed by extremely large intervals of conventional evolution and vice versa. After these, the model shares both the advantages and disadvantages of cyclic cosmology.\\\\   \\paragraph*{{\\bf{Acknowledgements:}}} The author  acknowledges partial financial support through the research program ``Pythagoras'' of the EPEAEK II (European Union and the Greek Ministry of Education)."
    },
    "0710/0710.5096_arXiv.txt": {
        "abstract": "We  investigate various galaxy occupation  statistics of dark matter halos using  a large galaxy group catalogue  constructed from the Sloan Digital Sky Survey Data Release 4 (SDSS DR4) with an adaptive halo-based group finder.   The  conditional  luminosity  function (CLF),  which  describes  the luminosity  distribution of galaxies  in halos  of a  given mass,  is measured separately for all, red and blue galaxies,  as well as in terms of central and satellite galaxies.  The  CLFs for central and satellite  galaxies can be well modelled  with  a  log-normal  distribution  and a  modified  Schechter  form, respectively.   About 85\\%  of  the central  galaxies  and about  80\\% of  the satellite  galaxies  in  halos  with  masses $M_h\\ga  10^{14}\\msunh$  are  red galaxies. These numbers decrease to 50\\% and 40\\%, respectively, in halos with $M_h \\sim 10^{12}\\msunh$.  For halos of  a given mass, the distribution of the luminosities of central galaxies, $L_c$, has a dispersion of about $0.15$ dex. The mean  luminosity (stellar mass) of  the central galaxies  scales with halo mass as $L_c \\propto M_h^{0.17}$ ($M_{*,c} \\propto M_h^{0.22}$) for halos with masses $M\\gg  10^{12.5}\\msunh$, and  both relations are  significantly steeper for less  massive halos.   We also measure  the luminosity (stellar  mass) gap between the first  and second brightest (most massive)  member galaxies, $\\log L_1  -  \\log  L_2$  ($\\log  M_{*,1}-\\log  M_{*,2}$).   These  gap  statistics, especially  in  halos with  $M_h  \\la  10^{14.0}  \\msunh$, indicate  that  the luminosities  of central  galaxies are  clearly distinct  from those  of their satellites.  The fraction of fossil groups, defined as those groups with $\\log L_1 -  \\log L_2\\ge  0.8$, ranges  from $\\sim 2.5\\%$  for groups  with $M_h\\sim 10^{14}\\msunh$ to 18-60\\% for groups with $M_h\\sim 10^{13}\\msunh$.  The number distribution of satellite galaxies in groups of a given mass follows a Poisson distribution,  in agreement  with  the occupation  statistics  of dark  matter sub-halos.  This provides strong support  for the standard lore that satellite galaxies reside in sub-halos.  Finally, we measure the fraction of satellites, which  changes  from $\\sim  5.0\\%$  for  galaxies  with $\\rmag\\sim  -22.0$  to $\\sim40\\%$ for galaxies with $\\rmag\\sim -17.0$. ",
        "introduction": "In recent  years, the halo occupation distribution  and conditional luminosity function have become  powerful statistical measures to probe  the link between galaxies  and their  hosting dark  matter halos.   Although  these statistical measures themselves do not give physical explanations of how galaxies form and evolve, they provide important  constraints on various physical processes that govern  the  formation  and  evolution  of  galaxies,  such  as  gravitational instability,  gas  cooling,  star  formation,  merging,  tidal  stripping  and heating, and  a variety of  feedback processes. In particular,  they constrain how their efficiencies scale with halo mass.  The halo occupation distribution (hereafter  HOD), $P(N \\vert M)$, which gives the probability of finding $N$  galaxies (with some specified properties) in a halo of mass  $M$, has been extensively used to  study the galaxy distribution in dark matter halos and galaxy  clustering on large scales (e.g.  Jing, Mo \\& B\\\"orner 1998;  Peacock \\&  Smith 2000; Seljak  2000; Scoccimarro  \\etal 2001; Jing, B\\\"orner  \\& Suto 2002; Berlind  \\& Weinberg 2002;  Bullock, Wechsler \\& Somerville  2002; Scranton  2002; Kang  \\etal 2002;  Marinoni \\&  Hudson 2002; Zheng \\etal 2002;  Magliocchetti \\& Porciani 2003; Berlind  \\etal 2003; Zehavi \\etal  2004, 2005;  Zheng \\etal  2005;  Tinker \\etal  2005).  The  conditional luminosity function (CLF),  $\\Phi(L \\vert M) {\\rm d}L$,  which refines the HOD statistic by considering the average number of galaxies with luminosity $L \\pm {\\rm  d}L/2$ that reside  in a  halo of  mass $M$,  has also  been extensively investigated (Yang, Mo \\& van den Bosch  2003; van den Bosch, Yang \\& Mo 2003; Vale \\& Ostriker 2004, 2006; Cooray 2006; van den Bosch et al.  2007a) and has been  applied to  various  redshift surveys,  such  as the  2dFGRS, the  Sloan Digital Sky Survey  (SDSS) and DEEP2 (e.g.  Yan, Madgwick  \\& White 2003; Yang \\etal 2004; Mo et al. 2004; Wang  \\etal 2004; Zehavi \\etal 2005; Yan, White \\& Coil  2004).   These  investigations  demonstrate  that  the  halo  occupation statistics  are very  useful  in establishing  and  describing the  connection between galaxies and  dark matter halos. Furthermore, they  also indicate that the galaxy/dark halo connection can provide important constraints on cosmology (e.g.,van den  Bosch, Mo \\& Yang  2003; Zheng \\& Weinberg  2007). Finally, the HOD/CLF framework also  allows one to split the  galaxy population in centrals and  satellites, and  to describe  their properties  separately  (e.g.  Cooray 2005; White \\etal 2007; Zheng \\etal 2007).  As has been pointed out in Yang et al.  (2005c; hereafter Y05c), a shortcoming of  the  HOD/CLF  models  is   that  the  results  are  not  completely  model independent. Typically,  assumptions have to be made  regarding the functional form of either $P(N \\vert M)$  or $\\Phi(L \\vert M)$.  Moreover, in all HOD/CLF studies  to date,  the occupation  distributions  have been  determined in  an indirect  way:  the  free  parameters  of  the  assumed  functional  form  are constrained  using {\\it  statistical}  data on  the  abundance and  clustering properties  of  the  galaxy  population.   One may  hope  to  circumvent  this shortcoming by directly measure  the dark matter distribution around galaxies. Such measurements  can in principle be obtained  through gravitational lensing and X-ray observations.  However, both methods are hampered by requirements on the data quality and uncertainties  related to the interpretation of the data. For instance, weak lensing  measurements, which requires high-quality imaging, typically needs  to resort to the stacking  of many lens galaxies  in order to get  a   detectable  signal,  but  this  stacking   severely  complicates  the interpretation in terms of the halo  masses of the lens galaxies.  In the case of X-ray  observations, robust  constraints can only  be obtained  for massive clusters, but even here the interpretation  of the data can be complicated due to the  presence of substructure and deviations  from hydrostatic equilibrium. An alternative method  to directly probe the galaxy -  dark halo connection is to use galaxy groups as a representation of dark matter halos and to study how the galaxy population  changes with the properties of  the groups (e.g., Y05c; Zandivarez et al. 2006; Robotham et al. 2006; Hansen \\etal 2007).  Recently,  we have constructed  a large  galaxy group  catalogue based  on the Sloan  Digital  Sky  Survey Data  Release  4  (SDSS  DR4), using  an  adaptive halo-based group finder  (Yang \\etal 2007; Paper I;  Y07 hereafter).  Detailed tests with  mock galaxy catalogues have  shown that this group  finder is very successful  in associating  galaxies  according to  their  common dark  matter halos.  In  particular, the group finder  performs reliably not  only for rich systems, but also for poor systems, including isolated central galaxies in low mass halos.  This makes  it possible to  study the galaxy-halo  connection for systems  covering a  large  dynamic range  in  masses.  Various  observational selection effects  have been  taken into account,  especially the  survey edge effects and  fiber collisions.  The halo  masses for the  groups are estimated according to  the abundance match,  using the characteristic  group luminosity and stellar masses (see \\S\\ref{sec_data} below).  According to tests with mock galaxy catalogues, the halo masses  are estimated with a standard deviation of about 0.3 dex.   With these well-defined galaxy group  catalogues, one can not only study  the properties of galaxies  in different groups  (e.g.  Y05c; Yang \\etal 2005d; Collister \\& Lahav 2005; van den Bosch \\etal 2005; Robotham \\etal 2006;  Zandivarez \\etal  2006; Weinmann  \\etal  2006a,b; van  den Bosch  \\etal 2007b; McIntosh  \\etal 2007), but also  probe how dark matter  halos trace the large-scale  structure of  the Universe  (e.g.  Yang  \\etal 2005b,  2006; Coil \\etal 2006;  Berlind \\etal 2007; Wang  et al.  2007 in  preparation).  In this paper, which is the second in a series, we use the SDSS DR4 group catalogue to probe  various  occupation  statistics  and  measure the  CLFs  for  different populations of galaxies.  This paper is organized as  follows: In Section~\\ref{sec_data} we describe the data (galaxy and group catalogues) used in this paper.  Section~\\ref{sec_CLFs} presents  our  measurement  of  the  CLFs  for all,  red  and  blue  galaxies. Sections~\\ref{sec_central}, ~\\ref{sec_HOD} and ~\\ref{sec_satfrac} describe the properties  of  central  galaxies,  the  halo occupation  statistics  and  the fraction  of  satellite galaxies,  respectively.   Finally,  we summarize  our results  in  Section~\\ref{sec_summary}.   Throughout  this  paper,  we  use  a $\\Lambda$CDM `concordance' cosmology whose  parameters are consistent with the three-year   data  release   of  the   WMAP  mission:   $\\Omega_m   =  0.238$, $\\Omega_{\\Lambda}=0.762$,  $n_s=0.951$, $h=0.73$ and  $\\sigma_8=0.75$ (Spergel et al. 2007).  \\begin{figure} \\plotone{f1.eps} \\caption{The color-magnitude relation for galaxies in our group sample.  The open circles indicate the Gaussian  peaks of the bi-normal distribution of galaxies in each luminosity bin. The solid dots indicate the corresponding averages  of the  two  Gaussian peaks.   The  solid line  is the  best-fit quadratic relation  to these averages (see  eq.~[\\ref{quadfit}]), which we use  to split  the  galaxies  into red  and  blue population  (color-coded accordingly).} \\label{fig:data_color} \\end{figure}  \\begin{figure*} \\plotone{f2.eps} \\caption{The conditional  luminosity functions (CLFs) of  galaxies in groups of different  mass bins.   Symbols correspond to  the CLFs  obtained using $M_L$  as   halo  mass  (estimated   according  to  the  ranking   of  the characteristic group  luminosity), with solid and  open circles indicating the contributions  from central and satellite  galaxies, respectively. The errorbars  reflect  the 1-$\\sigma$  scatter  obtained  from 200  bootstrap samples.  The solid lines  indicate the related best-fit parameterizations using  equation~[\\ref{eq:CLF_fit}].  For  comparison, we  also  show, with dashed  lines, the  CLFs  obtained  using $M_S$  as  halo mass  (estimated according  to the  ranking of  the group's  characteristic  stellar mass). Results shown in this plot are obtained from Sample II. } \\label{fig:CLF} \\end{figure*} \\begin{figure*} \\plotone{f3.eps} \\caption{Similar to  Fig.~\\ref{fig:CLF}, but here  we show the CLFs  for red (dashed  lines) and  blue (dotted  lines)  galaxies, for groups with  halo masses $M_L$. The solid  lines indicate the best-fit parameterizations for the  CLFs  of red  galaxies.   In both  cases  the  central and  satellite components  of  the  CLFs  are  indicated separately.   For  clarity,  the errorbars, again obtained using 200  bootstrap samples, are only shown for the red galaxies.  } \\label{fig:CLF_color} \\end{figure*}   \\begin{figure*} \\plotone{f4.eps} \\caption{The    best   fit    parameters   ($\\phi_s^{\\star}$    upper   row, $\\alpha_s^{\\star}$ second row, $L_c$ third row, and $\\sigma_c$ bottom row) to  the CLFs  shown  in Figs.~\\ref{fig:CLF}  and ~\\ref{fig:CLF_color},  as functions of  halo mass.  Panels  on the left,  in the middle, and  on the right show results for all, red, and blue galaxies, respectively. Since we have two  different halo  mass estimators ($M_L$  and $M_S$) and  two main group  samples (II  and III),  we have  obtained CLFs  for  four different combinations of sample and group mass estimator.  The results for all four combinations  are  shown  using  different  symbols  and  line-styles,  as indicated.  The errorbars  in the  first two  and last  rows  indicate the 1-$\\sigma$  variances obtained  from our  200 bootstrap  samples.   In the third row of  panels, however, the errorbars correspond  to the log-normal scatter, $\\sigma_c$, shown  in the bottom row of  panels.  For clarity the errorbars are only shown for the `$M_L$-Sample II' case, but they are very similar for the other four cases.  } \\label{fig:fit_CLF} \\end{figure*} ",
        "conclusions": "\\label{sec_summary}  Using a large galaxy group catalogue  constructed from the SDSS Data Release 4 (DR4)  by Y07,  we have  investigated  various halo  occupation statistics  of galaxies. In particular,  we have split the galaxy population  in red and blue galaxies,  and in  centrals  and satellites,  and  determined the  conditional luminosity  functions of  these  varies subsamples.   We  have also  presented luminosity gap  statistics, satellite  fractions, and halo  occupation numbers for  the galaxies  in our  group sample.  The main  results are  summarized as follows:  \\begin{enumerate}  \\item The conditional luminosity  functions for central and satellite galaxies can be well  modelled with a log-normal distribution  and a modified Schechter form, respectively.   The corresponding best fitting parameters  are listed in Table 1.  \\item The average scatter of the log-normal luminosity distribution of central galaxies decreases from $\\sim 0.15$ dex at the massive end ($\\log [M_h/\\msunh] \\ga 13.5$)  to $\\sim  0.1$ dex at  the low  mass end ($\\log  [M_h/\\msunh] \\sim 12.0$). However, due  to the method used to assign halo  masses to the groups, at the low mass end this should be considered a lower limit on the true amount of scatter.  \\item  The slope of  the relation  between the  average luminosity  of central galaxies (in the  $^{0.1}r$-band) and halo mass, ${\\rm  d}\\log L_c/{\\rm d}\\log M_h$, decreases  from $\\sim  0.68$ for $\\log  [M_h/\\msunh] \\ll 12.5$  to $\\sim 0.17$ for$\\log [M_h/\\msunh]  \\gg 12.5$. For the stellar  masses of the central galaxies we obtain that ${\\rm  d}\\log M_{*,c}/{\\rm d}\\log M_h$, decreases from $\\sim  1.83$  for  $\\log  [M_h/\\msunh]  \\ll  12.1$  to  $\\sim  0.22$  for$\\log [M_h/\\msunh] \\gg 12.1$.  \\item  The halo (group)  occupation numbers  of satellite  galaxies accurately follow Poisson  statistics. Since the  same applies to dark  matter sub-halos, this supports the standard picture that satellite galaxies are associated with dark matter sub-halos.  \\item In massive halos with  masses $M_h\\ga 10^{14}\\msunh$ roughly 85\\% (80\\%) of the central  (satellite) galaxies are red. These  red fractions decrease to 50\\% (40\\%) in halos with masses $M_h \\sim 10^{12}\\msunh$.  \\item By comparing  the scatter in the luminosities of  BCGs to the luminosity difference between the BCG and its  brightest satellite, we find that the BCGs form a  `special' subclass, in that  their luminosities can  not be considered the extreme  values of the distribution of  satellite luminosities, expecially in halos with masses $M_h \\la 10^{14.0} \\msunh$.  \\item  The fractions  of fossil  groups, which  are defined  as groups  a with luminosity gap  $\\log L_1 -  \\log L_2 \\ge  0.8$, decreases with  increasing of halo  mass from  18\\%-60\\% in  halos with  $M_h \\sim  10^{13}\\msunh$  to $\\sim 2.5\\%$ in halos with $M_h \\sim 10^{14}\\msunh$.  \\item The satellite  fractions obtained from our group  catalogue as functions of both  luminosity and stellar  mass (listed in Table~\\ref{tab:fsat})  are in good agreement with independent constraints from analyses of galaxy clustering and galaxy-galaxy lensing.  \\end{enumerate}   These results can be used  to constrain the various physical processes related to galaxy formation  and to interpret the various  statistics used to describe large scale structures (e.g.,  galaxy correlation functions, pairwise velocity dispersions, etc.).  Most of our  findings are in good agreement with previous studies (e.g.  Y05c,  Zandivarez et al.  2006; Robotham et  al.  2006) and can be linked to the semi-analytical modelling of galaxy formations (e.g., Kang et al. 2005; Zheng et al.  2005; Bower et al. 2006; Croton et al.  2006; De Lucia et al. 2007).   The luminosity and stellar  mass gap can be used  to probe the specific  formation properties  of central  galaxies (e.g.,  Vale  \\& Ostriker 2007). The fraction of the red  and blue populations for central and satellite galaxies can be  used to probe the color evolution  of satellite galaxies (ven den Bosch et al.  2007b)."
    },
    "0710/0710.3165_arXiv.txt": {
        "abstract": "Stellar-mass black holes are discovered in X-ray emitting binary systems, where their mass can be determined from the dynamics of their companion stars\\cite{rem06,cha06,oro03}.  Models of stellar evolution have difficulty producing black holes in close binaries with masses \\boldmath{$>10\\,M_{\\odot}$} (ref.\\ 4), which is consistent with the fact that the most massive stellar black holes known so far\\cite{cha06,oro03} all have masses within \\boldmath{$1\\sigma$} of \\boldmath{$10\\, M_{\\odot}$}. Here we report a mass of \\boldmath{$15.65 \\pm 1.45\\,M_{\\odot}$} for the black hole in the recently discovered system M33 X-7, which is located in the nearby galaxy Messier 33 (M33) and is the only known black hole that is in an eclipsing binary\\cite{pie06}. In order to produce such a massive black hole, the progenitor star must have retained much of its outer envelope until after helium fusion in the core was completed\\cite{bro01}.  On the other hand, in order for the black hole to be in its present 3.45 day orbit about its \\boldmath{$70.0 \\pm 6.9 M_{\\odot}$} companion, there must have been a ``common envelope'' phase of evolution in which a significant amount of mass was lost from the system\\cite{tau06}.  We find the common envelope phase could not have occured in M33 X-7 unless the amount of mass lost from the progenitor  during its evolution was an order of magnitude less than what is usually assumed in evolutionary models of massive stars\\cite{sch92,mey94,vaz07}. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.1998_arXiv.txt": {
        "abstract": "The concept of black hole entropy is one of the most important enigmas of theoretical physics. It relates thermodynamics to gravity and allows substantial hints toward a quantum theory of gravitation. Although Bekenstein conjecture --assuming the black hole entropy to be a measure of the number of precollapse configurations-- has proved to be extremely fruitful, a direct and conclusive proof is still missing. This article computes accurately the entropy evaporated by black holes in $(4+n)$ dimensions taking into account the exact greybody factors. This is a key process to constrain and understand the entropy of black holes as the final state is unambiguously defined. Those results allow to generalize Zurek's important argument, in favor of the Bekenstein conjecture, to multi-dimensional scenarios. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.3023_arXiv.txt": {
        "abstract": "Three-dimensional large eddy simulations of solar surface convection using realistic model physics are conducted. The thermal structure of convective motions into the upper radiative layers of the photosphere, the range of convection cell sizes, and the penetration depths of convection are investigated. A portion of the solar photosphere and the upper layers of the convection zone, a region extending $60\\times60$ Mm horizontally from 0 Mm down to 20 Mm below the visible surface, is considered. We start from a realistic initial model of the Sun with an equation of state and opacities of stellar matter. The equations of fully compressible radiation hydrodynamics with dynamical viscosity and gravity are solved. We use: 1) a high order conservative TVD scheme for the hydrodynamics, 2) the diffusion approximation for the radiative transfer, 3) dynamical viscosity from subgrid scale modeling. The simulations are conducted on a uniform horizontal grid of $600\\times600$, with 168 nonuniformly spaced vertical grid points, on 144 processors with distributed memory multiprocessors on supercomputer MVS-15000BM in the Computational Centre of the Russian Academy of Sciences. ",
        "introduction": "Convection near the solar surface has a strongly non-local and dynamical character. Hence, numerical simulations provide useful information on the spatial structures resulting from convection and help in constructing consistent models of the physical processes underlying the observed solar phenomena. We conduct an investigation of the temporal evolution and growth of convective modes on scales of mesogranulation and supergranulation in a three-dimensional computational box. In previous work by the author [\\citet{ustyugs06}] it was shown that collective motion of small convective cells of granulation expels weak magnetic field on the edges of cells at mesogranular scales. The average size of such cells is 15-20 Mm and the lifetime of order 8-10 solar hours. Simulation of solar photosphere convection [\\citet{stein06}] in a computational domain of size 48 Mm in the horizontal plane and 20 Mm in depth showed that the sizes of convective cells increase with depth. The purpose of this work is to investigate the development and scales of convection in a region of size 60 Mm in the horizontal plane and 20 Mm in depth. ",
        "conclusions": ""
    },
    "0710/0710.3037_arXiv.txt": {
        "abstract": "I present a new census of the stellar population in the Chamaeleon~I star-forming region. Using optical and near-IR photometry and followup spectroscopy, I have discovered 50 new members of Chamaeleon~I, expanding the census of known members to 226 objects. Fourteen of these new members have spectral types later than M6, which doubles the number of known members that are likely to be substellar. I have estimated extinctions, luminosities, and effective temperatures for the known members, used these data to construct an H-R diagram for the cluster, and inferred individual masses and ages with the theoretical evolutionary models of Baraffe and Chabrier. The distribution of isochronal ages indicates that star formation began 3-4 and 5-6~Myr ago in the southern and northern subclusters, respectively, and has continued to the present time at a declining rate. The IMF in Chamaeleon~I reaches a maximum at a mass of 0.1-0.15~$M_\\odot$, and thus closely resembles the IMFs in IC~348 and the Orion Nebula Cluster. In logarithmic units where the Salpeter slope is 1.35, the IMF is roughly flat in the substellar regime and shows no indication of reaching a minimum down to a completeness limit of 0.01~$M_\\odot$. The low-mass stars are more widely distributed than members at other masses in the northern subcluster, but this is not the case in the southern subcluster. Meanwhile, the brown dwarfs have the same spatial distribution as the stars out to a radius of $3\\arcdeg$ (8.5~pc) from the center of Chamaeleon~I. ",
        "introduction": "\\label{sec:intro}        The characteristics of the distributions of masses, ages, and positions in a newborn stellar population are determined by the process of star formation. As a result, measurements of these distributions in star-forming regions are potentially valuable for testing models of the birth of stars and brown dwarfs. For instance, the properties of the stellar initial mass function \\citep[IMF,][]{mey00} can constrain the relative importance of turbulent fragmentation \\citep{pn02}, gravitational fragmentation \\citep{lar85}, dynamical interactions \\citep{bon03}, and accretion and outflows \\citep{af96} in regulating the final masses of stars. Models of the star formation rates of molecular clouds (e.g., constant, accelerating, bursts) can be tested against the distributions of ages and positions of members of young clusters \\citep{fei96,pal97,har01}. The spatial distributions also provide insight into cloud fragmentation, binary formation, cluster dynamics, and the origin of brown dwarfs \\citep{lar95,hil98,har02,luh06tau}. To obtain measurements of this kind, one must identify the members of star-forming regions and estimate their masses and ages. Only a few young stellar populations have been characterized in detail, such as the Orion Nebula Cluster \\citep{hil97}, Taurus \\citep{kh95}, and IC~348 in Perseus \\citep{luh03ic}.  The Chamaeleon~I star-forming region is amenable to a thorough census of its stellar population for several reasons. It is among the nearest star-forming regions \\citep[$d=160$-170~pc,][]{whi97,wic98,ber99}, exhibits less extinction than many young clusters ($A_V\\lesssim5$), is compact enough that a large fraction of the region can be surveyed to great depth in a reasonable amount of time, and is sufficiently rich for a statistically significant analysis of its stellar population. In addition, because Chamaeleon~I is relatively isolated from other star-forming regions, confusion with other young populations is minimal. Previous surveys have already identified more than 150 young stars and brown dwarfs in Chamaeleon~I through photometric variability, H$\\alpha$ emission, X-ray emission, mid-infrared (IR) excess emission, and optical and near-IR color-magnitude diagrams \\citep[][references therein]{com04,luh04cha,luh07cha}. However, in the census of known members produced by these surveys, the completeness as a function of mass and position is unknown \\citep{luh04cha}. In this paper, I present a set of magnitude-limited surveys for members of Chamaeleon~I that have well-defined completeness limits (\\S~\\ref{sec:new}). I then use the new census of Chamaeleon~I to measure the star formation history (\\S~\\ref{sec:hr}), IMF (\\S~\\ref{sec:imf}), and spatial distribution of its stellar population (\\S~\\ref{sec:spatial}). ",
        "conclusions": "I have presented an extensive search for new members of the Chamaeleon~I star-forming region. Because the completeness limits of my survey are well-determined, I have been able to perform robust measurements of the distributions of members of Chamaeleon~I as a function of mass, position, and age. The primary results of this study are summarized as follows:  \\begin{enumerate}  \\item I have discovered 50 new members of Chamaeleon~I, which increases the census of known members to 226 objects. The new members include 14 objects that are later than M6 ($M\\lesssim0.08$~$M_\\odot$) and the two faintest known members of the cluster, which may have masses of only 0.005-0.01~$M_\\odot$. The current census now contains 28 members that are likely to be substellar.  \\item The distribution of isochronal ages for members of Chamaeleon~I between 0.1-1~$M_\\odot$ suggests that star formation has occurred for the past 3-4 and 5-6~Myr in the southern and northern subclusters, respectively, at rates that have declined with time.  \\item The IMF in Chamaeleon~I reaches a maximum at a mass of 0.1-0.15~$M_\\odot$, which is similar to the turnover mass observed in IC~348 and the Orion Nebula Cluster \\citep{hil97,hc00,mue02,mue03,luh03ic}. The substellar IMF is roughly flat in logarithmic units and shows no indication of reaching a minimum down to a completeness limit of 0.01~$M_\\odot$.  \\item Chamaeleon~I does not contain a widely-distributed population of brown dwarfs, which is contrary to the predictions of some embryo ejection models. Instead, the substellar members share the same spatial distribution as the stars. However, low-mass stars in the northern subcluster do appear to have a wider distribution than members at other masses, which resembles the mass segregation that has been previously observed in Orion and IC~348 \\citep{hc00,mue03}.  \\end{enumerate}"
    },
    "0710/0710.1518_arXiv.txt": {
        "abstract": "% After presenting three ways of defining a bulge component in disc galaxies, we introduce the various types of bulges, namely the classical bulges, the boxy/peanut bulges and the disc-like bulges. We then discuss three specific topics linked to bulge formation and evolution, namely the coupled time evolution of the bar, buckling and peanut strengths; the effect of velocity anisotropy on peanut formation; and bulge formation via bar destruction. ",
        "introduction": "Three ways of defining a bulge have been used so far, one morphological, the second photometrical and the third kinematical. Based on morphology, a bulge is a component of a disc galaxy that has a smooth light distribution that swells out of the central part of a disc viewed edge-on. This definition has the disadvantage of being applicable only to edge-on systems and the advantage of necessitating only an image of the galaxy. The second definition is based on photometry and defines a bulge as the extra light in the central part of the galaxy, above the exponential profile fitting the remaining (non central) part of the disc. In earlier papers this component was fitted with an $r^{1/4}$ law, while more recent ones use its generalisation to an $r^{1/n}$ law (S\\'ersic 1968). This definition has the advantage of being applicable to disc galaxies independent of their inclination. It has also the advantage of leading to quantitative results about the light distribution, but has the disadvantage of assigning to the bulge any extra central luminosity of the disc, independent of its origin. The third definition is based on kinematics, and in particular on the value of $V/\\sigma$, or more specifically on the location of the object on the ($V/\\sigma$, ellipticity) plot, which is often referred to as the Binney diagram (Binney 1978, 2005). This definition, potentially quite powerful, has unfortunately been very little used so far, due to the small number of galaxies for which the necessary data are available, a situation which is rapidly improving with large surveys, such as SAURON (Bacon \\tal 2001; de Zeeuw \\tal 2002; Peletier this volume).  The lack of a single, clear-cut definition of a bulge is due to the fact that disc galaxies are viewed in different orientations and also to the fact that not all types of data are available for all objects. Nevertheless, it has led to considerable confusion and to the fact that bulges are an inhomogeneous class of objects. Indeed, many different types of objects, with very different properties and formation histories are included in the general term `bulges'. To remedy this, Kormendy (1993) and Kormendy \\& Kennicutt (2004) distinguish classical bulges from pseudo-bulges, the latter category encompassing all bulges that are not classical. Athanassoula (2005a) argues that pseudo-bulges also are an inhomogeneous class of objects, and thus distinguishes three types of bulges.  {\\bf Classical bulges} are formed by gravitational collapse or hierarchical merging of smaller objects and corresponding dissipative gas processes. The material forming this bulge could be externally accreted, or could come from clumps in the proto-disc. In general, bulges of this type are formed before the actual disc (e.g. Steinmetz \\& M\\\"uller 1995; Noguchi 1998; Immeli \\tal 2004). Nevertheless, a bulge can also form from material externally accreted at much later stages (e.g. Pfenniger 1991; Athanassoula 1999; Aguerri, Balcells \\& Peletier 2001; Fu, Huang, Deng 2003). Their morphological, photometrical and kinematical properties are similar to those of ellipticals.  {\\bf Box/peanut bulges} (B/P) form from a vertical instability of the disc material. This has often been observed in $N$-body simulations of bar-unstable discs, where the initial stage of bar formation is followed by a puffing up of the inner parts of the bar (e.g. Combes \\& Sanders 1981; Combes \\tal 1990; Raha \\tal 1991; Athanassoula \\& Misiriotis 2002; Athanassoula 2003, 2005a; O'Neil \\& Dubinski 2003; Martinez-Valpuesta \\& Shlosman 2004; Debattista \\tal 2004, 2006; Martinez-Valpuesta \\tal 2006). Viewed side-on (i.e. edge-on with the line of sight along the bar minor axis), this structure protrudes from the disc and has a characteristic boxy or peanut shape whose size is of the order of a few disc scale-lengths. Thus, a box/peanut bulge is just {\\it part} of a bar seen side-on.  Finally {\\bf disc-like bulges} form from inflow of (mainly) gas material to the centre of the galaxy and subsequent star formation (e.g. Athanassoula 1992; Friedli \\& Benz 1993; Heller \\& Shlosman 1994; Wada \\& Habe 1995) . The torque exerted by the bar pushes gas, and to a lesser extent also stars, to the inner parts of the disc where they form an inner disc. Star formation can be very high there, due to the increased gas density. Thus the result of this process should be a central disc, or disc-like object, whose stellar component should include a sizable fraction of young stars and whose size should be less than, or of the order of a kpc. Such a component could harbour sub-structures such as spirals, or bars, as is indeed sometimes observed (Kormendy \\& Kennicutt 2004 and references therein). It is thus clear that disc-like bulges are very different objects from boxy/peanut bulges, since they are much smaller, have a different shape, different kinematics and provide a different type of excess on the radial photometric profiles. They also have different formation histories.  The different formation histories of these three types of objects lead to different properties -- morphological, photometrical and kinematical -- which in turn help classify observed bulges into one of the three above mentioned types. Nevertheless, as stressed by Athanassoula (2005a), different types of bulges often co-exist and it is possible to find all three types of bulges in the same simulation, or galaxy.  Realising the non-homogeneity of the objects classified as bulges and attempting to classify them is only the first step. Much more work is now necessary, particularly on two issues. The first one is the understanding of the formation and evolution of these types of objects. The second one is to predict the properties of these objects, starting from their formation scenarios. The latter is particularly important in order to bridge the gap between classification schemes based on formation histories and classification schemes based on observed properties. Here we make small contributions to both these issues, using $N$-body simulations. In Sect. 2 we present the time evolution of the bar, the buckling and the peanut strengths and their interplay. Sect. 3 discusses the velocity anisotropy and its link to the above strengths. Finally, in Sect. 4 we discuss the photometrical properties of a destroyed bar and boxy/peanut bulge. ",
        "conclusions": ""
    },
    "0710/0710.3347_arXiv.txt": {
        "abstract": "It has been suggested that fossil groups could be the canibalized remains of compact groups, that lost energy through tidal friction. However, in the nearby universe,  compact groups which are close to the merging phase and display a wealth of interacting features (such as HCG 31 and HCG 79) have very low velocity dispersions and poor neighborhoods, unlike the massive, cluster-like fossil groups studied to date. In fact, known z$=$0 compact groups are very seldom embedded  in massive enough structures which may have resembled the intergalactic medium of fossil groups.   In this paper we study the dynamical properties of CG6, a massive compact group at z$=$0.220 that has several properties in common with known fossil groups.  We report on new \\sloang~and \\sloani~imaging and multi-slit spectroscopic performed with GMOS on Gemini South. The system has 20  members, within a radius of 1 h$_{70}^{-1}$ Mpc, a velocity dispersion of 700~\\kms~and has a mass of 1.8 $\\times$ 10$^{14}$ h$_{70}^{-1}$~\\Msol, similar to that of the most massive fossil groups known. The merging of the four central galaxies in this group would form a galaxy with magnitude $M_{r'} \\sim -23.4$, typical for first-ranked galaxies of fossil groups. Although nearby compact groups with similar properties to CG 6 are rare, we speculate that such systems occurred more frequently in the past and they may have been the precursors of fossil groups. ",
        "introduction": "\\label{sec:intro}  Groups of galaxies are small systems of typically a few L$^{*}$ galaxies, which comprise over 55\\% of the nearby structures in the universe. A small fraction of galaxy groups are classified as {\\it compact} groups, which are responsible for $\\sim$ 1\\% of the luminosity density of the universe \\citep{mdoh91}. Although they are rare objects in the nearby universe, their high galactic densities and low velocity dispersions make them ideal systems for the study of galaxy transformation through galaxy-galaxy collisions.  As expected, these systems have a high fraction of interacting members, although merged objects are rare \\citep{zepf93}.  They are commonly believed to evolve through dynamical friction and finally merge to form one single galaxy \\citep{Barnes92}. \\citet{vikhlinin99} and \\citet{jones03} have suggested that the merging of compact groups can lead to the formation of  fossil groups. A fossil group (FG, hereafter) is a system with an extended and luminous X-ray halo (L$_X$ $>$ $10^{42}$ h$_{70}^{-2}$ erg s$^{-1}$), dominated by one single brighter than L$^{*}$ elliptical galaxy, surrounded by low-luminosity companions \\citep[where the difference in magnitude between the bright dominant elliptical and the next  brightest companion is $>$ 2 mag in the R-band;][]{jones03}.  One important goal of this article is to investigate if compact groups (CG, hereafter) as we known them today, could be the precursors of FGs. In order to answer this question, we summarize, in section 2, the properties of a few of the most strongly interacting nearby CGs known, which are about to merge, and in section 3, we describe the properties of the five FGs which have been studied spectroscopically so far. In section 4, we present new observations for a CG embedded in a cluster-size potential, at redshift z$=$0.22 and section 5 puts together all the observations described and discusses the CG-FG scenario. Throughout this paper we adopt when necessary a standard cosmological model: $H_{0}=70\\,h_{70}\\,$km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{m}=0.3$ and $\\Omega_{\\Lambda}=0.7$. At z$=$0.22, 1\\arcsec~corresponds to $ 3.5\\, h_{70}^{-1}\\,$kpc.  \\section {Interacting compact groups: HCG 31, HCG 79 and HCG 92} \\label{sec:hcg}  There is evidence from both observations and simulations that groups evolve through dynamical friction and coalesce to form more compact structures as the universe ages. A few of the most compact, and therefore most evolved groups known, from Hickson's catalogue \\citep{h92}  are HCG 31, HCG 79 (or Seyfert Sextet) and HCG 92 (or Stephan's quintet).  The study of these groups is important to help understanding processes common in merging systems, environments that may have occurred more often in the high-redshift universe.  HCG 31 is a group at z$\\sim$0.013 and with a velocity dispersion of $\\sigma$ $\\sim$ 60 \\kms. This is a gas-rich group with intense star forming activity \\cite[e.g. ][]{moetal06, amram07}, dominated by a central pair of interacting dwarf galaxies A$+$C. HCG 31 is thought to be in a pre-merging phase \\citep{amram05,vm05} and it has well developed tidal tails seen in H$\\alpha$ and HI. The group hosts two excellent candidates for tidal dwarf galaxies, namely member F, in the southern tail and member R, 50 h$_{70}^{-1}$ kpc to the north of the group (for an assumed distance modulus of DM=33.8).  HCG 79, also known as {\\it Seyfert Sextet}, was originally identified as a sextet of galaxies but it is now known to be a quartet at z$=$ 0.0145 (the 5th object is in the background and the 6th is a luminous tidal debris to the northeast of the group).  This is the most CG in Hickson's catalogue with a galaxy-galaxy distance below 10 kpc (for an adopted DM=34.0) and a velocity dispersion of $\\sigma$ = 138 \\kms. The four galaxies present morphological distortions and increased activity (tidal debris, bar in HCG 79B, dust lane in HCG 79A, radio and infrared emission, disturbed rotation curves and nuclear activity).  The group presents a prominent intra-group light envelope which contains 45\\% of the total light of the group \\citep{darocha05}  and irregular envelopes of HI \\citep{wil91} and X--rays \\citep{pil95}. These suggest that recent or on-going interaction is taking place within this system.  HCG 92, also known as Stephan's quintet, is in reality a  quartet with z$=$ 0.0215  and a foreground galaxy. It is the most well studied  CG -- multi wavelength data are available from radio to X-rays. Most of the gaseous material in Stephan's quintet is concentrated not in the galaxies but in the intragroup medium, suggesting that collisions among group members may have happened frequently. A number of tidal dwarf galaxies and intergalactic HII regions have been identified in this group \\citep[e.g.][]{moetal01,moetal04,xu05}. Of the three groups described above, HCG 92 is the only one to have detected X-rays, with a total bolometric luminosity of  2.96$\\times$10$^{42}$ h$^{-2}_{70}$ \\ergs~\\citep{Xue2000}.  These three spiral-rich groups are thought to be in their final stages of evolution -- they are, in fact, some of the most compact systems found in the Hickson's catalogue. Yet, they have members that can be clearly identified as individual galaxies, suggesting that once merging starts, it may proceed quickly, and the groups may no longer be recognized as such. The bright members of these  groups will almost certainly end up as a single galaxy pile. A discussion of whether these systems will most likely end up as FGs or as single isolated elliptical galaxies is deferred to section 5. In the following section, some of the optical properties of the FGs studied so far are summarized. ",
        "conclusions": "Dynamical friction and subsequent merging are probably the processes responsible for the lack of L$^{*}$ galaxies in FGs. Considering the merging scenario, it is possible that the overluminous central galaxy in a FG has been formed within a substructure, inside a larger structure.  In that case, one could think of a scenario where a CG was formed within a rich group, which would then have merged, leaving  behind the brightest elliptical galaxy of what today is seen as a FG. One weak argument against this scenario is that the nearby examples of CGs are not usually found within such massive structures, but instead are more often surrounded by very sparse structures. There are, however, examples such as CG6, surrounded by large numbers of lower-luminosity galaxies, inhabiting a deep potential well.  We would like to test the hypothesis that CGs, as observed in the nearby universe, could be the precursors of FGs. We may examine two aspects: (1) if the sum of the  luminosities of the brightest CG galaxies is similar to the luminosity of a first-ranked FG galaxy and; (2) if the  neighbourhoods of CGs are rich, i.e., if the system as a whole  (group plus environment) has a velocity dispersion/mass similar to that of a FG.  We compute the total luminosity of the galaxies in the soon-to-merge CGs,  HCG 31, HCG 79 and HCG 92, to check how these compare with the luminosities of first-ranked galaxies in FGs.  Adding up the luminosities of galaxies HCG 31 A to C, G and Q, which are the brightest in the group HCG 31, a magnitude of M$_R = -22.5$ is obtained (for a distance modulus, DM = 33.8).  Summing up the luminosities of galaxies HCG 79 A-D, an equal total magnitude of M$_R=-22.5$ is obtained (for DM = 34.0). These are upper limits on the luminosities of these objects given that several members are starburst galaxies. After fading, the merged central object in HCG 31 and HCG 79 will have somewhat lower magnitudes than that of a typical first-ranked galaxy in a FG.  Fossil groups first-ranked galaxies have luminosities well above L$^{*}$. For the five FGs studied by \\citet{jones03}, the first-ranked galaxies had a median luminosity of $M_R=-23.2$ and for the 34 FGs found in the SDSS DR5 by \\citet{dosSantos07} the median luminosity was $M_R=-23.5$.  Although for HCG 92, the final object (adding up luminosities of galaxies A-E) would have an absolute magnitude of $M_R=-24.2$ (for DM$=$34.8), which after allowing for some fading, could be similar to that of an FG first-ranked galaxy, HCG 92 would possibly still not resemble an FG when merged, because its neighbourhood is very sparse, i.e., it is not embedded in any larger structure, as it is often the case for the central galaxy in FGs. This is in agreement with its relatively low bolometric X-ray luminosity of 2.96 $\\times$ 10$^{42}$ h$_{70}^{-2}$ ergs s$^{-1}$ \\citep{Xue2000}.  The environments of nearby CGs have been surveyed by \\citet{ribeiro98,zabludoff98,carrasco06} among others. Spectroscopy of dozens of members in the neighbourhood of quite a number of groups was obtained, confirming in all cases that CGs have low velocity dispersions typical of the group regime (typically 200-300 km s$^{-1}$).  In fact, even for HCG 62, thought to be one of the most massive CGs in Hickson's catalogue, the velocity dispersion obtained from 45 members of the system showed that it is a bonafide group (376 km s$^{-1}$). HCG 62 was suggested by \\citet{ponman93} as an example of a system that could turn into a FG in a few Gygayears, but its velocity dispersion is still much lower than the value of $\\sim$ 600 km s$^{-1}$, typical for rich FGs. Two other massive nearby CGs in Hickson's catalogue are HCG 94 and HCG 65.  The first is known to have a very high bolometric X-ray luminosity of 2.35 $\\times$ 10$^{44}$ h$_{70}^{-2}$ \\citep{Xue2000} which may, however, be contaminated by the emission of a nearby cluster. The velocity dispersion obtained from 11 members in this system gives a value of  479 km s$^{-1}$. HCG 65 is the center of the cluster Abell 3559. It is in the heart of the Shapley supercluster and its location makes it hard to disentangle its dynamics and determine its mass. The three most massive Hickson CGs known, HCG 62, HCG 65 and HCG 94, are strongly early-type dominated, as expected from the velocity dispersion-morphology relation observed for CGs.  The conclusion is then that a typical CG, as observed at z=0, is unlike to turn into a FG. It is more likely to  merge into an isolated elliptical galaxy.  For the compact group CG 6, at $z=0.22$ and $\\sigma =$703 km s$^{-1}$,  if we merge the four central galaxies (A, B, C and D), we end up with a  galaxy with total magnitude \\sloani$=$16.31  and \\sloang$=$17.80 mag (M$_{i'}$=--23.87 and M$_{g'}$=--22.38, with no k-correction). Using the color relation for a galaxy at z$=$0.2 from \\cite{fuk95},  the magnitude in r$^{'}$ will be 16.83 or M$_{r}=-23.35$. This magnitude is similar to those of typical central galaxies in FGs and the velocity dispersion of the system is typical for the  studied FGs. However, no gap of at least two magnitudes between the first-ranked  relic and the remaining objects of the system would be observed because there is at least one other bright galaxy in the system within half the  virial radius of the group.  We point out that one other example of a possible massive system, at z=0.39, which may turn into a FG, has recently  been discovered by \\cite{rines07}. Spectroscopic  studies of CGs at medium redshifts may find many more of such objects."
    },
    "0710/0710.1497_arXiv.txt": {
        "abstract": "We present the results of an analysis of a well-selected sample of galaxies with active and inactive galactic nuclei from the Sloan Digital Sky Survey, in the range $0.01 < z < 0.16$. The SDSS galaxy catalogue was split into two classes of active galaxies, Type~2 AGN and composites, and one set of inactive, star-forming/passive galaxies. For each active galaxy, two inactive control galaxies were selected by matching redshift, absolute magnitude, inclination, and radius. The sample of inactive galaxies naturally divides into a red and a blue sequence, while the vast majority of AGN hosts occur along the red sequence. In terms of H$\\alpha$ equivalent width, the population of composite galaxies peaks in the valley between the two modes, suggesting a transition population. However, this effect is not observed in other properties such as colour-magnitude space, or colour-concentration plane.  Active galaxies are seen to be generally bulge-dominated systems, but with enhanced H$\\alpha$ emission compared to inactive red-sequence galaxies.  AGN and composites also occur in less dense environments than inactive red-sequence galaxies, implying that the fuelling of AGN is more restricted in high-density environments. These results are therefore inconsistent with theories in which AGN host galaxies are a `transition' population.  We also introduce a systematic 3D spectroscopic imaging survey, to quantify and compare the gaseous and stellar kinematics of a well-selected, distance-limited sample of up to 20 nearby Seyfert galaxies, and 20 inactive control galaxies with well-matched optical properties. The survey aims to search for dynamical triggers of nuclear activity and address outstanding controversies in optical/IR imaging surveys. ",
        "introduction": "Active Galactic Nuclei (AGN) have long been considered a curiosity in their own right \\citep{schmidt63, lybell69,blan74}, but are now recognised to be integral to galaxy formation and evolution. At early cosmological epochs, gas-rich galaxies are thought to form and collide in violent mergers \\citep{mihos96}, triggering vast bursts of star formation, and fuelling supermassive black holes in their cores \\citep{kriek06}.  Recent simulations of isolated and merging galaxies by \\cite{spring05} have incorporated feedback from star formation and black-hole accretion, and find that once an accreting supermassive black hole (SMBH) has grown to some critical size, the AGN feedback terminates its growth as a large fraction of the remaining nuclear gas is driven out by the powerful quasar.  In the current epoch, the peak of the quasar era is over and the galaxy merger rate has declined \\citep{struck97}.  Nevertheless, at least 20\\% of today's galaxies show scaled-down quasar activity in their centres, and direct measurements of active and inactive galaxy dynamics have revealed a tight correlation between the central black-hole mass and host galaxy stellar velocity dispersion or bulge mass. This $M_{\\bullet}-\\sigma$ relation \\citep{geb00,merr01} points to an intimate link between host galaxy evolution and central black-hole growth, suggesting that all bulge-dominated galaxies today harbour dead quasars and that a lack of nuclear activity cannot be attributed to the absence of a central black hole.  Therefore, given the ubiquity of supermassive black holes, what determines the degree of nuclear activity in todays galaxies and what role is played by the host galaxy in triggering and fuelling their dormant black holes remain open issues.  Previously, studies of nearby active galaxies were based on small galaxy samples, like the C$f$A Seyfert Sample \\citep{huchra92}, the 12 Micron Active Galaxy Sample \\citep{rush93} and the \\cite{ho95} galaxy sample of approximately 486 galaxies covering a range of activity types. Now, however, the standard has been set by the Sloan Digital Sky Survey (SDSS), from which a range of useful samples of both active and inactive galaxies can be selected in a well-defined, uniform way.  The SDSS \\citep{york00} provides, for the first time, 5-band photometry and spectroscopy of many thousands of low redshift AGN \\citep{kauf03c,hao05a}, enabling constraints to be placed on galaxy and AGN evolution over a wide range of galaxy masses, and acts as the definitive supporting data to all detailed follow-up galaxy studies. Initial investigations into the growth and evolution of black holes using these databases have yielded contrasting results --- the origin of the correlation between galaxy bulge and central black-hole masses is hotly debated in particular.  \\citet{heck04} find that the majority of present-day accretion occurs onto 10$^{8}$ solar-mass black holes in moderate mass galaxies, suggesting that bulge and black-hole evolution is still tightly coupled today, and that the evolution of AGN luminosity functions is driven by a decrease in the mass scale of accreting black holes. In contrast, the \\citet{hao05b} study of about 3\\,000 SDSS AGN concludes that evolution in AGN luminosity functions is driven by evolution in the Eddington ratio, rather than black-hole mass. Meanwhile, \\cite{grup04} and \\cite{grup05} use a sample of 75 X-ray selected AGN to argue that black holes in Narrow Line Seyfert 1 galaxies in particular, grow by accretion in well-formed bulges to produce the $M_{\\bullet}-\\sigma$ relation over time, refuting theories for the origin of black-hole / bulge relations in which black-hole mass is a constant fraction of bulge mass at all epochs or in which bulge growth is controlled by AGN feedback \\citep{king03,ferr02,kriek06}. The subclass of NLS1s has been further explored more recently by \\cite{komo07}, who find that NLS1s are accreting at a rate higher than the Eddington rate, confirming that their BHs must be growing. They suggest that either NLS1 galaxies evolve into Broad Line Seyfert 1 (BLS1) galaxies with respect to their black hole mass distribution, which would require some change in the bulge properties, possibly due to feedback, or that NLS1s are just low-mass extensions of BLS1 galaxies, and the high accretion rate could just be caused by a relatively short-lived accretion phase.  Another characteristic revealed by large automated surveys such as the SDSS, is that the galaxy population is found to be bimodal in colour \\citep{lilly95,strat01}, with it now being more natural to describe a galaxy as being on the ``red sequence'' or ``blue sequence'', rather than being ``early type'' or ``late type'' \\citep{bald06}.  A key goal of galaxy evolution theory is then to explain the colour bimodality of galaxies, the relationships within each sequence, and where active hosts fit into this picture. Therefore, the understanding of galaxy formation and evolution processes necessitates full inclusion of AGN and their hosts.  Recently efforts have been made to try to understand where AGN host galaxies fall in colour-magnitude space. The red sequence consists of mainly massive, passively evolving galaxies, while the majority of galaxies show blue colours, attributed to ongoing star formation. The two sequences are separated by a relatively narrow valley in colour space.  The emerging consensus is that intense star formation is fuelled by galaxy mergers at high redshift, which forms massive bulges, and at some point the star formation ceases, resulting in the galaxy migrating from the blue sequence to the red sequence. AGN have been singled out as the mechanism for star-formation quenching, leading to suggestions that AGN should occupy a distinct, `transition' region, of the colour-magnitude diagram (CMD). However this does not necessarily have to be the case, as the resulting galaxy could just be bluer due to increased star-burst activity, or redder due to enhanced dust (e.g., Luminous Infra-Red Galaxies), or some combination of the original colours of the merging galaxies. Young stellar populations and nuclear starbursts are therefore also important components in the dynamics of nuclear activty \\citep{gonza98,sarzi07}.  \\cite{nand07} studied the colour-magnitude relation for a sample of 50 X-ray selected AGN from the AEGIS survey (All-wavelength Extended Groth strip International Survey), in the range $0.6 < z < 1.4$. They conclude that AGN fall on the red sequence or on the red edge of the blue sequence, with many in between these two modes.  \\cite{mart07} further explored the idea of a `transition' region by exploiting the separation of the blue and red sequences in the ${\\rm UV} - r$ colour-magnitude diagram. Using a sample of UV selected galaxies from the GALEX Medium Imaging Survey, along with SDSS data, they explored the nature of galaxies in the transition zone. The AGN fraction of their sample was found to peak in the transition zone, and they also found circumstantial evidence that star-formation quenching rates were higher in higher luminosity AGN.  Higher image quality than the SDSS is provided by the Millennium Galaxy Catalogue (MGC; \\citealt{liske03,cross04}), which has obtained redshifts for 10\\,095 galaxies to $B < 20 {\\rm\\,mag}$, covering 37 deg$^{2}$ of equatorial sky. The MGC study of galaxy bimodality was in the colour-structure plane, (\\citealt{drive06}; hereafter D06) and expressed a contrasting theory for the bimodality.  The MGC survey has sufficient resolution to achieve reliable bulge-disc decomposition \\citep{allen06}, and in separating out the two components, D06 claim that galaxy bimodality is caused by the bulge-disc nature of galaxies, and not by two distinct galaxy classes at different evolutionary stages. In this case, they show that the bulge-dominated, early-type galaxies populate one peak and the bulge-less, late-type galaxies occupy the second. The early- and mid-type spirals sprawl across and between the peaks. They also propose that the reason for the dual structure of galaxies is due to galaxy formation proceeding in two stages; first, there is an initial collapse phase, which forms the centrally concentrated core and a black hole, and second, there is the formation of a planar rotating disc caused by accretion of external material building up the galaxy disk.  In this paper, we present the properties of a magnitude-limited sample of ``active'' and ``inactive'' galaxies carefully selected from the SDSS, which also acts as the parent sample for detailed followup studies using the IMACS-IFU on the Magellan 6.5m telescope. Our sample selection is described in \\S 3, the sample properties in \\S 4, and implications of these properties in \\S 5. We conclude by introducing the Magellan survey, which is now underway. ",
        "conclusions": "We have carried out a robust classification of `active' and `inactive' galaxies in the SDSS, based on their emission line properties and locations in the BPT diagram. We have compared the active galaxies with well selected control galaxies matched in redshift, absolute magnitude, aspect ratio and radius. We find that:  \\begin{itemize} \\item Type~2 AGN host galaxies occur mainly on the red sequence of the CMD, but show increased levels of H$\\alpha$ flux in their emission compared to inactive red-sequence galaxies (Fig.~\\ref{fig:H-alpha}(a) and Fig.~\\ref{fig:CMD}(a)); \\item A separate class of composite galaxies appears to peak on the blue edge of the red sequence on the CMD, whereas the peak of the H$\\alpha$ distribution places composite galaxies firmly in the valley between the blue and red sequences (Fig.~\\ref{fig:H-alpha}(b) and Fig.~\\ref{fig:CMD}(b)); \\item Colour-concentration relations, however, show a more complex, possibly double, morphology in the peak of the composite distribution, rather than being in a valley (Fig.~\\ref{fig:colour-conc-comp}); \\item AGN (and composites) are found in less dense environments on average then matched inactive red-sequence galaxies. The more clustered inactive galaxies are likely to be satellite galaxies in high-density environments (Fig.~\\ref{fig:mean-enviro} and Fig.~\\ref{fig:distrib-enviro}). \\end{itemize}  The key to understanding this in more detail now lies in the dynamics of the central regions of galaxies, to understand what activates some galaxies, but not others."
    },
    "0710/0710.4567_arXiv.txt": {
        "abstract": "We construct merger trees from the largest database of dark matter haloes to date provided by the Millennium simulation to quantify the merger rates of haloes over a broad range of descendant halo mass ($10^{12} \\la M_0 \\la 10^{15} M_\\odot$), progenitor mass ratio ($10^{-3} \\la \\xi \\le 1$), and redshift ($0 \\le z \\la 6$). We find the mean merger rate {\\it per halo}, $B/n$, to have very simple dependence on $M_0$, $\\xi$, and $z$, and propose a universal fitting form for $B/n$ that is accurate to 10-20\\%. Overall, $B/n$ depends very weakly on the halo mass ($\\propto M_0^{0.08}$) and scales as a power law in the progenitor mass ratio ($\\propto \\xi^{-2}$) for minor mergers ($\\xi \\la 0.1$) with a mild upturn for major mergers. As a function of time, we find the merger rate per Gyr to evolve roughly as $(1+z)^{n_m}$ with $n_m=2-2.3$, while the rate per unit redshift is nearly independent of $z$. Several tests are performed to assess how our merger rates are affected by, e.g. the time interval between Millennium outputs, binary vs multiple progenitor mergers, and mass conservation and diffuse accretion during mergers. In particular, we find halo fragmentations to be a general issue in merger tree construction from $N$-body simulations and compare two methods for handling these events.  We compare our results with predictions of two analytical models for halo mergers based on the Extended Press-Schechter (EPS) model and the coagulation theory.  We find that the EPS model overpredicts the major merger rates and underpredicts the minor merger rates by up to a factor of a few. ",
        "introduction": "In hierarchical cosmological models such as $\\Lambda$CDM, galaxies' host dark matter haloes grow in mass and size primarily through mergers with other haloes.  As the haloes merge, their more centrally concentrated baryonic components sink through dynamical friction and merge subsequently. The growth of stellar masses depends on both the amount of mass brought in by mergers and the star formation rates.  Having an accurate description of the mergers of dark matter haloes is therefore a key first step in quantifying the mergers of galaxies and in understanding galaxy formation and growth.  Earlier theoretical studies of galaxy formation typically relied on merger trees generated from Monte Carlo realisations of the merger rates given by the analytical extended Press-Schechter (EPS; \\citealt{1993MNRAS.262..627L, 1991ApJ...379..440B}) model (e.g. \\citealt{1993MNRAS.264..201K, 1999MNRAS.310.1087S, 2000MNRAS.319..168C}).  Some recent studies have chosen to bypass the uncertainties and inconsistencies in the EPS model by using halo merger trees from $N$-body simulations directly (\\citealt{1999MNRAS.303..188K, 2000MNRAS.311..793B, 2003MNRAS.338..903H, 2005ApJ...631...21K, 2005Natur.435..629S}).  As we find in this paper, obtaining robust halo merger rates and merger trees requires rich halo statistics from very large cosmological simulations as well as careful treatments of systematic effects due to different algorithms used for, e.g., assigning halo masses, constructing merger trees, removing halo fragmentation events, and choosing time spacings between simulation outputs.  The aim of this paper is to determine the merger rates of dark matter haloes as a function of halo mass, merger mass ratio (i.e. minor vs major), and redshift, using numerical simulations of the $\\Lambda$CDM cosmology. This basic quantity has not been thoroughly investigated until now mainly because large catalogues of haloes from finely spaced simulation outputs are required to provide sufficient merger event statistics for a reliable construction of merger trees over a wide dynamic range in time and mass. We achieve this goal by using the public database of the Millennium simulation \\citep{2005Natur.435..629S}, which follows the evolution of roughly $2\\times10^{7}$ dark matter haloes from redshift $z=127$ to $z=0$. This dataset allows us to determine the merger rates of dark matter haloes ranging from galaxy-mass scales of $\\sim 10^{12} M_\\odot$ over redshifts $z=0$ to $\\sim 6$, to cluster-mass scales up to $\\sim 10^{15} M_\\odot$ for $z=0$ to a few.  We are also able to quantify the merger rates as a function of the progenitor mass ratio $\\xi$, from major mergers ($\\xi \\ga 0.1$) down to minor mergers of $\\xi \\sim 0.03$ for galaxy haloes and down to $\\xi \\sim 3\\times 10^{-4}$ for cluster haloes.  The inputs needed for measuring merger rates in simulations include a catalogue of dark matter haloes and their masses at each redshift, and detailed information about their ancestry across redshifts, that is, the merger tree.  Unfortunately there is not a unique way to identify haloes, assign halo masses, and construct merger trees.  In this paper we primarily consider a halo mass definition based on the standard friends-of-friends (FOF) algorithm and briefly compare it with an alternative mass definition based on spherical overdensity.  For the merger trees, we investigate two possible algorithms for treating events in which the particles in a given progenitor halo end up in more than one descendant halo ('fragmentations').  We find that these events are common enough that a careful treatment is needed.  In the conventional algorithm used in the literature, the progenitor halo is linked one-to-one to the descendant halo that has inherited the largest number of the progenitor's particles.  The ancestry links to the other descendant haloes are severed (for this reason we call this scheme 'snipping').  We consider an alternative algorithm ('stitching') in this paper, in which fragmentations are assumed to be artefacts of the FOF halo identification scheme.  We therefore choose to recombine the halo fragments and stitch them back into the original FOF halo.  Earlier theoretical papers on merger rates either relied on a small sample of main haloes to estimate the overall redshift evolution over a limited range of halo masses, or were primarily concerned with the mergers of {\\it galaxies} or {\\it subhaloes}.  For halo mergers, for example, \\cite{1999AJ....117.1651G} studied $z < 1$ major mergers of galaxy-sized haloes in an open CDM and a tilted $\\Omega_m=1$ CDM model using $N$-body simulations in a 100 Mpc box and $144^3$ particles. \\cite{2001ApJ...546..223G} used a sample of $\\sim 4000$ haloes to study the environmental dependence of the redshift evolution of the major merger rate at $z < 2$ in $\\Lambda$CDM.  \\cite{2006ApJ...652...56B} studied major mergers of subhaloes in $N$-body simulations in a 171 Mpc box with $512^3$ particles and the connection to the observed close pair counts of galaxies. For galaxy merger rates, \\cite{2002ApJ...571....1M} and \\cite{2006ApJ...647..763M} are based on up to $\\sim 500$ galaxies formed in SPH simulations in $\\sim 50$ Mpc boxes with up to $144^3$ gas particles, while \\cite{2007arXiv0708.1814G} used the semi-analytical galaxy catalogue of \\cite{2006MNRAS.366..499D} based on the Millennium simulation.  This paper is organised as follows. Section~\\ref{MillenniumSection} describes the dark matter haloes in the Millennium simulation (\\S\\ref{DarkMatterHaloes}) and how we construct the merger trees (\\S\\ref{mergertrees}) .  We then discuss the issue of halo fragmentation and the two methods ('snipping' and 'stitching') used to treat these events in \\S\\ref{FOFLims}.  The notation used in this paper is summarised in \\S\\ref{Notation}.  Section~\\ref{StatsSection} describes how mergers are counted (\\S\\ref{MultiCount}) and presents four (related) statistical measures of the merger rate (\\S\\ref{stats}). The relation between these merger rate statistics and the analytical merger rate based on the Extended Press-Schechter (EPS) model is derived in Section~\\ref{EPSConnection}.  Our main results on the merger rates computed from the Millennium simulation are presented in Section~\\ref{Results}.  We first discuss the $z\\approx 0$ results and quantify the merger rates as a function of the descendant halo mass and the progenitor mass ratios using merger trees constructed from the stitching method (\\S\\ref{MRz0}).  The evolution of the merger rates with redshifts up to $z\\sim 6$ is discussed in Section~\\ref{redshiftdependence}.  We find a simple universal form for the merger rates and present an analytic fitting form that provides a good approximation (at the 10-20\\% level) over a wide range of parameters (\\S\\ref{FitSection}).  Section~\\ref{Tests} compares the stitching and snipping merger rates (\\S\\ref{SnipvsStitch}) and presents the key results from a number of tests that we have carried out to assess the robustness of our results.  Among the tests are: time convergence and the dependence of the merger rates on the redshift spacing $\\dz$ between the Millennium outputs used to construct the merger tree (\\S\\ref{TimeResolutionSection}); how the counting of binary vs multiple progenitor mergers affects the merger rates (\\S\\ref{MultiCountValid}); mass non-conservation arising from 'diffuse' accretion in the form of unresolved haloes during mergers (\\S\\ref{MassCons}); and how the definition of halo masses and the treatment of fragmentation events affect the resulting halo mass function (\\S\\ref{MassFunction}).  In Section~\\ref{DiscussionSection}, we discuss two theoretical frameworks that can be used to model halo mergers: EPS and coagulation.  A direct comparison of our merger rates and the EPS predictions for the Millennium $\\Lambda$CDM model shows significant differences over a large range of parameter space (\\S\\ref{eps}).  Section~\\ref{Coagulation} discusses Smoluchowski's coagulation equation and the connection between our merger rates and the coagulation merger kernel.  The appendix compares a third merger tree (besides snipping and stitching) constructed from the Millennium catalogue by the Durham group \\citep{2006MNRAS.370..645B, 2006MNRAS.367.1039H,2003MNRAS.338..903H}.  Two additional criteria are imposed on the subhaloes in this algorithm to reduce spurious linkings of FOF haloes.  We find these criteria to result in reductions in both the major merger rates and the halo mass function.  The cosmology used throughout this paper is identical to that used in the Millennium simulation: a $\\Lambda\\textrm{CDM}$ model with $\\Omega_m=0.25$, $\\Omega_b=0.045$, $\\Omega_\\Lambda=0.75$, $h=0.73$, an initial power-law index $n=1$, and $\\sigma_8=0.9$ \\citep{2005Natur.435..629S}.  Masses and lengths are quoted in units of $M_\\odot$ and Mpc without the Hubble parameter $h$. ",
        "conclusions": "In this paper we have computed the merger rates of dark matter FOF haloes as a function of descendant halo mass $M_0$, progenitor mass ratio $\\xi$, and redshift $z$ using the merger trees that we constructed from the halo catalogue of the Millennium simulation.  Our main results are presented in Figs.~\\ref{fig:B} to \\ref{fig:Rz}, which show very simple and nearly separable dependence on $M_0$, $\\xi$, and $z$. The mean merger rate per descendant FOF halo, $B/n$, is seen to depend very weakly on the halo mass $M_0$ (Fig.~\\ref{fig:B} right panel and Fig.~\\ref{fig:Rmass}). As a function of redshift $z$, the per halo merger rate in units of per Gyr increases as $(1+z)^\\alpha$, where $\\alpha\\sim 2$ to 2.3 (top panel of Fig.~\\ref{fig:Rz}), but when expressed in units of per redshift, the merger rate depends very weakly on $z$ (bottom panel of Fig.~\\ref{fig:Rz}). Regardless of $M_0$ and $z$, the dependence of $B/n$ on the progenitor mass ratio, $\\xi = M_i/M_1$, is a power law to a good approximation in the minor merger regime ($\\xi \\la 0.1$) and shows an upturn in the major merger regime (Fig.~\\ref{fig:B}). These simple behaviours have allowed us to propose a universal fitting formula in equation~(\\ref{fiteqn}) that is valid for $10^{12}\\leq M_0\\la 10^{15} M_\\odot$, $\\xi \\ga 10^{-3}$, and up to $z\\sim 6$.  Throughout the paper we have emphasised and quantified the effects on the merger rates due to events in which a progenitor halo fragments into multiple descendant haloes.  We have shown that the method commonly used to remove these fragmented haloes in merger trees -- the snipping method -- has relatively poor $\\dz$-convergence (Figs.~\\ref{fig:TR} and \\ref{fig:RzTR}).  Our alternative approach -- the stitching method -- performs well with regards to this issue without drastically modifying the mass conservation properties or the mass function of the Millennium FOF catalogue (Figs.~ \\ref{fig:DM} and \\ref{fig:MassFunction}).  We have computed the two predictions for merger rates from the analytical EPS model for the same $\\Lambda$CDM model used in the Millennium simulation.  At $z=0$, we find the EPS major merger rates to be too high by 50-100\\% (depending on halo mass) and the minor merger rates to be too low by up to a factor of 2-5 (Fig.~\\ref{fig:EPS}).  The discrepancy increases at higher $z$.  The coagulation equation offers an alternative theoretical framework for modelling the mergers of dark matter haloes.  We have discussed how our merger rate is related to the coagulation merger kernel in theory.  In practice, however, we find that mergers in simulations are not always mass-conserving binary events, as assumed in the standard coagulation form given by equation~(\\ref{eqn:coag}).  Equation~(\\ref{eqn:coag}) will therefore have to be modified before it can be used to model mergers in simulations.  \\cite{2001ApJ...546..223G} studied the rate of major mergers (defined to be $\\xi \\ge 1/3$ in our notation) in $N$-body simulations and found a steeper power law dependence of $\\propto (1+z)^3$ (at $z\\la 2$) for the merger rate per Gyr than ours.  Their simulations did not have sufficient mass resolution to determine the rate at $z \\ga 2$.  It is important to note, however, that our $B/n$ at redshift $z$ measures the instantaneous rate of mergers during a small $\\Delta z$ interval at that redshift.  By contrast, they studied the merging history of {\\it present-day} haloes and measured only the rate of major mergers for the most massive progenitor at redshift $z$ of a $z=0$ halo (see their paragraph 4, section 2).  A detailed comparison is outside the interest of this paper.  Mergers of dark matter haloes are related to but not identical to mergers of galaxies.  It typically takes the stellar component of an infalling galaxy extra time to merge with a central galaxy in a group or cluster after their respective dark matter haloes have been tagged as merged by the FOF algorithm.  This time delay is governed by the dynamical friction timescale for the galaxies to lose orbital energy and momentum, and it depends on the mass ratios of the galaxies and the orbital parameters (\\citealt{2008MNRAS.383...93B} and references therein).  In addition to this difference in merger timescale, the growth in the stellar mass of a galaxy is not always proportional to the growth in its dark matter halo mass.  A recent analysis of the galaxy catalogue in the Millennium simulation \\citep{2007arXiv0708.1814G} finds galaxy growth via major mergers to depend strongly on stellar mass, where mergers are more important in the buildup of stellar masses in massive galaxies while star formation is more important in galaxies smaller than the Milky Way. Extending the analysis of this paper to the mergers of {\\it subhaloes} in the Millennium simulation will provide the essential link between their and our results.  For similar reasons, our results for the evolution of the dark matter halo merger rate per Gyr ($(1+z)^{n_m}$ with $n_m\\sim2-2.3$) cannot be trivially connected to the observed merger rate of \\emph{galaxies}.  It is nonetheless interesting to note that a broad disagreement persists in the observational literature of galaxy merger rates.  The reported power law indices $n_m$ have ranged from 0 to 5 (see, for example, \\citealt{1994ApJ...429L..13B, 1994ApJ...435..540C, 1995ApJ...445...37Y, 1995ApJ...454...32W, 1997ApJ...475...29P, 2000MNRAS.311..565L, 2002ApJ...565..208P,2003AJ....126.1183C, 2004ApJ...601L.123B,2004ApJ...612..679L, 2004ApJ...617L...9L}). \\cite{2006ApJ...652...56B} followed the redshift evolution of subhalo mergers in N-body simulations and provided a more detailed comparison with recent observations by, e.g., \\cite{2004ApJ...617L...9L} that find $n_m<1$.  They attributed such a weak redshift evolution in the number of close companions per galaxy to the fact that the high merger rate per halo at early times is counteracted by a decrease in the number of haloes massive enough to host a galaxy pair.  The merger rates in this paper are global averages over all halo environments.  The rich statistics in the Millennium simulation allow for an in-depth analysis of the environmental dependence of dark matter halo merger rates, which we will report in a subsequent paper (Fakhouri \\& Ma, in preparation)."
    },
    "0710/0710.1945_arXiv.txt": {
        "abstract": "{The ARGO-YBJ experiment is a full coverage EAS-array installed at the YangBaJing Cosmic Ray Laboratory (4300 m a.s.l., Tibet, P.R. China). We present the results on the angular resolution measured with different methods with the full central carpet. The comparison of experimental results with MC simulations is discussed.} \\begin{document} ",
        "introduction": "The ARGO-YBJ detector is constituted by a single layer of Resistive Plate Chambers (RPCs). This carpet has a modular structure, the basic unit is a cluster, composed by 12 RPCs (2.8$\\times$1.25 m$^2$ each). Each chamber is read by 80 strips, logically organized in 10 independent pads\\cite{nim_argo}. The central carpet, constituted by 10$\\times$13 clusters with $\\sim$93$\\%$ of active area, is enclosed by a guard-ring partially instrumented ($\\sim$40$\\%$) in order to improve rejection capability for external events. A lead converter 0.5 cm thick will uniformly cover the apparatus in order to improve the angular resolution. Since December 2004 the pointing accuracy of the detector has been studied, during the detector setting-up, with 3 different carpet areas: 42 clusters (ARGO-42, $\\sim$1900 m$^2$), 104 clusters (ARGO-104, $\\sim$4600 m$^2$) and the full central carpet, 130 clusters (ARGO-130, $\\sim$5800 m$^2$), yet without any converter sheet. The data have been collected with a so-called {\\it \"Low Multiplicity Trigger\"}, requiring at least $20$ fired pads on the whole detector. ",
        "conclusions": "Since December 2004 increasing fractions of ARGO-YBJ detector have been put in data taking even with a reduced duty-cycle due to installation and debugging operations. In this paper we presented a measurement of the pointing accuracy of the ARGO-130 detector. The capability of reconstructing the primary shower direction has been investigated with the chessboard method and with a preliminary Moon shadow analysis. Studies are in progress in order to determine the final angular resolution."
    },
    "0710/0710.5427_arXiv.txt": {
        "abstract": "\\footnotesize\\ The efficient numerical solution of Non-LTE multilevel transfer problems requires the combination of highly convergent iterative schemes with fast and accurate formal solution methods of the radiative transfer (RT) equation. This contribution\\footnote{Published in 1999 in the book {\\it Solar Polarization}, edited by K.N. Nagendra \\& J.O. Stenflo. Kluwer Academic Publishers, 1999. (Astrophysics and Space Science Library ; Vol. 243), p. 219-230} begins presenting a method for the formal solution of the RT equation in three-dimensional (3D) media with horizontal periodic boundary conditions. This formal solver is suitable for both, unpolarized and polarized 3D radiative transfer and it can be easily combined with the iterative schemes for solving non-LTE multilevel transfer problems that we have developed over the last few years. We demonstrate this by showing some schematic 3D multilevel calculations that illustrate the physical effects of horizontal radiative transfer. These Non-LTE calculations have been carried out with our code MUGA 3D, a 3D multilevel Non-LTE code based on the Gauss-Seidel iterative scheme that Trujillo Bueno and Fabiani Bendicho (1995) developed for RT applications. ",
        "introduction": "To what extent can we trust diagnostic results obtained with the assumption that the solar atmospheric plasma is composed of {\\it homogeneous} plane-parallel layers or via approximations that neglect {\\it horizontal} radiative transfer (RT) effects? How important are the errors in the magnetic fields, temperatures and velocities inferred by confronting spectro-polarimetric observations with Non-LTE 1D RT model calculations? Clearly, to provide proper answers to questions like these requires to develop first efficient 3D RT methods that allow Non-LTE effects in complex atomic models with many levels and transitions to be rigorously investigated.  There is a second reason which makes the development of fast iterative methods for 3D Non-LTE RT so relevant. This is because processes of energy exchange by radiation play an important role in the structure and dynamical behaviour of the stellar magnetized plasma. Thus, for instance, if one wishes to perform time-dependent radiation hydrodynamics simulations similar to those carried out by Carlsson and Stein (1997), but in 3D instead of 1D, it turns out to be imperative to have first access to numerical methods capable of accurately yielding the self-consistent atomic level populations at the cost of only {\\it very} few formal solution times.  The efficient solution of multilevel transfer problems requires the combination of a highly convergent iterative scheme with a fast formal solver of the RT equation. In Section 2 we briefly comment on a hierarchy of iterative schemes that can be applied for solving multilevel Non-LTE problems with increasing improvements in the convergence rate and total computational work. The 3D multilevel transfer calculations that we present in this contribution have been obtained by combining a highly convergent iterative scheme based on Gauss-Seidel iteration (Trujillo Bueno and Fabiani Bendicho, 1995) with a fast 3D formal solver that has parabolic accuracy (see Section 3). As was the case with our 2D formal solver, our generalization to 3D is based on the ``short-characteristics'' method of Kunasz and Auer (1988). Our 3D multilevel code is called MUGA 3D (``Multi-level Gauss-Seidel Method'') and it is substantially {\\it faster} than our code MALI 3D, which is based on Jacobi iteration (see Rybicki and Hummer, 1991; Auer, Fabiani Bendicho and Trujillo Bueno, 1994).  Section 3 briefly describes our 3D formal solver as applied to the scalar transfer equation for the specific intensity (I). In order to be able to consider 3D atmospheric models where solar plasma structures repeat themselves along the horizontal directions we choose horizontal periodic boundary conditions along the Cartesian coordinates X and Y. Although we do not give any details here, we have also generalized to 3D the Stokes-vector 1D formal solver method developed by Trujillo Bueno (1998), which is based on the matrix exponential approximation to the evolution operator.  In Section 4 we show some illustrative 3D multilevel transfer calculations for a 5-level Ca II model atom where the H, K and infrared triplet lines are treated simultaneously, taking fully into account the {\\it interlockings} by which photons are converted back and forth between the different line transitions in the assumed 3D medium. Here we consider schematic 3D solar models characterized by horizontal sinusoidal temperature inhomogeneities. With the help of these 3D multilevel calculations we are able to illustrate some subtle effects of horizontal radiative transfer that are important for the correct interpretation of high spatial resolution observations. Finally, Section 5 gives our conclusions. ",
        "conclusions": "We have developed a 3D multilevel code (MUGA-3D) that combines the Gauss-Seidel iterative scheme of Trujillo Bueno and Fabiani Bendicho (1995) with a 3D formal solver that uses horizontal periodic boundary conditions. With this new code we have performed some 3D multilevel simulations that highlight the importance of carefully investigating the effects of horizontal radiative transfer using realistic atmospheric and atomic models.  We point out that our 3D formal solver can be used not only for solving ``unpolarized'' multilevel transfer problems, but also resonance line polarization and Hanle effect problems, like those considered in these Proceedings by Manso Sainz and Trujillo Bueno (1999), Paletou {\\it et. al.} (1999) or Dittmann (1999). This is because, for these polarization transfer cases, the absorption matrix is diagonal. As a result, we have similar equations for the Stokes I, Q and U parameters.  However, for the solution of more general polarization transfer problems, like the Non-LTE Zeeman line transfer case considered by Trujillo Bueno and Landi Degl'Innocenti (1996), but in 3D instead of 1D, one needs a 3D formal solution method of the Stokes-vector transfer equation. This is because here the absorption matrix turns out to be a full $4\\times4$ matrix and all the Stokes parameters are coupled together. To this end we have generalized to 3D the Stokes-vector formal solver developed by Trujillo Bueno (1998), which can be considered as a generalization to polarization transfer of the short-characteristics method.  We would like to end this paper by saying that over the last 10 years we have witnessed impressive progress concerning the development of highly convergent iterative schemes and accurate formal solvers for RT applications. Now it is time to apply them with physical intuition in order to improve our knowledge of the Sun, its magnetic field and its polarized spectrum."
    },
    "0710/0710.1338_arXiv.txt": {
        "abstract": "The two-body problem in general relativity is reviewed, focusing on the final stages of the coalescence of the black holes as uncovered by recent successes in numerical solution of the field equations. ",
        "introduction": "A black hole is one of the most fascinating and enigmatic predictions of Einstein's theory of general relativity. Its interior can have rich structure and is intrinsically dynamical, where space and time itself are inexorably led to a singular state. The exterior of an isolated black hole is, on the other hand, remarkably simple, described uniquely by the stationary Kerr solution. The dynamics of black holes are governed by laws analogous to the laws of thermodynamics, and indeed when quantum processes are included, emit Hawking radiation with a characteristic thermal spectrum. Most remarkable however, is that black holes, ``discovered'' purely through thought and the mathematical exploration of a theory far removed from every day experience, appear to be ubiquitous objects in our universe.  The evidence that black holes exist, though circumstantial, is quite strong~\\cite{Narayan:2005ie}. The high luminosity of quasars and other active galactic nuclei (AGN) can be explain by gravitational binding energy released through gas accretion onto supermassive ($10^6-10^9 \\msun$) black holes at the centers of the galaxies~\\cite{Rees:1984si,Ferrarese:2004qr}, several dozen X-ray binary systems discovered to date have compact members too massive to be neutron stars and exhibit phenomena consistent with matter interactions originating in the strong gravity regime of an inner accretion disk~\\cite{McClintock:2003gx}, and the dynamical motion of stars and gas about the centers of nearby galaxies and our Milky Way Galaxy infer the presence of very massive, compact objects there, the most plausible explanation being supermassive black holes~\\cite{Gebhardt:2000fk,Schodel:2002py,Ghez:2003qj}.  To conclusively prove that black holes exist one needs to ``see'' them, or conversely see the compact objects masquerading as black holes. The only direct way of observing black holes is via the gravitational waves they emit when interacting with other matter/energy (an isolated black hole does not radiate). The quadrupole formula says that the typical magnitude $h$ of the gravitational waves emitted by a binary with reduced mass $\\mu$ on a circular orbit measured a distance $r$ from the source is (for a review of gravitational wave theory see~\\cite{Flanagan:2005yc}) \\be h=\\frac{16 \\mu v^2}{r}, \\ee where $v$ is the average tangential speed of the two members in the binary (and geometric units are used---Newton's constant $G=1$ and the speed of light $c=1$). This formula suggests that the strongest sources of gravitational waves are simply the most massive objects that move the fastest. To reach large velocities in orbit, the binary separation has to be small; black holes, being the most compact objects allowed in the theory, can reach the closest possible separations and hence largest orbital velocities. Therefore, modulo questions about source populations in the universe, a binary black hole interaction offers one of most promising venues of observing black holes through gravitational wave emission.  Joseph Weber pioneered the science of gravitational wave detection with the construction of resonant bar detectors. Weber claimed to have detected gravitational waves~\\cite{Weber:1969bz}, though no similar detectors constructed following his claims were able to observe the putative (or any other) source, and the general consensus is that given the sensitivity of Weber's detector and expected strengths of sources it is very unlikely that it was a true detection~\\cite{Thorne:1980rt}. Note that the {\\em existence} of gravitational waves is not in doubt---the observed spin down rate of the Hulse-Taylor binary pulsar~\\cite{Hulse:1974eb} and several others discovered since, is in complete accord with the general relativistic prediction of spin down via gravitational wave emission. Today a new generation of gravitational wave detectors are operational, including laser interferometers (LIGO~\\cite{LIGO}, VIRGO~\\cite{VIRGO}, GEO600~\\cite{GEO}, TAMA~\\cite{TAMA}) and resonant bar detectors (NAUTILUS~\\cite{NAUTILUS}, EXPLORER~\\cite{EXPLORER}, AURIGA~\\cite{AURIGA}, ALLEGRO~\\cite{ALLEGRO}, NIOBE~\\cite{NIOBE}). A future space-based observatory is planned (LISA~\\cite{LISA}), and pulsar timing and cosmic microwave background polarization measurements also offer the promise of acting as gravitational wave ``detectors'' (for reviews see~\\cite{Maggiore:1999vm,Cutler:2002me}). The ultimate success of gravitational wave detectors, in particular with regards to using them as more that simply detectors, but tools to observe and understand the universe, relies on source modeling to predict the structure of the waves emitted during some event. Even if an event is detected with a high signal-to-noise ratio (SNR), there simply is not enough information contained in such a one dimensional time series to ``invert'' it to reconstruct the event; rather template banks of theoretical waveforms from plausible sources need to be built and used to decode the signal. In rare cases an electromagnetic counterpart may be detected, for example during a binary neutron star merger if this is a source of short gamma ray bursts, which could identify the event without the need for a template. Though even in such a case, to extract information about the event, its environment, etc. requires source modeling.  Gravitational wave detectors have therefore provided much of the impetus for trying to understand the nature of binary black hole collisions, and the gravitational waves emitted during the process. However, from a theoretical perspective black hole collisions are fantastic probes of the dynamical, strong-field regime of general relativity. What is already know about this regime---the inevitability of spacetime singularities in gravitational collapse via the singularity theorems of Penrose and Hawking~\\cite{Hawking:1969sw,hawking_ellis}; the spacelike, chaotic ``mixmaster'' nature of these singularities conjectured by Belinsky, Khalatnikov and Lifshitz (BKL)~\\cite{Berger:1998us}; the null, mass-inflation singularity discovered by Poisson and Israel~\\cite{Poisson:1990eh} that, together with regions of BKL singularities could generically describe the interiors of black hole; the rather surprising discovery of critical phenomena in gravitational collapse by Choptuik~\\cite{Choptuik:1992jv,Gundlach:1999cu}; etc---together with the sparsity of solutions (exact, numerical or perturbative), suggests there is potentially a vast landscape of undiscovered phenomena. Of particular interest, and potential application to high energy particle collision experiments, are ultra-relativistic black hole collisions. It is beyond the scope of this article to delve much into these aspects of black hole coalescence, though a brief overview of this will be given in Sec.~\\ref{sec_he}.  The two body problem in general relativity, introduced in more detail in Sec.~\\ref{sec_2bdy}, is a very rich and complicated problem, with no known closed-form solution. Perturbative analytic techniques have been developed to deal with certain stages of the problem, in particular the inspiral prior to merger and ringdown after merger. Numerical solution of the full field equations are required during the merger, and this aspect of the problem is the main focus of this article. Much effort has been expended by the community over the past 15-20 years to numerically solve for merger spacetimes, and within the last two years an understanding of this phase of the two body problem is finally being attained. Sec.~\\ref{sec_num} summaries the difficulties in discretizing the field equations, and describes the methods known at present that work for black hole collisions, namely {\\em generalized harmonic coordinates} and {\\em BSSN} (Baumgarte-Shapiro-Shibata-Nakamura) with moving punctures. Preliminary results are discussed in Sec.~\\ref{sec_res}, though given the rapid pace at which the field is developing much of this will probably be dated in short order. Sec.~\\ref{sec_imp} concludes with a discussion of some astrophysical and other implications of the results. ",
        "conclusions": "The two body problem in general relativity is a fascinating, rich problem that is just beginning to be fully revealed by recent breakthroughs in numerical relativity. At the same time, a new generation of gravitational wave detectors promise to offer us a view of the universe via the gravitational wave spectrum for the first time. Black hole mergers are a promising source for gravitational waves, and detecting them would provide direct evidence for these remarkable objects, while providing much information about their environments. Suggestions that there might be more than four spacetime dimensions offers the astonishing possibility that black holes could be produced by proton collisions at $TeV$ energies, which will be reached by the Large Hadron Collider, planned to begin operation within a year. Given all this, it is difficult not to be excited about what might be learnt about the universe from the smallest to largest scales during the next decade, and that black hole collisions could have something important to say at both extremes.  \\bigskip  \\noindent{\\bf{\\em Acknowledgments:}} I would like to thank Alessandra Buonanno, Matthew Choptuik, Gregory Cook, Charles Gammie, David Garfinkle, Steven Gubser, Carsten Gundlach, Luis Lehner, Jeremiah Ostriker, Don Page, David Spergel and Ulrich Sperhake for many stimulating conversations related to some of the discussion presented here."
    },
    "0710/0710.5388_arXiv.txt": {
        "abstract": "Recently, it has been shown that the standard Nambu-Jona-Lasinio (NJL) model is not able to reproduce the correct QCD behavior of the gap equation at large density, and therefore a different cutoff procedure at large momenta has ben proposed. We found that, even with this density dependent cutoff procedure, the pure quark phase in neutron stars (NS) interiors is unstable, and we argue that this could be related to the lack of confinement in the original NJL model. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.5207_arXiv.txt": {
        "abstract": "If massive black holes (BHs) are ubiquitous in galaxies and galaxies experience multiple mergers during their cosmic assembly, then BH binaries should be common albeit temporary features of most galactic bulges. Observationally, the paucity of active BH pairs points toward binary lifetimes far shorter than the Hubble time, indicating rapid inspiral of the BHs down to the domain where gravitational waves lead to their coalescence.  Here, we review a series of studies on the dynamics of massive BHs in gas-rich galaxy mergers that underscore the vital role played by a cool, gaseous component in promoting the {\\it rapid formation of the BH binary}.  The BH binary is found to reside at the center of a massive self-gravitating nuclear disc resulting from the collision of the two gaseous discs present in the mother galaxies.  Hardening by gravitational torques against gas in this grand disc is found to continue down to sub-parsec scales.  The eccentricity decreases with time to zero and when the binary is circular, accretion sets in around the two BHs.  When this occurs, each BH is endowed with it own small-size ($\\simlt 0.01$ pc) accretion disc comprising a few percent of the BH mass. Double AGN activity is expected to occur on an estimated timescale of $\\simlt 1$ Myr.  The double nuclear point--like sources that may appear have typical separation of $\\simlt 10$ pc, and are likely to be embedded in the still ongoing starburst.  We note that a potential threat of binary stalling, in a gaseous environment, may come from radiation and/or mechanical energy injections by the BHs.  Only short--lived or sub--Eddington accretion episodes can guarantee the persistence of a dense cool gas structure around the binary necessary for continuing BH inspiral. ",
        "introduction": "Dormant black holes (BHs) with masses in excess of $\\simgt 10^6\\msun$ are ubiquitous in bright galaxies today (Kormendy \\& Richstone 1995; Richstone 1998).  Relic of an earlier active phase as quasars, these massive BHs appear a clear manifestation of the cosmic assembly of galaxies.  The striking correlations observed between the BH masses and properties of the underlying hosts (Magorrian \\etal 1998; Ferrarese \\& Merritt 2000; Gebhardt et al. 2000; Graham \\& Driver 2007) indicates unambiguously that BHs evolve in symbiosis with galaxies, affecting the environment on large--scales and self-regulating their growth (Silk \\& Rees 1998; King 2000; Granato et al. 2004; Di Matteo, Springel \\& Hernquist, 2005).   According to the current paradigm of structure formation, galaxies often interact and collide as their dark matter halos assemble in a hierarchical fashion (Springel, Frenk \\& White 2006), and BHs incorporated through mergers into larger and larger systems are expected to evolve concordantly (Volonteri, Haardt \\& Madau 2003). In this astrophysical context, close BH {\\it pairs} form as natural outcome of binary galaxy mergers (Kazantzidis et al. 2005).   In our local universe, one outstanding example is the case of the ultra--luminous infrared galaxy NGC 6240, an ongoing merger between two gas--rich galaxies (Komossa \\etal 2003; for a review on binary black holes see also Komossa 2006). {\\it Chandra} images have revealed the occurrence of two nuclear X-ray sources, 1.4 kpc apart, whose spectral distribution is consistent with being two accreting massive BHs embedded in the diluted emission of a starburst. Similarly, Arp 299 (Della Ceca et al. 2002; Ballo et al. 2004) is an interacting system hosting an obscured active nucleus, and possibly a second less luminous one, distant several kpc away.  A third example is the elliptical galaxy 0402+369 where the cleanest case of a massive BH {\\it binary} has been recently discovered. Two compact variable, flat--spectrum active nuclei are seen at a projected separation of only 7.3 pc (Rodriguez et al. 2006). Arp 299, NGC 6240, and 0402+369 may just highlight different stages of the BH dynamical evolution along the course of a merger, with 0402+369 being the latest, most evolved phase (possibly related to a dry merger).  Energy and angular momentum losses due to gravitational waves are not yet significant in 0402+392, so that stellar interactions and/or material and gravitational torques are still necessary to bring the BHs down to the domain controlled by General Relativity.  From the above considerations and observational findings, it is clear that binary BH inspiral down to coalescence is a major astrophysical process that can occur in galaxies.  It is accompanied by a gravitational wave burst so powerful to be detectable out to very large redshifts with current planned experiments like the Laser Interferometer Space Antenna ({\\it LISA}; Bender et al. 1994; Vitale et al. 2002).  These extraordinary events will provide not only a firm test of General Relativity but also a view, albeit indirect, of galaxy clustering (Haehnelt 1994; Jaffe \\& Backer 2003; Sesana et al. 2005). With {\\it LISA}, BH masses and spins will be measured with such an accuracy (Vecchio 2004) that it will be possible to trace the BH mass growth across all epochs. Interestingly, {\\it LISA} will explore a mass range between $10^3\\msun$ and $10^7\\msun$ that is complementary to that probed by the distant massive quasars ($>10^7\\msun$), providing a complete census of the BHs in the universe.   Both minor as well as major mergers with BHs accompany galaxy evolution in environments that involve either gas-rich (wet) as well as gas-poor (dry) galaxies.  Thus, the dynamical response of galaxies to BH pairing should differ in many ways according to their properties.  Exploring the expected diversities in a self-consistent cosmological scenario is a major challenge and only recently, with the help of high-resolution N-body/SPH simulations, it has become possible to ``start'' addressing a number of compelling issues.  Galaxy mergers cover cosmological volumes (a few to hundred kpc aside), whereas BH mergers probe volumes of only few astronomical units or less. Thus, tracing the BH dynamics with scrutiny requires N-Body/SPH force resolution simulations spanning more than nine orders of magnitude in length. For this reason, two complementary approaches have been followed in the literature.  A statistical approach (based either on Monte Carlo realizations of merger trees or on N-Body/SPH large scale simulations) follows the collective growth of BHs inside dark matter halos.  Supplemented by semi-analytical modeling of BH dynamics (Volonteri et al. 2003) or/and by sub-grid resolution criteria for accretion and feedback (Springel \\& Hernquist 2003; Springel, Di Matteo \\& Hernquist 2005), these studies have proved to be powerful in providing estimates of the expected coalescence rates, and in tracing the overall cosmic evolution of BHs including their feedback on the galactic environment (Di Matteo, Springel \\& Hernquist 2005; Di Matteo \\etal 2007).  The second approach, that we have been following, looks at individual binary collisions, as it aims at exploiting in detail the BH dynamics and some bulk physics from the galactic scale down to and within the BH sphere of influence.  Both approaches, the collective and the individual, are necessary and complementary, the main challenge being the implementation of realistic input physics in the dynamically active environment of a merger.   Following a merger, how can BHs reach the gravitational wave inspiral regime?  The overall scenario was first outlined by Begelman, Blandford \\& Rees (1980) in their seminal study on the dynamical evolution of BH pairs in pure stellar systems.  They indicated three main roots for the loss of orbital energy and angular momentum: (I) dynamical friction against the stellar background acting on each individual BH; (II) hardening via 3--body scatterings off single stars when the BH binary forms; (III) gravitational wave back--reaction.  Early studies explored phase (I) simulating the collisionless merger of spherical halos (Makino \\& Ebisuzaki 1996; Milosavljevi\\'c \\& Merritt 2001; Makino \\& Funato 2004).  Governato, Colpi \\& Maraschi (1994) in particular first noticed that when two equal mass halos merge, the twin BHs nested inside the nuclei are dragged effectively toward the center of the remnant galaxy by dynamical friction and form a close pair, but that the situation reverses in unequal mass mergers, where the less massive halo tidally disrupted leaves its ``naked'' BH wandering in the outskirts of the main halo. Thus, depending on the halo mass ratio and internal structure, the transition from phase (I) to phase (II) can be prematurely aborted or drastically relented.  Similarly, the transit from phase (II) to phase (III) is not always secured, as the stellar content inside the ``loss cone\" may not be rapidly refilled with fresh low--angular momentum stars to harden the binary down to separations where gravitational wave driven inspiral sets in (see, e.g., Milosavljevic \\& Merritt 2001; Yu 2002; Berczik, Merritt \\& Spurzem 2005; Sesana, Haardt \\& Madau 2007). For an updated review on the last parsec problem and its possible solution (see Merritt 2006a; Gualandris \\& Merritt 2007).  Since BH coalescences are likely to be events associated with mergers of (pre--)galactic structures at high redshifts, it is likely that their dynamics occurred in gas dominated backgrounds, NGC 6240 being just the most outstanding case visible in our local universe. Other processes of BH binary hardening are expected to operate in presence of a dissipative gaseous component that we will highlight and study here.  Kazantzidis \\etal (2005) first explored the effect of gaseous dissipation in mergers between gas--rich disc galaxies with central BHs, using high resolution N--Body/SPH simulations.  They found that the merger triggers large--scale gas dynamical instabilities that lead to the gathering of cool gas deep in the potential well of the interacting galaxies.  In minor mergers, this fact is essential in order to bring the BHs to closer and closer distances before the less massive galaxy, tidally disrupted, is incorporated in the main galaxy. Moreover, the interplay between strong gas inflows and star formation leads naturally to the formation, around the two BHs, of a grand, massive ($\\sim 10^9\\msun$) gaseous disc on a scale smaller than $\\sim 100$ pc.  It is in this equilibrium circum--nuclear disc that the dynamical evolution of the BHs continues, after the merger has been subsided.  Escala \\etal (2005, hereinafter ELCM05; see also Escala \\etal 2004) have been the first to study the role played by gas in affecting the dynamics of massive ($\\sim 10^8\\msun$) twin BHs in equilibrium Mestel discs of varying clumpiness.  In both these approaches (i.e., in the large scale simulations of Kazantzidis et al., and in the equilibrium disc models of ELCM05) it was clear that the gas temperature is a key physical parameter and that a hot gas brakes the BHs inefficiently. Instead, when the gas is allowed to cool, the drag becomes efficient: the large enhancement of the local gas density relative to the stellar one leads to the formation of prominent density wakes that are decelerating the BHs down to the scale where they form a ``close'' binary.  Later, binary hardening occurs under mechanisms that are only partially explored, and that are now subject of intense investigation.  The presence of a cool circum--binary disc and of small--scale discs around each individual BH appear to be critical for their evolution down to the domain of gravitational waves driven inspiral. In this context there is no clear ``stalling problem'' that emerges from current hydrodynamical simulations but this critical phase need a more through, coherent analysis.    The works by Kazantzidis \\etal (2005) and ELCM05 have provided our main motivation to study the process of BH pairing along two lines: In gas-rich binary mergers, line (1) aims at studying the transit from state (A) of pairing when each BH moves individually inside the time-varying potential of the colliding galaxies, to state (B) when the two BHs dynamically couple their motion to form a binary. The transit from (A) $\\to$ (B) requires exploring a dynamic range of five orders of magnitude in length from the cosmic scale of a galaxy merger of 100 kpc down to the parsec scale for BHs of million solar masses (i.e., BHs in the LISA sensitivity domain).  After all transient inflows have subsided and a new galaxy has formed, the BH binary is expected to enter phase (C) where it hardens under the action of gas-dynamical and gravitational torques. Research line (2) aims at studying the braking of the BH binary from (B) $\\to$ (C) and further in, exploring the possibility that during phase (C) two discs form and grow around each individual BH. As first discussed by Gould \\& Rix (2000) the binary may later enter a new phase (D) controlled by the balance of viscous and gravitational torques in a circum--binary disc surrounding the BHs, in a manner analogous to the migration of planets in circum-stellar discs (a scenario particularly appealing when the BH mass ratio is less than unity). Phase (D) likely evolves into (E) when gravitational wave inspiral terminates the BH binary evolution.     There is a number of key questions to address:   \\noindent (i) How does transition from state (A) $\\to$ (B) depend on the gas thermodynamics? How do BHs bind?  \\noindent (ii) In the grand nuclear disc inside the remnant galaxy, how do eccentric orbits evolve? Do they become circular or highly eccentric?  \\noindent (iii) During the hardening through phase (B) and (C), do the BHs collect substantial amounts of gas to form cool individual discs?  \\noindent (iv) Can viscous torques drive the binary into the gravitational wave decaying phase?  \\noindent (v) Is there a threat of a {\\it stalling} problem when transiting from (C) $\\to$ (D) or from (D) $\\to$ (E)?  And, for which mass ratios and ambient conditions? ",
        "conclusions": ""
    },
    "0710/0710.0368_arXiv.txt": {
        "abstract": "We compute opacities for the electronic molecular band systems \\ACrH -- \\XCrH of CrH and CrD, and \\AMgH -- \\XMgH of MgH and MgD. The opacities are computed by making use of existing spectroscopic constants for MgH and CrH. These constants are adjusted for the different reduced masses of MgD and CrD. Frank-Condon factors are used to provide intensities for the individual vibronic bands. These results are used in the computation of synthetic spectra between \\Tef = 1800 and 1200 K with an emphasis on the realisation of ``deuterium test'', first proposed by Bejar et al. (1999) to distinguish brown dwarfs from planetary mass objects. We discuss the possible use of CrD and MgD electronic bands for the ``deuterium test\". We find CrD to be the more promising of the two deuterides, potentially, the most useful bands of CrH/CrD are the  $\\Delta v = +1$  and $\\Delta v = -1$ at 0.795 and 0.968 \\mum . ",
        "introduction": "The ``deuterium test'' was suggested as a method of identifying planetary mass objects among cool objects (\\Bejar et al. 1999, Chabrier et al. 2000). In practise, it was proposed to search for absorption lines of molecules containing deuterium (HDO, CrD, FeD, etc.). Deuterium is burnt in stellar interiors via the fusion reaction $^2$D(p,$\\gamma$)$^3$He at  temperatures ($T > 8 \\times 10^5$~ K). The interiors of substellar objects with M $<$ 13~ \\MJ, where \\MJ is a Jovian mass (0.001 \\Msun), do not reach temperatures high enough for deuterium to ignite (Saumon et al. 1996). As a result the deuterium abundance in the atmospheres of these objects is unchanged from the formation of these objects. This gives rise to the definition of a brown dwarf as an object which has insufficient mass to fuse $^1$H to $^4$He, but has sufficient mass to fuse $^2$H to $^3$He.  By comparison, a planetary mass object has insufficient mass to ignite fusion of any sort (Saumon et al. 1996).  In higher-mass objects such as stars, deuterium burning is completed comparatively quickly (t = 1 -- 3 million years) during the evolution prior to the star's period on the main sequence (D'Antona \\& Mazitelli 1998). The deuterium depletion rate depends on the mass of the star or brown dwarf and thus the ``deuterium test'' can be used in a number of ways. \\begin{itemize} \\item Discern planets from brown dwarfs in a population of low-mass objects.  \\item Determine the evolutionary status of young objects in open clusters with ages of several million years.  \\item Study the evolution of the abundance of deuterium in the atmospheres of low-mass substellar objects, a phenomenon which is poorly understood. The rate of  depletion of deuterium depends upon rotation, magnetic field strength, and other parameters which affect the efficiency of convection in low mass objects. \\end{itemize}  Such investigations can usefully be combined with the ``lithium test'' proposed by Rebolo et al. (1992), \\Magazzu et al. (1993). For stars (M $>$ 75 \\MJ), the burning of lithium, Li (p,$\\alpha) ^4$He, becomes efficient at evolutionary stages preceding the main sequence at interior temperatures of T $>$ $2.5\\times 10^6$ K (D'Antona \\& Mazitelli 1998). The ``lithium test'' has been successfully applied to identify brown dwarfs in a population of ultracool dwarf stars (Rebolo et al. 1996, Basri 2000). The lithium test is relatively easily applied to M-dwarfs. The reason for this is that the resonance lines of neutral atomic lithium lie in the optical region of the spectrum, at 0.6708 \\mum. One blemish with the ``lithium test'' is the severe blending of lithium lines with the background of TiO lines, nonetheless Li lines break through the molecular background. In the spectra of cooler M and L dwarfs Ti atoms are depleted onto dust particles (Tsuji et al. 1996, Jones \\& Tsuji 1997, Pavlenko 1998) so TiO absorption is weakened. In the L dwarf regime the appearance of lithium lines contends with dust opacity and the wings of K I and Na I resonance lines (see Pavlenko et al. 2000 and references therein).  The realisation of the ``deuterium test'' is considerably more difficult. Conventionally, the deuterium abundance of hot objects is determined from analysis of multicomponent features on the background of the $L_\\alpha$ 0.1215 \\mum ~H~I line seen in emission. Unfortunately, this method cannot be used in the case of ultracool dwarfs which are covered by a thick envelope of neutral hydrogen. Observations also yield H~I emission lines ($H_\\alpha$ 0.6563 \\mum) formed in the outermost layers of hot chromospheres or accretion disks around young stars.  However, such a plasma is variable and may be polluted by  interstellar material.  Taking into account these circumstances, it is logical to analyse the spectral lines of deuterated molecules formed in the photospheric layers of ultracool dwarfs. The first investigation and analysis of the combined spectra of \\HHO /HDO were carried out by Chabrier et al. (2000) and Pavlenko (2002). Due to the change in mass and the breakdown of molecular symmetry the vibration-rotation bands of HDO in the mid-infrared spectrum shift with respect to the \\HHO bands. There are a few spectral regions which can be used for detection of HDO lines in the IR spectra of ultracool dwarfs: 3.5 -- 4, $\\sim$5, 6-7 \\mum (Pavlenko 2002). The main problem is that despite the shift in wavelength such HDO lines will be on a background of far stronger \\HHO lines.  A possible alternative to HDO are the diatomic hydrides. Strong molecular bands of diatomic metal hydrides such as MgH and CrH can be observed in the optical spectrum of ultracool dwarf stars. The MgH band system \\AMgH - \\XMgH can be observed at 0.47--0.6 \\mum, and the CrH band system \\ACrH -- \\XCrH\\  show absorption features at 0.6--1.5 \\mum.  The molecules MgH and CrH have been known in astrophysics for a long time. MgH has been more extensively studied than CrH, because it can be observed in the spectra of G to M type stars. The dissociation energy of MgH is very low (1.285 eV) so lines of this molecule are very sensitive to the temperature and gravity variations in stellar atmospheres. MgH lines were used to determine temperatures in the atmospheres of cool giants (Wyller 1961) and the Sun (Sinha et al. 1979, Sinha \\& Joshi 1982), and for the determination of the surface gravity of stars (Bell \\& Gustaffson 1981, Bell at al. 1985, Berdyugina \\& Savanov 1992, Bonnel \\& Bell 1993).  The pure rotational spectrum of MgH and MgD radicals (\\XMgH) in their ground state $v$=0 and $v$=1 vibrational modes has been studied by Ziurys et al. (1993). The first MgH linelist was computed by Kurucz (1993). Recently, more extensive studies of MgH transitions were performed by Weck et al. (2003, 2003a, 2003b) and Skory et al. (2003).  Although CrH has been known since Gaydon \\& Pearse (1937), its electronic spectrum remained relatively unstudied for many years. Engvold et al. (1980) identified lines of CrH in a spectrum of sunspots as formed by \\XCrH -- \\XCrH transitions. They used the results of studies of multiplicity of $\\Sigma$-terms of CrH by Klehman \\& Uhler (1959) and O'Connor (1967). Later Ram, Jarman \\& Bernath (1993) performed a rotational analysis of 0--0 band of the \\ACrH -- \\XCrH electronic transition and obtained improved rotational constants for the $v'$ =0 vibrational state. Combining these results with those of Bauschlicher et al. (2001) and Lipus et al. (1991) for the vibrationally excited transitions, Burrows et al.(2002) computed an extended linelist for CrH.  Recently Shin et al. (2005) have measured radiative lifetimes of the $v$ = 0,1 levels of \\ACrH state of CrH. These measured lifetimes are about 16\\% -- 45 \\% longer that those obtained by Burrows at al. (2002). These results provide evidence that the oscillator strengths of Burrows et al. (2002) should be corrected by a factor of 0.8 for at least the transitions to the $v'$ = 0.  The submillimeter spectra of CrH and CrD formed by pure rotational transitions in the ground electronic state, have been observed in the laboratory by Halfen \\& Ziurys (2004). Electronic bands of MgD and CrD are likely to be located in the same spectral regions as the corresponding bands of MgH and CrH. In this paper we model the bands of these molecules to analyse the possibility of their use for the determination of the D/H ratio in the atmospheres of ultracool dwarfs.  In section 2 we present a description of the procedures used to compute the molecular bands of CrH, CrD, MgH and MgD. In section 3 we present the vibrational-rotational constants of MgH, MgD, CrH and CrD. In section 4 we present the results of the computation of molecular bands. In section 5 we discuss the possibility of using the electronic bands of diatomic molecules for a realisation of the ``deuterium test''. ",
        "conclusions": "The detection of deuterated molecules in the spectra of ultracool dwarfs provide a challenge for both theoreticians and observers. Indeed, in the atmospheres of planetary mass objects (M $<$ 13 M$_J$) we cannot expect ratio D/H $> 2\\times10^{-5}$. This means the lines of deuterated molecules should be about 5000 times weaker than those of the hydrides. The ideal case would be a spectral region where the molecular bands are not blended. So a crucial requirement is a large difference in the wavelengths of the band heads of the hydrides and deuterides.  CrH appears to be more useful than MgH in the search for deuterated species. The band heads of CrD are displaced significantly from the band heads of CrH. Bands of CrH are observed in the spectra of the latest L dwarfs (Kirkpatrick et al. 1999). The CrD bands are located in the ``near infrared'' spectra, where fluxes are much higher than in the ``optical'' spectral regions.   In this paper we show that the most useful bands for the realisation of the deuterium test are $\\Delta v = +1$  and $\\Delta v = -1$ ($\\lambda$ = 0.795 and 0.968 \\mum, respectively). The $\\Delta v = -1$ band looks especially promising. It is located in the near infrared region with an absence of strong background absorption features. However, portions of this CrD band will be swamped by the 0-0 FeH Wing-Ford band at 0.99$\\mu$m and possibly by water bands. Still, the case for CrD looks better than for HDO lines which are formed on a background of strong H2O lines (Pavlenko 2002). High quality line lists are required to test these possibilities fully.  It is worth noting, that a potential problem lies in the possibility that Cr atoms are absorbed onto dust particles. The depletion of Cr will reduce the strength of both CrD and CrH bands. Fortunately as CrH bands are located in the same spectral region, we can ``scale'' the CrD depletion processes by adjusting the strength of CrH bands.  For more precise studies, more accurate and detailed linelists of CrH and CrD are required. The calculation of such linelists would require new improved computations supported by new laboratory measurements. One problem that concerns us is that even once we have a good agreement with the model and experimental data for the MgH or CrH molecule, there are still perturbations about which we know very little from experiments done so far. Nonetheless, MgH is a non-starter for the deuterium test and as we note above, the model adopted in our paper is more likely to be reliable for CrH and CrD. It is worth noting that the use of the pure rotational-vibrational bands located in the mid and far infrared spectral region may offer an alternative. Indeed, the displacement between CrH and CrD rotation-vibration bands is even larger, than for the case of electronic bands.     Nevertheless, we cannot be certain that we have identified the best candidate systems for the deuterium test. Future investigations of deuterated molecules in different spectral regions are important to determine which offers the best possibility for the realisation of this test. The ideal solution would be to detect lines of deuterated molecule(s) in different spectral regions. This presents a serious observational challenge which can only be met in combination with careful laboratory measurement and the computation of high quality molecular spectra."
    },
    "0710/0710.1472_arXiv.txt": {
        "abstract": "Very recently, J~1128+5925 was found to show strong intraday variability at radio wavelengths and may be a new source with annual modulation of the timescale of its radio variability. Therefore, its radio variability can be best explained via interstellar scintillation. Here we present the properties of its optical variability for the first time after a monitoring program in 2007 May. Our observations indicate that in this period J~1128+5925 only showed trivial optical variability on internight timescale, and did not show any clear intranight variability. This behavior is quite different from its strong radio intraday variability. Either this object was in a quiescent state in optical in this period, or it is intrinsically not so active in optical as it is in radio regimes. ",
        "introduction": "Blazars are the most variable subset of AGN. They show a variety of variability timescales. The longest timescales can be far longer than one year, while the shortest may be less than one hour. The variability with a timescale less than one day is often called intraday variability or IDV, as first reported by \\citet{heeschen84}, \\citet{witzel86}, and \\citet{heeschen87}. Strong IDV phenomena have been observed in the radio domain in a large number of blazars. If interpreted as being source intrinsic, the short-timescale variability would require a very small emitting region and hence a very high apparent brightness temperature of $10^{16}\\sim10^{21}\\,\\rm{K}$, which is far beyond the inverse-Compton limit of about $10^{12}\\,\\rm{K}$ \\citep{keller69}.  Alternatively, the IDV can be explained via extrinsic origin, e.g., via interstellar scintillation (ISS). A strong support to the ISS origin is the so-called annual modulation of the variability timescale, which is the result of the annual changes of the relative velocity vector between the scattering screen and the Earth as the Earth orbits around the Sun \\citep[e.g.,][]{dennett02,dennett03}. Such annual modulation has been observed in a few IDV sources, as mentioned by \\citet{gabanyi07}.  Very recently, the flat-spectrum radio quasar J~1128+5925 was found to show strong IDV at centimeter wavelengths, and its IDV timescale displays an annual modulation \\citep{gabanyi07}. Therefore, its IDV may be caused by ISS. In optical, there is no previous report on its variability. In order to know the properties of its optical variability and to make a comparison to those of its radio variability, we performed an optical monitoring program on this object in 2007 May. Here we present the results. ",
        "conclusions": "We performed an optical monitoring program on J~1128+5925 in the $R$-band from 2007 May 5 to 29. Our monitoring results indicate that in this period J~1128+5925 only showed trivial optical variability on internight timescale, and did not show any clear intranight variability. Either this object was in a quiescent state in optical regimes in this period, or it is intrinsically not as active at optical as it is at radio wavelengths.  Some blazars that show strong IDV in radio regimes also display rapid and strong variability in optical regimes, such as S5~0716+714 \\citep[e.g.,][] {wu05,wu07,montagni06,pollock07}. This doesn't seem to be the case for J~1128+5925. This object exhibits strong IDV at radio wavelengths, but not at optical wavelengths. It is easy to explain this difference if the optical and radio variabilities come from different origins: The optical variability may be intrinsic to the source (the ISS cannot change the optical flux), while the radio variability is mainly the result of ISS, as implied by the observations of \\citet{gabanyi07}.  We present the first report on the optical variability of J~1128+5925 in this paper. However, because of the solar conjunction and observations of other targets with the telescope, our monitoring did not last long. More observations are needed to know whether or not this object is always optically quiescent. Multi-band optical monitoring is also necessary to constrain its optical variability in more detail. Of particular interest is to carry out simultaneous optical and radio monitorings on this object in order to make a more direct comparison between the variabilities at these two wavelengths. Future campaigns can investigate whether there is correlated optical IDV when strong radio IDV is observed. If such correlations are detected, it would be strong evidence that both the optical and radio variability structures are intrinsic to the source, as in the case of S5~0716+714 \\citep{quirren91,wagner95,wagner96}. The broadband variability is also helpful to derive for this object some basic parameters, such as the mass of the central supermassive black hole, the boosting factor of the relativistic jet, etc \\citep[e.g.,][]{fan05}."
    },
    "0710/0710.4582_arXiv.txt": {
        "abstract": "We compare the stellar structure of the isolated, Local Group dwarf galaxy Pegasus (DDO\\,216) with low resolution HI maps from \\cite{young2003}. Our comparison reveals that Pegasus displays the characteristic morphology of ram pressure stripping; in particular, the HI has a ``cometary'' appearance which is not reflected in the regular, elliptical distribution of the stars. This is the first time this phenomenon has been observed in an isolated Local Group galaxy. The density of the medium required to ram pressure strip Pegasus is at least $10^{-5} - 10^{-6}$\\,cm$^{-3}$. We conclude that this is strong evidence for an inter-galactic medium associated with the Local Group. ",
        "introduction": "\\cite{einasto1974} first highlighted that dwarf satellites of large galaxies tend to be gas deficient compared to isolated dwarfs. The former generally have little or no ongoing star formation and the stars are pressure supported (dwarf spheroidal, dSph). The latter generally have ongoing star formation and the gas dynamics show that rotational support is important (dwarf irregular, dIrr). ``Transition'' dwarfs are gas-rich and, unlike dIrr galaxies, have little or no detectable HII regions, although they usually show indications of recent star formation.  The processes by which dwarf galaxies loose their gas are not fully understood. Internal feedback, particularly winds from supernovae, are likely important (\\citealt{dekel1986}) and the existence of the position-morphology relation clearly indicates that environmental influences are significant. \\cite{mayer2006} show that it is possible for dwarf galaxies to be ram pressure stripped of some of their gaseous component in a hot halo of the Milky Way or M31. This idea was originally proposed by \\cite{lin1983}, who calculated the density of the medium required to be of order $10^{-6}$\\,cm$^{-3}$. There have been no direct detections of such a medium, although recently \\cite{nicastro2002,nicastro2003} and \\cite{sembach2003} have detected OVI absorption which they attribute to hot gas associated with either a Milky Way corona or a Local Group medium.  In this {\\it Letter}, we compare the stellar and gaseous structure of the isolated, transition-type, dwarf galaxy Pegasus (DDO216). We show that it displays the characteristic signature of ram pressure stripping and conclude that this is strong evidence for hot gas associated with the Local Group.  Table~1 summarises some of the observed properties of Pegasus. We adopt the distance estimate by \\cite{mcconnachie2005a}, $D \\simeq 919$\\,kpc, derived from the same photometry used in this Letter.  \\begin{table}[htdp] \\begin{center} \\caption{Summary of observed parameters for the Pegasus (DDO216) dwarf galaxy} \\begin{tabular}{lll} Parameter & Value & Reference \\\\ \\hline $\\alpha$ (J2000)  &  23h 28m 36.2s& --- \\\\ $\\delta$ (J2000) & +14$^\\circ$ 44$^\\prime$ 35$^{\\prime\\prime}$& --- \\\\ $(l, b)$ & $(94.8^\\circ, -43.6^\\circ)$ & --- \\\\ $M_V~{\\it(L_V)}$ & -12.9~$(1.24 \\times 10^7\\,L_\\odot)$ & \\cite{mateo1998a} \\\\ $M_{HI}$ & $4.06 \\times 10^6\\,M_\\odot$ & \\cite{young2003}\\\\ $v_\\odot$ & -183\\,km\\,s$^{-1}$ & \\cite{young2003} \\\\ $v_r/\\sigma$ & $1.7$ & \\cite{mateo1998a} \\\\ Distance & 24.82 $\\pm$ 0.07 (919\\,kpc) & \\cite{mcconnachie2005a} \\\\ & 24.4  $\\pm$ 0.2  & \\cite{gallagher1998} \\\\ & 24.9  $\\pm$ 0.1  & \\cite{aparicio1994} \\\\ \\hline \\end{tabular} \\end{center} \\label{distances} \\end{table} ",
        "conclusions": "\\subsection{Comparison of stellar and HI contours}  The top-right panel of Figure~1 shows the tangent-plane projection of the spatial distribution of objects identified as stellar from our INT~WFC observations of Pegasus. The dotted lines in this panel (and the remaining panels of Figure~1) correspond to the approximate edges of each CCD of the INT~WFC. Only objects which lie within $1-\\sigma$ of the stellar locus in both the $V^\\prime$ and $i^\\prime-$band observations are shown. The hole at the center of the main body of Pegasus is due to severe crowding which causes incompleteness. The bottom-left panel of Figure~1 shows a contour map of the density distribution of stars. The first contour is $2-\\sigma$ above the background, and the contours correspond to $2.2, 5.0, 8.6, 13.2, 19,0, 26.3, 35.7, 47.5$ and $62.5$\\,stars\\,arcmin$^{-2}$. The contour map was made in the standard way and follows exactly the methodology described in \\cite{mcconnachie2006b}. This panel shows that Pegasus is significantly more extended than suggested by the image in the first panel.  The bottom-right panel of Figure~1 shows the stellar density distribution as a grey-scale with square-root scaling. The red contours are the low-resolution HI distribution from \\cite{young2003}. The contours correspond to column densities of $0.1, 0.2, 0.4, 0.8, 1.6, 3.2, 6.4$ and $12.8 \\times 10^{20}$\\,cm$^{-2}$. Whereas the stars are distributed in a regular ellipse (typical of a flattened spheroid or an inclined disk) the HI has a ``cometary'' appearance; the contours in the south-east are more closely packed and do not extend as far as in the north-west.  \\subsection{Is Pegasus being ram pressure stripped?}  \\begin{figure*} \\begin{center} \\includegraphics[angle=270, width=8.9cm]{f1a.eps} \\includegraphics[angle=270, width=8.9cm]{f1b.ps} \\includegraphics[angle=270, width=8.9cm]{f1c.ps} \\includegraphics[angle=270, width=8.9cm]{f1d.ps} \\caption{Projections in the tangent plane $\\left(\\xi, \\eta\\right)$ of the structure of the Pegasus dwarf galaxy (DDO216) with the orientation of the field indicated. The dotted lines trace the approximate edges of the four CCDs of the INT~WFC. Top-left panel: The reduced $V^\\prime-$band image of Pegasus taken with the INT~WFC. Also marked are the positions of fields analysed in previous studies of this galaxy; the fields analysed by \\cite{gallagher1998} are shown in green as the largest rectangular field and the small WFPC2 footprint, the field analysed by \\cite{aparicio1994} is shown in blue as the smallest rectangular field, and the field studied by \\cite{hoessel1982} is shown in red as the medium sized rectangle. Top-right panel: The distribution of all sources confidently identified as stellar in both the $V^\\prime$ and $i^\\prime-$bands. The hole at the center of Pegasus is due to severe crowding causing the photometry to become seriously incomplete. Bottom-left panel: The stellar density distribution of Pegasus shown as a contour map. The first contour is $2-\\sigma$ above the background and the 9 contours correspond to $2.2, 5.0, 8.6, 13.2, 19,0, 26.3, 35.7, 47.5$ and $62.5$\\,stars\\,arcmin$^{-2}$. Bottom-right panel: the stellar density distribution is shown as a grey-scale with square-root scaling. The red contours show the low resultion HI distribution from \\cite{young2003}. Contours correspond to $0.1, 0.2, 0.4, 0.8, 1.6, 3.2, 6.4$ and $12.8 \\times 10^{20}$\\,cm$^{-2}$. Also shown are the projected directions to all galaxies within $\\sim 500$\\,kpc of Pegasus which have a significant gaseous content. Pegasus displays the characteristic morphology of ram pressure stripping.} \\label{fig} \\end{center} \\end{figure*}  The shape of the low resolution HI contours in Pegasus - the smooth, compressed contours in the south-east and the ``tail'' to the north-west - is very similar to the simulated morphology of gas undergoing ram pressure striping (e.g., \\citealt{stevens1999,mori2000,marcolini2003,roediger2005,mayer2006}). Observationally, the M81 group dwarf galaxy Holmberg~II is observed to have a similar morphology (\\citealt{bureau2002}), interpreted as evidence of an intra-group medium. In clusters of galaxies, ram pressure stripping of galaxies by an intra-cluster medium is used to explain various observations, including the deficit of HI in cluster spiral galaxies compared to field spirals (e.g., \\citealt{giovanelli1985}). Indeed, several individual galaxies in the Virgo Cluster have been shown to display gaseous morphologies indicative of ram-pressure stripping (\\citealt{vollmer2000,vollmer2004,vollmer2005}).  What else could explain the peculiar appearance of Pegasus? While tidal stripping by large galaxies can affect the structure of dwarf galaxies (e.g., \\citealt{penarrubia2007b}), the closest large galaxy to Pegasus is M31 at $\\sim 470$\\,kpc (all the distance estimates in Table~1 place Pegasus at $> 400\\,$kpc from M31).  Even if we assume Pegasus is a weakly-bound satellite of M31, tidal effects at this distance are minimal. If Pegasus was disrupted at pericenter, it is unlikely that the gas would still show signs of current disturbance. Further, tidal stripping tends to produce symmetrical distortions and both gas and stars should be affected. However, these are inconsistent with the structure of Pegasus that we observe.  Could the appearance of Pegasus be due to internal effects rather than external influences? Enhanced star formation in the south-east of Pegasus could produce winds which remove gas from this region. However, if this is the case then the densely packed contours in the south-east should have a more concave, rather than convex, shape. For example, \\cite{young2007} discuss a gas cloud associated with the Phoenix dwarf galaxy and conclude that it was blown out by supernovae winds based in part on the concave shape of its contours.  An alternative explanation for the HI morphology of Pegasus is that it consists of multiple HI clouds, the sum total of which has a cometary appearance. Figure~6 of \\cite{young2003} is a position-velocity diagram of Pegasus along its major axis. It shows a gradient in velocity and two main concentrations of HI which \\cite{young2003} interpret as two distinct HI clouds. The strength of the secondary feature ($v \\sim -200$\\,km\\,s$^{-1}$) is weaker than the main feature ($v \\sim -180$\\,km\\,s$^{-1}$) and they join at relatively high column density (between the $8 - 16 - \\sigma$ contour levels). An alternative explanation of the data is that the overall velocity gradient is a result of ram-pressure stripping. The velocity difference between the two features may be due to stripped gas leaving a ``hole'' in the distribution, making the secondary feature appear at a higher density than its immediate surroundings (we do not necessarily expect that the column density should smoothly vary over the entire cloud).  Henceforth, we adopt the hypothesis that Pegasus is being ram pressure stripped. Following \\cite{gunn1972}, material will be ram pressure stripped from a galaxy if the density of the surrounding medium, $n_{IGM} \\gtrsim \\left(2\\,\\pi\\,G\\,\\Sigma_{T}\\,\\Sigma_{HI}\\right)/\\left(\\mu v^2\\right)$. $\\Sigma_{T}$ is the total surface density (stars plus gas), $\\Sigma_{HI}$ is the column density of HI and $v$ is the relative velocity of the galaxy to the medium. Thus,  \\begin{eqnarray} n_{IGM} &\\simeq& 3.7 \\times 10^{-6}\\,\\rm{cm}^{-3} \\left(\\frac{\\rm{100\\,km\\,s^{-1}}}{v}\\right)^2 \\nonumber\\\\ &&\\left(\\frac{\\Sigma_{HI}}{0.1 \\times 10^{20}\\,\\rm{cm}^{-2}}\\right)^2~, \\end{eqnarray}  \\noindent where we take the mean particle mass $\\mu = 0.75\\,m_p$ for fully ionized media.  We approximate the Local Group space velocity of Pegasus as $v \\sim \\sqrt{3}\\,\\sigma_{LG} \\sim 100$\\,km\\,s$^{-1}$ where $\\sigma_{LG} \\sim 60$\\,km\\,s$^{-1}$ is the Local Group line-of-sight velocity dispersion (\\citealt{sandage1986}).  HI at a column density much lower than $\\Sigma_{HI} \\sim 0.1 \\times 10^{20}$\\,cm$^{-2}$ has been stripped from Pegasus, implying that this is a reasonable lower limit for use in this calculation. We adopt $\\Sigma_T = \\Sigma_{HI} \\left(1 + M_\\star/M_{HI}\\right)$, where $M_\\star \\sim 1.24 \\times 10^7\\,M_\\odot$ is the stellar mass of Pegasus (Table~1). This seems reasonable; the surface brightness of Pegasus is $25$\\,mags\\,arcsec$^{-2}$ at a radius of $r = 1.5^\\prime$ on the minor axis (\\citealt{nilson1973,devaucouleurs1991}), corresponding to a stellar surface density of $\\Sigma_\\star \\sim 4 \\times 10^{20}$\\,cm$^{-2}$. This is approximately equivalent to the stellar-to-gas mass ratio multiplied by the gas surface density ($M_\\star/M_{HI} \\times \\Sigma_{HI}$) at $r = 1.5^\\prime$.  These values yield $n_{IGM} \\sim 3.7 \\times 10^{-6}$\\,cm$^{-3}$. However, given the uncertainties involved, it is entirely plausible that the value of $n_{IGM}$ could be at least an order of magnitude larger than in Equation~1.  \\subsection{Consequences}  What is the source of the material that is stripping Pegasus? The bottom right panel of Figure~1 shows the distances of Pegasus to its nearest gas-rich neighbours. The dwarf neighbours are unlikely to be the source of the stripping medium; not only is the required mass of gas unrealistically large (an ejected spherical shell $\\sim 1$\\,kpc thick with a radius of $\\sim 300$\\,kpc would have a mass $> 3 \\times 10^6\\,M_\\odot$ at a density of $n_{IGM}$) but the energy required is too large for a dwarf galaxy to reasonably provide.  Alternatively, the gas could be associated with M31. From observations of the Magellanic stream, \\cite{murali2000} estimate that the density of the Milky Way halo at the stream must be $\\lesssim 10^{-5}$cm$^{-3}$, although \\cite{stanimirovic2002} estimate $\\sim 10^{-4}$ cm$^{-3}$. If the gas density in the halo of M31 is similar, then not only would M31 need to have a very extended corona, but its density would need to decrease very slowly with radius. Indeed, if the Milky Way has a similarly extended corona, then the two will overlap and the result may be observationally indistinguishable from a Local Group medium.  The isolation of Pegasus raises the strong possibility that the stripping medium is associated with the Local Group, rather than individual galaxies within the group. Clusters of galaxies have such media, and observations of Holmberg~II imply the presence of an intra-group medium in the M81 group (\\citealt{bureau2002}). The density of the intra-group medium implied in Equation~1 is of the same order as the density of the medium responsible for local OVI absorption detected by \\cite{nicastro2002,nicastro2003} and \\cite{sembach2003}, which they suggest is associated with either a Milky Way corona or a Local Group medium. Our result favors the latter interpretation. Theoretically, $\\sim 30\\,\\%$ of baryons in the Local Volume are expected to be in a warm/hot phase ($T \\sim 10^5 - 10^6\\,$K; \\citealt{kravtsov2002}); this is likely concentrated around galaxies and galaxy groups as an intra-group medium.  If the stripping medium pervades the Local Group, why do more dwarf galaxies not show evidence of ram pressure stripping? \\cite{lin1983} suggest that all the dSphs have been stripped in this fashion, (although \\cite{mayer2006} show that ram-pressure stripping by itself is insufficient to remove all the gas from a dIrr). It is possible that the Local Group medium will be clumpy and perhaps Pegasus is passing through a region of higher density compared to other dIrrs. Alternatively, Pegasus could be falling into and interacting with the Local Group for the first time, as has recently been speculated for two dSph galaxies at large radii from M31 (And~XII, \\citealt{chapman2007}; AndXIV, \\citealt{majewski2007}). However, the reason why only Pegasus currently shows signs of ram-pressure stripping is unlikely to be known until such time as the masses and orbits of the dIrrs have been determined. Given the distances of these galaxies, this will be some time yet."
    },
    "0710/0710.3477_arXiv.txt": {
        "abstract": "We use Monte-Carlo simulations, combined with homogeneously determined age and mass distributions based on multi-wavelength photometry, to constrain the cluster formation history and the rate of bound cluster disruption in the Large Magellanic Cloud (LMC) star cluster system. We evolve synthetic star cluster systems formed with a power-law initial cluster mass function (ICMF) of spectral index $\\alpha =-2$ assuming different {\\dts}s.  For each of these cluster disruption time-scales we derive the corresponding cluster formation rate (CFR) required to reproduce the observed cluster age distribution. We then compare, in a ``Poissonian'' $\\chi^2$ sense, model mass distributions and model two-dimensional distributions in log(mass) vs. log(age) space of the detected surviving clusters to the observations.  Because of the bright detection limit ($M_V^{\\rm lim} \\simeq -4.7$ mag) above which the observed cluster sample is complete, one cannot constrain the characteristic cluster disruption time-scale for a $10^4$ M$_\\odot$ cluster, $t_4^{\\rm dis}$ (where the disruption time-scale depends on cluster mass as $t_{\\rm dis} = t_4^{\\rm dis} (M_{\\rm cl} / 10^4 {\\rm M}_\\odot )^\\gamma$, with $\\gamma \\simeq 0.62$), to better than a lower limit, $t_4^{\\rm dis} \\ge 1$\\,Gyr. \\\\ We conclude that the CFR has been increasing steadily from 0.3 clusters Myr$^{-1}$ 5 Gyr ago, to a present rate of $(20-30)$ clusters Myr$^{-1}$, for clusters spanning a mass range of $\\sim 100-10^7$ M$_\\odot$. For older ages the derived CFR depends sensitively on our assumption of the underlying CMF shape. If we assume a universal Gaussian ICMF, then the CFR has increased steadily over a Hubble time from $\\sim 1$ cluster Gyr$^{-1}$ 15 Gyr ago to its present value. On the other hand, if the ICMF has always been a power law with a slope close to $\\alpha=-2$, the CFR exhibits a minimum some 5 Gyr ago, which we tentatively identify with the well-known age gap in the LMC's cluster age distribution. ",
        "introduction": "\\label{sec:intro}  The mass and age distributions of star cluster systems contain the (fossil) records of their formation conditions. They are therefore among the best tracers of the star-formation histories of their host galaxies available to observers. It is important to realise, however, that one needs to understand both the dominant internal and external evolutionary processes in order to disentangle this formation record, and hence obtain a glimpse of the initial conditions required for star cluster formation.  The effects of stellar evolution in a given star cluster (which can be approximated by a ``simple'' stellar population once the cluster has reached an age that is well in excess of its formation time-scale) are rather well understood, whereas we have only recently begun to make major {\\it quantitative} inroads into understanding the environmental effects leading to star cluster ``weight loss'' (i.e., the preferential depletion of the low-mass component of the cluster's stellar mass function caused by tidal stripping and the ejection of stars owing to internal two-body relaxation) and -- eventually -- disruption.  Estimates of the characteristic cluster disruption time-scales in various star cluster environments have been calculated by Boutloukos \\& Lamers (2003), de Grijs et al. (2003a,b,c), Gieles et al. (2005) and de Grijs \\& Anders (2006), among others (see also Lamers et al. 2005a,b). Specifically, Boutloukos \\& Lamers (2003) and Lamers et al. (2005b) show that the cluster age distribution and cluster mass function approximate a double power law when a cluster system is affected by both fading and secular dynamical evolution. Knowledge of the detection limit in terms of cluster mass vs. age, combined with the estimate of the cluster age or mass at the ``break point'' of the integrated cluster age (mass) distribution (see, e.g., fig. 1 in Boutloukos \\& Lamers 2003, where the ``break points'' are referred to as $t_{\\rm cross}$ and $M_{\\rm cross}$ in the age and mass distributions, respectively) then leads to the typical cluster disruption time-scale. Their analysis, however, explicitly builds on the assumption of a constant cluster formation rate (CFR, i.e. the number of clusters formed per linear time interval ${\\rm d}N/{\\rm d}t$ is constant in time) as a function of time. In that context, the (poorly-known) time-variable CFR of the LMC cluster system hampers such an analysis. If only the observed cluster age distribution is known, then we are left with a degeneracy between the CFR and the disruption time-scale, in the sense that one cannot distinguish between a low CFR combined with slow secular dynamical evolution on the one hand, and a vigourous CFR combined with cluster disruption occurring on a rapid time-scale on the other.  The star cluster system in the Large Magellanic Cloud (LMC) has the potential of providing strong constraints to the theory of star cluster disruption as a function of environment, since it is composed of the largest resolved cluster system spanning {\\it both} a reasonable mass range ($\\sim 10^2 - 10^6$ M$_\\odot$; cf. Hunter et al. 2003, hereafter H03; de Grijs \\& Anders 2006, and references therein) {\\it and} an age range from a few Myr to $\\sim 13$ Gyr available. In addition, thanks to the LMC's proximity, we have been able to obtain observations -- and derived the age and mass distributions -- of a sufficiently large cluster sample to allow a statistical approach to its evolution (e.g., H03; de Grijs \\& Anders 2006; see also Sect. 2).  On the basis of the few star clusters systems analysed in detail to date, including M51 and the Antennae interacting system, Bastian et al. (2005), Fall et al. (2005) and Fall (2006) suggest that the early evolution of star cluster systems is most likely characterised by a rapid, largely mass-independent ``infant mortality'' phase, at least for masses $\\ga 10^4$ M$_\\odot$ (see also de Grijs \\& Parmentier 2007; and references therein), combined with ``infant weight loss''(the loss of stars caused by rapid, early gas expulsion; cf. Weidner et al. 2007), the effects of which are enhanced by stellar evolutionary mass loss. In this scenario, this early phase, which ends when clusters reach an age of $\\sim 40$ to 50 Myr (e.g., Goodwin \\& Bastian 2006), would then be followed by (mass-dependent) secular evolution. The early, rapid cluster disruption process results from the expulsion of the intracluster gas due to adiabatic or explosive expansion driven by stellar winds or supernova activity (Mengel et al. 2005; Bastian \\& Goodwin 2006; Goodwin \\& Bastian 2006; see de Grijs \\& Parmentier 2007 for a review). Star clusters are expected to settle back into virial equilibrium $\\sim 40$ to 50 Myr after gas expulsion (Goodwin \\& Bastian 2006). In our analysis in this paper we will therefore exclude star clusters younger than 50 Myr, since our main purpose is to derive the characteristic (mass-dependent) time-scale of cluster disruption in the LMC driven by secular evolution only. In a follow-up paper (Goodwin et al., in prep.), we will discuss the evolution of the LMC cluster system on the shortest time-scales relevant to the infant mortality and infant weight loss scenarios.  In de Grijs \\& Anders (2006), we found that the LMC's CFR has been roughly constant outside of the well-known age gap between $\\sim 3$ and 13 Gyr, when the CFR was a factor of $\\sim 5$ lower (assuming a roughly constant rate during this entire period). Based on this observation as our main underlying assumption, we used a simple approach to derive the characteristic cluster disruption time-scale in the LMC, for which we found that $\\log(t_4^{\\rm dis} {\\rm yr}^{-1}) = 9.9 \\pm 0.1$, where $t_{\\rm dis} = t_4^{\\rm dis} (M_{\\rm cl}/10^4 {\\rm M}_\\odot)^{0.62}$ (Boutloukos \\& Lamers 2003; Baumgardt \\& Makino 2003; Gieles et al.~2005). We argued that this was consistent with earlier, preliminary work done on a smaller cluster sample: Boutloukos \\& Lamers (2002) found $\\log( t_4^{\\rm dis} {\\rm yr}^{-1} ) = 9.7 \\pm 0.3$ for a smaller sample of 478 clusters within 5 kpc from the centre of the LMC, in the age range $7.8 \\le \\log(\\mbox{age yr}^{-1}) \\le 10.0$. We also considered our result qualitatively consistent with Hunter et al. (2003), who noticed very little destruction of clusters at the high-mass end. This long characteristic disruption time-scale would imply that hardly any of our LMC sample clusters are affected by significant disruptive processes, so that we are in fact observing the {\\it initial} cluster mass function (CMF).  However, a close inspection of fig. 6 of de Grijs \\& Anders (2006) highlights an apparent contradiction. The ``crossing time'', $t_{\\rm cross}$, defined by the crossing point between the best-fitting lines describing the number of clusters per unit time-scale that are mostly affected by fading of their stellar populations and those that are undergoing significant secular disruption, seems to imply that a more appropriate time-scale for the disruption of the LMC cluster system may be of order $\\log( t_4^{\\rm dis} {\\rm yr}^{-1} ) \\simeq 8.9$. Since this implies a downward adjustment of the characteristic cluster disruption time-scale in the LMC by up to an order of magnitude, we decided to re-investigate the LMC's cluster disruption history.  Here, we approach this problem from a different angle, by running a large number of Monte-Carlo simulations in which we vary the cluster disruption time-scale. Meanwhile, for each of these cluster disruption time-scales we derive the corresponding CFR required to reproduce the observed cluster age distribution. We then match the observed cluster mass distribution, integrated over time, and the observed two-dimensional distribution of the detected surviving clusters in the log(mass) vs. log(age) plane to the model results. $\\chi^2$ fit estimates are used to quantify which cluster disruption time-scale and, therefore, which cluster formation history, best describes the presently available data.  This paper is organised as follows. In Section \\ref{sec:data} we justify our choices used in the data analysis leading to the cluster age and mass distributions used in the remainder of the paper. Section \\ref{sec:synth} discusses our basic assumptions in constructing synthetic cluster populations, which we then use in Section \\ref{sec:disr} to explore the range of characteristic cluster disruption time-scales allowed by the data. In Section \\ref{sec:comp} we highlight the importance of properly understanding the data's completeness limit, and use this in Section \\ref{sec:agap} to constrain possible variations in the CFR over time.  In Section \\ref{sec:1vs10}, we assess precisely what we would need to fully and unambiguously constrain $t_4^{\\rm dis}$. Our results and conclusions are summarized in Section \\ref{sec:conc}.  \\begin{figure*} \\begin{minipage}[t]{\\linewidth} \\centering\\epsfig{figure=obs_age_mass.eps, width=\\linewidth} \\end{minipage} \\caption{Distribution of the LMC star cluster sample of de Grijs \\& Anders (2006) in the [$\\log({\\rm age}),\\log(M_{\\rm cl})$] plane. The filled triangles correspond to the vertical dashed lines in the individual panels of Fig.~\\ref{fig:compl} (upright triangles: left-hand column; upside-down triangles: right-hand column). For subsequent cluster age ranges (in steps of 0.25 dex and 0.5 dex wide) they trace the mass limit below which the sample becomes incomplete (see section \\ref{sec:synth} for details). They are therefore considered tracers of the fiducial detection limit (thick dash-dotted line with squares), which corresponds to $M_V^{\\rm lim}=-4.7$ mag (based on the {\\sc galev} mass-to-light ratios for ``simple'' stellar populations). The three thin dash-dotted lines, labelled `[1]', `[2]' and `[3]', are the detection limits corresponding to $M_V^{\\rm lim}=-5.2$, $M_V^{\\rm lim}=-4.2$ and $M_V^{\\rm lim}=-3.5$ mag, respectively.  They are therefore equivalent to the thick dash-dotted line shifted vertically by, respectively, $\\Delta \\log (M_{\\rm cl})=0.2, -0.2$, and $-0.48$. The lower dash-dotted curve (`[3]') is the $M_V^{\\rm lim}=-3.5$\\,mag fading limit of H03. The thick dashed lines represent the cluster disruption limits for $\\log (t_4^{\\rm dis} {\\rm yr}^{-1})=8.1, 9.0$ and $9.9$ (labelled `[a]', `[b]' and `[c]', respectively). The age range on which we focus in this paper is bracketed by vertical solid lines, at $\\log (\\mbox{age yr}^{-1})=7.7$ and $\\log (\\mbox{age yr}^{-1})=9.2$; clusters brighter than $M_V^{\\rm lim}=-4.7$ mag in that age range are represented by crosses} \\label{fig:obs_am} \\end{figure*}  \\begin{figure*} \\begin{minipage}[t]{\\linewidth} \\centering\\epsfig{figure=MF_agebin_completeness.eps, width=\\linewidth} \\end{minipage} \\caption{Observed cluster mass functions for the age ranges included at the top of each panel. In each panel, the vertical dashed line is the mass limit bracketing 25 and 75 per cent of the cluster subsample on either side. This is a good proxy to the cluster mass at the turn-over of each CMF. For each age range, turn-over masses are indicated in ($\\log (\\rm age),\\log (M_{\\rm cl})$) space as filled triangles in Fig.~\\ref{fig:obs_am}. The mass limits defined by the dashed lines evolve with time following a line of constant luminosity, at $M_V=-4.7$ mag. This implies that the decrease in cluster numbers observed for each CMF below its turn-over mass (i.e. below the vertical dashed line) is mostly driven by incompleteness effects} \\label{fig:compl} \\end{figure*} ",
        "conclusions": "\\label{sec:conc}  In this paper, we have carried out numerous detailed Monte-Carlo simulations aimed at constraining the cluster formation history and the rate of bound cluster disruption in the LMC star cluster system. We considered only clusters older than 50\\,Myr in order not to determine erroneously short cluster disruption time-scale as a result of the inclusion of the infant mortality and infant weight loss evolutionary phases. Our data are based on the $UBVR$ photometry of H03 (obtained from Massey's (2002) survey of the Magellanic Clouds), for which de Grijs \\& Anders (2006) obtained homogeneously determined age and mass estimates. We estimated a fiducial detection limit above which the cluster sample is (fairly) complete from the CMF as a function of age, at $M_V^{\\rm lim} = -4.7 {\\rm mag}$.  This is significantly brighter than H03's fading limit. We consider only clusters brighter than this limit, in order to avoid severe statistical incompleteness effects.  We have evolved synthetic star cluster systems characterised by constant CFRs assuming 20 different {\\dts}s. The CFR was adjusted to reproduce the observed age distribution. By doing so, we are likely to loose sensitivity to CFR variations occurring on a time-scale shorter than the width of the age range corresponding to each logarithmic age bin (i.e., more likely at old age where our age bins span greater linear age ranges than at young age). However, the general behaviour of the CFR, as well as the CFR averaged over age are robustly recovered (see Section \\ref{sec:1vs10}). We then compare, in a ``Poissonian'' $\\chi^2$ sense the modelled mass distribution and the modelled [$\\log({\\rm age}),\\log(M_{\\rm cl})$] distribution to the observations. We show that because of the bright detection limit at $M_V^{\\rm lim}= -4.7$ mag, one cannot constrain $t_4^{\\rm dis}$ to better than a lower limit, $t_4^{\\rm dis} \\ge 1$\\,Gyr. The tightest constraints are set by the CMF integrated over age. The $\\chi^2$ test applied to the distribution of points in [$\\log({\\rm age}),\\log(M_{\\rm cl})$] space turns out to be a poor diagnostic tool. This is probably related to the low density of points in each cell of the [$\\log({\\rm age}),\\log(M_{\\rm cl})$], compared to the density of points in each bin of the integrated CMF, leading to smaller Poissonnian error bars and therefore a better constrained model in the latter case. Our range of {\\dts} estimates is robust with respect to model variations, such as of the upper limit of the initial cluster mass range, the location of the grid in [$\\log({\\rm age}),\\log(M_{\\rm cl})$], and the size of cells in this grid.  We have shown that should the detection limit be underestimated, artificially shortened {\\dts}s would result. This is so because there is a degeneracy between incompleteness and secular evolution, i.e., the fading-driven turn-over in a CMF in a given age bin is interpreted as resulting from secular evolution, leading to a shortened cluster disruption time-scale.  Having set the best possible constraints on $t_4^{\\rm dis}$, we explored the corresponding CFR, in particular considering $t_4^{\\rm dis}=1$ and 10 Gyr.  The CFR has been increasing steadily from about 0.3 clusters Myr$^{-1}$ 5 Gyr ago, to a present rate of $(20-30)$ clusters Myr$^{-1}$, for clusters spanning an initial mass range of $\\sim 100-10^7$ M$_\\odot$. The CFRs inferred for both disruption time-scales differ by at most a factor of three (Fig.~\\ref{fig:ad_all}). At older ages however, the situation becomes unclear. The uncertainty in the CFR as a result of the uncertainty on $t_4^{\\rm dis}$ increases. In addition, the overall temporal behaviour of the CFR depends on the shape of the ICMF of the oldest, globular cluster-like objects. If this is the universal Gaussian ICMF, then the CFR has increased steadily over a Hubble time from $\\sim 1$ cluster Gyr$^{-1}$ 15 Gyr ago to its present value. On the contrary, if the ICMF has always been a power law with a slope close to $\\alpha=-2$, the CFR exhibits a minimum some 5 Gyr ago. Our results may be related to the orbital history and dynamics of the LMC with respect to both the SMC and the Milky Way, although this remains poorly constrained because of a lack of proper motion measurements with the required accuracy (Besla et al.~2007).  Additionally, we note that interactions between the Clouds and between the Clouds and the Galaxy, while affecting their star-formation history, also affect the {\\it structure} of the Magellanic Clouds (Bekki \\& Chiba 2005). This may have induced temporal variations in the {\\dts} over the past Hubble time, rendering the {\\dts} estimate at old age even more uncertain.  Finally, we have investigated which strategy should be adopted in the future in order to better constrain the {\\dts} in the LMC. Specifically, we have generated synthetic cluster populations defined by a given cluster formation history, ICMF and various dissolution rates.  After the addition of Gaussian noise to mimic an observational situation, we processed these simulated data sets in the same way as the actual LMC data. We demonstrate that {\\it if} the {\\dts} is known, then the CFR can be derived accurately. We confirm our inability to distinguish $t_4^{\\rm dis}=1$ Gyr from 10 Gyr because of the bright detection limit. With such a bright detection limit, the expected turn-over in the CMF caused by dynamical evolution is not detected, for any cluster age range. As a result, only a lower limit to the {\\dts} can be retrieved, i.e., we can exclude all $t_4^{\\rm dis}$ that are short enough to lead to a turn-over above the detection limit. To obtain age and mass estimates for an LMC star cluster sample complete above $M_V^{\\rm lim}=-3.5$ is desirable to more easily distinguish between $t_4^{\\rm dis}=1$ Gyr and $t_4^{\\rm dis}=10$ Gyr."
    },
    "0710/0710.3180.txt": {
        "abstract": "We present new optical spectroscopy and photometry, 2MASS infrared observations and 24 years of combined AAVSO and AFOEV photometry of the symbiotic star candidate \\ae. The long-term light curve is characterized by outbursts lasting several years and having a slow decline of $\\sim 2 \\times 10^{-4}$ mag/day. The whole range of variability of the star in the $V$ band is about 4 magnitudes. The periodogram of the photometric data reveals strong signals at $\\sim$ 342 and 171 days. The presence of the emission feature at $\\lambda$ 6830 \\AA~ at minimum and the detection of absorption lines of a $\\sim$ K5 type star confirm the symbiotic classification and suggest that AE\\,Cir is a new member of the small group of s-type yellow symbiotic stars. We estimate a distance of 9.4 kpc. Our spectrum taken at the high state shows a much flatter spectral energy distribution, the disappearance of the $\\lambda$ 6830 \\AA~  emission feature and the weakness of the He\\,II 4686 emission relative to the Balmer emission lines. Our observations indicate the presence of emission line flickering in time scales of minutes in 2001. The peculiar character of \\ae~ is revealed in  the visibility of the secondary star at the high and low state, the light curve resembling a dwarf nova superoutburst and the relatively short low states.  The data are hard to reconciliate with standard models for symbiotic star outbursts.    %Another peculiarity is the detection of features of the cool star at the low and high states.  %We discuss the multi-wavelength data and suggest that the ellipsoidal-like variability is due to %fluorescence of the cool star atmosphere that is irradiated by high energy photons from the hot component. %We observe a 40-day length eclipse-like episode in %our photometry, that we interpret as self-occultation of the irradiated hemisphere of the red star. Then, assuming a white dwarf %hot component of 1 M$\\odot$ we find that the red giant probably fills their Roche lobe and has a radius $R_{g} \\approx$ 100 R$_{\\odot}$. %From independent methods we have found a distance to AE\\,Cir of 14 $\\pm$ 3 kpc. %Our observations suggest that accretion onto a compact object %in a semi-detached binary could play a major role in the high states and short-term variability observed in AE\\,Cir. ",
        "introduction": "AE\\,Cir (S32 = StHA32, $\\alpha_{2000}$ = 14:44:52.0, $\\delta_{2000}$ = -69:23:35.9) was classified as a RCB star in the Fourth General Catalog of Variable Stars. This classification was rejected by Kilkenny (1989), who found strong emission lines of H\\,I and weaker emission lines of He\\,II and other species in one spectrogram spanning a wavelength range of 3400-5100 \\AA. Kilkenny noted the lack of forbidden lines and the anomalous strong He\\,II 4686 emission for this object, and suggested a symbiotic classification, although no lines of the cool component were detected. He also mentions the photometric variability of the star, ranging from 12 to 14 mag. (from visual estimates) in hundreds of days. $B,V$ photoelectric measures by Lawson \\& Cottrell (1990) taken in an interval of 107 days show the star with $V$ magnitude between 13.52 and 14.51 and color $B-V$ between 0.95 and 1.43, the object being  redder when fainter. Based on the spectroscopy by Kilkenny, AE\\,Cir was listed as a suspected symbiotic in the catalog of Belczynski et al. (2000). Symbiotic stars have been reviewed recently by Miko{\\l}ajewska (2007). In this paper we present new photometry and low and high-resolution spectra of AE\\,Cir and investigate their emission line properties and symbiotic nature.  In Section 2 we give details of our spectroscopic observations and show the methods used in data reduction and analysis. In that section we also analyze available visual long-term photometric records along with our own CCD photometry. In Section 3 we analyze our data, giving the main results of our research. In Section 4 we discuss our results in the context of symbiotic stars and present our conclusions in Section 5. ",
        "conclusions": "In this paper we have investigated the nature of the symbiotic star candidate AE\\,Cir. We have analyzed new optical photometric and spectroscopic data, 2MASS infrared photometry and 24 years of  visual photometry in an integrated way in order to get a better understanding of the system. Our conclusions can be summarized as follows:\\\\  \\begin{itemize}  \\item The symbiotic nature is confirmed. This result is based on the detection of the $\\lambda$ 6830 \\AA~ emission feature and the spectral signatures of a cool stellar component of spectral type $\\sim$ K5.  \\item The spectral type $\\sim$ K5 and the spectral and photometric properties indicate that \\ae~ is a member of the small group of yellow symbiotic stars of the s-type.  \\item The light curve is characterized by outbursts lasting $\\sim$ 4000 days and overall amplitude of variability about 4 magnitudes. The outbursts show a slow decline of $\\sim 2 \\times 10^{-4}$ mag/day. The duration  of the low state is about 38\\% the high state.  \\item A strong signal at 342 $\\pm$ 15 days is detected in the light curve. The light curve folded with this period shows two broad minima with different amplitude.  \\item The spectrum in the low state is characteristic of a symbiotic star without forbidden lines and very strong He\\,I\\,4868 emission. At maximum the emission feature at $\\lambda$ 6380 \\AA~ disappears, the spectrum becomes flat and the relative intensity of the He\\,I 4686 \\AA~line, compared with the Balmer lines, becomes lower.  At the same time the overall emission line spectrum shows smaller equivalent widths. Signatures of the cool companion are observed at the high and low state.  \\item We observe H\\,I, He\\,I and He\\,II emission line variability in time scales of minutes in 2001 being larger in the line centers. The Na\\,D and $\\lambda$ 6820 \\AA~ lines do not follow this variability. We also observe larger broadening and asymmetry in the Balmer lines compared with others emission lines in this epoch.  \\item The data suggest that full visibility of the red giant should occur at the low state at $V \\sim 15.5$. For a typical K5 giant behind the galactic disc this would indicate a distance of 9.4 kpc. This should indicate that \\ae~  is about 1.4 kpc below the galactic plane.  \\item We suggest that an eclipse-like event could be interpreted as self-occultation of the irradiated red giant hemisphere. In this case we found that it is possible that the red giant fills its Roche-lobe. Assuming a white dwarf hot component of 1 M$_{\\odot}$, we estimate a radius of 99 R$_{\\odot}$, a mass of 1.1 M$_{\\odot}$ and M$_{V}$ = -2.4 for the red giant at the epoch of the SMARTS observations.   \\item The atypical character of \\ae~ is revealed in the long time passing in outburst and the dwarf nova superoutburst shape of the light curve. Another atypical feature is the rather strong TiO absorption band around $\\lambda$ 6300 \\AA~ at the high state.  This could indicate that the red giant and the hot component are partially obscured at minimum or the brightness of the secondary star is not constant. Obscuration hardly explains the large line emission observed at minimum and the long brightness decay after maximum whereas variations of the red giant brightness imply unrealistic changes in the stellar radius.  \\item  At present, the available data for AE\\,Cir are hard to reconciliate with canonical models for symbiotic outbursts.  %\\item {\\bf The similar visibility of the secondary star at the low and high state probably rules out the %hypothesis of the hot component outbursts and suggests an scenario where the %outbursts are due to thermal instabilities in an accretion disc fed by a Roche lobe filling giant star. %In this case the low state spectrum is dominated by the giant star and emission lines arising from wind accretion. %On the contrary, the hot state is characterized by a blue continuum arising from the hot and luminous  disc in outburst %and a diminishing of the spectral features arising from the accreted wind. Irradiation of the secondary star by the %outbursting disc, and eventually stellar pulsation, could explain its variable brightness. The disappearance of the $\\lambda$ %6830 \\AA~ feature at the high state could be explained if the O\\,VI photon forming region is occulted by the accretion disc %during outburst.}    \\end{itemize}"
    },
    "0710/0710.2317_arXiv.txt": {
        "abstract": "The goal of this research is to investigate how well various turbulence models can describe physical properties of the upper convective boundary layer of the Sun. An accurate modeling of the turbulence motions is necessary for understanding the excitation mechanisms of solar oscillation modes. We have carried out realistic numerical simulations using several different physical Large Eddy Simulation (LES) models (Hyperviscosity approach, Smagorinsky, and dynamic models) to investigate how the differences in turbulence modeling affect the damping and excitation of the oscillations and their spectral properties and compare with observations. We have first calculated the oscillation power spectra of radial and non-radial modes supported by the computational box with the different turbulence models.  Then we have calculated the work integral input to the modes to estimate the influence of the turbulence model on the depth and strength of the oscillation sources. We have compared these results with previous studies and with the observed properties of solar oscillations. We find that the dynamic turbulence model provides the best agreement with the helioseismic observations. ",
        "introduction": "Dominant acoustic sources within the Sun are generated by strong fluctuations in the outer convective layers. Turbulent motions stochastically excite the resonant modes via Reynolds stresses and entropy fluctuations. The dominant driving comes from the interaction of the nonadiabatic, incoherent pressure fluctuations with the coherent mode displacement \\citep{Nordlund2001}.  The modes excitation sources occur close to the surface, mainly in the intergranular lanes and near the boundaries of granules \\citep{Stein2001}. Thus an accurate modeling of the turbulence motions is necessary to understand the excitation mechanisms of solar oscillation modes. The correct choice of turbulence model is also important in many other astrophysical simulations.  The objective of this research is to study the influence of turbulence models on the excitation mechanisms by means of realistic numerical simulations. We have compared different physical Large Eddy Simulation (LES) models (Hyperviscosity approach, Smagorinsky, and dynamic models) to show the influence on the damping and excitation of the oscillations.  The organization of this paper is as follows. In \\S2, we describe the main lines of the code and the different turbulence models. The kinetic energy of the radial modes obtained with the different turbulence models are presented in \\S 3. Then a comparison of the results obtained with the different turbulence models for the non-radial modes is given in \\S 4. The work integral input to the modes is calculated in \\S5 in order to estimate the influence of the turbulence models on the depth of the oscillation sources. ",
        "conclusions": "The goals of this research was to investigate how well various turbulence models can describe the convective properties of the upper boundary layer of the Sun and to study the excitation and damping of acoustic oscillations. Results obtained with the hyperviscosity approach have been compared with those obtained with the Smagorinsky and dynamic turbulence models. We have seen that the dissipation is very high with the Smagorinsky model while the hyperviscosity approach and dynamic modes give similar results. Besides we find that the dynamic turbulence model provides the best agreement with observations."
    },
    "0710/0710.5701_arXiv.txt": {
        "abstract": "{} {We present the period analysis of unfiltered photometric observations of V5116~Sgr (Nova Sgr 2005 \\#2) and we search for superhump candidates in novae remnants.} {The PDM method for period analysis is used. The masses of the novae componets are estimated from the secondary mass -- orbital period and primary mass -- decline time relations.} {We found that 13 nights of V5116~Sgr observations in the year 2006 are modulated with a period of $0.1238 \\pm 0.0001$~d ($2.9712 \\pm 0.0024$~h). Following the shape of the phased light curves and no apparent change in the value of the periodicity in different subsamples of the data, we interpret the period as orbital in nature. The binary system then falls within the period gap of the orbital period distribution of cataclysmic variables. From the maximum magnitude -- rate of decline relation, we estimate the maximum absolute visual magnitude of $M_{\\rm Vmax} =  -8.85 \\pm 0.04$~mag using the measured value of decline $t_{\\rm 2} = 6.5 \\pm 1.0$~d. The mass-period relation for cataclysmic variables yields a secondary mass estimate of about $0.26 \\pm 0.05~{\\rm M}_{\\rm \\odot}$. We propose that V5116~Sgr is a high inclination system showing an irradiation effect of the secondary star. No fully developed accretion disc up to the tidal radius with the value lower than $3.5~10^{10}$~cm is probable. The mass ratio was estimated in a few novae and the presence or absence of superhumps in these systems was compared with the mass ratio limit for superhumps of about 0.35. We found that in the majority of novae with expected superhumps, this variability has not been found yet. Therefore, more observations of these systems is encouraged.} {} ",
        "introduction": "Novae are a subclass of cataclysmic variable stars. In these interracting binaries, the white dwarf is accreting the matter transfered from the secondary star. The accretion disc may be formed in the non magnetic case. The intermediate polar systems have a truncated disc and polar systems have magnetic field strong enough to prevent the disc formation (see Warner 1995 for review). The accumulation of critical amount of accreted material onto the white dwarf surface results in a nova explosion. The distribution of orbital periods in cataclysmic variables shows a period gap between about 2 and 3 hours. Novae do not show this lack of objects in the mentioned interval.  V5116~Sgr (Nova Sgr 2005 \\#2) was discovered by Liller (2005) on 2005 July 4.049 UT. The nova had a magnitude $\\sim 8.0$ on two red photographs. An unfiltered CCD image from 2005 July 5.085 UT showed the object at mag 7.2. The spectrum from 2005 July 5.099 UT showed H$_{\\rm \\alpha}$ with the FWHM of $\\sim 970$ km~s$^{-1}$. The expansion velocity derived from the sharp P~Cyg profile was $\\sim 1300$ km~s$^{-1}$. The position of the nova was measured by Gilmore and Kilmartin (2005) and Jacques (2005). Gilmore and Kilmartin (2005) searched for the nova precursor, but no convincing candidate has been found. Russell et al. (2005) performed a 0.8 -- 2.5 $\\mu$m spectroscopy of the nova on 2005 July 15. The object showed emission lines of H~I, He~I, C~I, N~I, Ca~II and O~I with a FWHM $\\sim$ 2200 km~s$^{-1}$. He~I showed P~Cyg profile at 1.0830 and 2.0581 $\\mu$m. No thermal dust emission was observed.  After the nova explosion the accretions disc is destroyed. The new disc is reformed by the stream of matter flowing from the secondary, interracting with itself and forming a ring in the circularisation radius. The disc starts to form by the viscous shearing. Matter losing angular momentum is moving invard and the excess of angular momentum is transported by the matter flowing outward. The invard moving matter touchs the white dwarf and the disc is reformed (Pringle 1981). If the white dwarf is magnetic, the matter interracts with the magnetosphere and is then conducted by the magnetic field to the magnetic poles. The interraction of the flow with the poles of the rotating star is observed as periodic signal with the spin period of the white dwarf. In the case of intermediate polars the spin period is usually much shorter than the orbital period (Patterson 1994, Hellier 1996). The polar systems without a disc are synchronous rotators, hence the spin period is equal or close to the orbital period (see e.g. Schmidt and Stockman 1991). Nova V1500 Cyg (polar system without a disc) changed the period from 0.141~d to 0.138~d and then stabilised at 0.140~d (Patterson 1978, 1979). The difference of 1.8\\% between the rotation period of the white dwarf and the orbital period is ascribed to the effects of the nova explosion in 1975. The synchronisation can be then corrupted by the nova explosion, but the spin period remain still very close to the orbital period. This is in contrast with intermediate polars with disc presence.  Searching for periods in novae allows to study the orbital distrubution and evolution of these systems. Currently, there are about 50 novae with known orbital periods (Warner 2002) with typical values ranging from 2 to 9 hours. The existence of the accretion disc or its renovation after the nova explosion is confirmed by the detection of the superhump period (see e.g. Retter et al. 1997, Kang et al. 2006a) or the spin period of a magnetic white dwarf in the case of intermediate polars (see e.g. Retter et al. 1998).  Superhumps are caused by a precessing accretion disc generally in systems with mass ratio $M_{\\rm 2}/M_{\\rm 1} < 0.35 \\pm 0.02$ (see Patterson et al. 2005 for review), in which the disc radius reaches the 3:1 resonance. The mass ratio indicates that systems with massive primaries and low mass secondaries are probable superhumpers. Novae in general possess high mass white dwarfs (see e.g. Warner 1995, Smith and Dhillon 1998, Webbink 1990) and short orbital periods sugest low mass secondaries. It is therefore meaningfull to search for superhump variability in novae with short orbital period.  We have an ongoing program to observe novae with small telescopes to search for periodicities in their optical light curves. In this paper we report the detection of one periodicity in our photometric data ($P = 0.1238 \\pm 0.0001$~d) of V5116~Sgr and we discuss the presence of superhumps in nova remnants. In Section~\\ref{obs} we present our observational material. The long-term light curve with the period analysis of the data is presented in Sec.~\\ref{data_anal} and in Sec.~\\ref{disc} we discuss the long-term light curve behaviour ({Sec.~\\ref{disc_v5116_l}}), the results of the period analysis (Sec.~\\ref{disc_v5116_s}) and superhump search in nova remnants (Sec.~\\ref{disc_SH}). ",
        "conclusions": "\\label{disc}  \\subsection{Long-term variations of V5116~Sgr} \\label{disc_v5116_l}  The long-term light curve of the nova V5116~Sgr (Fig.~\\ref{curve} -- bottom panel) shows a transition from smooth decline to probable oscillations. Several models has been suggested for such ``transient'' phase. One of them explains the transition as the time when the accretion disc is re-established (Retter 2002). The author propose a connection between this phase and intermediate polars. There is no indication of magnetic nature of the white dwarf in V5116~Sgr yet. Spin period of the white dwarf indicating the intermediate polar type of this ``transient'' phase in novae was detected for example in V4745~Sgr by optical observations (Dobrotka et al. 2006a), V1494~Aql (Drake et al. 2003) and V4743~Sgr by X-ray observations (Ness et al. 2003). If this ``transient'' phase interpretation of Retter (2002) is applicable in the case of V5116~Sgr, the accretion must be then re-established and the accretion disc should be present.  \\subsection{Short-term variations of V5116~Sgr} \\label{disc_v5116_s}  We have identified one periodicity in the light curve of the nova V5116~Sgr about 15 months after the maximum brightness. The value is $P = 0.1238 \\pm 0.0001$~d ($2.9712 \\pm 0.0024$~h). The upper limits of the period difference between two subset of data with $\\simeq 16$ days of mean distance analysed in Fig.~\\ref{power2} is $1.5~10^{-4}$~d which gives $|\\dot{P}| \\simeq 0.94~10^{-5}$. Following Patterson et al. (1993) the mean variation of the superhump period in SU~UMa superhumping sytems is $|\\dot{P}| \\simeq 3-10~10^{-5}$ (see their Table.~1). For the recurrent nova VY~Aqr the authors derived a variation of $|\\dot{P}| \\simeq 8.2~10^{-5}$. Our value of $|\\dot{P}|$ comes probably from the errors of period measurements rather than from real period changes. The period seems then to be constant during our observations and the shape of the folded light curve suggests a primary and a secondary eclipse. We therefore propose that this periodicity is the orbital period of the binary system. Such a period is at the lower end of the mostly populated region of orbital periods in novae (Warner 2002). The dominant first harmonic frequency in the power spectra is a result of the clear structure primary -- secondary eclipse in the folded light curve which suggests a high inclination angle of the binary system. Using the orbital period of the system and equation (9) from Smith and Dhillon (1998) we obtain a rough estimate for the secondary star mass of $0.26 \\pm 0.05$~M$_{\\rm \\odot}$. Using a mean white dwarf mass of $0.85 \\pm 0.05$~M$_{\\rm \\odot}$ from Smith and Dhillon (1998), we find a mass ratio (secondary/primary mass) of $M_{\\rm 2}/M_{\\rm 1} = 0.3 \\pm 0.1$.  After the nova explosion, the hot white dwarf may heat and irradiate the cooler companion. The observed orbital light curve of the nova can be the result of the aspect variations or eclipses of the secondary due to heating from the hot primary and the asymmetry in the pulse profiles could be produced once the shape of the secondary is of a tear drop model. The irradiation effects in classical novae can also be detected long after the outburst stage (e.g., V1500~Cyg; Sommers and Naylor 1999, DN~Gem; Retter et al. 1999). Two eclipse like features are present in the folded light curve of V5116~Sgr (Fig.~\\ref{folded}). The shape is similar to the light curve of V2540~Oph in 2003 (Ak et al. 2005) with a dip at phase $\\sim 0.5$ ($\\sim 33\\%$ amplitude of the primary minimum for V2540~Oph and $\\sim 70\\%$ for V5116~Sgr). The authors concluded that V2540~Oph is likely a high inclination system either showing an irradiation effect or having a spiral structure in its accretion disc. Woudt and Warner (2003a) noted that one of the following requirements must be fulfilled for the large amplitude orbital modulations to be seen in the light curve of a recent nova in which the accretion disc does not dominate the luminosity of the system: 1) the disc is foreshortened but the secondary is seen (a high inclination angle), 2) the disc has small dimensions, 3) no disc (case of polars). In the case of V5116~Sgr the polar interpretation is possible because of the orbital period distribution of polars which peaks below 5 hours (Warner 1995), but nothing else indicate this option. The small dimension of the disc is supported by; 1) the short orbital period, 2) the post nova stage when the disc is reforming after explosion.  A comparison to two novae within the period gap IM~Nor and DD~Cir with present irradiation effect (Woudt and Warner 2003b,c) can be made. The light curve of V5116~Sgr is different from those of IM~Nor and DD~Cir but similar to V2540~Oph. IM~Nor and DD~Cir showed very small dips at phase 0.5 interpretted as partial eclipses of the irradiated secondary by the disc or matter stream. The light curve shape depends on the disc radius and on the inclination angle. The deep secondary eclipses in V5116~Sgr require a large disc or a high inclination. The strength of the irradiation effect (white dwarf post-outburst temperature) can also play a significant role. The differences in the phase of the maxima in Fig.~\\ref{folded}b,c can be due to spiral structures in the disc as mentioned before in regard with the nova V2540~Oph.  The reconstruction of an accretion disc after the nova eruption is indicated by the discovery of the white dwarf spin or by the superhump period. A probable spin period of the white dwarf was detected $\\sim 1$ year after the outburst in V4745~Sgr (Dobrotka et al. 2006a), $\\sim 2.75$ years in V4743~Sgr (Kang et al. 2006b) and $\\sim 15$ months after the maximum in V1425~Aql (Retter et al. 1998). Several systems show superhumps as early as two and a half months after the outburst like V4633~Sgr (a spin period is another option in this case, Lipkin et al. 2001) or two years after the outburst like V1974~Cyg (Retter et al. 1997). In our V5116~Sgr light curves extending $\\sim 470$ days after outburst we did not find any photometric variations consistent with white dwarf spin modulation or superhump properties. We can not say anything about the magnetic nature of the white dwarf, but the components mass ratio indicates that the superhump existence is probable. The mass ratio using the primary mass average is not decisive. Using the derived secondary mass 0.26~M$_{\\rm \\odot}$ (Section~\\ref{disc_v5116_s}) the mass ratio is lower than 0.35 for primary mass higher than $0.74$~M$_{\\rm \\odot}$. Taking all known primary masses in cataclysmic variables (Ritter and Kolb 2003), 63\\% have higher or equal mass than $0.74$~M$_{\\rm \\odot}$. This probability value is not enought to make sure that superhumps are expected inV5116~Sgr, but the search for such variability could be fruitful. However V5116~Sgr is a very fast nova. According to equation (13) from Livio (1992) and taking $t_{\\rm 3} = 20.2 \\pm 1.9$~d derived in this paper (Section~\\ref{t2t3}) we obtain a mass estimate of the white dwarf of $1.04 \\pm 0.02$~M$_{\\rm \\odot}$ and thus $M_{\\rm 2}/M_{\\rm 1} = 0.25 \\pm 0.05$. The presence of superhumps is then expected if a disc is fully developed up to tidal radius.  The superhump period is a few percent longer or shorter than the orbital period. The possibility that the periodicity found in this paper is a superhump is rejected following the stability discussion and the shape of the folded light curve (Fig.~\\ref{folded}). The mean shape of superhumps is typically an asymetric sinusoid and our data show typical eclipse like features.  The presence of the disc in cataclysmic variables depends on the mass loss from the secondary. The matterial supplied from the secondary within the period gap depends on the strength of the magnetic braking driving the secondary out of thermal equilibrium (mass loss time scale shorter than the thermal time scale). The stronger the braking, the wider the gap will be and the higher is the upper end of the period gap. If the magnetic braking is low enought, the mass loss time scale may never become shorter than the thermal time scale. In the case of novae another condition is important. The nova explosion heats the secondary which leads to an enhanced mass transfer. Therefore the complete absence of accretion discs in novae within the period gap is not expected.  \\subsection{Searching for superhumps in other nova remnants} \\label{disc_SH}  The orbital period 2.462~h of the mentionned IM~Nor (Woudt and Warner 2003b) yields a secondary mass estimate of 0.20~M$_{\\rm \\odot}$ following Smith and Dhillon (1998). The decay time $t_{\\rm 3} \\simeq 50$~d (Kato et al. 2002) yields a primary mass estimate of $\\simeq 0.86$~M$_{\\rm \\odot}$ (Livio 1992) giving a mass ratio $\\simeq 0.23$. Woudt and Warner (2003b) concluded that they did not observe the superhump modulation. This conclusion with the mass ratio safely lower than 0.35 indicate that the existence of a fully developed disc is not probable. DD~Cir (Woudt and Warner 2003c) has a period of 2.339~h yielding a secondary mass estimate of 0.18~M$_{\\rm \\odot}$. The decay time $t_{\\rm 2} \\simeq 4.5$~d (Liller 1999, no other estimates are available) place the object in the class of very fast novae (similar to V5116~Sgr). The decay time $t_{\\rm 3} \\simeq 10$~d following the equation $t_{\\rm 3} \\simeq 2.75~t_{\\rm 2}^{0.88}$ from Warner (1995) gives a rought white dwarf mass estimate of $\\simeq 1.16$~M$_{\\rm \\odot}$ (mass ratio $< 0.2$). Woudt and Warner(2003b) interpretted the deep eclipse as obscuration of the disc and did not detect any periodicity near the orbital period in the period analysis. Therefore, following the tidal instability model it is again probable that the disc is not fully developed up to 3:1 resonance radius 3 years after the outburst. However the authors argued that the disc radius is 47 \\% of the orbital separation. Using the third Kepler law we obtain a disc radius of $\\sim 3.1~10^{10}$~cm. The disc in DD~Cir is then large enough to reach the 3:1 resonance radius calculated from the equation (3.39) from Warner (1995); $r_{\\rm 3:1}\\sim 3~10^{10}$~cm. The result is marginally in contrast with the absence of superhumps. Taking the derived values for V5116~Sgr we get $r_{\\rm 3:1}\\simeq 3.5~10^{10}$~cm as an upper limit of the disc radius taking the absence of superhumps into consideration. The boundaries of the eclipse in our light curve are not so clear as in the DD~Cir case, therefore not suitable to estimate the disc radius.  Comparison to other systems follows the same way as in the case of IM~Nor and DD~Cir. We took components masses from the Catalogue of Ritter and Kolb (2003) and other masses are estimated from the orbital period (using Smith and Dhillon 1998) and $t_{\\rm 3}$ time (using Livio 1992). We rejected systems with unknown or non measurable $t_{\\rm 2}$ or $t_{\\rm 3}$ time (peculiar light curve, strong oscilations in the decline) without catalogue mass values (no information on mass ratio, ex: RS~Car), magnetic novae (AM~Her systems without disc, ex: V1500~Cyg), systems with orbital period $> 10$ h (Smith and Dhillon fitting not adequate, systems with evolved secondary are then excluded too, ex: DI~Lac, V841~Oph, V723~Cas, GK~Per), systems with insuficient photometric observations to detailed period analysis (ex: V500~Aql, V446~Her, HZ~Pup, DY~Pup, CT~Ser). The final list of systems is in Table~\\ref{systems}. \\begin{figure} \\includegraphics[width=90mm]{q_p.eps} \\caption{Selected nova remnants in the mass ratio ($q$) -- orbital period ($P_{\\rm orb}$) plane. Open circle -- non detected superhumps, filled circle -- detected superhumps. Panel a -- mass ratio from combined sourcees (catalogue and estimates from $P_{\\rm orb}$, $t_{\\rm 3}$ time), panel b -- mass ratio only from catalogue values (Ritter and Kolb 2003). The dashed line is the 0.35 limit.} \\label{q_p} \\end{figure}  The results are depicted in Fig.~\\ref{q_p}. The lower panel only shows the systems with known catalogue mass values and the upper panel is after the addition of estimated mass values in this work. Filled circles are systems with detected superhumps. The critical mass ratio of 0.35 is also shown. The results presented in the lower panel are as expected from the mass ratio limit for superhumps, but the statistical set is small. Adding other mass estimates changes the situation. 7 systems below or close to the limit 0.35 show superhumps but 14 systems do not. The estimated mass ratio has a mean difference of 0.2 from the catalogue values. The differences are distributed randomly. Therefore there is no reason to suspect that all mass ratios are systematically underestimated. The 7 systems with detected superhumps are safely located in comparing to the 0.35 critical value as expected by theory and are only a third of all systems falling below or close to the mass ratio of 0.35. The majority of all systems occupy the orbital period interval of 2 -- 4 hours. Following Patterson (2005) this interval has a $\\sim$ 40 -- 90 \\% probability to observe superhumps. 7 systems versus 14 from our investigation yields only $\\sim 30$ \\%.  Superhumps are observed during the outbursts of short orbital period dwarf novae (SU~UMa stars) where the disc reaches the 3:1 resonance radius (superoutburst case). The disc during this active stage (outburst -- active and hot stage, quiescence -- not active and cold stage) reaches larger radius that in quiescence (see Lasota 2001 for review). The irradiation of the secondary by the central white dwarf in nova remnants can produce enhanced mass transfer rate strong enough to keep the disc in this stable hot active stage (as in the nova like stars) with the larger diameter (Hameury and Lasota 2005). In addition, superoutbursts in SU~UMa stars, when superhumps are observed, can be explained by the irradiation and the enhanced mass transfer (Smak 1995, 2004, Hameury et al. 2000, Schreiber et al. 2004). Therefore, permanent superhumpers are possible after a nova outburst. A few possible interpretation of the superhump lack can be: 1) not systematically applicable tidal instability theory; tidal torques are probably not the main and only condition of superhump existence as concluded by Hameury and Lasota (2005), 2) disc radius not developed up to tidal radius, 3) superhumps present but not detected, because of insufficient sets of data."
    },
    "0710/0710.5192_arXiv.txt": {
        "abstract": "The group of 7 thermally emitting and radio-quiet isolated neutron stars (INSs) discovered by ROSAT constitutes a nearby population which locally appears to be as numerous as that of the classical radio pulsars. So far, attempts to enlarge this particular group of INSs finding more remote objects failed to confirm any candidate. We found in the 2XMMp catalogue a handful of sources with no catalogued counterparts and with X-ray spectra similar to those of the ROSAT discovered INSs, but seen at larger distances and thus undergoing higher interstellar absorptions. In order to rule out alternative identifications such as an AGN or a CV, we obtained deep ESO-VLT and SOAR optical imaging for the X-ray brightest candidates. We report here on the current status of our search and discuss the possible nature of our candidates. We focus particularly on the X-ray brightest source of our sample, 2XMM J104608.7-594306, observed serendipitously over more than four years by the XMM-Newton Observatory. A lower limit on the X-ray to optical flux ratio of $\\sim$ 300 together with a stable flux and soft X-ray spectrum make it the most promising thermally emitting INS candidate. Beyond the finding of new members, our study aims at constraining the space density of this population at large distances and at determining whether their apparently high local density is an anomaly or not. ",
        "introduction": "One of the outstanding results of ROSAT is the discovery of seven X-ray bright isolated neutron stars (INSs). These slowly rotating and radio-quiet neutron stars display thermal emission with $kT \\sim$ 40 -- 110 eV, undergo little interstellar absorptions and are not associated with any supernova remnant (see reviews in \\cite{tre00} and \\cite{hab07}). Several have identified faint optical counterparts with $B \\sim$ 25.8 -- 28.6. Proper motion studies (see Motch et al., these proceedings, and references therein) have shown that they are most probably young cooling neutron stars, with ages of a few 10$^5$ years. Their proximity and the apparent absence of strong non-thermal activity turn them into unique laboratories for testing radiative properties of neutron star surfaces in extreme conditions of gravitational and magnetic fields. Moreover, the possibility of measuring their distances through parallaxes \\cite{ker07} or from the distribution of absorption on the line of sight \\cite{pos07} can eventually bring important constraints on the debated equation of state of matter in neutron star interiors.  In the solar neighborhood, ROSAT INSs are as numerous as young radio and $\\gamma$-ray pulsars. It is not clear whether this group is homogeneous. In particular, the absence of radio emission can be either due to the presence of intense magnetic fields, indeed inferred from the measured spin down rates and cyclotron X-ray spectral features, or due to the fact that the radio pencil beam, which narrows at long spin periods, does not sweep over the earth. Considering that ROSAT had not enough sensitivity and spatial resolution to detect the thermal emission of distant sources, the population of cooling radio-quiet INSs could represent a considerable fraction of the total neutron star population of the Galaxy, undetectable in radio surveys \\cite{mot07}. In any case, our knowledge of the overall population characteristics will remain highly unsatisfactory as long as only seven objects are known.  \\begin{table} \\begin{tabular}{l c r r c r r} \\hline \\tablehead{1}{l}{t}{Source}     & \\tablehead{1}{c}{t}{Cand}       & \\tablehead{1}{r}{t}{RA}         & \\tablehead{1}{r}{t}{DEC}        & \\tablehead{1}{c}{t}{$R_{90}$}   & \\tablehead{1}{r}{t}{Count Rate} & \\tablehead{1}{r}{t}{Mag}\\\\ & & (J2000) & (J2000) & arcsec & s$^{-1}$ & $R$\\\\ \\hline 2XMM J104608.7-594306 & 065 & 10 46 08.7 & -59 43 06.1 & 1.33 & 0.060(4)   & $>$ 25\\\\ 2XMM J121017.0-464609 & 164 & 12 10 17.1 & -46 46 11.2 & 2.60 & 0.027(6)   & 20.3  \\\\ 2XMM J010642.3+005032 & 318 & 01 06 42.4 & +00 50 31.3 & 3.80 & 0.020(5)   & 24.5  \\\\ 2XMM J214026.1-233222 & 364 & 21 40 26.2 & -23 32 22.3 & 1.90 & 0.0181(20) & 23.8  \\\\ 2XMM J125904.5-040503 & 604 & 12 59 04.6 & -04 05 02.3 & 1.90 & 0.0129(21) & 21.2  \\\\ 2XMM J125045.7-233349 & 681 & 12 50 45.7 & -23 33 47.7 & 3.30 & 0.0122(21) & 22.1  \\\\ \\hline \\end{tabular} \\caption{INS candidates selected for optical investigation which have already been observed. Count rates are for the EPIC pn camera in the full XMM-Newton energy band (0.2--12.0 keV). We list the $R$ magnitude of the brightest object present in the error circles of the X-ray sources.\\label{tab_cand}} \\end{table}  \\begin{figure} \\includegraphics[height=.28\\textheight]{besteso_0010.eps} \\caption{The positions of our INS candidates, selected from $\\sim$ 7.5$\\cdot$10$^4$ XMM sources, are shown in the HR$_1 \\times$ HR$_2$ diagram (squares). The known ROSAT INSs (crosses) occupy the lowest (less absorbed) part of the diagram. Dashed lines denote soft absorbed blackbodies of different temperatures (50-200 eV) and column absorptions. The nine INS candidates selected for optical follow-up during the present year are highlighted with circles. The X-ray brightest and most promising candidate (source 65) can be noticed by its somewhat smaller error bars. \\label{fig_bestcand}} \\end{figure} ",
        "conclusions": "Our search for new thermally-emitting INSs in the 2XMMp catalogue has revealed a handful of interesting and previously unknown soft X-ray sources among which source 65 is, by far, the most promising candidate. The analysis of its X-ray emission, although based on archival data obtained with non-optimal configurations, reveals an intrinsically soft energy distribution, apparently stable on long time scales. The derived $N_{\\textrm{H}}$ is consistent with that observed towards Eta Carinae and its cluster ($N_{\\textrm{H}} \\sim$ 3$\\cdot$10$^{21}$ cm$^{-2}$). Scaling from RX J0720.4$-$3125 yields a distance of $\\sim$ 3.9 kpc probably compatible with the 2.5 kpc assumed for the Carina Nebula \\cite{2007ApJ...656..462S}. Optical follow-up observations failed to reveal counterparts brighter than $R \\sim$ 25, supporting the idea that source 65 is a new thermally emitting INS.  We have already obtained optical data for some of the other candidates. Preliminary analysis of these data reveals the presence of faint optical candidates within the error circles in all cases. We plan to use optical colour indices and spectra to identify these sources, thus characterizing this sample of soft objects strictly selected from more than 120 thousand entries present in the 2XMMp catalogue.  \\begin{theacknowledgments} This work has been supported by Funda\\c c\\~ao de Amparo \\`a Pesquisa do Estado de S\\~ao Paulo (FAPESP), Coordena\\c c\\~ao de Aperfei\\c coamento de Pessoal de N\\'ivel Superior (CAPES), Brazil and by Universit\\'e Louis Pasteur, Strasbourg, France. \\end{theacknowledgments}"
    },
    "0710/0710.3816.txt": {
        "abstract": "Spectral studies of quiescent emission and bursts of magnetar candidates using XMM-Newton, {\\it Chandra} and {\\it Swift} data are presented. Spectra of both the quiescent emission and the bursts for most magnetar candidates are reproduced by a photoelectrically absorbed two blackbody function (2BB). There is a strong correlation between lower and higher temperatures of 2BB ($kT_{\\mathrm{LT}}$ and $kT_{\\mathrm{HT}}$) for the magnetar candidates of which the spectra are well reproduced by 2BB. In addition, a square of radius for $kT_{\\mathrm{LT}}$ ($R_{\\mathrm{LT}}^2$) is well correlated with a square of radius for $kT_{\\mathrm{HT}}$ ($R_{\\mathrm{HT}}^2$). % A ratio $kT_{\\mathrm LT}/kT_{\\mathrm HT} \\approx 0.4$ is nearly constant irrespective of objects and/or emission types (i.e., the quiescent emission and the bursts). This would imply a common emission mechanism among the magnetar candidates. The relation between the quiescent emission and the bursts might be analogous to a relation between microflares and solar flares of the sun. Three AXPs (4U\\,0142$+$614, 1RXS\\,J170849.0$-$400910 and 1E\\,2259$+$586) seem to have an excess above $\\sim$\\,7\\,keV which well agrees with a non-thermal hard component discovered by INTEGRAL. ",
        "introduction": "Among peculiar celestial objects in the universe, a dense highly magnetized neutron star ($\\rho \\sim 10^{14}$\\,g\\,cm$^{-3}$ and $B \\sim 10^{15}$\\,G), so-called ``magnetar'' \\citep{duncan1992, paczynski1992, thompson1995, thompson1996}, would be one of the most exotic objects. % Soft gamma repeaters (SGRs) and anomalous X-ray pulsars (AXPs) are well known as magnetar candidates. An apparent difference between the SGRs and the AXPs would be considered from their first detections. The SGRs were discovered as sporadically bursting objects, while the AXPs were regarded as peculiar pulsars with long spin periods. However, current observations unveil a lot of similarities between these objects. They have, for instance, long spin periods ($P \\sim $ 5-12\\,s) with spindown rates of $\\dot{P} \\sim 10^{-10}$-$10^{-13}$\\,s\\,s$^{-1}$, no signature of a companion star, a distribution around the galactic plane (two magnetar candidates are in other galaxies), quiescent soft X-ray emission. Several of these objects have non-thermal hard ($>20$\\,keV) components, some are associated with supernova remnants (SNRs), and bursting activity is not confined to the SGRs but is observed in the AXPs as well. Considering these similarities, the SGRs and the AXPs should be classified into a common class of objects.  So far, five SGRs (0501$+$4516, 0526$-$66, 1627$-$41, 1806$-$20 and 1900$+$14) are known \\citep{woods2006, barthelmy2008} as well as three candidates, SGR\\,1801$-$23 \\citep{cline2000}, SGR\\,1808$-$20 \\citep{lamb2003} and SGR/GRB\\,050925. SGR/GRB\\,050925 was regarded as a gamma-ray burst (GRB) when first detected, but soon after was recognized as a new SGR \\citep{holland2005}. On the other hand, ten AXPs (1E\\,2259$+$586, 1E\\,1048.1$-$5937, 4U\\,0142$+$614, 1RXS\\,J170849.0$-$400910, 1E\\,1841$-$045, XTE\\,J1810$-$197, AX\\,J1845$-$0258, CXOU\\,J010043.1$-$721134, CXOU\\,J164710.2$-$455216 and 1E\\,1547.0$-$5408) are known to date \\citep{woods2006, dib2008} with one AXP candidate, AXP\\,CXOU\\,J160103.1$-$513353 \\citep{park2006}. % A short burst from AXP\\,CXOU\\,J164710.2$-$455216 was detected by {\\it Swift} BAT \\citep{krimm2006} at 01:34:52 on 2006 September 21. The follow-up observations performed by {\\it Swift} XRT found a remarkable result in which the quiescent emission of post-burst became 190 times brighter than that of pre-burst \\citep{campana2006}. In addition to these objects, AX\\,J1818.8$-$1559 discovered by ASCA \\citep{sugizaki2001} recently exhibited a short burst \\citep{mereghetti2007b} similar to those from the magnetar candidates. Therefore AX\\,J1818.8$-$1559 could be a new SGR or AXP \\citep{mereghetti2007b}.  The most exciting phenomena among the magnetar candidates would be a sudden release of huge energy in rather short period, the so-called {\\it giant flares} from the SGRs. They typically have a short intense spike which last less than 1\\,s, and followed by a long pulsating tail which lasts a few hundred seconds. Their peak energy flux can be larger than $\\sim10^{6}$ times Eddington luminosity. Theoretical studies suggested that the giant flares were triggered by a catastrophic deformation of the neutron star crust due to a torsion of the strong magnetic field (e.g., \\cite{thompson2001}). Some different emission mechanisms have been proposed by several authors \\citep{yamazaki2005, lyutikov2006, cea2006}. % In the past three decades, three giant flares were recorded. The first detection, from the source now known as SGR\\,0526$-$66 in Large Magellanic Cloud (LMC), was made on March 5 in 1979 \\citep{mazets1979, cline1980, evans1980, fenimore1996}. % The second one from SGR\\,1900$+$14 was recorded on August 27 in 1998 \\citep{hurley1999b, feroci1999, mazets1999, feroci2001, tanaka2007}. More recently, the most energetic giant flare from SGR\\,1806$-$20 was observed on December 27 in 2004 \\citep{cameron2005, gaensler2005, hurley2005, mazets2005, palmer2005, terasawa2005, tanaka2007}. The fluence of its initial intense spike with 600\\,ms was evaluated to be $\\sim 2$\\,erg\\,cm$^{-2}$ by the plasma particle detectors on the Geotail space probe \\citep{terasawa2005}.  Soft X-ray spectra of the quiescent emission of the SGRs and the AXPs were observed by a number of satellites. Although their spectral model is still under discussion, two two-component models are proposed. % One of them is a photoelectrically absorbed two blackbody function (2BB). Spectral parameters of 2BB are reported by some authors for the SGRs \\citep{mereghetti2006a} and the AXPs \\citep{tiengo2002, morii2003, gotthelf2004, gotthelf2005, halpern2005, tiengo2005, israel2006, gotthelf2007}. % Typical lower and higher temperatures are $\\sim 0.5$\\,keV and $\\sim 1.4$\\,keV, respectively. % The other model is a photoelectrically absorbed power law plus a blackbody (PL$+$BB). Some authors report spectral parameters of PL$+$BB for the SGRs \\citep{marsden2001, kurkarni2003, mereghetti2005, mereghetti2006a, mereghetti2007a} and the AXPs \\citep{morii2003, patel2003, rea2003, gotthelf2004, mereghetti2004, woods2004, tiengo2005, gavriil2006, israel2006}. % A typical power law index and a blackbody temperature are $\\sim 3$ and $\\sim 0.5$\\,keV. % At present it is still unclear which model is more reliable or physically suitable.  Recent observations by the International Gamma-Ray Astrophysics Laboratory (INTEGRAL) discovered a non-thermal hard component in the spectra of the quiescent emission above 20\\,keV for 5 magnetar candidates \\citep{molkov2005, gotz2006b, kuiper2006}. The non-thermal hard component is well reproduced by a power law model, $E^{-\\Gamma}$, where $\\Gamma$ ranges from 1.0-1.8, while the soft X-ray emission below $\\sim12$\\,keV, mentioned above, clearly indicates steeper power-law index of $\\sim 3$ if the PL$+$BB is applied as the model spectrum. Hence, the non-thermal hard emission seen by INTEGRAL is a different component and presumably has a different origin than the soft X-ray emission. Since some magnetar candidates have two different emission mechanisms, there seems to be more complex physics than expected before. % Moreover, the non-thermal hard component shows pulsations for three AXPs, 1RXS\\,J170849.0$-$400910, 4U\\,0142$+$614 and 1E\\,1841$-$045, through the INTEGRAL and RXTE observations \\citep{kuiper2006}, which is related to a neutron star rotation, and hence there are particle acceleration processes in the vicinity of neutron stars \\citep{kuiper2006}.  If the energy source of the quiescent emission and the % short bursts is the magnetic field as thought to be, at least very similar physical process would govern both of them and their spectra could emerge alike. It is claimed based on High Energy Transient Explorer 2 (HETE-2) data that the most acceptable spectral model of the short bursts from two SGRs 1806\u00a1\u00dd20 and 1900+14 is 2BB even though it should be regarded just as an empirical model  (Nakagawa et al. 2007). It would be also preferred to represent spectra of quiescent emissions by 2BB rather than BB+PL for SGRs, and even for AXPs if it is  the same class of object. In this paper, we present a comprehensive spectral study with 2BB for both the quiescent emission and the for the magnetar candidates.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "The spectral studies using the photoelectrically absorbed two blackbody function (2BB) were presented for the quiescent emission and the burst of the magnetar candidates. % The spectra of the quiescent emission were well reproduced by a 2BB with some exceptions. The spectra of three AXPs (4U\\,0142$+$614, 1RXS\\,J170849.0$-$400910 and 1E\\,2259$+$586) seem to have an excess which might be due to a non-thermal hard component discovered by INTEGRAL. % The spectrum of the burst from the SGR candidate SGR\\,2013$+$34 was also well reproduced by 2BB.  A strong linear correlations between $kT_{\\mathrm{LT}}$ and $kT_{\\mathrm{HT}}$ was found using 2BB spectra. The ratio $kT_{\\mathrm{LT}}/kT_{\\mathrm{HT}} \\sim 0.4$ is almost constant irrespectively of the objects and/or emission types (burst or quiescent emission). The relationship between $R_{\\mathrm{LT}}^2$ and $R_{\\mathrm{HT}}^2$ seems to have a linear correlation. Considering these correlations, there seems to be a common emission mechanism among these objects, and between the quiescent emission and the burst. The relationship between the quiescent emission and the burst might be similar to the relationship between microflares and an ordinary solar flares of the sun. The quiescent emission might be due to very frequent small activity similar to the microflares. On the other hand, the burst might be due to a relatively large activity similar to the ordinary solar flare.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % Acknowledgement \\bigskip We would like to thank an anonymous referee for helpful comments and suggestions to improve our paper. This work is based on observations obtained with XMM-Newton, an ESA science mission with instruments and contributions directly funded by ESA Member States and NASA. % We would like to thank public data archive of {\\it Chandra}. % This research has made use of software provided by the Chandra X-ray Center (CXC) in the application packages CIAO, ChIPS, and Sherpa. % We acknowledge the use of public data from the Swift data archive. % YEN is supported by the JSPS Research Fellowships for Young Scientists. This work is supported in part by a special postdoctoral researchers program in RIKEN.   %\\appendix  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % Figures %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %\\clearpage % \\begin{figure} \\begin{center} \\FigureFile(80mm,129mm){figure1.eps} \\end{center} \\caption{Time history of the (a) flux in 2-10 keV in units of ergs cm$^{-2}$ s$^{-1}$, (b) photoelectric absorption $N_{\\mathrm{H}}$ in units of cm$^{-2}$, (c) temperature of the lower blackbody $kT_{\\mathrm{LT}}$ in units of keV, (d) radius of the lower blackbody $R_{\\mathrm{LT}}$ in units of km, (e) temperature of the higher blackbody $kT_{\\mathrm{HT}}$ in units of keV and (f) radius of the higher blackbody $R_{\\mathrm{HT}}$ in units of km for the emissions of AXP\\,CXOU\\,J164710.2$-$455216 observed by XRT/{\\it Swift}.}\\label{fig:lc_persistent_cxou_j16} \\end{figure}  \\begin{figure} \\begin{center} \\FigureFile(80mm,80mm){figure2.eps} \\end{center} \\caption{A schematic view of $\\nu{F}_{\\nu}$ spectra using 2BB$+$PL (a) and BB$+$2PL (b) for AXP\\,4U\\,0142$+$614. Spectral parameters of X-ray spectra (i.e., $\\lesssim$12\\,keV) are derived from our analyses using the XMM-Newton observation of 0112781101, while those of the non-thermal hard component (i.e., $\\gtrsim$20\\,keV) is obtained by INTEGRAL observations \\citep{kuiper2006}. The circles denote data derived from our analyses using the XMM-Newton observation of 0112781101. The squares represent INTEGRAL observations taken from Fig.7 in \\citet{kuiper2006} by eye. The dashed, dot-dash, dotted lines in (a) show the $kT_{\\mathrm{LT}}$ component, the $kT_{\\mathrm{HT}}$ component and the PL component for the hard spectrum, respectively. Those lines in (b) show the $kT$ component, the PL component for an X-ray spectrum and the PL component for the hard spectrum, respectively.}\\label{fig:eeufspec_summary} \\end{figure}  \\begin{figure} \\begin{center} \\FigureFile(80mm,83mm){figure3.eps} \\end{center} \\caption{Relationship between the 2BB temperatures $kT_{\\mathrm{LT}}$ and $kT_{\\mathrm{HT}}$. The triangles and squares denote the previous work on the bursts \\citep{feroci2004, olive2004, gotz2006a, nakagawa2007} and the quiescent emission \\citep{morii2003, gotthelf2004, gotthelf2005, tiengo2005, mereghetti2006a}, respectively. The circles and stars denote our work on the bursts and the quiescent emission, respectively. The line represents the best-fit power law model.}\\label{fig:bb_relation}  \\end{figure}  \\begin{figure} \\begin{center} \\FigureFile(80mm,80mm){figure4.eps} \\end{center} \\caption{Relationship between the square of the blackbody radii $R^2_{\\mathrm{LT}}$ and $R^2_{\\mathrm{HT}}$. The triangles and squares denote the previous work on the the bursts \\citep{olive2004, nakagawa2007} and the quiescent emission \\citep{morii2003, tiengo2005, mereghetti2006a}, respectively. The stars denote our work on the quiescent emission. The solid line shows a ratio of $R_{\\mathrm{HT}}^{2}$ to $R_{\\mathrm{LT}}^{2}$ of 0.01.}\\label{fig:radius_relation} \\end{figure}  \\clearpage  \\begin{figure} \\begin{center} \\FigureFile(160mm,80mm){figure5.eps} \\end{center} \\caption{{\\it Left:} Relationship between the lower temperature of 2BB $kT_{\\mathrm{LT}}$ and the square of the blackbody radii of 2BB $R_{\\mathrm{LT}}^2$. {\\it Right:} Relationship between the higher temperature of 2BB $kT_{\\mathrm{HT}}$ and the square of the blackbody radii $R_{\\mathrm{HT}}^2$. The dotted and dashed lines correspond to bolometric fluences of $10^{-8}$ and $10^{-9}$\\,ergs\\,cm$^{-2}$, respectively.}\\label{fig:kt_rsquared_relation} \\end{figure}   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % Tables %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\clearpage  \\begin{table*} \\small \\caption{A summary of magnetar candidates which were employed in our study.}\\label{tab:magnetar_list} \\begin{center} \\begin{tabular}{lllll} \\hline\\hline Object\\footnotemark[$*$] & Satellite/Instrument\\footnotemark[$\\dagger$] & Period\\footnotemark[$\\ddagger$] & Distance\\footnotemark[$\\S$] & Ref.\\footnotemark[$\\|$] \\\\ \\hline SGR\\,0526$-$66 & {\\it Chandra} ACIS & 2000-2001 & 50 & (1) \\\\ SGR\\,1627$-$41 & XMM-Newton EPIC, {\\it Chandra} ACIS & 2001-2005 & 11 & (2), (3), (4) \\\\ SGR\\,1806$-$20 & XMM-Newton EPIC, {\\it Chandra} ACIS & 2000-2005 & 15\\footnotemark[$\\#$] & (5), (6) \\\\ SGR\\,2013$+$34 & {\\it Swift} BAT & 2005 & 10 & (7), (8) \\\\ SGR\\,1819$-$16 & XMM-Newton EPIC & 2003 & 10 & \\\\ AXP\\,CXOU\\,J010043.1$-$721134 & XMM-Newton EPIC, {\\it Chandra} ACIS & 2000-2005 & 57 & (9), (10) \\\\ AXP\\,4U\\,0142$+$614 & XMM-Newton EPIC & 2002-2004 & 3 & (11) \\\\ AXP\\,CXOU\\,J164710.2$-$455216 & {\\it Swift} BAT, {\\it Swift} XRT, {\\it Chandra} ACIS & 2005-2007 & 5 & (12), (13), (14), (15) \\\\ AXP\\,1RXS\\,J170849.0$-$400910 & XMM-Newton EPIC & 2003 & 5 & (16), (17) \\\\ AXP\\,1E\\,1841$-$045 & XMM-Newton EPIC & 2002 & 7\\footnotemark[$**$] & \\\\ AXP\\,1E\\,2259$+$586 & XMM-Newton EPIC & 2002-2005 & 3 & (18) \\\\ \\hline \\multicolumn{5}{@{}l@{}}{\\hbox to 0pt{\\parbox{180mm}{\\footnotesize \\footnotemark[$*$] Object name of magnetar candidates (SGR\\,2013$+$34 denotes SGR candidate SGR/GRB\\,050925). \\par\\noindent \\footnotemark[$\\dagger$] Instrument and satellite names from which obtained the data used in our analysis. \\par\\noindent \\footnotemark[$\\ddagger$] Interval which these observations were performed. \\par\\noindent \\footnotemark[$\\S$] Distances to each object in units of kpc (see \\cite{woods2006} and references there in). \\par\\noindent \\footnotemark[$\\|$] (1) \\citet{kurkarni2003}; (2) \\citet{kouveliotou2003}; (3) \\citet{wachter2004}; (4) \\citet{mereghetti2006b}; (5) \\citet{kaplan2002}; (6) \\citet{mereghetti2005}; (7) \\citet{holland2005}; (8) \\citet{markwardt2005}; (9) \\citet{lamb2002}; (10) \\citet{majid2004}; (11) \\citet{guver2007}; (12) \\citet{krimm2006}; (13) \\citet{campana2006}; (14) \\citet{muno2006}; (15) \\citet{muno2007}; (16) \\citet{oosterbroek2004}; (17) \\citet{rea2007}; (18) \\citet{woods2004} \\par\\noindent \\footnotemark[$\\#$] The latest distance estimate is $d = 8.7$\\,kpc \\citep{bibby2008}. \\par\\noindent \\footnotemark[$**$] The latest distance estimate is $d = 8.5$\\,kpc \\citep{tian2008}. }\\hss}} \\end{tabular} \\end{center} \\end{table*}  \\begin{table*} \\scriptsize \\caption{XMM-Newton observations of the quiescent emissions of the SGRs.}\\label{tab:sgr_obs_xmm} \\begin{center} \\begin{tabular}{lllllllllllll} \\hline\\hline Object\\footnotemark[$*$] & ObsID\\footnotemark[$\\dagger$] & \\multicolumn{2}{l}{Observation Date (MJD)\\footnotemark[$\\ddagger$]} & \\multicolumn{3}{l}{Observation Mode\\footnotemark[$\\S$]} & \\multicolumn{3}{l}{Exposure Time (ks)\\footnotemark[$\\|$]} & \\multicolumn{3}{l}{Source/Background Radii} \\\\ &                         & Start & End & pn & MOS1 & MOS2 & pn & MOS1 & MOS2 & pn & MOS1 & MOS2 \\\\ \\hline 1627$-$41 & 0204500201 & 53051.590 & 53051.992 & Full & Full & Full & 15.89 & 20.03 & 20.17 & $\\timeform{10''}$/$\\timeform{10''}$ & $\\timeform{10''}$/$\\timeform{10''}$ & $\\timeform{10''}$/$\\timeform{10''}$ \\\\ 1627$-$41 & 0204500301 & 53252.750 & 53253.131 & Full & Full & Full & 27.06 & 32.00 & 32.10 & $\\timeform{10''}$/$\\timeform{10''}$ & $\\timeform{10''}$/$\\timeform{10''}$ & $\\timeform{10''}$/$\\timeform{10''}$ \\\\ 1627$-$41 & 0202560101 & 53270.677 & 53271.281 & Small & P-W2 & P-W2 & 26.66 & 36.19 & 49.38 & $\\timeform{10''}$/$\\timeform{10''}$ & $\\timeform{10''}$/$\\timeform{10''}$ & $\\timeform{10''}$/$\\timeform{10''}$ \\\\ 1806$-$20 & 0148210101 & 52732.566 & 52733.209 & Full & P-W3 & P-W3 & 4.84 & 5.65 & 5.62 & $\\timeform{32''}$/$\\timeform{32''}$ & $\\timeform{28''}$/$\\timeform{28''}$ & $\\timeform{28''}$/$\\timeform{28''}$ \\\\ 1806$-$20 & 0148210401 & 52919.404 & 52919.663 & Full & P-W3 & P-W3 & 7.61 & 7.07 & 7.23 & $\\timeform{32''}$/$\\timeform{32''}$ & $\\timeform{28''}$/$\\timeform{28''}$ & $\\timeform{28''}$/$\\timeform{28''}$ \\\\ 1806$-$20 & 0205350101 & 53254.377 & 53254.978 & Small & P-W3 & P-W3 & 30.21 & 39.14 & 39.56 & $\\timeform{32''}$/$\\timeform{32''}$ & $\\timeform{28''}$/$\\timeform{28''}$ & $\\timeform{28''}$/$\\timeform{28''}$ \\\\ 1806$-$20 & 0164561101 & 53284.706 & 53284.925 & Small & Fast-U & Fast-U & 11.54 & $t$ & $t$ & $\\timeform{32''}$/$\\timeform{32''}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ \\\\ 1806$-$20 & 0164561301 & 53436.348 & 53436.636 & Small & Fast-U & Full & 7.37 & $t$ & 5.68 & $\\timeform{32''}$/$\\timeform{32''}$ & $\\cdot\\cdot\\cdot$ & $\\timeform{28''}$/$\\timeform{28''}$ \\\\ 1806$-$20 & 0164561401 & 53647.427 & 53647.809 & Small & Fast-U & Full & 22.11 & $t$ & 28.68 & $\\timeform{32''}$/$\\timeform{32''}$ & $\\cdot\\cdot\\cdot$ & $\\timeform{28''}$/$\\timeform{28''}$ \\\\ 2013$+$34 & 0212481201 & 53655.026 & 53655.334 & Full & Full & Full & 22.18 & 25.27 & 25.27 & $\\timeform{32''}$/$\\timeform{32''}$ & $\\timeform{28''}$/$\\timeform{28''}$ & $\\timeform{28''}$/$\\timeform{28''}$ \\\\ 1819$-$16 & 0152834501 & 52726.191 & 52726.310 & Full & Full & Full & 3.3 & 4.3 & 4.8 & $\\timeform{32''}$/$\\timeform{80''}$\\footnotemark[$\\#$] & $\\timeform{28''}$/$\\timeform{70''}$\\footnotemark[$\\#$] & $\\timeform{28''}$/$\\timeform{70''}$\\footnotemark[$\\#$] \\\\ \\hline \\multicolumn{13}{@{}l@{}}{\\hbox to 0pt{\\parbox{180mm}{\\footnotesize \\footnotemark[$*$] Object name of the SGRs (2013$+$34 denotes SGR candidate SGR/GRB\\,050925 and 1819$-$16 denotes SGR candidate AX\\,J1818.8$-$1559). \\par\\noindent \\footnotemark[$\\dagger$] XMM-Newton observation ID. \\par\\noindent \\footnotemark[$\\ddagger$] Start and end time of observations. \\par\\noindent \\footnotemark[$\\S$] Observation mode for each instrument; full-window mode (Full), small-window mode (Small), partial-w2 mode (P-W2), partial-w3 mode (P-W3) and fast-uncompressed mode (Fast-U). \\par\\noindent \\footnotemark[$\\|$] Net exposure time for each instrument. $t$ denotes the data sets obtained by the MOS cameras in timing mode and not utilized. \\par\\noindent \\footnotemark[$\\#$] The background regions were extracted from an annular region whose center was the source position. The first values are source radii, and the inner radii of the background regions. The second values are outer radii of the background regions. }\\hss}} \\end{tabular} \\end{center} \\end{table*}  \\begin{table*} \\small \\caption{{\\it Chandra} observations of the quiescent emissions of the SGRs.}\\label{tab:sgr_obs_cxo} \\begin{center} \\begin{tabular}{lllllll} \\hline\\hline Object\\footnotemark[$*$] & ObsID\\footnotemark[$\\dagger$] & \\multicolumn{2}{l}{Observation Date (MJD)\\footnotemark[$\\ddagger$]} & Observation Mode\\footnotemark[$\\S$] & Exposure Time (ks)\\footnotemark[$\\|$] & Source/Background Radii \\\\ &  & Start & End                                &                  &             &       \\\\ \\hline 0526$-$66 & 747 & 51547.017 & 51547.539 & FAINT & 39.86 & $\\timeform{1''}$/$\\timeform{1''}$ \\\\ 0526$-$66 & 1957 & 52152.937 & 52153.566 & FAINT & 48.45 & $\\timeform{1''}$/$\\timeform{1''}$ \\\\ 1627$-$41 & 1981 & 52182.205 & 52182.803 & FAINT  & 48.93 & $\\timeform{2''}$/$\\timeform{2''}$ \\\\ 1627$-$41 & 3877 & 52722.169 & 52722.494 & VFAINT & 25.67 & $\\timeform{2''}$/$\\timeform{2''}$ \\\\ \\hline \\multicolumn{7}{@{}l@{}}{\\hbox to 0pt{\\parbox{180mm}{\\footnotesize \\footnotemark[$*$] SGR names. \\par\\noindent \\footnotemark[$\\dagger$] {\\it Chandra} observation ID. \\par\\noindent \\footnotemark[$\\ddagger$] Start and end time of the observations. \\par\\noindent \\footnotemark[$\\S$] FAINT and VFAINT denote the imaging mode, and CC33\\_FAINT denotes the timing mode. \\par\\noindent \\footnotemark[$\\|$] Net exposure time. }\\hss}} \\end{tabular} \\end{center} \\end{table*}  \\begin{table*} \\scriptsize \\caption{XMM-Newton observations of the quiescent emissions of the AXPs.}\\label{tab:axp_obs_xmm} \\begin{center} \\begin{tabular}{lllllllllllll} \\hline\\hline Object\\footnotemark[$*$] & ObsID\\footnotemark[$\\dagger$] & \\multicolumn{2}{l}{Observation Date (MJD)\\footnotemark[$\\ddagger$]} & \\multicolumn{3}{l}{Observation Mode\\footnotemark[$\\S$]} & \\multicolumn{3}{l}{Exposure Time (ks)\\footnotemark[$\\|$]} & \\multicolumn{3}{l}{Source/Background Radii} \\\\ &                         & Start & End & pn & MOS1 & MOS2 & pn & MOS1 & MOS2 & pn & MOS1 & MOS2 \\\\ \\hline 0100$-$721 & 0110000201 & 51834.626 & 51834.867 & E-Full & Full & Full & 20.81 & 14.62 & $g$ & $\\timeform{32''}$/$\\timeform{32''}$ & $\\timeform{28''}$/$\\timeform{28''}$ & $\\timeform{28''}$/$\\timeform{28''}$ \\\\ 0100$-$721 & 0018540101 & 52233.983 & 52234.303 & Full & Full & Full & 21.16 & 25.73 & $g$ & $\\timeform{32''}$/$\\timeform{32''}$ & $\\timeform{28''}$/$\\timeform{28''}$ & $\\timeform{28''}$/$\\timeform{28''}$ \\\\ 0142$+$614 & 0112781101 & 52663.920 & 52663.995 & Small & Fast-U & Fast-U & 4.18 & $t$ & $t$ & $\\timeform{32''}$/$\\timeform{32''}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ \\\\ 1708$-$400 & 0148690101 & 52879.906 & 52880.426 & Full & P-W3 & P-W3 & 26.88 & $p$ & $p$ & $\\timeform{20''}$/$\\timeform{20''}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ \\\\ 1841$-$045 & 0013340101 & 52552.122 & 52552.192 & Large & Full & Full & 2.34 & 3.54 & 3.57 & $\\timeform{12''}$/$\\timeform{12''}$ & $\\timeform{12''}$/$\\timeform{12''}$ & $\\timeform{12''}$/$\\timeform{12''}$ \\\\ 1841$-$045 & 0013340201 & 52554.115 & 52554.193 & Large & Full & Full & 4.37 & 6.30 & 6.30 & $\\timeform{12''}$/$\\timeform{12''}$ & $\\timeform{12''}$/$\\timeform{12''}$ & $\\timeform{12''}$/$\\timeform{12''}$ \\\\ 2259$+$586 & 0038140101 & 52436.378 & 52436.986 & Small & Full & Full & 30.63 & $p$ & $p$ & $\\timeform{32''}$/$\\timeform{32''}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ \\\\ 2259$+$586 & 0155350301 & 52446.400 & 52446.759 & Small & P-W2 & Full & 17.65 & $p$ & $p$ & $\\timeform{32''}$/$\\timeform{32''}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ \\\\ 2259$+$586 & 0203550701 & 53579.965 & 53580.030 & Small & P-W2 & Fast-U & 2.66 & $p$ & $t$ & $\\timeform{32''}$/$\\timeform{32''}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ \\\\ \\hline \\multicolumn{13}{@{}l@{}}{\\hbox to 0pt{\\parbox{180mm}{\\footnotesize \\footnotemark[$*$] Object name of the AXPs; CXOU\\,J010043.1$-$721134 (0100$-$721), 4U\\,0142$+$614 (0142$+$614), 1RXS\\,J170849.0$-$400910 (1708$-$400), 1E\\,1841$-$045 and 1E\\,2259$+$586. \\par\\noindent \\footnotemark[$\\dagger$] XMM-Newton observation ID. \\par\\noindent \\footnotemark[$\\ddagger$] Start and end time of the observations. \\par\\noindent \\footnotemark[$\\S$] Observation mode for each instrument; extended full-window mode (E-Full), Full-window mode (Full), small-window mode (Small), fast-uncompressed mode (Fast-U), fast-timing mode (Fast-T), partial-w3 mode (P-W3), large-window mode (Large) and partial-w2 mode (P-W2). \\par\\noindent \\footnotemark[$\\|$] Net exposure time for each instrument. $g$ denotes that the source fell on a gap of the CCD chips, $t$ denotes observations in timing mode, and $p$ denotes that the data sets are affected by a photon pile-up. These data sets were not utilized. }\\hss}} \\end{tabular} \\end{center} \\end{table*}  \\begin{table*} \\small \\caption{{\\it Chandra} observations of the quiescent emissions of the AXPs.}\\label{tab:axp_obs_cxo} \\begin{center} \\begin{tabular}{lllllll} \\hline\\hline Object\\footnotemark[$*$] & ObsID\\footnotemark[$\\dagger$] & \\multicolumn{2}{l}{Observation Date (MJD)\\footnotemark[$\\ddagger$]} & Observation Mode\\footnotemark[$\\S$] & Exposure Time (ks)\\footnotemark[$\\|$] & Source/Background Radii \\\\ &  & Start & End                                &                  &          &          \\\\ \\hline 0100$-$721 & 1881 & 52044.080 & 52045.261 & FAINT  & 98.67 & $\\timeform{11''}$/$\\timeform{10''}$\\footnotemark[$\\#$] \\\\ 0100$-$721 & 4616 & 53031.791 & 53032.009 & VFAINT & 15.56 & $\\timeform{2''}$/$\\timeform{2''}$ \\\\ 0100$-$721 & 4617 & 53032.189 & 53032.399 & VFAINT & 15.27 & $\\timeform{2''}$/$\\timeform{2''}$ \\\\ 0100$-$721 & 4618 & 53033.904 & 53034.130 & VFAINT & 15.00 & $\\timeform{2''}$/$\\timeform{2''}$ \\\\ 0100$-$721 & 4619 & 53042.806 & 53043.023 & VFAINT & 15.04 & $\\timeform{2''}$/$\\timeform{2''}$ \\\\ 0100$-$721 & 4620 & 53089.185 & 53089.395 & VFAINT & 15.22 & $\\timeform{2''}$/$\\timeform{2''}$ \\\\ 1647$-$455 & 6283 & 53512.860 & 53513.102 & FAINT  & 18.81 & $\\timeform{2''}$/$\\timeform{2''}$ \\\\ 1647$-$455 & 5411 & 53539.673 & 53540.141 & FAINT  & 38.47 & $\\timeform{2''}$/$\\timeform{2''}$ \\\\ \\hline \\multicolumn{7}{@{}l@{}}{\\hbox to 0pt{\\parbox{180mm}{\\footnotesize \\footnotemark[$*$] Object name of the AXPs; CXOU\\,J010043.1$-$721134 (0100$-$721) and CXOU\\,J164710.2$-$455216 (1647$-$455). \\par\\noindent \\footnotemark[$\\dagger$] {\\it Chandra} observation ID. \\par\\noindent \\footnotemark[$\\ddagger$] Start and end time of observations. \\par\\noindent \\footnotemark[$\\S$] FAINT and VFAINT denote the imaging mode. \\par\\noindent \\footnotemark[$\\|$] Net exposure time. \\par\\noindent \\footnotemark[$\\#$] Since the source fell on an off-axis CCD chip, the source region was extracted from an elliptical region with major and minor axes of $\\timeform{11''}$ and $\\timeform{10''}$, respectively. }\\hss}} \\end{tabular} \\end{center} \\end{table*}  \\begin{table*} \\small \\caption{{\\it Swift} observations of the quiescent emission and the bursts of AXP\\,CXOU\\,J164710.2$-$455216.}\\label{tab:axp_j1647_obs_swift} \\begin{center} \\begin{tabular}{lllllllll} \\hline\\hline SeqNum\\footnotemark[$*$] & \\multicolumn{4}{l}{Observation Date (MJD)\\footnotemark[$\\dagger$]} & \\multicolumn{2}{l}{Exposure Time (ks)\\footnotemark[$\\ddagger$]} & \\multicolumn{2}{l}{Source/Background Radii} \\\\ & Start (WT) & End (WT) & Start (PC) & End (PC) & WT & PC & WT & PC \\\\ \\hline 00230341000\\footnotemark[$\\S$] & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ \\\\ 00030806001 & 53999.604 & 53999.924 & 53999.604 & 53999.936 & 1.92 & 7.74 & $\\timeform{36''}\\times\\timeform{18''}$/$\\timeform{36''}\\times\\timeform{18''}$\\footnotemark[$\\|$] & $\\timeform{30''}$/$\\timeform{30''}$ \\\\ 00030806002 & 54000.610 & 54000.619 & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 0.77 & $\\cdot\\cdot\\cdot$ & $\\timeform{36''}\\times\\timeform{18''}$/$\\timeform{36''}\\times\\timeform{18''}$\\footnotemark[$\\|$] & $\\cdot\\cdot\\cdot$ \\\\ 00030806003 & 54000.819 & 54001.212 & 54000.819 & 54001.073 & 4.91 & 1.84 & $\\timeform{36''}\\times\\timeform{18''}$/$\\timeform{36''}\\times\\timeform{18''}$\\footnotemark[$\\|$] & $\\timeform{30''}$/$\\timeform{30''}$ \\\\ 00030806004 & 54004.276 & 54004.483 & 54004.343 & 54004.489 & 1.25 & 2.48 & $\\timeform{36''}\\times\\timeform{18''}$/$\\timeform{36''}\\times\\timeform{18''}$\\footnotemark[$\\|$] & $\\timeform{30''}$/$\\timeform{30''}$ \\\\ 00030806006 & 54010.461 & 54010.716 & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 1.98 & $\\cdot\\cdot\\cdot$ & $\\timeform{36''}\\times\\timeform{18''}$/$\\timeform{36''}\\times\\timeform{18''}$\\footnotemark[$\\|$] & $\\cdot\\cdot\\cdot$ \\\\ 00030806007 & 54011.515 & 54011.594 & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 2.03 & $\\cdot\\cdot\\cdot$ & $\\timeform{36''}\\times\\timeform{18''}$/$\\timeform{36''}\\times\\timeform{18''}$\\footnotemark[$\\|$] & $\\cdot\\cdot\\cdot$ \\\\ 00030806008 & 54014.001 & 54014.077 & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 2.16 & $\\cdot\\cdot\\cdot$ & $\\timeform{36''}\\times\\timeform{18''}$/$\\timeform{36''}\\times\\timeform{18''}$\\footnotemark[$\\|$] & $\\cdot\\cdot\\cdot$ \\\\ 00030806009 & 54017.746 & 54017.954 & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 3.52 & $\\cdot\\cdot\\cdot$ & $\\timeform{36''}\\times\\timeform{18''}$/$\\timeform{36''}\\times\\timeform{18''}$\\footnotemark[$\\|$] & $\\cdot\\cdot\\cdot$ \\\\ 00030806010 & 54018.009 & 54018.094 & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 2.83 & $\\cdot\\cdot\\cdot$ & $\\timeform{36''}\\times\\timeform{18''}$/$\\timeform{36''}\\times\\timeform{18''}$\\footnotemark[$\\|$] & $\\cdot\\cdot\\cdot$ \\\\ 00030806011 & 54023.237 & 54023.517 & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 5.62 & $\\cdot\\cdot\\cdot$ & $\\timeform{36''}\\times\\timeform{18''}$/$\\timeform{36''}\\times\\timeform{18''}$\\footnotemark[$\\|$] & $\\cdot\\cdot\\cdot$ \\\\ 00030806012 & 54029.125 & 54029.342 & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 5.52 & $\\cdot\\cdot\\cdot$ & $\\timeform{36''}\\times\\timeform{18''}$/$\\timeform{36''}\\times\\timeform{18''}$\\footnotemark[$\\|$] & $\\cdot\\cdot\\cdot$ \\\\ 00030806013 & 54035.677 & 54035.825 & 54035.678 & 54035.774 & 2.82 & 0.42 & $\\timeform{36''}\\times\\timeform{18''}$/$\\timeform{36''}\\times\\timeform{18''}$\\footnotemark[$\\|$] & $\\timeform{30''}$/$\\timeform{30''}$ \\\\ 00030806014 & 54119.174 & 54119.253 & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 2.06 & $\\cdot\\cdot\\cdot$ & $\\timeform{36''}\\times\\timeform{18''}$/$\\timeform{36''}\\times\\timeform{18''}$\\footnotemark[$\\|$] & $\\cdot\\cdot\\cdot$ \\\\ 00030806015 & 54122.050 & 54122.267 & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 3.82 & $\\cdot\\cdot\\cdot$ & $\\timeform{36''}\\times\\timeform{18''}$/$\\timeform{36''}\\times\\timeform{18''}$\\footnotemark[$\\|$] & $\\cdot\\cdot\\cdot$ \\\\ \\hline \\multicolumn{9}{@{}l@{}}{\\hbox to 0pt{\\parbox{180mm}{\\footnotesize \\footnotemark[$*$] {\\it Swift} sequence number. \\par\\noindent \\footnotemark[$\\dagger$] Start and end time of the observations for each mode (WT denotes window timing mode, and PC denotes photon counting mode). \\par\\noindent \\footnotemark[$\\ddagger$] Net exposure time for each observation mode. \\par\\noindent \\footnotemark[$\\S$] The observation of a burst. \\par\\noindent \\footnotemark[$\\|$] The source and background regions were extracted from a rectangle region. Two background regions are utilized near both sides of the source region. }\\hss}} \\end{tabular} \\end{center} \\end{table*}  \\begin{table*} \\caption{Spectral parameters of the quiescent emissions of the SGRs observed by XMM-Newton.}\\label{tab:sgr_spc_xmm} \\begin{center} \\begin{tabular}{lllllllll} \\hline\\hline Object\\footnotemark[$*$] & ObsID\\footnotemark[$\\dagger$] & $N_{\\mathrm{H}}$\\footnotemark[$\\ddagger$] & $kT_{\\mathrm{LT}}$\\footnotemark[$\\S$] & $R_{\\mathrm{LT}}$\\footnotemark[$\\|$] &  $kT_{\\mathrm{HT}}$\\footnotemark[$\\S$] & $R_{\\mathrm{HT}}$\\footnotemark[$\\|$] & $F$\\footnotemark[$\\#$] & $\\chi^{2}$ (d.o.f.) \\\\ &       & (10$^{22}$ cm$^{-2}$) & (keV)         & (km)         & (keV)         & (km)         &   & \\\\ \\hline 1627$-$41 & 0204500201 & 15.98$_{-7.60}^{+15.51}$ & 0.58$_{-0.25}^{+0.35}$ & $<$19.25 & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & $\\sim$\\,0.05 & 44 (42) \\\\ 1627$-$41 & 0204500301 & 7.53$_{-3.56}^{+6.19}$ & 0.85$_{-0.22}^{+0.29}$ & $<$0.62 & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 0.08$\\pm0.06$ & 29 (65) \\\\ 1627$-$41 & 0202560101 & 9.00$_{-3.89}^{+6.71}$ & 0.94$_{-0.24}^{+0.31}$ & $<$0.42 & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 0.06$\\pm0.04$ & 54 (47) \\\\ 1806$-$20 & 0148210101 & 5.18$_{-0.73}^{+0.92}$ & 0.84$_{-0.17}^{+0.23}$ & 1.64$_{-0.53}^{+1.04}$ & 2.62$_{-0.38}^{+0.97}$ & 0.28$_{-0.12}^{+0.10}$ & 11.05$\\pm1.92$ & 312 (295) \\\\ 1806$-$20 & 0148210401 & 5.87$_{-0.55}^{+0.63}$ & 0.85$_{-0.11}^{+0.12}$ & 1.97$_{-0.43}^{+0.65}$ & 3.39$_{-0.58}^{+1.2}$ & 0.20$\\pm0.07$ & 12.29$\\pm2.4$ & 506 (459) \\\\ 1806$-$20 & 0205350101 & 5.75$_{-0.19}^{+0.20}$ & 0.96$\\pm0.05$ & 2.06$_{-0.17}^{+0.20}$ & 3.19$_{-0.21}^{+0.28}$ & 0.32$\\pm0.04$ & 25.43$\\pm0.52$ & 2243 (2224) \\\\ 1806$-$20 & 0164561101 & 5.43$_{-0.35}^{+0.39}$ & 1.02$\\pm0.1$ & 1.94$_{-0.27}^{+0.37}$ & 3.92$_{-0.69}^{+1.46}$ & 0.23$\\pm0.08$ & 24.65$\\pm2.74$ & 737 (741) \\\\ 1806$-$20 & 0164561301 & 5.64$_{-0.43}^{+0.50}$ & 0.96$\\pm0.1$ & 1.98$_{-0.33}^{+0.46}$ & 3.67$_{-0.68}^{+1.54}$ & 0.22$\\pm0.08$ & 18.91$\\pm4.21$ & 654 (531) \\\\ 1806$-$20 & 0164561401 & 5.91$_{-0.31}^{+0.34}$ & 0.87$\\pm0.06$ & 2.03$_{-0.25}^{+0.31}$ & 3.27$_{-0.34}^{+0.48}$ & 0.22$\\pm0.04$ & 13.30$\\pm0.5$ & 1026 (1069) \\\\ 2013$+$34 & 0212481201 & 0.29$_{-0.13}^{+0.15}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 0.13$_{-0.02}^{+0.02}$ & $<$7.54 & 20.78$\\pm20.12$ & 57 (64) \\\\ 1819$-$16 & 0152834501 & $1.6_{-0.5}^{+0.7}$ & $1.9_{-0.2}^{+0.3}$ & $0.11_{-0.02}^{+0.03}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 1.3$\\pm0.1$ & 59 (74) \\\\ \\hline \\multicolumn{9}{@{}l@{}}{\\hbox to 0pt{\\parbox{180mm}{\\footnotesize \\footnotemark[$*$] Object name of the SGRs (2013$+$34 denotes a SGR candidate SGR/GRB\\,050925). \\par\\noindent \\footnotemark[$\\dagger$] XMM-Newton observation ID. \\par\\noindent \\footnotemark[$\\ddagger$] $N_{\\mathrm{H}}$ denotes the column density with 90 \\% confidence level errors. \\par\\noindent \\footnotemark[$\\S$] $kT_{\\mathrm{LT}}$ and $kT_{\\mathrm{HT}}$ denote the blackbody temperatures with 90 \\% confidence level errors. \\par\\noindent \\footnotemark[$\\|$] $R_{\\mathrm{LT}}$ and $R_{\\mathrm{HT}}$ denote the emission radii with 90 \\% confidence level errors. \\par\\noindent \\footnotemark[$\\#$] $F$ denotes a flux in the energy range 2-10 keV in units of $10^{-12}$ ergs cm$^{-2}$ s$^{-1}$ with 68 \\% confidence level errors. }\\hss}} \\end{tabular} \\end{center} \\end{table*}  \\begin{table*} \\caption{Spectral parameters of the quiescent emissions of the SGRs observed by {\\it Chandra}.}\\label{tab:sgr_spc_cxo} \\begin{center} \\begin{tabular}{lllllllll} \\hline\\hline Object\\footnotemark[$*$] & ObsID\\footnotemark[$\\dagger$] & $N_{\\mathrm{H}}$\\footnotemark[$\\ddagger$] & $kT_{\\mathrm{LT}}$\\footnotemark[$\\S$] & $R_{\\mathrm{LT}}$\\footnotemark[$\\|$] &  $kT_{\\mathrm{HT}}$\\footnotemark[$\\S$] & $R_{\\mathrm{HT}}$\\footnotemark[$\\|$] & $F$\\footnotemark[$\\#$] & $\\chi^{2}$ (d.o.f.) \\\\ &  & (10$^{22}$ cm$^{-2}$) & (keV)         & (km)         & (keV)         & (km)         &   & \\\\ \\hline 0526$-$66 & 747 & 0.20$_{-0.02}^{+0.03}$ & 0.36$\\pm0.03$ & 10.87$_{-1.26}^{+1.66}$ & 0.98$_{-0.14}^{+0.22}$ & 1.06$_{-0.37}^{+0.46}$ & 0.49$\\pm0.06$ & 180 (181) \\\\ 0526$-$66 & 1957 & 0.26$_{-0.04}^{+0.05}$ & 0.30$\\pm0.04$ & 15.08$_{-2.87}^{+4.93}$ & 0.70$_{-0.07}^{+0.11}$ & 2.24$_{-0.72}^{+0.85}$ & 0.39$\\pm0.06$ & 176 (185) \\\\ 1627$-$41 & 1981 & 8.47$_{-4.86}^{+6.41}$ & 0.89$_{-0.28}^{+0.57}$ & $<$0.67 & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & $<$0.07 & 36 (34) \\\\ 1627$-$41 & 3877 & 17.34$_{-7.36}^{+9.89}$ & 0.42$_{-0.13}^{+0.2}$ & 3.56$_{-2.99}^{+33.39}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & $<$0.06 & 22 (21) \\\\ \\hline \\multicolumn{9}{@{}l@{}}{\\hbox to 0pt{\\parbox{180mm}{\\footnotesize \\footnotemark[$*$] Object name of the SGRs. \\par\\noindent \\footnotemark[$\\dagger$] {\\it Chandra} observation ID. \\par\\noindent \\footnotemark[$\\ddagger$] $N_{\\mathrm{H}}$ denotes the column density with 90 \\% confidence level errors. \\par\\noindent \\footnotemark[$\\S$] $kT_{\\mathrm{LT}}$ and $kT_{\\mathrm{HT}}$ denote the blackbody temperatures with 90 \\% confidence level errors. \\par\\noindent \\footnotemark[$\\|$] $R_{\\mathrm{LT}}$ and $R_{\\mathrm{HT}}$ denote the emission radii with 90 \\% confidence level errors. \\par\\noindent \\footnotemark[$\\#$] $F$ denotes the flux in the 2-10\\,keV band in units of $10^{-12}$ ergs cm$^{-2}$ s$^{-1}$ with 68 \\% confidence level errors. }\\hss}} \\end{tabular} \\end{center} \\end{table*}  \\begin{table*} \\caption{Spectral parameters of a burst of SGR\\,2013$+$34 (GRB/SGR\\,050925) observed by {\\it Swift}.}\\label{tab:sgr_050925_spc_swift} \\begin{center} \\begin{tabular}{llllllll} \\hline\\hline SeqNum\\footnotemark[$*$] & $N_{\\mathrm{H}}$\\footnotemark[$\\dagger$] & $kT_{\\mathrm{LT}}$\\footnotemark[$\\ddagger$] & $R_{\\mathrm{LT}}$\\footnotemark[$\\S$] &  $kT_{\\mathrm{HT}}$\\footnotemark[$\\ddagger$] & $R_{\\mathrm{HT}}$\\footnotemark[$\\S$] & $F$\\footnotemark[$\\|$] & $\\chi^{2}$ (d.o.f.) \\\\ & (10$^{22}$ cm$^{-2}$) & (keV)         & (km)         & (keV)         & (km)         &   & \\\\ \\hline 00156838000 & $\\cdot\\cdot\\cdot$ & 6.6$_{-3.9}^{+5.6}$ & 3.1$_{-1.7}^{+3.7}$ & 20$_{-4}^{+15}$ & $\\gtrsim$0.2 & 0.81$\\pm0.28$ & 27 (25) \\\\ \\hline \\multicolumn{8}{@{}l@{}}{\\hbox to 0pt{\\parbox{180mm}{\\footnotesize \\footnotemark[$*$] {\\it Swift} sequence number. \\par\\noindent \\footnotemark[$\\dagger$] $N_{\\mathrm{H}}$ denotes the column density with 90 \\% confidence level errors. \\par\\noindent \\footnotemark[$\\ddagger$] $kT_{\\mathrm{LT}}$ and $kT_{\\mathrm{HT}}$ denote the blackbody temperatures with 90 \\% confidence level errors. \\par\\noindent \\footnotemark[$\\S$] $R_{\\mathrm{LT}}$ and $R_{\\mathrm{HT}}$ denote the emission radii with 90 \\% confidence level errors. \\par\\noindent \\footnotemark[$\\|$] $F$ denotes a flux in the energy range 15-150 keV in units of $10^{-6}$ ergs cm$^{-2}$ s$^{-1}$ with 68 \\% confidence level errors. }\\hss}} \\end{tabular} \\end{center} \\end{table*}  \\begin{table*} \\small \\caption{Spectral parameters of the quiescent emissions of the AXPs observed by XMM-Newton.}\\label{tab:axp_spc_xmm} \\begin{center} \\begin{tabular}{lllllllll} \\hline\\hline Object\\footnotemark[$*$] & ObsID\\footnotemark[$\\dagger$] & $N_{\\mathrm{H}}$\\footnotemark[$\\ddagger$] & $kT_{\\mathrm{LT}}$\\footnotemark[$\\S$] & $R_{\\mathrm{LT}}$\\footnotemark[$\\|$] &  $kT_{\\mathrm{HT}}$\\footnotemark[$\\S$] & $R_{\\mathrm{HT}}$\\footnotemark[$\\|$] & $F$\\footnotemark[$\\#$] & $\\chi^{2}$ (d.o.f.) \\\\ &       & (10$^{22}$ cm$^{-2}$) & (keV)         & (km)         & (keV)         & (km)         &   & \\\\ \\hline 0100$-$721 & 0110000201 & $\\lesssim0.06$ & 0.39$\\pm0.02$ & 6.94$_{-0.67}^{+1.09}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 0.09$\\pm0.01$ & 264 (275) \\\\ 0100$-$721 & 0018540101 & $\\lesssim0.15$ & 0.29$\\pm0.06$ & 11.64$_{-3.35}^{+7.47}$ & 0.64$_{-0.11}^{+0.25}$ & 1.86$_{-1.08}^{+1.17}$ & 0.13$\\pm0.09$ & 97 (129) \\\\ 0142$+$614 & 0112781101 & 0.53$\\pm0.01$ & 0.36$\\pm0.01$ & 9.38$_{-0.31}^{+0.34}$ & 0.82$_{-0.02}^{+0.03}$ & 0.89$_{-0.08}^{+0.09}$ & 57.26$\\pm0.61$ & 1086 (819) \\\\ 1708$-$400 & 0148690101 & 0.95$\\pm0.01$ & 0.48$\\pm0.01$ & 4.46$_{-0.10}^{+0.11}$ & 1.49$\\pm0.04$ & 0.29$\\pm0.02$ & 27.42$_{-0.13}^{+0.08}$ & 1566 (1232) \\\\ 1841$-$045 & 0013340101 & 1.86$_{-0.13}^{+0.14}$ & 0.52$\\pm0.03$ & 3.79$_{-0.46}^{+0.57}$ & 1.99$_{-0.20}^{+0.25}$ & 0.21$_{-0.04}^{+0.05}$ & 17.39$\\pm0.83$ & 408 (391) \\\\ 1841$-$045 & 0013340201 & 1.90$\\pm0.11$ & 0.51$\\pm0.02$ & 3.97$_{-0.40}^{+0.48}$ & 1.81$_{-0.12}^{+0.14}$ & 0.25$_{-0.03}^{+0.04}$ & 17.16$\\pm0.51$ & 583 (641) \\\\ 2259$+$586 & 0038140101 & $0.59\\pm0.01$ & $0.353_{-0.004}^{+0.005}$ & $4.88_{-0.13}^{+0.15}$ & $0.75\\pm0.02$ & $0.50_{-0.04}^{+0.05}$ & $12.34\\pm0.07$ & 1167 (888) \\\\ 2259$+$586 & 0155350301 & $0.54\\pm0.01$ & $0.390\\pm0.006$ & $5.01_{-0.14}^{+0.15}$ & $0.85\\pm0.02$ & $0.67_{-0.04}^{+0.05}$ & $33.01\\pm0.21$ & 1531 (1053) \\\\ 2259$+$586 & 0203550701 & $0.59\\pm0.03$ & $0.356_{-0.015}^{+0.013}$ & $4.97_{-0.39}^{+0.47}$ & $0.82_{-0.07}^{+0.08}$ & $0.43_{-0.10}^{+0.13}$ & $13.77_{-0.44}^{+0.44}$ & 457 (458) \\\\ \\hline \\multicolumn{9}{@{}l@{}}{\\hbox to 0pt{\\parbox{180mm}{\\footnotesize \\footnotemark[$*$] Object name of the AXPs; CXOU\\,J010043.1$-$721134 (0100$-$721), 4U\\,0142$+$614 (0142$+$614), 1RXS\\,J170849.0$-$400910 (1708$-$400), 1E\\,1841$-$045 and 1E\\,2259$+$586. \\par\\noindent \\footnotemark[$\\dagger$] XMM-Newton observation ID. \\par\\noindent \\footnotemark[$\\ddagger$] $N_{\\mathrm{H}}$ denotes the column density with 90 \\% confidence level errors. \\par\\noindent \\footnotemark[$\\S$] $kT_{\\mathrm{LT}}$ and $kT_{\\mathrm{HT}}$ denote the blackbody temperatures with 90 \\% confidence level errors. \\par\\noindent \\footnotemark[$\\|$] $R_{\\mathrm{LT}}$ and $R_{\\mathrm{HT}}$ denote the emission radii with 90 \\% confidence level errors. \\par\\noindent \\footnotemark[$\\#$] $F$ denotes a flux in the energy range 2-10 keV in units of $10^{-12}$ ergs cm$^{-2}$ s$^{-1}$ with 68 \\% confidence level errors. }\\hss}} \\end{tabular} \\end{center} \\end{table*}  \\begin{table*} \\caption{Spectral parameters of the quiescent emissions of the AXPs observed by {\\it Chandra}.}\\label{tab:axp_spc_cxo} \\begin{center} \\begin{tabular}{lllllllll} \\hline\\hline Object\\footnotemark[$*$] & ObsID\\footnotemark[$\\dagger$] & $N_{\\mathrm{H}}$\\footnotemark[$\\ddagger$] & $kT_{\\mathrm{LT}}$\\footnotemark[$\\S$] & $R_{\\mathrm{LT}}$\\footnotemark[$\\|$] &  $kT_{\\mathrm{HT}}$\\footnotemark[$\\S$] & $R_{\\mathrm{HT}}$\\footnotemark[$\\|$] & $F$\\footnotemark[$\\#$] & $\\chi^{2}$ (d.o.f.) \\\\ &       & (10$^{22}$ cm$^{-2}$) & (keV)         & (km)         & (keV)         & (km)         &   & \\\\ \\hline 0100$-$721 & 1881 & 0.06$_{-0.05}^{+0.06}$ & 0.33$_{-0.04}^{+0.04}$ & 9.65$_{-1.74}^{+2.98}$ & 0.65$_{-0.11}^{+0.23}$ & 1.42$_{-0.85}^{+1.25}$ & 0.12$_{-0.05}^{+0.05}$ & 172 (150) \\\\ 0100$-$721 & 4616 & $<0.04$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 0.39$\\pm0.02$ & 6.84$_{-0.56}^{+1.06}$ & 0.09$\\pm0.01$ & 63 (68) \\\\ 0100$-$721 & 4617 & $<0.41$ & 0.18$_{-0.06}^{+0.25}$ & 21.90$_{-20.07}^{+173.99}$ & 0.44$_{-0.05}^{+21.23}$ & 5.53$_{-5.52}^{+1.24}$ & 0.12$\\pm0.04$ & 66 (64) \\\\ 0100$-$721 & 4618 & $<0.18$ & 0.28$\\pm0.1$ & 10.88$_{-3.73}^{+15.89}$ & 0.51$_{-0.09}^{+1.04}$ & 3.02$_{-2.85}^{+2.67}$ & 0.10$\\pm0.09$ & 50 (65) \\\\ 0100$-$721 & 4619 & $<0.04$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 0.41$_{-0.2}^{+0.02}$ & 6.40$_{-0.51}^{+0.88}$ & 0.11$\\pm0.01$ & 47 (67) \\\\ 0100$-$721 & 4620 & $<0.03$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 0.40$_{-0.02}^{+0.01}$ & 6.67$_{-0.41}^{+0.79}$ & 0.10$\\pm0.01$ & 70 (68) \\\\ 1647$-$455 & 6283 & 2.54$_{-0.69}^{+0.81}$ & 0.49$\\pm0.06$ & 0.52$_{-0.18}^{+0.31}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 0.15$\\pm0.04$ & 23 (21) \\\\ 1647$-$455 & 5411 & 1.44$_{-0.28}^{+0.32}$ & 0.58$\\pm0.05$ & 0.26$_{-0.05}^{+0.07}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 0.13$\\pm0.01$ & 54 (44) \\\\ \\hline \\multicolumn{9}{@{}l@{}}{\\hbox to 0pt{\\parbox{180mm}{\\footnotesize \\footnotemark[$*$] Object name of the AXPs; CXOU\\,J010043.1$-$721134 (0100$-$721) and CXOU\\,J164710.2$-$455216 (1647$-$455). \\par\\noindent \\footnotemark[$\\dagger$] {\\it Chandra} observation ID. \\par\\noindent \\footnotemark[$\\ddagger$] $N_{\\mathrm{H}}$ denotes the column density with 90 \\% confidence level errors. \\par\\noindent \\footnotemark[$\\S$] $kT_{\\mathrm{LT}}$ and $kT_{\\mathrm{HT}}$ denote the blackbody temperatures with 90 \\% confidence level errors. \\par\\noindent \\footnotemark[$\\|$] $R_{\\mathrm{LT}}$ and $R_{\\mathrm{HT}}$ denote the emission radii with 90 \\% confidence level errors. \\par\\noindent \\footnotemark[$\\#$] $F$ denotes a flux in the energy range 2-10 keV in units of $10^{-12}$ ergs cm$^{-2}$ s$^{-1}$ with 68 \\% confidence level errors. }\\hss}} \\end{tabular} \\end{center} \\end{table*}  \\begin{table*} \\caption{Spectral parameters of the quiescent emission and the bursts of AXP\\,CXOU\\,J164710.2$-$455216 observed by {\\it Swift}.}\\label{tab:axp_j1647_spc_swift} \\begin{center} \\begin{tabular}{llllllll} \\hline\\hline SeqNum\\footnotemark[$*$] & $N_{\\mathrm{H}}$\\footnotemark[$\\dagger$] & $kT_{\\mathrm{LT}}$\\footnotemark[$\\ddagger$] & $R_{\\mathrm{LT}}$\\footnotemark[$\\S$] &  $kT_{\\mathrm{HT}}$\\footnotemark[$\\ddagger$] & $R_{\\mathrm{HT}}$\\footnotemark[$\\S$] & $F$\\footnotemark[$\\|$] & $\\chi^{2}$ (d.o.f.) \\\\ & (10$^{22}$ cm$^{-2}$) & (keV)         & (km)         & (keV)         & (km)         &   & \\\\ \\hline 00230341000\\footnotemark[$\\#$] & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 8.92$_{-1.62}^{+1.34}$ & 1.21$_{-0.35}^{+0.59}$ & 0.35$\\pm0.12$ & 8 (12) \\\\ 00030806001 & 1.72$_{-0.20}^{+0.31}$ & 0.63$_{-0.12}^{+0.07}$ & 2.40$_{-0.44}^{+0.94}$ & 2.08$_{-0.93}^{+38.74}$ & 0.14$_{-0.13}^{+0.36}$ & 27.94$\\pm10.25$ & 253 (242) \\\\ 00030806002 & 1.76$_{-0.55}^{+0.70}$ & 0.73$_{-0.08}^{+0.08}$ & 1.79$_{-0.44}^{+0.7}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 20.01$\\pm4.36$ & 16 (17) \\\\ 00030806003 & 1.95$_{-0.53}^{+0.86}$ & 0.42$_{-0.15}^{+0.26}$ & 3.03$_{-1.9}^{+9.21}$ & 0.81$_{-0.09}^{+0.35}$ & 1.04$_{-0.81}^{+0.43}$ & 13.98$\\pm6.28$ & 101 (92) \\\\ 00030806004 & 1.79$_{-0.7}^{+0.91}$ & 0.51$_{-0.21}^{+0.28}$ & 2.20$_{-1.66}^{+4.56}$ & 0.93$_{-0.38}^{+3.06}$ & 0.68$_{-0.67}^{+0.5}$ & 13.79$\\pm12.1$ & 91 (74) \\\\ 00030806006 & 0.99$_{-0.32}^{+0.37}$ & 0.75$\\pm0.06$ & 1.24$_{-0.23}^{+0.31}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 11.99$\\pm1.2$ & 31 (33) \\\\ 00030806007 & 1.65$_{-0.39}^{+0.49}$ & 0.67$_{-0.06}^{+0.07}$ & 1.52$_{-0.33}^{+0.48}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 9.85$\\pm1.34$ & 27 (29) \\\\ 00030806008 & 1.78$_{-0.42}^{+0.55}$ & 0.60$\\pm0.05$ & 2.02$_{-0.46}^{+0.69}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 9.18$\\pm1.33$ & 73 (64) \\\\ 00030806009 & 2.09$_{-0.46}^{+0.57}$ & 0.62$\\pm0.06$ & 1.69$_{-0.39}^{+0.58}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 6.86$\\pm1.38$ & 64 (55) \\\\ 00030806010 & 1.51$_{-0.32}^{+0.4}$ & 0.71$\\pm0.06$ & 1.31$_{-0.25}^{+0.34}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 9.48$\\pm0.81$ & 47 (59) \\\\ 00030806011 & 1.35$_{-0.24}^{+0.28}$ & 0.67$\\pm0.04$ & 1.38$_{-0.20}^{+0.26}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 7.93$\\pm0.61$ & 90 (93) \\\\ 00030806012 & 1.30$_{-0.24}^{+0.29}$ & 0.67$\\pm0.05$ & 1.22$_{-0.20}^{+0.25}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 6.60$\\pm0.45$ & 66 (75) \\\\ 00030806013 & 1.79$_{-0.40}^{+0.49}$ & 0.63$\\pm0.06$ & 1.57$_{-0.36}^{+0.54}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 7.00$\\pm0.97$ & 63 (70) \\\\ 00030806014 & 1.57$_{-0.51}^{+0.72}$ & 0.63$\\pm0.09$ & 1.26$_{-0.38}^{+0.68}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 4.64$\\pm1.43$ & 22 (31) \\\\ 00030806015 & 1.80$_{-0.42}^{+0.53}$ & 0.61$_{-0.06}^{+0.07}$ & 1.38$_{-0.34}^{+0.53}$ & $\\cdot\\cdot\\cdot$ & $\\cdot\\cdot\\cdot$ & 4.77$\\pm0.84$ & 49 (47) \\\\ \\hline \\multicolumn{8}{@{}l@{}}{\\hbox to 0pt{\\parbox{180mm}{\\footnotesize \\footnotemark[$*$] {\\it Swift} sequence number. \\par\\noindent \\footnotemark[$\\dagger$] $N_{\\mathrm{H}}$ denotes the column density with 90 \\% confidence level errors. \\par\\noindent \\footnotemark[$\\ddagger$] $kT_{\\mathrm{LT}}$ and $kT_{\\mathrm{HT}}$ denote the blackbody temperatures with 90 \\% confidence level errors. \\par\\noindent \\footnotemark[$\\S$] $R_{\\mathrm{LT}}$ and $R_{\\mathrm{HT}}$ denote the emission radii with 90 \\% confidence level errors. \\par\\noindent \\footnotemark[$\\|$] $F$ denotes fluxes in the energy ranges 15-150 keV in units of $10^{-6}$ ergs cm$^{-2}$ s$^{-1}$ for the burst observation of 00230341000 and 2-10 keV in units of $10^{-12}$ ergs cm$^{-2}$ s$^{-1}$ for other observations with 68 \\% confidence level errors. \\par\\noindent \\footnotemark[$\\#$] Results for the burst. }\\hss}} \\end{tabular} \\end{center} \\end{table*}  \\clearpage  %%% % See the manual for the detail. %%%"
    },
    "0710/0710.1778_arXiv.txt": {
        "abstract": "Using Magellan/IMACS images covering a 1.2 x 1.2 sq. degree FOV with seeing of 0.4\"-0.6\", we have applied convolution techniques to analyse the light distribution of 364 confirmed globular cluster in the field of NGC 5128 and to obtain their structural parameters. Combining these parameters with existing Washington photometry from Harris et al. (2004), we are able to examine the size difference between metal-poor (blue) and metal-rich (red) globular clusters. For the first time, this can be addressed on a sample of confirmed clusters that extends to galactocentric distances about 8 times the effective radius, R$_{eff}$, of the galaxy. Within 1 R$_{eff}$, red clusters are about $30\\%$ smaller on average than blue clusters, in agreement with the vast majority of extragalactic globular cluster systems studied.  As the galactocentric distance increases, however, this difference becomes negligible. Thus, our results indicate that the difference in the clusters' effective radii, r$_e$, could be explained purely by projection effects, with red clusters being more centrally concentrated than blue ones and an intrinsic r$_e$--R$_{gc}$ dependence, like the one observed for the Galaxy. ",
        "introduction": "\\label{sec:intro} Since sizes and structural parameters of globular clusters (GCs) in different GC systems (GCSs) have first been obtained, it has become clear that some of these properties correlate with global properties of their host galaxies \\citep[see for example][]{jordan05,brodie06}. The existence of the so called fundamental plane relation for an increasing number of studied GCSs seems to confirm that GCs populate a narrow region in this parameter space \\citep{djorgovski95,mclaughlin00,mclaughlin05,barmby07}. However, there are puzzling trends that are still awaiting confirmation and need to be addressed using larger samples of GCs.  It is necessary to study structural parameters of GCs and GC-like objects in different environments before definitive statements can be made regarding their formation. Among the structural parameters that can be studied, the effective (or half-light) radius is of particular importance. Models have shown that this quantity remains fairly constant throughout the entire GC lifetime \\citep{spitzer72,aarseth98}, making it a good indicator of proto-GC sizes that are still observable today. A decade ago, HST observations unveiled a systematic size difference between red and blue GCs \\citep{kundu98}.  Since then, multiple studies have found that the blue GCs are between $17\\%-30\\%$ larger than their metal-rich counterparts in both spirals and early--type galaxies \\citep{kundu99,puzia99,larsen01,larsen_fb01,kundu01,barmby02,jordan05}. However, most of these studies have made use of HST observations and examine only the innermost regions of the galaxy or small fields in regions at galactocentric distances greater than the galaxy's effective radius.  According to \\cite{larsen03}, the systematic size difference between red and blue GCs is caused merely by a projection effect. Since red (metal-rich) GCs are found to be more centrally concentrated than blue (metal-poor ones) in early type galaxies \\citep[][among others]{cote01,dirsch03,woodley05}, the red GCs will appear to lie, on average, at a smaller galactocentric distance.  The red clusters will on average be smaller than the blue clusters assuming that both types shares the same relation between the GC size and galactocentric distance.  The relation r $\\sim \\sqrt{\\rm{R}_{gc}}$ was first found in the Milky Way by \\cite{vandenbergh91}.  In this scenario, the difference between the cluster sizes should be most apparent at small galactocentric distance and should decrease strongly beyond 1 galaxy effective radius \\citep{larsen03}.  Alternatively, \\cite{jordan04} suggests that this effect could be explained by an intrinsic difference between metal-rich and metal-poor GCs. Assuming half-mass radii that are independent of metallicity, effects of mass segregation combined with a metallicity-dependent stellar lifetime should lead to different sizes between the blue and red clusters.  The brightest stars would be more massive and more centrally concentrated for the metal-rich GCs.  This scenario should have little to no dependence on a cluster's distance from the center of its parent galaxy.  In a recent study, \\cite{spitler06} analysed the GCS of NGC 4594 (Sombrero, at a distance of $ \\sim 9$ Mpc) using a six-image mosaic from HST/ACS. They confirm that within the inner 2 arcmin (2.2 R$_{eff}$), the metal-rich GCs are, on average, $17\\%$ smaller than the metal-poor clusters. However, the size difference becomes negligible at $\\sim 3$ arcmin, corresponding to $\\sim 3.4$ R$_{eff}$, where R$_{eff} =0.89$ arcmin \\citep{baggett98}.  To further understand the sizes of red and blue clusters, we need a homogeneous survey of a GCS with the ability to eliminate contaminating sources, high resolution to measure structural parameters, and over a large range in galactocentric distance.  NGC 5128 is the nearest giant elliptical galaxy, at a distance of 3.8 Mpc \\citep{mclaughlin07}.  Its GCs are thus easily resolvable with sub-arcsecond seeing \\citep{harris06}. In this paper we present effective radius results for 337 GCs from the \\cite{woodley07} catalog that are confirmed GCs by either radial velocity measurement from various studies \\citep[see the references in ][]{woodley07} or are resolved by HST/ACS images \\citep{harris06}.  We also present the effective radii of 27 GCs newly confirmed through radial velocity measurements using the Baade 6.5-m telescope with the instrument LDSS-2 (data in preparation for publication). This list represents a clean sample of confirmed clusters.  All of these also have ellipticities less than 0.4 and effective radii less than 8 pc, both of which are consistent with {\\it normal} GC properties in NGC 5128.  We find that only an additional $2.4\\%$ of GCs from the \\cite{woodley07} catalog have effective radii greater than the 8 pc boundary we have imposed here (to be discussed in detail in  G\\'omez \\& Woodley, 2008, in preparation). Those few GCs are not considered here as our purpose is to establish the effective radius trends within the bulk of the GC population. ",
        "conclusions": "\\label{sec:conclusions}  Using a contaminant-free sample of 364 GCs in NGC 5128, confirmed with radial velocity measurements or by resolved HST images, we have measured effective radii using ISHAPE.  Our results indicate that the blue or metal-poor clusters do not show any significant r$_{e}$--R$_{gc}$ relation.  However, the red or metal-rich GCs do show a steep relation in which red clusters within 1 R$_{eff}$ of the galaxy's light are $30\\%$ smaller than the blue clusters.  Beyond this distance there is no indication for a size difference between the two metallicity populations.  This finding in NGC 5128, not previously seen in any other early-type galaxy, supports the more tentative findings of the Sombrero galaxy's GCS \\citep{spitler06}.  Both studies support the idea that the size differences are most likely caused by projection effects \\citep{larsen03} and not by intrinsic physical differences between the two subgroups.   Acknowledgements: M.G. and K.A.W. thank Dean McLaughlin for use of HST structural parameters in advance of publication.  M.G. thanks the Dept. of Physics and Astronomy at McMaster University and especially Bill and Gretchen Harris for their hospitality.  K.A.W. thanks NSERC and Bill Harris for financial support, and also the Depto. de F{\\'i}sica at the Universidad de Concepci{\\'o}n, especially Doug Geisler, for their hospitality. We thank the anonymous referee for her/his valuable suggestions and comments.   \\clearpage"
    },
    "0710/0710.3850_arXiv.txt": {
        "abstract": "The first map of interstellar acetylene (C$_2$H$_2$) has been obtained with the infrared spectrograph onboard the {\\it Spitzer Space Telescope}. A spectral line map of the $\\nu_5$ vibration-rotation band at 13.7 $\\mu$m carried out toward the star-forming region Cepheus A East, shows that the C$_2$H$_2$ emission peaks in a few localized clumps where gas-phase CO$_2$ emission was previously detected with {\\it Spitzer}. The distribution of excitation temperatures derived from fits to the C$_2$H$_2$ line profiles ranges from 50 to 200 K, a range consistent with that derived for gaseous CO$_2$ suggesting that both molecules probe the  same warm gas component. The C$_2$H$_2$ molecules are excited via radiative pumping by 13.7 $\\mu$m continuum photons emanating from the HW2 protostellar region. We derive column densities ranging from a few $\\times$ 10$^{13}$ to $\\sim$ 7 $\\times$ 10$^{14}$ cm$^{-2}$, corresponding to C$_2$H$_2$ abundances of 1 $\\times$ 10$^{-9}$ to 4 $\\times$ 10$^{-8}$ with respect to H$_2$. The spatial distribution of the C$_2$H$_2$ emission along with a roughly constant $N$(C$_2$H$_2$)/$N$(CO$_2$) strongly suggest an association with shock activity, most likely the result of the sputtering of acetylene in icy grain mantles. ",
        "introduction": "Acetylene constitutes a key ingredient in the production of large complex hydrocarbon molecules in the dense interstellar medium (Herbst 1995). Because acetylene has no permanent dipole moment, it lacks a rotational spectrum that could be observed at radio wavelengths; observations of interstellar acetylene have therefore been limited to mid-infrared studies of rovibrational bands, carried out from ground-(e.g., Evans, Lacy \\& Carr 1991; Carr et al. 1995) and space-based (e.g., Lahuis \\& van Dishoeck 2000; Boonman et al. 2003; Lahuis et al. 2007) observatories. Acetylene has been detected in the gas-phase - either in absorption (e.g., Carr et al. 1995; Lahuis \\& van Dishoeck 2000) or in emission (e.g. Boonman et al. 2003, this paper) - mostly toward young stellar objects. C$_2$H$_2$ can be used as a tracer of warm (100 K to 1000 K) molecular gas along often complicated sightlines. C$_2$H$_2$ abundance estimates, which were sometimes a few orders of magnitude higher than the predictions of cold gas-phase steady-state chemical models, have led to a better understanding of the role that warm gas-phase chemistry (e.g., Doty et al. 2002; Rodgers \\& Charnley 2001) and/or grain mantle processing (e.g., Ruffle \\& Herbst 2000) can play in star-forming regions, both locally and in extra-galactic objects (Lahuis et al. 2007).  In this Letter, we present the first detection of the $\\nu_5$ band of acetylene (C$_2$H$_2$) at 13.7 $\\mu$m toward the star forming region Cepheus A East using the Infrared Spectrograph (IRS) onboard the {\\it Spitzer Space Telescope}. This is the first map of C$_2$H$_2$ obtained toward any interstellar gas cloud. Section 2 describes the observations and data analysis. Sections 3 and 4 compare the spatial distribution of C$_2$H$_2$  to those of gaseous CO$_2$ and H$_2$ $S$(2), and discuss the C$_2$H$_2$-emitting gas in the context of shock chemistry and local outflow activity. The presence of C$_2$H$_2$ on interstellar dust grains will also be discussed in the context of cometary ices composition. ",
        "conclusions": "Acetylene was previously observed toward low-to-high mass star-forming regions either in absorption (e.g., Lahuis \\& van Dishoeck 2000) or in emission (e.g., Boonman et al. 2003) using the {\\it Infrared Space Observatory} ({\\it ISO}). Excitation temperatures ranging from $\\sim$ 10 to 900 K and abundances with respect to H$_2$ ranging from a few $\\times$ 10$^{-8}$ to a few $\\times$ 10$^{-7}$ were derived. Steady-state models of gas-phase chemistry in cold (10-50K) dense (10$^3$-10$^5$ cm$^{-3}$) molecular clouds predict abundances for acetylene between a few $\\times$ 10$^{-10}$ and 1$\\times$ 10$^{-8}$ with respect to H$_2$ depending on the role that neutral-neutral destruction reactions may play (Bettens, Lee \\& Herbst 1995; Lee et al. 1996). Similar abundances are predicted by models of gas-grain chemistry in quiescent clouds with the highest values obtained only after 10$^6$ years (Ruffle \\& Herbst 2000). While such models could account for the observed abundances in the cold gas, they were unable to reproduce the enhancements observed toward much warmer gas components in those objects. Mechanisms such as C$_2$H$_2$ ice sublimation from grain mantles and/or C$_2$H$_2$ enhancements via warm gas-phase chemistry were then invoked (e.g., Carr et al. 1995; Doty et al. 2002; Boonman et al. 2003).  For the NE outflow region in Cepheus A East, we derive C$_2$H$_2$/H$_2$ abundance ratios in the range 1 $\\times$ 10$^{-9}$ to 4 $\\times$ 10$^{-8}$, for an assumed H$_2$ column density of 1.5 $\\times$ 10$^{22}$ cm$^{-3}$ (G\\'omez et al. 1999).  These values are averages along the observed sight-lines.  The highest abundances are localized around and to the south of, the NE position as well as around the positions of the HW5/6 sources (see Fig.~2).  This is precisely where the interaction between the {\\it northeast} outflow and the ambient molecular clouds occurs, and it is in these regions that we observe strong H$_2$ $S$(2) emission, a tracer of warm shocked gas. Thus the spatial variation of the C$_2$H$_2$ abundance again strongly suggests an association with shock activity, perhaps as a result of (1) production in the gas phase via high temperature reactions, or (2) grain mantle sputtering.  Models for chemistry in hot cores (Rodgers \\& Charnley 2001; Doty et al. 2002) indicate that enhanced abundances of C$_2$H$_2$ ($\\sim$ few $\\times$ 10$^{-8}$) are expected in warm regions with $T \\ge$ 200 K. The good correlations between H$_2$ $S$(2), gaseous CO$_2$ and C$_2$H$_2$ indicate that such high temperatures were reached in the warm gas at the passage of the non-dissociative shock. However, the chemical pathways leading to the production of C$_2$H$_2$ are slow, and enhanced abundances only occur after $\\sim$ 10$^4$ years, a time scale much greater than that expected for shock heating of the gas ($\\sim$ 300 years; e.g. Kaufman \\& Neufeld 1996). Hence, enhanced production of C$_2$H$_2$ by high temperature gas-phase chemistry is unlikely to be predominant in the observed region. Thus the correlations shown in Fig.~3 argue in favor of grain mantle sputtering over gas-phase production as the origin of the C$_2$H$_2$ in Cepheus A East. This is the same production mechanism that we favored for gaseous CO$_2$ (Sonnentrucker et al. 2006). While models for the production of C$_2$H$_2$ in shocks are not available to our knowledge, our results further suggest that both C$_2$H$_2$ and CO$_2$ are released into the gas phase under very similar physical conditions.  The gaseous $\\rm C_2H_2/CO_2$ ratio is roughly constant, with a mean value of 0.08 for all sight-lines where we detected acetylene at the 1.5 $\\sigma$ level.  If this value reflects the composition of the grain mantle, and given a $N$(CO$_2$)$_{ice}$/$N$(H$_2$O)$_{ice}$ ratio in this source of 0.22 (Sonnentrucker et al. 2007, ApJ in press), then the required $N$(C$_2$H$_2$)$_{ice}$/$N$(H$_2$O)$_{ice}$ ratio is 0.02.  This value is at least a factor 4 larger than those derived toward other star-forming regions (0.1-0.5\\%, Evans et al. 1991; Lahuis \\& van Dishoeck 2000) and those predicted by theoretical models (0.1-0.5\\%, Hasegawa \\& Herbst 1993; Ruffle \\& Herbst 2000), and at least a factor 2 larger than the gaseous C$_2$H$_2$/H$_2$O ratios obtained in observations of cometary comae (0.1-0.9\\%, Brooke et al. 1996; Weaver et al. 1999).  We speculate that these discrepancies might result from (1) the destruction of CO$_2$ by reaction with atomic hydrogen in shocks faster than $\\sim$ 30 km s$^{-1}$ (predicted by Charnley \\& Kaufman, 2000) and/or (2) a greater efficiency for sputtering of C$_2$H$_2$ in slow shocks.\\footnote{Although both these effects would be strongly dependent upon the shock velocity, the relative constancy of the $\\rm C_2H_2/CO_2$ would not necessarily require any fine tuning of the shock velocity.  In reality, any sight-line typically samples an ensemble of shocks with a {\\it range} of shock velocities, and the constancy of the $\\rm C_2H_2/CO_2$ would simply indicate that the admixture of shock velocities varies little from one sight-line to another (e.g. Neufeld et al.\\ 2006).}  In either case, the gaseous C$_2$H$_2$/CO$_2$ ratios we observed may exceed the solid C$_2$H$_2$/CO$_2$ ratio, and the $N$(C$_2$H$_2$)$_{ice}$/$N$(H$_2$O)$_{ice}$ ratio could be less than 0.02.  Unfortunately, direct measurements of acetylene ice are not possible, the weak features expected from solid C$_2$H$_2$ being blended with much stronger features of CO and H$_2$O (Boudin et al. 1998).  However, further observations of gaseous C$_2$H$_2$ at a higher signal-to-noise ratio would be very valuable as a probe of any variations in the $\\rm C_2H_2/CO_2$ ratio which might provide important clues to the shock physics."
    },
    "0710/0710.1064_arXiv.txt": {
        "abstract": "In this paper we report on the gas-phase abundance of singly-ionized magnesium (Mg II) in 44 lines of sight, using data from the {\\it Hubble Space Telescope} ({\\it HST}).  We measure Mg II column densities by analyzing medium- and high-resolution archival STIS spectra of the 1240 \\AA{} doublet of Mg II.  We find that Mg II depletion is correlated with many line of sight parameters (e.g.~$\\fHmol$, $E_{B-V}$, $\\ebvdist$, $A_V$, and $\\avdist$) in addition to the well-known correlation with $\\nHavg$.  These parameters should be more directly related to dust content and thus have more physical significance with regard to the depletion of elements such as magnesium.  We examine the significance of these additional correlations as compared to the known correlation between Mg II depletion and $\\nHavg$.  While none of the correlations are better predictors of Mg II depletion than $\\nHavg$, some are statistically significant even assuming fixed $\\nHavg$.  We discuss the ranges over which these correlations are valid, their strength at fixed $\\nHavg$, and physical interpretations. ",
        "introduction": "\\label{s:MgII_intro} Magnesium is both a relatively abundant element in the Galaxy and an important component in most interstellar dust models.  Mg I has an ionization potential of only 7.65 eV, while Mg II has an ionization potential of 15.04 eV.  In H I regions, the dominant form of gas-phase interstellar magnesium should be Mg II.  In H II regions, magnesium should be found primarily in the form of Mg III.  Gas-phase Mg I is rare, even in $\\Hmol$ regions.  As is the case for other elements such as silicon and iron, the average gas-phase abundance of magnesium is much smaller than the assumed overall cosmic abundance of magnesium, implying that the majority of interstellar magnesium is tied up in dust.  Therefore, variations in the gas-phase magnesium abundance are only capable of having minor effects on grain composition.  Nevertheless, observed variations in the gas-phase magnesium abundance may still shed light on the physical conditions of interstellar clouds.    Two major {\\it Copernicus} surveys that included measurements of Mg II abundances and depletions were \\citet{Murray} and \\citet{Jenkins1986}.  Both of these studies confirmed that magnesium depletion increases with increased average hydrogen volume density, $\\nHavg=\\NHtot/r$ (where $r$ is the line-of-sight pathlength), in the line of sight.  \\citet{Jenkins1986} also found several correlations between the depletions of magnesium and other elements, which they cited as secondary to the correlation between those depletions and $\\nHavg$.  In both studies, the depletion of magnesium was not strongly correlated to other line of sight parameters, such as $\\ebv$ and the magnitude of the 2175 \\AA{} extinction ``bump''.  We also note that there is a systematic difference in the absolute values of the abundances and depletions between these studies and more recent studies that is fairly substantial.  This is due to a difference in the assumed $f$-values of the 1240 \\AA{} doublet of magnesium.  We comment on our chosen $f$-values in \\S \\ref{ss:MgII_lines}.  More recently, \\citet{Cartledge2006} also examined the abundances and depletions of Mg II and several other elements in the interstellar medium using STIS data.  (Hereafter, this paper will be referred to as CLMS, for the initials of the authors.)  Similar to the {\\it Copernicus} studies, CLMS concluded that the line of sight parameter with the clearest connection to elemental depletions is the average hydrogen volume density, $\\nHavg$.  CLMS also explored potential correlations between depletions and other line of sight parameters such as the molecular fraction of hydrogen, $\\fHmol$; selective extinction, $\\ebv$; and selective extinction divided by line-of-sight pathlength, $\\ebvdist$.  CLMS concluded that $\\nHavg$ was the parameter that best identified warm vs.~cold clouds, and that no other parameters produced correlations with magnesium depletion that were both as strong and with as little scatter.  We began this study before CLMS was published.  In spite of the similarities between that study and this one, we have proceeded with this study to provide an independent analysis of Mg II, and also because we have analyzed potential trends with respect to parameters not analyzed in CLMS, e.g.~$\\av$ and $\\rv$.  The Mg II column densities of 11 out of the 44 lines of sight in our sample have also been analyzed previously by CLMS.  We still report our results in this paper for two main reasons:  (1) these lines of sight provide a basis of comparison for the methods of this paper and those of CLMS and (2) these lines of sight can be analyzed with respect to the aforementioned line of sight parameters not analyzed in CLMS.  A comparison of the column density measurements for these common lines of sight is found in \\S \\ref{ss:MgII_coldensities}.  In \\S \\ref{s:MgII_obsdata} we discuss our observations and data reduction, including comments on the 1240 \\AA{} doublet and our derivation of column densities and abundances.  In \\S \\ref{s:results} we discuss our results, including observed correlations and a review of the Galactic abundance of magnesium.  In \\S \\ref{s:MgII_summary} we summarize our findings. ",
        "conclusions": "\\label{s:results}  \\subsection{Correlations} \\label{ss:correlations} Our results for the column densities of Mg II are given in Table \\ref{MgII_coldensities}.  The major correlation found by both \\citet{Jenkins1986} and CLMS is that the depletion of magnesium and many other elements increases with overall line of sight density, $\\nHavg$.  Both interpreted their models in light of the model by \\citet{Spitzer1985}.  The \\citeauthor{Spitzer1985} model contends that the ISM contains two distinct varieties of cloud (warm and cold), each with a distinct depletion level.  In this model, lines of sight with average densities of a few-tenths of a particle per cm$^{-3}$ are largely sampling the warm ISM, while lines of sight with average densities of a few particles per cm$^{-3}$ are largely sampling the cold ISM.  This explains the observed plateaus of depletion at both low and high densities, with a transition near average densities of $\\sim1$ cm$^{-3}$, where lines of sight are sampling both types of clouds.  This model, however, does not take into account the transition to a much different regime of cloud chemistry expected for translucent clouds \\citep[for a recent review, see][]{SnowMcCall}.  \\citet{Jenkins1986} also found evidence of lesser correlations between the depletion of various elements and other parameters of reddening and extinction, but concluded that these correlations were secondary to the correlation with $\\nHavg$.  CLMS examined the depletion of Mg II compared to $\\fHmol$, $\\ebv$, and $\\ebvdist$.  CLMS concluded that only $\\ebvdist$ provided a correlation with magnesium depletion that was as at least roughly as significant as the correlation with $\\nHavg$, but with increased scatter.  However, in our recent work \\citep{JensenFeII}, we explored the possibility of correlations between iron depletion and various other line of sight parameters.  We found very similar correlations between iron depletion and measures of dust density ($\\ebvdist$ and $\\avdist$) and the molecular fraction of hydrogen ($\\fHmol$).  We also examined iron depletion as a function of $\\rv$, the ratio of total visual extinction to selective extinction, though a conclusive trend did not present itself.  Using many of the same lines of sight in this paper as in \\citet{JensenFeII}, we find the same correlations generally hold with respect to magnesium depletion.  In what follows, we discuss the nature of these correlations and possible interpretations.  \\subsubsection{Correlations with Hydrogen} \\label{sss:hydrogen_correlations} First, we note that magnesium depletion is clearly correlated with total hydrogen volume density $\\nHavg$.  Five lines of sight (HD 27778, HD 37021, HD 37061, HD 37903, and HD 147888) have substantially larger depletions than any of the other lines of sight in this sample; HD 27778, HD 37021, HD 37061, and HD 147888 are the four densest lines of sight (in terms of $\\nHavg$), while the density of HD 37903 is in the top 20\\% of our sample.  This correlation is plotted in Figure \\ref{fig:logMgIIHlognh}.  Since these dense lines of sight are also some of the lines of sight that are found in both our sample and that of CLMS, we are not probing significantly higher average density and cannot expand on their conclusions regarding $\\nHavg$.  We also find that magnesium depletion is correlated with total hydrogen column density.  However, we conclude that this is primarily a secondary correlation due to the correlation with $\\nHavg$.  However, we can attempt to analyze whether or not the correlation is independently significant.  To do this, we follow the methods described in \\citet{Jenkins1986}.  First, we calculate Pearson correlation coefficients between depletion and the two variables of interest (in this case $\\logHtot$ and $\\lognHavg$), as well as those two variables with each other.  The partial correlation coefficient is then given by $\\rho_{12.3} = (\\rho_{12}-\\rho_{13}\\rho_{23})/[(1-\\rho_{13}^2)(1-\\rho_{23}^2)]^{-1/2}$, where the subscripts on the correlation coefficients indicate the two variables being correlated, and $\\rho_{12.3}$ is the correlation coefficient between the first two variables if the third variable is held fixed.  However, what is really being calculated is $r$, the sample correlation coefficient(s), as opposed to $\\rho$, the population correlation coefficient(s).  Once $r_{12.3}$ has been calculated, we examine the significance level for a $t$-test of the appropriate number of degrees of freedom (in this case, the number of data points minus three) to determine the probability that $\\rho$ is non-zero (i.e.~a true correlation exists).  Using these methods, we find a 49\\% chance of the null hypothesis (i.e.~the two-sided probability that $\\rho=0$) for a correlation between Mg II depletion and $\\NHtot$ when $\\nHavg$ is held fixed.  Because this technique of trivariate analysis uses Pearson correlation coefficients, a few caveats apply, namely that linear relationships and normal distributions are implicitly assumed.  We used $\\logMgIIH$, $\\logHtot$, and $\\lognHavg$ in the analysis just discussed because the correlation between all combinations of those variables are stronger than when in linear form.  CLMS also explored possible correlations between magnesium depletion and $\\fHmol$, but only briefly commented on the results.  They concluded that any correlation was less significant than the correlation between magnesium depletion and $\\nHavg$, in that it did not as effectively discriminate between the distinct depletion levels expected in the \\citet{Spitzer1985} model.  The upper right panel of Figure 9 in CLMS shows a very clear correlation between magnesium depletion and $\\fHmol$ for $\\fHmol \\gtrsim 0.1$, superimposed with a scatter plot in depletion for a smaller subset of lines with $0.01 \\lesssim \\fHmol \\lesssim 0.1$.  Two additional lines of sight with $\\fHmol \\lesssim 10^{-4}$ conform to the main correlation in that they show minimal depletion.  We also see a clear correlation between magnesium depletion and the molecular fraction of hydrogen, $\\fHmol$, plotted in Figure \\ref{fig:logMgIIHHf}, with depletion increasing with increasing $\\fHmol$.  There is the exception of one discrepant point, HD 147888, that exhibits substantial depletion at $\\fHmol \\sim 0.1$.  We note that CLMS found a larger column density for this line of sight than we do ($0.2\\dex$); however, even if we adopt the CLMS Mg II column density for this line of sight, its abundance is still $0.26\\dex$ smaller than any line of sight in our sample with $\\fHmol < 0.4$.  We concur with CLMS that the correlation between magnesium depletion and $\\fHmol$ is not as rigorous as the correlation with $\\nHavg$.  We examine the data using various combinations of the variables in their logarithmic (where correlations are the strongest) and linear forms.  In all cases, we find that the probability of the null hypothesis (i.e.~$\\rho=0$, as discussed above) between Mg II depletion and $\\fHmol$ with $\\nHavg$ held constant is less than 5\\%; in most cases, it is $\\lesssim1$\\%.  Conversely, the probability that there is no correlation between Mg II depletion and $\\nHavg$ with $\\fHmol$ held constant is less than 0.1\\%.  Therefore, we conclude that while $\\nHavg$ is clearly the dominant correlation, the correlation with $\\fHmol$ has some independent significance.  A question of interest is where scatter seems to be introduced and why.  Between Figure 9 of CLMS and Figure \\ref{fig:logMgIIHHf} of this paper, scatter is only observed at $\\fHmol \\lesssim 0.1$.  In our recent related work on Fe II \\citep{JensenFeII}, we note that Fe II depletion is also correlated with $\\fHmol$, but there are a few exceptions to the trend at both high and low values of $\\fHmol$.  The most severe exceptions noted in that paper were HD 147888 and HD 164740 (with large depletions but $\\fHmol \\lesssim 0.1$) and HD 210121 (with less depletion despite $\\fHmol \\sim 0.7$).  For Mg II, however, we do not see any outlying points at values of $\\fHmol$ larger than $\\sim0.1$ within either this sample or CLMS.  Points such as HD 147888, however, still require explanation.  \\citet{Snow1983} put forth the possibility that in some dense environments an increased average grain size, which decreases the grain surface area per unit volume, may suppress $\\Hmol$ formation (as $\\Hmol$ is thought to form on grain surfaces).  This scenario was specifically discussed in the context of the $\\rho$ Oph cloud, for which there is indepedent evidence (in part, a value of $\\rv$ greater than the interstellar average of 3.1) that grain coagulation has occurred.  Our main outlying line of sight, HD 147888 (with $\\rv$ of 4.06), passes through the $\\rho$ Oph cloud, so observing the combination of a dense, depleted environment with small $\\fHmol$ is not surprising in this case.  Two of the other lines of sight with large depletions are HD 37021 and HD 37061.  Reliable far-ultraviolet data sets do not exist for these lines of sight; therefore, they do not have measurements of the molecular hydrogen column densities or subsequently derived molecular fractions of hydrogen.  However, as stated above in \\S \\ref{ss:hydrogen}, \\citet{Cartledge2001} has argued that these lines of sight have small values of $\\fHmol$ based on a lack of Cl I.  These two lines of sight also have large values of $\\rv$ (5.54 and 4.23, respectively), implying a larger average grain size.  The possible effect of grain size on depletions and $\\Hmol$ formation, and other interpretations, will be discussed further in \\S \\ref{sss:extinction_correlations}.  Barring such outlying points, the trend of increased Mg II depletion with increased $\\fHmol$ has a relatively clear interpretation.  $\\Hmol$ is formed in the same dense, dusty environments that foster large depletions.  Within the context of this sample and CLMS, this seems to hold for lines of sight with $\\fHmol \\gtrsim 0.1$.  While the ubiquity of $\\Hmol$ even in diffuse regions complicates the issue \\citep[see conclusions of][]{Rachford2002}, there is still a physical argument that $\\fHmol$ should be a good diagnostic of the local conditions of interstellar clouds, and therefore depletions.  It is worth noting that the similar trend between iron depletion and $\\fHmol$ exhibits scatter up to $\\fHmol \\sim 0.3$ in the work of \\citet{SavageBohlin}, in addition to the outlying points mentioned above from \\citet{JensenFeII}.  Whether the range in $\\fHmol$ over which there is scatter in the abundances is truly different for magnesium and iron or is simply a selection effect is unclear.   \\subsubsection{Correlations with Extinction and Reddening Parameters} \\label{sss:extinction_correlations} We find that the depletion of Mg II is correlated to both selective extinction, $\\ebv$, and total visual extinction, $\\av$.  However, there is significant scatter in these correlations.  Because these are integrated line of sight parameters, it makes sense to divide by line-of-sight pathlength.  Both $\\ebv$ and $\\av$ are strongly correlated with $\\NHtot$, and both are thought to be rough measures of the total dust column density.  Therefore, $\\ebvdist$ and $\\avdist$ should be strongly correlated with $\\nHavg$ and be approximations of the total dust volume density.  When we look for correlations between magnesium depletion and $\\ebvdist$ and $\\avdist$ we find that the correlations are substantially increased when compared to the integrated line of sight parameters.  Therefore, we can say, with reasonable confidence, that magnesium depletion is increased in increasingly dusty environments.  The correlations with $\\ebvdist$ and $\\avdist$ are plotted in Figures \\ref{fig:logMgIIHebv_dist} and \\ref{fig:logMgIIHav_dist}.  As with the correlation between depletion and $\\fHmol$, we examine partial correlation coefficients to determine the independent significance of these correlations.  The partial correlation coefficient between Mg II depletion and $\\log{\\ebvdist}$ with $\\lognHavg$ held fixed implies that the probability of the null hypothesis is less than 6\\%.  The same partial correlation coefficient with $\\log{\\avdist}$ implies that there is less than a 1\\% chance of the null hypothesis.  (Note that both probabilities are two-sided to a $t$-distribution.)  We consider these variables in their logarithmic forms because these are the versions of the variables that exhibit the strong correlations (for all combinations of the variables in question).  However, if the variables are considered in non-logarithmic forms, the probability of the null hypothesis generally increases.  Again we note that this type of trivariate statistical measure implicitly assumes that the correlation is linear with normally distributed scatter.  As with the correlation between depletion and $\\fHmol$, we conclude that while these correlations do not improve on $\\nHavg$ as a predictor of depletions, this is limited evidence that they are significant in their own right.   We have briefly explored the possibility that the additional (or ``missing'') magnesium that is depleted from these lines of sight is found in the form of gas-phase Mg I in regions that are presumably shielded from radiation by dust.  While the STIS data do not cover Mg I absorption lines in many cases, our results indicate that gas-phase Mg I column densities are far too small to account for the order-of-magnitude increase in depletion seen in gas-phase Mg II.  Therefore, it likely that the missing gas-phase magnesium is tied up in the additional grains found in these environments, supporting the conclusion above that the correlations between Mg II depletion and the parameters $\\log{\\ebvdist}$ and $\\log{\\avdist}$ are physically significant.  It is also worth noting that in no case do we see Mg/H less than $\\sim1\\ppm$, even in the densest environments, with values of $\\ebvdist$ and $\\avdist$ several times larger than the average of the sample.  This suggests that these lines of sight are not probing what might be considered ``translucent clouds'' \\citep[though some may be ``translucent lines of sight''; see][]{SnowMcCall}.   The Mg II abundance is plotted against the ratio of total visual to selective extinction, $\\rv \\equiv \\av / \\ebv$, in Figure \\ref{fig:logMgIIHRv}.  The plot shows significant scatter; however, some statistical measures show the possibility of a slight correlation.  A Pearson correlation coefficient between $\\logMgIIH$ and $\\rv$ implies that the probability of the null hypothesis is about 31\\%.  A Spearman's $\\rho$ rank correlation coefficient, which does not depend on the functional form assumed (including whether variables are considered linearly or logarithmically) beyond assuming that the correlation is either monotonically increasing or monotonically decreasing, shows a negative correlation (decreasing abundance/increasing depletion as $\\rv$ increases) and is significant to approximately 1.3-$\\sigma$.  Therefore, there is evidence of a possible slight correlation between depletion and $\\rv$.  However, this is far from a certain conclusion.  Significant selection effects are also a possibility, as the correlations are dominated by some of the points that have greater depletion and large values of $\\rv$.  In fact, if these lines of sight are excluded, the trend begins to reverse toward a positive correlation between increasing Mg II abundance and increasing $\\rv$.  In general, we conclude that $\\rv$ is a poor predictor of depletions; as one anecdotal counterexample, HD 91597 has a very large value of $\\rv=4.9$ but does not exhibit particularly large Mg II depletion.  However, there are a few lines of sight that present interesting interpretive challenges where the value of $\\rv$ may provide insight.  As discussed above, the possibility of large magnesium depletion but small $\\fHmol$ exists for three lines of sight:  HD 37021, HD 37061, and HD 147888.  In the latter case the effect is clear, while in the former two cases the small value of $\\fHmol$ is merely inferred.  What is interesting, as noted above, is that these three lines of sight all have values of $\\rv>4$ which is a fairly significant deviation from the interstellar average of 3.1.  Because large grains contribute to $\\av$ (i.e.~grey extinction) but less so to $\\ebv$, $\\rv$ is thought to be correlated to average grain size.  Explaining why depletion should increase in a line of sight with large grains is difficult.  As mentioned above in our discussion of $\\Hmol$ and iron depletion, increased grain size decreases dust surface area per unit volume, and therefore reduces $\\Hmol$ formation rates.  However, decreased surface area per unit volume also implies a reduction in rates of sticking between dust grains and gas-phase atoms.  It seems we can reasonably conclude that the large values of $\\rv$, i.e.~the larger average grain populations, are not responsible for the large depletions by way of atoms and ions sticking to the grains.  One possibility is that the large depletions are instead ``locked in'' prior to grain coagulation.  Another possibility is the effect of a high-radiation field:  this is known for the line of sight toward HD 147888 ($\\rho$ Oph D) as well as HD 37021 and HD 37061 which are in Orion (radiation is presumed to be responsible for the relative lack of Cl I, and thus also $\\Hmol$, as mentioned in \\S \\ref{ss:hydrogen}).  However, \\citet{Snow1983} argues that the increased radiation is unlikely to be entirely responsible for the low $\\fHmol$ in the $\\rho$ Oph cloud.  Whether or not this is the case for HD 37021 and HD 37061 is unclear.  More details about the radiation field and the exact nature of the grain population (we have only considered the crude measure of $\\rv$) are probably necessary to fully understand these lines of sight.     \\subsubsection{Anticorrelation with Distance} \\label{sss:distance_anticorrelation} We find that magnesium depletion is generally anticorrelated with distance to the background star, that is, line of sight pathlength; this relationship is shown in Figure \\ref{fig:logMgIIHdist}.  Depletion decreases by nearly an order of magnitude between very short lines of sight and those up to about 2 kpc or so, and then is relatively constant (to within about 0.3-0.4$\\dex$) out to about 6 kpc.  As we concluded for a similar anticorrelation seen between iron depletion and distance \\citep{JensenFeII}, the long pathlengths are likely sampling a variety of cloud conditions, resulting in the constant depletion for long-pathlength lines of sight.  On the other hand, given the comparable hydrogen column densities of all the lines of sight in this study ($\\logHtot \\approx21-22$), the shorter lines of sight are generally the denser lines of sight.  \\subsubsection{Spatial Variations} \\label{sss:spatial_variations} We find one very interesting correlation with Galactic location:  the five stars with the largest depletions reside at higher Galactic latitudes of $|b|>15^{\\circ}$.  However, with pathlengths of less than 1 kpc, these lines of sight are still primarily in the Galactic disk.  When we analyze magnesium depletions against the height from the center of the Galactic disk, $z=r \\sin{b}$, we do not see a strong correlation.  The variation with respect to Galactic latitude is most likely a coincidence, given that these are some of the densest and most reddened lines of sight.  We do not see any other evidence of significant spatial variations.   \\subsection{Mg/H of Galactic Stars} \\label{ss:stellarMgH} In the last several years, three major papers have attempted to analyze the cosmic abundance ``standards'' in the ISM through studies of stellar abundances and meteoritic abundances---\\citet{SnowWitt}, \\citet{SofiaMeyer}, and \\citet{Lodders}.  The importance of these standards is to compare them with the observed gas-phase abundances and infer an absolute value for depletions---and therefore absolute values for the amount of these elements in phases other than atomic gas, i.e.~dust grains and molecules.  Of the major elements relevant to dust, the element with the best determined cosmic abundance is iron.  The four major measurements of the cosmic Fe/H ratio---solar, B stars, F and G stars, and CI chondrites---all agree very closely, largely within the errors.  The situation is somewhat more complex for other elements.  Carbon and oxygen show apparent overabundances in the Sun compared to F and G stars (whether or not the solar and F/G star abundances potentially agree within the errors depends on the choice of solar abundances, regarding which there is still some uncertainty), while B stars show relative deficits in these abundances compared to the Sun and other F and G stars.  The chondritic abundances of C and O are even smaller.  Nitrogen seems to be somewhat less abundant in B stars than in the Sun, though the two are reconciliable within the errors; F and G nitrogen abundances are generally unknown, and the chondritic abundances are substantially lower.  Silicon seems to be most abundant in F and G stars, slightly less abundant in the Sun, and about half as abundant in B stars.  The errors, however, do not rule out agreement between all three measurements.  However, the chondritic abundance tightly matches the solar abundance.  Both \\citet{SofiaMeyer} and \\citet{Lodders}, cite \\citet{Holweger} for the solar abundance of magnesium, $\\logMgH=-4.46$; \\citet{SnowWitt} report a slightly older value from \\citet{AndersGrevesse} of $\\logMgH=-4.42$, though these values are consistent within the errors.  The chondritic abundances in \\citet{Lodders} of $\\logMgH=-4.44$ are also very consistent with these solar values.  The discrepancy arises when various stellar abundances are considered.  Both \\citet{SnowWitt} and \\citet{SofiaMeyer} found significantly smaller abundances of Mg for B stars.  \\citet{SnowWitt} found $\\logMgH=-4.63$ for field B stars and $\\logMgH=-4.68$ for cluster B stars; \\citet{SofiaMeyer}, making no distinction between cluster and field stars, found $\\logMgH=-4.64$ for all B stars.  Though there is marginal agreement within the very large errors in these numbers, the B star abundances are $\\approx60\\%$ smaller than the solar abundances.  \\citet{SnowWitt} and \\citet{SofiaMeyer} also  disagree on the Mg abundance in F and G stars ($\\logMgH=-4.52$ and $-4.37$, respectively) due most likely to \\citet{SofiaMeyer} restricting their sample to stars with ages of $\\leq2$ Gyr.  Again, however, these values have relatively large errors and are reconciliable with the B star abundances, though just barely.  What effect does the choice of a cosmic magnesium abundance have for the implied dust-phase abundances to be used in dust models?  The differences between the cosmic abundances just discussed leads to nontrivial differences in the dust-phase abundances.  Our weighted interstellar average of Mg II/H is $2.7\\pm0.1\\ppm$ (parts per million), though the median value in our 44 lines of sight is somewhat larger at $6.2\\ppm$.  Taking the extremes of the above numbers, anywhere between $\\sim20$ and $\\sim40\\ppm$ of Mg is available for creating dust.  Examining the various models compared in Table 3 of \\citet{SnowWitt}, we find that most models require much more Mg than the lower value of $\\sim20$ implied by a B star abundance standard.  That a B star abundance is less likely to represent the cosmic abundance was also found by \\citet{ZDA2004}, who had more difficulty fitting dust models to observations using the dust-phase abundances implied by assuming B star abundances as the cosmic standard.  In fact, this is true even though \\citeauthor{ZDA2004} assumed $\\approx0$ ppm of magnesium to be in the gas-phase.  If the few ppm of magnesium in the gas-phase as measured by this paper and CLMS were included, the \\citet{ZDA2004} fits would become even more strained (the best fits for B star abundances were at the limit of the error in those abundances and inferior to the fits obtained using other abundances).  Therefore, the major conclusion that we can make regarding cosmic abundances and the incorporation of magnesium into dust is to add to the evidence that B star abundances, despite B stars being younger and therefore potentially good tracers of the current ISM, are a poor cosmic standard.  Whether or not the solar or an F and G star abundance standard for Mg is a better fit is a test that is too sensitive for us to comment on, given the uncertainties in those abundances.    We have analyzed the abundance of Mg II in 44 lines of sight.  Our study does not probe substantially larger average hydrogen volume densities than previously observed by CLMS; therefore, we observe the same correlations between Mg II and the $\\nHavg$ and $\\ebvdist$.  We also note a correlations between magnesium depletion and $\\avdist$, a different measure of dust density.  Correlations between $\\NHtot$ and the reddening and extinction parameters $\\ebv$ and $\\av$ mean that correlations between Mg II depletion and dust density measures are expected.  However, these latter correlations, while not strong than the correlation between depletion and $\\nHavg$, show some evidence of being significant even at fixed $\\nHavg$ and should be more directly related to the line-of-sight dust content.  We also note a correlation between magnesium depletion and $\\fHmol$ in our data; combined with the results of CLMS, this correlation seems to be valid for $\\fHmol \\gtrsim 0.1$ but not at smaller $\\fHmol$.  A question that is related to the trend with $\\fHmol$ is why so little $\\Hmol$ forms in certain high-density lines of sight.  Our results are consistent with the \\citet{Snow1983} suggestion that the reduced grain surface area per unit volume of large grains plays a role in reducing $\\Hmol$ formation rates.  For similar reasons, we can conclude that the grain coagulation probably occurs after depletions are already ``locked into'' the dust, rather than depletion of gas-phase atoms onto grain surfaces."
    },
    "0710/0710.1913_arXiv.txt": {
        "abstract": "We report on the H$_2$O maser distributions around IRAS 22480+6002 (=IRC+60370) observed with the Japanese VLBI Network (JVN) at three epochs spanning 2 months. This object was identified as a K-type supergiant in 1970s, which was unusual as a stellar maser source. The spectrum of H$_2$O masers consists of 5 peaks separated roughly equally by a few km s$^{-1}$ each. The H$_2$O masers were spatially resolved into more than 15 features, which spread about 50 mas along the east--west direction. However, no correlation was found between the proper motion vectors and their spatial distributions; the velocity field of the envelope seems random. A statistical parallax method applied to the observed proper-motion data set gives a distance of $1.0\\pm 0.4$ kpc for this object, that is considerably smaller than previously thought. The distance indicates that this is an evolved star with $L\\sim 5800\\ L_{\\odot}$. This star shows radio, infrared, and optical characteristics quite similar to those of the population II post-AGB stars such as RV Tau variables. ",
        "introduction": "\\label{sec:introduction} H$_2$O maser emission has been observed in circumstellar envelopes of evolved stars such as O-rich Mira variables and OH/IR stars with large mass loss rates of $\\dot{M} \\geq 10^{-7}M_{\\odot}$~yr$^{-1}$ \\citep{rei81,eli92}. Most of these stars are asymptotic giant branch (AGB) stars or red supergiants both with the spectral type M, with a few exceptions for transient stars at pre-planetary nebula phase (or supposedly a few pre-main sequence stars such as Ori KL; \\cite{mor98}). For the central stars with spectral types earlier than M, UV radiation from stellar chromosphere eventually dissociates most of molecules (except CO) in the inner envelope (e.g., \\cite{wir98}). Therefore, H$_2$O (or SiO) masers are usually not expected for these stars, except for the case that the molecules in dense circumstellar clumps shield themselves from UV radiation. In fact, H$_2$O masers found in a young planetary nebula \\citep{mir01,sue07} must be such an exceptional case. OH masers have been found in yellow hypergiants with spectral types F and G (such as IRC+10420 and V1427 Aql) \\citep{gig76,ned92,hum02}. However, H$_2$O  and SiO masers have never been detected in these objects \\citep{nak03}, though thermal emission of a few other molecules have been observed in the outer circumstellar shell \\citep{cas01,tey06}.  \\begin{table*}[ht] \\caption{Status of the telescopes, data reduction, and resulting performances in the individual epochs of the JVN observations.}\\label{tab:status} \\begin{center} \\footnotesize \\begin{tabular}{lccccccc} \\hline \\hline & Epoch in & & & Reference & 1-$\\sigma$ level & Synthesized & Number of \\\\ Observation & the year & Duration & Used & velocity\\footnotemark[2] & noise & beam\\footnotemark[3] & detected \\\\ code & 2005 & (hr) & telescopes\\footnotemark[1] & (km s$^{-1}$) & (Jy beam$^{-1}$) & (mas) & features \\\\ \\hline r05084b \\dotfill & March 25 & 7.3 & MZ, IR, OG, IS, KS, NB\\footnotemark[4] & $-52.3$  & 0.22 & 1.7$\\times$1.6, $-$37$^{\\circ}$ & 20 \\\\ r05116b \\dotfill & April 26 & 7.3 & MZ, IR, OG, IS\\footnotemark[5], KS, NB & $-52.0$  & 0.15 & 3.8$\\times$2.0, $-$14$^{\\circ}$ & 17 \\\\ r05151a \\dotfill & May 31 & 8.1 & MZ, OG\\footnotemark[5], IS\\footnotemark[5], KS, NB & $-$52.6 & 0.15 & 3.2$\\times$2.8, $-$66$^{\\circ}$ & 14 \\\\ \\hline \\end{tabular} \\end{center} \\footnotemark[1] Telescopes that were effectively operated and whose recorded data were valid: MZ: the VERA 20~m telescope at Mizusawa, IR: the VERA 20~m telescope at Iriki, OG: the VERA 20~m telescope at Ogasawara Is., IS: the VERA 20~m telescope at Ishigakijima Is., KS: the NiCT 34-m telescope at Kashima, NB: the NRO 45-m telescope at Nobeyama. \\\\ \\footnotemark[2] Velocity channel used for the phase reference in data reduction. \\\\ \\footnotemark[3] The synthesized beam made in natural weight; major and minor axis lengths and position angle. \\\\ \\footnotemark[4] Ceasing operation for 2.5~hr due to strong winds and pointing correction. \\\\ \\footnotemark[5] High system temperature ($>$300~K) due to bad weather conditions. \\end{table*} \\ \\\\   The optical counterpart of IRAS 22480+6002 ($=$AFGL~2968, or IRC$+$60370) was identified as a K-type supergiant (K0Ia; \\cite{hum74}, or K4.5Ia; \\cite{faw77}). Therefore, the detections of H$_2$O and SiO masers \\citep{han98,nym98} were surprising. Though a search for OH 1612 MHz emission was negative \\citep{les92}, CO emission was detected toward this star \\citep{jos98}. From the CO $J=2$--1 line profile, the systemic stellar velocity and the expansion velocity of this star were obtained to be $V_{\\rm lsr}$=$-49.3$ km s$^{-1}$ and $V_{exp}=26.4$ km~s$^{-1}$, respectively \\citep{jos98, gro99}. They are consistent with those obtained from the H$_2$O\\ and SiO maser spectra, and the expansion velocity of the envelope of this star is typical for OH/IR stars. The radial velocity gives a kinematic distance of 5.0~kpc. It suggests a large luminosity $L_{\\ast}=$140 000~$L_{\\odot}$ of the central star \\citep{gro99}, but it is consistent with the supergiant interpretation of this object. A blue nearby star, a B5II star, is seen by about 12$''$ east of this object. Though it is cataloged as a visual binary \\citep{wor97}, a physical association of this object with the maser source is questionable because of the large velocity difference of about 40 km s$^{-1}$ \\citep{hum74}. \\citet{win94} gave a new spectral classification of M0I for IRAS 22480+6002 from the low-resolution spectrum between 6000 and 8800 A, which was significantly different from the previous type assignment of this star. For a long-period variable, optical spectral classification may vary with light variations. However, this star has not been reported as a variable star, though it is optically not very faint ($V\\sim 8.3$).  In this work, we report three-epoch VLBI observations of H$_2$O\\ masers of IRAS 22480+6002 to rectify the entangled situation associated with this object. From the spatio-kinematics of the masers, we diagnose a probable anomaly of a hot wind from the K-type star. We estimated the distance to this star using the statistical parallax method based on the proper motion data of H$_2$O masers. Our result gives a much smaller distance for this star than previously thought. The new estimation of the distance demands to reconsider various properties of this star. Based on the arguments presented in section 3, we conclude that this star is a population II post-AGB star. ",
        "conclusions": "\\begin{table*}[ht] \\caption{Parameters of the H$_2$O  maser features identified by proper motion toward IRAS 22480$+$6002.} \\label{tab:pmotions} \\begin{center} \\begin{tabular}{lrrrrrrrrrrr} \\hline \\hline Feature\\footnotemark[1] & \\multicolumn{2}{c}{Offset} & \\multicolumn{4}{c}{Proper motion\\footnotemark[2]} & \\multicolumn{2}{c}{Radial motion\\footnotemark[3]} & \\multicolumn{ 3}{c}{Peak intensity at  3 epochs} \\\\ & \\multicolumn{2}{c}{(mas)} & \\multicolumn{4}{c}{(mas yr$^{-1}$)} & \\multicolumn{2}{c}{(km s$^{-1}$)} & \\multicolumn{ 3}{c}{(Jy beam$^{-1}$)} \\\\ & \\multicolumn{2}{c}{\\hrulefill} & \\multicolumn{4}{c}{\\hrulefill} & \\multicolumn{2}{c}{\\hrulefill} & \\multicolumn{ 3}{c}{\\hrulefill} \\\\ & $\\Delta$R.A. & $\\Delta$decl. & $\\mu_{x}$ & $\\sigma(\\mu_{x})$ & $\\mu_{y}$ & $\\sigma(\\mu_{y})$ & $V_{z}$ & $\\Delta V_{z}$\\footnotemark[4] & Epoch  1& Epoch  2& Epoch  3 \\\\ \\hline 1   \\ \\dotfill \\  &$    -6.56$&$   -10.69$&$   4.20$&   1.46 &$   5.86$&   1.40 &$ -59.78$&   0.26  &        0.19 &        0.17 &    ...        \\\\ 2   \\ \\dotfill \\  &$     0.00$&$     0.00$&$   0.00$&   0.36 &$   0.00$&   1.29 &$ -58.48$&   2.74  &        2.33 &        2.72 &        3.48   \\\\ 3   \\ \\dotfill \\  &$   -45.23$&$     7.26$&$  -0.87$&   0.42 &$   0.51$&   0.66 &$ -55.01$&   2.04  &        3.77 &        2.47 &        2.42   \\\\ 4   \\ \\dotfill \\  &$    -4.95$&$     7.17$&$   0.25$&   0.29 &$   0.47$&   1.44 &$ -54.33$&   1.05  &        0.41 &        0.53 &        0.73   \\\\ 5   \\ \\dotfill \\  &$   -26.70$&$    -2.86$&$  -1.03$&   0.97 &$  -0.72$&   0.61 &$ -52.34$&   3.37  &        8.11 &        8.73 &        8.16   \\\\ 6   \\ \\dotfill \\  &$    -0.37$&$     1.50$&$   3.96$&   1.29 &$   0.86$&   1.91 &$ -50.37$&   0.56  &        0.53 &        0.73 &        0.49   \\\\ 7   \\ \\dotfill \\  &$     0.16$&$     2.74$&$   0.05$&   0.61 &$   0.98$&   1.56 &$ -49.29$&   2.18  &        8.44 &        7.96 &        6.20   \\\\ 8   \\ \\dotfill \\  &$    -0.04$&$    13.96$&$   0.90$&   0.67 &$  -0.29$&   0.78 &$ -47.88$&   1.33  &        1.92 &        1.69 &        1.03   \\\\ 9   \\ \\dotfill \\  &$   -34.92$&$    -7.63$&$  -1.95$&   1.66 &$  -3.07$&   2.97 &$ -47.39$&   0.52  &        0.25 &        0.26 &    ...        \\\\ 10   \\ \\dotfill \\  &$   -48.23$&$    11.08$&$   2.27$&   1.18 &$   0.71$&   0.88 &$ -47.24$&   0.56  &        0.24 &        0.20 &        0.18   \\\\ 11   \\ \\dotfill \\  &$    -2.24$&$     8.18$&$   0.06$&   1.10 &$  -0.16$&   0.64 &$ -46.17$&   0.84  &        0.41 &        0.28 &        0.26   \\\\ 12   \\ \\dotfill \\  &$    -1.02$&$    13.29$&$   0.06$&   0.41 &$   0.56$&   0.74 &$ -45.37$&   1.97  &        3.34 &        3.40 &        2.44   \\\\ 13   \\ \\dotfill \\  &$     0.13$&$     7.77$&$  -1.86$&   3.19 &$   0.26$&   1.32 &$ -44.40$&   0.77  &        0.24 &        0.33 &        0.32   \\\\ \\hline \\end{tabular} \\end{center} \\noindent \\footnotemark[1] H$_2$O maser features detected toward IRAS 22480+6002. The feature is designated as IRAS 22480+6002:I2007 {\\it N}, where {\\it N} is the ordinal source number given in this column (I2007 stands for sources found by Imai et~al. and listed in 2007). \\\\ \\footnotemark[2] Relative value with respect to the motion of the position-reference maser feature: IRAS 22480+6002:I2007 {\\it 2}. \\\\ \\footnotemark[3] Relative value with respect to the local stand of rest. \\\\ \\footnotemark[4] Mean full velocity width of a maser feature at half intensity. \\end{table*}  \\subsection{Spatial distribution and proper motions of maser features}  Figure \\ref{fig:I2248_spectrum} shows cross-power spectra of the H$_2$O masers of IRAS 22480+6002. The H$_2$O maser emission spread in a velocity range of 15 km~s$^{-1}$, which is typical for Mira-type AGB stars (e.g., \\cite{tak94}). Five spectral peaks were seen in roughly equal separations of 2--3 km~s$^{-1}$; the second highest peak was near the systemic velocity ($V_{\\rm lsr}=-49.3$ km s$^{-1}$). The correlated powers of these peaks equally increased by a factor of two during 2 months in our observing run, except for the second peak for which the intensity increased only by about 20\\%. However, the peak flux densities of individual features were found not to vary much (Table 2). This fact indicates that extended emissions were partially resolved in the shortest baseline between NRO and NICT (197.4 km), but resolved-out in the longer baselines. % Correlated flux densities were estimated to be about 30--40\\% of the total-power intensities. Figure \\ref{fig:1st-epoch} shows the distribution of maser features at the first epoch. The extent of 50 mas corresponds to 50~AU ($D$/1 kpc), which is somewhat larger than those seen around Mira variables at $D=1$ kpc (e.g., \\cite{bow94}). In this figure, one of the low-velocity (blue-shifted) components ($V_{\\rm lsr}=-58.48$ km s$^{-1}$: the position reference) is located at the origin and the other low velocity components (in grey and blue colors) are located both near the eastern and western edges. Many of the higher velocity (red-shifted)  components (shown in yellow, orange, and red colors) fall at the eastern edge, but a few of them are scattered at the west side too. The overall distribution of water maser features are characterized by the elongation to the east-west direction. But, no clear correlation is found between the velocities and spatial positions. If the circumstellar envelope of the K supergiant interacts with the wind from the eastern BII star (though this is unlikely), maser spots and features could be aligned perpendicularly to the wind direction, i.e., in the north--south direction (e.g., \\cite{ima02b}). We find no such N--S alignment of the H$_2$O maser features. \\citet{mei99} noted that the MIR image of this star with the NASA 3-m telescope showed a northeast-southwest elongation, but concluded that it was likely to be an artifact caused by astigmatism.  \\begin{figure}[htb] \\begin{center} \\FigureFile(8cm,8cm){fig2.eps} \\end{center} \\caption{Distribution of H$_2$O masers on 2005 March 25. The color code indicates the radial velocity of the feature, and the size of the filled circle indicates the flux density of the feature. Note that the $-58.48$ km s$^{-1}$ reference component is located at the origin (light blue), but is almost overlapped with the systemic-velocity component ($\\sim -49$ km s$^{-1}$) shown in green.} \\label{fig:1st-epoch} \\end{figure}  We detected 14--20 H$_2$O maser features though all epochs  (the last column of table \\ref{tab:status}). Note that the H$_2$O masers were persistent in velocity and in spatial distribution during the two months; 65--90\\% of the detected maser features survived during our observing run. Therefore, we identified the same maser features at three epochs relatively easily, and measured the proper motions of the individual features during two months. Table \\ref{tab:pmotions} gives the measured proper motions. Figure \\ref{fig:PM-I2248} shows the linear fits to the relative positions of the individual maser features. The fitted proper motions look significant with the second epoch contributing little for most features, which warrants our selecting the same maser features at the different epochs. % Circumstellar H$_2$O masers can amplify the radiation of the central star (for example, see the case of U Her; \\cite{vle02}). % In the present case, we may speculate that the $-52.34$ km s$^{-1}$ feature (No. 5 in Table 2 and Figure 4), which is one of the strongest components and located near the center of the maser distribution, is such a maser amplifying the steller radiation. However, it is hasty to draw any conclusion from this observation, since we have no information on the central-star position in this scale.  \\begin{figure}[ht] \\FigureFile(8cm,14cm){fig3.eps} \\caption{Observed relative proper motions of H$_2$O  maser features in IRAS 22480$+$6002 in R.A. (a) and decl. (b) directions. The number on the left indicates the feature name designated in Table 2. The vertical bar attached to each data point indicates the position uncertainty. The least-square-fitted line is also shown.} \\label{fig:PM-I2248} \\end{figure}  \\begin{figure}[ht] \\FigureFile(8cm,6cm){fig4.eps} \\caption{Doppler velocities (colorfully displayed) and relative proper motion vectors (indicated by arrows) of H$_2$O  masers in IRAS 22480$+$6002. The displayed proper motion vector is that subtracted by a velocity bias ($\\overline{\\mu _{x}}, \\overline{\\mu _{y}})=(0.97, 0.72)$ [mas yr$^{-1}$] from the original vector to cancel out the average motions of all the features. A number added for each feature with a proper motion shows the assigned one after the designated name form ``IRAS 22480$+$6002: I2007\". The map origin is set to the location of the feature IRAS 22480$+$6002: I2007 {\\it 2}.} \\label{fig:I2248-velocity} \\end{figure}   Figure \\ref{fig:I2248-velocity} shows the proper motion vectors of the individual H$_2$O maser features. Note that the largest two proper-motion vectors at the lower left and lower right, i.e., features 1 and 9 of the $V_{\\rm lsr}=-59.78$ and $-47.39$ km s$^{-1}$ components, respectively, were determined by two-epoch detections, so that they are slightly inaccurate. The proper motions of all other features with 3-epoch detections are within a few mas per year (relative to the reference component at $V_{\\rm lsr}=-58.48$ km s$^{-1}$). We cannot find any systematic trend of motions in this diagram. For example, features 3 and 10 at the western edge move in opposite directions, and features 6 and 13 at the eastern edge also move in opposite directions.  Figure \\ref{fig:expansion} shows the RA-offsets and $\\mu_x$ plots against $V_{\\rm lsr}$. The ellipse is a plot of expected offset and proper motion from a thin spherical-shell model with a constant velocity (in the Right Ascension direction because the maser features are spread mainly in this direction). If the shell model is correct, all of the maser features should fall between these ellipses. However, the right panel does not show such a tendency. Rather, the observed points seem to distribute randomly.  The randomness of the proper motions may partially originate from the large random errors in the position measurements. % In order to check this issue, we made a Monte Carlo simulation of the 3-epoch proper motion fitting with the same positional uncertainties but without real motions (i.e., position jitters only due to the measurement errors). We obtained the mean velocity dispersions ($0.79\\pm 0.25$ mas yr$^{-1}$, $0.80\\pm 0.20$ mas yr$^{-1}$) for the 13 proper motions in R.A. and Dec. directions for the present case from the simulations; here, the number after the '$\\pm$' sign is a standard deviation of dispersions obtained in the simulations. The observed velocity dispersions (1.95 mas yr$^{-1}$, 1.93 mas yr$^{-1}$), are significantly larger than the simulated mean dispersions (more than $4 \\sigma$). Therefore, the observed proper motions are substantially real motions of masing features.  \\begin{figure}[ht] \\FigureFile(8cm,6cm){fig5.eps} \\caption{Plot of relative R.A. offset (left) and proper motion (right) against radial velocity. The filled and unfilled circles indicate the three-epoch and two-epoch detections. The large and small ellipses in both panels indicate the position and proper motion curves expected from thin spherical shell models (one-dimensional in the R.A. direction) with a constant expansion velocity of 12.5 km s$^{-1}$ and a radius of $3.5\\times 10^{14}$ cm, and with 7 km s$^{-1}$ and $8\\times 10^{13}$ cm, respectively both at a distance of 0.9 kpc. The observed points in the right panel do not fit to these ellipses. } \\label{fig:expansion} \\end{figure}  The observed random motions may originate from the intrinsic random ballistic motions of matters ejected from or infalling into the atmosphere of the central supergiant. This may be a characteristic of mass outflow of a supergiant originating from the extended atmosphere which is considerably turbulent \\citep{lev05,jos07}. The motion of 1 mas yr$^{-1}$ corresponds a transverse motion of $\\sim 4.7 (D/$kpc) km s$^{-1}$. In order to obtain the distance, we applied the statistical parallax method to the obtained maser proper motions; % for example, see \\citet{sch81}. % Assuming random motions of maser features, we obtained a velocity dispersion in radial motion (with respect to the average velocity of maser features, $V_{\\rm lsr}=-50.6$ km s$^{-1}$) to be $\\sigma_v \\simeq 5.0$ km s$^{-1}$, which is smaller than the outflow velocity estimated from CO emission.  The dispersion in the maser proper motions can be obtained by subtracting the dispersion involved in the measurements; % see equation (3) of \\citet{sch81}. % We obtain $\\sigma _{\\mu} \\sim 1.40$ mas yr$^{-1}$ and get a distance to IRAS 22480+6002 to be $D=\\sigma_{v}/\\sigma_{\\mu}= 0.76 \\ (\\pm 0.25)$ kpc. The formal uncertainty involved in the distance estimation was computed using equation (4) of \\citet{sch81}. If we exclude the largest two proper-motion features with two-epoch detections from the sample, we get the distance $1.02 \\ (\\pm 0.38) $ kpc. Later on, we adopt this distance for IRAS 22480+6000, because large motions detected by two-epoch observations are somewhat dubious. Note that this distance is derived based on the assumption that the velocity field of masers is random and isotropic, and that the proper motions appeared in maser features are real motions of gas clumps. The distance 1.0 kpc gives a radius of water maser shell approximately $3.7\\times 10^{14}$ cm, which is compatible with the radii of water maser shells of miras, but considerably smaller than those of M-supergiants \\citep{yat94,cot04}. Though the obtained distance still involve a considerable uncertainty, it excludes the possibility of a very large distance of 5 kpc (a kinematic distance).  The luminosity of this star is re-evaluated to be $5.8 \\times 10^3~L_{\\odot}$ (reestimated from \\cite{gro99}). It is considerably small for a supergiant. However, from the distance of 1.0 kpc, we can compute the absolute V magnitude of this star from 2MASS K magnitude ($K=2.8$) using $V-K=3.7$ (for K5III; \\cite{zom90}), and with reddening corrections, we get $M_V\\sim -4.4$. This value falls near the absolute magnitude of K5Ib (or M0Ib) \\citep{zom90}. Therefore, the luminosity is still in a range of supergiants.  The radial velocity of $\\sim -50$ km s$^{-1}$ is typical for young objects in the Perseus spiral arm  in the direction of this star (for example, see \\cite{sit03}). If we take into account the large uncertainties involved in the obtained distance, we cannot completely exclude the possibility that IRAS 22480+6002 belongs to the Perseus arm at $D\\sim 3.0$ kpc at $l=108^{\\circ}$ \\citep{xux06}. However,  there are  several other bright stellar maser sources with similar radial velocities in the same direction, e.g., CU Cep ($-50$ km s$^{-1}$), IRC+60374 ($-52$ km s$^{-1}$), and AFGL 2999 ($-50$ km s$^{-1}$). Luminosity distances to these stars are inferred to be smaller than 3 kpc from their high IRAS flux densities. In addition, MY Cep is an M supergiant with $V_{\\rm lsr}=-56$ km s$^{-1}$ in the star cluster NGC 7419. The distance to this cluster has been well estimated to be about 2.3 kpc from luminosities of member stars of the cluster (e.g., see \\cite{bea94,sub06}). These example indicates that the stars  with $V_{\\rm lsr}\\sim -50$ km s$^{-1}$ do not necessarily belong to the Perseus arm, but they may be located much closely. Because the radial velocity expected by the galactic rotation is only $\\sim -10$ km s$^{-1}$ at 1 kpc at $l=108^{\\circ}$, and because the radial-velocity dispersion of stellar maser sources is as small as $\\sim 25$ km s$^{-1}$ at the solar neighborhood (see Appendix 2 of \\cite{deg05}), IRAS 22480+6002 is possibly kinematically anomalous.  \\begin{table*}[ht] \\begin{center} \\caption{Comparison of the catalogued positions of IRAS 22480$+$6002.} \\label{tab:positions} \\footnotesize \\begin{tabular}{llllllll} \\hline \\hline Catalog  & Band & Assignment & epoch & R.A.(J2000) & decl..(J2000) & error & flux density \\\\ &      &            &       &  \\ h \\ m \\ s & \\ \\ ${\\circ} \\ \\ ' \\  \\ ''$ & $''$  &  or magnitude     \\\\ \\hline  IRAS   & MIR & 22480+6002        & 1983   & 22 49 59.2   & +60 17 55    &  $11''\\times 5''$ (19$^{\\circ}$) &  $F_{\\rm 12}=142$ Jy  \\\\ MSX6   & MIR & G108.4255+00.8939 & 1995   & 22 49 58.89  & +60 17 56.8  &  0.3 & $F_{\\rm C}=123$ Jy\\\\ &     &                   &        &              &              &      &             \\\\ 2MASS  & NIR & 2495897+6017567   & 1997   & 22 49 58.97  & +60 17 56.8  &  0.29   & K=2.78      \\\\ GSC1.2 & optical& 0426500695        & 1954 & 22 49 59.43  & +60 17 55.8  &  0.3 & R=12.29     \\\\ GSC2.2 & optical& N012302336407     & 1989.6 & 22 49 58.900 & +60 17 57.17 &  0.3 & B=12.29     \\\\ &     &                   &        &              &              &      &             \\\\ USNO-B1.0 & optical&   1502-0356025    & 1971.7 & 22 49 59.75  & +60 17 56.7  & (0.7, 1.0)     & R=8.87 \\\\ USNO-B1.0 & optical&   1502-0356023    & 1979.7 & 22 49 59.44  & +60 17 55.9  & (0.7, 0.2)     & R=8.73 \\\\ USNO-B1.0 & optical&   1502-0356019    & 1979.7 & 22 49 59.15  & +60 17 57.5  & (0.5, 0.7)     & B=12.63 \\\\ &     &                   &        &              &              &      &             \\\\ JVN (this work) & radio &   22480+6002     & 2005.5 & 22 49 58.876 & +60 17 56.65 &  $0.1''$ &  $F_{\\rm H_2O}\\sim 8$ Jy  \\\\ \\hline  \\hline \\end{tabular} \\end{center} \\end{table*}   \\subsection{Past optical/infrared data of IRAS 22480+6002 (=J22495897+6017568).} Though this star is relatively bright at optical wavelengths ($V\\sim 8.30$; Tyco Input catalog), the star was not recorded in major optical catalogs, for example, not in Henry Draper (HD) Catalogue, The Hipparcos and Tycho Catalogue, nor the General Catalog of Variable Stars, possibly because of  confusion by the nearby B5II star (TYC 4265-870-1; $V\\sim 10.74$), located by about 12$''$ east. This was involved in The Washington Double Star Catalog\\footnote{available at http://ad.usno.navy.mil/wds/wdstext.html.}, giving a separation of 10.9$''$ in 1901 and 12.0$''$ in 2006 with a small position angle variation (by $\\sim 7^{\\circ}$) to the B star. From this data, we obtain the proper motion of 17 mas yr$^{-1}$ to the west for IRAS 22480+6002 relative to this B star. As noted by \\citet{hum74}, this B5II star is probably not a binary counterpart because of the large radial velocity difference. The ACT Reference Catalog gave a very small proper motion of this B5II star (less than 3 mas yr$^{-1}$); though The Hipparcos and Tycho Catalogue gave a large proper motion in declination due to position uncertainty, but this was corrected in ACT catalog. We also checked the past catalogs recording the position of this star and summarized the results in table 3.\\footnote{ all the data except JVN are available in the VizieR database ({\\it http://vizier.nao.ac.jp/viz-bin/VizieR}).} The GSC 1.2 catalog (which remeasured the POSS1 plate taken in 1950s) gave a different position by about 4.2$''$, which leads a large proper motion of 84 mas yr$^{-1}$ if compared with GSC 2.0. This value is much larger than the above-mentioned proper motion computed from the Washington Double Star Catalog, though the proper motion vectors are roughly in the same direction. We believe the direct measurements of binary separation gives better values. Therefore, we adopt the proper motion of 17 mas yr$^{-1}$ for this star, and get $U_0=-71$  km s$^{-1}$ and $V_0=-29$ km s$^{-1}$ for IRAS 22480+6002. This motion is considerably peculiar for a population I disk star. It is likely that IRAS 22480+6002 belongs to one of kinematical streaming groups of stars as population II G and K giants \\citep{fam05}.  In the past, OH and H$_2$O  masers have been found in a few planetary and preplanetary nebulae \\citep{zil89,mir01,sue07}, where central stars of these objects have spectral types earlier than M. These masers are a remnant of circumstellar material which was ejected at the AGB phase of the central star. The molecules responsible for masers are eventually to be dissociated. In contrast, \\citet{fix84} found OH 1665/1667 MHz emission toward several warm stars as RV Tau variables with spectral type of G and K, but so far only one case (TW Aql, a semi-regular variable of K7III) was confirmed to be a circumstellar maser \\citep{pla91}. SiO masers, which are emitted within a few stellar radii of the central star (much closer than H$_2$O masers are emitted), were not detected in these warm objects before. An exceptional case is the RV Tau variable, R Sct, with spectral type K0Ib. This object exhibits strong SiO and weak H$_2$O masers (I. Yamamura, 2004 private communication) as well as the 4 $\\mu$m SiO first overtone bands \\citep{mat02}. The RV Tau variables are believed to be low-mass post-AGB stars \\citep{jul86} with low metal abundances (population II; \\cite{gir00}), though these are spectroscopically classified as supergiants. Their spectral types change between K and M-type with light variation \\citep{pol97}. The atmosphere of late-type supergiants are not in hydrostatic equilibrium; effective temperature increases with decreasing metalicity \\citep{lev05}. The RV Tau variables are enshrouded by dust shell, and CO emission has been detected in two of these variables \\citep{buj88}. Although  the optical counterpart of IRAS 22480+6002 is not reported to show any strong light variability (e.g., TASS; The Amateur Sky Survey\\footnote{data available at http://www.tass-survey.org/}), the optical spectroscopic classification, middle infrared properties,  and maser characteristics of IRAS 22480+6002 indicate a close similarity to the properties of the RV Tau variables."
    },
    "0710/0710.4525_arXiv.txt": {
        "abstract": "We report a sensitive search for the \\bhcn\\ emission line towards SDSS\\,J114816.64+525150.3 (hereafter:\\ J1148+5251) at $z$=6.42 with the Very Large Array (VLA).  HCN emission is a star formation indicator, tracing dense molecular hydrogen gas ($n({\\rm H_2}) \\geq 10^4\\,$cm$^{-3}$) within star-forming molecular clouds.  No emission was detected in the deep interferometer maps of J1148+5251.  We derive a limit for the HCN line luminosity of $L'_{\\rm HCN} < 3.3 \\times 10^{9}\\,$K \\kms pc$^2$, corresponding to a HCN/CO luminosity ratio of $L'_{\\rm HCN}$/$L'_{\\rm CO}$$<$0.13. This limit is consistent with a fraction of dense molecular gas in J1148+5251 within the range of nearby ultraluminous infrared galaxies (ULIRGs; median value:\\ $L'_{\\rm HCN}$/$L'_{\\rm CO}$=0.17$^{+0.05}_{-0.08}$) and HCN-detected $z$$>$2 galaxies (0.17$^{+0.09}_{-0.08}$).  The relationship between $L'_{\\rm HCN}$ and $L_{\\rm FIR}$ is considered to be a measure for the efficiency at which stars form out of dense gas.  In the nearby universe, these quantities show a linear correlation, and thus, a practically constant average ratio.  In J1148+5251, we find $L_{\\rm FIR}$/$L'_{\\rm HCN}$$>$6600. This is significantly higher than the average ratios for normal nearby spiral galaxies ($L_{\\rm FIR}$/$L'_{\\rm HCN}$=580$^{+510}_{-270}$) and ULIRGs (740$^{+505}_{-50}$), but consistent with a rising trend as indicated by other $z$$>$2 galaxies (predominantly quasars; 1525$^{+1300}_{-475}$).  It is unlikely that this rising trend can be accounted for by a contribution of active galactic nucleus (AGN) heating to $L_{\\rm FIR}$ alone, and may hint at a higher median gas density and/or elevated star-formation efficiency toward the more luminous high-redshift systems.  There is marginal evidence that the $L_{\\rm FIR}$/$L'_{\\rm HCN}$ ratio in J1148+5251 may even exceed the rising trend set by other $z$$>$2 galaxies; however, only future facilities with very large collecting areas such as the Square Kilometre Array (SKA) will offer the sensitivity required to further investigate this question. ",
        "introduction": "High redshift galaxy populations are now being detected back to 780 million years after the Big Bang (spectroscopically confirmed:\\ $z$=6.96; Iye \\etal\\ \\citeyear{iye06}), probing into the epoch of cosmic reionization (e.g., Fan \\etal\\ \\citeyear{fan06}; Hu \\& Cowie \\citeyear{hu06}). Many of these very distant galaxies show evidence for star formation activity (e.g., Taniguchi \\etal\\ \\citeyear{tan05}). Some are even found to be hyperluminous infrared galaxies (HLIRGs; Bertoldi \\etal\\ \\citeyear{ber03a}; Wang \\etal\\ \\citeyear{wan07}) with far-infrared (FIR) luminosities exceeding 10$^{13}$\\,\\lsol, suggesting vigorous star formation and/or AGN activity.  To probe the earliest stages of galaxy formation and the importance of AGN in this process, it is necessary to study the star formation characteristics of these galaxies.  A good diagnostic to examine the star-forming environments in distant HLIRGs are observations of molecular gas, the fuel for star formation. The by far brightest and most common indicator of molecular gas in galaxies is line emission from the rotational transitions of carbon monoxide (CO), which was detected in $\\sim$40 galaxies at high redshift ($z$$>$1; see Solomon \\& Vanden Bout \\citeyear{sv05} for a review). These observations have revealed molecular gas reservoirs with masses of $>$10$^{10}$\\,\\msol\\ in these galaxies, even in the highest redshift quasar known, J1148+5251 at $z$=6.42 (Walter \\etal\\ \\citeyear{wal03}, \\citeyear{wal04}; Bertoldi \\etal\\ \\citeyear{ber03b}).  Although CO is a good tracer of the total amount of molecular gas in a galaxy, due to the relatively low critical density of $n_{\\rm H_2} \\sim 10^2-10^3\\,$cm$^{-3}$ required to collisionally excite its lower $J$ transitions, it is not a reliable tracer of the dense molecular cloud cores where the actual star formation takes place.  Recent studies of nearby actively star-forming galaxies have shown that hydrogen cyanide (HCN) is a far better tracer of the dense ($n_{\\rm H_2} \\sim 10^5-10^6$\\,cm$^{-3}$) molecular gas where stars actually form (e.g.\\ Gao \\& Solomon \\citeyear{gao04a}, \\citeyear{gao04b}, hereafter:\\ GS04a, GS04b).  In the local universe it was found that the HCN luminosity ($L'_{\\rm HCN}$) scales linearly (unlike $L'_{\\rm CO}$) with the FIR luminosity ($L_{\\rm FIR}$) over 7--8 orders of magnitude, ranging from Galactic dense cores to ULIRGs (Wu et al.\\ \\citeyear{wu05}).  As $L_{\\rm FIR}$ traces the massive star formation rate (unless AGN heating is significant), this implies that HCN is also a good tracer of star formation.  HCN has now also been detected in five galaxies at $z$$>$2 (Solomon \\etal\\ \\citeyear{sol03}; Vanden Bout \\etal\\ \\citeyear{vdb04}; Carilli \\etal\\ \\citeyear{car05}, hereafter:\\ C05; Wagg \\etal\\ \\citeyear{wag05}; Gao \\etal\\ \\citeyear{gao07}, hereafter:\\ G07). Adding a number of upper limits obtained for other high-$z$ galaxies, these observations indicate that the more luminous, higher redshift systems systematically deviate from the linear $L'_{\\rm HCN}$--$L_{\\rm FIR}$ correlation found in the local universe (G07), and hint at a rising slope of the relation toward high $L_{\\rm FIR}$ and/or $z$.  To further investigate this apparent non-linear, rising trend, our aim has been to extend the range of existing HCN observations beyond redshift 6 and to higher $L_{\\rm FIR}$.  In this letter, we report sensitive VLA\\footnote{The Very Large Array is a facility of the National Radio Astronomy Observatory, operated by Associated Universities, Inc., under cooperative agreement with the National Science Foundation.} observations of \\bhcn\\ emission toward the $z$=6.42 quasar J1148+5251, the highest redshift source detected in CO.  A previous, less sensitive search for \\bhcn\\ emission in this source has yielded no detection (C05).  We use a concordance, flat $\\Lambda$CDM cosmology throughout, with $H_0$=71\\,\\kms\\,Mpc$^{-1}$, $\\Omega_{\\rm M}$=0.27, and $\\Omega_{\\Lambda}$=0.73 (Spergel \\etal\\ \\citeyear{spe06}). ",
        "conclusions": "\\subsection{Median Gas Density and Star Formation Efficiency}  Krumholz \\& Thompson (\\citeyear{kt07}) argue that $L_{\\rm FIR}$/$L'_{\\rm HCN}$ is expected to be higher for galaxies with a median (molecular) gas density $n_{\\rm med}$ close to or higher than the critical density $n_{\\rm crit}^{\\rm HCN}$ required for excitation of the observed HCN transition than for galaxies with lower $n_{\\rm med}$.  In their case, they define star formation efficiency as the fraction of the mass that is converted into stars per dynamical time of the system. Note that this is different than the star formation rate per unit total gas mass.  They argue that the non-linear relation between $L_{\\rm FIR}$ and $L'_{\\rm CO}$ (e.g., Kennicutt \\citeyear{ken98a}, \\citeyear{ken98b}; GS04b; Riechers \\etal\\ \\citeyear{rie06b}) arises due to the fact that CO traces all gas.  The star-formation rate is then dictated by the density $n$ divided by the free-fall time $\\tau_{\\rm ff}$ ($\\tau_{\\rm ff} \\propto n^{-0.5}$), giving the standard Schmidt-law: star formation rate $\\propto n^{1.5}$, or $L_{\\rm FIR}$ $\\propto (L'_{\\rm CO})$$^{1.5}$.  For molecules like HCN, which only trace the small fraction of dense gas clouds directly associated with star formation in normal galaxies, $\\tau_{\\rm ff}$ is roughly fixed by $n_{\\rm crit}$.  Hence the star formation rate shows a linear relationship with $n$, or $L_{\\rm FIR}$ $\\propto (L'_{\\rm HCN})$$^{1.0}$. However, in extreme galaxies, where $n_{\\rm med}$ in the molecular ISM approaches $n_{\\rm crit}^{\\rm HCN}$, $\\tau_{\\rm ff}$ again becomes relevant (i.e., HCN emission no longer selects just the rare, dense peaks whose density is fixed by $n_{\\rm crit}^{\\rm HCN}$, but instead traces the bulk of the ISM, whose density can vary from galaxy to galaxy, and thus the variation of $n$ and $\\tau_{\\rm ff}$ re-enter the calculation), and the relationship approaches $L_{\\rm FIR}$ $\\propto (L'_{\\rm HCN})$$^{1.5}$ (and $L'_{\\rm HCN}$ $\\propto L'_{\\rm CO}$).  Interestingly, current data show a marginal trend for a changing power-law index at the highest luminosities of the type proposed by Krumholz \\& Thompson. This change in power-law index from 1 to 1.5 would suggest that, in these extreme luminosity systems, $n_{\\rm med}$ approaches $n_{\\rm crit}^{\\rm HCN}$.  More systems at high luminosity are required to confirm this trend of changing power-law index.   \\subsection{The Role of AGN Heating for $L_{\\rm FIR}$}  Like most of the $z$$>$2 HCN-detected sources, J1148+5251 is a quasar. It has been found that, even for such strong AGN galaxies, the bulk of $L_{\\rm FIR}$ is likely dominatly heated by star formation in most cases (e.g., Carilli \\etal\\ \\citeyear{car01}; Omont \\etal\\ \\citeyear{omo01}; Beelen \\etal\\ \\citeyear{bee06}; Riechers \\etal\\ \\citeyear{rie06b}). However, based on radiative transfer models of the dust SED of J1148+5251, Li \\etal\\ (\\citeyear{li07}) argue that this source may currently undergo a `quasar phase', in which AGN heating of the hot and warm dust contributes significantly to $L_{\\rm FIR}$. If correct, this may be an alternative explanation for the elevated $L_{\\rm FIR}$/$L'_{\\rm HCN}$ in this galaxy. The (rest-frame) IR properties (tracing emission from hot dust) of J1148+5251 are similar to those of other $z$$>$6 quasars with much lower $L_{\\rm FIR}$ (tracing emission from warm dust), and even to local quasars (Jiang \\etal\\ \\citeyear{jia06}).  This supports the assumption that the hot dust in J1148+5251 is dominantly heated by the AGN; however, the lack of a correlation between $L_{\\rm IR}$ and $L_{\\rm FIR}$ in quasars indicates that the warm dust may still be dominantly heated by star formation.  Moreover, J1148+5251 follows the radio-FIR correlation for star-forming galaxies (Carilli \\etal\\ \\citeyear{car04}), which also suggests a starburst origin for the dominant fraction of $L_{\\rm FIR}$.  Furthermore, one of the $z$$>$2 HCN detections and some of the meaningful limits are submillimeter galaxies without a known luminous AGN, but are still offset from the local $L_{\\rm FIR}$/$L'_{\\rm HCN}$ relation.  It thus appears unlikely that AGN heating alone can account for the higher average $L_{\\rm FIR}$/$L'_{\\rm HCN}$ in the high-$z$ galaxy sample.  \\subsection{Implications for Future Studies}  Even when assuming the highest $L_{\\rm FIR}$/$L'_{\\rm HCN}$ of 2835 found among all HCN-detected galaxies in Table \\ref{tab-2}, the depth of our observations is sufficient to detect a galaxy with the redshift and $L_{\\rm FIR}$ of J1148+5251 ($z$=6.42) in HCN emission at a signal-to-noise ratio of $>$4.5. To first order, our lower limit thus is consistent with previous suggestions (G07) that $L_{\\rm FIR}$/$L'_{\\rm HCN}$ ratios in high redshift sources lie systematically above those for nearby galaxies.  The scatter around this trend is still significant, and will primarily be improved by increasing the number of HCN-detected galaxies at high $z$.  In addition, it will be important to improve on the main sources of error for the individual high-$z$ detections (e.g., signal-to-noise limited HCN/CO linewidth ratio, accuracy of the FIR SED fit, AGN bias of $L_{\\rm FIR}$).  The statistical and individual results, so far, would even be consistent with an even stronger increase in $L_{\\rm FIR}$/$L'_{\\rm HCN}$ toward the highest $z$ and/or $L_{\\rm FIR}$. Our study of J1148+5251 may hint at such an effect.  Clearly, it is desirable to obtain more sensitive observations of this source to further investigate this issue.  Due to its superior collecting area and high calibrational stability, the VLA is ideally suited for such a sensitive study.  Although J1148+5251 is the most CO- and FIR-luminous $z$$>$6 galaxy known, 80\\,hr of VLA observations were necessary to obtain the current limit.  In a favourable case, the \\bhcn\\ line may have a strength of about 1.5 times the current rms. To obtain a solid 5\\,$\\sigma$ detection of such a line, of order 1000\\,hr of observations with the VLA would be required.  Due to improved receivers and antenna performance, the fully operational EVLA will be by a factor of two more sensitive to spectral lines of several 100\\,\\kms\\ width (such as in J1148+5251), but will still require long integration times. Studies of dense gas at $z$$>$6 thus appear to require an order of magnitude increase in collecting area, such as offered by future facilities like the SKA phase I demonstrator (e.g., Carilli \\citeyear{car06}), which can serve as a low frequency counterpart to the Atacama Large Millimeter/submillimeter Array (ALMA)."
    },
    "0710/0710.4149_arXiv.txt": {
        "abstract": "{The Solar Tower Atmospheric Cherenkov Effect Experiment (STACEE) is an atmospheric Cherenkov telescope (ACT) that uses a large mirror array to achieve a relatively low energy threshold.  For sources with Crab-like spectra, at high elevations, the detector response peaks near 100 GeV.  Gamma-ray burst (GRB) observations have been a high priority for the STACEE collaboration since the inception of the experiment.  We present the results of 20 GRB follow-up observations at times ranging from 3 minutes to 15 hours after the burst triggers.  Where redshift measurements are available, we place constraints on the intrinsic high-energy spectra of the bursts.}     \\begin{document} ",
        "introduction": "The Solar Tower Atmospheric Cherenkov Effect Experiment (STACEE) is a showerfront-sampling Cherenkov telescope sensitive to gamma rays above 100 GeV. It is located at the National Solar Thermal Test Facility (NSTTF) at Sandia National Laboratories outside Albuquerque, New Mexico, USA. The NSTTF is located at 34.96$^{\\circ}$N, 106.51$^{\\circ}$W and is 1700 m above sea level. The facility has 220 heliostat mirrors designed to track the sun across the sky, each with 37 m$^{2}$ area. STACEE uses 64 of these heliostats to collect Cherenkov light produced by air showers.  STACEE employs five secondary mirrors on the solar tower to focus the Cherenkov light onto photomultiplier tube (PMT) cameras, as shown in Figure \\ref{concept}. The light from each heliostat is detected by a separate PMT and the waveform of the PMT signal is recorded by a flash ADC. A programmable digital delay and trigger system\\cite{IEEENSS2000} selects showers for acquisition while eliminating most random coincidences of night sky background photons. Under good observing conditions, STACEE operates with a threshold around 5 photoelectrons per channel. A detailed description of the instrument can be found in D.M. Gingrich et al.\\cite{Gingrich05}.  \\begin{figure}[t] \\begin{minipage}{0.5\\textwidth} \\begin{center} \\includegraphics[width=.9\\textwidth]{concept5.eps} \\end{center} \\end{minipage} \\hfill \\begin{minipage}{0.45\\textwidth} \\caption{Conceptual drawing of STACEE.} \\label{concept} \\vskip 1in \\end{minipage} \\end{figure}  The large mirror area of the STACEE detector leads to an energy threshold lower than those attainable by most single-dish imaging telescopes or water-Cherenkov telescopes. The energy threshold - defined as the energy at which the trigger rate peaks - is determined by the spectrum of the source and the effective area of the detector at the target position. For targets above 60$^{\\circ}$ in elevation with power-law spectral indices between 2 and 3, the energy threshold is typically between 150 and 200 GeV. For targets near zenith, STACEE has significant effective area at energies as low as 50 GeV.  A low energy threshold opens up the possibility of detecting more distant sources\\cite{Primack05}. Collisions of gamma rays with extragalactic background light (EBL) photons produce electron-positron pairs, attenuating the gamma-ray flux from distant sources. The extinction becomes more severe with increasing energy, producing an energy-dependent horizon for gamma-ray observations.  Thus, a low energy threshold is advantageous when attempting to characterize the high-energy emission of GRBs, for which the mean measured redshift (for Swift bursts) is 2.7\\cite{Le07}.  STACEE observations are typically taken in pairs of on-source and off-source runs.  The off-source runs serve as measurements of the background event rate produced by cosmic-ray showers.  Under normal conditions, the cosmic-ray trigger rate is $\\sim$5 Hz.  Background rejection techniques have been explored and the most effective technique is described elsewhere\\cite{Kildea05,Kildea05_2,Lindner06}. After background rejection cuts, STACEE typically obtains a 5$\\sigma$ detection of the Crab with approximately 10 hours of on-source observations.  Under good observing conditions, STACEE would obtain a 5$\\sigma$ detection in 30 minutes for a source with a spectrum equal to 4.5 times that of the Crab. ",
        "conclusions": ""
    },
    "0710/0710.1841_arXiv.txt": {
        "abstract": "We report the discovery of a planet transiting a moderately bright ($V=12.00$) G star, with an orbital period of $2.788491\\pm0.000025$ days. From the transit light curve we determine that the radius of the planet is $\\rpl = 1.257\\pm0.053\\,\\rjuplong$. \\hatcurb\\ has a mass of $\\mpl = \\hatcurm\\,\\mjuplong$, similar to the average mass of previously-known transiting exoplanets, and a density of $\\rhopl = \\hatcurrho\\,\\gcmc$. We find that the center of transit is $T_{\\mathrm{c}} = \\hatcurT$ (HJD), and the total transit duration is \\hatcurdur\\,days. ",
        "introduction": "\\label{sec:intro}  To date about 20 extrasolar planets have been found which transit their parent stars and thus yield values for their mass and radius\\footnote{Extrasolar Planets Encyclopedia: http://exoplanet.eu}. Masses range from 0.07\\,\\mjup\\ \\citep[GJ436;][]{gillon07} to about 9\\,\\mjup\\ \\citep[HAT-P-2b;][]{bakos07}, and radii from 0.7\\,\\rjup\\ (GJ436) to about 1.7\\,\\rjup\\ \\citep[TRES-4;][]{mandushev07}. These data provide an opportunity to compare observations with theoretical models of planetary structure across a wide range of parameters, including those of the host star \\citep[e.g.][and references therein]{burrows07,fortney07}. Transits also yield precise determination of other physical parameters of the extrasolar planets, for instance the surface gravity. Interesting correlations between these parameters have been noted early on, such as that between masses and periods \\citep{mazeh05} or periods and surface gravities \\citep{southworth07}. Classes of these close-in planets have also been suggested, such as very hot Jupiters (VHJs; P=1--3 days) and hot Jupiters \\citep[HJs; P=3--9 days;][]{gaudi05}, or a possible dichotomy based on Safronov numbers \\citep{hansen07}. However, the small ensemble of transiting exoplanets (TEPs) does not allow robust conclusions, thus the addition of new discoveries is valuable.  Over the past year the HATNet project\\footnote{www.hatnet.hu} \\citep{bakos02,bakos04}, a wide-angle photometric survey, has announced four TEPs. In this Letter we report on the detection of a new transiting exoplanet, which we label \\hatcurb, and our determination of its parameters, such as mass, radius, density and surface gravity. ",
        "conclusions": "\\label{sec:disc}  \\hatcurb\\ is an ordinary hot Jupiter (P = 2.788 days) with slightly inflated radius (\\rpl = 1.26\\rjup) for its mass of 1.06\\mjup, orbiting a slightly metal rich solar-like star. The $\\sim$20\\% radius inflation is what current models predict for a planet with equilibrium temperature of $\\sim$1500K \\citep{burrows07,fortney07}.  \\hatcurb\\ is more massive than any of the known TEPs with similar period ($2.5\\lesssim P \\lesssim3$\\,d), such as XO-2b, WASP-1, HAT-P-3b, TRES-1, and HAT-P-4b, with the exception of TRES-2. The latter is fairly similar in mass, radius, orbital period, and stellar effective temperature.  However, \\hatcurb\\ is interesting in that it falls between Class I and II, as defined by the Safronov number and $T_{eq}$ of the planet \\citep{hansen07}. \\hatcurb\\ has a Safronov number of $0.059\\pm0.005$ , while Class I is defined as $0.070\\pm0.01$, especially at $T_{eq}\\sim 1500$K. It seems that the additional discovery and characterization of transiting planets of Jupiter and higher masses would be very helpful in order to understand these new correlations and their reality."
    },
    "0710/0710.1246_arXiv.txt": {
        "abstract": "We present an algorithm for solving the linear dispersion relation in an inhomogeneous, magnetised, relativistic plasma. The method is a generalisation of a previously reported algorithm that was limited to the homogeneous case.  The extension involves projecting the spatial dependence of the perturbations onto a set of basis functions that satisfy the boundary conditions (spectral Galerkin method).  To test this algorithm in the homogeneous case, we derive an analytical expression for the growth rate of the Weibel instability for a relativistic Maxwellian distribution and compare it with the numerical results. In the inhomogeneous case, we present solutions of the dispersion relation for the relativistic tearing mode, making no assumption about the thickness of the current sheet, and check the numerical method against the analytical expression. ",
        "introduction": "Dissipation of the energy carried by relativistic plasma outflows is important for the physics of pulsar winds and gamma-ray bursts (for recent reviews see \\cite{2007astro.ph..3116K} and \\cite{2005RvMP...76.1143P}).  In these objects, the plasma is probably composed of electrons, positrons and protons. As well as being in relativistic bulk motion with respect to the observer, the random thermal energy of the plasma may also be relativistic, i.e., comparable to the rest mass energy of the constituent particles.  In this paper we concentrate on two instabilities: these are the two-stream or Weibel instability~\\cite{1959PhRvL...2...83W} and the tearing modes in a relativistic neutral pair plasma current sheet, thought to play a role in the formation process of relativistic shocks \\cite{2007arXiv0706.3126S} and the dissipation of magnetic energy in pulsar winds \\cite{2005PPCF...47B.719K}. They are investigated by generalising and extending to the inhomogeneous case the method presented in \\cite{2007PPCF...49..297P}. Motivated primarily by the need for code verification, we have derived some analytical expressions for the linear growth rates of these instabilities.  The Weibel instability is very important in astrophysical processes because it is able to generate a magnetic field by extracting the free energy from an anisotropic momentum distribution in an unmagnetised plasma or from the kinetic drift energy. There is an extensive literature on the Weibel instability: general conditions for the existence of the relativistic Weibel instability for arbitrary distribution functions are discussed in~\\cite{2006PhPl...13b2107S}, and wave propagation in counter-streaming magnetised nonrelativistic Maxwellian plasmas are studied in~\\cite{2005PhPl...12.2901T,2006PhPl...13f2901T}.  Dispersion curves have been found in some special cases such as, for example, the fully relativistic bi-Maxwellian distribution function, (Yoon~\\cite{1989PhFlB...1.1336Y}), and the water-bag distribution function, in which case closed-form analytical expressions can be derived not only for the Weibel instability~(Yoon~\\cite{1987PhRvA..35.2718Y}), but also for the cyclotron maser and whistler instabilities (Yoon~\\cite{1987PhRvA..35.2619Y}). However, finding an analytical expression for the dispersion relation for a given equilibrium distribution function is a complicated or even impossible task.  It involves a four-dimensional integration (3D in momentum space and 1D in time) of the equilibrium distribution function which is difficult to perform in closed form.  For this reason, the water-bag distribution is the preferred profile to analyse magnetic field generation in fast ignitor scenarios, (Silva et al.~\\cite{2002PhPl....9.2458S}) and in relativistic shocks, (Wiersma and Achterberg~\\cite{2004A&A...428..365W}, Lyubarsky and Eichler~\\cite{2006ApJ...647.1250L}). The Weibel instability in a magnetised electron-positron pair plasma has been investigated by Yang et al~\\cite{1993PhFlB...5.3369Y} using two model distributions: the water bag, and one with a power-law dependence at high energy. A general covariant description has been formulated by Melrose~\\cite{1982AuJPh..35...41M} and by Schlickeiser~\\cite{2004PhPl...11g5532S}. In the present work, we focus on equilibrium configurations given by a relativistic Maxwellian distribution function, which allows one to reduce the four-dimensional integral to a simple one-dimensional integral, as we demonstrate in Section~\\ref{sec:weibel_analytical}. The growth rates are then found by solving this equation using a single numerical quadrature, and are compared to the results found using our extended algorithm in Section~\\ref{sec:weibel_numerical}.  In the inhomogeneous case, the stability properties of a nonrelativistic Harris current sheet also have a substantial literature, with notable recent studies by Daughton~\\cite{1999PhPl....6.1329D} and Silin et al.~\\cite{2002PhPl....9.1104S}.  In the relativistic case, the tearing mode instability has been investigated by Zelenyi \\& Krasnoselskikh~\\cite{1979SvA....23..460Z}, by integrating first order perturbations of the relativistic Maxwellian distribution function along approximate, straight-line particle trajectories, in the thick layer limit (in which the Larmor radius is much smaller than the thickness of the current sheet).  In Section~\\ref{sec:tearing_analytical} we lift these restrictions to present new results for the tearing mode instability in a neutral current sheet of arbitrary temperature and thickness, and compare these with the results found using the generalised algorithm in Section~\\ref{sec:tearing_numerical}. In this work, no assumption is made about the thickness of the current sheet, and the particle trajectories are found numerically in the background magnetic field.  The numerical method, which is an extension of our previous algorithm, \\cite{2007PPCF...49..297P}, that computes the linear dispersion relation of waves within a Vlasov-Maxwell description, is described in Section~\\ref{sec:algorithm}. It is based on the approach of Daughton~\\cite{1999PhPl....6.1329D} for non relativistic Maxwellians, and involves explicit time integration of particle orbits along the unperturbed trajectories. We modify and extend our former code to include inhomogeneities in the plasma equilibrium configuration. Moreover, we generalise it to a fully relativistic approach, i.e., one that allows for relativistic temperatures as well as relativistic drift speeds. ",
        "conclusions": "\\label{sec:conclusion}  We present a generalisation and extension of our previous algorithm~\\cite{2007PPCF...49..297P} to solve the linear dispersion relation for relativistic multi-component inhomogeneous and magnetised plasmas. The code is validated by comparing the results with two standard configurations: the relativistic Weibel instability in a homogeneous plasma, and the tearing mode instability in a relativistic neutral Harris sheet.  To effect the comparison, we derived useful analytical expressions, Eq.~(\\ref{eq:RelDispWeibel}) and Eq.~(\\ref{eq:Harris2}), for the dispersion relations in these configurations and solved them numerically.  We conclude that this code is a suitable tool for the study of stability properties of more general configurations of interest in gamma-ray burst and pulsar wind theories.   \\ack{This research was supported by a grant from the G.I.F., the German-Israeli Foundation for Scientific Research and Development.}"
    },
    "0710/0710.3769_arXiv.txt": {
        "abstract": "{We describe a broad class of time-dependent exact wave solutions to 6D gauged chiral supergravity with two compact dimensions. These 6D solutions are nontrivial warped generalizations of 4D pp-waves and Kundt class solutions and describe how a broad class of previously-static compactifications from 6D to 4D (sourced by two 3-branes) respond to waves moving along one of the uncompactified directions. Because our methods are generally applicable to any higher dimensional supergravity they are likely to be of use for finding the supergravity limit of time-dependent solutions in string theory. The 6D solutions are interesting in their own right, describing 6D shock waves induced by high energy particles on the branes, and as descriptions of the near-brane limit of the transient wavefront arising from a local bubble-nucleation event on one of the branes, such as might occur if a tension-changing phase transition were to occur.}  \\preprint{PI-COSMO-65}  \\begin{document} ",
        "introduction": "Understanding time-dependent dynamics is central to applications of higher-dimensional supergravity theories \\cite{HiDSugra} to cosmology and to particle physics. Inasmuch as higher-dimensional supergravities provide the low-energy limit of string theories, any understanding of time-dependence in the supergravity limit also provides a guide for the thornier issue of understanding these same issues in string theory. For these reasons there is considerable interest in finding time-dependent solutions to higher-dimensional supergravity \\cite{tdepsusy} (as well as of non-supersymmetric gravity \\cite{tdepnonsusy}, since this can also sometimes capture similar physics).  Much of this activity has focussed on 10D and 5D theories (motivated by string applications, and Randall-Sundrum \\cite{RS} constructions), although these can either be more difficult to solve or they can have features which are specific to the relative simplicity of co-dimension one spaces. Six-dimensional supergravity has more recently emerged as being a useful intermediate workshop within which to investigate phenomena which can often generalize to still higher-dimensional contexts. Interest in 6D supergravity has been further sharpened by the recognition that it can provide insights into the nature of the cosmological constant problem \\cite{SLED1}--\\cite{SLEDpheno}, by building on the observation that higher-dimensional theories can break the link between the 4D vacuum energy density and the curvature of 4D spacetime \\cite{5DSelfTune}--\\cite{6DNonSUSYSelfTunex} (see \\cite{Burgess:2007ui} for a review). There has also been considerable recent interest in 6 dimensional models more generally \\cite{Closelyrelatedworks}. Including branes is notoriously difficult due the necessity to regularize UV divergences which arise \\cite{regularizations}, however for recent work on understanding this more deeply in the context of effective field theory see \\cite{claudiaeft}.  In this paper we further the program of understanding time-dependent solutions to higher-dimensional supergravity in two ways. In \\S2, we present a method for constructing explicit exact solutions to the supergravity field equations of 6D gauged chiral supergravity, to derive a new class of exact solutions to these equations which describe a class of gravitational waves passing through spacetimes for which two dimensions are compactified in response to the stress-energy of two space-filling 3-branes. In \\S3 we discuss the applications of these solutions, including how to construct the shock wave metric corresponding to ultra-relativistic particles and BH's moving on one of the branes, and also how these solutions provide some insight into the transient part of the dynamics describing outgoing waves which would arise shortly after a phase transition on one of the source branes. ",
        "conclusions": "Six dimensional supergravity provides a fruitful laboratory for investigating the issues which underly higher-dimensional physics in general, and brane approaches to the cosmological constant problem in particular. It does so because 6D is rich enough to exhibit many of the properties of still-higher dimensions --- like moduli-stabilization through fluxes \\cite{Susha}, brane back-reaction on internal geometries, chiral fermions and Green-Schwarz anomaly cancellation \\cite{6Danomalies}. Yet it is also on the one hand simple enough to allow the development of techniques for obtaining physically interesting exact solutions, but on the other hand not so simple as to be misleading about what happens in higher dimensional (in a way which co-dimension one physics sometimes can be).  We have used these properties to explore solution-generation techniques which we believe to be applicable to a wide variety of higher-dimensional supergravities. We do so by using these techniques to construct a new class of time-dependent exact solutions to the field equations of 6D chiral gauged supergravity. These solutions describe the physics of nonlinear gravitational waves passing through compactified spacetimes for which two dimensions are self-consistently compactified in response to the presence of two space-filling branes and a bulk Maxwell flux.  We believe these methods to merit more detailed exploration."
    },
    "0710/0710.0901_arXiv.txt": {
        "abstract": "The diffuse interstellar bands (DIBs) probably arise from complex organic molecules whose strength in local galaxies correlates with neutral hydrogen column density, \\nhi, and dust reddening, \\ebmv. Since Ca{\\sc \\,ii} absorbers in quasar (QSO) spectra are posited to have high \\nhi\\ and significant \\ebmv, they represent promising sites for the detection of DIBs at cosmological distances.  Here we present the results from the first search for DIBs in 9 Ca{\\sc \\,ii}-selected absorbers at $0.07 < z_{\\rm abs} < 0.55$.  We detect the 5780\\,\\AA\\ DIB in one line of sight at $z_{\\rm abs} = 0.1556$; this is only the second QSO absorber in which a DIB has been detected. Unlike the majority of local DIB sight-lines, both QSO absorbers with detected DIBs show weak 6284\\,\\AA\\ absorption compared with the 5780\\,\\AA\\ band.  This may be indicative of different physical conditions in intermediate redshift QSO absorbers compared with local galaxies. Assuming that local relations between the 5780\\,\\AA\\ DIB strength and \\nhi\\ and \\ebmv\\ apply in QSO absorbers, DIB detections and limits can be used to derive \\nhi\\ and \\ebmv.  For the one absorber in this study with a detected DIB, we derive \\ebmv\\ = 0.23\\,mag and $\\log$\\nhi\\ $\\ge$ 20.9, consistent with previous conclusions that Ca{\\sc \\,ii} systems have high H{\\sc \\,i} column densities and significant reddening.  For the remaining 8 Ca{\\sc \\,ii}-selected absorbers with 5780\\,\\AA\\ DIB non-detections, we derive \\ebmv\\ upper limits of 0.1--0.3\\,mag. ",
        "introduction": "Damped Lyman alpha (DLA) systems are usually considered to be the class of QSO absorber with the highest neutral hydrogen column densities [\\nhi\\ $\\ge 2 \\times 10^{20}$ \\cm].  Nonetheless, the DLAs are characterised by generally low metallicities and gas-phase depletion fractions \\citep[e.g.][]{KhareP_04a,AkermanC_05a,ProchaskaJ_07a} and low reddening due to dust (\\citealt{MurphyM_04c}; \\citealt*{EllisonS_05a}).  The DLAs are also poor in molecules, as demonstrated by both their generally low fractions of H$_2$ \\citep[e.g.][]{LedouxC_03a} and the lack of a detection for any other molecular species, such as OH or CO \\citep[e.g.][]{CurranS_06a}. Although the handful of DLAs which do exhibit molecular H$_2$ absorption may be biased, e.g.  towards high metallicities \\citep{PetitjeanP_06a}, such systems can offer a novel insight into the physical conditions of the galactic interstellar medium \\citep[ISM, e.g.][]{SrianandR_05a,NoterdaemeP_07a}.  In addition to the study of H$_2$, one avenue that is just starting to be explored is how the diffuse interstellar bands \\citep[DIBs; see reviews by][]{HerbigG_95a,SarreP_06a} may be used to probe the intermediate redshift ISM.  Although lacking definitive identifications, the strength (both absolute and relative) of these broad absorption features in the Milky Way (MW) and other nearby galaxies exhibit dependencies on (and sometimes, tight correlations with) neutral gas content, dust reddening, metallicity and local radiation field \\citep[e.g.][]{HerbigG_93a,CoxN_06a,WeltyD_06a,CoxN_07a}.  Moreover, if DIBs are as strong in DLAs as they are in the MW [i.e.~for a given \\nhi], then they should be relatively easy to detect at intermediate redshifts.  The first systematic search for DIBs in DLAs has recently been carried out by Lawton et al.~(in preparation) in 7 $z < 1$ absorbers.  In only one case were DIBs detected: the 4428, 5705 and 5780\\,\\AA\\ features\\footnote{We cite all DIBs with reference to their normal air wavelengths, although their vacuum values have been used in practice in order to be consistent with our spectral wavelength calibration; see Section 2.} were all detected in the $z \\sim 0.5$ DLA towards AO 0235$+$164 \\citep{JunkkarinenV_04a,YorkB_06a}.  Lawton et al.~showed that for the 6 non-detections in their DLA sample, the strength of the 5780\\,\\AA\\ DIB [which shows one of the tightest correlations with \\nhi\\ in the MW] is often at least 3 times weaker in DLAs for a given \\nhi\\ compared with Galactic sight-lines.  The 6284\\,\\AA\\ DIB is even more under-abundant in DLAs for a given \\nhi: 4--10 times weaker than towards Galactic sight-lines.  A similar result has been found for DIBs in the Large and Small Magellanic Cloud \\citep[LMC and SMC;][]{WeltyD_06a} where the 5780\\,\\AA\\ DIB is typically 10--30 times weaker than expected from the Galactic relation.  On the other hand, the 5780\\,\\AA\\ DIB strength correlates well with \\ebmv\\ in both Galactic and Magellanic Cloud sight-lines, and the detection towards AO 0235$+$164 also lies on the same relationship \\citep{YorkB_06a}. These results hint that DIB formation/survival and high dust content are closely linked and that DIBs are therefore most likely to be detected in galaxies with high reddening.  \\citet*{WildV_06a} have recently suggested that absorbers identified via high equivalent widths (EWs) of Ca{\\sc \\,ii} may select the highest \\nhi\\ and highest \\ebmv\\ absorbers.  For example, whereas DLAs have been constrained to have \\ebmv\\ $\\le$ 0.04 \\citep{MurphyM_04c,EllisonS_05a}, \\citet{WildV_06a} find that absorbers with Ca{\\sc \\,ii} $\\lambda$3934 EWs $>$0.7\\,\\AA\\ have \\ebmv\\ values up to $\\sim$0.1 mag.  Ca{\\sc \\,ii} absorbers may therefore be promising sites for the detection of DIBs. ",
        "conclusions": "\\begin{figure} \\centerline{\\rotatebox{0}{\\resizebox{9cm}{!} {\\includegraphics{EBMV_NaI_5780.ps}}}} \\caption{\\label{ebmv_fig} Correlations of Na{\\sc \\,i} (top panel), \\nhi\\ (middle panel) and \\ebmv\\ (bottom panel) versus the log of the EW (in m\\AA) of the 5780\\,\\AA\\ DIB.  Open squares are Galactic points from \\citet{HerbigG_93a}; open triangles/diamonds are SMC/LMC points from \\citet{WeltyD_06a}, \\citet{VladiloG_87a} and \\citet{Vidal-MadjarA_87a}; crosses are other extra-galactic data points taken from \\citet{SollermanJ_05a}, \\citet{DOdoricoS_89a} and \\citet{HeckmanT_00b}; solid stars are DLAs from \\citet{YorkB_06a} and Lawton et al.~(in preparation).  The best (least squares) fit to the data in each panel is shown with a solid line and the fit solution given at the bottom of each panel.  The fit of the 5780\\,\\AA\\ DIB with \\ebmv\\ uses all available data; the fits with of 5780 with \\nhi\\ and Na{\\sc \\,i} use just the Galactic data.  In the middle and lower panels, the solid vertical line indicates the 5780\\,\\AA\\ DIB detection towards J0013$-$0024 and the dotted lines are the 3$\\sigma$ upper limits for 6 other sight-lines where we have a 5780\\,\\AA\\ limits (as given in Table \\ref{dib_table}).  } \\end{figure}  In local (e.g.~Galactic, LMC, SMC) sight-lines, the 6284\\,\\AA\\ DIB is typically 2--3 times stronger than the 5780\\,\\AA\\ DIB (e.g., York et al. 2006a and references therein).  The one exception is the unusual SMC wing sight-line towards Sk~143 where the 6284\\,\\AA\\ DIB has an EW less than half that of the 5780\\,\\AA\\ DIB \\citep{WeltyD_06a}. \\citet{YorkB_06a} also found that in the one DLA sight-line with a 5780\\,\\AA\\ band detection out of the 7 studied by Lawton et al.~(in preparation), the EW of the 6284\\,\\AA\\ feature was also constrained to be less than the EW of the 5780\\,\\AA\\ line. \\citet{YorkB_06a} suggested that these unusual line ratios could be an indication of an ISM that is more protected from the ambient UV radiation field.  In the Ca{\\sc \\,ii}-selected absorber towards J0013$-$0024, we constrain the EW of the 6284\\,\\AA\\ DIB to be at least $\\sim$20\\% weaker than the 5780\\,\\AA\\ band.  The DIB ratios in this absorber are therefore consistent with those in the DLA detection of \\citet{YorkB_06a} and the SMC wing sight-line Sk~143 but inconsistent with other local sight-lines, including starburst galaxies \\citep{HeckmanT_00b} and the Magellanic Clouds \\citep{WeltyD_06a}.  In the Galaxy, many DIBs show correlations with \\nhi\\ and $N$(Na{\\sc \\,i}) \\citep[e.g.][]{HerbigG_93a,HerbigG_95a}.  In Table \\ref{dib_table} we tabulate the EWs of the Na{\\sc \\,i} doublet for our Ca{\\sc \\,ii}-selected absorbers.  However, we do not calculate the Na{\\sc \\,i} column density because, if the lines are strong enough to be detected in our moderate resolution spectra, they are likely to be saturated.  The Galactic correlations of \\nhi\\ and $N$(Na{\\sc \\,i}) with the 5780\\,\\AA\\ DIB, which is one of the tightest of the DIB relations, is shown in Fig. \\ref{ebmv_fig}.  We also show data for the Magellanic Clouds \\citep{WeltyD_06a} and DLAs (\\citealt{YorkB_06a}; Lawton et al.~in preparation), where it can be seen that the DIBs are weak for their \\nhi\\ compared with the Galactic correlation.  As shown in Fig. \\ref{ebmv_fig}, the DIBs in extra-galactic sight-lines are also weak for their Na{\\sc \\,i} column densities.  These departures from the Galactic relations are probably due to a combination of effects including ambient radiation field, metallicity and dust-to-gas ratios \\citep{CoxN_06a}.  Assuming that the Galactic 5780--\\nhi\\ relation provides a lower limit for the H{\\sc \\,i} column density, DIB detections may be useful for constraining \\nhi\\ in the absence of \\lya\\ observations.  For example, \\citet{WildV_05a} and \\citet{WildV_06a} have argued that Ca{\\sc \\,ii} absorbers represent the high column density end of the DLA distribution.  Our detection of the 5780\\,\\AA\\ DIB in the $z_{\\rm abs} = 0.1556$ absorber towards J0013$-$0024 supports this hypothesis, and we derive $\\log$\\nhi\\ $\\ge$ 20.9 for this absorber.  Unlike correlations with \\nhi\\ and $N$(Na{\\sc \\,i}), \\citet{WeltyD_06a} have shown that the 5780\\,\\AA\\ DIB strength follows a single relationship with \\ebmv\\ in both Galactic and Magellanic Cloud sight-lines.  \\citet{YorkB_06a} found that the single DLA 5780\\,\\AA\\ DIB detection towards AO 0235$+$164 fell on the same relationship.  It is not yet clear whether the apparent universality of this correlation is driven by a tight physical connection between dust properties and DIB formation \\citep{CoxN_07a} or whether it is coincidence of different physical drivers working in different directions \\citep{CoxN_06a}.  However, if the 5780--\\ebmv\\ is applicable to QSO absorbers, we can use our DIB detection limits to constrain their reddening.  \\citet{WeltyD_06a} derive a best fit correlation between the 5780\\,\\AA\\ DIB (in m\\AA) and the \\ebmv\\ for Galactic sight-lines: log \\ebmv\\ = $-$2.70 + 1.01 log EW(5780).  We derive the best fit relation to the 5780--\\ebmv\\ data points of the Galactic plus Magellanic Cloud plus AO 0235$+$164 DLA sight-lines and find log \\ebmv\\ = $-$2.19 + 0.79 log EW(5780) (see Figure \\ref{ebmv_fig}).  The range in log \\ebmv\\ values around the best fit relation is $\\sim \\pm$ 0.4 dex.  This correlation gives a reddening for the Ca{\\sc \\,ii} absorber towards J0013$-$0024 of \\ebmv$\\,\\,\\sim0.23$\\,mag and upper limits for the other 8 Ca{\\sc \\,ii} absorbers in our sample of 0.1--0.3\\,mag.  These values provide independent estimates of reddening associated with Ca{\\sc \\,ii}-selected absorbers that do not depend directly on the choice of extinction law and can be applied for individual absorbers and not just in a statistical fashion \\citep[e.g.][]{MurphyM_04c,WildV_05a,WildV_06a}.  The Ca{\\sc \\,ii} EWs of our sample are typically $<0.7$\\,\\AA\\ (see Table \\ref{dib_table}); for this range of EWs, \\citet{WildV_06a} determine average reddenings of \\ebmv\\ = 0.02, 0.03 and 0.03\\,mag for MW, LMC and SMC extinction curves respectively."
    },
    "0710/0710.5523_arXiv.txt": {
        "abstract": "We have performed a high mass and force resolution simulation of an idealized galaxy forming from dissipational collapse of gas embedded in a spherical dark matter halo. The simulation includes star formation and effects of stellar feedback. In our simulation a stellar disk forms with a surface density profile consisting of an inner exponential breaking to a steeper outer exponential. The break forms early on and persists throughout the evolution, moving outward as more gas is able to cool and add mass to the disk. The parameters of the break are in excellent agreement with observations.  The break corresponds with a rapid drop in the star formation rate associated with a drop in the cooled gas surface density, but the outer exponential is populated by stars that were scattered outward on nearly circular orbits from the inner disk by spiral arms. The resulting profile and its associated break are therefore a consequence of the interplay between a radial star formation cutoff and redistribution of stellar mass by secular processes. A consequence of such evolution is a sharp change in the radial mean stellar age profile at the break radius. ",
        "introduction": "\\label{sec:intro}  Since the early work of \\citet{de-Vaucouleurs:1958} it has been recognized that the disks of spiral galaxies generally follow an exponential radial surface brightness profile, and various theories have explored the physical causes and consequences of this property \\citep[e.g.][]{Fall:1980lr, Lin:1987pb, Dalcanton:1997bh, Mo:1998mi, van-den-Bosch:2001aa}. However, since \\citet{van-der-Kruit:1979gb, van-der-Kruit:1987fk} it has been known that the outer regions of disks exhibit more varied behavior. This has been confirmed by an abundance of recent data \\citep[e.g.,][hereafter PT06]{Pohlen:2000ff, Pohlen:2002ec, Bland-Hawthorn:2005ys, Erwin:2005kl, Pohlen:2006lh}. In a sample of nearby late-type galaxies from the Sloan Digital Sky Survey, PT06 found that about 60\\% have an inner exponential followed by a steeper outer exponential (downward-bending), $\\sim$30\\% have the inner exponential followed by a shallower outer exponential (upward-bending), while only $\\sim10$\\% have no detectable breaks.  Therefore breaks are a common feature of disk galaxies that any complete theory of galaxy formation must be able to explain. Furthermore, the discovery of UV emission at radii well beyond the H$\\alpha$ emission cutoff \\citep{Gil-de-Paz:2005, Thilker:2005, Thilker:2007a}, the observational evidence for inside-out disk growth \\citep{Munoz-Mateos:2007}, and detections of disk breaks in the distant universe \\citep{Perez:2004, Trujillo:2005}, suggest that the outer disks provide a direct view of disk assembly.  Several theories for the formation of breaks have been investigated. \\Citet{van-der-Kruit:1987fk} proposed that angular momentum conservation in a collapsing, uniformly rotating cloud naturally gives rise to disk breaks at roughly 4.5 scale radii. \\Citet{van-den-Bosch:2001aa} suggested that breaks are due to angular momentum cutoffs of the cooled gas.  More commonly breaks have been attributed to a threshold for star formation (SF), whether due to low gas density which stabilizes the disk \\citep{Kennicutt:1989bs}, or to a lack of a cool equilibrium ISM phase \\citep{Elmegreen:1994eu, Schaye:2004kb}. Using a semi-analytic model, \\citet{Elmegreen:2006oq} demonstrate that a double-exponential profile may result from a multi-component star formation prescription.  The existence of extended UV disks \\citep[e.g][]{Thilker:2007a} and the lack of a clear correlation of H$\\alpha$ cut-offs and optical disk breaks \\citep{Pohlen:2004, Hunter:2006aa} further complicates the picture. Regardless, while a sharp SF cutoff may explain the disk truncation, it does not provide a compelling explanation for extended outer exponential components.  Alternatively, \\citet{Debattista:2006wd} demonstrated that the redistribution of angular momentum by spirals during bar formation also produces realistic breaks in collisionless $N$-body simulations.  In this Letter we present the first results from a series of high-resolution Smooth Particle Hydrodynamics (SPH) simulations of isolated galaxy formation aimed at exploring the formation and evolution of disk breaks and outskirts in a massive, high surface brightness galaxy without a strong central bar. Resulting breaks are analogous to downward-bending breaks seen in observations. The clear advantage of our approach over past attempts is that we use a fully self-consistent physical model of the system, making no a priori assumptions about the distribution of material in the disk.  Rather, we allow the disk to grow spontaneously under the effects of gravity and gas hydrodynamics, itself influenced by star formation and feedback. The $N$-body approach (at sufficiently high mass and force resolution) ensures that we capture the dynamical processes contributing to disk evolution. Furthermore, the inclusion of prescriptions for SF and feedback allows us to make observational predictions across the break region. We show that (1) the break forms rapidly ($\\lesssim$ 1Gyr) and persists throughout the evolution of the system, moving outward as the disk mass grows; (2) the break is seeded by a sharp decrease in star formation which is caused in our simulation by a rapid decrease in the surface density of cool gas; (3) the outer disk is populated by stars that have migrated, on nearly circular orbits, from the inner disk, and consequently the break is associated with a sharp change in the radial mean stellar age profile; (4) break parameters agree with current observations. ",
        "conclusions": "\\label{sec:conclusions}  We have shown that in a self-consistent model, where the stellar disk forms through gas cooling and subsequent star formation within a dark matter halo, breaks in the stellar surface density form through the combination of different effects. A rapid drop in the SFR seeds the break and secular evolution populates the outer exponential.  In our model the SFR drop is due wholly to a drop in the surface density of gas.  However a break in SFR induced by other means ({\\it e.g.} a volume density threshold or perhaps warps) would lead to similar behavior of the outer disk and stellar density break parameters. Our model properties satisfy current observational constraints, both in the statistical sense of break properties from galaxies in SDSS,  as well as the much more detailed observations of breaks in stellar populations of NGC~4244. Though our model does not account for the effects of evolution in a full cosmological context, its simplicity assures that this is the minimal degree of evolution (with no interactions or formation of a significant bar) and should therefore be rather generic provided that a disk is massive enough for strong transient spirals to form.  The model predicts that there should be an abrupt change in the radial mean stellar age profile coincident with the break, which can be tested with future observations."
    },
    "0710/0710.2135_arXiv.txt": {
        "abstract": "IC 4406 is a large (about 100$''$ $\\times$ 30$''$) southern bipolar planetary nebula, composed of two elongated lobes extending from a bright central region, where there is evidence for the presence of a large torus of gas and dust. We show new observations of this source performed with IRAC (Spitzer Space Telescope) and the Australia Telescope Compact Array. The radio maps show that the flux from the ionized gas is concentrated in the bright central region and originates in a clumpy structure previously observed in H$\\alpha$, while in the infrared images filaments and clumps can be seen in the extended nebular envelope, the central region showing toroidal emission. Modeling of the infrared emission leads to the conclusion that several dust components are present in the nebula. ",
        "introduction": "\\label{introduction} IC 4406 is a well-studied southern planetary nebula. It has been imaged with several telescopes at different wavelength ranges. Near-IR images show two H$_2$ lobes \\citep{storey},  orthogonal to the nebula's major axis and  $\\sim$25$''$ away from each other. These peaks are approximately coincident with the two blobs observed in H$\\alpha$+[\\ion{N}{2}] and [\\ion{O}{3}] \\citep{sahai}, interpreted as indicative of the presence of a dense equatorial torus of dust. The optical images show a central ionized region about 32$''$ in diameter. CO maps  show the presence of a collimated high velocity outflow in the polar direction and with  [CO]/[H$_2$]$\\approx 5\\times 10^{-6}$ and a total molecular mass in the range 0.16--3.2 M$_\\odot$ \\citep{sahai}. Hubble Space Telescope (HST) WFPC2 images in [\\ion{N}{2}], H$\\alpha$ and [\\ion{O}{3}] have revealed the existence of an intricate system of dark lane features, which led to the name of \\lq\\lq Retina Nebula\\rq\\rq~for this object \\citep{odell}. The nebula appears to be chemically homogeneous, as \\citet{corradi} found no evidence of radial variation for He, O, N, Ne, and Ar. \\citet{cox} have detected several C-rich features at mm wavelengths, such as CN, HCO$^+$, HCN and HNC, which indicate the nebula is C-rich, although a C/O ratio of 0.6 is reported by \\citet{cohen}.  IC 4406 is a relatively low electron density nebula. Values in the 400-2000 cm$^{-3}$ range have been estimated using several different optical and infrared lines, with values derived by [\\ion{S}{2}] and [\\ion{O}{3}] doublets matching around 540 cm$^{-3}$ \\citep{liu,wang}. Its central star has a \\ion{He}{2} Zanstra temperature of 96800 K \\citep{phillips} and its distance is probably around 1.6 kpc \\citep{sahai}, although some authors claim it may be overestimated \\citep{odell}.  \\citet{gruenwald} have modeled IC 4406 with a 3-D photoionization code and fit many observed line intensities assuming there is a torus around the central star. They find as a best fit a central star temperature of $8\\times10^4$ K, luminosity of 400 L$_\\odot$, torus density 1500 cm$^{-3}$ and nebular density 100 cm$^{-3}$.  In general, comparisons of IR images of planetary nebulae, which trace the molecular gas and warm dust emission, to optical  line images, which trace the ionized gas, have shown the presence of similar structures \\citep{latter}, leading to the conclusion that molecular and ionized gas spatially coexist in planetary nebulae, as well as dust grains, despite the different physical conditions these components are presumed to survive in. We have observed IC 4406 in the radio range to inspect the distribution of the ionized gas in its envelope and in the infrared to check for emission from the equatorial dust and molecular gas.  In \\S\\ref{observations} we explain how we performed our observations and reduced the data; in \\S\\ref{results} we show our results and in particular in \\S\\ref{sed} how we modeled the emission in the radio and infrared ranges; \\S\\ref{nebular} compares our model results to the nebular parameter values obtained directly from the observational data; in \\S\\ref{summary} we summarize the present work. ",
        "conclusions": "\\label{summary} We have observed IC 4406 in the cm and 3--10 micron ranges. Our radio observations have confirmed the presence of the complicated maze of lanes already observed in H$\\alpha$ in the central region of the nebula and have enabled us to calculate several nebular parameters, whose values match the classification for this target as an evolved planetary nebula, in particular its low dust to gas mass ratio and density. IRAC imaging has revealed the presence of filaments in the nebula that were not detected in previous observations.  Our IRAC measurements, combined with literature data at longer and shorter wavelengths, have enabled us to study the SED of the PN IC 4406 and reproduce it with DUSTY. This has revealed that three different dust components are needed to model the data, with temperatures ranging from 57 to 700 K. It has also been necessary to include in the model slightly larger grains than in the standard MRN composition (up to 6.5 $\\mu$m) to account for the calculated 60 $\\mu$m optical depth. The main limits of our modeled curve are the spherical geometry assumed in DUSTY and the lack of data in the mm and sub-mm ranges, which would give a constraint on the slope of the curve. As we have observed during our trials with DUSTY, the slope of the SED in the sub-mm range changes with the maximum size of the grains included in the model. Unfortunately, in this range observations are available only for a few stars so far: none for our target.  We can speculate that in such a diversified dust environment, as we find in IC 4406, further lower temperature components may exist and future high sensitivity, high angular resolution observations will give a fundamental contribution to understand the physics of circumstellar envelopes in planetary nebulae."
    },
    "0710/0710.2904_arXiv.txt": {
        "abstract": "Compared to planets around Sun--like stars, relatively little is known about the occurrence rate and orbital properties of planets around stars more massive than 1.3~\\msun. The apparent deficit of planets around massive stars is due to a strong selection bias against early--type dwarfs in Doppler--based planet searches. One method to circumvent the difficulties inherent to massive main--sequence stars is to instead observe them after they have evolved onto the subgiant branch. We show how the cooler atmospheres and slower rotation velocities of subgiants make them ideal proxies for F-- and A--type stars. We present the early results from our planet search that reveal a paucity of planets orbiting within 1~AU of stars more massive than 1.5~\\msun, and evidence of a rising trend in giant planet occurrence with stellar mass. ",
        "introduction": "A planet--bearing star can be thought of as a very bright, extremely dense remnant of a protoplanetary disk. After all, a star inherits its defining characteristics---its mass and chemical composition---from the same disk material that forms its planets. The physical characteristics of planet host stars therefore provide a crucial link between the planets we detect today and the circumstellar environments from which they formed long ago. Studying the relationships between the observed occurrence rate and orbital properties of planets as a function of stellar characteristics informs theories of planet formation, and also helps guide the target selection of future planet searches.  A wealth of recent work has demonstrated that planet occurrence is strongly correlated with chemical composition \\citep{gonzalez97, santos04}; metal--rich stars are 3 times more likely to host planetary companions compared to stars with solar abundances \\citep{fischer05b}. This finding can be understood in the context of the core accretion model. Increasing the metallicity of a star/disk system increases the surface density of solid material at the disk midplane, which in turn leads to an enhanced growth rate for protoplanetary cores \\citep{ida04b, kornet05}.  \\begin{figure}[ht!] \\plotone{mass_hist.eps} \\caption {\\footnotesize{Distribution of stellar masses for the target stars of the California and Carnegie Planet Search. The majority of the stars have masses between 0.7~\\msun\\ and 1.3~\\msun. \\label{mass_hist}}} \\end{figure}  Another factor that enhances the surface density of material in the disk midplane is its total mass. If the mass of circumstellar disks scales with the mass of the central star, then there should be an observed correlation between planet occurrence and stellar mass \\citep{laughlin04, ida05b, kennedy07}. In principle, testing this hypothesis is fairly simple: one need only measure the fraction of stars with planets over a wide range of stellar masses. However, in practice such a study is not so straight forward given the limited range of stellar masses encompassed by most planet searches.  The difficulty can be seen in Figure~\\ref{mass_hist}, which shows the distribution of stellar masses for the target stars in California and Carnegie Planet Search \\citep[CCPS; ][]{valenti05}, which is representative of most Doppler--based planet searches. Most of the stars in Figure~\\ref{mass_hist} have masses between 0.7~\\msun\\ and 1.3~\\msun. In a decidedly non--Copernican twist of nature, it turns out that stars like our Sun are ideal planet search targets. Solar--mass G and K dwarfs are slow rotators, have stable atmospheres, and are relatively bright. The fall--off toward lower stellar masses is simply because late K-- and M--type dwarfs are faint, making the acquisition of high--resolution spectra difficult without large telescope apertures \\citep{butler06b, bonfils05b, endl03}.  The sharp drop at higher stellar masses is due to a separate observational bias. Stars with spectral types earlier than F8 tend to have rotationally broadened spectral features ($V\\sin{i} > 50$~\\ks; Galland et al. 2005), have fewer spectral lines due to high surface temperatures, and display a large amount of atmospheric ``jitter.'' Stellar jitter is excess velocity scatter due to surface inhomogeneities and pulsation, which can approach 50--100~\\ms\\ for mid--F stars \\citep{saar98, wright05}. These features conspire to limit the attainable radial velocity precision for early--type dwarfs to $> 50$~\\ms, rendering exceedingly difficult the detection of all but the most massive and short--period planets.  One method to circumvent the observational limitations inherent to high--mass dwarfs is to observe these stars after they have evolved away from the main--sequence. After stars have expended their core hydrogen fuel sources their radii expand, and their atmospheres cool leading to an increase in the number of metal lines in the star's spectrum. Stars crossing the subgiant branch also shed a large amount of angular momentum through the coupling of stellar winds to rotationally generated magnetic fields \\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 them much better suited for precision Doppler surveys. ",
        "conclusions": ""
    },
    "0710/0710.4133_arXiv.txt": {
        "abstract": "If light scalar fields are present at the end of inflation, their non-equilibrium dynamics such as parametric resonance or a phase transition can produce non-Gaussian density perturbations. We show how these perturbations can be calculated using non-linear lattice field theory simulations and the separate universe approximation. In the massless preheating model, we find that some parameter values are excluded while others lead to acceptable but observable levels of non-Gaussianity. This shows that preheating can be an important factor in assessing the viability of inflationary models. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.4075_arXiv.txt": {
        "abstract": "% Motivated by the emergence of multicore architectures, and the reality that parallelism is rarely used for analysis in observational astronomy, we demonstrate how general users may employ tightly-coupled multiprocessors in scriptable research calculations while requiring no special knowledge of parallel programming. Our method rests on the observation that much of the appeal of high-level vectorized languages like IDL or MatLab stems from relatively simple internal loops over regular array structures, and that these loops are highly amenable to automatic parallelization with OpenMP. We discuss how ISIS, an open-source astrophysical analysis system embedding the \\slang\\ numerical language, was easily adapted to exploit this pattern.  Drawing from a common astrophysical problem, model fitting, we present beneficial speedups for several machine and compiler configurations.  These results complement our previous efforts with PVM, and together lead us to believe that ISIS is the only general purpose spectroscopy system in which such a range of parallelism -- from single processors on multiple machines to multiple processors on single machines -- has been demonstrated. ",
        "introduction": "As noted in Noble et al (2006), parallel computation is barely used in astronomical analysis. For example, models in XSPEC (Arnaud 1996), the de facto standard X-ray spectral analysis tool, still run serially on my dual-CPU desktop. In this situation scientists tend to either turn away from models which are expensive to compute or just accept that they will run slowly.  Analysis systems which do not embrace parallelism can process at most the workload of only 1 CPU, resulting in a dramatic $1/n$ underutilization of resources as more CPU cores are added. At the same time, however, astronomers are well versed in scripting, particularly with very high-level, array-oriented numerical packages like IDL, PDL, and \\slang, to name a few. They combine easy manipulation of mathematical structures of arbitrary dimension with most of the performance of compiled code, with the latter due largely to moving array traversals from the interpreted layer into lower-level code like this C fragment {\\small \\begin{verbatim} case SLANG_TIMES: for (n = 0; n $<$ na; n++) c[n] = a[n] * b[n]; \\end{verbatim} } \\noindent which provides vectorized multiplication in \\slang. This suggests that much of the strength and appeal of numerical scripting stems from relatively simple loops over regular structures.  Having such loops in lower-level compiled codes also makes them ripe for parallelization with \\omp\\ on shared memory multiprocessors.  Proponents contend that conceptual simplicity makes \\omp\\ more approachable than other parallel programming models, e.g. message-passing in MPI or PVM, and emphasize the added benefit of allowing single bodies of code to be used for both serial and parallel execution.  For instance, preceding the above loop with \\verb+#pragma omp parallel for+ parallelizes the \\slang\\ multiplication operator; the pragma is simply ignored by a non-conforming compiler, resulting in a sequential program. ",
        "conclusions": "We are witnessing the arrival of serious multiprocessing capability on the desktop: multicore chip designs are making it possible for general users to access many processors. At the granularity of the operating system it will be relatively easy to make use of these extra cores, say by assigning whole programs to separate CPUs.  As noted with increasing frequency of late, though, it is not as straightforward to exploit this concurrency {\\em within} individual desktop applications.  In this paper we demonstrated how we have helped our research colleagues prepare for this eventuality. We have enhanced the vectorization capabilities of \\slirp, a module generator for the \\slang\\ numerical scripting language, so that wrappers may be annotated for automatic parallelization with \\omp.  This lets \\slang\\ intrinsic functions be replaced with parallelized versions, at runtime, without modifying a single line of internal \\slang\\ source.  We have shown how \\slang\\ operators may also be parallelized with relative ease, by identifying key loops within the interpreter source, tagging them with \\omp\\ directives and recompiling.  These simple adaptations, which did not require any changes to the \\isis\\ architecture or codebase, have yielded beneficial speedups for computations actively used in astrophysical research, and allow the same numerical scripts to be used for both serial and parallel execution -- minimizing two traditional barriers to the use of parallelism by non-specialists: learning how to program for concurrency and recasting sequential algorithms in parallel form.  By transparently using \\omp\\ to effect greater multiprocessor utilization we gain the freedom to explore on the desktop more challenging problems that other researchers might avoid for their prohibitive cost of computation. The \\omp\\ support now available in GCC makes the techniques espoused here a viable option for many open source numerical packages, opening the door to wider adoption of parallel computing by general practitioners.  \\begin{figure*}[t] \\centering \\includegraphics[angle=-90,scale=0.22]{P10.5_3.eps} \\hspace*{6mm} \\includegraphics[angle=-90,scale=0.22]{P10.5_4.eps} \\caption{Aggregate speedup of the Weibull fit function due to the parallelized operators and functions discussed above.  Left: \\lint, with inflection point at 1907 elements.  Right: \\solf, with inflection point at 384 elements. } \\label{P10.5-fig-2} \\end{figure*}"
    },
    "0710/0710.1650_arXiv.txt": {
        "abstract": "We report on an investigation of the SBS 1520+530 gravitational lens system. We have used archival HST imaging, Keck spectroscopic data, and Keck adaptive-optics imaging to study the lensing galaxy and its environment. The AO imaging has allowed us to fix the lens galaxy properties with a high degree of accuracy when performing the lens modeling, and the data indicate that the lens has an elliptical morphology and perhaps a disk. The new spectroscopic data suggest that previous determinations of the lens redshift may be incorrect, and we report an updated, though inconclusive, value $z_{lens} = 0.761$. We have also spectroscopically confirmed the existence of several galaxy groups at approximately the redshift of the lens system. We create new models of the lens system that explicitly account for the environment of the lens, and we also include improved constraints on the lensing galaxy from our adaptive-optics imaging. Lens models created with these new data can be well-fit with a steeper than isothermal mass slope ($\\alpha = 2.29$, where $\\rho \\propto r^{-\\alpha}$) if $H_0$ is fixed at 72\\hunit; isothermal models require $H_0 \\sim 50$\\hunit.  The steepened profile may indicate that the lens is in a transient perturbed state caused by interactions with a nearby galaxy. ",
        "introduction": "The strong gravitational lens system SBS 1520+530 (hereafter SBS1520) was first investigated by \\citet{chavushyan}, who found that the system consists of a pair of images of a broad absorption line quasar ($z_{src} = 1.855$) separated by 1\\farcs6. The lensing galaxy was soon detected with adaptive optics (AO) imaging \\citep{crampton} and was assumed to be at the redshift of one of two absorption line systems seen in the spectra of the quasar images. \\citet{burud} attempted to deconvolve the lens spectrum from the quasar spectra and found the redshift of the lens to be $z_{lens} = 0.717$. This redshift is broadly consistent with a photometric redshift determined by \\citet{faure}. Furthermore, the lens was found to lie along the line of sight to a photometrically identified cluster of galaxies that is expected to be at approximately the same redshift as the lens \\citep{burud,faure}.  Optical monitoring campaigns have led to the measurement of a time delay between the quasar images of $\\sim 130$ days \\citep{burud,khamitov}. This time delay provides an additional constraint for determining the mass slope of the lens galaxy \\citep[e.g.,][]{rusin} and allows a value of the Hubble Constant to be determined for the system \\citep[$H_0 = 51$\\hunit assuming an isothermal mass profile;][]{burud}. Note, however, that the mass slope and, thus, the value of $H_0$ depend on the environment surrounding the lens system \\citep[e.g.,][]{dobke,auger}. An incorrect understanding of the mass distribution and environment of the lens might account for the anomalously low value of $H_0$ obtained for SBS1520 compared to other lens systems \\citep[e.g.,][]{koopmans03,york} and the WMAP \\citep[][]{spergel} and \\em{Hubble Space Telescope} (\\em{HST}) Key Project \\citep[][]{freedman} results.  In this paper we present new Keck AO and archival \\emph{HST} imaging of the lens system that indicates the lensing galaxy may have a disk component. We also present a spectroscopic investigation of the lens environment and provide a new analysis of the lensing galaxy's spectrum which results in an updated lens redshift of $z_{lens} = 0.761$. We discuss the implications of these new observational data on previous analyses performed with SBS1520 and suggest that, in spite of some complexity, this lens provides an interesting platform to investigate dark matter interactions in dense environments. ",
        "conclusions": "We have obtained deep AO imaging and optical spectroscopy of the time-delay lens SBS1520. The AO imaging has allowed us to fix the lens galaxy properties with a high degree of precision when performing the lens modeling, and the data indicate that the lens has an elliptical morphology and perhaps a disk. The new spectroscopic data suggest that previous determinations of the lens redshift may be incorrect, and the data also allow us to quantify the lensing contribution of several groups in the immediate foreground and background of the lens. Lens models created with these new data can be well-fit with a steeper than isothermal mass slope ($\\alpha = 2.29$) if $H_0$ is fixed at 72\\hunit; isothermal models require $H_0 \\sim 50$\\hunit. \\citet{dobke} found that galaxies in overdense environments might have steeper than isothermal mass slopes caused by interactions with other galaxies \\citep[also see][]{augerslacs}. This suggests an interpretation that we are observing transient steepening of the mass profile due to galaxy-galaxy interactions and indicates that other lens systems that have obtained anomalously low values of $H_0$ may lie in overdense regions and near an interacting galaxy \\citep[e.g., B1600+434;][]{koopmans1600,auger}. Alternatively, SBS1520 can be modeled in a manner consistent with an isothermal profile and $H_0 = 64$\\hunit~if if the lens is modeled by a pixelated mass distribution and jointly modeled with other lens systems \\citep{read}. These models indicate that twisting ellipticity, triaxial projection effects, or other shape degeneracies may be effecting the parametric analyses of SBS1520 \\citep{saha}.  However, there are still several ambiguities in the data that need to be resolved before making definitive claims about the profile of SBS1520, particularly in the context of the interaction-induced steepening scenario. While we have argued that the lens redshift is likely to be $z = 0.761$ and not $z = 0.717$, the data are not conclusive. Furthermore, our modeling has assumed that all of the neighbor galaxies are at the group redshift; if this is the case, the $z = 0.76$ group centroid would be pulled closer to the lens and the group would therefore provide a larger contribution to the lens model. If this is not the case, the neighboring galaxies might have a smaller impact on the lens model. This is particularly important for Galaxy M, as this is the neighboring galaxy that most affects the lens model but also has colors least like the lens galaxy compared to the other field galaxies. It is also important to verify that at least one of the neighboring galaxies is at the same redshift as the lens because this is a requirement of the interaction-driven steepening hypothesis. Finally, obtaining a dynamical estimate of the lens mass would help to further constrain models and potentially distinguish between shape degeneracies and the mass slope degeneracy."
    },
    "0710/0710.5073_arXiv.txt": {
        "abstract": "Due to the non-commutation of spatial averaging and temporal evolution, inhomogeneities and anisotropies (cosmic structures) influence the evolution of the averaged Universe via the cosmological backreaction mechanism. We study the backreaction effect as a function of averaging scale in a perturbative approach up to higher orders. We calculate the hierarchy of the critical scales, at which $10\\%$ effects show up from averaging at different orders. The dominant contribution comes from the averaged spatial curvature, observable up to scales of $\\sim 200~\\mbox{Mpc}$. The cosmic variance of the local Hubble rate is $10\\%$ $(5\\%)$ for spherical regions of radius 40 $(60)~\\mbox{Mpc}$. We compare our result to the one from Newtonian cosmology and Hubble Space Telescope Key Project data. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.5303_arXiv.txt": {
        "abstract": "We have updated predictions for high energy neutrino and antineutrino charged current cross-sections within the conventional DGLAP formalism of NLO QCD using a modern PDF fit to HERA data, which also accounts in a systematic way for PDF uncertainties deriving from both model uncertainties and from the experimental uncertainties of the input data sets. Furthermore the PDFs are determined using an improved treatment of heavy quark thresholds. A measurement of the neutrino cross-section much below these predictions would signal the need for extension of the conventional formalism as in BFKL resummation, or even gluon recombination effects as in the colour glass condensate model. ",
        "introduction": "\\label{sec:intro}  Predictions of neutrino cross-sections at high energies have sizeable uncertainties which derive largely from the measurement uncertainties on the parton distribution functions (PDFs) of the nucleon. In the framework of the quark-parton model, high energy scattering accesses very large values of $Q^2$, the invariant mass of the exchanged vector boson, and very small values of Bjorken $x$, the fraction of the momentum of the incoming nucleon taken by the struck quark. Thus when evaluating uncertainties on high energy neutrino cross-sections it is important to use the most up to date information from the experiments at HERA, which have accessed the lowest-$x$ and highest $Q^2$ scales to date. The present paper uses the formalism of the ZEUS-S global PDF fits~\\cite{Chekanov:2002pv}, updated to include {\\em all} the HERA-I data.  Conventional PDF fits use the Next-to-leading-order (NLO) Dokshitzer-Gribov-Lipatov-Altarelli-Parisi (DGLAP) formalism~\\cite{Altarelli:1977zs,Gribov:1972ri,Lipatov:1974qm,Dokshitzer:1977sg} of QCD to make predictions for deep inelastic scattering (DIS) cross-sections of leptons on hadrons. At low-$x$ where the gluon density is rising rapidly it is probably necessary to go beyond the DGLAP formalism in order to sum $\\ln(1/x)$ diagrams, as in the Balitsky-Fadin-Kuraev-Lipatov (BFKL) formalism~\\cite{Kuraev:1977fs,Balitsky:1978ic,Lipatov:1985uk} (for recent work see~\\cite{White:2006xv,Altarelli:2005ni,Ciafaloni:2006yk}). An alternative approach is to consider non-linear terms which describe gluon recombination as in the colour glass condensate model~\\cite{Iancu:2003xm} which has had considerable success in explaining RHIC data~\\cite{JalilianMarian:2005jf}. A recent suggestion is to use a structure function consistent with HERA data that saturates the Froissart unitarity bound and thus predicts a $\\ln^2 s$ dependence of the cross-section~\\cite{Berger:2007ic}. Such approaches are beyond the scope of the present paper, which is concerned with the more modest goal of estimating the uncertainties on high energy neutrino cross-sections which are compatible with the conventional NLO DGLAP formalism. The motivation is to provide an update on the neutrino cross-sections in the literature \\cite{Gandhi:1998ri} which are widely used e.g. for estimating event rates in neutrino telescopes such as Baikal \\cite{Antipin:2007zz}, ANTARES \\cite{Aslanides:1999vq} and IceCube \\cite{icecube}, cosmic ray observatories such as HiRes \\cite{Martens:2007ff} and Auger \\cite{auger}, and radio detectors such as GLUE \\cite{Gorham:2003da}, FORTE \\cite{Lehtinen:2003xv}, RICE \\cite{Kravchenko:2002mm} and ANITA \\cite{Barwick:2005hn}. As a corollary, if cross-sections much outside these limits are observed, it would be a clear signal of the need for extensions to conventional formalism. To date no unambiguous signals which require such extensions have been observed. The prospect for measuring the cross-section using very high energy cosmic neutrinos in order to distinguish between theoretical suggestions for gluon dynamics at low $x$ has been discussed by us elsewhere \\cite{Anchordoqui:2006ta}.  Previous work on estimating high energy neutrino cross-sections~\\cite{Gandhi:1998ri} used PDF sets which no longer fit modern data from HERA~\\cite{Tung:2004rw} and an {\\em ad hoc} procedure for estimating PDF uncertainties.  The present work improves on this in several respects.  Firstly, we use a recent PDF analysis which includes data from all HERA-I running~\\cite{Chekanov:2002pv}. Secondly, we take a consistent approach to PDF uncertainties --- both model uncertainties and, more importantly, the uncertainties which derive from the correlated systematic errors of the input data sets \\cite{CooperSarkar:2002yx}.  Thirdly, we use NLO rather than LO calculations throughout. Fourthly, we use a general-mass variable flavour number scheme~\\cite{Thorne:1997ga,Thorne:2006qt} to treat heavy quark thresholds. ",
        "conclusions": "\\label{sec:conc}  We have calculated the charged current neutrino cross-section at NLO in the Standard Model using the best available DIS data along with a careful estimate of the associated uncertainties.  As mentioned earlier, there are further uncertainties associated with QCD effects at very low $x$ which are not addressed in the DGLAP formalism. When $x$ is sufficiently small that $\\alpha_\\mathrm{s}\\,\\ln (1/x) \\sim 1$, it is necessary to resum these large logarithms using the BFKL formalism. Whereas such calculations at leading-log suggest an even steeper rise of the gluon structure function at low $x$ (which would imply a higher $\\nu-N$ cross-section), this rise is not so dramatic at next-to-leading-log; for a recent application of NLL BFKL resummation to deep inelastic scattering see~\\cite{White:2006yh}.  Moreover both the DGLAP and the BFKL formalisms neglect non-linear screening effects due to gluon recombination which may lead to saturation of the gluon structure function. This has been modelled in the colour dipole framework in which DIS at low $x$ is viewed as the interaction of the $q \\bar q$ dipole to which the gauge bosons fluctuate. An unified BFKL/DGLAP calculation \\cite{Kwiecinski:1998yf} supplemented by estimates of screening and nuclear shadowing effects, predicts a {\\em decrease} of the $\\nu-N$ cross-section by $20-100\\%$ at very high energies $E_\\nu \\sim 10^8-10^{12}$~GeV~\\cite{Kutak:2003bd}.  An alternative approach uses the colour glass condensate formalism~\\cite{Iancu:2003xm} and predicts a similar suppression when a dipole model~\\cite{Kharzeev:2004yx} which fits data from RHIC is used~\\cite{Henley:2005ms}. The predicted cross-section is even lower~\\cite{Henley:2005ms} if a different dipole model~\\cite{Bartels:2002cj} developed to fit the HERA data is used and the gluon distribution is assumed to decrease for $x < 10^{-5}$. Other possibilities for the behaviour of the high energy $\\nu-N$ cross-section have also been discussed~\\cite{Berger:2007ic,Jalilian-Marian:2003wf}.  Detectors for UHE cosmic neutrinos would be able to probe such new physics if they can establish deviations from the perturbative DGLAP prediction. Hence we recommend our calculated values for estimation of the baseline event rates in neutrino telescopes and for use in event generators such as ANIS \\cite{Gazizov:2004va}.  While the expected neutrino fluxes (e.g. from the sources of the observed high energy cosmic rays) are rather uncertain, experiments can in principle exploit the different dependence on the cross-section of the rate of Earth-skimming and quasi-horizontal events \\cite{Anchordoqui:2006ta}."
    },
    "0710/0710.5629_arXiv.txt": {
        "abstract": "The specifications of the Atacama Large Millimeter Array (ALMA) have placed stringent requirements on the mechanical performance of its antennas. As part of the evaluation process of the VertexRSI and Alcatel EIE Consortium (AEC) ALMA prototype antennas, measurements of the path length, thermal, and azimuth bearing performance were made under a variety of weather conditions and observing modes. The results of mechanical measurements, reported here, are compared to the antenna specifications. ",
        "introduction": "\\PARstart{T}{he} Atacama Large Millimeter Array (ALMA) for astronomical observations at millimeter and sub--millimeter wavelengths (up to the Terra--Hertz region) needs antennas of high mechanical precision and of understandable and predictable behaviour.  This behaviour must be established for structural deformations due to gravity, temperature changes, and wind loads. This means, in particular, that a high reflector surface precision, pointing and phase stability must be maintained under all motions of tracking and mapping. We present a summary of tests of the mechanical and thermal behaviour of the 12m diameter VertexRSI and AEC ALMA prototype antennas, built at the VLA site (2000m altitude), New Mexico, USA. The tests were made at several intervals between March 2003 and April 2004, and concentrated primarily on the verification of the antenna specifications, of path length variations and parameters which influence the pointing. The data were also analyzed to understand the general behaviour of the antennas. In the investigation we have paid attention to the fact that variations in the behaviour of the antennas may be predictable or sporadic. We believe that repeatable and/or predictable variations can to a large extent be considered in the pointing model, or any other correction device. The antennas were tested in stationary position, under sidereal tracking, On--the--Fly (OTF) mapping, and in Fast--Switching mode (FSW). A large amount of data was collected during commissioning and thus refer to all types of tracking, OTF, FSW, and unintended 'shaking'.  A more extended summary of these test results was reported by the Antenna Evaluation Group to the National Radio Astronomy Observatory (NRAO) and European Southern Observatory (ESO) (which forms the ALMA partnership) for selection of the ALMA production antenna(s).  An overview of these performance results was presented in \\cite{Mangum2006}.  The current paper presents a more detailed analysis of the mechanical performance measurements made of the ALMA prototype antennas. ",
        "conclusions": "The measurements indicate that the path length specifications are fulfilled on both antennas, at least during time intervals of 1/2 to 1 hour. The path length variation is primarily due to unavoidable residual thermal dilatation of the (insulated) antenna steel components, and may span $\\sim$\\,200\\,$\\mu$m within a day. The path length variation can be predicted with high precision from temperature measurements at a few positions of the steel components, either used in empirical relations or the finite element model. It is expected that identical antennas will experience similar temperature variations of the ambient air so that the differential effect may even be smaller than stated here. Wind at speeds below the specification limit (9\\,m/s), and OTF and FSW motions of the antennas, do not degrade the path length stability.  As far as possible to measure, the antennas show similar behaviour of damping of the thermal environment, \\ie\\ the ambient air temperature and the solar radiation. The BUS of the VertexRSI antenna shows a good temperature homogeneity, even under full exposure to Sun shine.  Altough the AZ bearings have a higher order azimuth dependent wobble, the effect can be considered in the pointing model with an accuracy better than 0.6 arcsecond. On the VertexRSI prototype antenna, the wobble was very stable with time."
    },
    "0710/0710.2009_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:intro} There has been much interest in recent years in cosmological applications of string theory. The availability of precision data relating to the cosmic microwave background (CMB) opens up the possibility of constraining the phenomenology of string--inspired models, in particular regarding the dynamics of the inflationary epoch. Essentially, signatures of the extreme conditions governing the behaviour of the early universe are amplified by inflationary expansion and can be accessible by modern day observations. The typical energy scales involved far exceed those obtainable in particle collider experiments. In this paper we consider a particular inflation scenario, the Dirac--Born--Infeld (DBI) model of~\\cite{Silverstein:2003hf,Alishahiha:2004eh}. In the usual formulation of this setup, a D3 brane rolls down a warped throat towards a $\\overline{\\text{D3}}$ brane where the interbrane separation is identified with the inflaton field. The brane moves relativistically, but its speed is curtailed by the warped geometry so that the potential energy dominates and inflation can occur. The observable consequences of this model depend of course upon the choice of solution for the warped throat. The original DBI proposal used a AdS--like geometry, with an artificially imposed cutoff where the throat joins the bulk geometry, which in principle is unknown although irrelevant for cosmological implications. Such a geometry is unstable, and more properly one should consider a solution of the supergravity field equations with non--trivial D--brane fluxes to stabilise the compact geometry by dynamically generating a cut--off. Not many such solutions are known analytically, but two examples are the Klebanov--Strassler~\\cite{Klebanov:2000hb} (KS) and Maldacena--Nu{\\~n}ez~\\cite{Maldacena:2000yy} (MN) results.\\\\  Several investigations of DBI inflation from the KS geometry abound in the literature~\\cite{Kecskemeti:2006cg,Shiu:2006kj,Easson:2007dh,Shandera:2006ax,Bean:2007hc}, including systematic comparisons to current cosmological data. In this paper we show that it is also possible to examine the inflationary consequences of the BGMPZ geometries introduced in~\\cite{Butti:2004pk}. These are a one--parameter family of geometries describing a deformation of the KS solution\\footnote{The BGMPZ solutions originally arose in a different context, that of the AdS / CFT correspondence, as the baryonic branch of the KS solution.}, and interpolating smoothly between the KS and MN backgrounds. Although the metric is not known fully analytically, one can solve for it numerically and thus obtain results for numerous observables. The hope is that these can be compared with experimental data for quantities such as spectral indices, which may then allow one to extract geometrical information about the throat--geometry. To this aim, we show examples of how these quantities vary as a function of the BGMPZ parameter $\\xi$ which characterises the family of solutions, at typical points in parameter space. We note that inflationary consequences of these geometries were also considered in~\\cite{Dymarsky:2005xt}, for slow-roll brane inflation rather than the DBI setup considered in this paper.\\\\  The structure of the paper is as follows. In section \\ref{sec:dbi} we summarise the salient details of DBI inflation relevant for our purposes, followed by a brief explanation of BGMPZ backgrounds. In section \\ref{sec:calc} we explain the details of the numerical procedure adopted in solving the BGMPZ equations. Example results are presented in section \\ref{sec:results}, and in section \\ref{sec:discuss} we discuss our results before concluding. ",
        "conclusions": "\\label{sec:discuss} In this paper we have explicitly demonstrated the possibility of calculating cosmological observables in the DBI inflation setup using a one parameter family of type IIB supergravity solutions that describe the geometry of a warped throat, the BGMPZ solutions of Butti et al.~\\cite{Butti:2004pk} that interpolate smoothly between the Klebanov--Strassler and the Maldacena--Nu{\\~n}ez solution. We have provided examples of cosmological parameters, namely spectral indices, that can be calculated from the underlying geometry.\\\\  The solution for the metric of the geometries in question is not possible analytically. Therefore numerical methods have been used and shown to provide an adequate representation of the metric in terms of numerical precision. Instabilities in the derivatives entering the equations have been dealt with by matching to known power series expansions of the metric functions at asymptotically small and large values of the radial coordinate.\\\\  We presented warp factors for two almost extremal ends of the family of solutions parametrised by~$\\xi$: the Klebanov--Strassler throat ($\\xi=1/2$), and a geometry close to the Maldacena--Nu{\\~ n}ez solution ($\\xi=0.167$). The qualitative behaviour of the warp factors is seen to be different. Both solutions show a flat warp factor at low values of the rescaled radial coordinate $\\phi$, corresponding to a dynamically generated cutoff, and an asymptotically AdS behaviour at large $\\phi$ values as expected. However, the warp factor as it moves away in $\\xi$ from the KS solution develops a shoulder at intermediate values, whose slope is in general different from the eventual asymptotic behaviour. This can be effectively parameterised by a function of form (\\ref{fpar2}), thus explicitly demonstrating that phenomenological fits of DBI inflation to cosmological data that assume a warp factor of this form indeed correspond to exactly realisable warped throat geometries in a known string theory compactification.\\\\  We presented examples of scalar spectral indices and the ratio of tensor to scalar modes calculated using the different geometries. The different qualitative behaviour observed in the warp factors carries through to the spectral indices, and a quantitative difference between the values of the indices from different solutions at 55 $e$--folds before the end of inflation is potentially measurable.  Constraints on non-Gaussianity and measured values of the scalar spectral index could already be used to rule out some regions of $\\xi$, when it is considered alongside the rest of the DBI parameter space. Very generically, we have found that the amount of non--Gaussianity increases as one moves away from the KS solution (all other parameters being equal).\\\\  Since the warping changes qualitatively as one varies $\\xi$, it is also conceivable that that certain constraints on the parameter space may be relaxed by including $\\xi$ as an additional parameter with respect to existing analyses such as~\\cite{Lorenz:2007ze}. This issue certainly deserves further study.  \\subsection*{Acknowledgements} We would like to thank Gary Shiu for initiating this project and for very helpful correspondence. We acknowledge informative discussions with Marieke Postma. Our research is supported by the Dutch Foundation for Fundamental Research of Matter (FOM)."
    },
    "0710/0710.3112_arXiv.txt": {
        "abstract": "We study CP-violation effects when neutrinos are present in dense matter, such as outside the proto-neutron star formed in a core-collapse supernova. Using general arguments based on the Standard Model, we confirm that there are no CP-violating effects at the tree level on the electron neutrino and anti-neutrino fluxes in a core-collapse supernova. On the other hand significant effects can be obtained for muon and tau neutrinos even at the tree level. We show that CP violating effects can be present in the supernova electron (anti)neutrino fluxes as well, if muon and tau neutrinos have different fluxes at the neutrinosphere. Such differences could arise due to physics beyond the Standard Model, such as the presence of flavor-changing interactions. ",
        "introduction": "Recent results from solar, atmospheric and reactor experiments have significantly improved our knowledge of the neutrino mass differences and of  two of the mixing angles. If the remaining mixing angle, $\\theta_{13}$, is relatively large there is a possibility that violation of CP symmetry may be observable in the neutrino sector. Currently planned and future experiments will have improved sensitivities to the value of this angle (see e.g. \\cite{Guo:2007ug,Ardellier:2006mn,Kato:2007zz}). Effects of CP-violation in accelerator neutrino oscillation experiments have been extensively investigated \\cite{Dick:1999ed,Peres:2003wd,Ishitsuka:2005qi,Barger:1980jm,Karagiorgi:2006jf,Volpe:2006in}. The discovery of a non zero Dirac delta phase might help our understanding of the observed matter-antimatter asymmetry of the universe \\cite{Pascoli:2006ie,Pascoli:2006ci,Mohapatra:2006se,Barger:2003gt}. Besides studies on terrestrial experiments with man-made sources, a few recent works have addressed CP-violation with neutrinos from astrophysical sources (see e.g. \\cite{Akhmedov:2002zj,Winter:2006ce}). The purpose of the present paper is to explore possible effects coming from the CP-violating phase in dense matter, such as that encountered in core-collapse supernovae.  Core-collapse supernovae occur following the stages of nuclear burning during stellar evolution after an iron core is formed. The iron cores formed during the evolution of massive stars are supported by the electron degeneracy pressure and hence are unstable against a collapse during which most of the matter is neutronized. Once the density exceeds the nuclear density this collapse is halted. Rebounding pressure waves break out into a shock wave near the sonic point where the density reaches the nuclear density. Evolution of this shock wave and whether it produces an explosion is a point of current investigations. However, it is observed that the newly-formed hot proto-neutron star cools by neutrino emission. Essentially the entire gravitational binding energy of eight or more solar mass star is radiated away in neutrinos. Although the initial collapse is a very orderly (i.e. low entropy) process, during the cooling stage at later times a neutrino-driven wind heats the neutron-rich material to high entropies \\cite{Woosley:1992ek,Woosley:1994ux,Takahashi:1994yz}.  Neutrino interactions play a very important role in the evolution of core-collapse supernovae and in determining the element abundance \\cite{Balantekin:2003ip}. Neutrino heating is a possible mechanism for reheating the stalled shock \\cite{Bethe:1984ux}. A good fraction of the heavier nuclei were formed in the rapid neutron capture (r-process) nucleosynthesis scenario \\cite{Burbidge:1957vc}. Core-collapse supernovae are one of the possible sites for the r-process nucleosynthesis. A key quantity for determining the r-process yields is the neutron-to-seed nucleus ratio, determined by the neutron-to-proton ratio, which is controlled by the neutrino fluxes. In addition, recent work indicates that neutrino-neutrino interactions plays a potentially very significant role in supernovae \\cite{Samuel:1993uw,Sigl:1992fn,Balantekin:2004ug,Duan:2006jv,Balantekin:2006tg,Hannestad:2006nj}.  In this paper we study CP violation aspects in the core-collapse supernova environment. We first analyze analytically and in general terms, how the neutrino propagation equations and the evolution operator are modified in matter, in presence of a non-zero Dirac delta phase. We obtain a general formula which is valid for any matter density profile\\footnote{Such findings are in agreement with what was found in Ref. \\cite{Yokomakura:2002av}.}. In particular we demonstrate that, as in vacuum, the electron (anti)neutrino fluxes are independent of the phase $\\delta$, if mu and tau neutrinos have the same fluxes at the neutrinosphere in the supernova\\footnote{A remark on this aspect was made in \\cite{Yoshida:2006sk}.}. On the other hand the electron (anti)neutrino fluxes will depend on $\\delta$, if mu and tau neutrinos have different fluxes at the neutrinosphere, at variance with what was found in \\cite{Akhmedov:2002zj}. We present numerical calculations on possible CP violation effects on the mu and tau neutrino fluxes as well as on the electron (anti-)neutrino flux and the electron fraction. The latter can only appear if physics beyond the Standard Model, such as flavor changing interactions, induces differences on the mu and tau neutrino initial total luminosities and/or temperatures. Finally we calculate the effects from the CP violating phase on the number of events in an observatory on earth.  The plan of this paper is as follows. In Section II we present the general formalism to describe the neutrino evolution in presence of the $\\delta$ phase. The formulas concerning neutrino fluxes and the electron fraction in the supernova environment are recalled in Section III. Numerical results on these quantities are presented in Section IV. Conclusions are made in Section V. ",
        "conclusions": "\\noindent In this work we have analyzed possible effects induced by the CP violating Dirac phase in a dense environment such as the core-collapse supernovae. Our major result are that in matter: i) significant effects are found on the muon and tau neutrino fluxes for a non-zero CP violating phase; ii) important effects are also found on the electron (anti)neutrino fluxes if the $\\nu_{\\mu}$ and $\\nu_{\\tau}$ neutrino fluxes differ at the neutrinosphere. On the other hand the usual assumption of ignoring the CP violating phase made in the literature is justified if contributions from physics beyond the Standard Model is small and the $\\nu_{\\mu}$ and $\\nu_{\\tau}$ fluxes are equal at emission. We have calculated the events in an observatory on earth and shown that effects at the level of $5 \\%$ are present on the number of events as a function of neutrino energy.  Recent calculations have shown that the inclusion of neutrino-neutrino interaction introduces new features in the neutrino propagation in supernovae. A detailed study of the neutrino evolution with the CP violating phase, the neutrino-neutrino interaction as well as loop induced neutrino refractive indices will be the object of further work."
    },
    "0710/0710.3748_arXiv.txt": {
        "abstract": "We used multiwavelength data (H{\\sc i}, FUV, NUV, R) to search for evidence of star formation in the intragroup medium of the Hickson Compact Group 100. We find that young star-forming regions are located in the intergalactic H{\\sc i} clouds of the compact group which extend to over 130 kpc away from the main galaxies. A tidal dwarf galaxy candidate is located in the densest region of the H{\\sc i} tail, 61 kpc from the brightest group member and its age is estimated to be only 3.3 Myr. Fifteen other intragroup H{\\sc ii} regions and TDG candidates are detected in the GALEX FUV image and within a field 10$'$$\\times$10$'$ encompassing the H{\\sc i} tail. They have ages $<$200 Myr, H{\\sc i} masses of 10$^{9.2-10.4}$ M$_{\\odot}$, 0.001$<$ SFR $<$0.01 M$_{\\odot}$ yr$^{-1}$, and stellar masses 10$^{4.3}$--10$^{6.5}$ M$_{\\odot}$. The H{\\sc i} clouds to which many of them are associated have column densities about one order of magnitude lower than  N(H{\\sc i})$\\sim$10$^{21}$ cm$^{-2}$. ",
        "introduction": "The environment of galaxies play an important role on determining their overall properties. One of the most interesting environmental effects is seen in interacting systems which contain stripped H{\\sc i} gas in the intragroup/intergalactic medium, instead of around galaxies. Three independent studies, Mendes de Oliveira et al. (2004, for the Stephan's Quintet), Mendes de Oliveira et al. (2006, for HCG 31), Ryan-Weber et al. (2004, for NGC 1533, HCG 16, ESO 149-G003) and Oosterloo et al. (2004, for the loose group around NGC 1490) have surveyed systems with H{\\sc i} intergalactic clouds and have shown that these are associated with actively star forming regions, the so-called {\\it intergalactic H{\\sc ii} regions}. These objects seem to be similar to the H{\\sc ii} regions in our Milky Way, but are located in the intragroup medium and have high metallicities (close to solar in most cases). The fate of these types of objects and their importance in galaxy formation and evolution, enrichment of the intergalactic medium and globular cluster formation is still debatable. In order to better address these issues, we embarked on a multiwavelength analysis of the environment of interacting systems and present here the results based on FUV, NUV, optical, and H{\\sc i} data of a Hickson compact group, HCG100.  The last group in Hickson's catalog of compact groups of galaxies (Hickson 1982) is at v$_{\\rm R}$ = 5336 km/s (z=0.0178, Hickson et al. 1992) and it is formed by four late-type galaxies with accordant redshifts: a bright central Sb galaxy (HCG100a), an irregular galaxy with an optical tidal tail (HCG100b), a late-type barred spiral (HCG100c) with central H$\\alpha$ emission (Vilchez \\& Iglesias-P\\'aramo 1998) and a late-type edge-on spiral (HCG100d). HCG100a, b and c show strong evidence of interaction, demonstrated not only by peculiarities in their morphologies but also in their velocity fields (Plana et al. 2003). H100b shows a tidal tail connecting with a faint galaxy not originally classified as a member of the group by Hickson. Past encounters are also confirmed by an extended H{\\sc i} tail (Verdes-Montenegro et al. 2006). Therefore, the well-advanced dynamical state of HCG100 makes it an ideal target for searching for intergalactic H{\\sc ii} regions that might have been triggered by tidal interaction.  This paper is organized as follows: \\S 2 describes the data, \\S 3 discusses the age, masses and star formation rate estimates, \\S 3.1 and \\S 3.2 compare the intragroup regions in HCG100 with other intergalactic H{\\sc ii} regions and high$-z$ UV-luminous galaxies. Finally, \\S4  summarizes the main conclusions. Throughout this paper, we use a cosmology with $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$~and $h=0.7$. Magnitudes are given in the AB-system and the adopted distance to HCG100 is 76.3 Mpc.   \\section {The data}  HCG100 was observed with the Galaxy Evolution Explorer (GALEX) mission in the far and near ultraviolet (FUV $\\lambda$$_{\\rm eff}$=1528\\AA, NUV $\\lambda$$_{\\rm eff}$=2271\\AA) as part of our Cycle 1 program (GI\\#73) to observe compact groups of galaxies. In Fig. \\ref{fuvhi_10arcmin} and Fig.\\ref{nuvhi_10arcmin} we show a cutout of a 10$'$$\\times$ 10$'$ window of the FUV and NUV images (GALEX field of view is 1$^{\\circ}$.28 and 1$^{\\circ}$.24 in FUV and NUV, respectively and pixel scale is 1.5 arcsec pixel$^{-1}$) together with the distribution of atomic hydrogen H{\\sc i} (PI Verdes-Montenegro).  We have also obtained an R-band image with the CTIO Blanco 4m telescope and a mosaic II CCD imager. Three 300s images of the group were taken, at a median seeing of 1.1$''$.  Each frame covered a 40 arcmin field on a side, at a pixel scale of 0.27$''$/pixel, but we only used a 10 arcmin field around the object, given that this was the area of interest. The data was bias subtracted and flat fielded using standard procedures (using the package mscred in IRAF\\footnote{IRAF is distributed by the National Optical Astronomy Observatories, which are operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation.}).  A flatfield was constructed from a combination of dark sky frames and twilight flats which worked well, given that the background, after flatfielding was done, showed an rms variation no greater than 1\\% over the whole field. The night was not photometric and therefore no calibration stars were observed.  Instead, we chose to use two galaxies previously observed by Rubin et al.  (1991), HCG 100c and HCG 100d, for obtaining the zero point of the photometry (see below).  \\subsection{Flux Calibration}  FUV and NUV fluxes were calculated using Morrissey et al. (2005) m$_{\\lambda}$=-2.5 log[F$_{\\lambda}$/a$_{\\lambda}$] + b$_{\\lambda}$, where a$_{FUV}$ = 1.4 $\\times$ 10$^{-15}$ erg s$^{-1}$ cm$^{-2}$ \\AA$^{-1}$, a$_{NUV}$=2.06$\\times$ 10$^{-16}$ erg s$^{-1}$ cm$^{-2}$ \\AA$^{-1}$, b$_{FUV}$=18.82 and b$_{NUV}$=20.08 for FUV and NUV, respectively. Fluxes were multiplied by the effective filter bandpass ($\\Delta$$\\lambda$$_{FUV}$=269\\AA\\, and $\\Delta$$\\lambda$$_{NUV}$=616\\AA) to give units of erg s$^{-1}$ cm$^{-2}$ and luminosities were calculated for a distance D=76.3 Mpc.  The R-band image was calibrated to match Rubin et al. (1991) magnitudes for galaxies HCG100c and HCG100d. Total integrated R magnitudes of these objects are given in their Table 3 (no errors were quoted), and these were compared with the total instrumental magnitudes we obtained for these galaxies using two methods: (1) through aperture photometry in IRAF (using the task daophot.phot) and (2) using the program SExtractor (Bertin and Arnouts et al. 1996). The zero points obtained for each of these objects, and for each method, agreed within 0.1 mag, and we used an average of the values as our final zero point. The lack of a proper calibration, does not significantly change our results as shown in \\S~3.  The H{\\sc i} data taken with the VLA (configuration D, PI Verdes-Montenegro) is not able to resolve the intragroup objects since the beamwidth is 61.0$''$$\\times$55.23$''$ (22.6~kpc $\\times$ 20.4~kpc at 76.3~Mpc) and objects have sizes $<$16$''$. Therefore, our measurements of the H{\\sc i} fluxes were centered on the FUV detections but measured within one beam. Fluxes were converted into mass by using the relation:  \\begin{equation} M_{HI}=2.356 \\times 10^5 F_{HI} D^2 \\end{equation}  where $D$ is the distance to the group in Mpc, $F_{HI}$ is the H{\\sc i} flux in Jy km/s.  Column Densities were calculated by applying the following relation \\begin{equation} N_{HI}=1.82\\times 10^{18} (\\lambda ^2 / (2.65 \\Theta ^2) F \\end{equation}  where $\\Theta$=beamwidth in arcmin$^{2}$, $\\lambda$ is the wavelength (21cm) and F is the H{\\sc i} flux density in Jy/beam km/s (Rohlfs \\& Wilson 2000).   \\subsection{Source Detection}  We used SExtractor (SE, Bertin \\& Arnouts 1996) to detect sources in the FUV and matched that catalog with the NUV and R-band catalogs. Therefore, only objects with FUV detections were included in our final catalog. In this paper we concentrate on all objects which were detected in a region of 10$'$$\\times$ 10$'$ centered on the H{\\sc i} tail (Verdes-Montenegro et al. 2006).  We used SE's Kron elliptical apertures to measure magnitudes (mag$_{-}$auto) in FUV, NUV and R-bands. Although mag$_{-}$auto is often used to measure total magnitudes in galaxy surveys (e.g. Bell et al. 2004,  de Mello et al. 2006, Zucca et al. 2006), high uncertainties might be expected at the faint magnitudes due to the assumption that the sky background has Gaussian random noise without source confusion (Brown et al. 2007). However, since we are comparing data taken with different resolutions and the fact that UV light does not necessarily peak at the same coordinate as the optical light, mag$_{-}$auto performs better than the other SExtractor choices, as long as a careful match between the different catalogs is performed. We matched the coordinates of the sources in the 3-bands within 1$''$ to 5$''$ and after visually checking each identification we chose 3$''$ as our best matched catalog with FUV, NUV and R-detections. The cross-match in the three bands resulted not only in the four most luminous members of HCG 100 (Table 1) but also  16 other objects, within the 10$'$$\\times$10$'$ field (Table 2).  Nine of these sources are within the H{\\sc i} tail: \\# 3 is the most distant object from the group (381 arcsec or 137.2 kpc from HCG100a) and is located in the southern tip of the H{\\sc i} tail, \\#4 is far from all of the brightest members of the group and located in the densest H{\\sc i} region, \\#5 and \\#6 are close to HCG100c, \\#8 is isolated and 150 arcsec (53.9 kpc) from HCG100a, \\#9 is close to HCG100a, \\#13 is a small galaxy not originally classified as a member of the HCG100 and located in the optical tidal tail of HCG100b, \\#14 and \\#15 are also close to the tidal tail. The other six objects have no H{\\sc i} detection but are still within the chosen field.  SE's magnitudes in FUV, NUV and R (mag$_{-}$auto) were corrected for foreground Galactic extinction using E(B-V)=0.081 and A$_{R}$=E(B-V)$\\times$2.634 (Schlegel et al. 1998), A$_{FUV}$=E(B-V)$\\times$8.29 and A$_{NUV}$=E(B-V)$\\times$8.18 (Seibert et al. 2005).  In Table 3 we list the H{\\sc i} mass per beam in the vicinity of nine of the intragroup objects detected, the other 7 objects are below the detection limit ($<1.7 \\times 10^{8}$). However, it is important to note that, due to the low resolution of the data, the HI masses of the objects near large galaxies are contaminated by the HI masses of the latter. Only objects \\#3, \\#4 and \\#8 are far away from bright members of the group and this contamination could be avoided.  Objects \\#13, \\#14, \\#15 besides being close to H100b are within the same beam, therefore the HI masses listed in Table 3 are not the individual masses of each object but the total mass within one beam centered in that region.  As seen in Fig.\\ref{fuvhi_10arcmin} most of the intragroup objects are located in the outskirts of the H{\\sc i} contours where column densities range from 7.5 $\\times 10^{19}$ to 5$\\times 10^{20}$ cm$^{-2}$, except for object \\#4 which is located in a peak of H{\\sc i} where NH{\\sc i} is 1.2$\\times 10^{21}$ cm$^{-2}$. Therefore, H{\\sc i} clouds to which many of them are associated have column densities about one order of magnitude lower than the N(HI)$\\sim$10$^{21}$ cm$^{-2}$, value thought to be required for triggering star formation (e.g. Skillman et al. 1988). Due to the low resolution of the HI data, all values are lower limits to the true HI column density.   \\begin{deluxetable}{cccccccc} \\tablecaption{HCG100 Properties} \\tablewidth{0pt} \\tablehead{ \\colhead{ID} & \\colhead{Morphology$^{\\rm a}$} & \\colhead{Velocity$^{\\rm b}$} & \\colhead{FUV$^{\\rm c}$} & \\colhead{NUV} & \\colhead{R$^{\\rm d}$} & \\colhead{FUV$_{\\rm corr}$$^{\\rm e}$} & \\colhead{NUV$_{\\rm corr}$} } \\startdata HCG100 a &  Sb & 5323 &17.30 & 16.78 & 12.5 & 16.12 & 16.63\\\\ HCG100 b &  Sm & 5163 &17.49 & 17.10 & 14.1 & 16.44 & 16.82\\\\ HCG100 c &  SBc& 5418 &18.46 & 17.91 & 14.7 & 17.25 & 17.79\\\\ HCG100 d &  Scd&      &19.08 & 18.65 & 15.5 & 17.99 & 18.41\\\\ \\enddata \\tablenotetext{a}{Morphology from Plana et al. 2003} \\tablenotetext{b}{Systemic Velocity in kms$^{-1}$ from Plana et al. 2003} \\tablenotetext{c}{FUV and NUV magnitudes were obtained using IRAF task ellipse.} \\tablenotetext{d}{R-band magnitudes are from Rubin et al. 1991.} \\tablenotetext{e}{Extinction corrections using Seibert et al. (2005) for FUV and NUV.} \\end{deluxetable}    \\begin{deluxetable}{ccccrrrr} \\tabletypesize{\\scriptsize} \\tablecaption{FUV sources within HCG100 $10' \\times 10'$ field} \\tablewidth{0pt} \\tablehead{ \\colhead{ID} & \\colhead{RA} & \\colhead{Dec} & \\colhead{R$^{\\rm a}$} & \\colhead{NUV} & \\colhead{FUV} & \\colhead{FUV-R} & \\colhead{FUV-NUV} } \\startdata 1 &    0.2410&  13.0647& 18.80 $\\pm$  0.01&    20.56 $\\pm$   0.04&  20.44  $\\pm$   0.06&   1.64 $\\pm$ 0.06 & -0.12 $\\pm$ 0.07\\\\ 2 &    0.2526&  13.1161& 19.54 $\\pm$  0.01&    21.61 $\\pm$   0.08&  21.48  $\\pm$   0.10&   1.94 $\\pm$ 0.10 & -0.13 $\\pm$ 0.13\\\\ 3 &    0.2722&  13.0236& 19.09 $\\pm$  0.01&    21.61 $\\pm$   0.08&  21.21  $\\pm$   0.09&   2.12 $\\pm$ 0.09 & -0.40 $\\pm$ 0.12\\\\ 4 &    0.2929&  13.0859& 20.25 $\\pm$  0.04&    21.26 $\\pm$   0.07&  21.09  $\\pm$   0.09&   0.84 $\\pm$ 0.10 & -0.17 $\\pm$ 0.12\\\\ 5 &    0.2933&  13.1377& 19.42 $\\pm$  0.01&    22.43 $\\pm$   0.15&  22.79  $\\pm$   0.22&   3.37 $\\pm$ 0.22 &  0.36 $\\pm$ 0.27\\\\ 6 &    0.3076&  13.1630& 17.48 $\\pm$  0.00&    20.19 $\\pm$   0.03&  20.20  $\\pm$   0.05&   2.72 $\\pm$ 0.05 &  0.01 $\\pm$ 0.06\\\\ 7 &    0.3176&  13.0290& 18.43 $\\pm$  0.00&    22.85 $\\pm$   0.11&  21.70  $\\pm$   0.12&   3.27 $\\pm$ 0.12 & -1.15 $\\pm$ 0.16\\\\ 8 &    0.3188&  13.0724& 20.39 $\\pm$  0.01&    22.49 $\\pm$   0.13&  21.18  $\\pm$   0.11&   0.79 $\\pm$ 0.11 & -1.31 $\\pm$ 0.17\\\\ 9 &    0.3429&  13.1208& 20.35 $\\pm$  0.01&    22.24 $\\pm$   0.11&  22.10  $\\pm$   0.16&   1.75 $\\pm$ 0.16 & -0.15 $\\pm$ 0.19\\\\ 10&    0.3535&  13.0704& 19.75 $\\pm$  0.01&    21.92 $\\pm$   0.09&  21.66  $\\pm$   0.11&   1.91 $\\pm$ 0.11 & -0.26 $\\pm$ 0.14\\\\ 11&    0.3627&  13.0532& 19.42 $\\pm$  0.01&    21.63 $\\pm$   0.07&  21.62  $\\pm$   0.10&   2.20 $\\pm$ 0.10 & -0.01 $\\pm$ 0.12\\\\ 12&    0.3704&  13.1375& 18.88 $\\pm$  0.00&    22.15 $\\pm$   0.09&  21.98  $\\pm$   0.14&   3.10 $\\pm$ 0.14 & -0.18 $\\pm$ 0.17\\\\ 13$^{\\rm b}$&\t 0.3735&  13.0986& 17.69 $\\pm$  0.00&\t 19.63 $\\pm$   0.02&  19.59  $\\pm$ 0.04&   1.90 $\\pm$ 0.04 & -0.04 $\\pm$ 0.04\\\\ 13A&   0.3738&  13.0987& 17.83 $\\pm$  0.01&    19.78 $\\pm$   0.05&  19.81  $\\pm$   0.17&   1.98 $\\pm$ 0.17 &  0.03 $\\pm$ 0.18\\\\ 13B&   0.3719&  13.0984& 17.77 $\\pm$  0.01&    19.62 $\\pm$   0.04&  19.83  $\\pm$   0.04&   2.06 $\\pm$ 0.04 &  0.22 $\\pm$ 0.06\\\\ 14&    0.3773&  13.1041& 20.42 $\\pm$  0.01&    21.63 $\\pm$   0.08&  21.10  $\\pm$   0.09&   0.69 $\\pm$ 0.09 & -0.53 $\\pm$ 0.12\\\\ 15&    0.3796&  13.0950& 19.52 $\\pm$  0.01&    22.04 $\\pm$   0.12&  21.79  $\\pm$   0.14&   2.27 $\\pm$ 0.14 & -0.25 $\\pm$ 0.19\\\\ 16&    0.3816&  13.0817& 18.89 $\\pm$  0.00&    21.55 $\\pm$   0.06&  21.30  $\\pm$   0.11&   2.41 $\\pm$ 0.11 & -0.24 $\\pm$ 0.12\\\\ \\enddata \\tablenotetext{a}{Magnitudes (AB) in all bands were obtained with SExtractor Mag$_{-}$auto. Galactic extinction corrections were done using Schlegel et al. (1998) for R and Seibert et al. (2005) for FUV and NUV.} \\tablenotetext{b}{Object 13 was separated into two objects, A and B, using IRAF polyphot task.} \\end{deluxetable}    \\begin{deluxetable}{cccccccccc} \\tabletypesize{\\scriptsize} \\tablecaption{Derived Properties} \\tablewidth{0pt} \\tablehead{ \\colhead{ID} & \\colhead{L$_{\\rm FUV}$ (erg/s)$^{\\rm a}$} & \\colhead{L$_{\\rm NUV}$ (erg/s)} & \\colhead{age$^{\\rm b}$} & \\colhead{SFR$_{\\rm FUV}$ $^{\\rm c}$} & \\colhead{SFR$_{\\rm NUV}$} & \\colhead{log(M$_{\\rm *}$)$^{\\rm d}$} & \\colhead{log(M$_{\\rm HI}$)$^{\\rm e}$} & \\colhead{Sep$^{\\rm f}$} & \\colhead{log I$_{1530}$$^{\\rm g}$} } \\startdata 1  &  5.90E+040  & 5.70E+040  & 3.9\t & 0.005&  0.007 & 5.0 &   & 131.2\t  & 6.02 \\\\ 2  &  2.27E+040  & 2.16E+040  & 3.8\t & 0.002&  0.003 & 4.6 &   & 102.2\t  & 5.91 \\\\ 3  &  2.91E+040  & 2.17E+040  & $<$1\t & 0.002&  0.003 & 4.7 & 9.6 & 137.2   & 5.90 \\\\ 4  &  3.25E+040  & 2.98E+040  & 3.3\t & 0.003&  0.004 & 4.7 & 10.4 & 60.7    & 5.35 \\\\ 5  &  6.79E+039  & 1.02E+040  & 194.1\t & 0.001&  0.001 & 6.5 & 9.9$\\dagger$ & 61.2    & 4.99 \\\\ 6  &  7.41E+040  & 8.01E+040  & 26.9\t & 0.006&  0.010 & 6.3 & 9.2$\\dagger$ & 74.7    & 6.31 \\\\ 7  &  1.84E+040  & 6.90E+039  & $<$1\t & 0.001&  0.001 & 4.5 &     & 108.5\t & 5.51 \\\\ 8  &  2.98E+040  & 9.62E+039  & $<$1\t & 0.002&  0.001 & 4.7 & 10.0 & 53.9    & 5.31 \\\\ 9  &  1.28E+040  & 1.21E+040  & 3.5\t & 0.001&  0.001 & 4.6 & 9.8$\\dagger$ & 17.5    & 5.80 \\\\ 10 &  1.92E+040  & 1.63E+040  & 2.9\t & 0.002&  0.002 & 4.4 &     & 58.8\t  & 5.97 \\\\ 11 &  2.00E+040  & 2.13E+040  & 21.0\t & 0.002&  0.003 & 5.6 &     & 83.7\t  & 5.85 \\\\ 12 &  1.43E+040  & 1.31E+040  & 3.2\t & 0.001&  0.002 & 4.3 &     & 58.0\t  & 5.95 \\\\ 13 &  1.29E+041  & 1.35E+041  & 13.7\t & 0.010&  0.016 & 6.2 & 9.5$\\dagger$$\\dagger$ & 53.3    & 6.67 \\\\ 13 A& 1.03E+041  & 1.36E+041  & 118.7\t & 0.008&  0.017 &     &  &  & 6.51 \\\\ 13 B& 1.06E+041  & 1.17E+041  & 33.8\t & 0.008&  0.014 &     &  &  & 6.50 \\\\ 14 &  3.22E+040  & 2.13E+040  & $<$1\t & 0.003&  0.003 & 4.8 & 9.5$\\dagger$$\\dagger$ & 56.3\t& 6.30 \\\\ 15 &  1.70E+040  & 1.46E+040  & 2.9\t & 0.001&  0.002 & 4.3 & 9.5$\\dagger$$\\dagger$ & 62.1    & 5.07 \\\\ 16 &  2.67E+040  & 2.29E+040  & 2.9\t & 0.002&  0.003 & 4.5 &     & 72.0\t  & 5.26 \\\\ \\enddata \\tablenotetext{a}{FUV and NUV luminosities are in erg/s, divide by the FUV and NUV bandwidths (269\\AA\\ and 616\\AA, respectively) to obtain in L erg/s/\\AA.} \\tablenotetext{b}{Age (Myr) from FUV-NUV using Thilker et al. (2007) assuming a Milky Way internal extinction (E(B-V)=0.2).} \\tablenotetext{c}{SFR (M$_{\\odot}$/yr) from Iglesias-P\\'aramo et al. (2006) using FUV and NUV without correcting for internal extinction.} \\tablenotetext{d}{Stellar mass (M$_{\\odot}$) obtained from Starburst99 using ages (column 4) and L$_{\\rm 1530}$ (erg/s/\\AA).} \\tablenotetext{e}{MHI (M$_{\\odot}$) was calculated using 2.36 $\\times$ 10$^{5}$ F$_{\\rm HI}$ D$^{2}$, where D is in Mpc and F$_{\\rm HI}$ in Jy Km/s. F$_{\\rm HI}$ was measured within one beam size 61.0$''$ $\\times$ 55.23$''$ and reflects the HI mass in the vicinity of each object.} \\tablenotetext{f}{Distance (kpc) between each object and HCG100a.} \\tablenotetext{g}{FUV surface brightness (L$_{\\odot}$ kpc$^{-2}$) is defined following Hoopes at al. (2007).} \\tablenotetext{\\dagger}{HI masses of Objects \\#5, \\#6, and \\#9 should be taken with caution due to contamination by large galaxies in the vicinity.} \\tablenotetext{\\dagger\\dagger}{HI masses of Objects \\#13, \\#14, and \\#15 are not individual masses of each object but the total mass within one beam centered in that region. Contamination from H100b mass is also possible.} \\end{deluxetable} ",
        "conclusions": "We have analyzed the UV and optical light in combination with the H{\\sc i} gas within the compact group of galaxies HCG100 and identified 16 star-forming regions in the intragroup region. The young age ($<200$~Myr) of these objects and the proximity to the tidal tail connects the OB stars formation time scale ($\\sim$10$^{8}$ yr) with the dynamic time scale of the tidal features. Moreover, the H{\\sc i} clouds to which many of them are associated have column densities about one order of magnitude lower than the N(HI)$\\sim$10$^{21}$ cm$^{-2}$ thought to be required for triggering star formation. So, in these cases, we have a strong indication that the H{\\sc i} clouds must have suffered recent collisions which could have then triggered the star formation process. Based on their ages, stellar masses and H{\\sc i} masses in their vicinities, we conclude that some of these objects are tidal dwarf galaxies with ongoing star formation and some are intergalactic HII regions or conglomeration of stellar clusters.   Although we have an estimate of the amount of neutral gas in these objects, the main ingredient in the star-formation process, the molecular gas, is unknown for these objects. From CO observations of tidal dwarf galaxies (TDGs), Braine et al. (2001) provided strong evidence that TDGs are self-gravitating entities with large amounts of atomic gas which will transform into molecular gas and subsequently form stars. CO observations together with spectroscopy of these objects will be of great value in understanding their nature. For instance, we cannot exclude the possibility that some of these objects are not associated with HCG100, but at least those falling within HI density peaks are likely to be associated with the group. Moreover, we will be able to estimate metallicity of these newly formed regions and evaluate the enrichment level of the processed gas in the intragroup medium.   \\begin{deluxetable}{lcccc} \\tablecaption{Comparison of Galaxy Properties} \\tablewidth{0pt} \\tablehead{ \\colhead{Parameter$^{\\rm a}$} & \\colhead{Large UVLGs} & \\colhead{Compact UVLGs} & \\colhead{LBGs} & \\colhead{HCG100 IGs} } \\startdata log L$_{1530}$ (L$_{\\odot}$)                     & 10.3 to 11.2  & 10.3 to 11.0  & 10.3 to 10.9 &  8 to 9 \\\\ log I$_{1530}$ (L$_{\\odot}$ kpc$^{-2}$) $^{\\rm b}$& 6.0 to 8.0    & 8.0 to 10.3   & 9.0 to 10.3 &  5.6 to 8.3 \\\\ log SFR (M$_{\\odot}$ yr$^{-1}$) $^{\\rm c}$       & 0 to 1.5      &  0.2 to 2.0   & 0.5 to 2.0   & -2 to -0.4   \\\\ FUV-R   $^{\\rm d}$                               & 1.0 to 3.5    &  0.2 to 2.8   & 0.2 to 1.7   &  0.7 to 3.4 $^{\\rm f}$ \\\\ \\enddata \\tablenotetext{a}{Parameters for UVLGS and LBGs are from Hoopes et al. (2007). The lowest value for HCG100 intragroup objects was calculated using average values A$_{1500}$=1.6 and the highest value using A$_{1500}$=4.1 from Calzetti (2001).} \\tablenotetext{b}{I$_{1530}$ is the FUV surface brightness as described in Hoopes et al. (2007), except for HCG100 intragroup objects where we used R as the FUV petrosian radius.} \\tablenotetext{c}{SFR for the HCG intragroup objects using FUV equation from Iglesias-P\\'aramo et al. (2006).} \\tablenotetext{d}{r magnitudes are from SDSS and corrected for extinction, except for HCG100 which is the same as in Table 1 (Cousin R magnitude) and is not corrected for extinction.} \\tablenotetext{e}{FUV-R shown is not corrected for internal extinction. Using Calzetti (2001) A$_{1500}$=1.6 and  A$_{r}$=0.48 FUV-R=-0.3 to 2.2.} \\end{deluxetable}  \\clearpage  \\begin{figure} \\plotone{f1.eps} \\caption{GALEX FUV image 10$'$$\\times$10$'$ with H{\\sc i} contours. Four HCG100 members are marked as `a, b, c and d'. Intragroup objects are circled (radius=8$''$) and numbered. VLA NH{\\sc i} contours are 0.6, 1.2, 2.1, 3.6, 4.4, 5.1, 5.9, 6.6, 7.4 $\\times$10$^{20}$ cm$^{-2}$, beam size (61.0$''$$\\times$55.23$''$) is shown on left corner. North is to the top and East to the left. \\label{fuvhi_10arcmin}} \\end{figure}  \\begin{figure} \\plotone{f2.eps} \\caption{GALEX NUV image 10$'$$\\times$10$'$ with H{\\sc i} contours. Intragroup objects are circled (radius=8$''$) and numbered. VLA NH{\\sc i} contours are 0.6, 1.2, 2.1, 3.6, 4.4, 5.1, 5.9, 6.6, 7.4 $\\times$10$^{20}$ cm$^{-2}$. Beam size (61.0$''$$\\times$55.23$''$) is shown in Fig.\\ref{fuvhi_10arcmin}. North is to the top and East to the left. \\label{nuvhi_10arcmin}} \\end{figure}   \\begin{figure} \\plotone{f3.eps} \\caption{GALEX FUV-NUV versus age from Thilker et al. models (2007) are shown as solid line (no extinction correction) and dotted line (internal extinction of the Milky Way E(B-V)=0.2). 4 objects are not included, they have ages $<$ 1~Myr old. \\label{colorsfuvnuvthilker}} \\end{figure}   \\begin{figure} \\plotone{f4.eps} \\caption{GALEX FUV-NUV versus FUV-R of the intragroup objects. Squares and triangles are FUV--R colors assuming an error of 0.2 magnitudes in the R-magnitude. Models from Thilker et al. 2007 are shown as solid line (no extinction correction) and dotted line (internal extinction of the Milky Way E(B-V)=0.2). Crosses mark ages 10, 50, 100, 200, and 500 Myr, from the bottom left to the top right. \\label{colorsfuvnuvfuvr}} \\end{figure}  \\begin{figure} \\plotone{f5.eps} \\caption{R images of all intragroup objects, circle has radius=8$''$, each cutout is 1$'$$\\times$1$'$. NH{\\sc i} contours (white) are the same as in Fig.\\ref{fuvhi_10arcmin}. \\label{all_obj}} \\end{figure}  \\clearpage"
    },
    "0710/0710.4397_arXiv.txt": {
        "abstract": "We describe a finite-volume method for solving the Poisson equation on oct-tree adaptive meshes using direct solvers for individual mesh blocks. The method is a modified version of the method presented by Huang and Greengard (2000), which works with finite-difference meshes and does not allow for shared boundaries between refined patches. Our algorithm is implemented within the FLASH code framework and makes use of the PARAMESH library, permitting efficient use of parallel computers. We describe the algorithm and present test results that demonstrate its accuracy. ",
        "introduction": "\\label{Sec:intro}  Astrophysical simulations commonly need to solve the Poisson equation, \\begin{equation} \\label{Eqn:Poisson} \\nabla^2 \\phi({\\bf x}) = 4 \\pi G \\rho({\\bf x})\\ , \\end{equation} for the gravitational potential $\\phi({\\bf x})$ given a density distribution $\\rho({\\bf x})$. Similar equations also arise in other contexts, such as incompressible flow problems and divergence-cleaning methods for magnetohydrodynamics. Self-gravitating problems offer special challenges because they frequently develop structure spanning large spatial dynamic ranges. The problem of spatial dynamic range is particularly acute for grid-based schemes for solving the Euler equations of hydrodynamics.  Within the context of grid-based methods for solving the Poisson equation, several approaches to the problem of spatial dynamic range have arisen. The simplest approach is to use Fourier transforms, multigrid methods, or sparse iterative solvers on uniform Eulerian grids. The maximum dynamic range is then limited by the available memory. Recently Trac and Pen (2006) have demonstrated an out-of-core uniform-grid Poisson solver that exceeds this limit by making use of disk space; the largest published calculations with this solver have used $4000^3$ zones. However, storage resource consumption still increases with the third power of the resolution, putting grids with $10^4$ zones on a side or larger out of reach for now.  If high resolution is not needed everywhere in the domain, as is frequently the case in cosmological structure formation simulations, it is also possible to employ nonuniform Eulerian or Lagrangian grids. Examples include COSMOS (Ricker, Dodelson, \\& Lamb 2000), which uses a nonuniform multigrid solver, and MMH (Pen 1998), which uses a deformable mesh. These methods work best when the region to be resolved is known beforehand, although fully Lagrangian codes like Pen's can follow the development of structures and adjust zone spacing appropriately. Nonuniform grids, however, introduce complicated position-dependent stencils and generally cannot be used with fast transform-based solvers. In addition, coupled numerical hydrodynamics methods generally place constraints on the allowed mesh anisotropy and nonuniformity, since numerical dissipation increases with zone spacing.  The greatest spatial dynamic ranges in grid-based astrophysical simulations have been achieved using adaptive mesh refinement (AMR) techniques. Modern AMR techniques for solving hyperbolic systems of equations were first developed by Berger and Oliger (1984) and Berger and Colella (1989). In the Berger and Colella formulation, AMR involves the construction of a hierarchical set of mesh ``patches'' with decreasing zone spacing. The coarsest mesh covers the entire computational domain, while more highly refined meshes cover only a portion. Generally refined meshes are taken to be nested; that is, each refined mesh lies completely within its coarser parent mesh. Examples of astrophysical codes employing patch-based AMR meshes include the code of Truelove et al.\\ (1998), AMRA (Plewa \\& M\\\"uller 2001), RIEMANN (Balsara 2001), Enzo (O'Shea et al.\\ 2004), and CHARM (Miniati \\& Colella 2007). To date self-gravitating AMR calculations have achieved effective spatial resolutions greater than $10^{15}$.  A considerable simplification of the Berger and Colella method was introduced by Quirk (1991) and de~Zeeuw and Powell (1993). Known as ``oct-tree'' AMR, this method requires that each refined patch contain the same number of zones, that each refinement level have zones a factor of two smaller in each dimension than the next coarser level, and that each refined patch be no more than one level removed from its immediate neighbor. Mesh data can then be stored in an oct-tree data structure, allowing for extremely efficient parallel implementations, even on high-latency systems (Warren \\& Salmon 1993). Also, because each mesh patch (often called a ``block'' in this context) contains the same number of zones, it is possible to achieve high levels of cache re-use when iterating over zones. Unless each block contains a very small number of zones, this efficiency comes with the price that refined blocks often must cover more of the volume than they would in a patch-based method. The primary astrophysical simulation code employing oct-tree AMR is FLASH (Fryxell et al.\\ 2000), which uses the PARAMESH library (MacNeice et al.\\ 2000) to handle its AMR mesh. (The ART (Kravtsov, Klypin, \\& Khoklov 1997), MLAPM (Knebe, Green, \\& Binney 2001), and RAMSES (Teyssier 2002) codes also use tree structures to manage AMR meshes, but in these codes the refined blocks consist of a single zone each, and the base mesh generally contains a large number of zones, so these codes do not employ oct-trees. A block-based AMR approach that allows for ``incomplete families'' has also recently been implemented within the VAC code (van der Holst \\& Keppens 2007); the oct-trees discussed in the current paper require complete ``families'' of child blocks.) By default PARAMESH uses blocks containing $8^3$ interior zones as a compromise between adaptive flexibility and memory efficiency, but any size larger than the differencing stencil and small enough to fit in the memory attached to a single processor can be employed.  Poisson solvers on oct-tree meshes generally employ some type of multigrid or sparse linear solver iteration scheme. Transform methods cannot be employed directly because of the varying mesh resolution and non-tensor-product character of the composite mesh.  For example, Matsumoto and Hanawa (2003) describe a relaxation-based method for solving the Poisson equation on nested grids. Within the FLASH framework, we have employed the Martin and Cartwright (1996) multigrid algorithm for several years. However, the speed and scalability of this algorithm have been limited because of the need to apply multiple relaxation iterations on each level, together with the communication of block boundary data that such iterations require. Because the cost of this algorithm dominates the cost of most self-gravitating simulations with FLASH, we are motivated to develop more efficient methods that require less communication.  In this paper we describe one such method, based on the direct multigrid algorithm of Huang and Greengard (2000, hereafter HG). This algorithm improves considerably upon relaxation-based multigrid solvers by allowing refined patches to be solved directly using ``black-box'' uniform-grid solvers. Unlike the direct method described by Couchman (1991), the HG algorithm properly minimizes the global residual by allowing information to flow back from fine meshes to coarse meshes. However, it is formulated on a finite-difference mesh in which refined patches are not permitted to touch. Here we describe a modified version of the HG algorithm suitable for finite-volume oct-tree AMR meshes. We have implemented this algorithm within the FLASH framework, and we present test results that demonstrate the solver's accuracy. The present paper should be regarded as a companion to Fryxell et al.\\ (2000) and a methodological description for future FLASH-based simulation papers in the areas of cosmic structure formation, galaxy cluster physics, star cluster formation, and binary star evolution, among others.  The paper is organized as follows. In \\S~\\ref{Sec:algorithm} we give a precise description of the algorithm. In \\S~\\ref{Sec:tests} we present the results of test problems run with the new solver. We conclude in \\S~\\ref{Sec:conclusions} with some remarks on performance. All calculations described in this paper were performed using version 2.4 of FLASH. ",
        "conclusions": "\\label{Sec:conclusions}  We have detailed the modifications needed in order to use the Huang \\& Greengard (2000) algorithm for Poisson's equation on oct-tree AMR meshes. This method allows us to use a local direct Poisson solver with Dirichlet boundary conditions on each block, yet it correctly minimizes the residual across the composite AMR mesh. The HG algorithm must be modified to take into account the fact that mesh quantities represent zone-averaged values and the fact that block boundaries can coincide on an oct-tree mesh. Because adjacent coarse and fine blocks do not share points in common for oct-tree meshes, additional interpolation (as compared to HG) is needed to set boundary values. The resulting errors at block corners reduce convergence to first order, but a fixed small number of boundary-zone relaxation steps restores the desired second-order convergence rate for block corners in uniformly refined regions. A higher-order scheme may be able to eliminate the need for such relaxation. Our test results show that, while jumps in refinement degrade convergence somewhat in comparison with solving on a uniform mesh, the effect is manageable because oct-tree meshes are usually fully refined up to some level. Thus first-order convergence takes over only once the error has been significantly reduced at second order on fully refined levels.  The parallel scaling of this solver, as implemented using the Message Passing Interface (MPI) in the FLASH code, is comparable to or slightly better than that of the relaxation multigrid solver distributed with FLASH 2.$x$. With constant total work and increasing processor count, parallel efficiency is close to 100\\% up to 8 -- 16 times the smallest number of processors on which a run can fit. When gasdynamics is included, the Poisson solver requires $\\sim$50\\% of the execution time; for particle-only simulations the amount is $\\sim $70 -- 80\\%. In comparison with the older solver, a factor of 2 -- 5 improvement in performance is often seen. The solver described here will be made available to the public as part of FLASH~3.0.\\footnote{FLASH is freely available at http://flash.uchicago.edu/.}  The primary scaling bottleneck for this and other multigrid algorithms is the fact that on the coarsest level there are too few zones to distribute among all of the processors. Since each V-cycle descends to the coarsest level, processors controlling blocks on finer levels must wait until the coarsest level is finished before proceeding. To counteract this work starvation, we are investigating the use of a uniform-grid parallel FFT solver (T.~Theuns, private communication) to handle the coarsest level in a distributed fashion. Additional strategies may be necessary to make optimal use of the petaflop computing resources now beginning to become available."
    },
    "0710/0710.4674_arXiv.txt": {
        "abstract": "We present a study of large-scale bars in the local Universe, based on a large sample of $\\sim3692$  galaxies,  with $-18.5\\leq M_g < -22.0$ mag and redshift $0.01\\leq z<0.03$, drawn from the Sloan Digitized Sky Survey. While most studies of bars in the local Universe have been based on relatively small samples that are dominated by bright galaxies of early to intermediate Hubble types with prominent bulges, the present sample is $\\sim$~10 times larger, covers a larger volume, and includes many galaxies that are disk-dominated and of late Hubble types. Both color cuts and S\\'ersic cuts yield a similar sample of $\\sim2000$ disk galaxies. We characterize bars and disks by ellipse-fitting  $r$-band images and applying quantitative criteria. After excluding highly inclined ($>60^{\\circ}$) systems, we find the following results. (1)~The optical $r$-band fraction ($f_{\\rm opt-r}$) of barred galaxies, when averaged over the whole sample, is $\\sim48\\%-52\\%$. The bars have diameters $d$ of 2 to 24 kpc, with most ($\\sim72\\%$) having $d\\sim$ 2 to 6 kpc. (2)~When galaxies are separated according to half light radius ($r_{\\rm e}$), or normalized $r_{\\rm e}$/$R_{\\rm 24}$, which is a measure of the bulge-to-disk ($B/D$) ratio, a remarkable result is seen: $f_{\\rm opt-r}$ rises sharply, from $\\sim$~40\\% in galaxies that have small $r_{\\rm e}$/$R_{\\rm 24}$ and visually appear to host prominent bulges, to $\\sim$~70\\% for galaxies that have large $r_{\\rm e}$/$R_{\\rm 24}$ and appear disk-dominated. Visual classification, performed for $\\sim900$ galaxies, confirms our result that disk-dominated galaxies with no bulge or a very low $B/D$ display a significantly higher optical bar fraction ($>70$\\% vs 40\\%) than galaxies with prominent bulges. It also shows that barred galaxies host a larger fraction (31\\% vs 5\\%) of quasi-bulgeless disk-dominated galaxies than do unbarred galaxies. (3)~$f_{\\rm opt-r}$ rises for galaxies with bluer colors (by $\\sim30\\%$) and lower masses (by $\\sim15\\%-20\\%$). (4) The significant rise in the optical bar fraction toward late-type galaxies is discussed in terms of their higher gas mass fraction, higher dark matter fraction, and lower bulge-to-disk ratio. (5) While hierarchical $\\Lambda$CDM models of galaxy evolution models fail to produce galaxies without classical bulges, our study finds that $\\sim20\\%$ of disk galaxies  appear to be ``quasi-bulgeless''. (6)~Our study of bars at $z\\sim0$ in the optical $r$-band provides the $z\\sim0$ comparison point for $HST$ ACS surveys (e.g., GEMS, GOODS, COSMOS) that measure the rest-frame optical bar fraction in bright galaxies out to $z\\sim1$. After applying the same cutoffs in magnitude, bar ellipticity ($e_{\\rm bar}\\ge0.4$), and bar size ($a_{\\rm bar}\\ge1.5$ kpc), which are applied in $z\\sim0.2-1.0$ studies in order to trace strong bars with adequate spatial resolution in bright disks, we obtain an optical $r$-band bar fraction of $34\\%$. This is comparable to the value reported at $z\\sim0.2-1.0$, implying that the optical fraction of strong bars does not suffer a dramatic order of magnitude decline in bright galaxies out to $z\\sim1$. ",
        "introduction": "The majority ($\\sim$~60\\%) of bright disk galaxies are barred, when observed in the near-infrared \\citep[][hereafter MJ07]{kna99,esk00,lau04,men07,mar07} and a significant fraction of these ($\\sim$~45\\%) also appear barred in the optical \\citep[][MJ07]{esk00}. Earlier studies suggested a striking or order of magnitude decline in the optical fraction of bars out to $z\\sim1$ \\citep{abr99,vdb00}, but subsequent studies have ruled out an order of magnitude decline and find that the optical fraction of strong bars remains fairly constant or show a moderate decline of a factor of $\\sim2$ \\citep[][Sheth et al. 2003,2004,2007; see $\\S$ \\ref{discu}]{jog04,elm04a,zhe05}.  Bars are believed to be very important with regard to the dynamical and secular evolution of disk galaxies, particularly in redistributing the angular momentum of the baryonic and dark matter components of disk galaxies \\citep{com81,wei85,com90,deb00}. The interaction between the bar and the disk material can lead to the inflow of gas from the outer disk to the central parts, which can trigger starbursts \\citep{elm94,kna95,reg04,jog05,sht05} and might contribute to the formation of disky bulges \\citep[or 'pseudobulges',][]{kor93,kor04,ath05a,jog05,deb06}. Additional evidence for secular evolution is provided by box- or peanut-shaped bulges in inclined galaxies. These features are commonly attributed to the orbital structure, resonances, and vertical instabilities in a barred potential \\citep{com90,kui95,bur99,bur05}.  From a theoretical perspective, it is possible to model some aspects of the evolution of disks and bars, and their interactions (e.g., the corresponding simulations are able to reproduce certain broad features of barred disks). However, it remains unclear why a specific galaxy has a bar, but a seemingly similar galaxy is unbarred; or why some barred galaxies have a classical bulge, whereas others harbor a disky bulge, etc. This might indicate that specific properties of the disks or the particular processes involved in their formation have a strong impact on their ability to form a bar. In order to investigate, how disk and bar formation are related, it is not only important to determine the fraction of disk galaxies that are barred, but also to relate bar and disk properties. There are different methods to find and characterize bars. The Third Reference Catalog of Bright Galaxies \\citep[][hereafter RC3]{dev91} uses three bar strength families (SA, SAB, and SB) to characterize bars based on a visual inspection of blue light images. Using this classification \\cite{ode96} showed that the optical fraction of strong bars in disk galaxies rises from Sc galaxies towards later-types. More quantitative measures, such as the gravitational torque method \\citep{blo02,lau02,but05}, or Fourier dissection \\citep{but06,lau06}, were also used, not only to find bars, but to quantitatively determine bar strengths and bar lengths. Similarly, the method of fitting ellipses to galaxy isophotes provides a tool to characterize the length and shape of bars \\citep{fri96,jog99,kna00,sht00,lai02,why02,jog02a,jog02b,sht03,elm04a,ree07,mar07,men07,sht07}.  These efforts were able to shed light on the fraction, shapes, and structures of bars  in local disk galaxies of early to intermediate Hubble types. First attempts were made to relate the presence of a bar or its structural properties to other galaxy characteristics. However, there were three important limitations. Firstly, samples used in earlier studies were small ($\\sim100$ to 200 objects) and mostly composed of bright galaxies of early to intermediate Hubble types (Sa to Sc), with fairly prominent bulges. One could barely get decent number statistics for bars in early-type disk galaxies, while the bins of disk-dominated late Hubble types were dominated by Poisson noise (e.g., see Figure 16 in MJ07). Secondly, with such small samples, it was difficult to bin galaxies in terms of the galaxy host properties. Thirdly, earlier samples were drawn from a very small volume, and could be highly impacted by cosmic variance.  In the present study, we use a sample of $\\sim2000$ galaxies, at $z=0.01-0.03$ with $M_r \\sim -18.5$ to $-22.0$ mag. The first advantage of this study is that it provides a factor of 10 improvement in number statistics and reduces the effect of cosmic variance by selecting galaxies drawn from a larger volume. Secondly, with $\\sim2000$ galaxies, we can for the first time have  $100-200$ galaxies per bin, while binning galaxies in terms of host galaxy parameters, such as  luminosity, measures of bulge-to-disk ($B/D$) ratios, size, colors, surface brightness, etc. This allows us to  conduct a comprehensive study of barred and unbarred galaxies as a function of host galaxy properties. Thirdly, our sample has a large number of galaxies, which are relatively faint ($M_g>-19.5$ mag) or/and appear disk-dominated, characteristic of late Hubble types. This allows us to shed light on what happens to bars at the fainter end of the luminosity function and in the regime of disk-dominated galaxies.  A fourth goal of our study is to provide a reference baseline for bars at $z\\sim0$ in the {\\it rest-frame optical} for intermediate redshift $HST$ surveys using the Advanced Camera for Surveys (ACS), such as the Tadpole field \\citep{tra03}, the Galaxy Evolution from Morphologies and SEDs \\citep[GEMS,][]{rix04}, the Great Observatories Origins Deep Survey \\citep[GOODS,][]{gia04}, and COSMOS \\citep{sco06}, which trace bars in the rest-frame optical band at $z\\sim0.2-1.0$ (look-back times of 3--8 Gyr). We use SDSS to provide the reference point at $z=0$ in the $r$-band, complementing the one in the $B$-band of MJ07. Our $B$- and $r$-band results can be directly compared to $HST$ ACS optical studies of bars in bright disks at $z\\sim0.2-1.0$ \\citep{elm04a,jog04}. The validity of this comparison is reinforced by the fact that we use the same procedure of ellipse fits ($\\S$ \\ref{meth}) that were used by these studies. We also note that the reference $z=0$ point for bars in the near-infrared band \\citep{men07} is not appropriate for comparison with the above $HST$ ACS surveys, which trace the  rest-frame optical rather than  the rest-frame near-infrared.  The outline of the paper is as follows: In $\\S$ \\ref{samsel} we present our sample selection. The method used to find and characterize bars is explained in $\\S$ \\ref{meth}. In $\\S$ \\ref{resol} we discuss the detection limits. Our results and more detailed assessments of specific findings are presented in $\\S$ \\ref{resul}. We discuss our results in $\\S$ \\ref{discu} and summarize our conclusions in $\\S$ \\ref{sum}. Throughout the paper we assume a flat cosmology with $\\Omega_M=1-\\Omega_{\\Lambda}=0.3$ and $H_0=70$ km~s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "\\label{sum} We have used the $r$-band images from the NYU-VAGC of a sample of 3692 galaxies with $-18.5\\leq M_g < -22.0$ mag and redshift $0.01\\leq z<0.03$ to find and characterize bars. While most studies of bars in the local Universe have been based on relatively small samples that are dominated by bright early type (Sa to Sc) galaxies with bulges, the present sample also includes many galaxies that are disk-dominated and of late Hubble types. Furthermore, the sample is $\\sim$~10 times larger and samples a larger volume than earlier local samples We used a color cut in the color-magnitude diagram to select $\\sim2000$ disk galaxies.  We cross-check that  S\\'ersic cuts  would yield a similar sample. We identify and characterize bars and disks using $r$-band images and a method based on ellipse fits and quantitative criteria. The typical seeing ($1\\farcs4$ or 290 to 840 pc over $0.01\\leq z<0.03$) is adequate for resolving large-scale bars, whose typical diameters are $\\ge$ 2 kpc. Smaller nuclear bars are not the focus of this study. After the standard procedure of excluding highly inclined ($>60^{\\circ}$) systems, we find the following results.  \\begin{enumerate} \\item The average optical $r$-band bar fraction ($f_{\\rm opt-r}$) in our sample, which primarily consists of late-type disk-dominated galaxies, is $\\sim48\\%-52\\%$. The bars have diameters $d$ of 2 to 24 kpc, with most ($\\sim72\\%$) having $d\\sim$ 2 to 6 kpc (Figure \\ref{barp}a). The bar length is typically much smaller than $R_{24}$ (Figure \\ref{r24}a) and most galaxies have a $a_{bar}/R_{24}$ in the range 0.2 to 0.4 (Figure \\ref{r24}b).  \\item When galaxies are separated according to normalized $r_{\\rm e}$/$R_{\\rm 24}$, which is  a measure of  the bulge-to-disk ($B/D$) ratio, a remarkable result is seen: the optical $r$-band fraction rises sharply, from $\\sim$~40\\% in galaxies that have small $r_{\\rm e}$/$R_{\\rm 24}$ and visually appear bulge-dominated, to $\\sim$~70\\% for galaxies that have large $r_{\\rm e}$/$R_{\\rm 24}$. Visual classification of $\\sim80\\%$ of our sample (with $i<60^{\\circ}$) confirms our result that {\\it late-type disk-dominated galaxies with no bulge or a very low $B/D$ display a significantly higher optical bar fraction ($>70$\\% vs 40\\%) than galaxies with prominent bulges.} It also shows that barred galaxies host a larger fraction (31\\% vs 5\\%) of quasi-bulgeless disk-dominated galaxies than do unbarred galaxies. The bar ellipticities or strengths are on average higher in faint disk-dominated galaxies than in bulge-dominated galaxies (Figure \\ref{visp2}d).  \\item Similar trends in the optical bar fraction are found using the central surface brightness and color. Bluer galaxies have higher  bar fractions ($\\sim58\\%$ at $g-r=0.3$) than the redder objects ($\\sim32\\%$ at $g-r=0.65$) (Figure \\ref{barf}d). The optical $r$-band fraction also shows a slight rise for galaxies with fainter luminosities (Figure \\ref{barf}c) and  lower masses (Figure \\ref{massp}). This is expected from (2), given that late-type galaxies are fainter, bluer, and less massive.  \\item The significant rise in the optical bar fraction toward disk-dominated galaxies is discussed in terms of their higher gas mass fraction, higher dark matter fraction, and lower bulge-to-disk ratio.  \\item While many hierarchical $\\Lambda$CDM models of galaxy evolution models fail to produce galaxies without classical bulges, our study finds that in the range  $-18.5\\leq M_g < -22.0$  mag and redshift $0.01\\leq z<0.03$, $\\sim20\\%$ of the 1144 moderately inclined disk galaxies appear to be ``quasi-bulgeless'',  without a classical bulge.  \\item Our study of bars at $z\\sim0$ in the optical $r$ band provides a reference $z\\sim0$ baseline for intermediate redshift $HST$ ACS surveys that trace bars in {\\it bright} disks in the rest-frame optical bands ($BVRI$) out to $z\\sim1$. By applying the same cutoffs in magnitude, bar ellipticity ($e_{\\rm bar} \\geq0.4$), and bar size ($a_{\\rm bar} \\geq1.5$ kpc), which are applied in $z\\sim0.2-1.0$ studies in order to trace strong bars with adequate spatial resolution in bright disks, we obtain an optical $r$-band fraction for strong bars of $34\\%$. This is comparable to the values of $\\sim30\\%$ at $z\\sim0.2-1.0$, $\\sim$~36\\%~$\\pm$~6\\% at $z\\sim0.2-0.7$, and $\\sim$~24\\%~$\\pm$~4\\% at $z\\sim0.7-1.0$. Our result implies that the optical fraction of strong bars in bright galaxies does not suffer any dramatic order of magnitude decline out to $z\\sim$~1. \\end{enumerate}"
    },
    "0710/0710.4442_arXiv.txt": {
        "abstract": "We construct for the first time, the sequences of stable neutron star (NS) models capable of explaining simultaneously, the glitch healing parameters, $Q$, of both the pulsars, the Crab ($Q \\geq 0.7$) and the Vela ($Q \\leq 0.2$), on the basis of starquake mechanism of glitch generation, whereas the conventional NS models cannot give such consistent explanation. Furthermore, our models also yield an upper bound on NS masses similar to those obtained in the literature for a variety of modern equations of state (EOSs) compatible with causality and dynamical stability. If the lower limit of the observational constraint of (i) $Q \\geq 0.7$ for the Crab pulsar and (ii) the recent value of the moment of inertia for the Crab pulsar (evaluated on the basis of time-dependent acceleration model of the Crab Nebula) , $I_{\\rm Crab,45} \\geq 1.93$ (where $I_{45}=I/10^{45}\\,{\\rm g.cm}^2$), both are  imposed together on our models, the models yield the value of matching density, $E_b = 9.584 \\times 10^{14}{\\rm\\,g\\,cm}^{-3}$ at the core-envelope boundary. This value of matching density yields a model-independent upper bound on neutron star masses, $M_{\\rm max} \\leq 2.22 M_\\odot$, and the strong lower bounds on surface redshift $z_R \\simeq 0.6232$ and mass $M \\simeq 2.11 M_\\odot$ for the Crab ($Q \\simeq 0.7$) and the strong upper bound on surface redshift $z_R \\simeq 0.2016 $, mass $M \\simeq 0.982 M_\\odot$ and the moment of inertia  $I_{\\rm Vela,45} \\simeq 0.587$ for the Vela ($Q \\simeq 0.2$) pulsar. However, for the observational constraint of the `central' weighted mean value $Q \\approx 0.72$, and $I_{\\rm Crab,45} > 1.93$, for the Crab pulsar, the minimum surface redshift and mass of the Crab pulsar are slightly increased to the values $z_R \\simeq 0.655$ and  $M \\simeq 2.149 M_\\odot$ respectively, whereas corresponding to the `central' weighted mean value $Q \\approx 0.12$ for the Vela pulsar, the maximum surface redshift, mass and the moment of inertia for the Vela pulsar are slightly decreased to the values $z_R \\simeq 0.1645,\\, M \\simeq 0.828 M_\\odot$ and $I_{\\rm Vela,45} \\simeq 0.459$ respectively. These results set an upper and lower bound on the energy of a gravitationally redshifted electron-positron annihilation line in the range of about 0.309 - 0.315 MeV from the Crab and in the range of about 0.425 - 0.439 MeV from the Vela pulsar. ",
        "introduction": "The data on the glitch healing parameter, $Q$, provide the best tool for testing the starquake (Ruderman 1972; Alpar et al 1996) and Vortex unpinning (Alpar et al 1993) models of glitch generation in pulsars. Both of these mechanisms of glitch generation, in fact, consider NSs, in general, a two component structure: a superfluid interior core surrounded by a rigid crust (in the present study we shall use the term `envelope' which includes the solid crust and other interior portion of the star right up to the superfluid core). In the starquake model, $Q$ is defined as the fractional moment of inertia, i.e. the ratio of the moment of inertia of the superfluid core, $I_{\\rm core}$, to the moment of inertia of the entire configuration, $I_{\\rm total}$, as (Pines et al 1974) \\begin{equation} Q = \\frac{I_{\\rm core}}{I_{\\rm total}}. \\end{equation} Recently, Crawford \\& Demia\\'{n}ski (2003) have collected the all measured values of the glitch healing parameter $Q$ for Crab and Vela pulsars available in the literature and found that for 21 measured values of $Q$ for Crab glitches, a weighted mean of the values yields $Q = 0.72 \\pm 0.05$, and the range of $Q \\geq 0.7$ encompasses the observed distribution for the Crab pulsar. In order to test the starquake model for the Crab pulsar, they have computed $Q$ (as given by Eq.(1)) values for seven representative EOSs of dense nuclear matter, covering a range of neutron star masses. Their study shows that the much larger values of $Q(\\geq 0.7)$ for the Crab pulsar is fulfilled by all the six EOSs (out of seven considered in the study) corresponding to a `realistic' neutron star mass range $1.4\\pm 0.2M_\\odot$. By contrast, a weighted mean value of the 11 measurements for Vela yields a much smaller value of $Q(= 0.12 \\pm 0.07)$ and the all estimates for Vela agree with the likely range of $Q \\leq 0.2$. Thus, their results are found to be consistent with the starquake model predictions for the Crab pulsar. They have also concluded that the much smaller values of $Q \\leq 0.2$ for the Vela pulsar are inconsistent with the starquake model predictions, since the implied Vela mass based upon their models corresponds to a value $ \\leq 0.5M_\\odot$ for $Q \\leq 0.2$, which is too low as compared to the `realistic' NS mass range.  Thus, in the literature, the starquake is considered as a viable mechanism for glitch generation in the Crab and the vortex unpinning, the another mechanism, is considered suitable for the Vela pulsar, since it can avoid some other problems associated with the starquake explanation of the Vela glitches (see, e.g. Crawford \\& Demia\\'{n}ski (2003); and references therein). However, it seems surprising that if the internal structure of NSs are described by the same two component conventional  models (as mentioned above), different kinds of glitch mechanisms are required for the explanation of a glitch! Furthermore, it also follows from the above discussion that the main reason for not considering the starquake, the feasible mechanism for glitch generation in the Vela, lies in the fact that there exists none of the sequence of NS models in the literature which could  explain simultaneously, on the basis of starquake model, both the extreme limiting cases of glitch healing parameter, $Q$, corresponding to the Vela ($Q \\leq 0.2$) and the Crab ($Q \\geq 0.7$) pulsars in the range $0 \\leq Q \\leq 1.0$  for the `realistic' NS mass values for both the pulsars. The present study, therefore, deals with the construction of such models \\footnote{however, the other problems associated with the starquake explanation of the Vela glitches (see, e.g. Crawford \\& Demia\\'{n}ski (2003); and references therein) are not considered in the present paper. The future study in this regard may provide some explanation, provided the correlation between various parameters of the Crab and the Vela pulsar, obtained in the present study, can be utilized.}   We assume that {\\em all} the NSs belong to the same family of NS sequence which terminates at the {\\em maximum} value of mass. Certainly, this {\\em maxima} should correspond to an {\\em upper bound} on NS masses. In order to construct such a sequence, we have to set the extreme causal EOS (in geometrized units), $dP/dE = 1$ (where $P$ is the pressure and $E$ the energy-density) to describe the core. Firstly, because various observational studies like - the gamma-ray burst data, X-ray burst data and the glitch data etc., and their explanation (see, e.g. Lindblom 1984; Cottam et al 2002; Datta \\& Alpar 1993) favour the stiffest EOSs. The latest estimate of the moment of inertia for the Crab pulsar (based upon the `newest' observational data on the Crab nebula mass) rules out most of the existing EOSs of the dense nuclear matter, leaving only the stiffest ones (Bejger \\& Haensel 2002; Haensel et al 2006). Secondly, because of the fact that the `real' EOS of the dense nuclear matter beyond the density range $\\sim 10^{14}\\,{\\rm g\\, cm}^{-3}$ are largely unknown due to the lack of knowledge of nuclear interactions (see, e.g. Dolan 1992; and references therein; Haensel et al 2006), and the various EOSs available in the literature (see, e.g. Arnett \\& Bowers 1977) for NS matter represent only an extrapolation of the results far beyond this density range. Though, the status of the `real' EOS for NS matter is not certain, one could impose some well-known physical principle, independent of the EOS, such as the `causality condition' ($dP/dE = 1$) throughout the core of the star beyond a fiduciary density, $E_b$, at the core-envelope boundary to ascertain a definite upper bound on NS masses (see, e.g., Rhoades \\& Ruffini 1974; Hartle 1978; Lindblom 1984; Friedman \\& Ipser 1987; Kalogera \\& Baym 1996). In this connection this is also to be pointed out here that the maximum mass for {\\em any} EOS describing the core, beyond the density $E_b$,  with a subluminal sound velocity turns out to be less than that of the upper bound obtained by using the extreme causal EOS (see, e.g., Haensel et al 2006).  The envelope of our models (below the density $E_b$ at the core-envelope boundary) may be characterized by the well-known EOS of classical polytrope ${\\rm d}$ln$P/{\\rm d}$ln$\\rho = \\Gamma_1$ (where $\\rho$ denotes the density of the rest-mass and $\\Gamma_1$ is a constant known as the adiabatic index) for different values of the constant $\\Gamma_1 = (4/3), (5/3)$ and 2 respectively. The reason for considering the polytropic EOS for the entire envelope lies in the fact that with this EOS, our models yield an upper bound on NS masses {\\em independent} of the value of $\\Gamma_1$, and this upper bound (for a fiduciary choice of $E_b$) is found fully consistent with those of the values cited in the literature (Kalogera \\& Baym 1996; Friedman \\& Ipser 1987). Thus, the choice of the said polytropic EOS for the entire envelope may be regarded entirely equivalent to the choice of the various EOSs like WFF (Wiringa, Fiks \\& Fabrocini 1988), FPS (Lorenz, Ravenhall \\& Pethick, 1993), NV (Negele \\& Vautherin 1973), or BPS (Baym, Pethick \\& Sutherland 1971) in an appropriate sequence below the density range $E_b$, adopted by various authors in the conventional models of NSs (see, e.g., Kalogera \\& Baym 1996; Friedman \\& Ipser 1987), so that the constant $\\Gamma_1$ appearing in the polytropic EOS may be looked upon as an `average' $\\Gamma_1$ for the density range below $E_b$, specified by the sequence of various EOSs in the conventional models of NSs. The choice of the constant $\\Gamma_1 = 4/3, 5/3$ and 2 thus become obvious, since this choice can cover almost the entire range of density discussed in the literature for NS matter which is also applicable for the envelope region - the polytropic EOS with $\\Gamma_1 = 4/3$ represents the EOS of extreme relativistic degenerate electrons and non-relativistic nuclei (Chandrasekhar 1935), $\\Gamma_1 = (5/3)$ represents the well-known EOS of non-relativistic degenerate `neutron gas' (Oppenheimer \\& Volkoff 1939), and $\\Gamma_1 = 2$ represents the case of extreme relativistic baryons interacting through a vector meson field (Zeldovich 1962) (The value of $\\Gamma_1  > 2$ is also  possible for some EOS describing the NS matter, e.g.,  Malone,  Johnson \\&  Bethe  1975;  Clark, Heintzmann \\& Grewing 1971, however, the  results  obtained  in this paper remain unaffected for the choice of $\\Gamma_1  > 2$), and the outcome of this study (in terms of explaining the glitch healing parameter for various pulsars and predicting the upper bound on the compactness of NSs (since the upper bound on mass is independent of the value of $\\Gamma_1$)) would finally decide, among the chosen values, the `appropriate' value of $\\Gamma_1$ for the NS envelope. The validity of assuming the extreme causal EOS in the core and a polytropic EOS in the envelope of the present models, in view of the various modern EOS of dense nuclear matter, is also discussed in the last section of the present paper.  We have noted that in all conservative models of NSs, the choice of the core-envelope boundary, $r_b$ (corresponding to a density denoted by $E_b$), is somewhat {\\em arbitrary} in the sense that there are no criteria available for the choice of a particular matching density, $E_b$, below which the EOS of the NS matter is assumed to be known and unique. One can freely choose somewhat lower values of $E_b$ (which will increase the core size) to obtain higher values of $Q$ (see, e.g. Shapiro \\& Teukolsky 1983; Datta \\& Alpar 1993). To avoid such a procedure, we choose the core-envelope boundary of our models on the basis of the `compatibility criterion' which asserts that for an assigned value of the ratio ($\\sigma$) of central pressure, $P_0$, to central energy-density, $E_0$, the compactness parameter $u(\\equiv M/R$; total mass to radius ratio in geometrized units) of any {\\em regular} configuration should not exceed the compactness parameter $u_h$ of the homogeneous density sphere, in order to assure the compatibility with the hydrostatic equilibrium (Negi \\& Durgapal 2001; Negi 2004a). This criterion is capable of constraining the core-envelope boundary of any physically realistic NS model. A combination of this criterion with those of the observational data on the glitch healing parameter and the recently estimated minimum value of the moment of inertia for the Crab pulsar (based on the newly estimated `central value' of the Crab nebula mass $M {\\rm (nebula)} \\simeq 4.6 M_\\odot$ in the time-dependent acceleration model), $I_{\\rm Crab,45} = 1.93$; where $I_{\\rm Crab,45} = I_{\\rm Crab}/10^{45}$ g\\,cm$^2$ (Bejger \\& Haensel 2003) can provide the desired NS models discussed above, since both the theory (criterion) and the observations (stated above) are being used to construct the NS models. ",
        "conclusions": "This study constructs the stable sequences of NS models terminate at the value of maximum mass, $M_{\\rm max} \\simeq 2.22 M_\\odot$, independent of the EOSs of the envelope, for the matching density, $E_b = 9.584 \\times 10^{14}$ g\\, cm$^{-3}$, at the core-envelope boundary. This value of `matching density' is a consequence of the observational constraints $Q \\simeq 0.7$ and $I_{\\rm Crab,45} \\simeq 1.93$ (associated with the Crab pulsar) imposed together on the $\\Gamma_1 = 2$ envelope model and in this sense does not represent a fiduciary quantity. The upper bound of the surface redshift, $z_R \\simeq 0.77$ (corresponding to a $u$ value $\\simeq 0.34$), however, belongs to the model with a $\\Gamma_1 = 2$ envelope which is consistent with the absolute upper bound on the surface redshift of NS models compatible with causality and pulsational stability (Negi 2004b). This special feature, together with some other remarkable ones, discussed in the present study underline the appropriateness of the $\\Gamma_1 = 2$ envelope model.  Since among the variety of modern EOSs discussed in the literature, the upper bound on NS mass compatible with causality and dynamical stability can reach a value up to $2.2M_\\odot$ (in this category, the SLy (Douchin \\& Haensel 2001) EOS yields a maximum mass of $2.05M_\\odot$, whereas the BGN1 (Balberg \\& Gal 1997) and the APR (Akmal et. al. 1998) EOSs yield the maximum masses of $2.18M_\\odot$ and $2.21M_\\odot$ respectively (see, e.g. Haensel et al 2006)). In view of this result, the model-independent maximum mass, $M_{\\rm max} \\simeq 2.22M_\\odot$, obtained in this study may be regarded as good as those obtained on the basis of modern nuclear theory.   In addition to this result, the appropriate sequences of stable NS models obtained in this study can explain the glitch healing parameter, $Q$, of any glitching pulsar, provided the weighted mean values of $Q$ lie in the range $0 < Q \\leq 0.78$. This finding also reveals that if the starquake is considered to be a viable mechanism for glitch generation in all pulsars, then the envelope of `real' NSs may be well approximated by a polytropic EOS corresponding to a polytropic index, $n$, closer to 1.  For the value of matching density, $E_b = 9.584 \\times 10^{14}$ g\\, cm$^{-3}$, the $\\Gamma_1 = 2$ envelope model yields the minimum values of mass $M \\simeq 2.11 M_\\odot$ and surface redshift $z_R \\simeq 0.6232$ for the Crab ($Q \\simeq 0.7$) and the maximum values of mass $M \\simeq 0.982 M_\\odot$ and surface redshift $z_R \\simeq 0.2016$ for the Vela pulsar ($Q \\simeq 0.2$). The minimum mass and surface redshift for the Crab pulsar are slightly increased up to the values $M \\simeq 2.149 M_\\odot$ and $z_R \\simeq 0.655$ respectively, if the `central' weighted mean value of $Q \\approx 0.72$ and the moment of inertia $I_{\\rm Crab,45} > 1.93 $ are also imposed on these models. However, for the `central' weighted mean value of $Q \\simeq 0.12$ corresponding to the Vela pulsar, the maximum mass and surface redshift are somewhat decreased to the values $M \\simeq 0.828 M_\\odot$ and $z_R \\simeq 0.1645$ respectively. This value of mass and surface redshift for the Vela pulsar can further decrease up to the values $M \\simeq 0.685 M_\\odot$ and $z_R \\simeq 0.1312$ respectively, if the lower weighted mean value of $Q \\simeq 0.05$ for the Vela pulsar is imposed. These results predict the upper and lower bounds on the energy of a gravitationally redshifted electron-positron annihilation line in the range of about 0.309 - 0.315 MeV from the Crab and in the range of about 0.425 - 0.439 MeV from the Vela pulsar respectively.  For a comparison, if the observational constraint of the minimum value of $I_{\\rm Crab,45} \\simeq 3.04$ (the value of moment of inertia for the Crab pulsar obtained earlier by Bejger \\& Haensel (2002), on the basis of the constant-acceleration model for the Crab nebula) together with $Q \\simeq 0.7$ is imposed on the models studied in the present paper, the $\\Gamma_1 = 2$ envelope model yields the value of matching density, $E_b = 7.0794 \\times 10^{14}$ g\\, cm$^{-3}$. This value of $E_b$ yields a model-independent upper bound on NS mass $M_{\\rm max} \\simeq 2.59 M_\\odot$. This value of maximum mass, however, represents an `average' of the maximum NS masses in the range $2.2M_\\odot \\leq M_{\\rm max} \\leq 2.9M_\\odot$ obtained by Kalogera \\& Baym 1996 (and references therein) on the basis of other EOSs for NS matter, fitted to experimental nucleon-nucleon scattering data and the properties of light nuclei. For this lower value of matching density, the $\\Gamma_1 = 2$ envelope models yield the minimum value of mass $M \\simeq 2.455 M_\\odot$ for the Crab ($Q \\simeq 0.7$) and the maximum value of mass $M \\simeq 1.142 M_\\odot$ for the Vela pulsar ($Q \\simeq 0.2$). The minimum mass for the Crab pulsar is slightly increased up to the value $M \\simeq 2.5 M_\\odot$, if the `central' weighted mean value of $Q \\approx 0.72$ and the moment of inertia $I_{\\rm Crab,45} > 3.04 $ are also imposed on these models. However, corresponding to the `central' weighted mean value of $Q \\simeq 0.12$, the maximum mass of the Vela pulsar is somewhat decreased to the value $M \\simeq 0.964 M_\\odot$. This value of mass for the Vela pulsar can further decrease up to the value $M \\simeq 0.796 M_\\odot$, if the lower weighted mean value of $Q \\simeq 0.05$ for the Vela pulsar is imposed.  Furthermore, the study can also explain some special features associated with the extraordinary gamma-ray burst of 5 March 1979."
    },
    "0710/0710.4168_arXiv.txt": {
        "abstract": "We have applied the torus fitting procedure described in Ng \\& Romani (2004) to PWNe observations in the \\emph{Chandra} data archive. This study provides quantitative measurement of the PWN geometry and we characterize the uncertainties in the fits, with statistical errors coming from the fit uncertainties and systematic errors estimated by varying the assumed fitting model. The symmetry axis $\\Psi$ of the PWN are generally well determined, and highly model-independent. We often derive a robust value for the spin inclination $\\zeta$. We briefly discuss the utility of these results in comparison with new radio and high energy pulse measurements. ",
        "introduction": "One of the greatest success of the Chandra X-ray Observatory (CXO) is the discovery of equatorial tori and polar jet structures in many pulsar wind nebula (PWN) systems. It is now believe that these features are common among young neutron stars. In the \\citet{ree74} picture, when the highly relativistic pulsar wind decelerates in the external medium, a termination shock is formed at a characteristic scale \\[r_t =\\left(\\frac{\\dot E}{4\\pi c \\eta P_{\\rm ext}} \\right )^{1/2} \\ , \\] where $\\dot E$ is the pulsar spin-down power and $\\eta$ is the filling factor. In general, if the pulsar is subsonic in the ambient medium, i.e. $ P_{\\rm ext} \\geq P_{\\rm ram}= 6\\times 10^{-10}nv^2_7\\;\\mathrm {g\\;cm^{-1}s^{-2}}$ for a pulsar speed $10^7v_7\\;\\mathrm{cm\\;s^{-1}}$ through a density of $n\\,m_p\\;\\mathrm{cm^{-3}}$, a toroidal shock structure is expected; faster objects produce bow shock nebulae. Indeed, many young pulsars still inside their high pressure supernova remnant birth sites do show such toroidal symmetry. The best-known example is the PWN around the \\object{Crab pulsar} as observed by the \\emph{CXO} \\citep{wei01}. Recently, several relativistic MHD models, e.g.\\ \\citet{kom03} and \\citet{del06}, have shown how such toroidal structure can form if the pulsar wind has a latitudinal variation.  \\citet{ng04} (hereafter \\citetalias{ng04}) developed a fitting procedure to measure the 3D orientation of the pulsar wind torus and applied to a few X-ray observations. While this simple geometrical model does not capture the fine details of the MHD simulations, is does allow one to extract the torus (and hence pulsar spin) orientation from relatively low signal-to-noise ratio (S/N) data. \\citetalias{ng04} also gave quantitative estimates for the statistical errors arising from Poisson statistics. However the systematic errors due to unmodeled components such as jets or background were neglected. For bright objects e.g.\\ the Crab and Vela pulsars, the S/N is high and such systematic errors dominate. In this study, we attempt to quantify these systematic errors and apply the fitting to more \\emph{CXO} PWN observations, thus providing a more comprehensive study. ",
        "conclusions": "In conclusion, we have applied the torus fitting technique to more PWN observations in the \\emph{Chandra} data archive and characterized the uncertainties in the fits. This study provides a better understanding of the systematic errors, giving quantitative estimates of the measurement uncertainties. We argue that these robust position angle $\\Psi$ and inclination $\\zeta$ values are particularly useful for comparison with the radio and high energy pulse data. If new observations can fill in more measurements from these energy bands in Table 3, we should be able to make substantial progress in understanding the emission zones and viewing geometries of young pulsars."
    },
    "0710/0710.1860_arXiv.txt": {
        "abstract": "We present a comprehensive analysis of structure in the young, embedded cluster, NGC~1333 using members identified with {\\it Spitzer} and 2MASS photometry based on their IR-excess emission. In total, 137 members are identified in this way, composed of 39 protostars and 98 more evolved pre-main sequence stars with disks. Of the latter class, four are transition/debris disk candidates. The fraction of exposed pre-main sequence stars with disks is $83\\% \\pm 11\\%$, showing that there is a measurable diskless pre-main sequence population. The sources in each of the Class~I and Class~II evolutionary states are shown to have very different spatial distributions relative to the distribution of the dense gas in their natal cloud. However, the distribution of nearest neighbor spacings among these two groups of sources are found to be quite similar, with a strong peak at spacings of 0.045~pc. Radial and azimuthal density profiles and surface density maps computed from the identified YSOs show that NGC~1333 is elongated and not strongly centrally concentrated, confirming previous claims in the literature. We interpret these new results as signs of a low velocity dispersion, extremely young cluster that is not in virial equilibrium. ",
        "introduction": "Observations of embedded, star-forming clusters and groups show that the stellar distributions are often elongated, clumpy, or both \\citep[cf.][]{ll03,gute05,teix06,alle07}, and the structure seems tied to the distribution of dense gas in the clusters' natal molecular clouds. Most clouds have some non-spherical structure, yet current results suggest that the associated clusters have varying degrees of agreement with the cloud's structure depending on how deeply the members are embedded \\citep{gute05,teix06}. Given the high frequency of asymmetric structure in clouds, it seems reasonable to assume that the exposed, relatively structureless clusters we observe may have been asymmetric in a previous epoch. With the ejection of the majority of their natal gas and adequate time to migrate from their birth sites, the imprint of the underlying cloud structure could be lost rather quickly in a recently exposed young cluster. Thus the structure we measure in a distribution of young stellar objects (YSO) relative to the dense gas in the associated cloud may be a reasonable proxy for the current dynamical state of an embedded cluster.   The launch of {\\it Spitzer} has provided a potent new facility for studies of the structure of young and embedded clusters. While some members of these young clusters are diskless pre-main sequence stars (Class~III), the majority of the membership are sources with excess emission at mid-IR wavelengths \\citep{hll01}, made up of a mixture of protostars (Class~I), still embedded and accreting from dense spherical envelopes, and the slightly more evolved pre-main sequence stars with circumstellar disks (Class~II). The mid-IR spectral energy distributions (SED) of these objects are dominated by the emission from their dusty circumstellar material, making them easily distinguishable from pure photospheric sources such as unrelated field stars and indistinguishable diskless cluster members. {\\it Spitzer}'s sensitivity to mid-IR emission makes it the best tool currently available for identifying {\\it and characterizing} YSOs with IR-excess, and that sensitivity is sufficient to detect these sources down to the Hydrogen--burning mass limit for regions within the nearest kiloparsec \\citep{gute04}. The complete census of YSOs with disks in a young cluster represents a high-confidence sample of bona fide cluster members, and for the youngest regions, such a sample includes a high fraction of the total number of members \\citep{hll01}. Furthermore, by separating the two canonical evolutionary classes of YSOs that {\\it Spitzer} so effectively detects, we are able to probe both the recent overall star-forming activity in the region as traced by the stars with disks, and the immediate star formation as represented by the protostars, since this phase is expected to be short-lived relative to the former.  Many young clusters are dominated by heavy and spatially variable extinction from dust within their natal molecular cloud environment. Another advantage {\\it Spitzer} brings to studies of these regions is that the mid-IR wavelengths targeted by {\\it Spitzer}'s imaging instruments (Infrared Array Camera, or IRAC, at 3.6-8.0~$\\mu$m and Multiband Imaging Photometer for {\\it Spitzer}, or MIPS, at 24-160~$\\mu$m) are less affected by extinction from dust in comparison to near-IR or visible wavelengths \\citep[e.g.][]{ccm89}. Recent papers have done an excellent job of characterizing the reddening law in the IRAC bandpasses \\citep{inde05,flah07,huar07} and have made a first attempt at doing the same for the 24~$\\mu$m channel of MIPS \\citep{flah07,huar07}. Given the wealth of information {\\it Spitzer} can bring to the study of embedded clusters, we have surveyed over thirty clustered star-forming regions within the nearest kiloparsec through the Guaranteed Time Observations (GTO) program for the IRAC and MIPS instrument teams. With these data, we can not only provide a full census of the YSOs with IR-excess in each region, but also perform a detailed examination and comparative analysis of structure in young, embedded clusters.  NGC~1333 has been a popular target for observations of deeply embedded protostars via radio wavelengths due to the presence of an unprecendented number of molecular outflows \\citep[e.g.][]{ks00} associated with several bright IRAS sources considered Class~0 protostars \\citep{jenn87}, all in relatively close proximity to the Sun \\citep[250~pc,][]{enoc06}. Many of the outflows are traced by shock-induced emission, a clear sign that they are indeed affecting the local, quiescent cloud material \\citep[e.g.][]{wala05}. Because of this, some studies have claimed that the NGC~1333 dense molecular cloud core is in the process of being destroyed by influence from outflows \\citep[e.g.][]{wari96,sk01,quil05}. In addition to the numerous protostars, there is a cluster of pre-main sequence stars identified by number counts analysis in the near-IR \\citep{svs76,asr94,lal96,wilk04}. \\citet{lal96} suggested that the distribution was well-described as a ``double cluster'', having two distinct surface density maxima.      Here we present a {\\it Spitzer} IRAC and MIPS imaging and photometric analysis of the NGC~1333 young cluster, one of the most nearby large membership ($N>100$) clusters in the {\\it Spitzer} Young Cluster Survey. We achieve a near-complete census of the cluster membership that possesses circumstellar material, significantly surpassing the sensitivity of ground-based mid-IR surveys \\citep[cf.][and references therein.]{rebu03}. In addition, we statistically infer the population of pre-main sequence stars that lack disks, enabling an estimate of the fraction of members with disks. Finally, using the identified YSOs, we apply both established and recently developed methods for characterizing the structure of this cluster, in specific reference to previous claims of structure found in the literature. ",
        "conclusions": "We have employed the mid-IR sensitivity of {\\it Spitzer} to achieve a census of the YSO members of the NGC~1333 embedded cluster that was previously unattainable. Furthermore, we have shown that the sources identified in our {\\it Spitzer} census represent a large fraction of the total cluster membership (83\\%). With this penetrating view of the cluster, we have performed several measurements of different aspects of the structure of the cluster that are not confused by field star contamination and are less biased by extinction effects than previous studies.  We confirm the double-peaked surface density morphology of the cluster reported in previous work \\citep{lal96}, with the caveat that this morphology is traced only by the more evolved Class~II population. The two main density peaks are located in local minima of the gas density distribution. In contrast, the protostars trace the dense gas in this region closely, as do a fraction of the Class~II. That gas distribution appears to be a network of filamentary structures, connecting the two density peaks into a single, albeit complex, structure. Despite the difference in spatial distributions, however, the Class~II and protostars have similarities in their nearest neighbor distance distributions, particularly in the location of the peak at spacings of 0.045~pc. These two evolutionary states are expected to differ in duration by an order of magnitude, thus a low overall velocity dispersion in the cluster and a very young Class~II population are needed to account for the similarity remaining in the spacings among the two populations.   A further implication of the lack of evidence for dynamical evolution of the cluster is the need for the dispersal of the molecular gas by the Class~II sources on very short timescales. Given that the more evolved YSOs have not moved significantly from their birthsites, regions dominated by more evolved sources where gas densities are preferentially low, such as the double-peaks in the stellar distribution here, were once locations of high gas density. Thus it seems plausible that groups of low-mass stars are able to disrupt their dense natal gas {\\it locally}. In this sense, we concur with previous work that has suggested that the active outflows in NGC~1333 are destroying the cloud \\citep{sk01,quil05}. With a virtual lack of low or high mass stars and dense gas in the center of the ring-like structure of the cloud and cluster core, we argue that this is in fact a primordial structure and not outflow-evacuated.  Considering all sources regardless of evolutionary class, the cluster is clearly elongated, an expected result given the double-peaked nature of the stellar surface density distribution and embedded nature of the cluster as a whole \\citep{gute05}. We have presented radial density profiles measured via several different methods, and all suggest a roughly uniform density distribution within a 0.3~pc radius, with a steep decline ($\\alpha=-3.3$) beyond. The flat central distribution and sharp radial decline in the surface density profile at radii larger than 0.3~pc also suggests a rather limited amount of dynamical interaction in this cluster. If we analyze this within the framework of dynamics-generated structures like those of King or EFF models, we have to choose extreme fitting parameters just to get marginal agreement with the measured profiles. The radial profile we have measured here has two clear regimes were a simple power law matchs the density profile well, and the sharp knee transition between them is poorly matched by either the King or EFF profiles. Furthermore, given the asymmetry of the NGC~1333 cluster and the presense of a considerable amount of structured, dense gas, the dynamical information inferred from either of these models is unlikely to be accurate. Cloud geometry is the most tenable cause for the structure we have observed. As such, the relative motions of the stars must be fairly slow and the Class~II population must be quite young in order to have preserved that structure into the Class~II evolutionary phase."
    },
    "0710/0710.3054_arXiv.txt": {
        "abstract": "We report on the spatial relationship between solar flares and coronal mass ejections (CMEs) observed during 1996-2005 inclusive.  We identified 496 flare-CME pairs considering limb flares (distance from central meridian $\\ge 45^\\circ$) with soft X-ray flare size $\\ge$ C3 level. The CMEs were detected by the Large Angle and Spectrometric Coronagraph (LASCO) on board the {\\it Solar and Heliospheric Observatory} ({\\it SOHO}). We investigated the flare positions with respect to the CME span for the events with X-class, M-class, and C-class flares separately. It is found that the most frequent flare site is at the center of the CME span for all the three classes, but that frequency is different for the different classes. Many X-class flares often lie at the center of the associated CME, while C-class flares widely spread to the outside of the CME span. The former is different from previous studies, which concluded that no preferred flare site exists. We compared our result with the previous studies and conclude that the long-term LASCO observation enabled us to obtain the detailed spatial relation between flares and CMEs. Our finding calls for a closer flare-CME relationship and supports eruption models typified by the CSHKP magnetic reconnection model. ",
        "introduction": "A solar flare is sudden flash of electromagnetic radiation (suggesting plasma heating) in the solar atmosphere, and a coronal mass ejection (CME) is an eruption of the atmospheric plasma into interplanetary space. Both phenomena are thought to be different manifestations of the same process which releases magnetic free energy stored in the solar atmosphere. The spatial relation between flares and CMEs contains information on the magnetic field configurations involved in the eruptive process and hence is important for modeling them. Many flare-CME models are based on the CSHKP (Carmichael, Sturrock, Hirayama, Kopp \\& Pneuman) magnetic reconnection model. The model requires that a flare occurs just underneath of an erupting filament which eventually becomes the core of the CME associated with the flare. Normally the core corresponds to the center of the CME, thus the CSHKP model requires that the flare occurs near the center of the CME span.  Full-scale studies on the flare-CME relationship started in the 70s and 80s with the CME observations obtained by the {\\it Solwind} coronagraph on board {\\it P78-1} and the Coronagraph/Polarimeter telescope on board the {\\it Solar Maximum Mission} ({\\it SMM}). \\citet{harri86} carried out a detailed analysis of three flare-CME events observed by {\\it SMM} and reported that flares occurred near one foot of an X-ray arch, which is supposed to become a CME. He also analyzed 48 flare-CME events observed by {\\it SMM} and {\\it Solwind} and reported that many flares occurred near one leg of the associated CMEs. This result, called the flare-ejection asymmetry, is inconsistent with the CSHKP flare-CME model. \\citet{kahle89} examined 35 events observed by the {\\it Solwind} and reported that flare positions did not peak neither at the center nor at one leg of the CMEs. They concurred with Harrison at the point that the observations do not match with the CSHKP model, while disagreeing with the result that flares are likely to occur at one leg of CMEs.  They pointed out that the parameter employed by Harrison was biased, and concluded that both observations are compatible with the fact that there is no preferred flare site with respect to the CME span.  It should be noted that the two studies applied different criteria for the event selection. Harrison did not apply any criteria on flare X-ray intensity, flare location, and CME span, while Kahler et al. used only strong limb flares ($\\ge$ M1 level; central meridian distance (CMD) $\\ge 40^\\circ$) and wide CMEs (angular span $\\ge 40^\\circ$). Different criteria might produce different spatial distributions, but the results in both the studies were inconsistent with the schematic view of the CSHKP type flare-CME model.  The Large Angle and Spectrometric Coronagraph \\citep[LASCO;][]{bruec95} on board the {\\it Solar and Heliospheric Observatory} ({\\it SOHO}) mission has observed more than 11,000 CMEs from 1996, which provides a great opportunity to investigate the flare-CME relationship. \\citet{harri06} reviewed several flare-CME studies and stated that \"the pre-SOHO conclusions about relative flare-CME locations and asymmetry are consistent with many recent studies.\" However, systematic statistical study is needed before reaching a firm conclusion. In this paper we revisit this issue using the large CME data obtained by {\\it SOHO} LASCO.  \\begin{figure*} \\epsscale{0.90} \\plotone{f1.eps} \\caption{Three CMEs observed by {\\it SOHO} LASCO to illustrate the measurement of CME span. The top row shows direct images used to measure the main CME body, and the bottom row shows corresponding running difference images used to measure the whole CMEs. $\\phi_1$ and $\\phi_2$ indicate the PAs of side edges of the main CME body, and $\\phi_A$ and $\\phi_B$ indicate those of the whole CME. Arrows point to the position of the flares associated with the CMEs.} \\end{figure*} ",
        "conclusions": ""
    },
    "0710/0710.1098_arXiv.txt": {
        "abstract": "% A planetary transit produces both a photometric signal and a spectroscopic signal. Precise observations of the transit light curve reveal the planetary radius and allow a search for timing anomalies caused by satellites or additional planets. Precise measurements of the stellar Doppler shift throughout a transit (the Rossiter-McLaughlin effect) place a lower bound on the stellar obliquity, which may be indicative of the planet's migration history. I review recent results of the Transit Light Curve project, and of a parallel effort to measure the Rossiter effect for many of the known transiting planets. ",
        "introduction": "%  I have great admiration for the people who discover transiting planets. Identifying the candidate transit signals from among a hundred thousand light curves, and flushing out the numerous astrophysical false positives, are impressive feats. This article, however, is not about transit discovery, but rather about the next step: performing high-precision photometry and spectroscopy of exoplanetary transits. The goal of this step is to determine the planetary and stellar properties well enough to allow for meaningful comparisons with the familiar properties of the Solar system, and to inform our theories of planet formation.  The most immediate result of transit photometry is a measurement of the planetary radius. In combination with the planetary mass, which can be inferred from the Doppler orbit of the star, these data give the first clues about the composition, interior structure, and atmospheric energy balance of the planet. An accurate radius is also needed to interpret the results of other observations, such as the detection of thermal emission or reflected light based on secondary-transit photometry. The timings of the transits can be used to refine the measurement of the orbital period and search for additional bodies in the system. In \\S~2, I describe the Transit Light Curve (TLC) project, an effort to gather high-precision photometry during exoplanetary transits.  The most prominent spectroscopic signal during a transit is the Rossiter-McLaughlin (RM) effect. This effect is an anomalous Doppler shift that arises from stellar rotation. Measuring this effect allows one to assess the alignment between the planetary orbital axis and the stellar spin axis, a fundamental system property that provides clues about the process of planet migration. I describe some recent measurements of the RM effect in \\S~3.  \\begin{figure}[!h] \\plotone{Lightcurves.eps} \\caption{{\\bf The Transit Light Curve project.} Each panel shows time-binned photometry from a campaign (2-5 separate transits) for a particular planet.  Most of the data are $z$ band observations with the FLWO 1.2m telescope and Keplercam detector.} \\end{figure} ",
        "conclusions": ""
    },
    "0710/0710.0866_arXiv.txt": {
        "abstract": "We discuss oxygen and iron abundance patterns in K and M red-giant members of the Galactic bulge and in the young and massive M-type stars inhabiting the very center of the Milky Way. The abundance results from the different bulge studies in the literature, both in the optical and the infrared, indicate that the [O/Fe]-[Fe/H] relation in the bulge does not follow the disk relation, with [O/Fe] values falling above those of the disk. Based on these elevated values of [O/Fe] extending to large Fe abundances, it is suggested that the bulge underwent a rapid chemical enrichment with perhaps a top-heavy initial mass function. The Galactic Center stars reveal a nearly uniform and slightly elevated (relative to solar) iron abundance for a studied sample which is composed of 10 red giants and supergiants.  Perhaps of more significance is the fact that the young Galactic Center M-type stars show abundance patterns that are reminiscent of those observed for the bulge population and contain enhanced abundance ratios of $\\alpha$-elements relative to either the Sun or Milky Way disk at near-solar metallicities. ",
        "introduction": "Abundance patterns in different populations in the Milky Way can shed light on Galaxy formation and its chemical evolution. Studies carried out over the last decade have provided accurate abundance patterns for stars in the Milky Way disk and halo so that these populations are now fairly-well mapped. There is significantly less information, however, on the Galactic bulge population and, in particular the Galactic Center, due to difficulties associated with heavy extinction.  The stellar content belonging to the bulge population is old, with an age of 10--12 Gyr (e.g. Zoccali et al. 2003) and chemical abundance studies in both low and high-resolution find that the bulge is overall metal-rich with a large abundance spread that spans $\\sim$1.5 dex (e.g. Fulbright et al. 2006). The Galactic Center, on the other hand, contains a significant young population, including many luminous and massive stars. Those known to date are concentrated in three clusters within 60 pc of the Galactic Center: the Central Cluster, the Arches Cluster and the Quintuplet Cluster. The youngest stars in these clusters were formed recently with ages ranging roughly between 1 -- 9 Myr. In this contribution we summarize the abundance results obtained recently for the bulge and Galactic center and briefly discuss their implications in light of Galactic chemical evolution. ",
        "conclusions": "Abundance patterns are discussed in two distinct populations: $\\sim$ 12 Gyr old red-giants from the bulge and the young and massive M-type giants and supergiants residing mostly within 2 pcs of the central Black Hole. The abundance results indicate that the iron abundance distribution for the admittedly small sample of Galactic Center targets studied so far shows very little abundance spread in contrast to the large metallicity spread found for the old bulge population. Most importantly, both populations show enhancements in the $\\alpha$-element abundances  relative to the Galactic disk trend that might be explained invoking a top-heavy IMF."
    },
    "0710/0710.2114_arXiv.txt": {
        "abstract": "{ We present new determination of the birth rate of AXPs and SGRS and their associated SNRs. We find a high birth rate of $1/(500\\ {\\rm yr})$ to $1/(1000\\ {\\rm yr})$  for AXPs/SGRs and their associated SNRs. These high rates suggest that all massive stars (greater than $\\sim (23$-$32) M_{\\odot}$) give rise to remnants with magnetar-like fields. Observations indicate a limited fraction of high magnetic fields in these progenitors thus our study necessarily implies  magnetic field amplification. Dynamo mechanisms during the birth  of the neutron stars require spin rates much faster than either observations or theory indicate. Here, we propose that neutron stars form with normal ($\\sim 10^{12}$ G) magnetic fields, which are then amplified to $10^{14}$-$10^{15}$ G after a delay of  hundreds of years. The  amplification is speculated to be a consequence of color ferromagnetism and to occur  after  the neutron star core reaches quark-deconfinement density.  This delayed amplification alleviates many  difficulties in interpreting simultaneously the high birth rate,  high magnetic fields, and state of isolation of AXPs/SGRs and their link to massive stars. ",
        "introduction": "Early studies of association of Anomalous X-ray Pulsars (AXPs) with  supernova remnants (SNRs) suggested that  5\\% of core-collapse SN results in AXPs (Gaensler et al. 1999). This was based on 3 SNR associations out of a total of 6 AXPs. Since then evidence has mounted that AXPs and soft gamma-ray repeaters (SGRs) are the same type of objects (Gavriil et al. 2002) and more AXPs, SGRs and associated SNRs have been identified. Thus it it timely to revisit the issue of AXPs/SGRs birthrates.  In this study  we present an updated investigation of the birth rate of AXPs/SGRs and in addition, for the first time,  the birth rate of associated SNRs is given. Since AXPs/SGRs ages rely on spin-down age estimates whereas SNRs ages are based  on shock expansion models,  this constitutes two independent estimates for birth rates.    Both samples yield a high birth rate for AXPs/SGRs\\footnote{An independent study by Gill\\&Heyl (2007), based on  a population synthesis of AXPs detected in the ROSAT All-Sky Survey, yields a birth rate of $\\sim 0.22$ per century.} of $(1/5)$-$(1/10)$ of all core-collapse SNe, higher than previously appreciated. This high frequency of occurrence of AXPs/SGRs brings into focus issues related to the origin of the strong magnetic fields which we address here. This paper is presented as follows: \\S 2 describes the methods and presents the birth rate results, and \\S 3  discusses the implications. Our model, based on a delayed amplification of magnetic field, is presented in \\S 4 before concluding in \\S 5. ",
        "conclusions": "Our study of  the birth rate of AXPs and SGRS  and their associated SNRs suggest that  about $1/5$ to $1/10$ of  all core-collapse SN lead to AXPs/SGRs. These high rates suggest that all massive stars (greater than $ M_{\\rm low}$) give rise to remnants with magnetar-like fields.  This raises these issues: (i) how  do all progenitors  with $M\\ge M_{\\rm low}$ generate $>10^{14}$ G fields in their compact remnants?; (ii)  why is there a dichotomy in magnetic field strength between compact remnants from progenitors with mass greater than $M_{\\rm low}$ (i.e. $B\\sim 10^{14}$ G) and those with mass less  than $\\sim M_{\\rm low}$ ($B\\sim 10^{12}$ G); (iii) why are all AXPs/SGRs isolated while many progenitors with $M>M_{\\rm low}$ are in binaries?   In this study, we introduce the notion of delayed magnetic field amplification to resolve these issues. We propose that  neutron stars from  progenitor masses $M> 9M_{\\odot}$ are born  with normal ($\\sim 10^{12}$ G) magnetic fields. A neutron star from  a progenitor with an approximate mass range $M_{\\rm low}< M <60 M_{\\odot}$ will experience an explosive transition to  a quark star (the QN) in which its magnetic field  is  amplified to $10^{14}$-$10^{15}$ G by color ferromagnetism (Iwazaki 2005). The second explosion (QN) and related mass loss helps to reduce the  surviving compact binary fraction  thus explaining the state of isolation of AXPs/SGRs. The transition occurs with a delay of several hundred years (Staff et al. 2006). This delayed amplification alleviates many  difficulties in interpreting simultaneously the high birth rate and high magnetic fields of AXPs/SGRs and their link to massive stars."
    },
    "0710/0710.0920_arXiv.txt": {
        "abstract": "Electric currents $j$ flow along the open magnetic field lines from the polar caps of neutron stars. Activity of a polar cap depends on the ratio $\\alph=j/c\\rhoGJ$, where $\\rhoGJ$ is the corotation charge density. The customary assumption $\\alph\\approx 1$ is not supported by recent simulations of pulsar magnetospheres and we study polar caps with arbitrary $\\alph$. We argue that no significant activity is generated on field lines with $0<\\alph<1$. Charges are extracted from the star and flow along such field lines with low energies. By contrast, if $\\alph>1$ or $\\alph<0$, a high voltage is generated, leading to unsteady $e^\\pm$ discharge on a scale-height smaller than the size of the polar cap. The discharge can power observed pulsars. Voltage fluctuations in the discharge imply unsteady twisting of the open flux tube and generation of Alfv\\'en waves. These waves are ducted along the tube and converted to electromagnetic waves, providing a new mechanism for pulsar radiation. ",
        "introduction": "Corotation of a plasma magnetosphere is impossible beyond the light cylinder of a star, and magnetic field lines that extend beyond this cylinder are twisted, $\\nabla\\times\\bB\\neq 0$. Thus, currents $\\bj_B=(c/4\\pi)\\nabla\\times\\bB$ are induced in the open magnetic flux tubes that connect the star (its ``polar caps'') to the light cylinder (Sturrock 1971). These currents are approximately force-free and flow along the magnetic field $\\bB$. A basic question of pulsar theory is what voltage develops along the open tube to maintain these currents. It determines the dissipated power and $e^\\pm$ creation that feeds the observed activity of pulsars.  The customary pulsar model assumes that the electric current $\\jB$ extracted from the polar cap nearly matches $c\\rhoGJ$, where $\\rhoGJ=-{\\bf \\Omega}\\cdot\\bB/2\\pi c$ is the corotation charge density (Goldreich \\& Julian 1969). The deviation of current from $c\\rhoGJ$ was calculated as an eigen value of an electrostatic problem and found to be small (Arons \\& Scharlemann 1979). This is in conflict with recent global models of pulsar magnetospheres, which report $|\\jB-c\\rhoGJ|\\sim\\jB$ (e.g.  Contopoulos et al. 1999; Spitkovsky 2006; Timokhin 2006; see Arons 2008 for a review). A significant mismatch between $\\jB$ and $c\\rhoGJ$ can be expected on general grounds (Kennel et al. 1979). Currents $\\jB$ are determined by the magnetic-field twisting near the light cylinder, while $\\rhoGJ$ is a local quantity at the polar cap that is practically independent of $\\jB$.  The open tube is surrounded by the grounded closed magnetosphere\\footnote{ The closed magnetosphere with $\\jB=0$ is expected to have $\\rho=\\rhoGJ$ and $\\Epar\\approx 0$ (e.g. Arons 1979). } and may be thought of as a waveguide, filled with magnetized (1D) plasma. Compared to usual plasma-filled waveguides, it has two special features: (1) Current $\\jB$ is imposed on the tube. The twisted tube extends into the star, which is a good conductor and maintains $j_B$. ${\\rm sign}(\\jB)=\\pm 1$ corresponds to $\\pm$ charge flowing outward along the magnetic field lines. (2) Vacuum has effective charge density $\\rho_0=-\\rhoGJ$ as Gauss law in the rotating frame reads $\\nabla\\cdot\\bE=4\\pi(\\rho-\\rhoGJ)$. The key dimensionless parameter (which can vary along and across the tube) is \\be \\alph=\\frac{\\jB}{c\\rhoGJ}. \\ee In this Letter, we discuss basic properties of the polar-cap accelerator with arbitrary $\\alph$. First, we discuss what happens without $e^\\pm$ creation: \\S~\\ref{sec:slab} studies how the current is extracted from the polar cap of a radius $\\rpc$ and flows at small heights $z\\ll\\rpc$ (this region is called ``slab zone'' below), and \\S~\\ref{sec:tube} discusses how the flow extends to the region $z>\\rpc$ (``thin-tube zone''). We argue that the accelerator is inefficient if $0<\\alph<1$.  The value of $\\alph$ depends largely on the angle $\\chi$ between ${\\bf\\Omega}$ and $\\bB$ at the polar cap. For aligned dipole rotators ($\\chi=0$), magnetospheric models predict $\\alph<1$ everywhere on the polar cap and $\\alpha<0$ near its edge (see Fig.~4 and 5 in Timokhin 2006). For orthogonal rotator ($\\chi\\approx\\pi/2$), $|\\alph|\\gg 1$ throughout most of the polar cap.  Generally, the polar cap has three regions where $\\alph>1$, $0<\\alph<1$, and $\\alph<0$. We propose that pulsar activity originates in the polar-cap regions where $\\alph^{-1}<1$ (i.e. $\\alph>1$ or $\\alph<0$) as a high voltage is generated in these regions.  We emphasize that the voltage is generated because $\\nabla\\times\\bB\\neq 0$, not because $\\rhoGJ\\neq 0$. The accelerator works as well if $\\rhoGJ=0$ (i.e. if $\\bf\\Omega\\perp\\bB$ at the polar cap), which corresponds to $\\alph\\rightarrow \\pm\\infty$. The accelerator height $h\\simlt\\rpc$ is regulated by unsteady $e^\\pm$ discharges (\\S~\\ref{sec:discharge}). \\S~\\ref{sec:Alfven} describes a mechanism for radio emission from unsteady discharges. ",
        "conclusions": ""
    },
    "0710/0710.5734_arXiv.txt": {
        "abstract": "Big Bang nucleosynthesis requires a fine balance between equations of state for photons and relativistic fermions. Several corrections to equation of state parameters arise from classical and quantum physics, which are derived here from a canonical perspective. In particular, loop quantum gravity allows one to compute quantum gravity corrections for Maxwell and Dirac fields. Although the classical actions are very different, quantum corrections to the equation of state are remarkably similar.  To lowest order, these corrections take the form of an overall expansion-dependent multiplicative factor in the total density.  We use these results, along with the predictions of Big Bang nucleosynthesis, to place bounds on these corrections and especially the patch size of discrete quantum gravity states. ",
        "introduction": "\\label{sec:INTRODUCTION}  Much of cosmology is well-described by a space-time near a spatially isotropic Friedmann--Robertson--Walker models with line elements \\begin{equation} \\md s^2= -\\md\\tau^2+a(\\tau)^2\\left(\\frac{\\md r^2}{1-kr^2}+r^2(\\md\\vartheta^2+ \\sin^2\\vartheta\\md\\varphi^2)\\right) \\end{equation} where $k=0$ or $\\pm 1$, sourced by perfect fluids with equations of state $P=w\\rho$. Such an equation of state relates the matter pressure $P$ to its energy density $\\rho$ and captures the thermodynamical properties in a form relevant for isotropic space-times in general relativity. Often, one can assume the equation of state parameter $w$ to be constant during successive phases of the universe evolution, with sharp jumps between different phases such as $w=-1$ during inflation, followed by $w=\\frac{1}{3}$ during radiation domination and $w=1$ during matter domination.  Observationally relevant details can depend on the precise values of $w$ at a given stage, in particular if one uses an effective value describing a mixture of different matter components. For instance, during big bang nucleosynthesis one is in a radiation dominated phase mainly described by photons and relativistic fermions. Photons, according to Maxwell theory, have an exact equation of state parameter $w=\\frac{1}{3}$ as a consequence of conformal invariance of the equations of motion (such that the stress-energy tensor is trace-free). For fermions the general equation of state is more complicated and non-linear, but can in relativistic regimes be approximately given by the same value $w=\\frac{1}{3}$ as for photons. In contrast to the case of Maxwell theory, however, there is no strict symmetry such as conformal invariance which would prevent $w$ to take a different value. It is one of the main objectives of the present paper to discuss possible corrections to this value.  For big bang nucleosynthesis, it turns out, the balance between fermions and photons is quite sensitive. In fact, different values for the equation of state parameters might even be preferred phenomenologically \\cite{FermionBoson}. One possible reason for different equations of state could be different coupling constants of bosons and fermions to gravity, for which currently no underlying mechanism is known. In this paper we will explore the possibility whether quantum gravitational corrections to the equations of state can produce sufficiently different values for the equation of state parameters. In fact, since the fields are governed by different actions, one generally expects different, though small, correction terms which can be of significance in a delicate balance.  Note that we are not discussing ordinary quantum corrections of quantum fields on a classical background. Those are expected to be similar for fermions and radiation in relativistic regimes. We rather deal with quantum gravity corrections in the coupling of the fields to the space-time metric, about which much less is known a priori. Thus, different proposals of quantum gravity may differ at this stage, providing possible tests.  An approach where quantum gravitational corrections can be computed is loop quantum cosmology \\cite{LivRev}, which specializes loop quantum gravity \\cite{Rov,ALRev,ThomasRev} to cosmological regimes. In such a canonical quantization of gravity, equations of state must be computed from matter Hamiltonians rather than covariant stress-energy tensors. Quantum corrections to the underlying Hamiltonian then imply corrections in the equation of state. This program was carried out for the Maxwell Hamiltonian in \\cite{MaxwellEOS}, and is done here for Dirac fermions. There are several differences between the treatment of fermions and other fields, which from the gravitational point of view are mainly related to the fact that fermions, in a first order formulation, also couple to torsion and not just the curvature of space-time. After describing the classical derivation of equations of state as well as steps of a loop quantization and its correction terms, we use big bang nucleosynthesis constraints to see how sensitively we can bound quantum gravity parameters.  We have aimed to make the paper nearly self-contained and included some of the technical details. Secs.~\\ref{sec:canonical Formulation} on the canonical formulation of fermions, \\ref{sec:MODIFICATIONS} on quantum corrections from loop quantum gravity and \\ref{sec:BBN} on the analysis of big bang nucleosynthesis can, however, be read largely independently of each other by readers only interested in some of the aspects covered here.  We will start with general remarks on the physics underlying the problem. ",
        "conclusions": "Big bang nucleosynthesis is a highly relativistic regime which, to a good approximation, implies identical equations of state for fermions and photons. There are, however, corrections to the simple equation of state $w=\\frac{1}{3}$ for fermions even classically. One observation made here is that the interaction term derived in \\cite{FermionImmirzi} leads to such a correction and might be more constrained by nucleosynthesis than through standard particle experiments \\cite{FermionTorsion}. We have not analyzed this further here because more details of the behavior of the fermion current would be required.  A second source of corrections arises from quantum gravity. Remarkably, while quantum gravity effects on an isotropic background do correct the equations of state, they do so equally for photons and relativistic fermions. Initially, this is not expected for both types of fields due to their very different actions. Thus, quantum gravity effects do not spoil the detailed balance required for the scenario to work and bounds from big bang nucleosynthesis obtained so far are not strong.  But there are interesting limits for the primary parameter, the patch size of a quantum gravity state. It is dimensionally expected to be proportional to the Planck length $\\ell_{\\rm P}$ but could be larger. In fact, current bounds derived here already rule out a patch size of exactly the elementary allowed value in loop quantum gravity. With more precise estimates, these bounds can be improved further.  We have made use of quantum gravity corrections in a form which does not distinguish fermions from radiation. Although the most natural implementation, quite unexpectedly, provides equal corrections as shown here, there are several possibilities for differences which suggest several further investigations.  Small deviations in the equations of state and thus energy densities of fermions and radiation are possible. First, there are always quantization ambiguities, and so far we tacitly assumed that the same basic quantization choice is made for the Maxwell and Dirac Hamiltonians.  Such ambiguity parameters can be explicitly included in specific formulas for correction functions; see e.g.\\ \\cite{Ambig,ICGC,QuantCorrPert}. Independent consistency conditions for the quantization may at some point require one to use different quantizations for both types of fields, resulting in different quantum corrections and different energy densities. Such conditions can be derived from an analysis of anomaly-freedom of the Maxwell field and fermions coupled to gravity, which is currently in progress. As shown here, if this is the case it will become testable in scenarios sensitive to the behavior of energy density such as big bang nucleosynthesis.  Moreover, assuming the same quantization parameters leads to identical quantum corrections for photons and fermions only on isotropic backgrounds. Small-scale anisotropies have different effects on both types of fields and can thus also be probed through their implications on the equation of state.  For this, it will be important to estimate more precisely the typical size of corrections, which is not easy since it requires details of the quantum state of geometry. The crucial ingredient is again the patch size of underlying lattice states. On the other hand, taking a phenomenological point of view allows one to estimate ranges for patch sizes which would leave one in agreement with big bang nucleosynthesis constraints.  Interestingly, corrections studied here provide upper bounds to the patch size, and other corrections from quantum gravity are expected to result in lower bounds. A finite window thus results, which can be shrunk with future improvements in observations."
    },
    "0710/0710.3732_arXiv.txt": {
        "abstract": "% Edge-on spiral galaxies offer a unique perspective on disks. One can accurately determine the height distribution of stars and ISM and the line-of-sight integration allows for the study of faint structures. The Spitzer IRAC camera is an ideal instrument to study both the ISM and stellar structure in nearby galaxies; two of its channels trace the old stellar disk with little extinction and the 8 micron channel is dominated by the smallest dust grains (Polycyclic Aromatic Hydrocarbons, PAHs). \\cite{Dalcanton04} probed the link between the appearance of dust lanes and the disk stability. In a sample of bulge-less disks they show how in massive disks the ISM collapses into the characteristic thin dust lane. Less massive disks are gravitationally stable and their dust morphology is fractured. The transition occurs at 120 km/s for bulgeless disks. Here we report on our results of our Spitzer/IRAC survey of nearby edge-on spirals and its first results on the NIR Tully-Fischer relation, and ISM disk stability. ",
        "introduction": "For 32 edge-on galaxies, spanning Hubble type and mass, we fit the edge-on infinite disk model by \\cite{vdKruit81a} on the IRAC mocaics; the stellar dominated 4.5 $\\mu$m and the PAH emission at 8 $\\mu$m, with the stellar contribution subtracted \\citep[][]{Pahre04a}. The disk's total luminosity is inferred from the fitted model: $L_{disk} = 2 \\pi h^2 \\mu_0$, with $h$ the scale-length and $L_0$ the face-on central surface brightness \\citep[][]{Kregel02}.  Figure \\ref{f:tf} shows the inferred Tully-Fischer relation for these disks for the stars (4.5 $\\mu$m). Notably, the slope ($\\alpha$) is 3.5, similar to what \\cite{Meyer06a} found but contrary for the increasing trend of slope with redder filters. The effects of age and metallicity of the stellar population become independent and opposite effects on the color-M/L relation in NIR \\citep[See][Fig. 2d]{BelldeJong}. Hence, the shallower slope in the  IRAC stellar channels, could be the metallicity effect starting to dominate. ",
        "conclusions": ""
    },
    "0710/0710.4262_arXiv.txt": {
        "abstract": "The accuracy of position measurements on stellar targets with the future Space Interferometry Mission (SIM) will be limited not only by photon noise and by the properties of the instrument (design, stability, etc.) and the overall measurement program (observing strategy, reduction methods, etc.), but also by the presence of other ``confusing'' stars in the field of view (FOV). We use a simple ``phasor'' model as an aid to understanding the main effects of this ``confusion bias'' in single observations with SIM. This analytic model has been implemented numerically in a computer code and applied to a selection of typical SIM target fields drawn from some of the Key Projects already accepted for the Mission. We expect that less than 1\\% of all SIM targets will be vulnerable to confusion bias; we show that for the present SIM design, confusion may be a concern if the surface density of field stars exceeds 0.4 star/arcsec$^2$. We have developed a software tool as an aid to ascertaining the possible presence of confusion bias in single observations of any arbitrary field. Some \\textit{a priori} knowledge of the locations and spectral energy distributions of the few brightest stars in the FOV is helpful in establishing the possible presence of confusion bias, but the information is in general not likely to be available with sufficient accuracy to permit its removal. We discuss several ways of reducing the likelihood of confusion bias in crowded fields. Finally, several limitations of the present semi-analytic approach are reviewed, and their effects on the present results are estimated. The simple model presented here provides a good physical understanding of how confusion arises in a single SIM observation, and has sufficient precision to establish the likelihood of a bias in most cases. We close this paper with a list of suggestions for future work on this subject. ",
        "introduction": "\\label{intro} The Space Interferometry Mission (SIM\\footnote{also currently called SIM--PlanetQuest}) is being designed by NASA's Jet Propulsion Laboratory to provide a facility-class instrument for measuring the positions and proper motions of stars at optical wavelengths with micro-arc-second (\\muas) precision. This represents an improvement by several orders of magnitude over the precision of all existing astrometric instruments. For faint sources (V $\\gtrsim 15$ mag.), SIM will also be more than a factor 10 better than any other future planned space mission, and will therefore uniquely permit many new classes of problems to be addressed. Such problems include: the direct measurement (for the first time) of the masses of earth-like planets in orbit around nearby stars; determining the distances to stars by direct triangulation over the whole Galaxy and out to the Magellanic Clouds; measuring the transverse motions of galaxies in the Local Group; and, establishing the shape of the dark matter distribution in the Galaxy. A description of the instrument and its current science program is available at the JPL/SIM web site.\\footnote{See {\\it http://planetquest.jpl.nasa.gov/SIM/sim\\_index.cfm} for descriptions of the current set of Key Projects. Almost half of the available 5-year Mission time is still unallocated.} The mission is now at the end of the detailed design phase. After more than 15 years of development, all the major technical questions have been answered. New devices have been invented in order to provide metrology internal to the spacecraft at a level of a few tens of picometers, a fraction of the inter-atomic distance in a molecule of oxygen. The next major step is to begin construction of the instrument.  SIM will be the second optical interferometer in space devoted to astrometry, following the Fine Guidance Sensors (FGS) on the Hubble Space Telescope (HST). However, SIM is a Michelson interferometer using separated collectors, quite different from the filled-aperture white-light shearing interferometer design of the FGS. SIM includes three long-baseline interferometers housed on a common truss, each formed by two $\\approx 0.3$ m apertures which compress their light beams and guide them through delay lines to beam combiners. During operation, two interferometers are used for precision guiding of the spacecraft, and the third views the target of interest. Data are then accumulated by tracking the target until enough photons have been recorded to achieve the particular science goal.  The precision with which astrometry can be done on a specific stellar target with SIM will be limited by photon noise, the design and stability of the instrument, and by the data calibration and processing. We have some control over the instrument properties, which depend on the specific choices made when implementing the design in hardware. Also, the operation and calibration of the instrument can be optimized so as to maximize the achievable precision. However, there is another source of error which may be present, and which is largely out of our control; it does not reduce the ultimate \\textit{precision} with which a given \\textit{single measurement} can be made with SIM, but it may reduce the final \\textit{accuracy} of that measurement. This source of error arises because of the presence of other stars in the SIM field of view (FOV) which can ``confuse'' any single observation made on the target star. The light from these extraneous stars perturbs the measurement, so that the measured target position can differ from the true position. The difference is a \\textit{bias} which can reach a level of many times the single-measurement precision estimated from the instrument parameters alone. It is this \\textit{single-measurement confusion bias} which concerns us in this paper.  It is important to emphasize that the final \\textit{accuracy} of the astrometric parameters (position, parallax, and proper motion) determined by SIM for any given target star will be a result of carrying out a complex program of several single measurements on that target, plus repeated measurements on many other stars for the determination of calibration and baseline orientation parameters \\citep{bod04,miltur03}. Since we are concerned only with possible bias in a \\textit{single measurement}, the details of the entire observing program are not directly relevant here; however, it also means that we can not quantify the consequences of confusion bias to the final accuracy with which the ``end-of-mission'' astrometric parameters can be obtained on any specific target. It is clear that the effects of single-measurement confusion bias will generally diminish as more observations are combined. But this also means that projects which involve only a few observations of a target (e.g.\\ ``targets of opportunity'', single parallax measurements on nearby bright stars, etc.) may have a greater susceptibility to confusion bias.  Specific aspects of confusion in astrometric measurements with SIM have been considered by several authors in the recent past. \\citet{dalalgriest01} showed that the characteristic response of SIM's fixed-baseline interferometer as a function of wavenumber and delay can be used to refine a model of the distribution of confusing stars in the FOV. This model can then be used to correct the measured position of the target of interest, and in many cases where the level of confusion is not too great the astrometric accuracy can approach the measurement precision. Dalal and Griest successfully applied their method to models of confused fields in the LMC in which $\\approx 16$ faint stars are scattered over the FOV around the $\\approx 19$ magnitude target star. Photon noise is included in these models. These authors then go on to consider the additional complication if one of the stars in the FOV changes brightness owing to a micro-lensing event, and show that an extension of their fitting algorithm to include the precision photometry provided by SIM's detectors permits even this apparently-intractable case to be handled almost as well. However, their method fails when the angular separation between any pair of sources in the FOV (as projected on the interferometer baseline) corresponds to a delay difference of $\\lesssim 2$ coherence lengths for the full bandpass of the detection system. This is a projected separation of 25 milli-arcseconds (mas). Indeed, this is a \\textit{general limit} for SIM observations. Our approach is somewhat simpler than that of Dalal and Griest, and yields some improvement in the minimum angular separation which can be measured, but the basic limitation can not be overcome. We will compare our approach to theirs in more detail in a future paper.  \\citet{rajbokall01} also considered a number of specific cases of confusion on SIM astrometry. These authors introduced a graphical analogy using phasors as an aid to understanding how errors in the target position arise from confusing sources in the FOV. Typical target fields were constructed on a simulated image with grid spacing of 5 mas, and the amplitude and phase of the fringe which would be measured with SIM for a given wavelength on that image field was computed with a Fourier transform. Diffraction effects at the edges of the (presumed $\\approx 1''$ square) SIM FOV were included in constructing the model image, and vector averaging of the individual (narrow) SIM wavelength channels was used to simulate the 1-dimensional apodization of the fringes over the FOV caused by the decreasing coherence of the fringes as the bandwidth increases. \\citet{rajbokall01} were particularly interested in modeling the effects of mispointing of the FOV in subsequent visits to the same target field when target proper motions were being measured; in that case the actual distribution of field stars changes as some disappear from one side of the FOV and others appear at the other side. They included photon noise, and also addressed the issue of how the size of the FOV defined by the field stop influenced the level of confusion. Their source field models were constructed to simulate SIM observations of the position and proper motion of target stars in M31, the LMC, and the Galactic bulge. They concluded that the confusion-induced errors in position can often be significant (several times the measurement precision) for faint target stars but the proper motion errors are likely to be small. The errors are, as expected, smaller for wider measurement bands.  In one other confusion-related study, \\citet{takvellin05} considered the effect of circumstellar disks on the measurement of stellar wobble during observations aimed at detecting extra-solar planets. Their models showed that neither the motion of the disk mass center nor the contamination by disk light is a serious threat to detecting planets around pre-main sequence stars; the basic reason for the insensitivity of these observations to confusion from circumstellar disks is that interferometers tend to resolve such an extended source, reducing its influence on the astrometry of the parent star.  The studies summarized above have shown that confusion poses limits to the accuracy of any single SIM measurement, and have therefore succeeded in raising our awareness of this problem.  However, the detailed design of SIM has changed since those studies were done, and many of the changes will affect on the modeling results. The size of the collector and its central obscuration, the entrance aperture field stop defining the geometrical FOV, the transmission efficiency of the optics, the fringe disperser design (which defines the bandwidth and central wavelength of the spectral channels), and the QE and spectral response of the detector are now all much better defined. Furthermore, previous studies have focused on specific science programs which were already suspected to be pushing the capabilities of the instrument; they have not provided us with any general ``tools'' for understanding and recognizing confusion, or for dealing with it. Previous studies have also often taken a statistical approach which is less suitable for answering direct questions about specific fields, such as: is a SIM observation of this particular target embedded in that particular field of stars likely to be confused? And, can the observation be done in such a way so as to reduce the confusion bias?  What \\textit{a priori} information about the target and the field would help? And, if the observation has already been taken, can we identify the effects of confusion in the data? These questions have provided the motivation for the work described in this paper.  This paper is organized as follows. In the next section, we summarize the basic Michelson interferometer response as it applies to SIM. We then recall the phasor model introduced by \\citet{rajbokall01} and elaborate upon it as a tool for understanding the behavior of confusion in SIM astrometry. Using this analytic model, together with updated knowledge of SIM's instrument parameters, we have constructed a simulation code for evaluating the likelihood of confusion bias in any specific field; details are presented in Section~\\ref{simu}. In Section~\\ref{limit_values}, we present single measurement confusion bias as a function of magnitude difference and projected separation of an additional star present within the SIM FOV. In Section~\\ref{applications}, we apply this semi-analytic model to a number of target fields drawn from the Key Projects which have already been chosen for the initial SIM science program. From this experience we then consider how the single-measurement confusion bias might be reduced through the addition of other information. The most useful additions appear to be knowledge of the approximate locations and spectral energy distributions (SEDs) of the target and of the most troublesome confusing stars in the SIM FOV. Finally, in Section \\ref{limitations}, we describe the limitations of our current approach. These limitations are primarily related to the simplified model of the focal plane of SIM  which we have used here. In particular, in this paper, we have not modeled the detailed mechanism by which the spectral dispersion is implemented,\\footnote{A thin prism disperser turns SIM into an objective prism spectrograph on the CCD detector.} nor have we considered the pixellation of the focal plane by the CCD detector. We have explored these points with the aid of a more elaborate model that includes these effects, and find that the biases estimated using this more elaborate model differ only by small amounts from those provided by the approach described here.\\footnote {However, consideration of this more sophisticated model does suggest additional ways to reduce any confusion bias.} We have therefore chosen to present the main issues relevant to SIM confusion with a minimum of complication, and leave the discussion of the more elaborate instrument model to a future paper.  Binary stars will be an important class of targets for SIM, and are the topic of one of the major Key Projects. In these cases, the two stars are in a bound orbit and are physically close to each other. Typical binaries to be studied with SIM will have separations from about a few mas to 1000 mas and orbital periods from a few days to several years.  Stars in crowded fields can occasionally mimic the effects of binaries if their projected separations become small for particular baseline orientations, but the effects of confusion from such ``apparent'' binaries can be reduced (or even eliminated) by rotating the interferometer baseline and repeating the observation. However, for ``real'' binaries, rotation of the baseline is an integral part of the measurement process. The goal of the binary observation is to obtain the characteristics of the orbit, and this is done by measuring the positions of the components for a number of baseline orientations. These targets are sufficiently specialized that we have removed them from the list of crowded-field problems treated in this paper. Such observations treat binaries as ``signal\", whereas here we treat them as ``noise\".  A discussion of astrometry on binary stars with SIM is planned for a future publication. ",
        "conclusions": "\\label{sum}  We have examined the bias that can occur in a single measurement with SIM owing to the presence of field stars within the FOV. In order to accomplish this task we have presented a model for the SIM interferometer, and a description of how SIM carries out a single measurement of the position of an isolated target star. The measured instrument response is then perturbed by adding a field star to the model FOV; the difference in the angles measured in the two cases is called the ``confusion bias''. The extremes of this bias are calculated for the specific (but common) case of a binary system in order to illustrate its main properties.  A number of source models are then developed which resemble the fields to be studied by SIM in several of the Key Projects already selected for inclusion in the initial mission science program. An unconfused version of the source model consisting only of the main target star is used as a reference measurement, and the results compared with a measurement made on the fully-populated field. The difference is the confusion bias in a single SIM measurement. Observations are simulated at various orientations of the interferometer baseline, and variants of the full field are examined in order to understand the sensitivity of the bias to structural details in the field.  The magnitude of the confusion bias is found to depend on a number of factors, some obvious, others perhaps less so: \\begin{itemize} \\item the relative brightnesses of the target and the field stars; \\item the shapes of the SEDs of the target and field stars; \\item the angular separation of the stars from the center of the FOV; \\item the angular separation of the field stars from the target star as projected on the interferometer baseline; and, \\item the baseline orientation. \\end{itemize} The largest contributions to the confusion bias in a crowded field come from a small number of stars having small projected angular separations from the target, but these stars may actually be located outside of the FOV. Field stars which are less than 4 mag fainter than the target and which have projected separations within 100 mas of the target are potentially the most troublesome.  The results of this study provides the understanding and the tools required to examine the likelihood of confusion bias in any single measurement with SIM. Unfortunately, data on the field stars in any specific FOV (especially their positions) is not likely to be available with sufficient accuracy to actually remove this bias.\\footnote{There are some possible exceptions one could imagine, but we have not explored them further here.} Our study nevertheless suggests some strategies for recognizing the presence of confusion bias and for dealing with it, both in the observation planning stage and in the data reduction stage. These strategies might include the following: \\begin{itemize} \\item While dealing with crowded fields, avoid fields with star densities in excess of 0.4 stars per square arcsec. \\item If avoidance is impossible, evaluate the likelihood of confusion in the field by using the tools developed here. \\item If confusion is likely, try to reduce your sensitivity to it by planning the observing program so that data is taken at the least sensitive orientations of the interferometer baseline. \\item If too little is known about the specific field, plan to distribute the available observing time over a number of orientations of the interferometer baseline which differ by a few degrees from each other. Inconsistent values in the data set can then be rejected with motivation. \\item If confusion is suspected in a given set of observations for which no prior data exists, acquiring new imagery from e.g., speckle or adaptive optics imaging would be useful for building a model. \\end{itemize} There is one additional strategy suggested by our more accurate model of the SIM focal plane. The CCD detector in the focal plane of SIM's camera is planned to have pixels which are smaller than the diameter of the FOV. If the data in the individual pixels can be made available, it would be possible to choose e.g.\\ only the central pixel, thereby effectively reducing the FOV and possibly attenuating an offending field star. The penalty of fewer target photons could then be offset by a reduction in the level of confusion bias. This possibility will be discussed in more detail in a future paper \\citep{sriron07}.  We wish to emphasize that the results of this paper refer to a bias which may be present in a single measurement of angular position with SIM. The determination of the full set of astrometric parameters (position, parallax, proper motion) on any SIM target will be done with a number of measurements, reducing the effects of any single anomalous point. Furthermore, the ultimate accuracy of the results depends on an extensive calibration program to determine the instrumental parameters, including the baseline length and orientations for each field observed."
    },
    "0710/0710.4054_arXiv.txt": {
        "abstract": "Boxy/peanut bulges in disk galaxies have been associated to stellar bars. In this talk, we discuss the different properties of such bulges and their relation with the corresponding bar, using a very large sample of a few hundred numerical N-body simulations. We present and inter-compare various methods of measuring the boxy/peanut bulge properties, namely its strength, shape and possible asymmetry. Some of these methods can be applied to both simulations and observations. Our final goal is to get correlations that will allow us to obtain information on the boxy/peanut bulge for a galaxy viewed face-on as well as information on the bars of galaxies viewed edge-on. ",
        "introduction": "Simulations have shown that bars are not vertically thin morphological features, but have a considerable vertical extent and a vertical structure, known as the Boxy/Peanut bulges (hereafter B/P; Combes \\& Sanders 1981, Combes et al. 1990). Comparisons between observations and $N$-body simulations have established this direct connection firmer (Athanassoula 2005 and references therein). Furthermore, observations have shown that both bars and B/P bulges are quite predominant in disc galaxies and that the corresponding frequencies are in good agreement with the link between the two structures (L\\\"utticke, Dettmar \\& Pohlen 2000).  We measure the peanut properties in a large sample of several hundred $N$-body simulations ran by one of us (EA) for different purposes. More information on these simulations and on their properties can be found in Athanassoula \\& Misiriotis (2002) and Athanassoula (2003, 2007). In particular, we seek correlations between the properties of the bar and the properties of the B/P bulge. ",
        "conclusions": "We presented several methods to calculate the strength of the bar and the strength, shape and asymmetry of the B/P bulge and found strong correlations between their results. The most important correlation relates the strength of the bar with the strength of the B/P bulge, the strongest bars having the strongest peanuts. We also find that the strength of the peanut depends on the number of buckling episodes it underwent, the strongest bars having undergone more buckling episodes (Fig. 4). Finally, we find a very interesting result about $C_{2,z}(R)$, i.e. about the shape of the radial density profiles along cuts perpendicular to the equatorial plane. For strong bars, having a strong peanut or X-shaped bulge, this profile is more flat-topped, while for weaker bars, with more boxy-like bulges, it is more peaked. All the results summarised here are discussed in length by Athanassoula \\& Martinez-Valpuesta (2007, in preparation).  \\begin{figure} \\begin{center} \\includegraphics[scale=0.35]{Martinez-Valpuesta_fig4.eps} \\caption{Correlations between bar and peanut properties. Each symbol corresponds to one simulation. The type of symbol is related to the number of buckling events suffered by the bar during its evolution. {\\it Left panel}: Strength of the B/P bulge measured with our Fourier based method vs. the strength of the bar. {\\it Right panel}: Shape of the B/P bulge (i.e. shape of the radial density profiles along cuts perpendicular to the equatorial plane, measured by the minimum of the kurtosis) plotted as a function of bar strength. } \\end{center} \\end{figure}"
    },
    "0710/0710.0441_arXiv.txt": {
        "abstract": "We study the non-thermal emissions in a solar flare occurring on 2003 May 29 by using RHESSI hard X-ray (HXR) and Nobeyama microwave observations. This flare shows several typical behaviors of the HXR and microwave emissions: time delay of microwave peaks relative to HXR peaks, loop-top microwave and footpoint HXR sources, and a harder electron energy distribution inferred from the microwave spectrum than from the HXR spectrum. In addition, we found that the time profile of the spectral index of the higher-energy ($\\gsim 100$ keV) HXRs is similar to that of the microwaves, and is delayed from that of the lower-energy ($\\lsim 100$ keV) HXRs. We interpret these observations in terms of an electron transport model called {\\TPP}. We numerically solved the spatially-homogeneous {\\FP} equation to determine electron evolution in energy and pitch-angle space. By comparing the behaviors of the HXR and microwave emissions predicted by the model with the observations, we discuss the pitch-angle distribution of the electrons injected into the flare site. We found that the observed spectral variations can qualitatively be explained if the injected electrons have a pitch-angle distribution concentrated perpendicular to the magnetic field lines rather than isotropic distribution. % ",
        "introduction": "\\label{sec1} Observations of hard X-rays (HXRs), microwaves, and occasionally gamma-rays in solar flares tell us that a significant amount of non-thermal particles are produced. Among them, HXR and microwave observations are believed to provide the most direct information on electrons. Because HXRs below $\\sim 100$ keV are emitted primarily by electrons with energy below several hundred keV via {\\bremss} radiation \\citep{1971SoPh...18..489B}, whereas microwaves above $\\sim 10$ GHz are emitted by electrons above several hundred keV via {\\gyros} \\citep{1969ApJ...158..753R,1999spro.proc..211B}, these two sources of emission give us information on electrons in two different energy ranges. Therefore, a comparative study by using both HXR and microwave observations is useful for discussing the physics of flare non-thermal electrons over a wide range of energies.  Impulsive behavior is commonly seen in both HXR and microwave lightcurves \\citep{1974IAUS...57..105K}, but the two emissions do not necessarily behave identically. Temporally, higher-energy HXR and microwave emissions tend to be delayed from lower-energy HXRs \\citep[e.g.,][]{1978ApJ...223..620C,1983Natur.305..292N,1985ApJ...292..699B,1997ApJ...487..936A}. \\cite{1997ApJ...487..936A} statistically analyzed the low-pass filtered HXR lightcurves for 78 flares observed with the {\\it Compton Gamma Ray Observatory} ({\\it CGRO}) and find a systematic increase of time delay toward higher energy. They interpreted these time delays in terms of electron precipitation under Coulomb collisions.  Spatially, microwave sources do not always coincide with HXR sources. HXRs are typically emitted at the footpoint regions of the flare loop \\citep{1994PhDT.......335S} whereas microwaves are emitted mainly at the loop-top region \\citep{2002ApJ...580L.185M}. \\cite{2002ApJ...580L.185M} suggested that only electrons with a pancake pitch-angle distribution concentrated transverse to the magnetic field lines can explain the observed loop-top microwave source.  Spectrally, \\cite{2000ApJ...545.1116S} statistically studied the correlation of the HXR and microwave spectral indices for 57 peaks of the non-thermal emission in 27 flares. They found that the electron energy distribution inferred from the microwave spectrum is systematically harder than that inferred from the HXR spectrum, and suggested that the electron energy distribution becomes harder towards higher energy. There are three probable explanations for such spectra: (1) two (or more) different electron populations with distinct physical characteristics, (2) ``second-step acceleration'' \\citep[e.g.,][]{1976SoPh...49..343B}, and (3) ``{\\TPP} (TPP)'' \\citep[e.g.,][]{1976MNRAS.176...15M}.  \\cite{1976MNRAS.176...15M} presented analytic solutions of the electron energy continuity equation under two conditions: {\\it strong} and {\\it weak diffusion limits} \\citep{1966JGR....71....1K}. In the strong diffusion limit, electrons injected into the loop undergo significant scattering and then are quickly isotropized during the loop transit. They can escape from the loop with a precipitation rate proportional to their velocity, $\\nu_{\\rm p} \\propto v$. In the weak diffusion limit, on the other hand, electrons are less scattered during the transit. When the loss cone distribution is formed, the pitch-angle diffusion time $\\tau_{\\rm d}$, which is longer than the transit time, controls the electron precipitation, yielding $\\nu_{\\rm p} \\propto 1/\\tau_{\\rm d}$. The precipitation rate and the evolution of electrons vary, depending on which condition applies.   There have been many observations that can be explained in terms of the TPP model \\citep[e.g.,][]{2000ApJ...531.1109L}. \\cite{2002ApJ...576L..87Y} reported the Nobeyama Radioheliograph (NoRH; \\citealt{1994PROCIEEE...82..705}) observation of a flare occurring on 1999 August 28. The NoRH observation showed clear flare-loop structure and propagating features along the loop. They showed that the microwave spectrum in the optically-thin regime is hard (with spectral index $\\sim 1.5$) around the loop-top, and then becomes softer (spectral index $\\sim 3.5$) toward the footpoints. Their observation indicates that the higher-energy electrons are efficiently trapped within the loop, supporting the TPP model. For this event, however, no HXR observation was available for comparison with the microwave observation. \\cite{2000ApJ...545.1116S} pointed out that the discrepancy of the energy distribution between the HXR and microwave emitting electrons found in their study could be explained by the TPP model. In their study, there was no imaging observation to confirm their suggestion. If the HXR and microwave sources do not coincide spatially, the discrepancy of the energy distribution between the HXR and microwave emitting electrons can be explained by the different spatial distribution between the HXR and microwave emitting electrons as a result of the TPP model. Imaging as well as spectral data at both HXR and microwave wavelengths are essential to confirm the role of TPP on the parent electrons  In this paper we analyze the non-thermal emissions of the 2003 May 29 flare by using the {\\it Reuven Ramaty High Energy Solar Spectroscopic Imager} \\citep[RHESSI;][]{2002SoPh..210....3L}, the Nobeyama Radio Polarimeters \\citep[NoRP,][and references theirin]{1985PASJ...37..163N} and NoRH. RHESSI has superior spectroscopic ability from $\\sim 3$ keV to $\\sim 17$ MeV, providing the HXR spectrum from $\\sim 3$ keV to $\\sim 300$ keV with a spectral resolution of $\\sim 1$ keV and arbitrary energy bands. In previous studies, the temporal evolution of the (HXR) spectrum has been considered in less detail, probably due to instrumental limitations. However, the temporally-resolved analysis of the spectrum is important because non-thermal emissions and thus non-thermal electrons are the most ``time-varying'' objects in solar flares. RHESSI enables us to analyze an accurate, temporally-resolved HXR spectrum above $\\sim 100$ keV. Because the HXRs above $\\sim 100$ keV are mainly emitted by electrons above $\\sim 200$ keV \\citep{1996ApJ...464..974A}, RHESSI's well-resolved spectral data below $\\sim 300$ keV provides us more accurate information on electrons from tens to hundreds of keV than before. Combining the RHESSI HXR and NoRH/NoRP microwave spectral data allows us to fully cover the electrons from tens to thousands of keV.  For a physical interpretation of the observations, we use a numerical model of TPP which treats the pitch-angle diffusion more generally than the analytic solutions developed for the weak and strong diffusion limits. \\cite{2000ApJ...543..457L} performed a similar treatment of the electron transport to explain their microwave observation of a flare on 1993 June 3. We also predict the microwave and HXR emissions from the calculated electron distribution. Comparing these model results with the observations, we discuss electron injection and transport, and address how the pitch-angle distribution of the injected electrons affects the evolution of the trapped and precipitating electrons, and their resultant emissions.  The paper proceeds as follows. In {\\S}~\\ref{sec2} we present a comparative study of the non-thermal emissions of a solar flare occurring on 2003 May 29, by using the RHESSI HXR and Nobeyama microwave observations. Temporally-resolved spectra of the HXRs and microwaves are analyzed in detail. We discuss energy-dependent delays of the time profiles of the spectral indices, which have not been discussed in previous studies. In {\\S}~\\ref{sec3} we present our treatment of the TPP model. We numerically solve the spatially-homogeneous {\\FP} equation \\citep{1990A&A...230..213M} with the Coulomb interaction \\citep[e.g.,][]{1981ApJ...251..781L} and a time-dependent injection. In {\\S}~\\ref{sec4} we describe the time evolution of the trapped and precipitating electron distribution and the predicted microwave and HXR emissions. The behavior of the HXR and microwave emissions predicted by the model are compared with the observations, allowing us to give some constraints on the properties of flare non-thermal electrons. In {\\S}~\\ref{sec5} we conclude our study. % ",
        "conclusions": "\\label{sec5} We presented the comparative study of the non-thermal emissions of the flare occurring on 2003 May 29 using the RHESSI HXR and Nobeyama microwave observations. Further, we considered the electron transport model, TPP, to explain the observations.  The 2003 May 29 flare showed two non-thermal HXR sources at the footpoints and a microwave source at the loop-top, as observed with RHESSI and NoRH. We interpreted this in terms of the TPP model. We presented the time profiles of the spectral indices of the higher-energy HXRs $\\gammat{H}{obs}$ as well as the lower-energy HXRs $\\gammat{L}{obs}$ and microwaves $\\alphat{obs}$. The spectra of microwaves and HXRs imply that the microwave emitting electrons have a harder energy distribution than the HXR emitting ones. % We found that the time profile of $\\gammat{H}{obs}$ shows similarity with that of $\\alphat{obs}$ rather than with $\\gammat{L}{obs}$, and is delayed from that of $\\gammat{L}{obs}$.  We numerically solved the spatially-homogeneous {\\FP} equation for the TPP model to describe the evolution of electrons. Precipitating electrons have a softer energy distribution than the trapped ones in the weak diffusion regime. Differences of the injection pitch-angle distribution especially affect the evolution of the precipitating electrons.  We calculated the microwave and HXR emissions from the calculated trapped electron distribution and precipitation flux for comparison with the observations. The TPP model in the weak diffusion regime can yield a soft HXR spectrum and a hard microwave spectrum. The calculated difference of the spectral indices between the HXRs and microwaves, $\\sim 1.5$, is in agreement with the observations. We further found that a pancake pitch-angle distribution for the injected electrons rather than an isotropic distribution is more adequate to qualitatively explain the temporal variation of $\\gammat{H}{obs}$. By comparing the model calculation with the observation, we can constrain the pitch-angle distribution of the injected electrons, which is crucially important for understanding the electron acceleration mechanism in solar flares.  Currently, we are improving our treatment of the TPP model to include the spatial inhomogeneity in the Fokker-Planck equation. Using this, a systematic investigation of the best parameter set to explain the observation is in progress, and will be reported in the future."
    },
    "0710/0710.0993_arXiv.txt": {
        "abstract": "The Alpha Magnetic Spectrometer (AMS), whose final version AMS-02 is to be installed on the International Space Station (ISS) for at least 3 years, is a detector designed to measure charged cosmic ray spectra with energies up to the TeV region and with high energy photon detection capability up to a few hundred GeV, using state-of-the art particle identification techniques. It is equipped with several subsystems, one of which is a proximity focusing Ring Imaging \\CK\\ (RICH) detector equipped with a dual radiator (aerogel+NaF), a lateral conical mirror and a detection plane made of 680 photomultipliers and light guides, enabling precise measurements of particle electric charge and velocity ($\\Delta \\beta / \\beta \\sim$ 10${}^{-3}$ and 10${}^{-4}$ for $Z=$~1 and $Z=$~10~$-$~20, respectively) at kinetic energies of a few GeV/nucleon. Combining velocity measurements with data on particle rigidity from the AMS-02 Tracker ($\\Delta R / R \\sim$ 2\\% for $R=$~1~$-$~10 GV) it is possible to obtain a reliable measurement for particle mass. One of the main topics of the AMS-02 physics program is the search for indirect signatures of dark matter. Experimental data indicate that dark, non-baryonic matter of unknown composition is much more abundant than baryonic matter, accounting for a large fraction of the energy content of the Universe. Apart from antideuterons produced in cosmic-ray propagation, the annihilation of dark matter will produce additional antideuteron fluxes. Detailed Monte Carlo simulations of AMS-02 have been used to evaluate the detector's performance for mass separation, a key issue for $\\bar{D}/\\bar{p}$ separation. Results of these studies are presented. ",
        "introduction": "The Alpha Magnetic Spectrometer (AMS)\\cite{bib:ams}, whose final version AMS-02 is to be installed on the International Space Station (ISS) for at least 3 years, is a detector designed to study the cosmic ray flux by direct detection of particles above the Earth's atmosphere using state-of-the-art particle identification techniques. AMS-02 is equipped with a superconducting magnet cooled by superfluid helium. The spectrometer is composed of several subdetectors: a Transition Radiation Detector (TRD), a Time-of-Flight (TOF) detector, a Silicon Tracker, Anticoincidence Counters (ACC), a Ring Imaging \\CK\\ (RICH) detector and an Electromagnetic Calorimeter (ECAL). Fig.~\\ref{amsdet} shows a schematic view of the full AMS-02 detector. A preliminary version of the detector, AMS-01, was successfully flown aboard the US space shuttle Discovery in June 1998.   \\begin{figure}[htb]  \\center   \\mbox{\\epsfig{file=ams2.eps,width=0.48\\textwidth,clip=}}  \\caption{Exploded view of the AMS-02 detector.\\label{amsdet}}   \\end{figure}   The main goals of the AMS-02 experiment are:  \\begin{itemize}  \\item A precise measurement of charged cosmic ray spectra in the rigidity region between \\mbox{$\\sim$ 0.5 GV} and \\mbox{$\\sim$ 2 TV}, and the detection of photons with energies up to a few hundred GeV;  \\item A search for heavy antinuclei ($Z \\ge$ 2), which if discovered would signal the existence of cosmological antimatter;  \\item A search for dark matter constituents by examining possible signatures of their presence in the cosmic ray spectrum.  \\end{itemize}  The long exposure time and large acceptance (0.5 m${}^2$sr) of AMS-02 will enable it to collect an unprecedented statistics of more than $10^{10}$ nuclei.   \\begin{figure}[htb]  \\center   \\mbox{\\epsfig{file=ev702white.eps,width=0.48\\textwidth,clip=}}  \\vspace{-0.3cm}  \\caption{A simulated proton event as seen in the AMS-02 display.\\label{amsdisplay1}}  \\end{figure} ",
        "conclusions": "AMS-02 will provide a major improvement on the current knowledge of cosmic rays. A total statistics of more than 10${}^{10}$ events will be collected during its operation. Detailed simulations have been performed to evaluate the detector's particle identification capabilities, in particular those of the RICH, which might be crucial for the identification of an antideuteron flux resulting from neutralino annihilation. Simulation results show that the separation of light isotopes is feasible. Using a set of simple cuts based on event data, relative mass resolutions of $\\sim$~2 \\% and rejection factors up to 10${}^4$ have been attained in D/p separation at energies of a few GeV/nucleon.   \\newpage"
    },
    "0710/0710.5578_arXiv.txt": {
        "abstract": "We investigate properties of iron fluorescent line arising as a result of illumination of a black hole accretion disc by an X-ray source located above the disc surface. We study in details the light-bending model of variability of the line, extending previous work on the subject.  We indicate bending of photon trajectories to the equatorial plane, which is a distinct property of the Kerr metric, as the most feasible effect underlying  reduced variability of the line observed in several objects.  A model involving an X-ray source with a varying radial distance, located within a few central gravitational radii around a rapidly rotating black hole, close to the disc surface, may explain both the elongated red wing of the line profile and the complex variability pattern observed in MCG--6-30-15 by {\\it XMM-Newton}.  We point out also that  illumination by radiation which returns to the disc (following the previous reflection) contributes significantly to formation of the line profile  in some cases.  As a result of this effect, the line profile always has a pronounced blue  peak (which is not observed in the deep minimum state in MCG--6-30-15), unless the reflecting material is absent within the innermost 2--3 gravitational radii. ",
        "introduction": "Broad iron lines observed from many black hole systems most likely originate from the innermost regions of an accretion disc and their profiles are shaped by gravitational redshift and Doppler shifts. Modelling of the lines observed in several objects requires strongly enhanced fluorescent emission from a few gravitational radii (e.g., Wilms  et al.~2001, Fabian et al.~2002, Miller et al.~2002, Miller et al.~2004;  see review in Reynolds \\& Nowak 2003), which in turn indicates that  a primary source of hard X-ray emission must also be located close to the  black hole. \u00a0Thus, both the primary and reflected emission should be  subject to strong gravity effects. These effects are also tentatively considered as an explanation of the complex variability  pattern characterising radiation reflected from disc, including the iron  line, and the primary hard X-ray continuum emission.  Namely, weak  variability of the reflected component, uncorrelated with the variability  of the primary emission, has been reported in a number of sources (e.g.,  Vaughan \\& Fabian 2004; see review in Fabian \\& Miniutti 2005). This is contrary to expectations from a  simple geometrical model of a hard X-ray source located close to the  reflecting disc, where a strict correlation between variations of the  primary and reflected emission should be observed.  Fabian \\& Vaughan (2003) \u00a0first argued that such a reduced variability may be explained by relativistic effects, in particular by light bending and focusing the primary emission towards the accretion disc. Qualitatively, variations of the reflected emission should be much weaker as  changes of the height of the X-ray source cause variations of its observed luminosity at infinity, while changes of the flux received by the disc (enhanced by the gravitational focusing; e.g., Matt et al.\\ 1992; Martocchia \\& Matt 1996; Petrucci \\& Henri 1997) are much weaker.  However, for a static primary source located on the symmetry axis, the illuminating radiation is focused into the innermost part of the disc and the reflected emission is subject to similar light bending as the primary.  As a result, similar variability characterises the primary and reflected emission, at least for observers with low inclination angles.  Miniutti et al.\\ (2003) and Miniutti \\& Fabian (2004) have developed further  this model  to include rotation of the primary source around the axis and the resulting beaming of its emission toward outer regions of the disc.  Then, variations of the reflected emission are reduced because of the more extended region it is produced. Predictions of their model are found to be consistent with observations of black-hole systems, e.g.\\ by Miniutti et al.\\ (2003), Miniutti, Fabian \\& Miller (2004), Fabian et al.\\ (2004).  However, a number of important effects were not systematically studied.  In particular, a specific pattern of motion of a primary source is assumed (including both location relative to the symmetry axis and azimuthal motion) but it is not discussed how strongly the resulting properties depend on these assumptions.  For example,  corotation of the primary source with the disc is assumed by Miniutti  \\& Fabian (2004), which is likely  close  to the disc surface, but at most approximate at high latitudes.  As the azimuthal motion of the primary source relative to the disc may affect significantly the reflected emission (see e.g.\\ Reynolds \\& Fabian 1997), it is not clear whether predictions of this model are specific to the underlying assumptions or generic to models involving the light bending.  In this paper we systematically analyse the light bending model, concentrating on strong gravity effects. We neglect some other effects which may contribute to the original variability problem, for example, ionization of the disc surface (Nayakshin \\& Kazanas 2002).  We focus here on the iron K$\\alpha$ line; an analysis of reflected emission including the Compton reflected radiation will be presented in our next paper.  \u00a0 We study in details a number  of geometrical scenarios of the source location  and  motion. We point out certain inadequacies in the  original computations of Miniutti \\& Fabian (2004), and how they influence  their quantitative results.  We find also a novel scenario in which reduction of the line variability follows directly from properties of photon transfer in the Kerr metric.  Namely, we find that a source located close to the disc surface,  with a varying radial distance from a Kerr black hole, gives rise to both an approximately constant illumination of the  surrounding disc and  a very strong variability of the  primary emission observed at infinity. Some aspects of variability in similar models  have been considered recently e.g.\\ by Czerny et al.\\ (2004) and Pech\\'acek et al.\\  (2005), however, these studies considered only emission emerging locally from the region under the source. On the other hand, we find that transfer  of primary emission to more distant regions of the disc is  crucial for variability effects  in such a model.  We concentrate on low inclination objects,  with application to Seyfert 1 galaxies in mind. Obviously, the Doppler  shifts are more pronounced at high inclinations, but observational studies  of high inclination objects are less advanced, either because they are obscured (as Seyfert 2 galaxies), or because the studies on dynamical time  scale are not possible (as in stellar mass black hole systems). Furthermore, we consider only  emission averaged  over at least an orbital period. Effects resulting from varying azimuthal location of an off-axis source, with respect to observer, are studied, e.g., in Ruszkowski (2000), Yu \\& Lu (2000) and Goyder \\& Lasenby (2004).  In Section 3 we analyse various physical effects relevant to formation of relativistic line profiles and variability effects; in Section 4 we apply our results to a Seyfert 1 galaxy MCG--6-30-15 which is the best studied object with clear signatures of strong gravity effects. ",
        "conclusions": "\\subsection{Variability models}  We have extended the model, formulated by Miniutti \\& Fabian (2004), relating reduced variability of the reflected emission to changing magnitude  of relativistic effects as location of the primary X-ray source changes.  We find that original computations of Miniutti \\& Fabian (2004) - for a model involving a vertically moving source - overestimate  the reduction effect by assuming a value of the outer radius of the disc which is too small for the range of the source heights considered in that model.  On the other hand, we find a significant reduction of the variability of reflected emission in a model with a rapidly rotating black hole and a source moving radially, low above  the disc surface. The reduced variability occurs then for the innermost range  of radial distances, $\\le 4R_{\\rm g}$.  We find also that - only in this range of parameters  - the GR effects give rise to  a significant decline around 6.5 keV in rms spectra.  Note that generation of strong X-ray emission at these distances ($\\le 4R_{\\rm g}$) is consistent with the condition of a high value of $a$, as - for a rapidly rotating black hole - most accretion power is dissipated within a few innermost $R_{\\rm g}$.  \\subsection{The black hole spin}  Determination of the value of black hole spin  or, even more fundamentally, verification of effects predicted  by the Kerr metric solution of GR equations -  remains a major issue of black hole astrophysics. In this context, several effects are taken into account for X-ray spectroscopy. Strong redshift of photons forming the observed red wings, at $E<4$ keV,   is considered as evidence of rapid rotation of a black hole (e.g., Brenneman \\&  Reynolds 2006). However, similar redshift can be obtained for $a=0$ if emission from $r_{\\rm d}<6$  is taken into account (Reynolds \\& Begelman 1997).  Then, the derived high value of $a$ relies on assumption of no neutral iron emission  from within the radius of marginal stability.   Response of the line to increase of primary emission has been suggested for future studies.  In particular, Reynolds et al.\\ (1999) indicate a bump occurring in the line profile and  progressing to lower energies, with proceeding time, as a feature characteristic for a rapidly rotating black hole. Again, a similar - redshifted and Shapiro delayed - bump should appear  for a non-rotating black hole if fluorescence inside $r_{\\rm ms}$ was taken into account.  A straightforward analysis of the space-time metric could be performed for systems observed close to edge-on, for which effects due to lensing by a black hole would be directly seen in the line profile (e.g., Zakharov \\& Repin 2003).  However, as noted in Narayan \\& McClintock (2005), there seems to be  a selection effect preventing such systems from being observed.  Then, we point out that a (largely unambiguous) analysis of imprints of  the space-time metric would be possible in profile of the line resulting from  irradiation by a strong flare just above the disc surface.  If such a flare occurred at $r_{\\rm s} \\la 4$, properties of the space-time related with  the black hole rotation would result in $\\ga 10$ per cent in magnitude effects in the line profile. The effects related to the value of $a$ are rather subtle but in principle possible to establish observationally.   A compact flare dominating total emission would be required to make such an analysis feasible and viability of such scenario is uncertain.  Such flares are indeed occasionally observed (e.g., Ponti et al.\\ 2004).  Interestingly, however, the Fe K$\\alpha$ line was actually very weak during the flare analysed by Ponti et al.\\ (2004), and a strong line appeared in the spectrum with significant time delay after the flare.  Finally, we emphasise that the reduced variability of reflected component -  basing on mechanism  advocated in this paper - is itself a direct manifestation of the nature of the Kerr space-time. Note that another effect resulting from  properties of the Kerr metric,  namely anisotropic emission of hard X-rays generated close to a rotating black hole, is qualitatively consistent with inclination-angle dependence of intrinsic spectra of Seyfert galaxies (Nied\\'zwiecki 2005).  In support for the tentative relation of these two effects to properties of the Kerr metric, note also that comparison of the total mass in black holes in the local  universe with the total luminosity produced by active galactic nuclei indicates that most supermassive black holes should rotate rapidly, e.g.\\  Elvis, Risaliti \\& Zamorani (2002).  \\subsection{Modelling the red wing of relativistic Fe lines}  As discussed above, detection of photons with $E < 4$ keV is considered  as the evidence of rapid rotation of the black hole. Such strongly redshifted photons were revealed in the Fe line profiles observed in MCG--6-30-15 and several other objects (e.g., Miller et al.\\ 2004, Miniutti et al.\\ 2004).   When fitted by models assuming a power-law radial emissivity, the observed profiles require $q>4$ within $(6-10) R_{\\rm g}$.  Considering physical scenarios for generation of such profiles, we find that they \\newline (i) can be produced   as a result of illumination by a hard X-ray source located close to a black hole (preferably at $r_{\\rm s} =2$--3) and rather close to the disc surface ($h_{\\rm s} \\la 1$); \\newline (ii) cannot be explained by models involving a source at a height of several   $R_{\\rm g}$, or higher, close to the axis; \\newline if they are to come from a disc, with small inclination, surrounding a rotating black hole. Regarding (ii) we find, in agreement with previous studies, that such location of the source yields a  steep $\\epsilon_{\\rm Fe}$, with $q>3$, in the innermost part of the disc. However, this steep $\\epsilon_{\\rm Fe}$ occurs only within $r_{\\rm d} \\le 2$ (and due to gravitational blueshift rather than light bending). The majority of emission from this region of the disc is either captured or bent toward the disc plane; moreover,  the observed photons are strongly redshifted. As a result, this emission gives only a minor contribution to the line profile and it is not relevant to the shape and strength of the observed red wings.  This conclusion is not affected by effects related with azimuthal motion of the source, returning radiation or angular emission law of fluorescence; these effects influence mostly the strength of the blue peak. On the other hand, the red wing is primarily related with position of the X-ray source.  \\subsection{Returning radiation, non-local irradiation}  A further challenge for modelling certain line profiles results from illumination of the surrounding disc (i.e. beyond a few $R_{\\rm g}$) by returning radiation. We find that this effect may be strong  for the range of parameters (small $r_{\\rm s}$ and $V$) previously not explored.  A similarly strong enhancement of Compton reflection by returning radiation, for very small distances ($r_{\\rm s} \\la 2$) of primary source, was noted recently by Suebsewong et al.\\ (2006).  Another effect, related to bending of photon trajectories to the disc plane,  affects models relating the Fe radial profile to some physical processes.  These models typically assume that the radial profile of Fe emission is the same as some physically motivated profile of energy dissipation.  E.g., Reynolds et al.\\ (2004) make such assumption in their analysis of the {\\it XMM} observation of MCG--6-30-15, applying model of the torqued-disc emission (where emission results mostly from extraction of rotational energy of the black hole and thus is very centrally concentrated). Moreover, they assume that there is no dissipation, and thus no reprocessing, beyond rather  small (several $R_{\\rm g}$) outer radius.  We emphasise that even for X-rays generated very low above the disc surface,  a significant illumination of more distant regions of the disc  should occur.  Both the returning radiation and the non-locality of irradiation result primarily in enhancement  of the blue peak of the line.   \\subsection{Azimuthal motion}  Impact of strong gravity is usually studied under specific assumptions on the X-ray source motion. Typically, angular velocity of the source is related to that of the disc (e.g. Ruszkowski 2000; Miniutti \\& Fabian 2004), while Dabrowski \\& Lasenby (2001) assume a static primary source.  We note that change of $V$ significantly affects flux and profile of the iron line through combination of SR and GR effects. Moreover, the Kerr metric terms yield a non-trivial dependence  between $V$ and $\\Omega$ (equation (\\ref{v})). Then, the assumed parametrisation  of motion may be crucial for derived properties, e.g.\\ in some variability models.  Obviously, additional effects may result from relativistic vertical or radial motion of the source (e.g.\\ Yu \\& Lu 2001), which were neglected in this paper.   \\subsection{Applicability of our results}  We focused above on MCG--6-30-15, for which several observed properties may be explained by our model. Similar effects, including reduced variability of reflection, pronounced red wings or strongly reflection dominated spectra, have been revealed in some Narrow Line Seyfert 1 galaxies (e.g. Fabian et al.\\ 2004, Ponti et al.\\ 2006), as well as in stellar mass black-hole systems (e.g.\\ Miniutti et al.\\ 2004) and in galactic nuclei with much higher  black hole masses (Fabian et al.\\ 2005).  Note that our scenario implies that the X-ray luminosity   is underestimated by an order of magnitude.  in these objects.  On the other hand, most black-hole systems do not show signatures  of such extreme effects. Then, similar reduction of the X-ray flux does not necessarily characterise  any low inclination system.  Note, however that such property would be inevitable for rapidly rotating black holes, where major fraction  of accretion power is dissipated at very small  $r_{\\rm s}$, if this power is converted  into hard X-rays in situ.  As in virtually all previous studies of that subject, we analysed impact of GR  effects for the simplest case of an isotropic point source of hard X-ray emission, which approach follows from our lack of understanding of the X-ray source nature.  The derived properties result from transfer of radiation from source to disc and observer and from disc to observer.  Additional (but smaller in magnitude) effects may occur in more realistic models.  E.g., Comptonization close to a Kerr black hole gives rise to anisotropic emission, cf.~Nied\\'zwiecki (2005), therefore radiation with different spectral index  may irradiate various parts of the disc,  while the transfer effects affect only normalisation and cut-off energy of the primary emission.  The simplified description of X-ray source considered in this paper approximates  most closely scenario with magnetic flares above the disc surface. Then, our results may be directly applicable to a model with flares occurring randomly at various radial distances. Qualitatively, we may assess similar variability effects  for continuous spacial distributions of the hard X-ray source, e.g.\\ an extended corona covering the disc surface or a small hot torus replacing the disc within a few innermost $R_{\\rm g}$. Varying size of such an extended, hot plasma, in the Kerr metric, should give rise reduced variability of reflected component. In particular, a decline of direct emission from a  shrinking corona or torus would be observed, even if its intrinsic luminosity remained unchanged, while increasing fraction of its emission would be bent to the disc plane giving rise to strong reflection component.  The major shortcoming in our study results from the neglect of ionization effects. The strongest ionization should occur below the source, especially in models with low height above the disc surface. In general, strong ionization of this region should suppress the redshifted Doppler horns formed in the red-wing.  On the one hand, this would make studies of effects related to value of $a$ less feasible. On the other hand, such depletion of the variable contribution to the red wing would reduce variations of the line profile at various flux states.  However, details of ionization structure, and of the related reduction of the variable contribution to the lien, depend on additional assumptions, in particular, on azimuthal distribution of flares. Obviously, a strong single flare would give rise to stronger ionization  than  many weak flares uniformly distributed at a given $r_{\\rm s}$.  \\subsection{Summary}  If attributed to strong-gravity effects, both the time-averaged line profile  and the variability pattern observed  in MCG--6-30-15 by {\\it XMM-Newton\\/}  independently indicate that a primary  hard X-ray source must be located very close  ($\\la 4R_{\\rm g}$) to a black hole, i.e.\\ in the region where Kerr metric effects  become crucial.  Rapid rotation of the black hole is necessary (and vastly sufficient) to account for  reduction of variability of reflected emission for such spacial location of the  source.  Bending to the equatorial plane, underlying this reduction effect,  appears to be the most pronounced effect  of the Kerr metric to be studied in the X-ray spectra of black-hole systems."
    },
    "0710/0710.5052_arXiv.txt": {
        "abstract": "Magnetic fluctuations generated by a tangling of the mean magnetic field by velocity fluctuations are studied in a developed turbulent convection with large magnetic Reynolds numbers. We show that the energy of magnetic fluctuations depends on magnetic Reynolds number only when the mean magnetic field is smaller than $B_{\\rm eq} / 4 {\\rm Rm}^{1/4}$, where $B_{\\rm eq}$ is the equipartition mean magnetic field determined by the turbulent kinetic energy and ${\\rm Rm}$ is magnetic Reynolds number. Generation of magnetic fluctuations in a turbulent convection with a nonzero mean magnetic field results in a decrease of the total turbulent pressure and may cause formation of the large-scale inhomogeneous magnetic structures even in an originally uniform mean magnetic field. This effect is caused by a negative contribution of the turbulent convection to the effective mean Lorentz force. The inhomogeneous large-scale magnetic fields are formed due to the excitation of the large-scale instability. The energy for this instability is supplied by the small-scale turbulent convection. The discussed effects might be useful for understanding the origin of the solar nonuniform magnetic fields, e.g., sunspots. ",
        "introduction": "Magnetic fields in astrophysics are strongly nonuniform (see, e.g., \\cite{M78,P79,KR80,ZRS83,RSS88,RH04,O03,BS05}). Large-scale magnetic structures are observed in the form of sunspots, solar coronal magnetic loops, etc. There are different mechanisms for the formation of the large-scale magnetic structures, e.g., the magnetic buoyancy instability of stratified continuous magnetic field \\cite{P79,P66,G70,P82}, the magnetic flux expulsion \\cite{W66}, the topological magnetic pumping \\cite{DY74}, etc.  Magnetic buoyancy applies in the literature for different situations (see \\cite{P82}). The first corresponds to the magnetic buoyancy instability of stratified continuous magnetic field (see, e.g., \\cite{P79,P66,G70,P82}), and magnetic flux tube concept is not used there. The magnetic buoyancy instability of stratified continuous magnetic field is excited when the scale of variations of the initial magnetic field is less than the density stratification length. On the other hand, buoyancy of discrete magnetic flux tubes has been discussed in a number of studies in solar physics and astrophysics (see, e.g., \\cite{P82,P55,S81,SB82,FS93,SC94}). This phenomenon is also related to the problem of the storage of magnetic fields in the overshoot layer near the bottom of the solar convective zone (see, e.g., \\cite{SW80,T01,TH04,B05}).  A universal mechanism of the formation of the nonuniform distribution of magnetic flux is associated with a magnetic flux expulsion. In particular, the expulsion of magnetic flux from two-dimensional flows (a single vortex and a grid of vortices) was demonstrated in \\cite{W66}. In the context of solar and stellar convection, the topological asymmetry of stationary thermal convection plays very important role in the magnetic field dynamics. In particular, the topological magnetic pumping is caused by the topological asymmetry of the thermal convection \\cite{DY74}. The fluid rises at the centers of the convective cells and falls at their peripheries. The ascending fluid elements (contrary to the descending fluid elements) are disconnected from one another. This causes a topological magnetic pumping effect allowing downward transport of the mean horizontal magnetic field to the bottom of a cell but impeding its upward return \\cite{ZRS83,DY74,GP83}.  Turbulence may form inhomogeneous large-scale magnetic fields due to turbulent diamagnetic and paramagnetic effects (see, e.g., \\cite{KR80,Z57,VK83,K91,KR92,RKR03}). Inhomogeneous velocity fluctuations lead to a transport of mean magnetic flux from regions with high intensity of the velocity fluctuations. Inhomogeneous magnetic fluctuations due to the small-scale dynamo cause turbulent paramagnetic velocity, i.e., the magnetic flux is pushed into regions with high intensity of the magnetic fluctuations. Another effects are the effective drift velocities of the mean magnetic field caused by inhomogeneities of the fluid density \\cite{K91,KR92} and pressure \\cite{KP93}. In a nonlinear stage of the magnetic field evolution, inhomogeneities of the mean magnetic field contribute to the diamagnetic or paramagnetic drift velocities depending on the level of magnetic fluctuations due to the small-scale dynamo and level of the mean magnetic field \\cite{RK04}. The diamagnetic velocity causes a drift of the magnetic field components from the regions with a high intensity of the mean magnetic field.  The nonlinear drift velocities of the mean magnetic field in a turbulent convection have been determined in \\cite{RK06}. This study demonstrates that the nonlinear drift velocities are caused by the three kinds of the inhomogeneities, i.e., inhomogeneous turbulence; the nonuniform fluid density and the nonuniform turbulent heat flux. The nonlinear drift velocities of the mean magnetic field cause the small-scale magnetic buoyancy and magnetic pumping effects in the turbulent convection. These phenomena are different from the large-scale magnetic buoyancy and magnetic pumping effects which are due to the effect of the mean magnetic field on the large-scale density stratified fluid flow. The small-scale magnetic buoyancy and magnetic pumping can be stronger than these large-scale effects when the mean magnetic field is smaller than the equipartition field determined by the turbulent kinetic energy \\cite{RK06}. The pumping of magnetic flux in three-dimensional compressible magnetoconvection has been studied in direct numerical simulations in \\cite{OS02} by calculating the turbulent diamagnetic and paramagnetic velocities.  Turbulence may affect also the Lorentz force of the large-scale magnetic field (see \\cite{KRR89,KRR90,KR94,KMR96}). This effect can also form inhomogeneous magnetic structures. In this study a theoretical approach proposed in \\cite{KRR89,KRR90,KR94,KMR96} for a nonconvective turbulence is further developed and applied to investigate the modification of the large-scale magnetic force by turbulent convection and to elucidate a mechanism of formation of inhomogeneous magnetic structures.  This paper is organized as follows. In Sect.~II we discuss the physics of the effect of turbulence on the large-scale Lorentz force. In Sect.~III we formulate the governing equations, the assumptions, the procedure of the derivations of the large-scale effective magnetic force in turbulent convection. In Sect.~IV we study magnetic fluctuations and determine the modification of the large-scale effective Lorentz force by the turbulent convection. In Sect.~V we discuss  formation of the large-scale magnetic inhomogeneous structures in the turbulent convection due to excitation of the large-scale instability. Finally, we draw conclusions in Sect.~VI. In Appendix~A we perform the derivation of the large-scale effective Lorentz force in the turbulent convection. ",
        "conclusions": "In the present study we investigate magnetic fluctuations generated by a tangling of the mean magnetic field in a developed turbulent convection. When the mean magnetic field $B \\ll B_{\\rm eq} / 4 {\\rm Rm}^{1/4}$, the energy of magnetic fluctuations depends on magnetic Reynolds number. We study the modification of the large-scale magnetic force by turbulent convection. We show that the generation of magnetic fluctuations in a turbulent convection results in a decrease of the total turbulent pressure and may cause formation of the large-scale magnetic structures even in an originally uniform mean magnetic field. This phenomenon is due to a negative contribution of the turbulent convection to the effective mean magnetic force.  The large-scale instability causes the formation of inhomogeneous magnetic structures. The energy for these processes is supplied by the small-scale turbulent convection, and this effect can develop even in an initially uniform magnetic field. In contrast, the Parker's magnetic buoyancy instability is excited when the density stratification scale is larger than the characteristic scale of the mean magnetic field variations (see \\cite{P66,P79,G70}). The free energy in the Parker's magnetic buoyancy instability is drawn from the gravitational field. The characteristic time of the large-scale instability is of the order of the Alfv\\'{e}n time based on the large-scale magnetic field.  We study an initial stage of formation of the large-scale magnetic structures for horizontal and vertical mean magnetic fields relative to the vertical direction of the gravity field. In the turbulent convection there are two ranges for the large-scale instability of the horizontal mean magnetic field. The first range for the instability is related to the negative contribution of turbulence to the effective magnetic pressure for the case of $L_{\\rho} < L_{B}$, while the second range for the instability is mainly caused by the anisotropic contribution of the turbulent convection to the effective magnetic force. The large-scale instability of the vertical uniform mean magnetic field is caused by the modification of the mean magnetic tension by small-scale turbulent convection. The discussed effects in the present study might be useful for the understanding of the origin of the sunspot formation.  Since in the present study we neglect very small Brunt-V\\\"{a}is\\\"{a}l\\\"{a} frequency based on the gradient of the mean entropy, we do not investigate the large-scale dynamics of the mean entropy. This problem was addressed in \\cite{KM00} whereby the modification of the mean magnetic force by the turbulent convection was not taken into account.  In order to study magnetic fluctuations and the modification of the large-scale Lorentz force by turbulent convection we apply the spectral $\\tau$ approximation (see Sect.~III). The $\\tau$ approach is an universal tool in turbulent transport that allows to obtain closed results and compare them with the results of laboratory experiments, observations and numerical simulations. The $\\tau$ approximation reproduces many well-known phenomena found by other methods in turbulent transport of particles and magnetic fields, in turbulent convection and stably stratified turbulent flows (see below).  In turbulent transport, the $\\tau$ approximation yields correct formulae for turbulent diffusion, turbulent thermal diffusion and turbulent barodiffusion (see, e.g., \\cite{EKR96,BF03}). The phenomenon of turbulent thermal diffusion (a nondiffusive streaming of particles in the direction of the mean heat flux), has been predicted using the stochastic calculus (the path integral approach) and the $\\tau$ approximation. This phenomenon has been already detected in laboratory experiments in oscillating grids turbulence \\cite{EEKR04} and in a multi-fan turbulence generator \\cite{EEKR06} in stably and unstably stratified fluid flows. The experimental results obtained in \\cite{EEKR04,EEKR04} are in a good agreement with the theoretical studies performed by means of different approaches (see \\cite{EKR96,PM02}).  The $\\tau$ approximation reproduces the well-known $k^{-7/3}$-spectrum of anisotropic velocity fluctuations in a sheared turbulence (see \\cite{EKRZ02}). This spectrum was found previously in analytical, numerical, laboratory studies and was observed in the atmospheric turbulence (see, e.g., \\cite{L67}). In the turbulent boundary layer problems, the $\\tau$-approximation yields correct expressions for turbulent viscosity, turbulent thermal conductivity and the classical heat flux. This approach also describes the counter wind heat flux and the Deardorff's heat flux in convective boundary layers (see \\cite{EKRZ02}). These phenomena have been studied previously using different approaches (see, e.g.,  \\cite{MY75,Mc90,Z91}).  The theory of turbulent convection \\cite{EKRZ02} based on the $\\tau$-approximation explains the recently discovered hysteresis phenomenon in laboratory turbulent convection \\cite{EEKRM06}. The results obtained using the $\\tau$-approximation allow also to explain the most pronounced features of typical semi-organized coherent structures observed in the atmospheric convective boundary layers (\"cloud cells\" and \"cloud streets\") \\cite{ET85}. The theory \\cite{EKRZ02} based on the $\\tau$-approximation predicts realistic values of the following parameters: the aspect ratios of structures, the ratios of the minimum size of the semi-organized structures to the maximum scale of turbulent motions and  the characteristic lifetime of the semi-organized structures. The theory \\cite{EKRZ02} also predicts excitation of convective-shear waves propagating perpendicular to the convective rolls (\"cloud streets\"). This waves have been observed in the atmospheric convective boundary layers with cloud streets \\cite{ET85}.  A theory \\cite{ZEKR07} for stably stratified atmospheric turbulent flows based on the $\\tau$-approximation and the budget equations for the key second moments, turbulent kinetic and potential energies and vertical turbulent fluxes of momentum and buoyancy, is in a good agrement with data from atmospheric and laboratory experiments, direct numerical simulations and large-eddy simulations (see detailed comparison in Sect. 5 of \\cite{ZEKR07}).  The detailed verification of the $\\tau$ approximation in the direct numerical simulations of turbulent transport of passive scalar has been recently performed in \\cite{BK04}. In particular, the results on turbulent transport of passive scalar obtained using direct numerical simulations of homogeneous isotropic turbulence have been compared with that obtained using a closure model based on the $\\tau$ approximation. The numerical and analytical results are in a good agreement.  In magnetohydrodynamics, the $\\tau$ approximation reproduces many well-known phenomena found by different methods, e.g., the $\\tau$ approximation yields correct formulae for the $\\alpha$-effect \\cite{KR80,RK93,RK00,RKR03}, the turbulent diamagnetic and paramagnetic velocities \\cite{Z57,VK83,K91,KR92,RKR03}, the turbulent magnetic diffusion \\cite{KR80,VK83,KRP94,RKR03,RK04}, the ${\\bf \\Omega} {\\bf \\times} {\\bf J}$ effect and the $\\kappa$-effect \\cite{KR80,RKR03}, the shear-current effect \\cite{RK03,RK04,RK07}.  Generation of the large-scale magnetic field in a nonhelical turbulence with an imposed mean velocity shear has been recently investigated in \\cite{BH05} using direct numerical simulations. The results of these numerical simulations are in a good agreement with the theoretical predictions based on the $\\tau$ approximation (see \\cite{RK03,RK04,RK07}) and with the numerical solutions of the nonlinear dynamo equations performed in \\cite{BS05B,RKL06} (see detailed comparison in \\cite{RK07}).  The validity of the $\\tau$ approximation has been tested in the context of dynamo theory, in direct numerical simulations in \\cite{BSM05}. The alpha effect in mean field dynamo theory becomes proportional to a relaxation time scale multiplied by the difference between kinetic and current helicities. It is shown in \\cite{BSM05} that the value of the relaxation time is positive and, in units of the turnover time at the forcing wavenumber, it is of the order of unity. Kinetic and current helicities are shown in \\cite{BSM05} to be dominated by large scale properties of the flow. Recent studies in \\cite{SSB07} of the nonlinear alpha effect showed that in the limit of small magnetic and hydrodynamic Reynolds numbers, both the second order correlation approximation (or first-order smoothing approximation) and the $\\tau$ approximation give identical results. This is also supported by simulations \\cite{BS07} of isotropically forced helical turbulence whereby the contributions to kinetic and magnetic alpha effects are computed. The study performed in \\cite{BS07} provides an extra piece of evidence that the $\\tau$ approximation is a useable formalism for describing simulation data and for predicting the behavior in situations that are not yet accessible to direct numerical simulations."
    },
    "0710/0710.0663_arXiv.txt": {
        "abstract": "The intrinsic anisotropy $\\delta$ and flattening $\\epsilon$ of simulated merger remnants is compared with elliptical galaxies that have been observed by the SAURON collaboration, and that were analysed using axisymmetric Schwarzschild models. Collisionless binary mergers of stellar disks and disk mergers with an additional isothermal gas component, neglecting star formation cannot reproduce the observed trend $\\delta = 0.55 \\epsilon$ (SAURON relationship). An excellent fit of the SAURON relationship for flattened ellipticals with $\\epsilon \\geq 0.25$ is however found for merger simulations of disks with gas fractions $\\geq 20\\% $, including star formation and stellar energy feedback. Massive black hole feedback does not strongly affect this result. Subsequent dry merging of merger remnants however does not generate the slowly-rotating SAURON ellipticals which are characterized by low ellipticities $\\epsilon  < 0.25$ and low anisotropies. This indicates that at least some ellipticals on the red galaxy sequence did not form by binary mergers of disks or early-type galaxies. We show that stellar spheroids resulting from multiple, hierarchical mergers of star-bursting subunits in a cosmological context are in excellent agreement with the low ellipticities and anisotropies of the slowly rotating SAURON ellipticals and their observed trend of $\\delta$ with $\\epsilon$. The numerical simulations indicate that the SAURON relation might be a result of strong violent relaxation and phase mixing of multiple, kinematically cold stellar subunits with the angular momentum of the system determining its location on the relation. ",
        "introduction": "A popular formation scenario for early-type galaxies is the collision and merger of two roughly equal-mass galaxies with mass ratios between 1:1 and 4:1. This famous major merger scenario \\citep{1972ApJ...178..623T} has been very successful in explaining observed properties of ellipticals, like their kinematics, surface density profile or isophotal shape \\citep{1981MNRAS.197..179G,1983MNRAS.205.1009N,1988ApJ...331..699B, 1990Natur.344..379B,1992ARA&A..30..705B,1992ApJ...400..460H, 1993ApJ...409..548H,1993A&A...278...23B,1999ApJ...523L.133N,2001ApJ...554..291C, 2003ApJ...597..893N,2005MNRAS.357..753G,2005MNRAS.360.1185J, 2005A&A...437...69B,2006MNRAS.369..625N,2006MNRAS.372..839N, 2006ApJ...641...21R,2006astro.ph..7446C}. Numerical simulations for example showed that the family of disky, fast rotating ellipticals could result from stellar disk galaxy mergers with unequal mass ratios of 3:1 to 4:1 \\citep{1998giis.conf..275B,1998ApJ...502L.133B,1999ApJ...523L.133N, 2003ApJ...597..893N,2005A&A...437...69B,2006MNRAS.372..839N} or from gas-rich 1:1 to 2:1 disk mergers where the gas subsequently settles into the equatorial plane of the merger remnant and produces a new stellar disk component \\citep{2005MNRAS.359.1379K, 1996ApJ...471..115B,2001ApJ...555L..91N,2002MNRAS.333..481B, 2005ApJ...622L...9S}. Boxy, slowly rotating ellipticals, on the other hand form in stellar disk-disk mergers with mass ratios of 1:1 to 2:1 \\citep{1994ApJ...427..165H,2003ApJ...597..893N} or from multiple major disk mergers \\citep{1996ApJ...460..101W}. \\citet{2007arXiv0709.3439B} showed that repeated minor mergers result in remnant properties very similar to one corresponding to major mergers.  A serious problem of the major merger scenario was the fact that cosmological models do not predict a dependence of the mass ratio of mergers on total galaxy mass or luminosity \\citep{2003ApJ...597L.117K}. If, on the other hand, the mass ratio determines the isophotal shape and rotational properties of merger remnants one would expect that the ratio of the number of fast rotating, disky ellipticals to the number of slowly rotating, boxy systems should be independent of luminosity \\citep{2006ApJ...636L..81N}. This is in contrast with observations that show a strong dependence of isophotal shape and rotational properties on galaxy mass. While massive galaxies are preferentially boxy, slow rotators, lower-mass ellipticals are predominantly disky and fast rotators (for a summary see \\citealp{1996ApJ...464L.119K}). \\citet{2003ApJ...597L.117K} argued that this mass dependence of galaxy properties could be explained as a result of differences in the morphologies of the merging progenitors (see e.g. \\citealp{2007MNRAS.tmp..808K}). Using semi-analytical models they showed that gas-rich disk-disk mergers dominate at the low-mass end of ellipticals while intermediate mass ellipticals should have formed preferentially from mixed mergers involving a disk and an elliptical galaxy. Finally, the most massive early-type galaxies should have experienced a last elliptical-elliptical merger (dry merger) \\citep{2006ApJ...636L..81N}. Mixed mergers have not yet been studied in details, although, according to \\citet{2003ApJ...597L.117K} they should be more frequent than dry mergers (see however \\citealp{2007arXiv0706.1243H}). Dry mergers and their implications for the formation of the red galaxy sequence have however received a lot of attention recently \\citep{2005MNRAS.359.1379K,2005MNRAS.361.1043G,2005astro.ph..6044F, 2006ApJ...636L..81N,2006MNRAS.369.1081B}.  Further refinement of theoretical models has recently been achieved by including black-hole physics in simulations of galaxy mergers and elliptical galaxy evolution. Energetic feedback from central black holes might solve some pending problems of major merger models, like the suppression of late inflow of cold gas and star formation that would make ellipticals look much bluer than observed \\citep{2005ApJ...620L..79S}.  In summary, despite several still unsolved questions (e.g. \\citealp{2007astro.ph..2535N}), the major merger scenario has  become a popular model in order to explain the origin of bulge-dominated, spheroidal galaxies (see e.g. \\citealp{2007arXiv0706.1246H,2007arXiv0706.1243H}).  Progress in understanding galactic evolution is often driven by strong interactions between observers and theorists. Increasingly more sophisticated theoretical/numerical models are confronted with continuously improving observations that lead to new theoretical challenges. One example is the SAURON project \\citep{2004MNRAS.352..721E} which aims to determine the 2-dimensional structural and kinematical properties of early-type galaxies using a panoramic integral-field spectrograph. In order to interpret the observations and study the intrinsic galaxy structure, axisymmetric Schwarzschild models are applied to the observations \\citep{2006MNRAS.366.1126C,2007MNRAS.379..418C}. The results, published so far, have revealed interesting fine structures and physical properties which provide new and deeper insight into the origin of galaxies, in particular when compared to simulation \\citep{2000MNRAS.316..315B,2007MNRAS.376..997J}.  In this paper we confront a recently published SAURON analysis of preferentially axisymmetric elliptical galaxies \\citep{2007MNRAS.379..418C} with the predictions of numerical merger simulations and cosmological models of galaxy formation. Section 2 summarizes the observations. Section 3 shows that simulations of collisionless and gaseous disk-disk mergers, neglecting star formation and stellar feedback cannot reproduce the observational results. We demonstrate that star formation and stellar energy feedback has a strong effect on the final structure of merger remnants, leading to a good agreement with the SAURON observations of fast rotating ellipticals. The origin of the round, almost isotropic, slowly rotating SAURON ellipticals is explored in section 4. Isolated, dry mergers of ellipticals that formed as discussed in section 3 cannot explain these objects. We show however that cosmological initial conditions, leading to a series of multiple major and minor mergers, coupled with local star bursts generate spheroidal stellar systems in very good agreement with the observations. Conclusions follow in section 5. ",
        "conclusions": "The numerical simulations, discussed in the previous sections, have shown that interstellar gas dynamics, star formation and stellar feedback plays a crucial role in order to reproduce the observed kinematical and isophotal properties of fast rotating, early-type galaxies. The final structure of the merger remnants depends on the initial mass ratio and gas fraction. The remnants are more round, less anisotropic and more rotationally supported the smaller the mass ratio $M_1/M_2 \\geq 1$ of the progenitors and the larger the initial gas fraction. The dependence of $\\delta$ on $\\epsilon_{int}$ is in agreement with the observed trend found in the SAURON sample.  Subsequent dry re-merging of disk-disk merger remnants however does not generate the observed slowly-rotating red SAURON ellipticals with small anisotropies and ellipticities. This indicates that at least some early-type galaxies on the red galaxy sequence formed in a different way. We find that multiple mergers of stellar substructures that formed from cold gas infall into dark matter halos in cosmological simulations produce round, isotropic and slowly-rotating relaxed stellar systems that are in perfect agreement with the SAURON observations. Multiple mergers of stellar systems in dense group enviroments therefore appear to be a promising alternative scenario for the origin of the red, massive galaxy population.  Despite the fact that merger simulations with star formation lead to a correlation between anisotropy and ellipticity ($\\delta = 0.67 \\times \\epsilon_{int}$) that is very similar to that inferred from observations its origin is not understood yet. It is interesting that merger remnants appear to move closer towards this relation along lines of constant $V/\\sigma$ (i.e. roughly constant specific angular momentum) in the case of a strong relaxation process. Here, strong relaxation is defined as the merger of a system of kinematically cold systems of stars that lateron break up and generate a kinematically hot stellar remnant. Several conditions could lead to such a violent dynamical process. The cold stellar clumps could for example have formed in the star-bursting tidal tails of interacting, gas-rich disk galaxies. Another possibility is the cosmological multiple merging of dark matter substructures with an embedded stellar systems. The SAURON relation might represent the relaxed and phase-mixed end state of these complex mergers with the location of the remnant on the relation being determined by its specific angular momentum which is related to its value of $V/\\sigma$. More theoretical work will be required in order to better understand these interesting questions and their connection to early-type galaxy formation."
    },
    "0710/0710.2450_arXiv.txt": {
        "abstract": "Smooth double crossing of the phantom barrier $w_{\\Lambda} = -1$ has been found possible in cosmological model with Gauss-Bonnet-scalar interaction, in the presence of background cold dark matter. Such crossing has been observed to be a sufficiently late time phenomena and independent of the sign of Gauss-Bonnet-scalar interaction. The luminosity distance versus redshift curve shows a perfect fit with the $\\Lambda CDM$ model up to $z=3.5$. ",
        "introduction": "The puzzle associated with recent cosmic acceleration, triggered by $70\\%$ of dark energy or more \\cite{a1} is far from being resolved uniquely. In the mean time, cosmologists are being confronted with yet another more intriguing challenge to explain the crossing of the so called phantom divide line $(w_{\\Lambda}=-1)$, at sufficiently late time of cosmological evolution. Some recent analysis \\cite{a1},\\cite{b1} of the presently available observational data are in favour of the value $w_{de}<-1$, at present., $w_{de}$ being the dark energy equation of state. There are also a lot of evidence all around \\cite{c}, of a dynamical dark energy equation of state, which has crossed the so called phantom divide line $w_{\\Lambda}=-1$ recently, at the value of red-shift parameter $z \\approx 0.2$. Apparently though the problem turns out to be more serious and complicated, but then, the puzzle of crossing the phantom divide line has also rendered some sort of selection rule. $\\Lambda CDM$-model, which is known to suffer from the disease of fine tuning (see \\cite{d} for a comprehensive review) can now be ruled out due to the requirement of a dynamic state parameter. Further, if the analysis of Vikman \\cite{e} is correct, then it is not possible to cross the phantom divide line in a single minimally coupled scalar field theory, without violating the stability both at the classical \\cite{f} and also at the quantum mechanical levels \\cite{g}, (though it has recently been inferred \\cite{on} that quantum Effects which induce the $w<-1$ phase, are stable in the $\\phi^4$ model). Thus single minimally coupled scalar field models like quintessence $(w>-1)$ and phantom $(w<-1)$ are to be kept aside. Consequently, we are now left with some what more complicated models. One of these is a hybrid model, composed of two scalar fields, viz, quintessence and phantom - usually dubbed as quintom model \\cite{h}. Other models like non-minimal scalar tensor theory of gravity \\cite{i}, hessence \\cite{j} and models including higher order curvature invariant terms \\cite{k} also exist in the literature. \\\\ Gauss-Bonnet term is yet another candidate which may be pursued for the purpose. The possibility of crossing the phantom divide line through Gauss-Bonnet interaction has been explored in some recent works \\cite{l},\\cite{m}. But then, these models are even complicated in the sense that either brane-world scenario \\cite{l} or scalar field and matter coupling \\cite{m} are invoked. In this article the possibility of smooth crossing of the phantom divide line $w_{\\Lambda}=-1$ has been expatiated simply by introducing Gauss-Bonnet-Scalar coupling term in the 4-dimensional Einstein-Hilbert action.\\\\ Gauss-Bonnet term arises naturally as the leading order of the $\\alpha'$ expansion of heterotic superstring theory, where, $\\alpha'$ is the inverse string tension \\cite{n}. Gauss-Bonnet term is topologically invariant and thus does not contribute to the field equations in four dimensions. However, the low energy limit of the string theory gives rise to the dilatonic scalar field which is found to be coupled with various curvature invariant terms \\cite{o}. The leading quadratic correction gives rise to Gauss-Bonnet term with a dilatonic coupling \\cite{p}. Therefore it is reasonable to consider Gauss-Bonnet interaction in four dimension with dilatonic-scalar coupling. Several works with Gauss-Bonnet-dilatonic coupling are already present in the literature \\cite{q}. In particular, important issues like - late time dominance of dark energy after a scaling matter era and thus alleviating the coincidence problem, crossing the phantom divide line and compatibility with the observed spectrum of cosmic background radiation have also been addressed recently \\cite{km}.\\\\ In a recent work with Gauss-Bonnet interaction \\cite{a}, a solution in the form $a=a_{0}e^{A\\sqrt t}$ ($a$ being the scale factor, and $A>0$) has is been found to satisfy the field equations with different forms (sum of exponentials, sum of inverse exponentials, sum of powers and even quadratic) of potentials. Solution in a more general form $(a=a_{0}e^{A t^f}), A>0, 0<f<1$, for Einstein's gravity with a minimally coupled scalar field was found in the nineties \\cite{r} and was dubbed as intermediate inflation. We \\cite{a}, on the other hand, observed that such solution depicts a transition from decelerated to accelerated expansion at sufficiently later epoch of cosmic evolution, which asymptotically goes over to de-Sitter expansion. Thus, it appeared that such solution may construct viable cosmological models of present interests. Under this consequence, a comprehensive analysis has been carried out \\cite{b} with such solution in the context of a generalized k-essence model. It has been observed that it admits scaling solution with a natural exit from it at a later epoch of cosmic evolution, leading to late time acceleration with asymptotic de-Sitter expansion. The corresponding scalar field has also been found to behave as a tracker field \\cite{s}, thus avoiding cosmic coincidence problem. \\\\ In the present work, we show that Gauss-Bonnet-Dilatonic scalar coupling with Einstein's gravity in four dimensions, admits solution in a general form $(a=a_{0}e^{A t^f}), A>0, 0<f<1$, which is viable of crossing the phantom divide line twice, once from above and the other from below in the recent epoch. Since the crossing is transient, so we may conclude that it does not show any pathological behaviour like Big-Rip \\cite{f}, at least in the classical level. ",
        "conclusions": "Altogether we have obtained a late time transient crossing of the phantom divide line first from above and more recently from below, starting from the inclusion of a Gauss-Bonnet-dilatonic scalar coupling term in the standard Einstein-Hilbert action in four dimension. Since the crossing is transient, so such double crossing is free from any sort of pathological behaviour both at the classical \\cite{f} and at the quantum mechanical \\cite{g} levels. The striking feature of the model lyes in it's indistinguishability with the standard $\\Lambda CDM$ model, in terms of the luminosity-redshift and more precisely for distance modulus-redshift curves. To identify between the two models we therefore require to observe dark energy equation of state $w_{de}$ independently. If $w_{de}$ is truly found to be dynamical and has really encountered a recent crossing of the phantom divide line, then only we can definitely distinguish the standard $\\Lambda CDM$ model with the present one. It is highly interesting to learn that smooth transient crossing of the phantom divide line is allowed for both negative and positive $(\\lambda \\lessgtr 0)$ type of Gauss-Bonnet-scalar interaction. Figures (7), (8) and (9) also reveal that it is true even for different forms of the potential. Such transient crossing independent of the value of $\\lambda$ also signals that it might be possible to carry out the same treatment even for a single scalar field model.    \\textbf{Acknowledgement}:{Acknowledgement is due to Dipartimento di Scienze Fisiche, Universit$\\grave{a}$  degli studi di Napoli, Federico II, for their hospitality, TRIL (ICTP) for financial assistance and Prof. Claudio Rubano for some illuminating discussion.}"
    },
    "0710/0710.1948_arXiv.txt": {
        "abstract": "{The primary proton spectrum up to 100 TeV has been investigated by balloon- and satellite-borne instruments. Above this energy range only ground-based air shower arrays can measure the cosmic ray spectrum with a technique moderately sensitive to nuclear composition. An array which exploits the full coverage approach at very high altitude can achieve an energy threshold well below the TeV region, thus allowing, in principle, the inter-calibration of the measured proton content in the primary cosmic ray flux with the existing direct measurements from balloons/satellites.  The capability of the ARGO-YBJ experiment, located at the YangBaJing Cosmic Ray Laboratory (4300 m a.s.l., Tibet, P.R. China), in selecting the surviving primary cosmic ray protons is discussed. A procedure looking for quasi-unaccompanied events with a very steep lateral distribution is also presented.}  \\begin{document} ",
        "introduction": "Despite large progresses in operating new multi-component Extensive Air Shower (EAS) experiments and in the analysis techniques to infer energy spectra and chemical composition, the key questions concerning the origin of the \"knee\" in the cosmic ray energy spectrum are still open. In particular, one of the most important questions to be solved is the position of the proton knee (see for example \\cite{pds} and reference therein). In fact, different experiments have claimed to see a proton knee at different energies: at 10 TeV the MUBEE collaboration \\cite{mubee}, at a few hundreds TeV the TIBET-AS$\\gamma$ experiment \\cite{tibet} and at a few PeV the KASCADE \\cite{kascade} and EAS-TOP \\cite{eastop} experiments. In addition, direct measurements carried out in the 100 TeV region by RUNJOB \\cite{runjob} and JACEE \\cite{jacee} do not exhibit any spectral break up to the highest measured energy ($\\sim$ 800 TeV). The knowledge of the primary proton spectrum is of great importance to understand the cosmic rays acceleration mechanisms and propagation processes in the Galaxy. A careful measurement of the proton spectrum over a wide range of primary energies (from 0.1 TeV to 10 PeV) using the same method is one of the main tasks of the future cosmic ray experiments.  The energy region up to about 100 TeV has been investigated by balloon- and satellite-borne instruments, and above these energies only ground-based air shower experiments may provide data which have however poor mass resolution. The detection of single hadrons at ground level, strongly related to the surviving primary cosmic ray protons, has been recognized for a long time as a method to investigate the proton spectrum over a large energy range and to derive the total inelastic cross section \\cite{yodh,eastop-psgl,kascade-psgl}.  The ARGO-YBJ experiment, an air shower array exploiting the full coverage approach at the YangBaJing Cosmic Ray Laboratory (Tibet, P.R. China, 4300 m a.s.l., 606 g/cm$^2$), offers the unique opportunity to investigate the cosmic ray spectrum over a large energy range (about 3 decades) because of its ability to operate down to a few TeV, thus overlapping the direct measurements, by measuring small size air showers (strip or digital read-out with high spatial granularity) and up to the PeV region by measuring the RPCs charge (analog read-out \\cite{argo_bigpad}). In this paper we will discuss how a direct measurement of the primary proton spectrum could be achieved up to the PeV energies with the ARGO-YBJ experiment. ",
        "conclusions": "The selection of surviving primary protons is possible with the ARGO-YBJ experiment by measuring events interacting a few g/cm$^2$ above the observational level. The granularity of the detector allows to investigate in great details the lateral and temporal features of hadronic jets just produced.  Calculations are in progress to evaluate the contribution of heavier primaries and to properly take into account the fluctuations in the shower development. A preliminary data analysis will be presented at the conference."
    },
    "0710/0710.3984_arXiv.txt": {
        "abstract": "We argue that giant flares in SGRs can be associated to the core conversion of an isolated neutron star having a subcritical magnetic field $\\sim 10^{12}$ G and a fallback disk around it. We show that, in a timescale of $\\lesssim 10^5$ yrs, accretion from the fallback disk can increase the mass of the central object up to the critical mass for the conversion of the core of the star into quark matter. A small fraction of the neutrino-antineutrino emission from the just-converted quark-matter hot core annihilates into $e^+e^-$ pairs above the neutron star surface originating the gamma emission of the spike while the further cooling of the heated neutron star envelope originates the tail of the burst. We show that several characteristics of the giant flare of the SGR 1806-20 of 27 December 2004 (spike and tail energies, timescales, and spectra) can be explained by this mechanism. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.4871_arXiv.txt": {
        "abstract": "We present the detection of low-amplitude, long-period $g$-modes in two individual sdBV stars which are known to be $p$-mode pulsators. Only few of these hybrid objects, showing both $p$- and $g$-modes, are known today. We resolve the $g$-mode domain in HS\\,0702+6043 and add HS\\,2201+2610 to the list of hybrid pulsators. To discover the low-amplitude $g$-modes, a filtering algorithm based on wavelet transformations was applied to denoise observational data. ",
        "introduction": "Subdwarf B stars are evolved and compact objects which are thought to be core helium burning. They populate the extreme horizontal branch (EHB) at effective temperatures of $20\\,000$\\,--\\,$40\\,000$~K and surface gravities $\\log (g/\\rm{cm\\,s^{-2}})$ between 5.0 and 6.2. Variable sdB stars (sdBVs) can be divided into the following subclasses: The $p$-mode pulsators show higher amplitudes and shorter periods at higher temperatures. \\citet{rl_kil97} discovered EC 14026-2647 as the class prototype. The $g$-mode pulsators have lower amplitudes and longer periods at lower temperatures and the class prototype PG\\,1716+426 was discovered by \\citet{rl_gre03}. Hybrid pulsators show both $p$- and $g$-modes. These are particularly exciting objects, since the two mode types probe different regions in the star. Both mode types are thought to be driven by the same $\\kappa$-mechanism \\citep{rl_cha97,rl_fon03,rl_jes06}. Figure 1 shows the location of some sdB pulsators in the $\\log g$\\,--\\,$T_{\\rm eff}$ diagram. \\begin{figure}[!ht] \\centering \\includegraphics[scale=0.45,angle=-270]{loggteff_proceedings_color.eps} \\caption{Location of some known sdB pulsators in the $\\log g$-$T_{\\rm eff}$ diagram. Evolutionary tracks \\citep{rl_dor93} for three stellar masses are indicated, as well as helium main sequence (HeMS) and zero age (ZAEHB) and terminal age (TAEHB) extreme horizontal branches. The positions (taken from E. M. Green and G. Fontaine; for a discussion see also Green et al., these proceedings) clearly reveal two instability regions. The three known hybrid pulsators are labelled 1, 2 and 3.} \\end{figure} ",
        "conclusions": ""
    },
    "0710/0710.3383_arXiv.txt": {
        "abstract": "As part of our on-going investigation of magnetic activity in ultracool dwarfs we present simultaneous radio, X-ray, UV, and optical observations of \\lsr\\ (M8.5), and simultaneous X-ray and UV observations of \\vb\\ (M8), both with a duration of about 9 hours. \\lsr\\ exhibits persistent radio emission and H$\\alpha$ variability on timescales of $\\sim 0.5-2$ hr.  The detected UV flux is consistent with photospheric emission, and no X-ray emission is detected to a deep limit of $L_X/L_{\\rm bol}\\lesssim 10^{-5.7}$.  The H$\\alpha$ and radio emission are temporally uncorrelated, and the ratio of radio to X-ray luminosity exceeds the correlation seen in F--M6 stars by $\\gtrsim\\!2\\times 10^4$.  Similarly, $L_{\\rm H\\alpha}/L_X\\gtrsim 10$ is at least 30 times larger than in early M dwarfs, and eliminates coronal emission as the source of chromospheric heating.  The lack of radio variability during four rotations of \\lsr\\ requires a uniform stellar-scale field of $\\sim 10$ G, and indicates that the H$\\alpha$ variability is dominated by much smaller scales, $\\lesssim 10\\%$ of the chromospheric volume.  \\vb, on the other hand, shows correlated flaring and quiescent X-ray and UV emission, similar to the behavior of early M dwarfs.  Delayed and densely-sampled optical spectra exhibit a similar range of variability amplitudes and timescales to those seen in the X-rays and UV, with $L_{\\rm H\\alpha}/L_X\\sim 1$. Along with our previous observations of the M8.5 dwarf \\tvlm\\ we conclude that late M dwarfs exhibit a mix of activity patterns, which points to a transition in the structure and heating of the outer atmosphere by large-scale magnetic fields.  We find that rotation may play a role in generating the fields as evidenced by a tentative correlation between radio activity and rotation velocity.  The X-ray emission, however, shows evidence for super-saturation at $v{\\rm sin}i\\gtrsim 25$ km s$^{-1}$, which could be the result of secondary effects such as inefficient heating or centrifugal stripping of extended coronal loops.  These effects may underlie the severe violation of the radio/X-ray correlation in ultracool dwarfs. Upcoming observations of L dwarfs will reveal whether these trends continue in substellar objects. ",
        "introduction": "\\label{sec:intro}  Over the past several decades, observations of magnetic activity in stars of spectral type F--M have uncovered a variety of correlations between the different activity indicators, as well as between the level of activity and properties such as stellar rotation and age. Coronal X-ray emission, transition region UV emission, chromospheric H$\\alpha$ emission, and non-thermal radio emission increase with both rotation and youth (e.g., \\citealt{kra67,pgr+81,sis+88}), and are temporally and energetically correlated in quiescence and during flares (e.g., \\citealt{cra82,gb93,bg94,hgr96}).  These observations have led to a general paradigm of magnetic field amplification at the shearing interface between the radiative and convective zones -- the so-called $\\alpha\\Omega$ dynamo \\citep{par55}.  This dynamo operates through a combination of stretching by differential rotation ($\\Omega$) and twisting by convective motions ($\\alpha$).  In quiescence, the dynamo-generated fields provide a heating source for the chromospheres and coronae primarily through magnetic waves and field dissipation on small scales (e.g., \\citealt{nu96,apr01}).  The interplay between the chromosphere, transition region, and corona is not fully understood, but is known to involve a complex combination of radiation, conduction, and mass flows.  In dMe stars it has been proposed that the transition region and chromosphere are instead heated by coronal X-rays, leading to the observed typical luminosity ratios of $L_{\\rm CIV}/L_X\\sim 10^{-1.5}$ and $L_{\\rm H\\alpha}/L_X\\sim 10^{-0.5}$, respectively (e.g., \\citealt{cra82,hj03}).  In addition to quiescent emission, sudden and large-scale dissipation of the field through large-scale magnetic reconnection may lead to particle acceleration and evaporation of the lower atmosphere into the chromosphere, transition region, and corona, giving rise to correlated radio, X-ray, UV, and H$\\alpha$ flares (e.g., \\citealt{neu68}).  The level of activity increases with faster rotation, but eventually saturates at $L_X/L_{\\rm bol}\\sim 10^{-3}$ and $L_{\\rm H\\alpha}/L_{\\rm bol}\\approx 10^{-3.5}$ for rotation periods of $\\lesssim 3$ d in F--M stars.  It remains unclear whether this trend reflects a saturation of the dynamo itself or secondary centrifugal effects such as coronal stripping or sweeping of the field toward the poles (e.g., \\citealt{vil84,ju99,ssv01}).  These activity trends continue to hold even beyond spectral type $\\sim {\\rm M3}$, where the stellar interiors become fully convective and the $\\alpha\\Omega$ dynamo can no longer operate.  Indeed, the level of X-ray and H$\\alpha$ activity peaks in mid M dwarfs at saturated levels \\citep{vw87,gmr+00,mb03,whw+04}.  This suggests that whatever dynamo mechanism operates in these low mass stars, it is already present in higher mass objects, and becomes increasingly dominant beyond spectral type M3.  However, the level of X-ray and H$\\alpha$ activity drops precipitously beyond spectral type M7, reaching mostly non-detectable levels by spectral type L5.  This decrease is accompanied by a clear transition from quiescent emission to flares in the few percent of active objects \\citep{rkg+99,gmr+00,rbm+00,lkc+03,whw+04}.  These trends point to a change in the dynamo mechanism, the field configuration, and/or the field dissipation process.  Radio observations, however, have uncovered a substantial fraction of active late M and L dwarfs ($\\sim 10\\%$ or higher; \\citealt{ber06}), which exhibit both quiescent and flaring emission \\citep{bbb+01,ber02,brr+05,bp05,ber06,ohb+06,adh+07,pol+07,hbl+07,aob+07,bgg+07}. Unlike the trend in H$\\alpha$ and X-rays, the level of radio activity appears to {\\it increase} with later spectral type \\citep{ber02,ber06}.  The radio emission is orders of magnitude brighter than expected based on the radio/X-ray correlations that are observed in stars down to spectral type M6, and requires magnetic fields of $\\sim 0.1-3$ kG with order unity filling factor sustained over timescales of at least several years \\citep{brr+05,ber06,bgg+07}. Thus, contrary to evidence from X-rays and H$\\alpha$, a substantial fraction of ultracool dwarfs continue to generate and dissipate magnetic fields.  In order to investigate the field generation and dissipation in detail, we have recently initiated a program of simultaneous, multi-wavelength observations of ultracool dwarfs.  Such observations are required to trace the temporal evolution of flares across the corona, transition region, and chromosphere, and to study the relation between particle acceleration and heating, particularly in the context of the known correlations in F--M stars.  In a previous paper (\\citealt{bgg+07}; hereafter Paper I) we presented observations of the M8.5 rapid rotator \\tvlm, which exhibited a wide range of temporally uncorrelated emission in the radio, H$\\alpha$, and X-rays.  These included quiescent radio emission from a large-scale field, radio flares from a tangled field component with $B\\approx 3$ kG, and periodic H$\\alpha$ emission matching the stellar rotation period with an inferred hot spot covering about $50\\%$ of the stellar photosphere. In addition, the quiescent H$\\alpha$ emission exceeded the X-ray luminosity by about a factor of two, likely ruling out coronal X-ray emission as the source of chromospheric heating.  Here we present observations of the M8 and M8.5 dwarfs \\vb\\ and \\lsr, both of which are known to exhibit magnetic activity.  We find substantially different behavior in each of these two objects, with correlated X-ray/UV flares and quiescent emission in \\vb, and uncorrelated radio/H$\\alpha$ emission in \\lsr.  Along with \\tvlm, the mixed behavior in late M dwarfs thus indicates a transition in the properties of the magnetic field and its dissipation in this spectral type range.  We show that rotation may play at least a partial role in explaining these trends. ",
        "conclusions": "\\label{sec:conc}  We presented simultaneous multi-wavelength observations of two late M dwarfs to trace the magnetic activity in their outer atmospheres.  In the case of \\lsr\\ (M8.5) we detect persistent radio emission and highly variable H$\\alpha$ emission.  No excess UV emission is detected.  Similarly, we detect no X-ray emission to a deep limit of $L_X/L_{\\rm bol}\\lesssim 10^{-5.7}$.  The ratio of radio to X-ray flux exceeds the average value measured in a wide range of active stars by more than 4 orders of magnitude.  Similarly, the ratio $L_{\\rm H\\alpha}/L_X\\gtrsim 10$ exceeds the values measured for M0--M6 dwarfs by at least an order of magnitude, and rules out heating of the quiescent chromosphere by coronal emission.  Temporally, we find no correspondence between the radio and H$\\alpha$ light curves.  This indicates that even if the source of chromospheric heating is magnetic reconnection, it occurs on sufficiently small scales that the overall radio emission does not change by more than $\\sim 10\\%$.  Alternatively, the radio emission produced in conjunction with the H$\\alpha$ variability may peak at a lower frequency than our 8.5 GHz observations.  In the former scenario, our limit on the radio variability during the brightest H$\\alpha$ episode can be interpreted as a limit of $\\lesssim 10\\%$ on the chromospheric volume involved in the variable H$\\alpha$ emission.  In the latter scenario, it is possible that the coincident radio emission is dominated by coherent emission at $\\nu\\approx 2.8\\times 10^6\\,B$, with $B\\lesssim 3$ kG.  With our measured rotation velocity of $v{\\rm sin}i\\approx 50$ km s$^{-1}$ for \\lsr, the radio observations sample nearly 4 full rotations.  The stability of the radio emission thus requires a large-scale and uniform magnetic field, with an inferred strength of $\\sim 10$ G.  The similarity in flux to a past observation suggests that the field is stable on year timescales.  This may be expected given the long convective turnover time for ultracool dwarfs, on the order of several years.  \\vb, on the other hand, exhibits bright and variable X-ray and UV emission, including a pair of flares and quiescent emission.  The behavior in both bands, which trace the corona and transition region, respectively, is highly correlated.  From the X-ray spectrum we infer a coronal temperature of about $3\\times 10^6$ K, with a possible second component at $\\sim 10^7$ K.  The optical emission lines exhibit a similar amplitude and timescale of variability to the X-rays and UV, with two distinct states of gradual and impulsive flaring.  The lack of causal relation between the two states points to emission from distinct plasma environments.  Moreover, the shallow Balmer decrement and \\ion{He}{1} emission in the impulsive Flare state require significantly denser and hotter plasma.  Our detailed study of \\lsr, \\vb, and \\tvlm\\ reveals mixed patterns of of behavior compared to the magnetic activity in F-M6 stars, and thus points to a transition in the atmospheric structure and heating process in the late M dwarf regime.  In particular, stellar-scale magnetic fields are present, as evidence by quiescent and uniform radio emission, but the coronae are generally weaker than in early M dwarfs.  The chromospheric activity declines less rapidly, and all three targets exhibit a similar range of highly variable H$\\alpha$ emission, $L_{\\rm H\\alpha}/L_{\\rm bol}\\sim 10^{-5}-10^{-4.5}$.  This range is about an order of magnitude less that the saturated value of mid M dwarfs.  In the standard picture of solar and stellar magnetic flares, the release of magnetic energy through processes such as reconnection leads to acceleration of electrons, and subsequently heating of the chromosphere, transition region, and corona through evaporation of the lower atmosphere.  In quiescence, the chromospheres of M0--M6 dwarfs may instead be heated by coronal X-ray emission \\citep{cra82}, as evidenced by the typical observed ratios of $L_{\\rm H\\alpha}/L_X\\sim 1/3$ \\citep{cra82,hgr96}.  Such a mechanism is clearly not at play in the case of \\lsr, and most likely \\tvlm, since the chromospheric emission is significantly more energetic than the non-detected corona.  It is possible instead that the source of chromospheric heating is continuous micro-flaring activity, leading to replenishment of chromospheric and transition region plasma by evaporation.  In this scenario the weak or absent coronal emission may be due to a limited temperature enhancement of $\\lesssim 10^5$ K, which is sufficient to produce chromospheric emission, but no significant soft X-ray emission.  The continuously variable H$\\alpha$ emission in \\lsr\\ appears to support this idea, and along with the lack of corresponding radio variability points to heating on scales much smaller than the stellar photosphere.  As we noted in \\S\\ref{sec:rot}, it is also possible that rapid rotation in ultracool dwarfs suppresses the X-ray emission through centrifugal stripping.  The apparent increase in radio activity with rotation indicates that the dynamo itself does not saturate, at least up to $v{\\rm sin}i\\sim 60$ km s$^{-1}$, which is roughly $1/3$ of the break-up velocity.  The effects leading to the apparent super-saturation in the X-rays may also underlie the severe violation of the radio/X-ray correlation since the level of violation appears to be correlated with rotation velocity (Figure~\\ref{fig:bgrot}).  We stress that the role of rotation is still speculative due to the small number of objects with X-ray, radio, and rotation measurements.  It is therefore crucial to increase the sample of ultracool dwarfs for which these quantities are measured.  To summarize, we conclude that late M dwarfs mark a change in the properties of the magnetic field and its dissipation, as well as the generation of high temperature plasma in the outer atmosphere.  The weak X-ray emission, both in relation to $L_{\\rm bol}$ and $L_{\\rm H\\alpha}$, points to inefficient heating of plasma to coronal temperatures, or alternatively to stripping of the most extended magnetic loops.  The lack of temporal correlation between the radio and H$\\alpha$ activity, however, indicates that even if atmospheric evaporation is taking place, it occurs on smaller physical scale than the overall magnetic field structure.  We end by noting that the use of simultaneous, multi-wavelength, and long duration observations to probe the magnetic activity of individual ultracool dwarfs provides a crucial complement to large single-band surveys.  In particular, the long duration and high time resolution data elucidate the range of timescales and amplitudes of gradual and impulsive flares, and can moreover uncover rotationally-induced modulations, as in the case of \\tvlm\\ (Paper I). In addition, the existence or absence of correlations between the various activity indicators (as in \\lsr\\ and \\tvlm), and their implications for the magnetic field and its dissipation, could not have been inferred from existing observations.  In the upcoming Chandra cycle we will expand our analysis with observations of several L0--L3 dwarfs.  These observations will reveal whether the transition in magnetic activity patterns seen in the late M dwarfs continues to lower mass objects, some of which are bona-fide brown dwarfs."
    },
    "0710/0710.2336_arXiv.txt": {
        "abstract": "{ We derive an expression for the entropy of a present dark matter halo described by a Navarro-Frenk-White modified model with a central core. The comparison of this entropy with the one of the halo at the freeze-out era allows us to obtain an expression for the relic abundance of neutralinos, which in turn is used to constrain the parameter space in mSUGRA models, when used with the WMAP observations.  Moreover, by joning these results with the ones obtained from the usual abundance criteria, we are able to clearly discriminate validity regions among $\\tb$ values of the mSUGRA model, by demanding both criteria to be consistent with the 2 sigma bounds of the WMAP observations for the relic density: $0.112<\\Omega h^2<0.122$.  We found that for $sgn~ \\mu =+$, small values of $\\tb$ are not favored; only for $\\tb \\sim 50$ are both criteria significantly consistent. The use of both criteria also allows us to put a lower bound on the neutralino mass, $m_{\\chi}\\geq151$GeV. } \\PACS{ {14.80.Ly}{Supersymmetric partners of known particles}   \\and {95.35.+d}{Dark matter} \\and {98.62.Gq}{Galactic halos} } % ",
        "introduction": "s^h-s^f=ln\\left[\\frac{n_{\\chi}^f}{n_{\\chi}^h}\\left(\\frac{x^f}{x^h}\\right)^{3/2}\\right]~, \\end{equation} where $x=m_{\\chi}/T$, $T$ is the temperature of the gas. A region that fits with the conditions associated with the intermediate scale is the central region of halos ($~10 pc^3$ within the halo core); evaluating the thermodynamical quantities at this region, using equation (\\ref{entropy}) and some extra assumptions (conservation of photon entropy), it is possible to construct a theoretical estimate for $s^h$ that depends on the nature of neutralinos ($m_{\\chi}$ and $\\langle\\sigma v\\rangle$), initial conditions (given by $x^f$), cosmological parameters ($\\Omega_{\\chi}$, the Hubble parameter, $h$) and structural parameters of the virialized halo (central values for temperature and density); for details of these and the following, see section IV of \\cite{Cabral-Rosetti:2004kd}.  An alternative estimate for $s^h$ can be made based on empirical quantities for observed structures in the present Universe using the microcanonical entropy definition in terms of phase space volume, but restricting this volume to the actual range of velocities accessible to the central particles. That is, restricting the escape velocity up to a maximal value $v_e(0)$ which is related to the central velocity dispersion of the halo ($\\sigma_h$) by an intrinsic parameter $\\alpha$: $v_e^2(0)\\sim\\alpha\\sigma_h^2(0)$. In a  recent work \\cite{nuevo}, we estimate the value of $\\alpha$ using an NFW modified model with a central core, and obtain $16.4<\\alpha <27.8$. The range of values allowed for this parameter is of the highest importance to determine the allowed region of the parameter space in the mSUGRA model as will be clear in the results presented on next section.  Equating the theoretical an empirical estimates for the entropy per particle it is obtained a relation for the relic abundance of neutralinos using the EC criterion\\footnote{This formula is a small modification to the one presented in \\cite{Cabral-Rosetti:2004kd}}: \\begin{equation}\\label{consistency} ln(\\Omega_{\\chi}h^2)=10.853-x^f+ln\\left[\\frac{(x^f\\alpha)^{3/2}m_{\\chi}}{f_g^*(x^f)}\\right] \\end{equation} where $f_g^*(x^f)$ is a function related to the degrees of freedom at the ``freeze-out'' time (see for example \\cite{gondolo}) that will be described elsewhere \\cite{nuevo}.  Modifying the program micrOMEGAs, we obtained the value for $x^f$ for any region of the parameter space and then $\\Omega_{\\chi}$ using Eq.~(\\ref{consistency}), therefore we were  able to discriminate regions that are consistent with the WMAP constraints for the EC criterion. ",
        "conclusions": "Using both the AC and EC that have been described in the preceeding sections, we can compute the total mass density of neutralinos present today and constrain the region in the mSUGRA parameter space where both criteria are fulfilled. Out of the five parameters, we will fix $\\mu >0$. Then our strategy is to explore wide regions for the values of the other four parameters, by means of a bi-dimensional analysis in the $m_0 - m_{1/2}$ plane for different fixed values of $A_0$ and tan$\\beta$. It is important to mention that we are not presenting an exhaustive search in all the possible regions, but we concentrated on those regions which have received more attention in the literature, see for example \\cite{belanger}.  In Fig.~(\\ref{tb10new}), we present the results for tan$\\beta=10$, for three values of $A_0$, namely $A_0=1000, 0, -1000$ GeV, shown in the top, middle and bottom panels respectiveley.  The yellow region (lower right corner) is where the $\\tilde{\\tau}$ is the LSP, the lighter and darker areas (red and blue for the online version in colours) define the allowed regions for the EC and AC respectively according to the observed DM density. The area of the EC region depends on the size of the interval of values of the parameter $\\alpha$,  the lower and upper bounds of $\\alpha$ determine the upper and lower boundaries of the EC region. As can be seen from the figure, the region where both criteria are fullfilled is very small, in fact, only for the highest values of $\\alpha$ there is an intersection between both criteria. This behavior holds for all values of $A_0$ in the interval $[-1000, 1000]$ GeV, here we are showing only the extreme and central values.  \\begin{figure}\\centering \\includegraphics[height=15cm, width=4.8cm, viewport=200 0 300 350]{mytb_10.pdf} \\caption{Allowed regions in the parameter space for AC (lighter gray/red) and EC (darker grey/blue) for the mSUGRA model with sgn$\\mu=+$, tan$\\beta=10$, and $A_0=1000$ GeV, top panel, $A_0=0$ GeV, middle panel, and $A_0=-1000$ GeV, bottom panel. The figures show the so called bulk and coannihilation regions. The yellow region shows where the stau is the LSP.} \\label{tb10new} \\end{figure}  Repeating the same procedure for larger values of tan$\\beta$, it is found that the intersection region for both criteria becomes larger, but it gets to be significant for the largest values of this parameter. This is clearly shown in Fig.~(\\ref{tb50new}), which is equivalent to Fig.~(\\ref{tb10new}), but for tan$\\beta=50$. In this case the bottom panel is for $A_0=-500$ GeV. It is clear from the figure that for these values of tan$\\beta$ both criteria are consistent, as shown by the large intersection area for values of $A_0$ in the interval $[0,1000]$ GeV. For negative values of $A_0$ the intersection region decreases with $A_0$, see the bottom panel of the figure. For even lower values of $A_0$ the intersection becomes insignificant.  \\begin{figure}\\centering \\includegraphics[height=15cm, width=4.8cm, viewport=200 0 300 350]{mytb_50.pdf} \\caption{The same as Fig.~(\\ref{tb10new}), but for tan$\\beta=50$, and now $A_0=-500$ GeV in the bottom panel.} \\label{tb50new} \\end{figure}  In Figs.~(\\ref{results2}) we present the same analysis but for the Focus Point region, and for the central value $A_0=0$.  The situation is consistent with the previous results, both criteria intersect for $\\tan \\beta=50$ and there is nearly no intersection for $\\tan \\beta=10$. \\begin{figure}\\centering \\includegraphics[width=8cm]{fig3.pdf} \\includegraphics[width=8cm]{fig4.pdf} \\caption{Allowed regions in the parameter space for AC (lighter gray/red) and EC (darker grey/blue) in the mSUGRA model with $A_0=0$, sgn$\\mu=+$, tan$\\beta=10$, top panel, and tan$\\beta=50$, bottom panel.  The region shows the so called Focus Point region.} \\label{results2} \\end{figure}  This analysis allows us to arrive to one of the main results of our work. The use of both criteria favours large values of tan$\\beta$.  In Figs.(\\ref{fig:Mhiggsvsm12}) and (\\ref{fig:chi-higgs1}) we show the allowed values for the LSP and the Higgs mass after constraining the parameter space with the abundance and entropy criteria.  As can be seen from Fig. (\\ref{fig:Mhiggsvsm12}), the current limit for the Higgs favour, combined with the AC and EC criteria, favours even more a large value of $\\tan \\beta$.  This, in turn, puts a constraint on the allowed SUSY mass spectra of the bulk and coannihilation regions: it gives an LSP of mass $m_{\\chi} \\sim 140$ GeV for $\\tb ~ 10$, and a lower bound for the LSP mass $m_{\\chi} \\gtrsim 150$ GeV for large $\\tb$.    \\begin{figure}\\centering \\includegraphics[width=8cm]{mhiggsvsm12.pdf}  \\caption{Allowed values for $M_{Higgs}$ as function of $m_{1/2}$. As can be seen from the figure, the present bound on the Higgs mass favours a large $\\tan \\beta$.} \\label{fig:Mhiggsvsm12} \\end{figure} \\begin{figure}\\centering \\centerline{\\includegraphics[width=8cm,angle=0]{fig6.pdf}} \\caption{The figure shows the LSP mass plotted versus the Higgs mass, points above the dashed line are the allowed values for the LSP. The points in blue correspond to the Focus Point region, the ones in red to the bulk and coannihilation regions.} \\label{fig:chi-higgs1} \\end{figure}  Further analysis, which is currently under way, is required to give more precise conclusions about this new method to constrain the parameter space of the mSUGRA model \\cite{nuevo}. \\vspace{0.3cm}  We acknowledge partial support by CONACyT\\\\ M\\'exico, under grants 32138-E, 34407-E and 42026-F, and PAPIIT-UNAM IN-122002, IN117803 and \\\\ IN115207 grants.  JZ acknowledges support from DGEP-UNAM and CONACyT scholarships."
    },
    "0710/0710.0619_arXiv.txt": {
        "abstract": "The large majority of EGRET point sources remain without an identified low-energy counterpart, and a large fraction of these sources are most likely extragalactic. Whatever the nature of the extragalactic EGRET unidentified sources, faint unresolved objects of the same class must have a contribution to the diffuse extragalactic gamma-ray background (EGRB).  Understanding this component of the EGRB, along with other guaranteed contributions from known sources, is essential if we are to use this emission to constrain exotic high-energy physics. Here, we follow an empirical approach to estimate whether a potential contribution of unidentified sources to the EGRB is likely to be important, and we find that it is. Additionally, we show how upcoming GLAST observations of EGRET unidentified sources, as well as of their fainter counterparts, can be combined with GLAST observations of the Galactic and extragalactic diffuse backgrounds to shed light on the nature of the EGRET unidentified sources even without any positional association of such sources with low-energy counterparts. ",
        "introduction": "\\label{intro}  The origin of the isotropic diffuse emission  (Sreekumar et al.\\ 1998) in energies between $100 {\\, \\rm MeV}$ and $20 {\\, \\rm GeV}$, detected by the {\\it Energetic Gamma-Ray Experiment Telescope} (EGRET) aboard the {\\it Compton Gamma-Ray Observatory}, remains one of the great unknowns of GeV-energy astrophysics. There are two major questions that still remain unanswered. 1.\\ How much of the diffuse emission detected by EGRET is, in fact, extragalactic, and what is the spectrum of this extragalactic background? And 2.\\ what fraction of the extragalactic emission can be attributed to each of the observationally established classes of gamma-ray emitters? Despite the associated uncertainties, these two issues are critical in any attempt to use gamma-ray observations to constrain exotic high-energy physics and yet-undetected classes of theorized gamma-ray emitters.  To answer the first question, a good understanding of the Galactic diffuse emission in the EGRET energy range is required. In order to obtain the intensity and spectrum of the extragalactic emission from the EGRET sky maps, the Galactic emission needs to be modeled and subtracted. This is made complicated by the discrepancy between the observed Milky Way spectrum in energies of $\\gtrsim 1 {\\rm GeV}$ and theoretical expectations (Hunter et al.\\ 1997). The observed spectrum is more shallow than model predictions based on the local demodulated cosmic ray spectrum. This deviation is known as the ``GeV excess'', and although various explanations have been proposed to account for part or all of the discrepancy, its origin remains a matter of debate (e.g., Pohl et al.\\ 1997; B\\\"{u}sching et al.\\ 2001; Strong et al.\\ 2004b; Kamae et al.\\ 2005; de Boer et al.\\ 2006; Strong 2006; Stecker \\etal 2007). As a result, determinations of the gamma-ray background using different Galactic emission models yield very different answers, both in intensity and in spectrum, despite being based on the same set of observations (e.g., Sreekumar \\etal 1998; Strong et al.\\ 2004a).  Attempts to answer the second question have been plagued by uncertainties in the cosmic density and evolution of the two established classes of extragalactic gamma-ray emitters: normal galaxies and blazars. Our observational knowledge of the gamma-ray properties of normal galaxies is very limited, as the sample of normal galaxies which have been observed in gamma rays consists of only the Milky Way and a marginal detection of the Large Magellanic Cloud (Sreekumar et al.\\ 1992; Hunter et al.\\ 1997; Hartman et al.\\ 1999). For this reason, the accuracy of theoretical estimates of the contribution of normal galaxies to the gamma-ray background is unavoidably at the order-of-magnitude level (e.g., Lichti et al.\\ 1978; Pavlidou \\& Fields 2002). But even in the case of blazars, which are by far the most numerous and best studied class of identified gamma-ray emitters, estimates of their contribution to the gamma-ray background vary from a few percent to $100\\%$ of the background originally reported by the EGRET team (e.g., Padovani \\etal 1993; Stecker \\& Salamon 1996a; Kazanas \\& Perlman 1997; Mukherjee \\& Chiang 1999; M\\\"{ucke} \\& Pohl 2000; Narumoto \\& Totani 2006; Dermer 2007).  The issue is further complicated by the existence of 171 sources which, at the time of publication of the 3rd EGRET catalog (hereafter 3EG; Hartman \\etal 1999), had not been positively or potentially associated with a lower-energy counterpart. These sources are collectively known as the unidentified EGRET sources, and they are more numerous than any established group of gamma-ray emitters. The distribution of these sources on the sky is such that a Galactic feature can be clearly distinguished - however a large number of sources are located away from the Galactic plane and the Galactic center\\footnote{Note however that the presence of sources at high latitudes does not, in itself, constitute proof that these objects are extragalactic (see e.g.\\ the Gould Belt discussion in Gehrels et al.\\ 2000).}. No more than a handful of sources can be associated with the Milky Way halo if the Milky Way is not many times brighter in gamma-rays than similar galaxies such as M31 (Siegal-Gaskins \\etal 2007). Hence, it is almost certain that the EGRET unidentified sources include a significant extragalactic component. Although the nature of these extragalactic sources remains unknown, it is reasonable to believe that there is a large number of fainter, unresolved objects of the same class, which are guaranteed to have {\\em some} contribution to the extragalactic gamma-ray background (EGRB). If these sources represent yet unidentified members of some known class of gamma-ray emitters (e.g.\\ blazars), then excluding them from any calculation of the contribution of the parent class to the diffuse background would lead to a significantly underestimated result due to an incorrect normalization of the bright-end of the gamma-ray luminosity function. If they represent an unknown class of gamma-ray emitters, then the contribution of their unresolved counterparts to the diffuse emission would significantly limit the diffuse flux left to be attributed to known classes, exotic processes, and  truly diffuse emission.  Hence, some contribution of unresolved unidentified sources to the EGRB is certain. It is therefore clear that until we either  answer the question of the nature of unidentified sources or derive some strong constraint indicating that a possible contribution of such unresolved objects to the EGRB would indeed be minor, we cannot hope to fully understand the origin of the EGRB.  Detailed predictions for the level of the unidentified source contribution to the EGRB involve important uncertainties: since no low-energy counterparts have been identified, we have no estimates of distance, and therefore no estimates of the gamma-ray luminosities of these sources. As a result, very few constraints can be placed on their cosmic distribution and evolution. However, very simple estimates can offer some guidance on whether ignoring this EGRB component may be a safe assumption to make.  For example, we can use the number of unidentified sources, the minimum flux resolvable by EGRET, and the observed intensity to the extragalactic gamma-ray background to place rough limits on the distance scales associated with resolved and unresolved unidentified sources so that unresolved sources do not overproduce the background. A population of unbeamed, non-evolving, single-luminosity sources uniformly distributed in Euclidian space are resolvable out to a distance $D$ by an instrument of number flux sensitivity $F_{\\rm min}$. The relation between $D$, $F_{\\rm min}$, and number luminosity $L$ in this case is simply $L=4\\pi D^2F_{\\rm min}$. If the instrument detects $N$ such sources, their number density $n_{\\rm source}$ can be estimated to be $n_{\\rm source} = 3N/4\\pi D^3$. If the same distribution of sources continues out to a distance $d>D$, the isotropic intensity (photons per unit area per unit time per unit solid angle) from the {\\em unresolved} members of this population will be $I_{\\rm unres} = \\int_D^d dI_{\\rm shell}$, where $dI_{\\rm shell} = (1/4\\pi) (n_{\\rm source}4\\pi r^2 dr) L/(4\\pi r^2)$ is the contribution from sources within a spherical shell located at a distance $r$ from the observer. Substituting our results for $n_{\\rm source}$ and $L$ above, and performing the integral, we obtain \\begin{equation} I_{\\rm unres} =  3NF_{\\rm min}(d-D)/4\\pi D\\,. \\end{equation} If we require that the unresolved emission from this population does not exceed the EGRB observed by EGRET ($I_{\\rm unres} \\leq I_{\\rm EGRB}$), we obtain  $I_{\\rm EGRB} \\geq 3NF_{\\rm min}(d-D)/4\\pi D$. Substituting $N\\sim 100$ for the population of extragalactic unidentified sources (see discussion in \\S \\ref{samples}), $F_{\\rm min} \\sim 10^{-7} {\\rm \\, ph \\, cm^{2} \\, s^{-1}}$ for the sensitivity of EGRET, and $I_{\\rm EGRB} \\sim 10^{-5}  {\\rm \\, ph \\, cm^{2} \\, s^{-1} \\, sr^{-1}}$ (Sreekumar et al. 1998), we obtain $d \\lesssim 6D$. This result implies that the largest distance out to which such a distribution of objects persists cannot be larger than a few times the distance out to which these objects are currently resolved, since in any other case these sources would overproduce the EGRB. It is therefore conceivable that the unidentified source contribution to the EGRB is significant, if not dominant.  In this work, we approach the problem from a purely empirical point of view. Instead of attempting to {\\em predict} the level of a diffuse component due to unresolved objects of the same class as unidentified EGRET sources, we try to assess whether there are any empirical indications that this component is, in fact, minor. We construct samples of unidentified sources which, based on their sky distribution,  are likely to consist mostly of extragalactic objects. Under the assumption that the majority of these sources can be treated as members of a single class of gamma-ray emitters, we seek to answer the following three questions:  (1) Is it likely that unresolved objects of the same class could have a significant contribution to the EGRB at least in some energy range?  (2) How would the collective spectrum of their emission compare to the measured spectrum of the EGRB deduced from EGRET observations?  (3) How are GLAST observations expected to improve our understanding of the nature of unidentified sources, based on the insight gained from our analysis?   This paper is structured as follows. In \\S \\ref{samples} we discuss the samples of resolved unidentified sources used in our analysis. Our formalism for constructing the collective emission spectrum of unresolved unidentified sources is presented in \\S \\ref{formalism}. Inputs from EGRET data used in our analysis are described in \\S \\ref{inputs}. In \\S \\ref{results} we describe our results, and in \\S \\ref{sec:GLAST} we discuss how these results are expected to improve once GLAST observations become available. Finally, we summarize our conclusions in \\S \\ref{sec:disc}. ",
        "conclusions": "\\label{sec:disc}  In this work, we have used a purely empirical model to explore the possibility that unresolved gamma-ray sources of the same class as unidentified EGRET sources have an appreciable contribution to the EGRB. We have argued that some unidentified source contribution to the gamma-ray background is guaranteed. We have additionally found that (1) if most high-latitude unidentified sources are assumed to be extragalactic, a one order of magnitude extension of the cumulative flux distribution to lower energies without breaks implies a significant contribution to the EGRB, at least at the lower part of the EGRET energy range; and (2) the spectrum of the cumulative emission of such unresolved sources would be very consistent with the observational determination of Strong et al.\\ (2004a)  of the EGRET EGRB within systematics.  We emphasize that the purpose of this study is not to estimate, even at the order-of-magintude level, the diffuse flux expected from extragalactic unresolved unidentified sources, but rather to place constraints on the flux distribution of these objects under the constraint that the measured background should not be exceeded in any energy interval. Our treatment is therefore different in purpose and spirit from most past work aiming to estimate the level of the contribution of different populations to the extragalactic gamma-ray background. Such work has been traditionally of two types. The first involves population models built from some understanding of the physics of the sources (such as, e.g., the models of M\\\"{u}cke \\& Pohl 2000 and Dermer 2007 in the case of blazars; Miniati 2003, Keshet et al.\\ 2003, Blasi, Gabici, \\& Brunetti 2007 in the case of clusters of galaxies; Lichti, Bignami \\& Paul 1978 and Pavlidou \\& Fields 2002 in the case of normal galaxies). The second involves models of the population luminosity function based on our knowledge of the source population from other wavelengths and normalized to fit EGRET data (such studies require a sample of detected, identified members, and are therefore applicable to blazars only, e.g. Chiang et al.\\ 1995, Stecker \\& Salamon 1996, Mukherjee \\& Chiang 1999, Narumoto \\& Totani 2006). In our case, lacking any knowledge of the physics of sources as well as of even the bright end of the luminosity function (since in absence of identifications no redshift and hence no luminosity can be derived for any of the sources), we have reversed the problem. Instead of using some assumed cosmic evolution for the sources to derive the expected level of contribution to the EGRB, as in all of the investigations mentioned above, we have used the tightest possible constraint on the allowable EGRB contribution of these sources (the observed EGRET background) to constrain the flux distribution of the unresolved sources. Our conclusions for the overall expected unresolved unidentified source intensity come from the observation that our constraints on the flux distribution are indeed very tight; hence, the contribution of the unidentified sources to the EGRB is likely to be high.  Our analysis suggests that  any model of the EGRB would be incomplete without some treatment of the unidentified source contribution. The results of our empirical model therefore motivate the pursuit of specific population models for the unidentified sources. Although such models involve  a more restrictive set of assumptions and increased uncertainties, they can provide more concrete predictions for the luminosity function of unresolved objects. Once a luminosity function model is assumed, and under the assumption that unresolved members of this class do indeed contribute most of the extragalactic diffuse emission, additional estimates can be made regarding the distance scales associated with this population.  The dynamical range in fluxes of the {\\em resolved} members of the population, as can be seen in  Fig. \\ref{lognlogs}, is smaller than an order of magnitude. For a single-luminosity population, this corresponds to a factor of 3 dynamical range in distance, which increases or decreases if the typical source luminosity increases or decreases with increasing redshift respectively. Such arguments can constrain not only the distance scales, but also the number density, intrinsic brightness, and evolution of the resolved and unresolved objects, if in a model-dependent fashion. Additional constraints and predictions, once a luminosity function has been assumed, can be derived through the multi-messenger and multi-wavelength approach.  If the emission from extragalactic unidentified sources is primarily of hadronic origin, it will be accompanied by neutrinos at comparable fluxes, and that may have potentially observable consequences in the TeV range for future km$^3$ neutrino detectors such as IceCube and Km3Net. If on the other hand the gamma-ray emission is primarily leptonic, it will be accompanied by X-ray emission, and their extragalactic background flux in the X-ray band may provide an additional constraint. We will pursue such models and calculations in an upcoming publication.  In this work we have tried, where possible, to make assumptions, which, if anything, {\\em underestimate} the possible contribution of unresolved unidentified sources to the EGRB. An exception to this general trend is our working assumption that the majority of the resolved, high-latitude 3EG unidentified sources belong to a single class. It is conceivable that instead, the resolved unidentified sources are a collection of members of several known and unknown classes of gamma-ray emitters. In this case, it is still likely that the summed contribution of unresolved members of all parent classes to the diffuse background is significant. However, the construction of a single cumulative flux distribution from all sources and its extrapolation to lower fluxes is no longer an indicative test for the importance of such a contribution.  Although we have argued that the contribution of unresolved unidentified sources to the EGRB is likely to be important or even dominant, the cumulative emission spectrum derived here and presented in Fig.\\ \\ref{ress} is only an upper limit, as it was derived demanding that the observed EGRB is not exceeded at any energy above 100 MeV. The contribution of unresolved unidentified sources is further constrained by allowing for the presence of the guaranteed contributions of unresolved normal galaxies and blazars.  Finally, it is noteworthy that we have found no evidence for an inconsistency between the population properties of high-latitude unidentified sources and those of blazars. The presence of a yet-unknown population of extragalactic high-energy emitters among the high-latitude EGRET unidentified sources remains one of the most tantalizing possibilities in GeV astronomy and one of the most exciting prospects for the GLAST era. However, the presently available data on the spectral index distribution and the cumulative flux distribution of these sources (this work), as well as their variability properties (e.g., Nolan et al.\\ 2003), are consistent with their being members of the blazar population. If indeed a large number of members of the blazar class are present among the entragalactic unidentified sources, then this could have potentially serious effects on our understanding of the redshift distribution of resolved blazars and consequently of the blazar luminosity function."
    },
    "0710/0710.0796_arXiv.txt": {
        "abstract": "We present the result of a photometric and Keck-LRIS spectroscopic study of dwarf galaxies in the core of the Perseus Cluster, down to a magnitude of M$_{\\rm B}$ $= -12.5$. Spectra were obtained for twenty-three dwarf-galaxy candidates, from which we measure radial velocities and stellar population characteristics from absorption line indices. From radial velocities obtained using these spectra we confirm twelve systems as cluster members, with the remaining eleven as non-members. Using these newly confirmed cluster members, we are able to extend the confirmed colour-magnitude relation for the Perseus Cluster down to M$_{\\rm B}$ $= -12.5$. We confirm an increase in the scatter about the colour magnitude relationship below M$_{\\rm B}$ $= -15.5$, but reject the hypothesis that very red dwarfs are cluster members. We measure the faint-end slope of the luminosity function between M$_{\\rm B}$ $= -18$ and M$_{\\rm B}$ $= -12.5$, finding $\\alpha$ $= -1.26$ $\\pm$  0.06, which is similar to that of the field. This implies that an overabundance of dwarf galaxies does not exist in the core of the Perseus Cluster. By comparing metal and Balmer absorption line indices with $\\alpha$-enhanced single stellar population models, we derive ages and metallicities for these newly confirmed cluster members. We find two distinct dwarf elliptical populations: an old, metal poor population with ages $\\sim$ 8 Gyr and metallicities $[{\\rm Fe/H}]$ $<$ $-0.33$, and a young, metal rich population with ages $<$ 5 Gyr  and metallicities $[{\\rm Fe/H}]$ $>$ $-0.33$. Dwarf galaxies in the Perseus Cluster are therefore not a simple homogeneous population, but rather exhibit a range in age and metallicity. ",
        "introduction": "Faint, low mass galaxies are the most numerous galaxy type in the universe, and are thus fundamental in understanding galaxy formation. However, their low luminosities and low surface brightness make detailed studies of them difficult. Because dwarfs are so common, any galaxy evolution/formation theory must clearly be able to predict and describe the properties of these galaxies.  The luminosity functions of nearby galaxies in all environments reveal that dwarf galaxies are far more common than brighter galaxies.  Within galaxy clusters there furthermore appears to be an overdensity of dwarf galaxies when compared to the field. The origin of these extra dwarfs, or if this excess is even real, remains unknown. The cluster luminosity function itself is not universal, but is strongly dependent on environment \\citep{sab03}, with the luminosity function often steeper in the more diffuse outer regions of clusters than in the denser cluster cores. The luminosity function also depends on the individual cluster, with each one having a characteristic luminosity function. For example, the Fornax Cluster is compact and rich in early-type galaxies, with a flat luminosity function faint-end slope of $\\alpha = -1.1$ \\citep{mie}. The Virgo Cluster has a high abundance of spiral galaxies and has a steeper luminosity function with a faint end slope  of $\\alpha = -1.6$ (\\citealt{trenth02,sab03}).  We know that both Local Group (LG) dwarf elliptical and dwarf spheroidal galaxies display varying star formation histories, with metal poor populations as old as the classical halo globular galaxies, but with evidence for star formation as little as 2-3 Gyr ago. Some low-mass LG galaxies such as Sagittarius, contain surprisingly high metal rich stellar populations considering their luminosities, as is also seen in clusters of galaxies. Previous spectroscopic and ground based imaging has revealed that dwarfs in the core of clusters are not a simple homogeneous population, with cluster core galaxies fainter than M$_{\\rm B}$ $= -15$ containing a mix of metallicities and ages (\\citealt*{Po,RSMP,c1}).  The formation scenarios for dwarf galaxies fall into two categories. The first of these is that dwarfs are old, primordial objects (hierarchical model) (e.g. \\citealt{wf}). The second is that they have recently evolved or transformed from a progenitor galaxy population (e.g. \\citealt*{moore}). Within hierarchical models of galaxy formation, dwarf galaxies are formed from small density fluctuations in the early universe. The lowest mass systems (i.e. dwarf galaxies) form first, and massive galaxies are built from these in mergers. If star formation occurred soon after the gravitational collapse of initial density perturbations, dwarf galaxy halos would be amongst the first objects to form (e.g. \\citealt{ds}). Theory also predicts that such systems would form first in the densest environments, although cluster dEs would not necessarily contain the first stars formed in the universe. These halos could have formed their stars quite recently, if star formation was suppressed in some way \\citep{tully}.  Dwarfs in clusters, on the other hand, could also be the remnants of stripped discs or dwarf irregulars. Through galaxy harassment and interactions, these progenitors become morphologically transformed into dEs. Cluster galaxies become stripped of their interstellar medium, and become dynamically heated by high-speed interactions with other galaxies and the gravitational potential of the cluster. To compensate, the galaxy loses stars,  and over time the spiral can morphologically transform into a dE \\citep*{c3}. This model is supported by recent observations of embedded discs in dEs in the Virgo and Fornax clusters (\\citealt*{Bar,DeR}). Distant clusters at z $\\approx$ 0.4 are also filled with small spiral galaxies, but this population is largely absent in nearby clusters, where dwarf spheroidals make up the faint-end of the luminosity function \\citep{moore}.  Downsizing \\citep{DeL} is the observed trend that star formation occurs later, and over more extended timescales, in smaller galaxies. In this scenario, dEs and lower mass galaxies formed or entered clusters after the giant galaxies. Low mass galaxies on average end their star formation after the giant elliptical galaxies. For example, the faint end of the red sequence in clusters is not formed until z $<$ 0.8 \\citep{DeL}, implying the luminosity weighted stellar populations of lower mass galaxies are younger than the stars in the giant elliptical galaxies. This seems to contradict the hierarchical method of galaxy formation, where dwarf systems form first. Also, dwarf elliptical galaxies are preferentially found in dense cluster environments, with few examples of isolated dEs, again contradicting the hierarchical model of galaxy formation. However, some dEs in clusters contain old stellar populations, so it is likely that multiple formation methods are needed to explain the origin of cluster dwarfs.  The colour-magnitude relation (CMR) for cluster galaxies is well defined at bright magnitudes and forms a tight sequence. However, it is unclear what the shape of the red sequence is at fainter magnitudes \\citep{an}. \\citet*{c2} find that galaxies at the brightest 4 magnitudes of Perseus obey the CMR, with the fainter candidate members showing significant scatter from the relation. In contrast, \\citet{an} and others find no deviation from the CMR at fainter magnitudes in other clusters. Instead, the faint red sequence is an extrapolation of what is observed at brighter magnitudes. With spectroscopy, we can test this directly by finding confirmed dwarf members of the cluster to establish the true form of the colour magnitude relation.  Spectroscopy can also be used to measure the faint-end of the luminosity function in galaxy clusters, which is fundamental in describing the galaxy population. The luminosity function contains important information on the formation and evolution of galaxies, and at low redshifts contains the combined influence of the galaxy initial mass function, and the effects of any evolutionary processes that have taken place in the cluster since its formation.  In this paper, we use deep Keck spectroscopy to determine cluster membership for dwarfs at M$_{\\rm B}$ $<$ $-16$, and present the ages and metallicities of these galaxies based on the strengths of  Balmer and metal absorption lines in their spectra. We confirm that twelve dwarfs are Perseus Cluster members, while the remaining objects are background galaxies. Using these results, we examine the colour-magnitude relation for the Perseus Cluster down to M$_{\\rm B}$ $= -12.5$. We also measure the luminosity function faint-end slope $\\alpha$ based on our cluster membership results. Our main conclusion is that some dwarf galaxies in Perseus have old, metal poor populations, whilst others are younger, metal rich systems. This suggests that dwarf galaxies in the core of the Perseus Cluster are not a simple, homogeneous population, but require multiple scenarios for their formation.  This paper is organised as follows: in $\\S$ 2 we discuss the observations and the selection criteria for the dwarf galaxies, $\\S$ 3.1 identifies the cluster members, $\\S$ 3.2 presents the colour magnitude relation, and the central luminosity function is presented in $\\S$ 3.3. In $\\S$ 3.4 we derive luminosity weighted ages and metallicities for the newly confirmed cluster members. A discussion of the main results is presented in $\\S$ 4 and these results are summarised in $\\S$ 5. We assume the distance to the Perseus Cluster is 77 Mpc throughout this paper. ",
        "conclusions": "We have analysed Keck/LRIS spectra for twenty-three dwarf galaxy candidates in the Perseus Cluster, from which we have confirmed cluster membership for twelve systems based on radial velocities measured from absorption lines.  We extend the confirmed member colour-magnitude relation for Perseus down to M$_{\\rm B}$ $= -12.5$, finding that the slope of the colour-magnitude relation becomes bluer when the low-surface brightness dwarf galaxies are included. The fainter dwarfs also scatter more from the colour-magnitude relation, following the trend observed by \\citet{RS} for low-mass galaxies in Fornax and Coma. This scatter can be interpreted as a spread in the metallicity distributions of dwarf galaxies, which has been inferred for dwarf galaxies in other clusters by \\citet{RSMP} and \\citet{Po}.  After removing non-members from the $B$-band luminosity function for the Perseus Cluster we find a faint-end slope $\\alpha$ $= -1.26$ $\\pm$ 0.06, similar to the field. Previous studies of other galaxy clusters have found that the dwarf-to-giant ratio is a function of local projected galaxy density. Extending this study to the outer regions of Perseus would enable us to see how the luminosity function changes with the local galaxy density.  Colours, morphologies, and central surface brightnesses are not sufficient criteria to confirm cluster membership, as this work has shown. Cluster membership cannot be confirmed without spectroscopy, so the faint-end luminosity functions calculated for galaxy clusters where membership has not been confirmed spectroscopically are most likely artificially steepened by non-members resulting in higher dwarf-to-giant ratios.  By comparing the observed dwarf spectral absorption indices with population synthesis models of \\citet{tho}, we derive luminosity-weighted ages and metallicities for the dwarf galaxies. A range of ages is observed, ranging from older than 8 Gyr to younger than 5 Gyr. The metallicity distribution of the faint cluster members is also not that of a simple homogeneous population, with the younger galaxies typically having higher metallicities.  More observations of dwarf galaxies in rich, nearby cluster environments are required in order to help improve our understanding of the formation scenarios for cluster dwarfs. Further spectroscopic work with a larger sample would also allow better constraints on the ages and metallicities of cluster dwarf galaxies and would help with modelling the formation of such galaxies."
    },
    "0710/0710.2793_arXiv.txt": {
        "abstract": " ",
        "introduction": "SGRs are galactic X-ray stars that emit, during sporadic times of high activity, a large number of short-duration (around 0.1 s) bursts of hard X-rays (Duncan and Thompson, 1992). A SGR is thought to be a magnetar, being a strongly magnetized neutron star powered by a very strong magnetic field ($\\ge$ 10$^{15}$ Gauss). On 27 December 2004 a powerful burst of X- and $\\gamma$-rays from one of the most highly magnetized neutron stars (SGR 1806-20) of our Galaxy reached the Earth's environment (Hurley et al., 2005). The Solar system received a shock, which is thought to be due to a cataclysm in the magnetar that caused it to emit as much energy in two-tenths of a second as the Sun gives off in 250,000 years. The signature of this event on the Earth's magnetic field has not previously been investigated. Here, we present the first results of the magnetar footprints on magnetic data recorded by near-Earth satellites. The magnetar SGR 1806-20 is the third such event ever recorded along with two others that were noted in 1979 and 1998 (Mazets et al., 1979; Hurley et al., 1999).  Several properties of this magnetar flare are relevant to our study. Firstly, a precursor of $\\sim$ 1 s was observed 142 s before the flare, with a roughly flat-topped profile (Hurley et al., 2005). The intensity of the main initial spike saturated all X- and $\\gamma$-ray detectors. However, particle detectors on board of RHESSI and Wind spacecraft (Boggs et al., 2004; Mazets et al., 2004) were able to record reliable measurements. Several instruments designed for other purposes provided important information, as Geotail (Terasawa et al., 2005) and Cluster/Double star (Schwartz et al., 2005). The first spike was followed by a tail lasting 380 s, during which 7.56 s pulsations were clearly observed, by the $\\gamma$-ray detectors on board of RHESSI (Hurley et al., 2005).  Secondly, a disturbance of the Earth's ionosphere was simultaneously observed with the detection of the burst from SGR 1806-20 (Inan et al., 2005). This sudden ionospheric disturbance (SID) was recorded as a change in the signal strength from very low frequency (VLF) radio transmitters, being noticed by stations around the globe (Campbell et al., 2005). These changes in the radio signal strength were caused by X-rays arriving from SGR 1806-20, which ionized the upper atmosphere and modified the radio propagation properties of the Earth's ionosphere (see clearing house of SID data associated with SGR 1806-20 flare at http://www.aavso.org/observing/programs /solar/sid-sgr1806.shtml). One such observation of this ionospheric signature resides within a 21.4 kHz signal that originates in Hawaii and propagates along an ionosphere wave guide to Palmer Station, Antarctica (Inan et al., 2005). This wave guide is some $\\sim$ 10,000 km in path length (Inan et al., 2005). As explained above, this is not a direct radio detection of SGR 1806-20 (see also http://gcn.gsfc.nasa.gov/gcn3/2932.gcn3). Moreover, due to the sub-burst longitude and latitude (Inan et al., 2005) and to the geographical distribution of LF/VLF beacons and monitoring stations, this burst was not detected by active monitoring stations in Germany, Australia, or Canada (Campbell et al., 2005). Here, we note that ionospheric disturbances were also reported in the case of the magnetar observed in 1998 (Inan et al., 1999). In the case of the 1998 magnetar the flare illuminated the nightside of the Earth and ionized the lower ionosphere to levels usually found only during daytime. The magnetar responsible for the 2004 burst was about the same distance as the magnetar responsible for the 1998 burst, but within 5.25$^\\circ$ of the Sun as viewed from Earth. Therefore its $\\gamma$-rays arrived on the dayside of our planet. The 2004 flare changed the ionic density at an altitude of 60 km by six orders of magnitude (Inan, 2006). It is thus plausible that this change in the ionospheric conductivity can cause oscillating perturbations in the current-generated magnetic field.  The thrust of this study is to find signatures associated with the explosion of the magnetar SGR 1806-20 within satellite measurements of Earth's electromagnetic field. Currently, the Earth's electromagnetic field is monitored by a number of Low Earth Orbit (LEO) satellite missions. After the launch of \\O rsted satellite in 1999, the knowledge of the near-Earth electromagnetic field has been dramatically improved (Hulot et al., 2002; L\\\"uhr et al., 2002; Maus et al., 2002; Tyler et al., 2003; Balasis et al, 2004). Since 2000, \\O rsted, CHAMP and SAC-C satellites have offered a continuous flow of high quality magnetic field measurements. Additionally, the DEMETER satellite provides 1 Hz energetic electron detector data. Finally, let us note that all these LEO magnetic missions are flying between the Earth's surface, where the temporal variations of the magnetic field are continuously monitored by geomagnetic observatories, and the magnetosphere, where an in-situ investigation of the three-dimensional and time-varying phenomena is done by the four identical spacecraft of Cluster II mission. ",
        "conclusions": "The effect of the SGR1806-20 flare on the Earth's magnetic field was not large, but it was detectable. This first attempt to find a magnetar signature in the geomagnetic field clearly indicates that the high resolution CHAMP magnetic data are optimal to capture the extremely bright flare from SGR 1806-20. Indeed, during the first half of the decay phase of the flare a 7.5 s periodicity is observed in the magnetic field over a magnetically quiet period, near the South Pole at 400 km altitude. This observation can be explained by a mechanism through which the oscillating flux of ionizing $\\gamma$-rays could alter the ionospheric conductivity and hence cause oscillating perturbations in the current-generated magnetic field.  An attempt to verify this hypothesis for the two previously recorded giant flares was not possible since no magnetic satellite missions were operating in LEO at that time (i.e., in March 1979 and August 1998). Of course, there are many spacecraft carrying magnetometers within the Solar system, but very few near planetary ionosphere. For example, the wavelet analysis was performed on magnetometer data from the Cluster II mission that probes the Earth's magnetosphere. In order to be able to visualize, in the wavelet power spectrum graph, any significant disturbances of the magnetic field, the power spectral density of the signal was amplified by a factor of $2^6$ (in comparison to the corresponding spectral density values of the CHAMP data). Although there are some indications for a weak pulsation-like signal at $\\sim$ 8 s, the fact that this signal is almost two orders of magnitudes weaker than the one observed in CHAMP data favors the hypothesis of an ionospheric origin for the signature found in CHAMP data.  Furthermore, our analysis can be extended to 1 Hz magnetic data provided by ground-based magnetic observatories, but only a small number of them provide such high resolution sampling nowadays. Data provided by 12 Canadian observatories, for which 1 Hz values are available over the period we are interested in, were also analyzed. For five of these observatories, missing data or high-level noise, made it difficult to apply the wavelet technique. For the others, no conclusive evidence for a signature related to SGR 1806-20 exists.  Analyzing other magnetic data with such a powerful tool as wavelets techniques could be relevant for understanding the impact that giant flares have on the terrestrial and other planetary magnetic fields. However, the main difficulty in such studies is due to the availability and quality of magnetic data. For instance, wavelet analysis of Mars Global Surveyor mission magnetic measurements on 27/12/04 was not able to detect any of the magnetar features due to the inadequate sampling rate: only 3 s data are now available (Michael Purucker, pers. comm. 2005)."
    },
    "0710/0710.3795_arXiv.txt": {
        "abstract": "Massive black hole binary coalescences are prime targets for space-based gravitational wave (GW) observatories such as {\\it LISA}. GW measurements can localize the position of a coalescing binary on the sky to an ellipse with a major axis of a few tens of arcminutes to a few degrees, depending on source redshift, and a minor axis which is $2 - 4$ times smaller.  Neglecting weak gravitational lensing, the GWs would also determine the source's luminosity distance to better than percent accuracy for close sources, degrading to several percent for more distant sources.  Weak lensing cannot, in fact, be neglected and is expected to limit the accuracy with which distances can be fixed to errors no less than a few percent.  Assuming a well-measured cosmology, the source's redshift could be inferred with similar accuracy.  GWs alone can thus pinpoint a binary to a three-dimensional ``pixel'' which can help guide searches for the hosts of these events.  We examine the time evolution of this pixel, studying it at merger and at several intervals before merger.  One day before merger, the major axis of the error ellipse is typically larger than its final value by a factor of $\\sim 1.5-6$.  The minor axis is larger by a factor of $\\sim 2-9$, and, neglecting lensing, the error in the luminosity distance is larger by a factor of $\\sim 1.5-7$.  This large change over a short period of time is due to spin-induced precession, which is strongest in the final days before merger.  The evolution is slower as we go back further in time. For $z = 1$, we find that GWs will localize a coalescing binary to within $\\sim 10\\ \\mathrm{deg}^2$ as early as a month prior to merger and determine distance (and hence redshift) to several percent. ",
        "introduction": "\\label{sec:intro}  Among the most important sources of gravitational waves (GWs) in the low-frequency band of space-based detectors are the coalescences of massive black hole binaries (MBHBs).  Binaries containing black holes with masses in the range $10^4 - 10^7\\,M_\\odot$ are predicted to form through the hierarchical growth of structure as dark matter halos (and the galaxies they host) repeatedly merge; see, for example, {\\citet{svh07}} and {\\citet{mhsa07}} for recent discussion.  The {\\it Laser Interferometer Space Antenna} ({\\it LISA}) is being designed to have a sensitivity that would allow detailed measurement of the waves from these binaries.  ``Intrinsic'' parameters --- the masses and spins of the black holes which compose the binary --- should be determined with very high accuracy, with relative errors typically $\\sim 10^{-3}$ to $10^{-1}$, depending on system mass and redshift; see Lang \\& Hughes (2006, hereafter Paper I) for recent discussion.  By measuring an ensemble of coalescences over a range of redshifts, MBHB GWs may serve as a kind of structure tracer, tracking the growth and spin evolution of black holes over cosmic time.  ``Extrinsic'' system parameters, describing a binary's location and orientation relative to the detector, are also determined by measuring its GWs.  In Paper I, we showed that a binary's position on the sky can be localized at $z = 1$ to an ellipse with a major axis of a few tens of arcminutes and a minor axis a factor of $2-4$ smaller.  At higher redshift ($z = 3-5$), these values degrade by a factor of a few, reaching a few degrees in the long direction and tens of arcminutes to a degree or two in the short one.  We also found that, neglecting errors due to weak gravitational lensing, a source's luminosity distance typically can be determined to better $1\\%$ at low redshift ($z = 1$), degrading to several percent at higher redshift ($z = 3-5$).  The intrinsic ability of GWs to determine the distance to a coalescing binary is phenomenal.  Coalescing MBHB systems constitute exquisitely well calibrated distance measures, with the calibration provided by general relativity.  Unfortunately, this percent-level or better accuracy could only be achieved if we measured MBHB coalescences in an empty universe.  In our universe, weak lensing will magnify or demagnify the GWs, and we will infer a luminosity distance smaller (for magnification) or larger (for demagnification) than the true value.  This phenomenon affects all high-redshift standard candles.  Its impact on Type Ia supernovae in particular has been discussed in detail {\\citep{frieman97,hw98,h98}}.  It will not be possible to correct for this effect {\\citep{dhcf03}} since much of the lensing ``noise'' arises from structure on subarcminute scales that is not probed by shear maps (which map the distribution of matter on scales greater than $1^\\prime$ or so).  Since we will not know the extent of the magnification when we measure MBHB waves, we must simply accept the fact that lensing introduces a dispersion of several percent in determining the distance to these GW events (see, e.g., Wang et al. [2002] to compute this dispersion as a function of redshift).  When we quote distance measurement errors, we will typically quote only the intrinsic GW measurement error, neglecting lensing's impact.  When the intrinsic GW distance error is $\\lesssim 5\\%$, lensing will blur it to the several percent level.  Note that a source's redshift $z$ {\\it cannot} be directly determined using only GWs.  Gravitational wave measurements infer system parameters through their impact on certain dynamical timescales, such as orbital frequencies and the rate at which these frequencies evolve.  Since these time scales all suffer cosmological redshift, $z$ is degenerate with other parameters.  For example, any mass parameter $m$ is actually measured as $(1 + z)m$.  However, if the binary's luminosity distance is determined, its redshift can then be inferred by assuming a cosmography.  For most binaries, the redshift can be determined to several percent (with an error budget dominated by gravitational lensing\\footnote{At redshifts $z \\lesssim 0.3$, the error is actually dominated by peculiar velocity effects {\\citep{kfhm06}}; however, the event rate is probably negligible at such low redshifts.  As such, we will focus on gravitational lensing as the main source of systematic redshift error.}).  We thus expect that GW measurements will locate a binary to within a three-dimensional ``GW pixel'' which at $z = 1$ has a cross-sectional area of $\\sim 10^{-2}$ to $10^{-1}\\ \\mathrm{deg}^2$ and a depth $\\Delta z/z \\sim $ several percent.  It is anticipated that there will be great interest in searching the GW pixel for electromagnetic (e.g., optical, X-ray, radio) counterparts to MBHB GW events.  Finding such a counterpart would be much easier if galactic activity were catalyzed in association with the coalescence {\\citep{kfhm06}}.  The nature of that activity is likely to depend rather strongly on the mass of the coalescing system \\citep{dssch06}.  For example, {\\citet{an02}} predict strong outflows and galactic activity prior to the final black hole merger as the smaller member of the binary drives gas onto the larger member, consistent with the high-mass ($M_{\\rm tot} \\gtrsim 10^7\\,M_\\odot$) predictions of \\citet{dssch06}.  {\\citet{mp05}} describe an X-ray afterglow that would ignite after gas refills the volume that is swept clean by the coalescing binary; Dotti et al.\\ predict this outcome for smaller systems ($M_{\\rm tot} \\lesssim \\mbox{several}\\times 10^6\\,M_\\odot$).  Recent work by {\\citet{bp07}} suggests that the final burst of radiation from a coalescing binary (which can convert $\\sim 10\\%$ of the system's mass to GWs very suddenly) may excite radial waves, and consequently electromagnetic variability, in an accretion disk due to the quick change in the disk's Keplerian potential. Such a signature may be essentially mass independent.  On the other hand, the coalescence may be electromagnetically quiet, in which case we face the potentially daunting task of searching the three-dimensional pixel for galaxies with morphology consistent with a (relatively) recent merger, or that have a central velocity dispersion $\\sigma$ consistent with the inferred final black hole mass (assuming that the $M_{\\rm BH}-\\sigma$ relation [Ferrarese \\& Merritt 2000; Gebhardt et al. 2000] holds at the redshift of these sources, and so soon after merger).  If the host galaxy or some other counterpart can be identified, we could then contemplate combining GW information with electromagnetic data.  For instance, combining {\\it LISA} mass measurements with the luminosity of the counterpart may allow us to directly measure the Eddington ratio $L/L_{\\mathrm{Edd}}$ {\\citep{kfhm06}}.  Identifying a counterpart would also allow us to more accurately characterize the system.  For example, much of the intrinsic luminosity distance error is due to correlations between distance and sky position.  Finding an electromagnetic counterpart essentially determines a binary's position precisely, breaking those correlations.  Previous studies have found that intrinsic distance error can be reduced by almost an order of magnitude if the position is known {\\citep{h02,hh05}}.  (Lensing errors still dominate in such a case, so that the distance remains determined only at the few percent level.)  A counterpart may also make it possible to simultaneously determine a source's luminosity distance and redshift.  Such a ``standard siren'' (the GW analog of a standard candle) would very usefully complement other high-redshift standard candles {\\citep{hh05}}, such as Type Ia supernovae {\\citep{p93, rpk95, wgap03}}.  A direct measurement of redshift will also break the mass-redshift degeneracy more accurately than can be done with just the luminosity distance and some assumed cosmological parameters.  Breaking this degeneracy is critical when studying the growth of black holes with cosmic time {\\citep{h02}}.  Many analyses {\\citep{c98,h02,v04,bbw05,hh05}} have quantified how well {\\it LISA} can determine MBHB parameters, including sky position and distance, using maximum likelihood Fisher matrix estimation. Our results from Paper I, given earlier, include the effects of ``spin-induced precession'' --- precession of both the orbital plane and the individual spins of the black holes due to post-Newtonian spin-interaction effects.  A significant result from that analysis is that spin-induced precession improves sky position accuracy by about half an order of magnitude in each direction versus previous analyses. This result can be partially understood as due to the breaking of a degeneracy between position and orientation angles: thanks to precession, the binary's orientation evolves with time and can be untangled from sky position.  This effect was already known due to pioneering work by {\\citet{v04}}; by taking the analysis to higher order and considering a broader range of sources, we were able to show that this improvement held for essentially all astrophysically interesting MBHB sources.  The purpose of this paper is to examine the localization of MBHB systems more thoroughly, in particular how the GW pixel evolves as the final merger is approached.  Paper I only presented results for measurements that proceed all the way to merger.  It will clearly be of some interest to monitor potential hosts for the binary event some time before the merger happens; if nothing else, telescopes will need prior warning to schedule observing campaigns.  Understanding the rate at which localization evolves can also have an important impact on the design of the {\\it LISA} mission, clarifying how often it will be necessary to downlink data about MBHB systems in order to effectively guide surveys.  Our main goal is to understand for what range of masses and redshifts prior localization of a binary using GWs will be possible.  A previous analysis by Kocsis et al.\\ (2007b; hereafter K07) also examined this problem in great detail, but without including the impact of spin-induced precession.  One of our goals is to see to what extent precession physics changes the conclusions of K07.  We find that precession has a fairly small impact on the time evolution of the GW pixel except in the last few days before the final merger, at which point its impact can be tremendous.  Precession typically changes the area of the sky position error ellipse by a factor of $\\sim 3-10$ (up to $\\sim 60$ in extreme cases) in just the final day.  This is in accord with the predictions of K07 (and even earlier predictions by N. Cornish 2005, unpublished).  The structure of this paper is as follows.  First, in \\S\\ {\\ref{sec:background}}, we briefly review the basics of the MBHB gravitational waveform and the parameter estimation formalism that we use; this section is essentially a synopsis of relevant material from Paper I.  Section {\\ref{sec:intrinsicGW}} reviews the form of the GWs that we use in our analysis, while \\S\\ {\\ref{sec:extrinsicGW}} describes how those waves are measured by the {\\it LISA} constellation.  We describe the measurement formalism we use in \\S\\ {\\ref{sec:formalism}}.  In \\S\\ {\\ref{sec:review}}, we summarize our results from Paper I regarding the final localization accuracy that {\\it LISA} can expect to achieve.  We turn to a detailed discussion of the time evolution of the GW pixel in \\S\\ {\\ref{sec:results}}.  We begin by summarizing the key ideas behind the ``harmonic mode decomposition'' of K07 in \\S\\ {\\ref{sec:khmf}}.  This technique cleverly allows calculation of the GW pixel and its time evolution with much less computational effort than our method (albeit without including the impact of spin precession).  Unfortunately, we have discovered that some of the approximations used by K07 introduce a systematic underestimate of the final sky position error by a factor of $2 - 4$ or more in angle; the approximations are much more reliable a week or more prior to the black holes' final merger.  Modulo this underestimate, the K07 results agree well with a version of our code which does not include spin precession (particularly a week or more in advance of merger, when their underestimate is not severe).  K07 thus serves as a useful point of comparison to establish the impact of precession on source localization.  Section {\\ref{sec:timeevolve}} is dedicated to our results, including comparison to K07 when appropriate.  We find that all relevant parameter errors decrease slowly with time until the last day before merger, when they drop more dramatically.  This sudden drop is not found in K07, nor is it present in a variant of our analysis that ignores spin precession.  It clearly can be attributed to the impact of precession on the waveform.  Before this last day, the major axis is $\\sim 1.5-6$ times, the minor axis $\\sim 2-9$ times, and the intrinsic error in the luminosity distance $D_L$ $\\sim 1.5-7$ times bigger than at merger for most binaries (i.e., all except the highest masses).  Going back to one week (one month) before merger, these numbers change to $2-9$ ($4-11$) for the major axis, $3-14$ ($5-24$) for the minor axis, and $3-14$ ($5-18$) for the error in the luminosity distance.  As a result, for $z = 1$, most binaries can be located within a few square degrees a week before merger and $10\\ \\mathrm{deg}^2$ a month before merger.  The intrinsic distance errors are also small enough this early that $\\Delta z/z$ remains dominated by gravitational lensing errors of several percent.  Advanced localization of MBHB coalescences thus seems plausible for these binaries; the situation is less promising for sources at higher redshift.  As a corollary to our study of the time evolution, we also examine the sky position dependence of errors (\\S\\ {\\ref{sec:angdependence}}). The errors depend strongly on the polar angle with respect to the ecliptic, increasing in the ecliptic plane to as much as $35\\%$ over the median for the major axis, $85\\%$ over the median for the minor axis, and $15\\%$ over the median for errors in the luminosity distance.  The errors have a much weaker dependence on the azimuthal angle.  When we convert to Galactic coordinates, we find that the best localization regions appear to lie fairly far out of the Galactic plane, offering hope that searches for counterparts will not be too badly impacted by foreground contamination.  We conclude this paper in \\S\\ {\\ref{sec:disc}}.  Besides summarizing our results, we discuss shortcomings of this analysis and future work which could help to better understand how well GWs can localize MBHB sources.  Throughout the paper we set $G = c = 1$; a convenient conversion factor in this system is $10^6 M_\\odot = 4.92$ s.  When discussing results, we always quote masses as they would be measured in the rest frame of the source.  These masses must be multiplied by a factor of $1 + z$ when used in any of the equations describing the system's dynamics or its GWs (particularly the equations of \\S\\ {\\ref{sec:background}}).  We convert between distance and redshift using a spatially flat cosmology with $\\Omega_\\Lambda = 0.75$ (and hence $\\Omega_m = 0.25$) and Hubble constant $H_0 = 75\\ {\\rm km}\\ {\\rm s}^{-1}\\ {\\rm Mpc}^{-1}$. ",
        "conclusions": "\\label{sec:disc}  As discussed at length in Paper I, accounting for the general relativistic precession of the angular momentum vectors in an MBHB system has a dramatic impact on what we can learn by observing the system's gravitational waves.  Spin-induced precession breaks degeneracies among different parameters, making it possible to measure them more accurately than they could be determined if precession were not present.  This has a particularly important impact on our ability to locate such a binary on the sky and to determine its luminosity distance --- the degeneracy between sky angles, distance, and orientation angles is severe in the absence of precession.  Our analysis shows that the improvement that precession imparts to measurement accumulates fairly slowly.  In using one code which includes the impact of spin precession and a second which neglects this effect, we find little difference in the accuracy with which GWs determine sky position and distance for times more than a few days in advance of the final merger.  The difference between the two codes grows quite rapidly in these final days.  In the last day alone, the localization ellipse area decreases by a factor of $\\sim 3-10$ (up to $\\sim 60$ in a few low-mass systems) when precession effects are included. Distance determination is likewise improved by factors of $\\sim 1.5-7$ in that final day.  Not all of the precession effects occur in the final days.  We saw in Figure \\ref{fig:axes_evol} that the long tail of small minor axes can be seen, to some degree, throughout the inspiral.  We could get lucky and find a binary with a very small value of $2b$ weeks before merger. But the improvement in the median that we found in Paper I appears to take effect only in the final days of inspiral.  Therefore, while precession may in fact help improve the {\\it final} localization of a coalescing binary by a factor of $\\sim 2-10$ in each direction, it will not be much help in {\\it advanced} localization of a typical binary.  Nevertheless, the pixel sizes that we find are small enough that future surveys should not have too much trouble searching the region identified by GWs, at least over certain ranges of mass and redshift. At $z = 1$, the GW localization ellipse is $\\sim 10\\ \\mathrm{deg}^2$ or smaller for most binaries as early as a month in advance of merger. (At high masses, the ellipse can be substantially larger than this a month before merger, but it shrinks rapidly, reaching a comparable size $1-2$ weeks before merger.)  This bodes well for future surveys with large fields of view that are likely to search the GW pixel for counterparts.  In addition, GWs determine the source luminosity distance so well that the distance errors we find are essentially irrelevant --- gravitational lensing will dominate the distance error budget for all but the highest masses.  As redshift increases, the GW pixel rapidly degrades, particularly for the largest masses.  Let us adopt $10\\ \\mathrm{deg}^2$ (the approximate LSST field of view) as a benchmark localization for which counterpart searches may be contemplated.  At $z = 3$, this benchmark is reached at merger for almost the entire range of masses we considered.  As little as a day in advance of merger, however, some of the least massive and most massive systems are out of this regime.  One week prior to merger, the most massive systems are barely located at all (ellipses hundreds of square degrees or larger).  The intermediate masses do best, but even in their cases the positions are determined with $\\sim 10\\ \\mathrm{deg}^2$ accuracy no earlier than a few days in advance of merger.  The resolution degrades further at higher redshift.  At $z = 5$, systems with $M \\gtrsim 6 \\times 10^6\\,M_\\odot$ are not located more accurately than $\\sim 30\\ \\mathrm{deg}^2$ even at merger.  Smaller systems are located within $\\sim 10\\ \\mathrm{deg}^2$ at merger, but very few are at this accuracy even one day in advance of merger.  The luminosity distance errors also increase, so much that they exceed lensing errors a few days to a week before merger at $z = 3$, and only a day before merger at $z = 5$.  This degradation hurts the ability to search for counterparts by redshift and subsequently use them as standard candles.  Our main conclusion is that future surveys are likely to have good advanced knowledge (a few days to one month) of the location of MBHB coalescences at low redshift ($z \\sim 1 - 3$), but only a day's notice at most at higher redshift ($z \\sim 5$).  This conclusion may be excessively pessimistic.  As mentioned earlier, recent work examining the importance of subleading harmonics of MBHB GWs is finding that including harmonics beyond the leading quadrupole has an important effect on the final accuracy of position determination {\\citep{arunetal,ts08}}.  For most masses, these analyses show a factor of a few improvement in position, comparable to the improvement that we find when spin precession is added to the waveform model.  For high-mass systems, the higher harmonics increase the (previously small) overlap with the {\\it LISA} band; consequently, the improvement can be much larger, up to 2 or 3 orders of magnitude in area. Since these two improvements arise from very different physical effects, it is likely that their separate improvements can be combined for an overall improvement significantly better than each effect on its own.  We plan to test this in future work (which is just now getting underway).  Finally, we have also studied the sky position dependence of {\\it LISA}'s ability to localize sources.  We have found that the regions of best localization lie fairly far out of the Galactic plane. However, as emphasized by N. Cornish (2007, private communication), a proper anisotropic confusion background might impact this dependence.  In our calculations, we have assumed an isotropic background, neglecting the likely spatial distribution of Galactic binaries.  Properly accounting for this background is likely to strengthen our conclusion that LISA's ability to ``see'' is best for MBHB sources out of the Galactic plane."
    },
    "0710/0710.3106_arXiv.txt": {
        "abstract": "Development of the Aarhus adiabatic pulsation code started around 1978. Although the main features have been stable for more than a decade, development of the code is continuing, concerning numerical properties and output. The code has been provided as a generally available package and has seen substantial use at a number of installations. Further development of the package, including bringing the documentation closer to being up to date, is planned as part of the HELAS Coordination Action. ",
        "introduction": "The goal of the development of the code was to have a simple and efficient tool for the computation of adiabatic oscillation frequencies and eigenfunctions for general stellar models, emphasizing also the accuracy of the results. Not surprisingly, given the long development period, the simplicity is now less evident. However, the code offers considerable flexibility in the choice of integration method as well as ability to determine all frequencies of a given model, in a given range of degree and frequency.  The choice of variables describing the equilibrium model and oscillations was to a large extent inspired by \\citet{Dziemb1971}. As discussed in Section~\\ref{sec:eqmodel} the equilibrium model is defined in terms of a minimal set of dimensionless variables, as well as by mass and radius of the model.  Fairly extensive documentation of the code, on which the present paper in part is based, is provided with the distribution package% \\footnote{The package is available at \\\\ {\\tt http://astro.phys.au.dk/$\\sim$jcd/adipack.n}}. \\citet{Christ1991} provided an extensive review of adiabatic stellar oscillations, emphasizing applications to helioseismology, and discussed many aspects and tests of the Aarhus package, whereas \\citet{Christ1994} carried out careful tests and comparisons of results on polytropic models; this includes extensive tables of frequencies which can be used for comparison with other codes. ",
        "conclusions": ""
    },
    "0710/0710.2954_arXiv.txt": {
        "abstract": "We have compiled the emission-line fluxes of \\ion{O}{1} $\\lambda$8446, \\ion{O}{1} $\\lambda$11287, and the near-infrared (IR) \\ion{Ca}{2} triplet ($\\lambda$8579) observed in 11 quasars. These lines are considered to emerge from the same gas as do the \\ion{Fe}{2} lines in the low-ionized portion of the broad emission line region (BELR). The compiled quasars are distributed over wide ranges of redshift (0.06 $\\le z \\le$ 1.08) and of luminosity ($-29.8 \\le M_{B} \\le -22.1$), thus providing a useful sample to investigate the line-emitting gas properties in various quasar environments. The measured line strengths and velocities, as functions of the quasar properties, are analyzed using photoionization model calculations. % We found that the flux ratio between the \\ion{Ca}{2} triplet and \\ion{O}{1} $\\lambda$8446 is hardly dependent on the redshift or luminosity, indicating similar gas densities in the emission region from quasar to quasar. On the other hand, a scatter of the \\ion{O}{1} $\\lambda$11287/$\\lambda$8446 ratios appears to imply the diversity of the ionization parameter. These facts invoke a picture of the line-emitting gases in quasars that have similar densities and are located at regions exposed to various ionizing radiation fluxes. The observed \\ion{O}{1} line widths are found to be remarkably similar over more than 3 orders of magnitude in luminosity, which indicates a kinematically determined location of the emission region and is in clear contrast to the case of \\ion{H}{1} lines. We also argue about the dust presence in the emission region since the region is suggested to be located near the dust sublimation point at the outer edge of the BELR. ",
        "introduction": "Active galactic nuclei (AGNs) are known to have strong emission lines of various ion species. Among them, the \\ion{Fe}{2} emission lines are one of the most prominent features in the ultraviolet (UV) to optical spectrum of many AGNs. They have long been hoped to provide significant information about some aspects of the AGNs and their host environments, e.g., energy budget of the line emission region % and the epoch of the first star formation in the host galaxies. The determination of the first star formation epoch is based on the standard theory that the iron enrichment in galaxies is delayed compared to that of the $\\alpha$-elements, such as magnesium, due to their different origins; Type Ia supernovae for iron and Type II supernovae for the $\\alpha$-elements \\citep{hamann93,yoshii98}. The delay corresponds to the difference in life-times of the progenitors of the two types of supernovae, and is estimated to be 0.3 -- 1 Gyr depending on the host galaxy environments \\citep{yoshii96, matteucci01}. Many observations have been devoted to the measurement of \\ion{Fe}{2}/\\ion{Mg}{2} line flux ratios in high-redshift quasars for this purpose over the last decade \\citep[e.g.,][]{elston94,kawara96,dietrich02,dietrich03,iwamuro02,iwamuro04,freudling03,maiolino03}. However, the observed \\ion{Fe}{2}/\\ion{Mg}{2} ratios show a large scatter, preventing a detection of any significant trend in the Fe abundance as a function of redshift. While a part of the scatter might be due to the difference in the intrinsic Fe/Mg abundance ratio, it is presumable that the diversity of the physical condition within the line-forming gas, affecting line emissivities, is the main cause \\citep{verner03, baldwin04}. In the same sense, a change of the observed \\ion{Fe}{2}/\\ion{Mg}{2} ratio as a function of redshift, if found, should be carefully examined to tell whether it reflects the abundance evolution or the systematic variation of the line emissivity. Thus establishment of a method to probe the line-emitting gas and estimate its physical parameters such as density and incident-ionizing radiation flux has been much awaited.   Unfortunately, the Fe$^+$ atom is characterized by an enormous numbers of possible electronic transitions, yielding the ``\\ion{Fe}{2} pseudocontinuum'' often observed in AGN spectra, which makes analysis of the observations extremely difficult from both the observational and theoretical viewpoints \\citep[e.g., ][hereafter T06]{tsuzuki06}. On the other hand, a promising approach is to use the emission lines emitted by simple atoms in the same region as the \\ion{Fe}{2} lines. The most potent lines are \\ion{O}{1} and \\ion{Ca}{2}, whose co-spatial emergence with \\ion{Fe}{2} is indicated by a resemblance of their profiles \\citep{rodriguez02a} and by a correlation between the line strengths \\citep{persson88}. Note that it is a natural consequence of similar ionization potentials of the relevant ions, i.e., 16.2 eV for \\ion{Fe}{2}, 13.6 eV for \\ion{O}{1}, and 11.9 eV for \\ion{Ca}{2}.    The first extensive study of the physical properties of \\ion{O}{1} emitting gases in AGNs was presented by \\citet{grandi80}, who observed the strongest \\ion{O}{1} line, $\\lambda$8446, as well as other weaker \\ion{O}{1} lines in Seyfert 1 (Sy1) galaxies. He found that \\ion{O}{1} $\\lambda$8446 lacks the narrow component that characterizes other permitted lines, and concluded that the line is purely a BELR phenomenon. He also suggested that \\ion{O}{1} $\\lambda$8446 is produced by Ly$\\beta$ fluorescence, which was later confirmed by the observation of I Zw 1, the prototype narrow-line Seyfert 1 (NLS1), by \\citet{rudy89}. \\citet{rodriguez02b} compiled the UV and near-IR \\ion{O}{1} lines, namely, $\\lambda$1304, $\\lambda$8446, and $\\lambda$11287, in normal Sy1s and NLS1s in order to investigate their flux ratios. They found that there must be an additional excitation mechanism for \\ion{O}{1} $\\lambda$8446---besides Ly$\\beta$ fluorescence---which they concluded is collisional excitation. As for the \\ion{Ca}{2} lines, extensive studies of Sy1s were presented by \\citet{persson88} and \\citet{ferland89}. We have observed seven quasars at redshifts up to $\\sim$1.0 for the purpose of obtaining the UV and near-IR \\ion{O}{1} and \\ion{Ca}{2} lines, thus extending the previous studies to include quasars at high redshifts. Photoionization model calculations were performed and compared with the observations, which led us to conclude that the \\ion{O}{1} and \\ion{Ca}{2} lines are formed in a gas with density $n_{\\rm H} = 10^{11.5} - 10^{12.0}$ cm$^{-3}$, illuminated by the ionizing radiation corresponding to the ionization parameter of $U = 10^{-3.0} - 10^{-2.5}$ \\citep[][the latter is referred to as Paper I hereafter]{matsuoka05, matsuoka07}.    Now that the general properties of the \\ion{O}{1} and \\ion{Ca}{2} emitting gas, as a whole, are thus being revealed, we should turn our eyes to the following question: how are these gas properties related to the quasar characteristics, such as redshift and luminosity? Here we present results of a compilation of the emission-line fluxes of \\ion{O}{1} $\\lambda$8446, \\ion{O}{1} $\\lambda$11287, and the near-IR \\ion{Ca}{2} triplet ($\\lambda$8498, $\\lambda$8542, and $\\lambda$8662) observed in 11 quasars. The quasars are distributed over wide ranges of redshift (0.06 $\\le z \\le$ 1.08) and of luminosity ($-29.8 \\le M_{B} \\le -22.1$), thus providing a useful sample to track the line-emitting gas properties in various quasar environments. The observational and theoretical data used in this work are described in \\S \\ref{sec:data}, and results and discussion appear in \\S \\ref{sec:results}. Our conclusions are summarized in \\S \\ref{sec:summary}. ",
        "conclusions": "}   \\subsection{A Picture of the Dust-Free Emission Region}   We plot the observed values of $n$(\\ion{Ca}{2})/$n$(\\ion{O}{1} $\\lambda$8446), \\ion{O}{1} $n$($\\lambda$11287)/$n$($\\lambda$8446), and EW (\\ion{O}{1} $\\lambda$8446) as functions of the redshift and of the $B$-band absolute magnitude $M_{B}$ in Figure \\ref{vsprop2}. One of the remarkable results is found in the top panels; the \\ion{Ca}{2}/\\ion{O}{1} $\\lambda$8446 ratio is hardly dependent on redshift or luminosity over the plotted range, while the ratio is predicted to be very sensitive to the density of the line-emitting gas in the photoionization models. We show the model predictions on the $n$(\\ion{Ca}{2})/$n$(\\ion{O}{1} $\\lambda$8446) -- \\ion{O}{1} $n$($\\lambda$11287)/$n$($\\lambda$8446) plane in Figure \\ref{oioi} ({\\it left}), as well as the observed values (see also Fig. 7 in Paper I\\footnote{ Note that Figure 7 in Paper I shows the predictions of Model 1, which is not adopted in this paper since the assumed microturbulent velocity, $v_{\\rm turb}$ = 0 km s$^{-1}$, could not reproduce the observed \\ion{Fe}{2} UV emissions (see \\S \\ref{data_model}). However, Model 3 adopted in this work predict very similar results to those shown in Figure 7 regarding the \\ion{O}{1} and \\ion{Ca}{2} emissions, while the whole pattern of contour is slightly ($\\sim$ 0.5 dex) shifted to the high-density regime; the best-fit parameters in Model 1 are ($n_{\\rm H}$, $U$) = (10$^{11.5}$ cm$^{-3}$, 10$^{-3.0}$). }). The gas density $n_{\\rm H}$ and the ionization parameter $U$ are changed around the reference grid point, ($n_{\\rm H}$, $U$) = (10$^{12.0}$ cm$^{-3}$, 10$^{-2.5}$), over 2 orders of magnitude in both parameters. It is clearly seen that the predicted $n$(\\ion{Ca}{2})/$n$(\\ion{O}{1} $\\lambda$8446) ratio increases monotonically with the increased gas density, and that all the observed values are marked with the density in the vicinity of the reference point, log $n_{\\rm H}$ = 12.0. Thus the similar density of the line-emitting gases are strongly indicated for the quasars distributed over these redshift and luminosity ranges.   The values of EW (\\ion{O}{1} $\\lambda$8446), both observed and calculated with the density of log $n_{\\rm H}$ = 12.0, are shown in Figure \\ref{oioi} ({\\it right}) as a function of the \\ion{O}{1} $n$($\\lambda$11287)/ $n$($\\lambda$8446) ratio. It shows that the predictions of EWs are also consistent with the observed data when log $n_{\\rm H}$ = 12.0 and a covering fraction (cf) of the line-emitting gas as seen from the central energy source of 0.2 -- 0.5 are assumed. Note that the covering fraction could be much smaller if we assumed oxygen overabundance relative to the solar value.    On the other hand, a scatter of the observed data in the \\ion{O}{1} $n$($\\lambda$11287)/ $n$($\\lambda$8446) axis seems to be related to the diversity of the ionization parameter (Fig. \\ref{oioi}, {\\it left}). Note that the diversity of other parameters, such as microturbulent velocity, gas column density, and chemical composition, could not explain this diagram since they significantly alter the $n$(\\ion{Ca}{2})/$n$(\\ion{O}{1} $\\lambda$8446) ratio, rather than \\ion{O}{1} $n$($\\lambda$11287)/ $n$($\\lambda$8446), and thus they are unable to explain the observed similarity of the former ratios (Paper I). It is also quite unlikely that the diversity of these parameter values is balanced out by the fine-tuned density in such a way that the $n$(\\ion{Ca}{2})/$n$(\\ion{O}{1} $\\lambda$8446) ratio is always kept to be $\\sim$1.0, unless these lines are the dominant heating or cooling sources of the emission region. As with the $n$(\\ion{Ca}{2})/$n$(\\ion{O}{1} $\\lambda$8446) ratio, \\ion{O}{1} $n$($\\lambda$11287)/ $n$($\\lambda$8446) is not clearly dependent on the redshift or luminosity (Fig. \\ref{vsprop2}, {\\it middle left and middle right}).    The above arguments invoke a picture of the line-emitting gases in quasars that have similar densities and are located at regions exposed to various ionizing radiation fluxes. It would be a consequence of the difference in distance to the central continuum source and/or in the intrinsic luminosity of the quasars. Note that it is in clear contrast to the well-studied case of H$\\beta$, whose emission regions in AGNs are known to be characterized by similar ionization parameters. In fact, reverberation mapping results for H$\\beta$ show the emission region size ($r$) -- luminosity ($L$) relation of $r \\propto L^{0.5}$, which is consistent with the constant ionization parameter regime \\citep{peterson02, bentz06}. Such a situation has long been expected in order to account for the remarkably similar AGN spectra over a broad range of luminosity, and was incorporated into the locally optimally emitting cloud (LOC) model suggested by \\citet{baldwin95}, that is, that the BELR is composed of gas with widely distributed physical parameters and each emission line arises from its preferable environment. On the other hand, the case in \\ion{O}{1} and \\ion{Ca}{2} lines apparently indicates that the location of the emission region is not radiation-selected.    In line with the above arguments, we found a clear difference of the velocity-luminosity relation between \\ion{O}{1} and H$\\beta$; the measured \\ion{O}{1} line widths are plotted versus $M_B$ in Figure \\ref{mag_width}, which shows that the \\ion{O}{1} line widths are remarkably similar, concentrated around 1500 -- 2000 km s$^{-1}$, over more than 3 orders of magnitude in the $B$-band luminosity. On the other hand, those for the hydrogen Balmer lines usually have a large scatter as shown by, e.g., \\citet{kaspi00}; their sample of 34 AGNs, spanning over 4 orders of magnitude in continuum luminosity, has the line widths of 1000 -- 10,000 km s$^{-1}$. Such a trend is also indicated by \\citet{persson88}, who reported that while the correlation between FWHM (\\ion{O}{1}) and FWHM (H$\\beta$) is good for the small FWHM regime, the \\ion{O}{1} lines grow systematically narrower than H$\\beta$ at large line width. \\citet{rodriguez02a} conducted a detailed study of the near-IR emission line profiles in NLS1s, and found that \\ion{O}{1}, \\ion{Ca}{2}, and \\ion{Fe}{2} lines are systematically narrower than the broad components of other low-ionization lines such as hydrogen Paschen lines and \\ion{He}{1} $\\lambda$10830. They also argued that these lines are produced in the outermost portion of the BELR, since their widths are just slightly broader than those of [\\ion{S}{3}] $\\lambda$9531 which they assumed is formed in the inner portion of the narrow emission line region (NELR). While the scattered widths of \\ion{H}{1} lines could be interpreted as a consequence of the radiation-selected locations of the emitting gases, regardless of the gas kinematics, the remarkable similarity of the \\ion{O}{1} line widths might imply the kinematically determined emission regions. In such a situation, the diversity of the ionization parameters, as discussed above, would be an inevitable result. As stated by \\citet{persson88}, there is clearly interesting information which could be deduced from studies of the \\ion{O}{1} and \\ion{H}{1} line profiles; especially, reverberation mapping of these \\ion{O}{1} and \\ion{Ca}{2} lines would be a powerful tool to reveal the underlying physics.     \\subsection{A Picture of the Dusty Emission Region}  The widely-accepted theory of the BELR describes its outer edge, where the \\ion{O}{1} and \\ion{Ca}{2} emission lines are likely to be formed, set by the dust sublimation \\citep[e.g.,][]{laor93, netzer93}. If we accept this picture, the dust grains are possibly mixed in the line-emitting gas and suppress the \\ion{Ca}{2} emission through the substantial Ca depletion. % Such a situation is in fact reported for the NELR by the absence or significant weakness of the observed [\\ion{Ca}{2}] $\\lambda$7291 line \\citep{kingdon95, villar97}. \\citet{ferguson97} presented the LOC model calculations of the narrow emission lines and argued that Ca is depleted relative to the solar value by factors of 3 -- 160.   It is hard to see an evidence of the dust presence in the emission region from our results, since Figure \\ref{oioi} appears to show that our dust-free models successfully reproduce the observations. However, it is noteworthy that % it is quite difficult to account for % the observed data at \\ion{O}{1} $n$($\\lambda$ 11287)/$n$($\\lambda$ 8446) $<$ 0.4 by the models with log $n_{\\rm H}$ = 12.0. The problem is that such small values of \\ion{O}{1} $n$($\\lambda$ 11287)/$n$($\\lambda$ 8446) could only be reproduced with the higher densities than log $n_{\\rm H}$ = 12.0 so that \\ion{O}{1} $\\lambda$8446 emission is exclusively enhanced by the collisional excitation, while such a dense gas produces intense \\ion{Ca}{2} emission that is much stronger than observed. One can clearly see this trend in Figure \\ref{oioi} ({\\it left}). If we assumed significant Ca depletion in the line-emitting gas, these difficulties are naturally resolved since it significantly suppresses the otherwise intense \\ion{Ca}{2} emissions. For example, the data point representing (log $n_{\\rm H}$, log $U$) = (13.0, $-$2.5) in Figure \\ref{oioi} ({\\it left}) could provide a plausible model for the observations at \\ion{O}{1} $n$($\\lambda$ 11287)/$n$($\\lambda$ 8446) $<$ 0.4 if the \\ion{Ca}{2} emission is suppressed by a factor of a few times 10.\\footnote{ The predicted EW for the model with (log $n_{\\rm H}$, log $U$) = (13.0, $-$2.5) is EW (\\ion{O}{1} $\\lambda$8446) = 30 \\AA, which is consistent with the observations if rather large covering factors of 0.5 -- 1.0 and/or oxygen overabundance relative to the solar value are assumed. }   The dust grains present in the line-emitting gas might also give the natural explanation to the observed lack of some UV emission lines relative to their optical or near-IR counterparts. \\citet{ferland89} mentioned the possibility of the dust survival in the BELR gas in order to explain the extreme weakness of the observed \\ion{Ca}{2} $\\lambda$3934 and $\\lambda$3639 lines relative to the near-IR triplet. At least a part of the long-standing \\ion{Fe}{2} UV/opt problem, in which the observed ratios of \\ion{Fe}{2} UV flux to the optical flux fall far below the photoionization model predictions \\citep[see, e.g.,][]{baldwin04}, could also be explained. However, it should be noted that the dust would affect the line formation processes in a very complicated manner, which should be precisely addressed when discussing the specific lines; for the \\ion{O}{1} and \\ion{Ca}{2} lines, one of the most apparent effects as well as the Ca depletion would be the destruction of Ly$\\beta$ photons which otherwise excite \\ion{O}{1} atoms, thus the suppression of the \\ion{O}{1} emissions. The resultant line flux ratios could be much different from those derived from the simple speculations.       }  We have compiled the emission-line fluxes of \\ion{O}{1} $\\lambda$8446, \\ion{O}{1} $\\lambda$11287, and the near-IR \\ion{Ca}{2} triplet ($\\lambda$8579) observed in 11 quasars. The quasars are distributed over wide ranges of redshift (0.06 $\\le z \\le$ 1.08) and of luminosity ($-29.8 \\le M_B \\le -22.1$), thus providing a useful sample to track the line-emitting gas properties in various quasar environments. The measured line strengths and velocities, as functions of the quasar properties, were analyzed with photoionization model calculations. Our findings and conclusions are as follows:  1. There is no sign of a significant change in the flux ratios of the \\ion{Ca}{2} triplet and \\ion{O}{1} $\\lambda$8446 over the redshift and luminosity ranges studied here. It strongly indicates similar gas densities in the line-emission region from quasar to quasar.  2. The observed scatter of the \\ion{O}{1} $\\lambda$11287/$\\lambda$8446 ratios appears to be related to the diversity of the ionization parameter, while the ratio is not clearly dependent on the redshift or luminosity. Combined with the similarity of the \\ion{Ca}{2}/\\ion{O}{1} $\\lambda$8446 ratios, it invokes the picture of the line-emitting gases in quasars that have similar densities and are located at regions exposed to various ionizing radiation fluxes.   3. The \\ion{O}{1} line widths are remarkably similar from quasar to quasar over more than 3 orders of magnitude in luminosity. It might imply a kinematically determined location of the line emission region and is in clear contrast to the case of \\ion{H}{1} lines, whose emission region is considered to be radiation-selected.   4. If we accept that the \\ion{O}{1} and \\ion{Ca}{2} emission lines are formed at the outer edge of the BELR and that the outer edge is set by the dust sublimation, the line-emitting gas is possibly mixed with the dust grains. In fact such a situation may better reproduce the observations than the dust-free case through the significant Ca depletion."
    },
    "0710/0710.5603_arXiv.txt": {
        "abstract": "We review the current status of accelerator, direct and indirect Dark Matter (DM) searches, focusing on the complementarity of different techniques and on the prospects for discovery. After taking a census of present and upcoming DM-related experiments, we review the motivations to go beyond an \"accelerator-only\" approach, and highlight the benefits of multidisciplinarity in the quest for DM. ",
        "introduction": "The evidence for non-baryonic dark matter is compelling at all observed astrophysical scales~\\cite{Bergstrom:2000pn,Bertone:2004pz}. Although alternative explanations in terms of modified gravity (see Ref.~\\cite{Bekenstein:2004ne} for a relativistic theory of the MOND paradigm) cannot be ruled out, they can hardly be reconciled with the most recent astrophysical observations~\\cite{Clowe:2006eq} without requiring additional matter beyond the observed baryons (e.g. Ref.~\\cite{Feix:2007zm} and references therein). It is therefore natural to ask {\\it how can we identify the nature of DM particles?}. We review here the main strategies that have been devised to attack this problem, namely accelerator, direct and indirect searches, focusing on the interplay between them and on their complementarity.  In fact, a tremendous theoretical and experimental effort is in progress to clarify the nature of DM, mostly devoted, but not limited, to searches for Weakly Interacting Massive Particles (WIMPs), that achieve the appropriate relic density by {\\it freezing-out} of thermal equilibrium when their self-annihilation rate becomes smaller than the expansion rate of the Universe. The characteristic mass of these particles is $\\cal{O}$$(100)$ GeV, and the most representative and commonly discussed candidates in this class of models are the supersymmetric neutralino, and the B$^{(1)}$ particle, first excitation of the hypercharge gauge boson, in theories with Universal Extra Dimensions.  A tentative census of present and upcoming DM experiments (WIMPs only) is shown in fig. 1. Shown in the figure are: two particle accelerators, viz. the Tevatron at Fermilab, and the upcoming Large Hadron Collider (LHC) at CERN; the many direct detection experiments currently taking data or planned for the near future, along with the names of the underground laboratories hosting them; high-energy neutrino telescopes; gamma-ray observatories; gamma-ray and anti-matter satellites. Light blue points denote gamma-ray experiments that are not directly related to indirect DM searches, as DM signals would be typically produced at energies below their energy threshold. Nevertheless, they may turn out to be useful to discriminate the nature of future unidentified high-energy gamma-ray sources. Three satellites are shown in the inset of figure 1: PAMELA, an anti-matter satellite that has already been launched and is expected to release the first scientific data very soon. ; GLAST, a gamma-ray satellite that is scheduled for launch in early 2008; and AMS-02, anti-matter satellite that should be launched in the near future.  We will discuss below the prospects for detecting DM with the various experiments shown in fig. 1, and we will focus our attention on the complementarity of the various detection strategies. The paper is organized as follows: we first discuss accelerator searches, and show that although the LHC has the potential to make discoveries of paramount importance for our understanding of DM, it may not be able to solve all problems. In Section 3 we discuss the information that can be extracted from direct detection experiments, in case of positive detection. Section 4 is then dedicated to indirect searches, and to the question of what astrophysical observations can tell us about the nature of DM, and how to combine this information with all other searches.  \\begin{figure*} \\centering \\includegraphics[width=\\textwidth]{dmmap2.eps} \\caption{2007 census of present and upcoming Dark Matter-related experiments. Black points denote the location of high energy neutrino telescopes; Dark-blue points are for gamma-ray Air Cherenkov Telescopes, while light-blue points are for other ground-based gamma-ray observatories. Red points are for underground laboratories hosting existing and upcoming direct detection experiments. Yellow points show the location of the Fermilab's Tevatron, and the upcoming Large Hadron Collider at CERN.} \\label{DMmap} \\end{figure*} ",
        "conclusions": ""
    },
    "0710/0710.2215_arXiv.txt": {
        "abstract": "{Sequences of Doppler images of the young, rapidly rotating late-type stars AB Dor and LQ Hya show that their equatorial angular velocity and the amplitude of their surface differential rotation vary versus time. Such variations can be modelled to obtain information on the intensity of the azimuthal magnetic stresses within stellar convection zones. We introduce a simple model in the framework of the mean-field theory and discuss briefly the results of its application to those  solar-like stars. } ",
        "introduction": "\\noindent Doppler imaging techniques allow us to measure the surface differential rotation in rapidly rotating late-type stars by tracking the longitudinal motion of starspots located at different latitudes (Collier Cameron 2007). Spe\\-ci\\-fi\\-cal\\-ly, the surface angular velocity $\\Omega$ at colatitude $\\theta$ is assumed to be given by a solar-like law: \\begin{equation} \\Omega(\\theta) = \\Omega_{\\rm eq} - d\\Omega \\cos^{2} \\theta, \\end{equation} where $\\Omega_{\\rm eq}$ is the equatorial angular velocity and  $d\\Omega$ is the pole-equator angular velocity difference. $\\Omega_{\\rm eq}$ and $d\\Omega$ can be measured by fitting the shear of starspots along sequences of photospheric images covering successive rotations. Alternatively, $\\Omega$ and $d\\Omega$ can be included as free parameters in a code that reproduces the line profile distortions due to starspots, thus obtaining their values by a suitable $\\chi^{2}$ minimization when a sufficiently long time series of line profiles is available.  The long-term monitoring of surface differential rotation of the two late-type dwarfs AB Doradus and LQ~Hydrae shows that their equatorial angular velocity and surface shear are functions of the time (see Donati et al. 2003a; Jeffers et al. 2007). It is interesting to note that the variations of $\\Omega_{\\rm eq}$ and $d\\Omega$ are compatible with an internal angular velocity uniform on cylindrical surfaces co-axial with the rotation axis (Donati et al. 2003a). ",
        "conclusions": "We modelled the observed time variation of the differential rotation in AB~Dor and LQ~Hya under the hypotheses that the azimuthal Maxwell stresses rule the changes of their surface shear and the internal angular velocity depends only on the distance from the rotation axis (Taylor-Proudman regime). We obtained that the average intensity of the mean field Maxwell stres\\-ses is $ |B_{s} B_{\\varphi} |  \\sim 0.03 - 0.14 $ T$^{2}$, implying azimuthal mean fields $B_{\\varphi} \\sim (3-10)$ T for $B_{\\rm s} \\sim 0.01$ T. Similar Maxwell stresses are obtained if the magnetic torque is assumed to be confined within the overshoot layer $0.67 \\leq (r/R) \\leq 0.71$ and no restrictions are imposed on the internal rotation law.  It is interesting to note that azimuthal magnetic fields of $3-10$ T, occupying a sizeable fraction of the convection zone, have been invoked to explain orbital period changes observed in late-type active binaries (Lanza, Rodon\\`o \\linebreak \\& Rosner 1998; Lanza \\& Rodon\\`o 2004; Lanza 2005, \\linebreak 2006b). An $\\alpha$-effect related to an instability of the magnetic field itself (e.g., magnetic buoyancy instability, Brandenburg \\& Schmitt 1998; or magneto-rotational instability, R\\\"udiger et al. 2007) seems to be necessary to produce such super-equi\\-par\\-ti\\-tion fields, possibly acting in combination with differential rotation. The energy dissipated by  turbulence, estimated according to standard mixing-length arguments, may exceed stellar luminosity in the case of the largest surface shear observed in LQ Hya. However, the thermal equilibrium of the convection zone can be  significantly affected only if those large shear episodes last more than $\\sim 10-20$ per cent of the time.  Note also that a mixing-length estimate for the turbulent viscosity may not be appropriate for a rapidly rotating star (see Lanza 2006a for details).  In the present work, we adopted a point of view analogous to that of Covas et al. (2005) who modelled the variation of stellar differential rotation considering only the \\linebreak torque exerted by the Lorentz force and neglecting the roles of meridional flow and $\\Lambda$-effect. As a matter of fact, it is difficult to evaluate the perturbations of the meridional flow and of the Reynolds stresses produced by the magnetic field because they depend critically on the approximations made in the treatment of stellar turbulence in a rotating star. Nevertheless,   alternative models for the variation of differential rotation have been investigated, such as those based on a time-dependent component of the meridional flow (Rempel 2006, 2007) or the magnetic quenching of the $\\Lambda$-effect (R\\\"udiger et al. 1986). The present approach can be generalized to obtain amplitudes of the perturbations of the corresponding terms in Eq.~(\\ref{tau}), but we shall not pursue this application here (see Lanza 2007).  Finally, it is interesting to note that some inference on the amplitude of variation of the surface shear in the case of very active stars can be obtained not only by means of Doppler imaging techniques, but also by  an appropriate ana\\-ly\\-sis of their long-term wide-band photometry  (e.g., \\linebreak Rodon\\`o et al. 2001; Messina \\& Guinan 2003). For example, in the case of LQ~Hya, it is worth comparing the Doppler imaging results by Donati et al. (2003a) with the photometrically determined seasonal rotation periods by Kov\\'ari et al. (2004)."
    },
    "0710/0710.0210_arXiv.txt": {
        "abstract": "The unified dark energy and dark matter model within the framework of a model of a continuous medium with bulk viscosity (dark fluid) is considered. It is supposed that the bulk viscosity coefficient is an arbitrary function of the Hubble parameter. The choice of this function is carried out under the requirement to satisfy the observational data from recombination ($z\\approx 1000$) till present time. ",
        "introduction": "Modelling of an accelerated expansion of the present Universe lies on the way of creation of phenomenological models which may explain the observational data on one set of parameters and compare them with predictions of the models on other set. For example, a theoretical model adapts for the correct description of the acceleration of the Universe in an accessible interval $z$. Further, the results of modelling are extrapolated for large $z$ which are not accessible for observations yet. The corresponding cosmological scenario defines growth of the large-scale structure which determines the present-day fluctuations of the microwave background radiation.  Certainly, the models under consideration should not contradict the available observational data within the framework of general relativity in a field of its applicability. For the specified purposes a number of cosmological models successfully applied in the past in the theory of the early Universe is used. These are cosmological models with various scalar and non-scalar fields filling the space together with cold dark matter (see, e.g., the reviews \\cite{Sahni}). Cosmological models of the present accelerated Universe within the framework of high-order theories of gravity (HOTG) are also quite popular \\cite{Carrol}.  A number of models have recently been suggested \\cite{Ren, Brevik,Odin} which describe the present Universe with use of models of a continuous medium in the presence of bulk viscosity. Consideration of effects of viscosity within the framework of HOTG was also carried out \\cite{Brevik1}. Note that such models were well-known in the theory of the early Universe (see, for example, \\cite{Murphy,Barrow}). In particular, in Ref. \\cite{Barrow} a few exact solutions with the constant bulk viscosity coefficient and with the bulk viscosity being an arbitrary power function of energy density were obtained. In Ref. \\cite{Ren} the model of viscous dark fluid is considered. The main result of this paper is the model with the constant bulk viscosity coefficient. The model fits the observational data on luminosity at an acceptable level. In Ref. \\cite{Brevik} the models both with the constant bulk viscosity coefficient and the bulk viscosity linearly proportional to the Hubble parameter are examined. The question about influence of viscosity on presence of a singularity in the model in the future (the so-called Big Rip) is investigated.  In this paper we consider a model of \"viscous dark fluid\" with the bulk viscosity coefficient $\\mu(H)$ which  depends on the Hubble parameter arbitrarily.  Unlike Ref. \\cite{Ren}, comparison of the model with the observational data is not restricted to the observational data on luminosity. The model is being compared with results of observations on change of the deceleration parameter $q$ and values of the Hubble parameter in the range $2>z>0$. It will be shown below that the model with the constant bulk viscosity coefficient does not provide a good description for $q(z)$ and $H(z)$ which follow from the observations. We propose such a dependence $\\mu(H)$ which is adequate to the mentioned observations. The proposed model is extrapolated for $z$ beyond the specified range $2>z>0$. ",
        "conclusions": "The bulk viscosity in our model is an example of a dynamic $\\Lambda$-term. However, our model is close to the $\\Lambda$CDM model.  As it was rightly noted in \\cite{Ren}, the model with the viscosity does not give possibility to divide the true dust filling the Universe, and dark matter generated by the bulk viscosity. That is why it is difficult to introduce a phenomenological equation of state $p=w \\varepsilon$ which is often used for interpretation of the observational data. Influence of the bulk viscosity on formation of the large-scale structure of the Universe demands special examination. But taking into account that the viscosity is being ''involved`` at rather small $z\\approx 2$, it is possible to expect its influence on the dynamics of galactic clusters only at later non-linear stage."
    },
    "0710/0710.5435_arXiv.txt": {
        "abstract": "{\\object{NRAO~150}  --a compact and bright radio to mm source showing core/jet structure-- has been recently identified as a quasar at redshift $z=1.52$ through a near-IR spectral observation.} {To study the jet kinematics on the smallest accessible scales and to compute the first estimates of its basic physical properties,} {we have analysed the ultra-high-resolution images from a new monitoring program at 86~GHz and 43~GHz with the GMVA and the VLBA, respectively. An additional archival and calibration VLBA data set, covering from 1997 to 2007, has been used.} {Our data shows an extreme projected counter-clock-wise jet position angle swing at an angular rate of up to $\\approx 11^{\\circ}/\\rm{yr}$ within the inner $\\approx 31$~pc of the jet, which is associated with a non-ballistic superluminal motion of the jet within this region.} {The results suggest that the magnetic field could play an important role in the dynamics of the jet in \\object{NRAO~150}, which is supported by the large values of the magnetic field strength obtained from our first estimates. The extreme characteristics of the jet swing make \\object{NRAO~150} a prime source to study the jet wobbling phenomenon.} ",
        "introduction": "\\label{int}  An increasing number of jets in active galactic nuclei (AGN) have been reported to show either regular or irregular swings of the innermost jet structural position angle in the plane of the sky (e.g., in \\object{OJ~287}, Tateyama \\& Kingham~\\cite{Tat04}; in \\object{3C~273}, Savolainen et al.~\\cite{Sav06}; in \\object{3C~345}, Lobanov \\& Roland~\\cite{Lob05}; in \\object{BL~Lac}, Stirling et al.~\\cite{Sti03}; in \\object{S5~0716+71}, Bach et al.~\\cite{Bac05}). Time scales between 2 and 15 years and structural position-angle oscillations with amplitudes from $\\sim 25^{\\circ}$ to $\\sim 45^{\\circ}$ are typical for the reported cases. We will call this phenomenon \\emph{jet wobbling} hereafter. Parsec scale AGN jet curvatures and helical-like structures also at larger distances from the central engine are also believed to be triggered by changes in direction at the jet ejection nozzle (e.g., in \\object{3C~84}, Dhawan et al.~\\cite{Dha98}).  The physical origin for the observed jet wobbling is still poorly understood. Among the various possibilities, regular precession of the accretion disk is frequently used for modeling at present. Most AGN precession models are driven either by a companion super-massive black hole or another massive object inducing torques in the accretion disk of the primary (e.g., Lister et al.~\\cite{Lis03}, for \\object{4C~+12.50}; Stirling et al.~\\cite{Sti03}, for \\object{BL~Lac}; Caproni \\& Abraham~\\cite{CapAb04} for \\object{3C~120}) or by the Bardeen-Peterson effect (e.g., Liu \\& Melia~\\cite{Liu02}; Caproni et al.~\\cite{Cap04}). However, other AGN scenarios that have yet to be explored extensively, such as the orbital motion of the jet nozzles (also involving binary systems)  or other kinds of more erratic disk/jet instabilities (e.g., similar to those thought to produce the quasi periodic oscillations [QPO] in X-ray binaries), can not be ruled out yet. Note that, in support of these erratic instabilities, it is still under debate whether the observed jet wobbling is strictly periodic or not (see Mutel \\& Denn~\\cite{Mut05} for the case of \\object{BL~Lac}).  There is still no paradigm to explain the phenomenon of jet wobbling in AGN, but it is rather likely that, as it is triggered in the innermost regions of the jets, it must be tied to fundamental properties of the inner regions of the accretion system. Hence, there is ample motivation to study the jet wobbling phenomenon to place our understanding of the jet triggering region and the super--massive accretion systems on firmer ground.  VLBI observations at millimetre wavelengths are a powerful technique to image the innermost regions of AGN jets -which are self-absorbed at longer wavelengths- with the highest angular resolutions; $\\sim 50\\,\\mu$as at 86~GHz (3.5~mm) and $\\sim 0.15$\\,mas at  43~GHz (7~mm). Here, we report the discovery of an extreme case of jet swing in the quasar \\object{NRAO~150}, through the first ultra-high-resolution VLBI set of images obtained from this source at 86~GHz and 43~GHz. \\object{NRAO~150} is a strong radio-mm source. At radio wavelengths, on VLBI scales, \\object{NRAO~150} displays a compact core plus a one-sided jet extending up to $r \\apgt 80$~mas with a jet structural position angle (PA) of $\\sim 30^\\circ$ (e.g., Fey \\& Charlot~\\cite{Fey00}). \\object{NRAO~150} was not detected by the early optical surveys,  most probably due to obscuration through the Milky Way (Galactic latitude $= -1.6^\\circ$). Almost 40~yr after its discovery at radio wavelengths (Pauliny-Toth, Wade \\& Heeschen~\\cite{Pau66}), \\object{NRAO~150} has been identified as a quasar at redshift $z=1.52$ through a near-IR spectroscopic project (Acosta-Pulido et al. in prep.).  Throughout this paper we assume $H_{\\circ} = 72$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{m}=0.3$, and $\\Omega_{\\Lambda}=0.7$. Under these assumptions, the luminosity distance of \\object{NRAO~150} is $d _{L}=11025$~Mpc, 1~mas corresponds to $8.5$~pc in the frame of the source, and an angular proper motion of 1~mas/yr translates into a speed of $69.4~c$. These parameters are used here for the first time to make quantitative estimates of the basic physical properties of the jet in \\object{NRAO~150}. ",
        "conclusions": "\\label{concl}  We have reported the results from the first 86~GHz and 43~GHz VLBI monitoring program of the recently identified quasar \\object{NRAO~150}. Our observations reveal {\\it a)} a large projected misalignment of the jet by $>100^{\\circ}$ within the inner 0.5~mas to 1~mas from the core, {\\it b)} an extremely fast counter-clockwise rotation of the projected jet axis at a rate of $\\sim6^{\\circ}/\\rm{yr}$ to $\\sim11^{\\circ}/\\rm{yr}$, {\\it c)} non-ballistic superluminal motions with mean speeds from 2.3~$c$ to 3.3~$c$ within the inner 0.5~mas from the core (deprojected distance $\\approx 31$~pc), {\\it d)} transverse (non-radial) speeds of 2.7~$c$, 1.4~$c$, and 2.0~$c$ for Q1, Q2 and Q3, respectively, {\\it e)} an extreme case of \\emph{superluminal non-ballistic jet swing}, and {\\it f)} the first approximations to quantitative estimates of the basic physical properties of the jet in \\object{NRAO~150}: $\\delta \\approx 6$, $B\\approx 0.7$\\,G, $\\gamma \\approx 4$, and $\\phi \\approx 8^{\\circ}$.  Whereas the ultimate origin of the jet swing must be an intrinsic change of the direction of the inner jet axis (either caused by changes produced at the injection region or by interaction with the medium surrounding the jet), possible causes for the non-ballistic nature of the emitting flow are either a small inertia of the jet compared to that of the impacted ambient medium, or by the fact that we are observing a jet instability propagating downstream. However, the superluminal non-ballistic motion of the jet features in \\object{NRAO~150} is perhaps too fast and systematic along tens of parsecs (deprojected, see above) to be induced solely by the ambient medium. In addition, the perturbation that must be produced either by the impact with such medium or by the changes at the injection region should imply the growth of disruptive instabilities (e.g., Perucho et al.~\\cite{Per06}), in contrast to the remarkable collimation of the jet up to the kiloparsec scale (see Fig.~1). Mizuno et al.~(\\cite{Miz07}) have shown that a magnetic field with the appropriate configuration helps keeping a jet collimated against the growth of instabilities. Hence, this possibility, together with the high value for the magnetic field intensity estimated in Section~\\ref{phys-par}, suggests that the magnetic field could play an important role in the dynamics of the jet in \\object{NRAO~150}.  It is still unclear whether the reported change of the direction of ejection in \\object{NRAO~150} is related to a regular (strictly periodic or not) behaviour or to a single event. This, together with the still unknown nature of the moving jet features (either propagating curvatures or shocks in a bent jet), does not allow us to specify the ultimate origin of this phenomenon in the source. This will be the matter of future studies. Nevertheless, the extreme characteristics of the jet swing make \\object{NRAO~150} a prime source for future studies of the origin of jet wobbling."
    },
    "0710/0710.0817_arXiv.txt": {
        "abstract": "The fast rotating star CU Virginis is a magnetic chemically peculiar star with an oblique dipolar magnetic field. The continuum radio emission has been interpreted as gyrosyncrotron emission arising from a thin magnetospheric layer. Previous radio observations at 1.4~GHz showed that a 100\\% circular polarized and highly directive emission component overlaps to the continuum emission two times per rotation, when the magnetic axis lies in the plane of the sky. This sort of radio lighthouse has been proposed to be due to cyclotron maser emission generated above the magnetic pole and propagating perpendicularly to the magnetic axis. Observations carried out with the Australia Telescope Compact Array at 1.4 and 2.5~GHz one year after this discovery show that this radio emission is still present, meaning that the phenomenon responsible for this process is steady on a timescale of years. The emitted radiation spans at least 1 GHz, being observed from 1.4 to 2.5~GHz. On the light of recent results on the physics of the magnetosphere of this star, the possibility of plasma radiation is ruled out. The characteristics of this radio lighthouse provides us a good marker of the rotation period, since the peaks are visible at particular rotational phases. After one year, they show a delay of about 15 minutes. This is interpreted as a new abrupt spinning down of the star. Among several possibilities, a quick emptying of the equatorial magnetic belt after reaching the maximum density can account for the magnitude of the breaking. The study of the coherent emission in stars like CU~Vir, as well as in pre main sequence stars, can give important insight into the angular momentum evolution in young stars. This is a promising field of investigation that high sensitivity radio interferometers such as SKA can exploit. ",
        "introduction": "CU Virginis (=HD124224) is an A-type magnetic chemically peculiar star (MCP) with a rotational period of 0.52 days, one of the shortest for this class of stars. As observed in all MCP stars, the variability of the light curve is correlated with the spectroscopic variations \\citep{deutsch, hardie} and with the effective magnetic field \\citep{borra80}. All those characteristics can be explained in the framework of the oblique rotator model, where the axis of the dipolar magnetic field is tilted with respect to the rotational one \\citep{babcock} and abundance of elements is not homogeneously distributed over the stellar surface. The observed variabilities are thus consequence of the stellar rotation.  \\begin{figure*} \\includegraphics[width=17cm]{spettri.ps} \\caption{Dynamical spectra of CU~Vir during the two days of observations (left panels: May 29, 1999) at 2.5~GHz (upper panels) and 1.4~GHz (lover panels). Some spectral channels have been removed, reducing the bandpass. The spectral resolution is 8~MHz; the spectra have been smoothed in time with a window 2 minutes wide. Strong enhancements of the radio emission are evident at two phases at 1.4~GHz, while only one peak is visible at 2.5~GHz. } \\label{spettri} \\end{figure*}   MCP stars are in general slow rotators in comparison with normal B and A-type main sequence stars. This behaviour can be explained as the result of the action of a magnetic breaking. But, at the present, only two MCP stars have been found to increase their rotational period; they are 56 Ari and CU Vir. While the former shows a continuous breaking down at the rate of few seconds per century \\citep{adelman}, CU Vir has been subject of an abrupt increase of the rotational period \\citep{pyper}. From the analysis of photometric light curves collected over 40 years, it seems that a change of period of about two seconds occurred abruptly in 1984. The mechanism responsible for this event is still under debate, and precise timing of the rotation is needed in order to detect any further slowing down of the star.  Radio emission has been observed in CU Vir \\citep{leo94}. The radio spectrum is quite flat and extends up to mm wavelengths \\citep{leo2004} with an high degree of circular polarization \\citep{leo96}. The variabilities of total intensity and polarization are both correlated with the rotation of the star, suggesting we are in presence of gyrosyncrotron emission from an optically thick source. The anisotropic stellar wind inferred from spectral line variations in MCP star \\citep{shore}, associated with the magnetic field justifies the radio emission.  \\citet{trig04} developed a three-dimensional model to explain the radio emission from MCP stars. The dipolar magnetic field interacts with the stellar wind, that can freely escape from the polar regions outside the Alfv\\'en surface (outer magnetosphere) but remains trapped in the equatorial belt (inner magnetosphere). Ionized particles flowing out in the transition region between outer and inner magnetosphere, called middle magnetosphere, stretch the magnetic field lines just outside the Alfv\\'en radius and open the field generating current sheets, where particles are accelerated up to relativistic energies. They eventually propagate back toward the stellar magnetic poles following the field lines and radiating for gyrosyncrotron process. As the magnetic field intensity increases traveling to the star, they are reflected back outward by the magnetic mirroring effect and are definitively lost from the magnetosphere. This model, used to explain the observed fluxes and variability of MCP stars (HD37479 and HD37017) has been successfully applied to CU Vir by \\citet{leto06}. On the basis of multiwavelengths radio observations, important physical parameters of the stellar magnetosphere, as the Alfv\\'en radius ($12-17\\,R_\\ast$) and the mass loss (about $10^{-12}M_\\odot \\rmn{yr}^{-1}$) have been inferred.  A further observational evidence supporting the picture outlined above is the discovery of coherent, highly directive, 100\\% polarized radio emission at 1.4~GHz \\citep{trig00}. The two peaks of radio emission have been observed at the rotational phases when the magnetic axis of the dipole is perpendicular to the line of sight. The two peaks have been observed in three observing runs spanning 10 days, indicating that the emission mechanism is persistent at least in timescales of weeks. This has been interpreted as Electron Cyclotron Maser Emission from a population of electrons accelerated in the current sheets at the Alfv\\'en point, that developed a loss cone anisotropy after the mirroring and masing in a direction almost perpendicular to the motion, so to the field lines, just above the magnetic pole.  In this paper we present radio observations of CU Vir at 1.4 and 2.5~GHz carried out with the aim to confirm the maser emission and to study its spectrum. The directivity of the radiation is used to check the rotational period as the beam point toward the Earth two times per rotation. ",
        "conclusions": "The observations of CU Vir carried out with the ATCA at 1.4 and 2.5 GHz confirm the presence of the coherent emission already reported by \\citet{trig00} after one year from the discovery, indicating that this is a steady phenomenon. The emitted radiation is visible only at rotational phases corresponding to the instant when the oblique axis of the magnetic dipole lies on the plane of the sky. This indicates the high directivity of this component of the radio emission. All those characteristics, and the fact that it is 100\\% right hand circularly polarized, are in agreement with the process of electron cyclotron maser from electrons accelerated in the current sheets out of the Alfv\\'en radius, propagating back to the stellar polar caps and reflected outward by the converging magnetic field. The lack of reflected electrons at small pitch angle, that fall in the stellar atmosphere, is the cause of the loss cone anisotropy that, in turn, generates this auroral radiation above the magnetic pole of the star. The polarization properties, the fact that the radiation is only right hand polarized, means that this process in efficient only above the north magnetic pole. This is possible as the magnetic field is not purely dipolar. The possibility of plasma radiation is ruled out since the high magnetic field strength in the region where the radiation is generated.  Since the magnetosphere rotates obliquely around the rotational axis, the cyclotron maser is visible only when it points toward the Earth, like a lighthouse. In this mode it is a good marker of the rotation of the star. The analysis done shows that the peaks are delayed of about 15 minutes with respect to the observations carried out one year before, indicating a possible change of the rotational period of the star, of the order of 1 second, occurred in the period 1998--1999. A similar change of period \\citep{pyper}, occurred around 1985, has been already reported. CU Vir is the unique single main sequence star with frequent abrupt spindown.  Different spinning down mechanisms are discussed: the possibility of a change of the moment of inertia of the star, the continuous spindown due to the wind flowing from the Alfv\\'en surface and the violent emptying of the inner magnetosphere. In this latter hypothesis the material accumulated in the closed field lines of the equatorial magnetic belt reaches a maximum density and opens the field lines in a violent event, releasing an angular momentum which may account the observed breaking.  Cyclotron maser emission from stars provides important information on the magnetospheres, as it has been observed in flare stars, in dMe and brown dwarfs. In the future, when high sensitivity radio interferometers such as SKA will allow to discover more and more radio lighthouse of the same kind of CU Vir, a new possibility to study with high precision the angular momentum evolution of young main sequence and pre main sequence stars will be opened."
    },
    "0710/0710.1811_arXiv.txt": {
        "abstract": "{}{X-ray Bright Optically Normal Galaxies (\\xbongs) constitute a small but not negligible fraction of hard X-ray selected sources in recent \\chandra\\ and \\xmm\\ surveys. Even though several possibilities were proposed to explain why a relatively luminous hard X-ray source does not leave any significant signature of its presence in terms of optical emission lines, the nature of \\xbongs\\ is still subject of debate. We aim to a better understanding of their nature by means of a multiwavelength and morphological analysis of a small sample of these sources. } {Good-quality photometric near-infrared data (ISAAC/VLT) of four low-redshift ($z=0.1-0.3$) \\xbongs, selected from the {\\rm HELLAS2XMM} survey, have been used to search for the presence of the putative nucleus, applying the surface-brightness decomposition technique through the least-squares fitting program GALFIT. } {The surface brightness decomposition allows us to reveal a nuclear point-like source, likely to be responsible of the X-ray emission, in two out of the four sources.  The results indicate that moderate amounts of gas and dust, covering a large solid angle (possibly 4$\\pi$) at the nuclear source, combined with the low nuclear activity, may explain the lack of optical emission lines. The third \\xbong\\ is associated with an X-ray extended source and no nuclear excess is detected in the near infrared at the limits of our observations. The last source is associated to a close (d$\\leq$ 1 arcsec) double system and the fitting procedure cannot achieve a firm conclusion.}{} ",
        "introduction": "Thanks to the \\chandra\\ and \\xmm\\ surveys, the hard X-ray sky is now probed down to a flux limit where the bulk of the X-ray background is almost completely resolved into discrete sources (Hasinger et al. 2001; Alexander et al. 2003; Bauer et al. 2004; Worsley et al. 2004, 2005). Extensive programs of multiwavelength follow-up observations showed that the large majority of hard X-ray selected sources are identified with Active Galactic Nuclei (AGN) spanning a broad range of redshifts and luminosities. At variance with optically selected quasars, X-ray selected AGN are characterized by a much larger spread in their optical properties, especially for what concerns the intensity of the emission lines. Indeed, a sizable fraction of relatively luminous X-ray sources hosting an active nucleus would not have been easily recognized as such on the basis of optical observations either because associated with very faint ($R >$ 24) counterparts (e.g., Fiore et al. 2003; Mignoli et al. 2004; Civano et al. 2005) or due to the lack of AGN emission lines in their optical spectra. The latter class of sources is variously termed as ``optically-dull\", ``optically normal\" or \\xbongs\\  (X-ray Bright Optically Normal Galaxies; Comastri et al. 2002). The common meaning of these definitions is that they lack evidence of accretion-driven activity in their optical spectra, in contrast with ``normal\" Seyfert galaxies and quasars. Their X-ray luminosities ($\\approx 10^{42}-10^{43}$ \\lum), X-ray spectral shape and X-ray-to-optical flux ratio (X/O\\footnote{Where X/O is defined as $X/O=\\log{\\frac{f_X}{f_R}}=\\log{f_X}+C+\\frac{R}{2.5}$.}$\\sim -1$) suggest AGN activity of  moderate strength. Originally discovered  in early {\\it Einstein} observations (Elvis et al. 1981) and named optically dull galaxies, the interest on the nature of these sources has gained a renewed attention after the discovery of several examples in XMM-{\\it Newton} and {\\it Chandra} surveys (Fiore et al. 2000; Comastri et al. 2002a,b; Georgantopoulos et al. 2005; Kim et al. 2006). Several possibilities were proposed in the literature in order to explain why a relatively luminous, hard X-ray source does not leave any significant signature of its presence in the form of emission lines.  A simple explanation favoured by Moran et al. (2002) and more recently by Caccianiga et al. (2007) for faint sources in the \\chandra\\ deep fields and brighter object in the \\xmm\\ XBS survey, respectively, is dilution by the host galaxy starlight. The combination of optical faintness and lack of strong emission lines in the observed wavelength range for distant \\chandra\\ sources  or the inadequate observing set-up among brighter nearby \\xmm\\ objects (Severgnini et al. 2003) may account for the \\xbong\\ properties. More in general, if the contrast between the host galaxy starlight and nuclear emission is high, AGN emission lines may easily be undetected. The physical reason may be ascribed to obscuration, most likely with a large covering factor, or to the fact that the lines are not efficiently produced by the central engine.  If \\xbongs\\ are merely obscured AGN, two hypotheses may be envisaged:  \\par $\\bullet$ In order to explain the multiwavelength  properties of the \\xbong\\ prototype PKS~0312018, also known as P3, Comastri et al. (2002) suggested heavy obscuration by Compton-thick gas covering almost 4$\\pi$ at the nuclear X-ray sources. In this way, no ionizing photons can escape to produce the narrow emission lines which are observed in ``normal\" Type 2 narrow-line AGN which are thought to have a lower covering fraction following the AGN Unified Scheme (but see Section \\ref{xray} for a recent re-analysis of the X-ray data of P3).  \\par $\\bullet$ According to a detailed multiwavelength analysis of  ``optically-dull\" galaxies in the \\chandra\\ deep fields, Rigby et al. (2006) conclude that extranuclear dust in the host galaxy plays an important role in hiding the emission lines.\\\\  Alternatively, \\xbongs\\ may be members of a class, or classes, of {\\it exotic} objects for which emission lines are either intrisically weak or absent:  \\par $\\bullet$ Radiatively Inefficient Accretion Flows (RIAFs) are expected at accretion rates well below those inferred for Seyferts and quasars. A distinctive property of low accretion-rate flows is that the standard Shakura-Sunyaev accretion disk is truncated at a relatively large inner radius. As a consequence, it cannot generate the ``big blue bump\" and enough UV photons to photoionize the line-emitting circumnuclear gas. The infalling gas is heated to high temperatures and emits a hard X-ray power-law by upscattering of low-energy seed photons. According to Yuan \\& Narayan (2004), the SED of source P3 could be reproduced by a RIAF model.  \\par $\\bullet$ \\xbongs\\ could be extreme BL Lac objects in which the featureless non-thermal continuum is much weaker than the host galaxy starlight. Following Fossati at al. (1998), \\xbongs\\ could belong to the low-luminosity tail of the blazar spectral sequence based on the anti-correlation between luminosity and frequency of the synchrotron peak.  \\par $\\bullet$ A highly speculative possibility is that of a transient AGN phenomenon in the process of tidally disrupting a star. If this were the case, the X-ray emission should be witnessing the transient accretion phenomenon (see Komossa et al. 2004 for extreme variability events in ROSAT observations; see also Gezari et al. 2006 for a luminous flare observed in the GALEX Deep Imaging Survey). The transient is most likely over in subsequent follow-up optical observations.  Finally, it should be noted that diffuse emission from a galaxy group, whose X-ray extended emission may have escaped detection in low signal-to-noise X-ray observations, is also possible and indeed observed in a few cases (Georgantopoulos et al. 2005).    While a unique solution may not necessarily hold for all the \\xbongs\\ observed in different surveys, they represent a useful benchmark for a better understanding of the AGN activity and, as such, deserve further studies. Ideally, one would need sensitive, high-spatial resolution, multiwavelength observations from radio to X-rays. As a first step, in the following we use good-quality photometric near-infrared data obtained with ISAAC at VLT of four low-redshift ($z=0.1-0.3$) \\xbongs, selected in the {\\rm HELLAS2XMM} survey to search for the presence of a putative nucleus which has escaped detection in the optical spectroscopy. The rather obvious advantage of near-infrared data is that the effects of dust reddening are minimized and that the nuclear emission in this band should rise more rapidly than the stellar light due to the reprocessing by hot dust. At the same time, the excellent quality of the near-IR images makes possible to apply a surface brightness decomposition technique, already successfully applied by several authors (e.g. S\\'anchez et al. 2004; Peng et al. 2006), to search for weak unresolved nuclear emission down to faint near-infrared magnitudes. We also discuss the broad-band properties of the four \\xbongs\\ using available multiwavelength data. Throughout the paper we assume a cosmological model with $H_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m$ = 0.3 and $\\Omega_{\\Lambda}= 0.7$. ",
        "conclusions": "We have presented a multiwavelength analysis of four \\xbongs\\ selected from the {\\rm HELLAS2XMM} survey. For these sources, deep near-infrared images taken with ISAAC at VLT, good-quality optical spectra and \\xmm\\ data are available.  Applying the morphological decomposition technique, we were able to detect the presence of a nuclear component in two out of the four sources (PKS~03120017 and PKS~03120018). There is no evidence of nuclear emission in the near-infrared in Abell~2690013; moreover, the X-ray appearance is consistent with an extended source. For source Abell~1835140, the near-infrared images reveal a complex morphology, where two sources are embedded in a common envelope. The main issue about the \\xbong\\ nature is whether they represent a truly distinct class or, rather, they are a mixed source population. Our results point towards the latter hypothesis.  The results regarding the nature of the \\xbongs\\  with nuclear component and the lack of optical emission lines can be summarized as follows. \\begin{itemize} \\item Source PKS~03120017 and PKS~03120018 are well described by a mildly obscured ($E(B-V)= 0.5-0.8$) optically weak nucleus, responsible for the X-ray emission, hosted by a bright galaxy (mag$_{K_s}^{nucl}$--mag$_{K_s}^{host}$ $\\sim 4$).  \\item The lack of optical emission lines cannot be attributed to observational limitations such as a inadequate observational setup or low signal-to-noise-ratio, at least for the objects in the present sample for which high-quality spectroscopic and photometric observations are available.  \\item We can safely discard for the two objects both a Compton-thick scenario (as found by Caccianiga et al. 2007 for a sample of elusive AGN in the XBSS) and an important blazar contribution. A RIAF solution seems to be supported by the estimated values of the Eddington ratios in the two objects and is not ruled out by the broad-band SED fitting, also because of the large number of free parameters in the RIAF model that can be tuned in order to reproduce the observed SEDs. A weak nuclear source, described by a standard accretion disk solution, but not powerful enough in the production of UV photons, would also provide an acceptable description of the observations.  \\item The presence of a thin nuclear gas and dust structure (as argued by Cocchia et al. 2007) covering 4$\\pi$ at the nuclear source, combined with the low level of activity of the BH, could prevent the ionization of the narrow-line regions and produce also the extinction we measured.\\\\  The analysis of already obtained VLT/VIMOS-IFU (Integral Field Unit) observations could help in the detection of the typical AGN emission lines down to a very faint flux limit and in the study of the gas kinematics.  \\end{itemize}"
    },
    "0710/0710.0671.txt": {
        "abstract": "Astronomy provides a laboratory for extreme physics, a window into environments at extremes of distance, temperature and density that often can't be reproduced in Earth laboratories, or at least not right away.  A surprising amount of the science we understand today started out as solutions to problems in astronomy.  Some of this science was key in the development of many technologies which we enjoy today.  This paper describes some of these connections between astronomy and technology and their history. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.3163_arXiv.txt": {
        "abstract": "This template file shows how to use the \\texttt{aipproc} class to produce a paper with the correct layout for \\emph{% AIP Conference Proceedings  8.5in x 11in double column}.  A full description of the features supported by the \\texttt{aipproc} class can be found in the \\texttt{aipguide.pdf} document accompanying the distribution.  Frequently asked questions can be found in the \\texttt{FAQ.txt} document. ",
        "introduction": "Infandum, regina, iubes renovare dolorem, Troianas ut opes et lamentabile regnum cruerint Danai; quaeque ipse miserrima vidi, et quorum pars magna fui. Quis talia fando Myrmidonum Dolopumve aut duri miles Ulixi temperet a lacrimis? Et iam nox umida caelo praecipitat, suadentque cadentia sidera somnos. ",
        "conclusions": ""
    },
    "0710/0710.5486_arXiv.txt": {
        "abstract": "{ Type~Ia supernovae are believed to be white dwarfs disrupted by a thermonuclear explosion.  Here we investigate the scenario in which a rather low-mass, carbon-oxygen (C~+~O) white dwarf accumulates helium on its surface in a sufficient amount for igniting a detonation in the helium shell before the Chandrasekhar mass is reached.  In principle, this can happen on white dwarfs accreting from a non-degenerate companion or by merging a C~+~O white dwarf with a low-mass helium one.  In this scenario, the helium detonation is thought to trigger a secondary detonation in the C~+~O core.  It is therefore called the ``double-detonation sub-Chandrasekhar'' supernova model.  By means of a set of numerical simulations, we investigate the robustness of this explosion mechanism for generic 1-\\msol\\ models and analyze its observable predictions.  Also a resolution dependence in numerical simulations is analyzed.  Hydrodynamic simulations of the double-detonation sub-Chandrasekhar scenario are conducted in two and three spatial dimensions.  The propagation of thermonuclear detonation fronts, both in helium and in the carbon-oxygen mixture, is computed by means of both a level-set function and a simplified description for nuclear reactions.  The decision whether a secondary detonation is triggered in the white dwarf's core or not is made based on criteria given in the literature.  In a parameter study involving different initial flame geometries for He-shell masses of 0.2 and 0.1~\\msol\\ (and thus 0.8 and 0.9~\\msol\\ of C~+~O), we find that a secondary detonation ignition is a very robust process.  Converging shock waves originating from the detonation in the He shell generate the conditions for a detonation near the center of the white dwarf in most of the cases considered.  Finally, we follow the complete evolution of three selected models with 0.2~\\msol\\ of He through the C/O-detonation phase and obtain \\nni-masses of about 0.40 to 0.45~\\msol.  Although we have not done a complete scan of the possible parameter space, our results show that sub-Chandrasekhar models are not good candidates for normal or sub-luminous type~Ia supernovae.  The chemical composition of the ejecta features significant amounts of \\nni\\ in the outer layers at high expansion velocities, which is inconsistent with near-maximum spectra.} ",
        "introduction": "\\label{sec:int}  One of the major uncertainties in modeling type~Ia supernovae (SN~Ia) originates from the unknown nature of the progenitor systems because neither observations nor theoretical models are yet conclusive \\citep{Branch1995,Ruiz-Lapuente2000,Livio2000,Nomoto2003,Han2004,Napiwotzki2005,Stritzinger2006,Parthasarathy2007}. Over the past years, most studies of thermonuclear supernova explosions focused on models in which a thermonuclear flame is formed by a runaway near the center of a white dwarf (WD), composed of carbon and oxygen, once it has reached the Chandrasekhar mass ($M_\\mathrm{Ch} \\sim 1.4~\\msol$) by accretion from a non-degenerate companion. Starting out in the sub-sonic deflagration mode of flame propagation, these models were shown to give rise to energetic explosions either in pure turbulence-boosted deflagrations \\citep{Reinecke2002b,Gamezo2003,Roepke2005,Roepke2006a,Schmidt2006a,Schmidt2006b,Roepke2007c} or with a delayed triggering of a supersonic detonation phase \\citep{Gamezo2004,Plewa2004,Roepke2007a}.  There seems to be reasonable agreement of such models with the gross features of observed SNe~Ia \\citep{Roepke2007c,Mazzali2007}.  However, it remains unclear whether the Chandrasekhar mass model can account for the full range of SN~Ia observations.  In particular, the mechanism of the sub-luminous events (\\object{SN~1991bg} being a prototypical example) remains a puzzle \\citep{Mazzali2007}.  \\citet{Stritzinger2006} even claim typical ejecta masses below $M_\\mathrm{Ch}$ for a sample of well observed SNe~Ia, with the trend that the low-luminosity explosions eject less mass.  Therefore, a long standing question is whether other progenitor channels contribute to the observed SN~Ia sample, and two alternatives have been suggested: the \\emph{WD merger} or \\emph{double degenerate scenario} (in contrast to single degenerate models, which consist of binaries with only one WD) and the \\emph{sub-Chandrasekhar model}.  In the present work we explore the latter.  The basic idea of the sub-Chandrasekhar model is that if the accretion rate onto a WD is lower than about $1$ - $4 \\cdot 10^{-8}~\\msol \\, \\mathrm{yr}^{-1}$ the accreted He (or the H processed to He) is not fused steadily into C and O.  Instead, after reaching a critical amount of He at relatively low densities, the He shell becomes unstable and detonates \\citep[cf.][and references therein]{Woosley1986}.  The detonation ignites most likely close to the bottom of the He shell, produces almost pure \\nni, and can occur long before the WD reaches the Chandrasekhar limit.  A second detonation may be triggered spontaneously when the He shell detonation shock wave hits the C~+~O core or with some delay after the shock has converged near the center of the WD (\\emph{double detonation} models). This way, in the sub-Chandrasekhar mass models the conditions for thermonuclear runaway are caused by the compression in the shock wave and not by the high degree of degeneracy as in the Chandrasekhar mass models.  Delayed double detonations have been studied extensively before. \\citet{Woosley1994}, and \\citet{Livne1990} carried out one-dimensional (1D) simulations and \\citet{Livne1990a,Livne1991} considered two-dimensional (2D) setups.  In 1D the He detonation ignites synchronously in a layer close to the bottom of the He shell and due to the spherical symmetry constraint the shock converges perfectly in the center.  However, this is not a very realistic model.  According to \\citet{Livne1990} an ignition most likely happens in a single point leading to an off-center convergence of oblique shock waves in the core.  Several other simulations predict successful directly ignited double detonations in two dimensions \\citep{Livne1997,Arnett1997,Wiggins1997,Wiggins1998}.  The smoothed particle hydrodynamics simulations by \\citet{Benz1997} and \\citet{Garcia-Senz1999} were carried out in 3D\\@.  \\citet{Livne1990a} first reported an increased probability of the direct core ignition if the He detonation does not happen directly at the core--shell interface but at a certain distance above it.  This way, the pressure jump can grow large enough before hitting the core.  In this work, the results of a (restricted) parameter study are presented investigating the possibility of triggering the second detonation in two and three dimensions.  In all our models the WD has a total mass of 1~\\msol, but the (C~+~O)-core and the He-shell masses differ.  We compute sequences of models with very different initial flame geometries in order to test the robustness of this explosion mechanism.  Since this latter question is the main focus of our paper, most of the simulations are stopped once the conditions for a detonation are matched.  However, for a few successful cases the energetics and nucleosynthesis of complete double detonations are computed.  In the following chapter the model details will be described.  Section~\\ref{sec:sim} shows the simulation results, which are summarized and discussed in the last part of the paper. ",
        "conclusions": "\\label{sec:con}  We have studied the sub-Chandrasekhar model of SNe~Ia by means of a series of two-dimensional and a few three-dimensional simulations with different initial conditions.  The numerical scheme used is based on the PPM algorithm, and the propagation of the detonation front is modeled applying the level set technique.  This novel implementation allowed for an accurate treatment of the hydrodynamic features of the sub-Chandrasekhar model (such as shock waves and thin detonation fronts) in multiple dimensions.  With our generic 1-\\msol\\ model we performed a parameter study involving different He masses and ignition geometries of the initial He detonation.  We find that a detonation in the He shell can clearly trigger a second detonation in the core and at least for the mass configurations studied, the double detonation seems to be a very robust process, which works without any ``fine-tuning'' of our model parameters.  In almost all of the simulations performed, the ignition conditions for a core detonation were reached at or near the center of the WD as a result of the convergence of the shock from the He shell detonation and in the few other cases the failure is likely a result of the finite numerical resolution.  The high maximum densities and temperatures that are needed for detonation ignition and the fact that the shock is not diluted too much when propagating inwards against the pressure gradient are made possible through a geometrical shock amplification effect that appears in spherical symmetry: As the surface of the shock decreases its strength has to increase.  The high compression that can be achieved theoretically in this way, especially if full convergence is reached and the shock is reflected (cf.\\ Sect.~\\ref{sec:res}), could even allow successful double detonations with considerably smaller He masses.  The shock amplification that is reached, however, turns out to be resolution dependent on the numerical grid as it is coupled to the smallest resolvable shock surface.  Taking this into account it is very likely that all our models would explode, if they were simulated with sufficiently high resolution.  The question of whether an incineration could also happen directly at the core--shell interface, was not addressed here.  This could of course prevent the spherical shock convergence mechanism from playing a role, at least in parts of the parameter space.  Therefore, edge-lit detonations are postponed to a separate study.  The complete double detonation simulations that were performed for the models z4.24A and s4.10A resulted in \\nni-masses of about $0.40$ to $0.45~\\msol$.  Are the studied events thus good candidates for normal SNe~Ia?  Most likely not.  Starting with a He-shell detonation, all of our models show a significant amount of rapidly expanding \\nni\\ in the outer layers.  In the observed spectra of normal and sub-luminous events, however, this has never been observed (cf.\\ \\citet{Branch1982,Branch1984a,Woosley1986,Arnett1997} and also the discussion in \\citet{Livne1995}).  The only exception is the super-luminous \\object{SN~1991T}.  Thus our solar-mass models most likely will not be able to explain normal or sub-luminous SNe~Ia and, given the robustness of the explosion mechanism of the model, the only conceivable explanation for the lack of observations of corresponding SN~Ia events is that the progenitors are not realized in nature -- or are very rare.  For a better agreement with observations a reduction of the He-shell mass would be the most obvious choice.  However, then a significantly larger core mass would most likely be required, because otherwise the He would not detonate.  In this case the core density would be much higher resulting in a very large Ni mass.  This would again not be a candidate for normal SNe~Ia, but it might be a promising candidate for objects like \\object{SN~1991T}.  Further work could cover model improvements like a more realistic treatment of nuclear reactions including the calculation of real reaction rates depending on the actual thermodynamic state of the burnt matter.  This would make an investigation of the onset and the explosion dynamics of the core detonation possible.  Also, more extended parameter studies, especially towards the lower end of possible He masses, would be interesting.  For a comparison with observations the calculation of synthetic light curves and spectra would also be desirable."
    },
    "0710/0710.5023_arXiv.txt": {
        "abstract": "{} {The principal aim of this project is to determine the structural parameters of the rapidly pulsating subdwarf B star PG 0911+456 from asteroseismology. Our work forms part of an ongoing programme to constrain the internal characteristics of hot B subdwarfs with the long-term goal of differentiating between the various formation scenarios proposed for these objects. So far, a detailed asteroseismic interpretation has been carried out for 6 such pulsators, with apparent success. First comparisons with evolutionary theory look promising, however it is clear that more targets are needed for meaningful statistics to be derived.} {The observational pulsation periods of PG 0911+456 were extracted from rapid time-series photometry using standard Fourier analysis techniques. Supplemented by spectroscopic estimates of the star's mean atmospheric parameters, they were used as a basis for the \"forward modelling\" approach in asteroseismology. The latter culminates in the identification of one or more \"optimal\" models that can accurately reproduce the observed period spectrum. This naturally leads to an identification of the oscillations detected in terms of degree $\\ell$ and radial order $k$, and infers the structural parameters of the target.} {The high S/N low- and medium resolution spectroscopy obtained led to a refinement of the atmospheric parameters for PG 0911+456, the derived values being $T_{\\rm eff}$=31,940$\\pm$220 K, $\\log{g}$=5.767$\\pm$0.029, and $\\log{\\rm He/H}$=$-$2.548$\\pm$0.058. From the photometry it was possible to  extract 7 independent pulsation periods in the 150$-$200 s range with amplitudes between 0.05 and 0.8 \\% of the star's mean brightness. There was no indication of fine frequency splitting over the $\\sim$68-day time baseline, suggesting a very slow rotation rate. An asteroseismic search of parameter space identified several models that matched the observed properties of PG 0911+456 well, one of which was isolated as the \"optimal\" model on the basis of spectroscopic and mode identification considerations. All the observed pulsations are identified with low-order acoustic modes with degree indices $\\ell$=0,1,2 and 4, and match the computed periods with a dispersion of only $\\sim$ 0.26 \\%, typical of the asteroseismological studies carried out to date for this type of star. The inferred structural parameters of PG 0911+456 are $T_{\\rm eff}$=31,940$\\pm$220 K (from spectroscopy) , $\\log{g}$=5.777$\\pm$0.002, $M_{\\ast}/M_{\\odot}$=0.39$\\pm$0.01, $\\log{M_{env}/M_{\\ast}}$=$-$4.69$\\pm$0.07, $R/R_{\\odot}$=0.133$\\pm$0.001 and $L/L_{\\odot}$=16.4$\\pm$0.8. We also derive the absolute magnitude $M_V$=4.82$\\pm$0.04 and a distance $d$=930.3$\\pm$27.4 pc.} {} ",
        "introduction": "Subdwarf B (sdB) stars are evolved extreme horizontal branch stars with atmospheric parameters in the 20,000 K $\\lesssim T_{\\rm eff} \\lesssim$ 40,000 K and  5.0 $\\lesssim \\log{g} \\lesssim$ 6.2 range \\citep{heber1986}. They are believed to be composed of helium-burning cores surrounded by thin hydrogen-rich envelopes and are characterised by a narrow mass distribution strongly peaked at $\\sim$ 0.48 $M_{\\odot}$ \\citep{dorman1993}.  While it is generally accepted that they evolved from the red giant branch (RGB) and constitute the immediate progenitors of low-mass white dwarfs \\citep{bergeron1994}, the details of their formation are not yet understood. It does however seem clear that sdB progenitors lost a significant fraction of their envelope mass near the tip of the first RGB, leaving them with insufficient fuel to ascend the asymptotic giant branch (AGB) after core helium exhaustion. Plausible formation channels were modelled in detail by, e.g., \\citet{han2002,han2003} and include binary evolution via a common envelope (CE) phase, stable Roche lobe overflow (RLOF), and the merger of two helium white dwarfs. These distinct evolutionary scenarios should leave a clear imprint not only on the binary distribution of sdB stars (CE evolution will produce sdB's in very close binary systems, RLOF gives rise to much longer period binaries, and the white dwarf merger results in single sdB stars), but also on their mass and hydrogen-envelope thickness distribution. Observational surveys focusing on radial velocity variations and the spectroscopic detection of companions have recently established that at least half of the sdB population reside in binaries (e.g. \\citet{allard1994,ulla1998}), a significant fraction of them having short orbital periods from hours to days (\\citet{green1997,maxted2001}). Accurate determinations of the stars' internal parameters on the other hand are harder to come by using traditional techniques; the mass has so far been measured only for the very rare case of an eclipsing binary \\citep{wood1999}, while the envelope thickness eludes direct study.  Fortunately, a small fraction ($\\sim$ 5 $\\%$) of sdB stars have been discovered to exhibit rapid, multi-periodic luminosity variations on a time-scale of hundreds of seconds, thus opening up the possibility of using asteroseismology to constrain their internal parameters. Since the near-simultaneous theoretical prediction \\citep{charp1996, charp1997} and observational discovery \\citep{kilkenny1997} of the so-called EC 14026 stars, both the modelling and measurement of their pulsational properties have come a long way (see \\citet{fontaine2006} for a review). Simulating the pulsation spectra of a large grid of sdB models in terms of low-degree, low-order $p$-modes and numerically determining the \"optimal\" model that best fits a series of observed periodicities has so far resulted in the asteroseismological determination of the internal parameters for six EC 14026 pulsators: PG 0014+067 \\citep{brassard2001}, PG 1047+003 \\citep{charp2003}, PG 1219+534 \\citep{charp2005a}, Feige 48 \\citep{charp2005b}, EC 20117-4014 \\citep{randall2006c}, and PG 1325+101\\citep{charp2006}. These first asteroseismic results show promising trends as far as matching the expected mass and hydrogen shell thickness distribution is concerned, however more targets are needed to start assessing the importance and accuracy of the proposed formation channels. Here we present an asteroseismological analysis of the subdwarf B star PG 0911+456 based on photometry obtained with the new Mont4kccd at Mt. Bigelow, Arizona. In the next sections we describe the observations and frequency analysis, followed by the asteroseismic exploitation of the target and a discussion of the internal parameters inferred. ",
        "conclusions": "We obtained 57 hours of broad-band time-series photometry as well as high S/N low- and medium resolution time-averaged spectroscopy for the EC 14026 pulsator PG 0911+456. Our observations led to refined estimates of the star's atmospheric parameters and the detection of 7 independent harmonic oscillations, 4 more than were known previously. There was no sign of frequency splitting over the 68-day period during which the photometry was obtained, indicating a slow rotation rate. Fixing the effective temperature to the spectroscopic value and conducting an asteroseismic search in 3-dimensional $\\log{g}-M_{\\ast}-\\log{q(\\rm H)}$ parameter space enabled the identification of several families of models that could reproduce the observed periods to within less than 1 \\%. While some of these were rejected from the outset due to obvious inconsistencies with the spectroscopic estimate of $\\log{g}$ or implausible associated mode identifications, we retained two promising solutions for closer inspection. Unlike in some previous asteroseismological studies, it was not immediately obvious which model was to be preferred on the basis of the structural parameters alone; instead we used the inferred mode identification, in particular the degree index $\\ell$, to discriminate between the two. The main difference  between the solutions was that one identified a relatively high amplitude peak with an $\\ell$=3 mode, while the other required only modes with $\\ell$=0,1,2 and 4. Since detailed computations reveal $\\ell$=3 modes to have extremely small disk-integrated amplitudes that would most likely not be detectable, we favoured the latter and adopted it as the optimal model.  The inferred structural parameters for PG 0911+456 include the total stellar mass and the thickness of the hydrogen-rich shell, two quantities that can normally not be derived using other means but are invaluable for a detailed understanding of subdwarf B stars' evolutionary history. The total mass determined is smaller than that found for any EC 14026 star to date and places our target at the low-mass end of the predicted distribution. Similarly, the hydrogen envelope is measured to be thinner than that of most sdB's studied so far. If PG 0911+456 is confirmed to be a single star, as is suspected from its slow rotation, negligible radial velocity variation, and absence of a companion's spectroscopic or near-IR photometric signature, it may be the product of a WD merger according to the evolutionary channels proposed by \\citet{han2002,han2003}. In this case, we would indeed expect the hydrogen envelope mass to be smaller than for an sdB having undergone a CE or RLOF phase. The low total mass derived would tend to support a non-canonical evolutionary history, even if it lies slightly below the mass distribution predicted from a WD merger.  An obvious follow-up study for PG 0911+456 is to verify the asteroseismological solution found on the basis of additional observations. These could aim for a higher S/N level, thus enabling the detection of further pulsations and strengthening the constraints on the asteroseismic model. One of the main challenges we faced in the search of parameter space was the relatively large number of models that could account for the oscillations observed quite accurately; it would be very instructive to see whether any newly found periods can also be fit by the optimal model isolated. Observations containing wavelength-dependent information from which modes may be partially identified are another option. Following recent theoretical investigations into the amplitude-wavelength dependence of a mode on its degree $\\ell$ (\\citet{ramachandran2004}, \\citet{randall2005}), there has been a surge in observational efforts to obtain multi-colour photometry of EC 14026 stars, most notably using the 3-channel CCD ULTRACAM (e.g. \\citet{jeffery2005}) and the Mont4kccd predecessor, LaPoune I (e.g. \\citet{fontaine2006}). While discriminating between low-degree modes with $\\ell$=0,1,2 has proved extremely challenging, $\\ell$=4 modes exhibit a more clearly distinguishable amplitude-wavelength behaviour (see e.g. Figure 26 of \\citet{randall2005}). Given the necessary data, it should therefore be possible to confirm the identification of the two $\\ell$=4 modes inferred for PG 0911+456 and thus verify the structural parameters computed.  The work presented here constitutes the 7th detailed asteroseismological analysis of a rapidly pulsating subdwarf B star. While we estimate that the structural parameters of around 20 targets are required to start detailed comparisons with evolutionary theory, first tentative efforts in this direction look promising. Nevertheless, it is clear that there is still ample room for improvement on the modelling front. Firstly, the fact that the dispersion between the observed and theoretical periods of the optimal model is generally an order of magnitude higher than the measurement accuracy indicates remaining shortcomings in the models. We are currently working on full evolutionary (rather than envelope) \"third-generation\" models to address this problem. The reliability of the \"optimal\" models identified during the search of parameter space is another issue. Although we do apply cross-checks such as compatibility with non-adiabatic theory and spectroscopic values of the atmospheric parameters, the latter (especially $T_{\\rm eff}$) must often be used to discriminate between, or constrain, regions of minimum $S^2$ and can no longer be employed as independent estimates. Moreover, the period ranges of unstable oscillations computed from our non-adiabatic pulsation code are sensitive mostly to $T_{\\rm eff}$ and $\\log{g}$ and are largely independent of the model mass and envelope thickness. It is therefore vital that additional checks are carried out with regard to the robustness of the \"forward\" approach if the structural parameters inferred from asteroseismology are to be compared with evolutionary predictions in a quantitative manner. The most obvious way of doing this is by detecting more frequencies from higher S/N observations or constraining the identification of the degree $\\ell$ of individual modes from multi-wavelength time-series data. Such efforts are ongoing, and will likely prove invaluable for the future of sdB star asteroseismology."
    },
    "0710/0710.2435_arXiv.txt": {
        "abstract": " ",
        "introduction": "Magnetic fields are observed to be associated with most structures in the universe. Observations indicate magnetic fields on stellar upto supergalactic scales. The field strengths vary from a few $\\mu$G on galactic scale, upto $10^3$ G for solar type stars and upto $10^{13}$ G for neutron stars. Furthermore, the magnetic field structure depends on the object it is associated with. Thus, e.g., magnetic fields observed in elliptical galaxies show a different structure from those associated with spiral galaxies \\cite{mag}.  Magnetic fields in stars can be explained by the formation of protostars out of condensed interstellar matter which was pervaded by a pre-existing large scale magnetic field (see, e.g., \\cite{rees}). An open problem remains to explain the origin of such large scale magnetic fields.  There are different types of proposals. Ranging from processes on small scales, such as vortical perturbations and phase transitions to models taking advantage of the possibility of amplifying perturbations in the electromagnetic field during inflation in the early universe (see, e.g., \\cite{rev-mag}).    Inflation provides a mechanism to amplify perturbations in some field to appreciable size. In order for this mechanism to lead to primordial magnetic seed fields of cosmologically interesting strength, the corresponding lagrangian should not be conformally invariant. The Maxwell lagrangian describing linear electrodynamics is conformally invariant. There have been already a multitude of proposals to break the conformal invariance of the Maxwell theory \\cite{tw}, e.g. by coupling to a scalar field \\cite{mag-sc}, breaking Lorentz invariance \\cite{mag-lor}, adding extra dimensions \\cite{mag-ex} or a coupling to curvature terms \\cite{mag-curv}.  Here nonlinear electrodynamics is considered. It has its origins in the search for a classical singularity-free theory of the electron by Born and Infeld \\cite{bi}. Later on it was realized that virtual electron pair creation induces a self-coupling of the electromagnetic field. For slowly varying, but arbitrarily strong electromagnetic fields the self-interaction energy was computed by Heisenberg and Euler (cf. \\cite{he}-\\cite{bb}).  The propagation of a photon in an external electromagnetic field can be described effectively by the Heisenberg-Euler langrangian. Moreover, the transition amplitude for photon splitting in quantum electrodynamics is nonvanishing in this case. In principle, this might lead to observational effects, e.g., on the electromagnetic radiation coming from neutron stars which are known to have strong magnetic fields. \\cite{bb,p-sp}. In particular, certain features in the spectra of pulsars can be explained by photon splitting \\cite{puls}.    Finally, Born-Infeld type actions also appear as a low energy effective action of open strings \\cite{bi-strings,gh}. As was shown in \\cite{dbi} the low energy dynamics of D-branes is described by the Dirac-Born-Infeld action.    The model of the cosmological background that will be considered consists of a stage of de Sitter inflation followed by reheating and a standard radiation dominated stage. Quantum fluctuations in the electromagnetic field are excited within the horizon during inflation. Once outside the horizon they become classical perturbations. As mentioned above, in general, the conformal invariance of the four dimensional Maxwell field has to be broken in order to amplify the perturbations in the electromagnetic field significantly. Thus, here electrodynamics is considered to be nonlinear during the de Sitter stage. This could be motivated by the presence of possible quantum corrections to quantum electrodynamics at high energies. However, once inflation ends electrodynamics is described by standard Maxwell electrodynamics. Thus the subsequent evolution described by the standard model of cosmology is unchanged. ",
        "conclusions": "Observations of magnetic fields on large scales provide an intriguing problem. A possible class of mechanisms to create such fields is provided by inflationary models. Fluctuations in the electromagnetic field are amplified during inflation and provide a seed magnetic field at the time of structure formation which might be further amplified by a dynamo process. In general a sufficiently strong initial field strength can only be achieved if the conformal invariance of electrodynamics is broken. This has been realized, for example, in models where the Maxwell lagrangian has been coupled to a scalar field, to curvature terms, etc. or by breaking Lorentz invariance or adding extra dimensions.  Here nonlinear electrodynamics has been considered. It has been assumed that whereas during the early universe electrodynamics is nonlinear it becomes linear at the end of inflation. In particular the evolution of the magnetic energy density has been studied in a model of nonlinear electrodynamics which is described by a lagrangian of the form $L\\sim -\\left[\\left(F_{\\mu\\nu}F^{\\mu\\nu}\\right)^2/\\Lambda^8\\right]^{\\frac{\\delta-1}{2}} F_{\\mu\\nu}F^{\\mu\\nu}$, where $\\Lambda$ and $\\delta$ are  parameters. Originally the nonabelian version of this model had been proposed to describe low energy QCD \\cite{pt}. Here this model has been chosen as it is a strongly nonlinear theory of electrodynamics which allows to study in a semi-analytical way the possible creation and amplification of a primordial magnetic field during de Sitter inflation. This is so since on the one hand the lagrangian only depends on one of the electromagnetic invariants, namely $X=\\frac{1}{4}F_{\\mu\\nu}F^{\\mu\\nu}$, which leads to a significant simplification of the equations. On the other hand the power-law structure of the lagrangian make the equations simpler.  Approximate solutions have been found in three regimes of approximation which describe the evolution of the ratio of the energy densities of the electric and magnetic fields during inflation. It is assumed that initially the energy density of the electric and magnetic field are of the same order. Furthermore, these initial fields are due to quantum fluctuations in the electromagnetic field during inflation. Whereas in the radiation dominated era, the energy density in the magnetic field decreases as $a^{-4}$, the electric field strength rapidly decays in the highly conducting plasma (see, e.g., \\cite{tw,d93}). Solutions in closed form have been found and the resulting primordial magnetic field estimated. It has been shown that depending on the regime of approximation and the value of the Pagels-Tomboulis parameter $\\delta$ primordial magnetic fields can be generated that are strong enough to seed a galactic dynamo. Thus we have provided an example of a theory of nonlinear electrodynamics where the nonlinearities act in a way as to amplify sufficiently an initial magnetic field.  The energy-momentum tensor of the electromagnetic field can be cast in the form of an imperfect fluid. This has been found explicitly for the particular model of nonlinear electrodynamics under consideration here. Moreover, this allows to find the expression for the energy density $\\rho$ of the fluid. Requiring that $\\rho$ should be positive provides the bound $\\delta\\geq\\frac{1}{2}$.  In  \\cite{gfc} the possible creation and amplification of magnetic fields was studied in an inflationary model coupled to a pseudo Goldstone boson (see also \\cite{tw}). In this case the lagrangian has the form $L\\sim\\frac{1}{2}\\partial_{\\mu}\\theta\\partial^{\\mu}\\theta -X+g_a\\theta Y$, where $\\theta$ is the axion field. This provides an example of a more general lagrangian having also an explicit dependence on $Y=\\frac{1}{4}F_{\\mu\\nu}\\;^{*}F^{\\mu\\nu}$. However, as it turns out the resulting primordial magnetic field is not strong enough in order to seed, for example, a galactic dynamo. In \\cite{as} the model of \\cite{gfc} was generalized to N axions. In this case it was found that at least the weaker bound of $r>10^{-57}$ can be satisfied. Here, in this work the creation of primordial magnetic fields in a particular model of nonlinear electrodynamics has been studied. It might be interesting to generalize this to lagrangians depending on both electromagnetic invariants $X$ and $Y$."
    },
    "0710/0710.1600_arXiv.txt": {
        "abstract": "Prior to the  incineration of a white dwarf (WD) that makes a Type Ia supernova (SN Ia), the star ``simmers'' for $\\sim 1000$ years in a convecting, carbon burning region. We have found that weak interactions during this time increase the neutron excess by an amount that depends on the total quantity of carbon burned prior to the explosion. This contribution is in addition to the metallicity ($Z$) dependent neutronization through the $^{22}$Ne abundance (as studied by Timmes, Brown, \\& Truran). The main consequence is that we expect a ``floor'' to the level of neutronization that dominates over the metallicity contribution when $Z/Z_\\odot\\lesssim2/3$, and it can be important for even larger metallicities if substantial energy is lost to neutrinos via the convective Urca process. This would mask any correlations between SN Ia properties and galactic environments at low metallicities. In addition, we show that recent observations of the dependences of SNe Ia on galactic environments make it clear that metallicity alone cannot provide for the full observed diversity of events. ",
        "introduction": "\\label{sec:introduction}  The use of Type Ia supernovae (SNe Ia) as cosmological distance indicators has intensified the need to understand white dwarf (WD) explosions. Of particular importance is the origin of the Phillips relation \\citep{phi99}, an essential luminosity calibrator. Recent models demonstrate that it can be explained by large variations in the abundance of stable iron group elements \\citep{woo07} with the dominant cause for  diversity likely residing in the explosion mechanism \\citep{maz07}.  One additional variable is the metallicity of the WD core, which yields excess neutrons relative to protons due to the isotope $^{22}$Ne. This is usually expressed as \\be Y_e = \\sum_i \\frac{Z_i}{A_i}X_i, \\ee where $A_i$ and $Z_i$ are the nucleon number and charge of species $i$ with mass fraction $X_i$. The neutronization is critical for setting the production of the neutron-rich isotopes \\citep{thi86}. If no weak interactions occur during the explosion, the mass fraction of $^{56}$Ni produced is simply $X(^{56}{\\rm Ni}) = 58Y_e-28$, assuming $^{56}$Ni and $^{58}$Ni are the only burning products \\citep{tbt03}. The neutronization also affects the explosive burning, including the laminar flame speed \\citep{cha07}. However, the metallicity range of progenitors is not large enough to account for the full SNe Ia diversity (see \\S 4), making it critical to explore all factors that determine $Y_e$.  A potential neutronization site is the convective carbon burning core that is active for $\\sim1000\\ {\\rm years}$ prior to the explosion. The hydrostatic evolution associated with this simmering phase terminates when the core temperature is sufficiently high that burning is dynamical \\citep{nom84,ww86,woo04,ww04,kwg06}, and a flame commences \\citep{tw92}. We show here that protons from the $^{12}$C($^{12}$C,$p$)$^{23}$Na reaction during simmering capture on $^{12}$C, and that subsequent electron captures on $^{13}$N and $^{23}$Na decrease $Y_e$. This process continues until either the explosion occurs or sufficient heavy elements have been synthesized that they capture the protons instead.  In \\S \\ref{sec:rates}, we present simmering WD core models and summarize the reaction chains that alter $Y_e$. We find that one proton is converted to a neutron for each six $^{12}$C nuclei consumed for  burning at $\\rho<1.7\\times 10^9\\ {\\rm g \\ cm^{-3}}$. At densities above this, an additional conversion occurs from an electron capture on  $^{23}$Na. Hence, the $Y_e$ in the core depends on the amount of carbon burned during simmering and the density at which it occurs, which we quantify in \\S 3. We find that neutronization during simmering  dominates for metallicities $Z/Z_\\odot\\lesssim2/3$. We conclude in \\S \\ref{sec:conclusion} by discussing the observations and noting where future work is needed. ",
        "conclusions": "\\label{sec:conclusion}  We have found that significant neutronization of the WD core occurs throughout the simmering stage of carbon burning until the onset of the explosion. If substantial energy is lost to the convective Urca process \\citep[][and references therein]{les05}, then the neutronization is truncated by proton captures onto freshly synthesized heavy elements (resulting in eq. [\\ref{eq:yemax}]). The main consequence is a uniform ``floor'' to the amount of neutronization that dominates over the metallicity dependent contribution for all progenitors with $Z/Z_\\odot\\lesssim2/3$.  Given the likely significance this has for SNe Ia, more work needs to be done. In particular, at high ignition densities, heavy element electron captures and a full reaction network are needed to follow the resulting diverse collection of elements (see the discussion in \\S 2.2). The convective Urca process % is another complication we have not addressed. In principle, if more energy is lost to neutrinos then more burning (and thus more neutronization) is required to make the burning dynamical. Assessing this will necessitate coupling a full nuclear network \\citep{cha07b} to convective calculations. Such models would accurately determine $\\eta$ rather than simply setting it to 1 or 2.  In closing, we highlight some important features exhibited by recent observations of SNe Ia. It is clear that the amount of $^{56}$Ni produced in SNe Ia has a dynamic range ($0.1-1M_\\odot$) larger than can be explained by metallicity or simmering neutronization.  However, since an intriguing trend is the prevalence of $^{56}$Ni rich events in star-forming regions it is interesting to quantitatively explore how large the observed discrepancy is. Using the SNLS sample of Sullivan et al. (2006), Howell et al. (2007) found that the average stretch is $s=0.95$ in passive galaxies (e.g. E/S0's) and $s=1.05$ in star-forming galaxies. Using Jha et al's (2006) $\\Delta M_{15}(B)-s$ relation and Mazzali et al.'s (2007) relation between $\\Delta M_{15}(B)$ and $^{56}$Ni mass we get $0.58M_\\odot$ ($s=0.95$) and $0.72M_\\odot$ ($s=1.05$). Hence, amongst the large diversity, there is a tendency for SNe in star-forming galaxies to produce $\\approx 0.13M_\\odot$ more $^{56}$Ni than those in large ellipticals.  Since the SN Ia rate scales with mass in ellipticals and star formation rate in spirals (Mannucci et al. 2005; Scannapieco \\& Bildsten 2005; Sullivan et al. 2006), SNe from passive galaxies in the SNLS survey are from more massive galaxies than the SNe in star-forming galaxies (Sullivan et al. 2006). Using the mass-metallicity relation of Tremonti et al. (2004), our integration of the separate samples in Sullivan et al. (2006) yield average $12+\\log({\\rm O/H})=8.87$ in active galaxies and $9.1$ in ellipticals (solar value is $12+\\log({\\rm O/H})=8.69$). Due to the increase of ${\\rm N/O}$ at high metallicities (Liang et al. 2006), the SNe in ellipticals have twice as much $^{22}$Ne content as those in spirals. From the relation of Timmes et al. (2003), this implies $\\approx 5\\%$ less $^{56}$Ni, whereas the observed decrement is $>15\\%$. Explaining the observed decrement would require a contrast of $\\Delta X(^{22}$Ne$)\\approx 0.06$, or nearly 3 times solar. Although we have found that simmering enhances neutronization, the effect is not great enough ($\\Delta Y_{e,\\rm max}$ would give the same change in neutronization as doubling a solar metallicity), and a diverse set of core conditions would still be required. A large enhancement could be present in the core if substantial gravitational separation had occurred \\citep{bh01,db02}, yet convective mixing during simmering will reduce it based on the fraction of the star that is convective. For a convection zone that extends out to $M_\\odot$, the resulting $^{22}$Ne enhancement would be at most $\\approx 30\\%$."
    },
    "0710/0710.1752_arXiv.txt": {
        "abstract": "The early evolution of dense stellar systems is governed by massive single star and binary evolution. Core collapse of dense massive star clusters can lead to the formation of very massive objects through stellar collisions ($M\\geq$ 1000\\,\\msun). Stellar wind mass loss determines the evolution and final fate of these objects, and decides upon whether they form black holes (with stellar or intermediate mass) or explode as pair instability supernovae, leaving no remnant. We present a computationaly inexpensive evolutionary scheme for very massive stars that can readily be implemented in an N-body code. Using our new N-body code 'Youngbody' which includes a detailed treatment of massive stars as well as this new scheme for very massive stars, we discuss the formation of intermediate mass and stellar mass black holes in young starburst regions. A more detailed account of these results can be found in \\cite{Belkusetal}. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.4946_arXiv.txt": {
        "abstract": "Thermal instability of partially ionized plasma is investigated by means of a linear perturbation analysis. According to the previous studies under the one fluid approach, the thermal instability is suppressed due to the magnetic pressure. However, the previous studies did not precisely consider the effect of the ion-neutral friction, since they did not treat the flow as two fluid which is composed of ions and neutrals. Then, we revisit the effect of the ion-neutral friction of the two fluid to the growth of the thermal instability. According to our study, the characteristic features of the instability are the following four points: (1) The instability which is characterized by the mean molecular weight of neutrals is suppressed via the ion-neutral friction only when the magnetic field and the friction are sufficiently strong. The suppression owing to the friction occurs even along the field line. If the magnetic field and the friction are not so strong, the instability is not stabilized. (2) The effect of the friction and the magnetic field is mainly reduction of the growth rate of the thermal instability of weakly ionized plasma. (3) The effect of friction does not affect the critical wavelength $\\lambda _{\\rm F}$ for the thermal instability. This yields that $\\lambda _{\\rm F}$ of the weakly ionized plasma is not enlarged even when the magnetic field exists. We insist that the thermal instability of the weakly ionized plasma in the magnetic field can grow up even at the small length scale where the instability under the assumption of the one fluid plasma can not grow owing to the stabilization by the magnetic field. (4) The wavelength of the maximum growth rate of the instability shifts shortward according to the decrement of the growth rate. This is because the friction is effective at rather larger scale. Therefore, smaller structures are expected to appear than those without the ion-neutral friction. Our results indicate the friction with the magnetic field affects the morphology and evolution of the interstellar matter. In summary, the ion-neutral friction is important for the evaluation of the thermal instability in weakly ionized plasma along and perpendicular to the magnetic field. ",
        "introduction": "The recent progress of the interstellar medium (ISM) observations has established that the small and tiny scale structures of ISM is very ubiquitous. Historically, the initial observational target is the neutral phase of ISM. Indeed, the tiny scale structures are discovered by 21 cm absorption line observations against quasars with VLBI techniques \\citep{Di76}. The results are confirmed by \\citet{Dia89}. To look for the small scale structures in the cold neutral phase of the ISM furthermore, the observations of 21 cm absorption lines have been also performed against pulsars \\citep{Fr94}. \\citet{Me96} observe optical spectral lines to find small structures against close binary stars. The images of the small-scale H \\textsc{i} have been taken with the MERLIN array \\citep{Da96}. These results are established by higher angular resolution observations with VLBA and VLA toward quasars \\citep{Fa98,FG01}, which show small clumps on the order of a few AU in neutral Galactic H \\textsc{i} clouds. Cold H \\textsc{i} clouds have significant structure in subparsec scales \\citep{Gi00,Br05}. Such small and tiny scale structures has been detected in the local interstellar medium as H \\textsc{i} absorption lines, although their column density is very small \\citep{BK05}.  In addition to the small and tiny scale structures of H {\\rm \\textsc{i}} cloud, the picture of the planetary nebula NGC 7293 shows fine small structures and knots \\citep{Ro02,OD04}. Even in the starforming regions, there are variety small scale structures. For example, \\citet{La95} observe some clumps with size from 0.007 to 0.021 pc in the Taurus Molecular Cloud 1 (TMC1), in particular, core D. The mass of these fragments is estimated to be $\\lesssim 0.01 - 0.15$ $M_{\\odot}$. SPITZER has begun producing higher spatial resolution mid-infrared maps \\citep{Ch04}, and revealed the fine structure of the starforming region. We think that it is possible for proto-brown dwarfs of $< 0.08$ $M_{\\odot}$ to exist there. A high resolution observation in future is expected to reveal hidden, small structures, as well as substellar objects with very small mass. In any ways, the tiny-scale structure seems to be ubiquitous, not associated with large extinction \\citep{He97}.  To study the origin of the small-scale structures of ISM is important to understand the evolution of and structure formation in ISM. Especially, in the starforming regions, the small, low-mass structures can relate to the low-mass cutoff of the initial mass function, the coagulation unit to form the massive stars, and so on. Then, we should investigate the physical origin of these small and tiny structures in the partially ionized medium, since the cold H {\\rm \\textsc{i}} and molecular clouds are weakly ionized. By the way, it is not easy for the small and tiny scale structure to form as a result of gravitational instability since their size is much smaller than the Jeans length. According to \\citet{La95}, indeed, the small-scale structures appear to be gravitationally unbound. This suggests that some fragmentation mechanisms but pure Jeans gravitational instability may be important in the clouds. We expect the thermal instability as this mechanism, and revisit it in this paper.  The basics of the thermal instability has been summarized in \\citet{F65}. When the following condition is satisfied, the system is thermally unstable: If the cooling becomes efficient as the temperature decreases, the thermal energy gets lost more and more. If the cooling becomes efficient during the fluid contraction, the system shrinks on and on. Importantly, the critical length scale of the thermal instability is smaller than that of the dynamical instability like the Jeans instability. That is, even if a system is stable against the gravitational instability, the system can become thermally unstable. Then, the thermal instability can be the physical origin of the smaller-scale formation than the dynamical instability. Indeed, \\citet{KI02} propose that the clumpiness in clouds emerges naturally from their formation through the thermal instability. Their two-dimensional calculations follow the fragmentation into small cloudlets that result from the thermal instability in a shock-compressed layer. \\citet{BL00} investigate the cooling and fragmentation of optically thin gas with a power-law cooling function. According to them, small-scale perturbations have the potential to reach higher amplitude than large-scale fluctuations. Thermal instability can be important for the structure formation.  We sketch ISM structure formation such as molecular clouds as follows. The hot ionized ISM cools to be cold and weakly ionized clouds. Ionization degree of ISM decreases as it evolves, and then the ISM fluid is often partially ionized. In partially ionized fluid, an ion component and neutral component interact each other, exchanging the momentum (i.e., ion-neutral friction). Especially, the partially ionized plasma in a magnetic field could not be treated as one fluid. This is because although the ion component in the partially ionized plasma is directly influenced by a magnetic field, the neutral component does not directly feel a magnetic field. In weakly ionized fluid, the neutral component is affected by a magnetic field via the ion-neutral friction. For example, ambipolar diffusion takes an important role in dynamical evolution owing to the gravitational instability of ISM into protostar \\citep{MeSp56, Na76}. If the fluid frozen in a magnetic field contracts without ambipolar diffusion, the magnetic field becomes too large and suppresses the growth by the magnetic pressure and tension. In addition, \\citet{Na79} points out the possibility that ISM with sufficient magnetic flux quasistatically evolves into protostar. He also notices that the time scale of plasma drift depends on the ionization degree. The relation among the ion-neutral friction, the magnetic field, and the amount of the ion component is important to understand the detailed process.   We here emphasize that the MHD approximation is not always applicable for our purpose in this paper. The thermal instability is effective in small structure formation, during which a neutral component is not always enough frozen in an ion component dynamically to be treated as one fluid. There is a possibility that the ion component can drift away owing to ambipolar diffusion, and that the system can contract owing to thermal instability. Then, in the present paper, to study how the ion-neutral drag and magnetic field influence the growth of the thermal instability, we treat the plasma as the two fluid of ions and neutrals. We assume that the ionization degree is very small, because we are interested in the final stage of the formation of a molecular cloud, in which the ionization degree decreases very much. We focus on the condensation mode of thermal instability, since it is important in structure formation, rather than the oscillation (overstable) mode. In \\S2, the problem is formulated. The property of the dispersion relation of the instability is presented in \\S3. Several discussions relating to the instability are found in \\S4. Applicability of our results to cold H {\\rm \\textsc{i}} medium is briefly examined in \\S5, and then we summarize the paper in \\S6. ",
        "conclusions": "\\label{S_Discuss}  We study the basic property of the thermal instability of the weakly ionized plasma in the previous sections. In this section, basing on our understanding the thermal instability, we shall discuss the structure formation.   \\subsection{Critical Wavelength and Thermal Conduction} \\label{subsec:critical_wavelength_conduction}  First in this section, the critical wavelength of the weakly ionized plasma in the magnetic field is discussed. As mentioned in section \\ref{S_Results}, the critical wavelength of the mode relating $\\mu_{\\rm n}$ is not changed by the magnetic field with the friction as long as the thermal conduction is not changed by the magnetic field. In fact, the thermal conductivity is varied owing to the magnetic field. In this meaning, we may oversimplify the problem in this paper.   However, the thermal conductivity will decrease as the magnetic field strength increases, and then the critical length become smaller, according to equation (\\ref{eq:appendix:criticallength:neutral}). Thermal conduction erases the temperature perturbation and thermally stabilizes the system, especially at smaller scale. The larger thermal conduction makes the critical wavelength longer, owing to the criterion \\ref{crite_cond} or \\ref{crite_ion}, according to \\citet{F65}. On the other hand, the magnetic field lessens the thermal conduction of the ion component, because the ion component winds around a magnetic field.   Therefore, we can insist that the critical length for the thermal instability of the weakly ionized plasma can not be enlarged by the friction and the magnetic field at least, while the effect of the magnetic field can influence not only the ion but the neutral component via the ion-neutral friction. This statement has universal validity for the thermal instability of the weakly ionized plasma. The effect of the friction and the magnetic field is mainly reduction of the growth rate of the thermal instability. On the other hand, the magnetic field directly enlarges the critical wavelength of the mode characterized by $\\mu_{\\rm i}$ as expected in the one fluid plasma model.      \\subsection{Spatial Distribution of Magnetized Interstellar Medium}  In the previous section, we learn that the suppression of the thermal instability is characterized by the ion-neutral friction, strength of the magnetic field, and the direction of the field. Interestingly, the suppression of the instability is different owing to the magnetic field direction. This suggests the anisotropic nature of the suppression of the instability is imprinted in the morphology of the structure. In this standing point, we shall discuss how the structure formation proceeds in the magnetized ISM.    \\subsubsection{Growth Rate Parallel to the Magnetic Field}  In an usual one fluid approach, the motion of the plasma parallel to the magnetic field is regarded not to be affected. However, Figure \\ref{B0C003} suggests that the growth rate of the thermal instability at larger scale decreases via the ion-neutral friction. This suppression of the instability yields the delay of the evolution in comparison to the evolution under the MHD approximation without the ion-neutral friction.    \\subsubsection{Distribution of Partially Ionized Plasma Perpendicular to the Magnetic Field Line} \\label{sss_spread_of_part}  Here, we suggest that the spatial distribution of the partially ionized plasma is elongated, whose semi-major axis is vertical to the magnetic field line. Figures \\ref{B0C003} and \\ref{B06C003}, both of which friction strengths are the same, show the feasibility. The thermal instability in Figure \\ref{B0C003} corresponds to the mode along the magnetic field line since $\\mbox{\\boldmath$B$} = 0$, while that in Figure \\ref{B06C003} is vertical to the magnetic field line. The growth of thermal instability vertical to the line is suppressed by both of the magnetic field and the friction, while the growth along the line is suppressed only by the friction. Thus, the growth rate along the magnetic field line is larger than that vertical to the line, and then the resultant morphology of the structure via the thermal instability is elongated.   \\subsubsection{Difference of Spatial Distribution of Ion and Neutral Components} \\label{ssS_Difference_in_Spread}  We insist the growth rate depending on the direction of the field results in the morphology of interstellar structure. Furthermore, we learn in \\S3 the mode characterized by $\\mu_{\\rm i}$ are stabilized by the friction and field, while that by $\\mu_{\\rm n}$ can grow as found in Figure \\ref{B06C003}. Figure \\ref{B600C5} shows that although the mode characterized by $\\mu_{\\rm i}$ is completely stabilized, the mode characterized by $\\mu_{\\rm n}$ is still unstable. Consequently, we can insist that the neutral component condenses owing to the thermal instability, while the ion component still keeps spreading.   It is noticed that the ion component may condense since it is dragged by the neutral component contracting owing to the thermal instability. However, once the distribution of the ion and the neutral components separates and the motion of the two becomes independent there, the ion component in the area is thermally stable owing to the magnetic field. This means the magnetic pressure lets the ion component be spread in the area even if the neutral component condenses.   This thin ion component left behind by the contracting neutral component will be observed as a halo around H {\\rm \\textsc{i}} clouds, and it emits, especially some forbidden lines. This halo possibly has wide distribution vertical to the magnetic field, as discussed in subsubsection \\ref{sss_spread_of_part}. This is because the ion component can condense, parallel to the magnetic field like the neutral. The magnetic pressure can not prevent the motion of the ion component efficiently as found in Figure \\ref{B0C003}. Thus, the morphology of the ion halo is also elongated according to the orientation of the field line.    \\subsubsection{Origin of the Small Clumpinesses of Partially Ionized Plasma} \\label{subsubsec:small_clumpinesses_of_partial}   It is suggested that there are small clumpy structures of ISM. Here, we examine whether the clumpiness originates via the thermal instability. According to Figures \\ref{B600C03} and \\ref{B600C5}, the growth rate of the thermal instability of partially ionized plasma depends on the strength of the friction. The scale length at the maximum growth rate of the instability becomes smaller, owing to the suppression via the friction which is effective at larger scale. Thus, we can expect the small clumpiness appears selectively via the thermal instability since the ion-neutral friction suppresses the large scale condensations.    If we are interested in the small clumpiness of ISM, we should mention about the thermal conduction in relation to the magnetic field. According to subsection \\ref{subsec:critical_wavelength_conduction}, the small clumpiness may the results of the decrement of the conduction. It is noticed that the conduction along the field line is not so affected, and then there is still ambiguity for the reduced conduction coefficient to perform the key role for the origin of the small clumpiness.   It is also noticed that the thermal instability is a slow process relative to the dynamical processes. If molecular clouds form from the H {\\rm \\textsc{i}} medium via the gravitational instability, the small clumpiness inside the molecular clouds appears after the formation of the clouds. This trend is enhanced if the effects of the magnetic field and friction are concerned. This is because the friction and the field suppress the large scale growth, while the short wavelength modes are still unstable. In a enough long span for the growth, the tiny structures can appear in ISM and affect the evolution of its structure especially as the non-linear effects. Thus, if we want to know how faint and small structures come to be observable, the growth time-scale of the thermal instability should be precisely examined. At this time, we never forget the effect of the magnetic field and the ion-neutral friction.   \\subsection{Ineffectiveness of the Magnetic Field} \\label{S_Discuss_Satu}    The modes in Figures \\ref{B600C03} and \\ref{B6000C03} have similar growth rate, while the magnetic fields are significantly different in each other. It implies that the suppression of the growth by increasing the magnetic field with the fixed friction has limits. By the way, a dispersion relation of only ionized fluid in a magnetic field, which is derived by equalizing the second bracket on the left-hand side in equation (\\ref{dis-rela}) to be zero, is reduced to \\begin{equation} n^3 + n^2 v_{\\rm si} k_{\\rm iT} = 0 , \\end{equation} when $k$ approaches $0$. This is solved to be \\begin{equation} n=0 ~ {\\rm or}~ n = -v_{\\rm si} k_{\\rm iT} . \\end{equation} Thus, we find that this solution is independent of a magnetic field. Even if the magnetic field becomes strong, the value of the growth rate at $k \\sim 0$ is scarcely changed, where the friction is more effective than at small wavelength. Therefore, once the effect of the friction is fixed, increasing the magnetic field strength beyond a threshold level apparently becomes ineffective for the suppression of the growth. Not only information of the magnetic field but the friction is necessary to investigate the evolution of the thermal instability of the weakly ionized plasma.         \\subsection{Comment on One Fluid Approximation}  The formulation of the fully one component plasma can not be always applicable to study the thermal instability of weekly ionized plasma. This is because there is the growing mode even when there exist the magnetic field and friction, whose growth rate is comparable to that of the neutral component when there are no magnetic field and friction. Indeed, this trend is found from the comparison among Figures \\ref{B06C003}, \\ref{B600C03} and \\ref{B6000C03}.  In addition, regarding weakly ionized plasma including the friction effect, the friction does not change the critical wavelength $\\lambda_{\\rm F}$. The friction tends to efficiently reduce the growth at much larger scale than the $\\lambda_{\\rm F}$. This feature yields that the critical wavelength in the weakly ionized plasma is not changed by the magnetic field via the friction, as mentioned in section \\ref{S_Results} and subsection \\ref{subsec:critical_wavelength_conduction}. Even if the magnetic field becomes strong, only the growth rate is reduced, keeping the critical wavelength constant. This is also the difference from the simple analysis of the one component magnetized plasma. The previous study of the one component plasma shows the stabilization at the smaller scales, while the current study insist $\\lambda_{\\rm F}$ is unchanged even if the effect of the field is concerned.         \\label{S_Summary}   The thermal instability of the weakly ionized fluid is investigated with a linear perturbation analysis. The plasma is assumed to consist of two fluids of the ion and the neutral components. With this approach, the effect of the friction between the ion and the neutral components and the magnetic field are important. Here, the properties of the thermal instability of weakly ionized plasma and the observational implications are summarized.  \\begin{enumerate}  \\item Modes relating to $\\mu_{\\rm n}$ are not stabilized by the magnetic field. This is because the neutral component feels the field via the ion-neutral friction.  \\item Modes relating to $\\mu_{\\rm i}$ are directly stabilized by the magnetic field. This means that when the magnetic field is large, the ion component is hard to follow the motion of the neutrals which condensate owing to the thermal instability.  \\item The growth rate of the mode vertical to the magnetic field is reduced by the magnetic field and the ion-neutral friction.  \\item The growth rate of the mode along the magnetic field decreases owing to the ion-neutral friction, especially at large scale. The suppression of the instability is more ineffective than that of the mode vertical to the magnetic field.  \\item The critical wavelength for the thermal instability is not affected by the friction. The effect owing to the two fluid approximation is the reduction of the growth rate of the thermal instability. The magnetic field makes the critical wavelength of modes relating $\\mu_{\\rm i}$ larger.  \\item To study the thermal instability of the partially ionized plasma, one fluid approximation is not always useful.  \\item The ion-neutral friction and the magnetic field affect the distribution or morphology of ISM, especially after the long time compared to the free-fall time. The difference between the growth rate along and perpendicular to the magnetic field is important. The partially ionized plasma possibly is elongated perpendicular to the magnetic field. In addition, the neutral component and the ion component of weakly ionized fluid in the magnetic field are possibly separated each other. Then, the neutral component condenses owing to thermal instability, while the ion component left behind by the contracting neutral component still keeps spreading. Therefore, it is possibly observed that the weakly ionized fluid is elongated vertically to the magnetic field and surrounded by a halo which is rare and diffuse ions and emits forbidden lines.  \\end{enumerate}  To conclude, the treating both of the independent motions of the neutral and the ion components yield the ion-neutral friction. The friction is important for the thermal instability of weakly ionized fluid, especially when studying the structure formation such as molecular cloud at its final stage of formation. Our study indicates that the fully ionized plasma approximation or totally neutral fluid approximation is not always applicable to weakly ionized plasma."
    },
    "0710/0710.1564_arXiv.txt": {
        "abstract": "{% In a self-absorbed synchrotron source with power-law electrons, rapid inverse Compton cooling sets in when the brightness temperature of the source reaches $T_{\\rm B}\\sim10^{12}\\,$K. However, brightness temperatures inferred from observations of intra-day variable sources (IDV) are well above the \"Compton catastrophe\" limit. This can be understood if the underlying electron distribution cuts off at low energy.} {% We examine the compatibility of the synchrotron and inverse Compton emission of an electron distribution with low-energy cut-off with that of IDV sources, using the observed spectral energy distribution of S5~0716+714 as an example.} {% We compute the synchrotron self-Compton (SSC) spectrum of monoenergetic electrons and compare it to the observed spectral energy distribution (SED) of S5~0716+714. The hard radio spectrum is well-fitted by this model, and the optical data can be accommodated by a power-law extension to the electron spectrum. We therefore examine the scenario of an injection of electrons, which is a double power law in energy, with a hard low-energy component that does not contribute to the synchrotron opacity.} {% We show that the double power-law injection model is in good agreement with the observed SED of S5~0716+714. For intrinsic variability, we find that a Doppler factor of $\\doppler\\geq30$ can explain the observed SED provided that low-frequency ($<32\\,$GHz) emission originates from a larger region than the higher-frequency emission. To fit the entire spectrum, $\\doppler\\ge65$ is needed. We find the constraint imposed by induced Compton scattering at high $T_{\\rm B}$ is insignificant in our model.  } {% We confirm that electron distribution with a low-energy cut-off can explain the high brightness temperature in compact radio sources. We show that synchrotron spectrum from such distributions naturally accounts for the observed hard radio continuum with a softer optical component, without the need for an inhomogeneous source. The required low energy electron distribution is compatible with a relativistic Maxwellian.} ",
        "introduction": "\\label{intro}  Observations of many extra-galactic radio sources have found rapid flux variations at radio frequency \\citep[e.g.][]{kedziorachudczeretal01}, some of which fluctuate over a time scale of a day or less. They are referred to as intra-day variable sources (IDV). The variability time scale is often used to constrain the size of the source based on causality arguments. Using this constraint, one can derive a variability brightness temperature \\citep{wagnerwitzel95} \\eqb T_{\\rm var}=4.5\\times10^{10}F_{\\nu}\\left({\\lambda d_{\\rm L}\\over t_{\\rm obs} (1+z)}\\right)^2\\, {\\rm K} \\eqe where the flux density $F_\\nu$, wavelength $\\lambda$, luminosity distant $d_{\\rm L}$, and observed variability time scale $t_{\\rm obs}$ are measured in Jy, cm, Mpc, and days, respectively.  The high radio flux frequently measured in IDV sources implies an extremely high brightness temperature, often many orders of magnitude above $10^{12}\\,$K. \\citet{kellermannpaulinytoth69} have shown that, assuming the electron distribution follows a single power law, the luminosity of the inverse Compton scattered photons exceeds that of the synchrotron photons when the brightness temperature of the source reaches $\\sim10^{12}\\,$K. Above this threshold, rapid cooling of the relativistic electrons due to inverse Compton scattering --- the  \\lq\\lq Compton catastrophe\\rq\\rq\\ --- forbids a further increase in the brightness temperature \\citep[see e.g.][for a recent review of the brightness temperature problem]{kellermann02}. The limiting value is even lower, $T_{\\rm B}<10^{11}\\,$K, if the magnetic field and particle energy density of the source is driven towards equipartition \\citep{readhead94}. The observed variability in some sources can be interpreted as the result of extrinsic effects, which, at first sight, relaxes the size constraint. For example, the flux variations of PKS~1519$-$273 and PKS~0405$-$385 are convincingly identified as interstellar scintillation. Nevertheless, all realistic models of the scintillation mechanism impose a new constraint on the size and require a brightness temperature of $T_{\\rm B}>10^{13}\\,$K in some cases \\citep{macquartetal00,rickettetal02}, far exceeding the limit imposed by the Compton catastrophe.  A prevalent feature associated with IDV sources is a flat or inverted spectrum ($\\alpha\\leq0$, with flux $F_\\nu\\propto\\nu^{-\\alpha}$) at radio-millimeter wavelengths \\citep[e.g.,][]{gearetal94,kedziorachudczeretal01}. Optically thick synchrotron emission from power-law electrons rises as $\\nu^{5/2}$, too fast to account for the observed spectra. Optically thin synchrotron emission in the scope of the conventional interpretation of the synchrotron theory has a flux $F_\\nu\\propto\\nu^{-(s-1)/2}$, where $s$ is the power-law index of the electrons ($\\diff\\nelec/\\diff\\gamma\\propto\\gamma^{-s}$). If $\\alpha=(s-1)/2\\leq0$, the number density of electrons diverges towards high $\\gamma$. Imposing a high-energy cut-off in the electron spectrum avoids the divergence and may account for the commonly observed spectral steepening at optical frequencies, but \\citet{marscher77} showed that electron spectra with $s\\leq1$ would result in a high flux between infrared and optical frequencies that is not supported by observations. The most common interpretation of the flat or inverted spectra is, therefore, a superposition of many synchrotron spectra within an inhomogeneous source \\citep[e.g.][]{debruyn76,marscher77,blandfordkoenigl79}.  In \\cite{kirktsang06}, we discussed a synchrotron self-Compton model in which the electron distribution is monoenergetic. The lack of low-energy electrons enables more GHz photons to emerge from the source, allowing a higher brightness temperature to be observed without initiating catastrophic cooling. We found that a temperature of up to $T_{\\rm B}\\sim10^{14}\\,$K at GHz frequencies is possible with only a moderate Doppler boosting factor of $\\sim10$. In \\citet{tsangkirk07}, we discussed the parameters of the monoenergetic model and showed that the assumption of equipartition of energy in the source does not prevent the Compton catastrophe. We also showed that an injection of highly relativistic electrons or strong acceleration in the source cannot produce temperatures much higher than our limit due to copious electron-positron pair production.  In this paper, we examine the spectral properties of synchrotron emission from monoenergetic electrons and from an electron distribution that is a double power law in energy, by comparing the model spectra with the observations of S5~0716+714, a BL~Lac object that is one of the brightest known IDV sources, as well as a gamma-ray blazar \\citep{hartmanetal99}. In doing so, we assume that the dominant targets for inverse Compton scattering are produced within the source (SSC model). The emission from gamma-ray blazars can also be interpreted in the context of models in which the target photons are created externally (EC model), for example in the broad line region, the accretion disk, or a molecular torus \\citep{sokolovmarscher05}. However, in many sources there is no observational evidence of a significant external photon source. This is the case for S5~0716+714, where, despite much effort over the past three decades, no emission lines have been detected \\citep[e.g.,][]{bychkova06}. Furthermore, XMM-Newton observations of S5~0716+714 in 2004 analysed by \\citet{ferreroetal06} and \\citet{foschinietal06} show two spectral components in the $0.5-10\\,$keV band, whose variability properties appear to favour the SSC interpretation. The recent extensive simultaneous observations of this object from radio to optical frequencies by \\citet{ostoreroetal06}, together with INTEGRAL pointings at GeV $\\gamma$-ray energies during the same period, provide the best test for our model.  In the following, we present the computation of the stationary electron distribution and the resulting synchrotron and inverse Compton spectra. The model spectra computed using the monoenergetic electron approximation, as described in \\citet{tsangkirk07}, are presented first. Although adequate for the radio emission, the monoenergetic model cannot reproduce the entire spectrum of S5~0716+714. We therefore investigate an electron distribution that is a double power law in energy --- a hard low-energy part that softens to a high-energy tail above a characteristic energy. In this way, the inverted optically thin radio emission is retained and complemented by nonthermal synchrotron emission from the high energy tail. In section~\\ref{parameter}, we briefly describe these injection models. The resulting stationary electron distribution is calculated in section~\\ref{stationarysoln} and used for the computation of the synchrotron and inverse Compton spectra. In section~\\ref{sed}, we compare the predictions of these models with the observed spectral energy distribution (SED) of the source to S5~0716+714. Our findings and some limitations of our approach are discussed in section~\\ref{discussion} and our conclusions presented in section~\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  Using the specific case of S5~0716+714 as an example, we confirm that it is possible to produce high brightness temperatures at GHz frequencies in compact radio sources without the onset of catastrophic cooling, provided that the radiating particles have a distribution that is sufficiently hard below a characteristic. In addition, we show qualitatively that induced Compton scattering is insignificant in sources with a low-energy electron cut-off despite the high brightness temperature, the underlying reason being the low occupation number of the photons that can couple with the electrons at the cut-off energy.  The model where an electron distribution that is a double power law in energy, peaking at $\\gammap$, is injected into the source offers more flexibility at higher frequencies in the synchrotron spectrum (from infrared to optical) at the expense of more free parameters, compared to either monoenergetic or single power-law distributions. These parameters should be constrained by simultaneous observations due to the highly variable nature of IDV sources. In the case of S5~0716+714 where such data is available, the spectral break at about $230\\,$GHz determines the value of $\\gammap$, the optical data at $5\\times10^{14}\\,$Hz gives the lower limit of $\\numax$, as well as constraining the spectral index $s_2$, and the INTEGRAL upper limits give the upper limit of $\\numax$ and also constrain the value of $r_{\\rm p}$, which in turn determines the electron density.  The example of S5~0716+714 illustrates several important spectral properties of an electron distribution with a low-energy cut-off, as described in the previous sections. The most noticeable feature is the hard, inverted, optically thin synchrotron spectrum, spanning a wide frequency range, which is a prevalent feature in compact radio sources at radio frequencies \\citep[e.g.,][]{gearetal94,kedziorachudczeretal01}. Other features are the spectral breaks at $\\nup=\\gammap^2\\nu_0$, $\\nucool=\\gcool^2\\nu_0$, and the exponential cut-off at $\\numax=\\gammamax^2\\nu_0$. This model, therefore, allows a simple homogeneous source to reproduce the common features shown by many IDV sources.  \\begin{acknowledgement} We thank Luisa Ostorero and Stefan Wagner for helpful discussions and for providing us with easy access to the observational data. We would also like to thank the anonymous referee for constructive comments and suggestions that we feel have led to a significant improvement in this paper. \\end{acknowledgement}"
    },
    "0710/0710.1939_arXiv.txt": {
        "abstract": "We present a new code aimed at the simulation of diffusive shock acceleration (DSA), and discuss various test cases which demonstrate its ability to study DSA in its full time-dependent and non-linear developments. We present the numerical methods implemented, coupling the hydrodynamical evolution of a parallel shock (in one space dimension) and the kinetic transport of the cosmic-rays (CR) distribution function (in one momentum dimension), as first done by Falle. Following Kang and Jones and collaborators, we show how the \\emph{adaptive mesh refinement} technique (AMR) greatly helps accommodating the extremely demanding numerical resolution requirements of realistic (Bohm-like) CR diffusion coefficients. We also present the \\emph{parallelization} of the code, which allows us to run many successive shocks at the cost of a single shock, and thus to present the first direct numerical simulations of linear and non-linear \\emph{multiple} DSA, a mechanism of interest in various astrophysical environments such as superbubbles, galaxy clusters and early cosmological flows. ",
        "introduction": "Diffusive Shock Acceleration (DSA) at supernova remnant blast waves is the favoured production mechanism for the production of the galactic cosmic-rays (CR). This theory, developed since the late 70s (see \\citealt{Drury1983a} for a review), has now both strong theoretical and observational supports. The theoretical grounds of the model lie in the early ideas of Fermi (\\citeyear{Fermi1949a}, \\citeyear{Fermi1954a}): the regular Fermi acceleration mechanism (known as Fermi~I, the stochastic one being known as Fermi~II) can naturally explain the formation of a power-law spectrum by a shock wave -- with a remarkable universal slope whose value $s$ depends solely on the shock compression ratio $r$ (which is always 4 for strong non-relativistic shocks). However, the acceleration process can easily be so efficient that the CR may back-react on the shock dynamics, modifying the acceleration process in a fully non-linear way, and requiring a much more detailed analysis (see \\citealt{Malkov2001c} for a review).  Thus the DSA mechanism has still received a lot of attention in the last 20 years, from both a theoretical and a numerical perspective. Analytical works have been mostly limited to the test-particle (linear) regime. The full non-linear time-dependent problem has been mostly investigated through numerical simulations, using several different approaches (see \\citealt{Jones2001a} for a short review). A first class is based on particle methods, from the early Monte-Carlo simulations developed by \\cite{Ellison1984a} to the recent Particle-In-Cells codes (eg \\citealt{Dieckmann2000a}). An alternate approach consists of solving the (Fokker-Planck) transport equation. This has first been done in the \"two fluid\" model (eg \\citealt{Jones1990a}), then dealing with the complete particle distribution function (\\citealt{Falle1987a}, \\citealt{Bell1987a}, \\citealt{Kang1991a}, \\citealt{Duffy1992a}).  Most of this work has been aimed at understanding the role of single (isolated) supernovae remnants, although in many contexts CR are likely to experience many shocks, most notably inside superbubbles (\\citealt{Parizot2004a}).  In this paper we present a new code for the study of DSA, named \\emph{Marcos} for \\emph{Machine {\\`a} Acc{\\'e}l{\\'e}rer les Rayons COSmiques}\\footnote{the French for \\emph{COSmic-Rays Acceleration Machine}}. In section~\\ref{sec-dsa} we present the basics of the numerical methods implemented in our code, which couples the hydrodynamical evolution of a fluid with the kinetic transport of the CR. In section~\\ref{sec-scales} we present the Adaptive Mesh Refinement (AMR) technique which allows us to resolve the (very) different scales induced by CR diffusion. In section~\\ref{sec-multi} we present parallelization of the code, to be able to study in reasonable wall-clock time the effects of multiple shocks. ",
        "conclusions": "We have presented a new code aimed at the simulation of time-dependent non-linear diffusive shock acceleration. It is based on the kinetic approach, coupling the hydrodynamical evolution of the plasma with the diffusive transport of the distribution function of the supra-thermal particles. As such it falls under the legacy of the pioneers (\\citealt{Falle1987a}, \\citealt{Duffy1992a}) and of the masters (\\citealt{Kang1991a}, \\citealt{Kang2001a}) of the genre. As the CRASH code it implements an efficient AMR technique to deal with the huge range of space- and time-scales induced by CR diffusion of Bohm-like type. To save even more on computing time we have also parallelized our code in momentum so that we can study acceleration by multiple shocks as fast as acceleration by a single shock. However in many aspects (high-Mach flows, shock tracking, self-consistent injection) our code remains numerically simpler than CRASH -- which can be both a limitation and an advantage. Regarding the physics we note that various mechanisms of importance could be included in the code: self-consistant diffusion coefficient (adding magnetic waves transport), second-order Fermi acceleration (especially between multiple shocks), electrons acceleration (adding radiative losses), CR radiation (hadronic and leptonic)\\ldots  We have presented a few tests that show that our code works well, with respect to both the physical accuracy and the numerical efficiency, even in realistic difficult situations. We are now able to investigate in details the various aspects of the DSA mechanism, which 30 years after its early developments still poses some difficulties. In particular we can address the non-linear \\emph{multiple} DSA mechanism, which we believe hasn't received so far all the attention it deserves. Our very first results suggest that the injection fraction plays a crucial role. We intend now to study in more details the situation in superbubbles. %"
    },
    "0710/0710.1108_arXiv.txt": {
        "abstract": "We study the limits of accuracy for weak lensing maps of dark matter using diffuse 21-cm radiation from the pre-reionization epoch using simulations.  We improve on previous ``optimal'' quadratic lensing estimators by using shear and convergence instead of deflection angles.  We find that non-Gaussianity provides a limit to the accuracy of weak lensing reconstruction, even if instrumental noise is reduced to zero.  The best reconstruction result is equivalent to Gaussian sources with effective independent cell of side length $2.0h^{-1}\\, \\rm Mpc$. Using a source full map from z=10-20, this limiting sensitivity allows mapping of dark matter at a Signal-to-Noise ratio (S/N) greater than 1 out to $l\\lesssim 6000$, which is better than any other proposed technique for large area weak lensing mapping. ",
        "introduction": "\\label{INTRO}  The lens mapping of dark matter is an essential cornerstone of modern precision cosmology.  Weak gravitational lensing has developed rapidly over the past years, which allows the measurement of the projected dark matter density along arbitrary lines-of-sight using galaxies as sources.  Recently, \\citet{2007PhRvD..76d3510S} have demonstrated the first CMB lensing detection.  The goal is now to achieve high precision cosmological measurements through lensing, at better than 1\\% accuracy.  Galaxies are plentiful on the sky, but their intrinsic properties are not understood from first principles, and must be measured from the data.  Future surveys may map as many as $10^{10}$ source objects. Using galaxies as lensing sources has several potential limits \\citep{2004PhRvD..70f3526H}, including the need to calibrate redshift space distributions and PSF corrections, to be better than the desired accuracy, say 1\\%.  This will be challenging for the next generation of experiments.  Some sources, such as the CMB, are in principle very clean, since its redshift and statistical properties are well understood. Unfortunately, there is only one 2-D CMB sky with an exponential damping at $l \\gg 1000$, which limits the number of source modes to $\\sim 10^6$.  The potential of detecting the 21-cm background from the dark ages will open a new window for cosmological detections. Studying the 21-cm background as high redshifts lensing source, as well as the physics of the 21-cm background itself, provide rich and valuable information to the evolution of universe.  The number of modes on the sky is potentially very large, with numbers of $10^{16}$ or more.  For this reason, 21-cm lensing has recently attracted attention. However, most of the reconstruction methods are based on a Gaussian assumption \\citep{2004NewA....9..417P,2004NewA....9..173C,2006ApJ...653..922Z, 2006astro.ph.11862B,2007arXiv0706.0849H}. In contrast to CMB lensing, where the Gaussian assumption works well, non-Gaussianity in 21-cm lensing may affect the results.  Non-linear gravitational clustering leads to non-Gaussianity, and ultimately to reionization. In this paper, we will address the problem of the lensing of pre-reionization gas.  21-cm emission is similar to CMB: both are diffuse backgrounds.  It is natural to apply the techniques used in CMB lensing. \\citet{2002ApJ...574..566H} expand the CMB lensing field in terms of the gravitational potential (or deflection angles), and construct a trispectrum based quadratic estimator of potential with maximum S/N. However, unlike CMB, the 21-cm background has a 3-D distribution and is intrinsically non-Gaussian.  A fully 3-D analysis is explored in \\citet{2006ApJ...653..922Z}, where they generalize the 2-D quadratic estimator of CMB lensing \\citep{2002ApJ...574..566H} to the 3-D Optimal Quadratic Deflection Estimator (OQDE).  A local estimator was proposed in \\citet{2004NewA....9..417P}, which assumed a power law density power spectrum. In this paper, we will design localized estimators for the lensing fields under the Gaussian assumption, and apply the derived reconstruction technique to Gaussian and non-Gaussian sources. The influence of non-Gaussianity can be measured by comparing the numerical results between the Gaussian sources and non-Gaussian sources.  Quadratic lensing reconstruction is a two point function of the lensed brightness temperature field of the 21-cm emission.  In the paper, 3-D quadratic estimators are constructed for the convergence ($\\kappa$), as well as the shear ($\\gamma$).  Our method recovers the $\\kappa$ and $\\gamma$ directly instead of gravitational potential or deflection angles. Our estimators have in principle the same form as the OQDE, consisting of the covariance of two filtered temperature maps. The OQDE reconstructs the deflection angle, while our estimators reconstruct the kappa and shear fields. Our filtering process can be written as a convolution of the observed fields. As presented in Appendix and section 4, our combined estimator is unbiased, and equally optimal as the OQDE for Gaussian sources, and has better performance for non-Gaussian sources, and recovers three extra (constant) modes.   Other authors also developed reconstruction methods from alternative approaches.  \\citet{2006astro.ph.11862B} give a estimator for shear. They choose the separate 2-D slices at certain redshift intervals, and then these slices can be treated as independent samples for the same lensing structure. As a result, the information between these slices are lost. \\citet{2004NewA....9..173C} expands the lensed field to a higher order of the gravitational potential, and investigates the higher order correction to the lensed power spectrum.  The paper is organized as follows: The basic framework of lensing and the reconstruction method is introduced in $\\S 2$. The numerical methods are presented in $\\S 3$. The results are discussed in $\\S 4$. We conclude in $\\S 5$. ",
        "conclusions": "\\label{CONC}  In this paper, we developed the maximum likelihood estimator for the large-scale structure from the 21-cm emission of the neutral gas before the epoch of re-ionization. The convergence and shears can be constructed independently.  To test the effects of non-Gaussianity, we applied our estimators to simulated data.  The sources were generated by N-body simulations, because gas is expected to trace the total mass distribution.  To investigate the influence of non-Gaussianity, we also use Gaussian sources which have the same power spectrum as the simulated sources.  We applied our estimator and the OQDE on both the Gaussian and non-Gaussian sources.  Though our estimators are derived in the simplified case of a constant convergence, the noise of our combined estimator of convergence and shear are the same as the OQDE for Gaussian sources.  For a finite survey area, three extra constant modes can be recovered.  The non-Gaussian nature of the source can increase the error bar by orders of magnitude, depending on the experimental cut off scale. Shear construction is affected less by non-Gaussianity than the convergence field, and the combined estimator with non-Gaussian noise weights is a better choice than reconstructing with the OQDE.  S/N can not be boosted infinitely by reducing the experimental noise, and achieves its maximum for a cut off around $k^{\\rm NG}_{\\rm c}\\approx 4h\\,\\Mpc^{-1}$. Below that scale the S/N start to saturate or even decrease. The maximum S/N for non-Gaussian sources is equal to Gaussian sources with $k^{\\rm G}_{\\rm c}\\approx2h\\,\\Mpc^{-1}$, where the power spectrum of source is $\\Delta^2\\approx0.2$ and the side length of the effectively independent cells is $ 2.0 h^{-1}\\,\\rm Mpc$. The maximum S/N is greater than unity for $l\\lesssim 6000$, which makes 21-cm lensing very competitive compared to optical approaches.  {\\it Acknowledgments} We thank  Oliver Zahn, Chris Hirata, Brice M\\'{e}nard and Mike Kesden for helpful discussions. T.T. Lu thanks  Pengjie Zhang, Zhiqi Huang, Hy Trac, and Hugh Merz for help in the early stage of the work."
    },
    "0710/0710.0562_arXiv.txt": {
        "abstract": "Several studies have correlated observations of impulsive solar activity --- flares and coronal mass ejections (CMEs) --- with the amount of magnetic flux near strong-field polarity inversion lines (PILs) in active regions' photospheric magnetic fields, as measured in line-of-sight (LOS) magnetograms.  Practically, this empirical correlation holds promise as a space weather forecasting tool. Scientifically, however, the mechanisms that generate strong gradients in photospheric magnetic fields remain unknown.  Hypotheses include: the (1) emergence of highly twisted or kinked flux ropes, which possess strong, opposite-polarity fields in close proximity; (2) emergence of new flux in close proximity to old flux; and (3) flux cancellation driven by photospheric flows acting fields that have already emerged.  If such concentrations of flux near strong gradients are formed by emergence, then increases in unsigned flux near strong gradients should be correlated with increases in total unsigned magnetic flux --- a signature of emergence.  Here, we analyze time series of MDI line-of-sight (LOS) magnetograms from several dozen active regions, and conclude that increases in unsigned flux near strong gradients tend to occur during emergence, though strong gradients can arise without flux emergence.  We acknowledge support from NSF-ATM 04-51438. ",
        "introduction": "\\label{sec:intro}  It has been known for decades that flares and filament eruptions (which form CMEs) originate along polarity inversion lines (PILs) of the radial photospheric magnetic field.  In studies using photospheric vector magnetograms, Falconer {\\em et al.}~(2003, 2006) \\nocite{BW_Falconer2003,BW_Falconer2006} reported a strong correlation between active region CME productivity and the total length of PILs with strong potential transverse fields ($>150$ G) and strong gradients in the LOS field (greater than 50 G Mm$^{-1}$).  They used a $\\pm$2-day temporal window for correlating magnetogram properties with CMEs. Falconer {\\em et al.}~(2003) \\nocite{BW_Falconer2003} noted that these correlations remained essentially unchanged for ``strong gradient'' thresholds from 25 to 100 G Mm$^{-1}$.  Using more than 2500 MDI (LOS) magnetograms, Schrijver (2007) \\nocite{BW_Schrijver2007} found a strong correlation between major (X- and M-class) flares and the total unsigned magnetic flux near (within $\\sim 15$ Mm) strong-field PILs --- defined, in his work, as regions where oppositely signed LOS fields that exceed 150 G lie closer to each other than the instrument's $\\sim$ 2.9 Mm resolution.  Schrijver's (2007) effective gradient threshold, {$\\sim$ 100 G Mm$^{-1}$}, is stronger than that used by Falconer {\\em et al.}~(2003, 2006). \\nocite{BW_Falconer2003,BW_Falconer2006}  Although these studies were published recently, the association between flares and $\\delta$ sunspots, which posses opposite-sign umbrae within the same penumbra --- and therefore also possess strong-field PILs --- has been well known for some time \\cite{BW_Kunzel1960,BW_Sammis2000}.  In particular, $\\beta \\gamma \\delta$ spot groups are most likely to flare \\cite{BW_Sammis2000}.  A $\\beta \\gamma$ designation means no obvious north-south PIL is present in an active region \\cite{BW_Zirin1988}.  We note that Cui {\\em et al.}~(2006) \\nocite{BW_Cui2006} found that the occurrence of flares is correlated with the maximum magnitude of the horizontal gradient in active region LOS magnetograms --- not just near PILs --- and that the correlation increases strongly for gradients stronger than $\\sim$ 400 G Mm$^{-1}$.  One would expect the measures of CME- and flare- productivity developed by both Falconer {\\em et al.}~(2003,2006) \\nocite{BW_Falconer2003,BW_Falconer2006} Schrijver (2007) \\nocite{BW_Schrijver2007} to be larger for larger active regions. Importantly, however, both studies showed that their measures of flux near strong-field PILs is a better predictor of flare productivity than total unsigned magnetic flux.  Evidently, more flux is not, by itself, as significant a predictor of flares as more flux near strong-field PILs.  These intriguing results naturally raise the question, ``How do strong-field PILs form?''  For brevity, we hereafter refer to strong-field PILs as SPILs.  Schrijver (2007) \\nocite{BW_Schrijver2007} contends that large SPILs form primarily, if not solely, by emergence.  But he also noted that flux emergence, by itself, does not necessarily lead to the formation of SPILs.  Rather, a particular type of magnetic structure must emerge, one containing a long SPIL at its emergence.  He suggests such structures are horizontally oriented, filamentary currents.  Beyond the ``intact emergence'' scenario presented by Schrijver (2007), \\nocite{BW_Schrijver2007} other mechanisms can generate SPILs. When new flux emerges in close proximity to old flux --- a common occurrence \\cite{BW_Harvey1993} --- SPILs can form along the boundaries between old and new flux systems.  Converging motions in flux that has already emerged can also generate SPILs.  If the convergence leads to flux cancellation by some mechanism --- emergence of U loops, submergence of inverse-U loops, or reconnective cancellation \\cite{BW_Welsch2006} --- then the total unsigned flux in the neighborhood of the SPIL might decrease as the SPIL forms.  We note that, while cancellation in already-emerged fields can occur via flux emergence (from upward moving U-loops), the emergence of a new flux system across the photosphere must increase the total unsigned flux that threads the photosphere.  If the emergence of new flux were primarily responsible for SPILs, then a straightforward prediction would be that an increase in total unsigned flux should be correlated with an increase in the amount of unsigned flux near SPILs.  Hence, observations showing that increases in the unsigned flux near SPILs frequently occur without a corresponding increase in total unsigned flux would rule out new flux emergence as the sole cause of these strong field gradients.  Our goal is to investigate the relationship between increases in the amount of unsigned flux near SPILs with changes in unsigned flux in the active regions containing the SPILs, to determine, if possible, which processes generate SPILs. ",
        "conclusions": "In Figure \\ref{fig:tres}, we show a scatter plot of changes in $R$ as a function of changes in ${\\cal B}$.  The plot does not show the full range in $\\Delta {\\cal B}$, but the $\\Delta R$ for outliers on the horizontal axes are near zero.  One striking feature of the plot is its flatness, i.e., that most changes in ${\\cal B}$ are not associated with any change in $R$.  In Table \\ref{tab:uno}, we tabulated the data points in each quadrant of this plot.  Clearly, increases in $R$, the unsigned flux near SPILs, usually occur simultaneously with increases in the unsigned flux over the entire active region. Increases in $R$ only occur less frequently when flux is decreasing, i.e., during cancellation.  \\begin{figure}[!ht] \\includegraphics[width=5.5in]{welb_fig03.eps} \\caption{A scatter plot of changes in $R$ as a function of changes in ${\\cal B}$.  Increases in $R$, the unsigned flux near SPILs, usually occur simultaneously with increases in the unsigned flux over the entire active region. Increases in $R$ only occur rarely when flux is decreasing, i.e., during cancellation.  For a breakdown of the data points in each quadrant, see Table \\ref{tab:uno}.} \\label{fig:tres} \\end{figure}  \\begin{table} \\caption{Breakdown of Flux Changes \\label{tab:uno}} \\begin{tabular}{l|c|c} & $\\Delta {\\cal B} < 0$ & $\\Delta {\\cal B} > 0$ \\\\ \\hline $\\Delta R > 0$ & 215 & 671 \\\\ \\hline $\\Delta R < 0$ & 363 & 371 \\\\ \\end{tabular} \\end{table}  We set out to answer the question, ``How do strong-field PILs form?'' We related changes in total, unsigned flux over whole active regions with changes in total, unsigned flux in subwindows of the same active regions --- defined by weighting maps. One might expect, therefore, that these quantities should be correlated, casting doubt about our ability to discrimintate between changes in total flux in active regions and in subwindows. If the two were strongly correlated, the excess of events with $\\Delta R > 0$ and $\\Delta {\\cal B} > 0$ might not be very meaningful.  In fact, however, $\\Delta R$ and $\\Delta {\\cal B}$ are poorly correlated: the two have a linear correlation coefficient $r = 0.29$, and a rank-order coefficient of $0.36.$ This suggests that the relationship between increases in $R$ and increases in total, unsigned active region flux is not an artifact of our approach.  Nonetheless, our active region sample is not ideally suited to address the origin of SPILs, generally.  Our sample was not unbiased with respect to active region morphology; we selected regions with well-defined PILs.  In addition, our sample included some decayed active regions that NOAA AR designations.  Consequently, we believe that a follow-up study, with a much larger, unbiased sample of active regions, is warranted.  With caveats, therefore, our study supports Schrijver's (2007) \\nocite{BW_Schrijver2007} contention that the emergence of new flux creates the strong-field polarity inversion lines that he found to be correlated with flares."
    },
    "0710/0710.5171_arXiv.txt": {
        "abstract": "We describe a method for computing the biases that systematic signals introduce in parameter estimation using a simple extension of the Fisher matrix formalism. This allows us to calculate the offset of the best fit parameters relative to the fiducial model, in addition to the usual statistical error ellipse. As an application, we study the impact that residual systematics in tomographic weak lensing measurements.  In particular we explore three different types of shape measurement systematics: (i) additive systematic with no redshift evolution; (ii) additive systematic with redshift evolution; and (iii) multiplicative systematic.  In each case, we consider a wide range of scale dependence and redshift evolution of the systematics signal. For a future DUNE-like full sky survey, we find that, for cases with mild redshift evolution, the variance of the additive systematic signal should be kept below $10^{-7}$ to ensure biases on cosmological parameters that are sub-dominant to the statistical errors.  For the multiplicative systematics, which depends on the lensing signal, we find the multiplicative calibration $m_0$ needs to be controlled to an accuracy better than $10^{-3}$. We find that the impact of systematics can be underestimated if their assumes redshift dependence is too simplistic. We provide simple scaling relations to extend these requirements to any survey geometry and discuss the impact of our results for current and future weak lensing surveys. ",
        "introduction": "\\label{intro} Weak gravitational lensing, or `cosmic shear', is undergoing  a phase of rapid expansion \\citep[see][for reviews]{2003ARA&A..41..645R,2006astro.ph.12667M,2003astro.ph.10908H} with many future surveys and instruments being planned (e.g. DUNE\\footnote{http://www.dune-mission.net}, PanSTARRS\\footnote{http://pan-starrs.ifa.hawaii.edu}, DES\\footnote{https://www.darkenergysurvey.org}, SNAP\\footnote{http://snap.lbl.gov} and LSST\\footnote{http://www.lsst.org}).  Central to the planning and designing of these instruments is our ability to predict the uncertainties that such measurements will achieve on the cosmological parameters.  To this end, the Fisher matrix has become a widely used tool in cosmology for calculating their covariance matrix.  However, a limitation of this approach is that it is only able to account for statistical errors, i.e. ones that cause an enlargement of the error bars, and is not well-suited for treatment of systematic errors, which can introduce biases that move the measured central value relative to its true value.   One approach that is commonly taken to overcome this limitation is to treat the systematic errors in the same way as statistical errors and to marginalise over possible values. This introduction of nuisance parameters, in addition to the cosmological parameters, causes the error ellipses to expand.  A more accurate approach, which we use here, is to directly calculate the bias that the systematic signals will introduce.  This bias will tend to offset the central value to the measurements from the true values, as shown in figure \\ref{fig:fig0}.  Since the computations needed for this calculation are very similar to those performed in the standard Fisher matrix analysis, extending the current Fisher matrix analysis to include a calculation of bias is relatively straightforward.  We apply this formalism to study the impact of  residual systematics on tomographic cosmic shear surveys. In particular, we consider systematics arising in the measurement of galaxy shapes after correction of instrumental effects (such as the Point Spread Function). We consider both additive and multiplicative systematics and explore a wide range of scale and redshift dependences. In an earlier work, \\cite{2006MNRAS.366..101H} considered the impact of photometric calibration errors and power law shape systematics using the Fisher matrix formalism. A similar bias formalism was introduced by \\cite{2005APh....23..369H} and applied to theoretical uncertainties in modeling the matter power spectrum with N-body simulations.  Our work expands upon these earlier works, by appying the bias formalism to a broad set of shape systematics and by studying the joint impact of systematic and statistical errors in current and future surveys.  This paper is organised as follows.  In section 2, we describe the formalism that we use to quantify systematic biases.  In section 3, we apply our formalism to cosmic shear surveys by exploring the effect of three types of shape measurement systematics: (i) additive with no redshift evolution; (ii) additive with redshift evolution; and (iii) multiplicative. For each type, we consider several possibilities for their scale dependence: (i) log-linear systematics; (ii) systematics that have the same shape as the lensing signal; and  (iii) systematics that mimic a small change in the cosmological parameters.  In section 4, we study the impact of the systematics in the design of future surveys. Our conclusions are summarised in section 5. ",
        "conclusions": "\\label{conclusion}  In this paper, we have outlined a method for computing the biases that residual systematics introduce. This approach involves a simple extension of the Fisher matrix formalism that is now widely used in cosmology to make error forecasts.  As an application, we have used it to study the impact that residual systematic signals will have on future tomographic cosmic shear measurements.  Specifically, we have explored  three different types of shape systematic signal affecting tomographic shear power spectra: (i) additive systematics with no redshift evolution; (ii) additive systematics with redshift evolution; and (iii) multiplicative systematics. The requirement target is then to have all types of systematics close to zero.  This defines a tolerance envelope for the systematics that allows the residual systematics errors in the power spectra to be positive or negative within its limits. It is important to note that it is the $worst$ systematic possible within this limit which drives the requirement, not a marginalised average over all systematics. To this end we have investigated a wide class of possible systematic shapes and used the most constraining ones to set our systematic requirements.  In doing this, we have found that, for both the additive and multiplicative parts, it is vital to consider systematics that have positive and negative power spectra $C_\\ell^{sys}$.  For instance we see, in the multiplicative case, that investigating only power-law behaviour for its redshift evolution (i.e. $m$ is always positive) can lead to a factor of 5 underestimation of the impact of a systematic within a given tolerance window.  From our calculation we are able to set the following requirements on the survey we have considered (a DUNE-like survey covering 20,000 sq. degrees with 35 galaxies per arcmin$^{-2}$ and a median redshift of $z_m=0.9$): \\begin{itemize} \\item For both the additive and the multiplicative signals, the redshift evolution needs to be weak $\\beta < 1.5$, where the errors  for a given galaxy scale as $(1+z)^\\beta$. This is not a trivial requirement since the shapes of more distant galaxies are harder to measure since they are smaller and fainter. \\item The power spectrum of the residual additive shear error, that is the part that is not correlated to the lensing signal, must be controlled such that its amplitude is $\\sigma^2_{sys} < 10^{-7}$ (as defined in equation \\ref{eq:var}) \\item The multiplicative part needs to be controlled to a precision of $m_0 <10^{-3}$, where $m_0$ is the shear calibration error.  This means that we need to be able calibrate shears to an accuracy of 0.1\\%, which is about one order of magnitude better than the current best measurement methods are able to achieve, as determined by the latest STEP simulations \\citep{2006MNRAS.368.1323H,2007MNRAS.376...13M}. \\end{itemize}  These specific requirements apply to our fiducial survey, but we provide scaling relations (Eqs \\ref{eq:scale1} and \\ref{eq:scale2}) which show the requirements for any survey geometry.  We have shown that for current survey covering $\\sim 100$ deg$^{2}$ \\citep{2007MNRAS.381..702B}, we need $\\sigma^2_{sys} < 3\\times10^{-6}$ and $m< 0.03$. This level of accuracy is at the limit of the performance of the current best shear measurement methods, as demonstrated by STEP. However, further systematics, such as that arising from PSF calibration and interpolation, are not accounted for by STEP and can dominate the error budget for future surveys. A discussion of the requirements for additive systematics and PSF modeling in the context of present and future surveys will be presented in a later paper (Paulin-Henriksson et al., 2007, in prep).  \\newpage \\appendix"
    },
    "0710/0710.3754.txt": {
        "abstract": "We present arcsecond-scale mid-ir photometry (in the 10.5 $\\mu$m N band and at 24.8 $\\mu$m), and low resolution spectra in the N band ($R\\simeq100$) of a candidate high mass protostellar object (HMPO) in IRAS 18151-1208 and of two HMPO candidates in IRAS 20343+4129, IRS 1 and IRS 3. In addition we present high resolution mid-ir spectra ($R\\simeq80000$) of the two HMPO candidates in IRAS 20343+4129.  These data are fitted with simple models to estimate the masses of gas and dust  associated with the mid-ir emitting clumps, the column densities of overlying absorbing dust and gas, the luminosities of the HMPO candidates, and the likely spectral type of the HMPO candidate for which [Ne II] 12.8 $\\mu$m\\ emission was detected (IRAS 20343+4129 IRS 3).  We suggest that IRAS 18151-1208 is a pre-ultracompact HII region HMPO, IRAS 20343+4129 IRS 1 is an embedded young stellar object with the luminosity of a B3 star, and IRAS 20343+4129 IRS 3 is a B2 ZAMS star that has formed an ultracompact HII region and disrupted its natal  envelope. ",
        "introduction": "Many open questions in high-mass star formation are related to the evolution of circumstellar envelopes, accretion disks,  and jets from high-mass protostellar objects (HMPOs). HMPOs are often bright sources in mid-ir continuum, but only in a few recent cases have images suggested specific structures such as disks, jets, or warm outflow cavity walls \\citep{sri05, deb05, deb06,deb07}.   Mid-ir ionic lines like [Ne II] and [S IV] have been used to map  compact and ultracompact HII regions (UC HII) and photodissociation regions, and to study their structure and excitation \\citep{lac82,oka01,zhu05,kas02,kas06}, but there is a lack of observations of HMPOs. A remarkable feature of high-mass star formation is that HMPOs, defined as actively accreting mass, can begin nuclear fusion and hence also be rapidly evolving massive young stellar objects (MYSOs) that have already formed hypercompact or ultracompact HII regions \\citep{beu07}.  This feature raises a possibility of determining the spectral type of an MYSO through the ionic lines excitation, or from the number of ionizing photons required for its observed centimeter continuum flux, separately from estimating its luminosity and spectral type from infrared emission. However there is also the possibility that the  ionization is collisionally excited by a jet.  It may be possible to distinguish between the two cases, depending on the Doppler velocities, the morphology of ionized gas, and the ratio of the required flux of ionizing photons to total luminosity.  Hoping to enlarge the sample of HMPOs that could be studied in detail despite potential limitations, in 2003 we made mid-ir observations on the IRTF of about a third of the survey of 69 HMPO candidates presented by \\citet{sri02} and \\citet{beu02a}. We chose objects from their survey that appeared compact  and/or bright in the MSX survey, and found that about 80\\% of them were unresolved or marginally resolved by MIRSI \\citep{deu02} on the IRTF in the broad N band at 10.5 $\\mu$m and in a narrow band filter at 24.8 $\\mu$m.  In addition, we obtained MIRSI grism low-resolution  spectra (R $\\simeq100$) of ten of them in the N band.  In 2006 on Gemini North, we obtained TEXES \\citep{lac02} high-resolution spectra (R $\\simeq80000$)  of two HMPO candidates for which we had grism spectra.   In this paper we present spectra and photometry of three candidate HMPOs including the two with TEXES spectra: IRAS 18151-1208, IRAS 20343+4129 IRS 1, and IRAS 20343+4129 IRS 3 (hereafter, 18151, 20343 IRS 1, and 20343 IRS 3). We will demonstrate that mid-ir emission from the dust  and gas near the HMPO candidates (where it is strongly heated) can be used as a useful probe of temperatures, masses, and luminosities, using simple isothermal clump models, even if each component (envelope, disk, jet, or cavity wall)  is not resolved.   In combination with observations cited below, we are able to use our new data to infer the nature of each candidate candidate HMPO (e.g. pre-UCHII region HMPO, ZAMS B2 star).  The objects chosen are near the centers of complex, large-scale massive molecular outflows mapped by \\citet{beu02b}.  20343 has an apparent large-scale N-S outflow whose red and blue lobes are both extended E-W \\citep{beu02b}, but IRS 1 also has a compact E-W velocity outflow in CO(2-1), while  IRS 3 presents an ambiguous situation \\citep{pal06}. All objects show near-ir emission from shocked H$_2$ \\citep{dav04,kum02}.  Two of them (18151 and 20343 IRS 3) were observed to have 0.5 and 1.8 mJy 3.6 cm emission, respectively \\citep{car99,sri02}.  We observed 20343 IRS 1 and IRS 3  with TEXES on Gemini North based on the 3.6 cm and H$_2$ observations, with the goal of studying the role of ionized gas in them.  The 10 $\\mu$m grism spectral shapes of the HMPO candidates fall into three classes: those with deep silicate absorption; those with moderate silicate absorption and an apparent peak at about 8.5 $\\mu$m; and those without an apparent silicate absorption feature but with continuum rising monotonically from short to long wavelengths (Campbell et al. 2007, in preparation).    Examples of these shapes can be seen in the UC HII spectra presented by \\citet{fai98}.  Since the HMPO candidates were chosen based on IRAS colors similar to UC HII regions \\citep{sri02}, one would expect the HMPO candidates to have similar 10 $\\mu$m spectra. IRAS 20343+4129, was observed with the IRAS LRS and has a silicate absorption feature \\citep{vol91}.  The three objects presented here include an example of each of the three classes of grism spectral shapes. Two of the ten candidate HMPO grism spectra show strong [Ne II] lines, IRAS 18247-1147 and 18530+0215.  \\citet{sri02} reported relatively strong 3.6 cm fluxes of 47 and 311 mJy, respectively, for them.  We did not observe them with TEXES on Gemini in order to focus on the earliest possible HMPO stage associated with ionized gas.    Deriving physical information from the continuum spectra is difficult because the actual geometry of the dust distribution is unknown, except that the sizes of the N band and 24.8 $\\mu$m images limit the extent of the emitting structures.  Experience with one-dimensional radiative transfer models of spectral energy distributions from UC HII regions \\citep{cam95,cam00,cam04}, and inspection of the new spectra themselves suggest that there are ranges of temperatures in the emitting regions.  However, the two-dimensional radiative transfer models of \\citet{deb05b} and \\citet{whi03a,whi03b,whi04} show that orientation of flattened envelopes and outflow cavities dramatically affects the depth of the silicate absorption feature, as does emission and absorption by individual clumps in a clumpy dust cloud in the three-dimensional models of  \\citet{ind06}. The recent observations of extended and complex near- and mid-ir emission from HMPOs \\citep{deb05,deb06,sri05} also indicate that one-dimensional models are unrealistic.  Nevertheless, a simple three part model will allow us to derive approximate temperatures, column densities, and masses of different dust components. The first component is hot dust that could be (part of)  a relatively compact disk, or a clump very near the HMPO candidate; the second is warm dust that could be a more extended (part of a) disk, a clump of dust further out, or perhaps  the inner wall of of an outflow cavity; and the third is cold dust in an outer envelope that creates the silicate absorption feature.  Deriving information from the ionic lines of an UC HII region is in principle straight-forward.  The line fluxes can be   corrected for local extinction based on the continuum models discussed above, and then the ratio of [Ne II] to [S IV] fluxes can be used to determine the exciting star's temperature \\citep{lac82,oka01}.  The numerical simulation code cloudy (http://www.nublado.org/) can also be used to deduce the star's temperature and luminosity by matching the simulated intensity and spatial extent of free-free, [Ne II], and [S IV] emission to the observations. The star's temperature can then be compared to that of the spectral type deduced from the cm continuum flux. In addition, comparison of Doppler velocities of the ionic lines to those of molecular lines can be used to indicate if the gas is in an UC HII region or a jet. ",
        "conclusions": "We have presented high resolution mid-ir observations made with MIRSI on the IRTF and TEXES on Gemini North of three  HMPO candidates taken from a partial follow up survey of HMPO candidates originally studied at 1.2 mm and radio wavelengths by \\citet{sri02} and \\citet{beu02a}.  They are typical for HMPO candidates observed in the follow up survey being compact at 1$\\arcsec$ resolution,  having low resolution spectra with  strong, moderate, or weak silicate absorption, and with  one  emitting the [Ne II] line.  A simple model of hot dust in emission, warm dust in emission, and cold dust in absorption was developed to fit our 8-13 $\\mu$m low resolution spectra and our 24.8 $\\mu$m photometric points.  Even an apparently flat 8-13 $\\mu$m spectrum requires an absorption component if the underlying emission is assumed to be due to hot or warm silicate dust.  The temperatures ranged from $\\sim$ 400-1000 K for the hot dust, and $\\sim$ 100-200 K for the warm dust.  Using \\citet{dra03a} $R_V$=5.5 model dust properties and gas-to-dust ratio, only small masses of gas and dust in the two emitting components are needed to fit the data.  The masses are less than about 1/10 solar mass (often much less) even though these are high or intermediate mass stars, and the mid-ir emission cannot be due the the bulk of the mass in massive accretion disks.   The mid-ir is likely to be emitted by the inner walls of outflow cavities and perhaps partly by the surfaces of accretion disks. On the other hand, high column densities, $10^{22} - 10^{23}$ H nucleons cm$^{-2}$, are required for the cold absorption components.   These column densities are less than derived from 1.2 mm 11$\\arcsec$ data using \\citet{dra03a} dust, but the discrepancy may be resolved if the slope of the absorption coefficient from far-ir to mm is flattened, as suggested by some observations.   Our three component model is not meant to fit either near-ir or far-ir to mm ends of SEDs.  Nevertheless, the dust we are modeling in the hot and warm components appears to absorb the bulk of an HMPO's or intermediate mass YSO's  photospheric emission, so that the integrated flux of the two model components without application of the cold dust's extinction matches the luminosity as measured including the far-ir by IRAS. The mid-ir measurements together with the model thus give a reasonable way to determine the luminosity for individual HMPOs.  The mid-ir emission of IRAS 18151-1208 together with weak 3.6 cm emission and other previous observations suggest that it is an early stage  pre-UC HII HMPO whose luminosity is that of a B0 star.  TEXES high resolution spectra that cover emission lines from ionized gas can be used to determine the nature of the emission (jet or HII region) and help determine the properties of the underlying star. In the case of IRAS 20343+4129 IRS 1,  a lack of [NeII] emission, a well defined compact CO outflow, a moderately strong silicate absorption feature, and  a dust model-based luminosity of 1400 L$_\\sun$ imply that it is an intermediate mass YSO whose luminosity is that of  a B3 star. For IRAS 20343+4129 IRS 3 observed  [Ne II] emission and 3.6 cm free-free emission  are consistent with a cloudy model indicating that the object is a B2 ZAMS star.  Its weak silicate absorption and small mid-ir based luminosity suggest that it has already disrupted much of its natal envelope. % For IRAS 20343+4129 IRS 3, observed  [Ne II] emission and % 3.6 cm free-free emission  are consistent with a cloudy model % indicating that the object is a B2 ZAMS star.  Its weak % silicate absorption and small mid-ir based luminosity suggest that it has already disrupted % much of its natal envelope.         %% If you wish to include an acknowledgments section in your paper, %% separate it off from the body of the text using the"
    },
    "0710/0710.0879_arXiv.txt": {
        "abstract": "We present the first fully relativistic longterm numerical evolutions of three equal-mass black holes in  a system consisting of a third black hole in a close orbit about a black-hole binary.  We find that these close-three-black-hole systems have very different merger dynamics from black-hole binaries.  In particular, we see complex trajectories, a redistribution of energy that can impart substantial kicks to one of the holes, distinctive waveforms, and suppression of the emitted gravitational radiation.  We evolve two such configurations and find very different behaviors. In one configuration the binary is quickly disrupted and the individual holes follow complicated trajectories and merge with the third hole in rapid succession, while in the other, the binary completes a half-orbit before the initial merger of one of the members with the third black hole, and the resulting two-black-hole system forms a highly elliptical, well separated binary that shows no significant inspiral for (at least) the first $t \\sim 1000M$ of evolution. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.1491_arXiv.txt": {
        "abstract": "{Mass loss from red giants in old globular clusters affects the horizontal branch (HB) morphology and post-HB stellar evolution including the production of ultraviolet-bright stars, dredge up of nucleosynthesis products and replenishment of the intra-cluster medium. Studies of mass loss in globular clusters also allows one to investigate the metallicity dependence of the mass loss from cool, low-mass stars down to very low metallicities. } {We present an analysis of new VLT/UVES spectra of 47 red giants in the Galactic globular clusters 47 Tuc (NGC 104), NGC 362, $\\omega$ Cen (NGC 5139), NGC 6388, M54 (NGC 6715) and M15 (NGC 7078). The spectra cover the wavelength region 6100--9900 \\AA\\ at a resolving power of $R = 110,000$. Some of these stars are known to exhibit mid-infrared excess emission indicative of circumstellar dust. Our aim is to detect signatures of mass loss, identify the mechanism(s) responsible for such outflows, and measure the mass-loss rates.} {We determine for each star its effective temperature, luminosity, radius and escape velocity. We analyse the H$\\alpha$ and near-infrared calcium triplet lines for evidence of outflows, pulsation and chromospheric activity, and present a simple model for estimating mass-loss rates from the H$\\alpha$ line profile. We compare our results with a variety of other, independent methods.} {We argue that a chromosphere persists in Galactic globular cluster giants and controls the mass-loss rate to late-K/early-M spectral types, where pulsation becomes strong enough to drive shock waves at luminosities above the RGB tip. This transition may be metallicity-dependent. We find mass-loss rates of $\\sim$10$^{-7}$ to $10^{-5}$ M$_{\\odot}$ yr$^{-1}$, largely independent of metallicity.} {} ",
        "introduction": "Understanding the evolution and mass loss in red giant branch (RGB) and asymptotic giant branch (AGB) stars is of substantial importance in uncovering the history of metal enrichment of the inter-stellar medium (ISM) and subsequent stellar and planetary formation; as well as understanding the evolution of the stars themselves. By observing mass-losing giant branch stars in globular clusters (GCs), we can simultaneously observe a co-eval set of objects within a cluster with similar progenitor mass and identical metallicity; while we can use different clusters to observe the effects of differing ages and masses, but more particularly metallicities, which can range from near-solar metallicity to as little as $\\sim$0.5\\% of solar.  Simultaneously, we also glean information about the globular cluster environment itself -- the mass lost from stars forms the intra-cluster medium (ICM). Studies of the ICM show it is removed from the cluster on a timescale much shorter than the time between Galactic Plane crossings (e.g. Boyer et al.\\ 2006; van Loon et al.\\ 2006a). This mass loss exacerbates cluster evaporation, possibly leading to their dispersal on a timescale of $10^{9-10}$ years (van Loon \\& McDonald 2007).  Both the RGB and AGB are associated with episodes of strong mass loss, eventually leading to stellar death. Globular cluster stars are thought to lose around 0.2 M$_{\\odot}$ over their time on the RGB (Rood 1973) with up to 0.1 M$_{\\odot}$ being lost in the helium flash at the RGB tip (Fusi-Pecci \\& Renzini 1975, and references therein). This is required to explain the horizontal branch (HB) morphology, which is dependent on the remaining mantle mass. Dupree et al.\\ (2007) have now observationally confirmed mass loss three magnitudes below the RGB tip, thought to be driven by hydrodynamically- or acoustically-dominated chromospheres.  Stars reaching the upper slopes of the AGB are even more extreme, losing mass at rates reaching well over $10^{-6}$ M$_{\\odot}$ yr$^{-1}$ (Wood et al.\\ 1983, 1992; van Loon et al.\\ 1999). This is facilitated by thermal pulses -- periodically occurring temporary ignition of a helium shell-burning source inside the star -- and shorter-timescale ($\\sim$10$^{2-3}$ days) radial pulsations at the surface which can lead to shock-wave emission in the extended stellar atmosphere (Fox \\& Wood 1985). In metal-rich AGB stars, dust forming in the outmost parts of these atmospheres is subject to radiation pressure, which forces it and the gas it collides with to be radially accelerated away from the star. It is not yet clear if dust-driving is the dominating effect in RGB and metal-poor AGB winds, and it is possible that Alfv\\'en waves or pressure (P-mode) waves dominate here (Hartmann \\& MacGregor 1980; Pijpers \\& Habing 1989). Indeed, Judge \\& Stencel (1991, hereafter JS91) argue that the mass-loss rate is largely invariant with the driving process and that most cool giants (RGB or AGB) do not have dust as the primary driver of their winds. More recently, Schr\\\"oder \\& Cuntz (2005, hereafter SC05) have suggested an improved Reimers relation (Reimers 1975), by assuming this invariance and using simple chromospheric energy considerations, which is calibrated using the total RGB mass-loss of GC stars.  Previous works have attempted to provide evidence for mass loss and calculate mass-loss rates by analysing optical spectral line profiles, either through asymmetries in the line cores (first performed by Deutsch (1956), for the field giant $\\alpha$ Her), line emission (Cohen 1976) or modelling. Line emission has also been used to model a chromosphere with a large bulk motion (Dupree et al.\\ 1984; Mauas et al.\\ 2006). Semi-empirical methods of finding mass-loss rates have also been used to define relationships based on stellar parameters. Mass-loss rates can also be derived from dust masses calculated from IR excess emission (e.g. Origlia et al.\\ 2002 -- hereafter OFFR02; van Loon 2007), though this must assume a value for the ratio of the masses of dust and gas lost. Mira-type long-period variability is only seen in stars with [Fe/H] $\\gtrsim -1$ (Frogel \\& Whitelock 1998), so one would expect that, assuming pulsation assists in driving mass-loss, the mass-loss rate and/or mechanism would vary with metallicity.  In this work, we use all of the above techniques to try to identify mass loss and, where possible, estimate mass-loss rates from a sample of 47 such objects from a range of clusters. Ultimately, we aim to assess the r\\^ole of dust, pulsation and chromospheric activity in driving mass loss, and its dependence on metallicity. Our observations are based on new VLT/UVES echelle spectra at very high spectral resolution. We present the spectra and literature photometry in Section 2 of the paper. In Section 3, we use the Kurucz {\\sc atlas9} models (Kurucz 1993) to estimate the effective temperature, metallicity and surface gravity of the stars, then use the near-IR photometry to calculate other physical parameters (luminosity, radius and escape velocity). A full abundance analysis has not been performed, but will become the subject of a subsequent paper. In Section 4, we analyse the H$\\alpha$ and near-IR Ca II triplet lines and from them provide evidence of outflow from the stellar surface. We present a simple model for estimating mass loss in the stellar wind from the H$\\alpha$ profile in Section 5, and use it to calculate mass-loss rates and wind velocities for our sample of stars. This is followed by a discussion in Section 6, in which we also estimate mass-loss rates via other procedures. Our conclusions are presented in Section 7. ",
        "conclusions": "In this study, we have presented VLT/UVES data of a set of giant branch stars in globular clusters, which are among the highest resolution, high signal-to-noise spectra of their kind. We use these to characterise and quantify the outflow from their atmospheres.  We have used Kurucz's {\\sc atlas9} models to determine stellar temperature and, from this, we have calculated basic stellar parameters, which appear typical for red and asymptotic giant branch stars.  Many of the stars we have investigated show clear emission in H$\\alpha$, most spectacularly when strong pulsation is present, which is seen to occur at luminosities above the RGB tip. It is also noted that many of the stars show strongly blue-shifted absorption cores, suggesting bulk outflow from the stellar surface. This is also mirrored in some cases in the near-infrared calcium triplet line profiles.  In an effort to quantify this outflow, we have constructed a simple model for the warmer and more metal-poor stars. We have calculated terminal velocities for the winds which are of order 10 km s$^{-1}$ -- much lower than the escape velocity of the stars -- and mass-loss rates of $\\approx$ 10$^{-7}$ to 10$^{-5}$ M$_\\odot$ yr$^{-1}$, which lie well above theoretical expectations, but are consistent with the mass-loss rates derived from IR emission from circumstellar dust. These models also suggest that an emissive shell exists close to the stellar surface in many stars, with the wind probably being largely isothermal beyond it.  Outflow velocities of the emission are of order 40 km s$^{-1}$, and mass-loss rates derived from this emission are similar to the ones we obtain from modelling the absorption profile. In the most extreme cases we may have either over-estimated the mass-loss rate from our model, or under-estimated the temperature of the emitting shell. Mass-loss rates correlate weakly with luminosity, but stars showing strong IR excesses (linked with dust production) do not necessarily exhibit higher gas mass-loss rates. We find no correlation between mass-loss rate and metallicity.  We suggest that the emission and mass-loss in early-type ($\\lsim$ K3--K5) giants are dominated by a warm chromospheric region, as suggested by some studies, though late-type ($\\gsim$ M3) giants have mass loss dominated by pulsation. It seems likely that the spectral type of the transition between the two regimes is metallicity-dependent, occurring at later spectral types for higher metallicities. The outward velocity of the warm emissive shells associated with the chromospheres is similar to the escape velocity at their anticipated radius (2 to 3 R$_\\ast$), allowing the chromosphere to be the sole driver of mass loss in these stars."
    },
    "0710/0710.5421_arXiv.txt": {
        "abstract": "\\footnotesize\\ This paper\\footnote{Published in 1999 in the book {\\it Solar Polarization}, edited by K.N. Nagendra \\& J.O. Stenflo. Kluwer Academic Publishers, 1999. (Astrophysics and Space Science Library ; Vol. 243), p. 73-96} addresses the modelling issue of the second solar spectrum. This is the name given to the linearly polarized spectrum which can be observed close to the solar limb using spectro-polarimeters of high polarimetric sensitivity (Stenflo and Keller, 1997). The second solar spectrum is due to scattering processes and offers a rich diagnostic potential for exploring solar magnetic fields via the Hanle effect. However, it is full of mysterious spectral features that cannot be understood with simplified polarization transfer theories, thus suggesting that the underlying scattering physics is more complex than previously thought. In this paper we argue that understanding the second solar spectrum requires the consideration of scattering processes in {\\it multilevel} atomic models, taking fully into account the transfer of {\\it atomic polarization} among {\\it all} the levels involved. To give support to this statement, we begin by pointing out the drastically different predictions, given by the standard resonance line polarization theory, with respect to the emergent polarization in three {\\it different} line transitions. This standard theory neglects the atomic polarization of the {\\it lower} level of the line transition under consideration, i.e. it assumes that there are no population imbalances among the lower-level sublevels. The density matrix polarization transfer theory is then applied to formulate the scattering polarization problem taking properly into account atomic polarization in both the upper and the lower line levels. The consideration of lower-level atomic polarization leads to coupled {\\it non-linear} and {\\it non-local} sets of equations, even for the two-level model atom case considered in this paper. The {\\it unknowns} of these equations are the irreducible tensor components of the atomic density matrix whose {\\it self-consistent} values have first to be obtained to be able to calculate the emergent Stokes profiles. To solve this non-LTE problem of the $2^{nd}$ kind we present some iterative methods that are very suitable for developing a general multilevel scattering polarization code. With these numerical methods some model calculations are performed in order to demonstrate that the inclusion of lower-level atomic polarization leads to similar emergent linear polarization signals in such three {\\it different} line transitions, as some observations show. After pointing out that the ``Na solar paradox'' (Landi Degl'Innocenti, 1998) might admit, in principle, a {\\it multilevel} solution, the paper ends establishing a new solar paradox: ``the Mg solar paradox'', for which no {\\it multilevel} solution seems to be possible. This new result demonstrates that there indeed exists {\\it ground} and {\\it metastable}-level atomic polarization in the solar chromosphere and it suggests that the solution to these ``solar paradoxes'' is to be found by carefully revising our current ideas about the chromospheric magnetic field. ",
        "introduction": "The ``second solar spectrum'' is a term adopted by Stenflo and Keller (1997) to refer to the remarkable observational discovery that the whole solar spectrum is linearly polarized, when observations are made close to the solar limb using novel spectro-polarimeters that allow the detection of very low amplitude polarization signals ($\\sim10^{-5}$ in the degree of linear polarization). This term is certainly adequate because, as pointed out by these authors, the linearly polarized solar spectrum has a structural richness that often exceeds that of the ordinary intensity spectrum. It is indeed as if the Sun has presented us with an entirely new spectrum to explore. In fact, the second solar spectrum contains a wealth of ``inexplicable'' spectral features which are the signature of physical processes that are presently challenging physicists working in the field.  This paper deals with the modelling issue of the second solar spectrum. This modelling requires the solution of a formidable numerical problem that is considered as one of the most challenging tasks of solar and stellar polarimetry. It consists in calculating, for {\\it multilevel} atomic models, the excitation and ionization states of chemical species of given abundance that are {\\it consistent} with the polarization properties of the radiation field produced by such species in any medium of given temperature, density, macroscopic velocity and magnetic field vector. Once this self-consistent atomic excitation is known along the line of sight, it is straightforward to solve the transfer equations for the Stokes parameters in order to calculate the emergent polarization profiles that are to be compared with spectro-polarimetric observations.  It is indeed a very complex problem because, in the polarization transfer case, one has to take into account that each level of total angular momentum value X has associated with it (2X+1) sublevels, with ${\\vec {\\rm X}}={\\vec {\\rm J}}= {\\vec {\\rm L}}+{\\vec {\\rm S}}$ if fine structure due to the spin-orbit LS coupling is assumed, or with ${\\vec {\\rm X}}={\\vec {\\rm F}}={\\vec {\\rm J}}+{\\vec {\\rm I}}$ if hyperfine structure due to the nuclear angular momentum $\\vec {\\rm I}$ is taken into account. The populations of these sublevels are sensitive to the polarization and anisotropy state of the radiation field at each point within the medium. Moreover, quantum interferences (or coherences) among the sublevels themselves may also appear, coherences that depend on the energy separation between the levels and on their splitting. These coherences must also be properly quantified to fully specify the excitation state. An additional complication stems from the fact that, in the polarized case, instead of the standard radiative transfer (RT) equation for the specific intensity, one has to solve, in general, a vectorial transfer equation for the four Stokes parameters.  Obviously, accounting for this complexity requires working within the framework of a robust theory for the generation and transfer of polarized radiation. I believe that the most rigorous (and suitable) theoretical framework to work with is that developed by Landi Degl'Innocenti (1983), which is based on the irreducible tensor components ($\\rho^K_Q$) of the atomic density matrix (see also Bommier and Sahal-Br\\'echot, 1978). According to this formalism, to each level of angular momentum value X, there correspond $\\rm (2X+1)^{2}$ density-matrix elements. These $\\rho^K_Q$-elements contain information about the populations of the atomic sublevels and about the coherences among them. For instance, for a level with total angular momentum J=1 we have that $\\rho_0^0=(N_1+N_0+N_{-1})/{\\sqrt{3}}$, $\\rho^1_0=(N_1\\,-\\,N_{-1})/{\\sqrt{2}}$ and $\\rho_0^2=(N_1-2N_0+N_{-1})/{\\sqrt{6}}$, where $N_{i}$ (with $i$=1, 0 and -1) are the populations of the three magnetic sublevels. Thus, $\\rho_0^0$ gives the total population of the level, while $\\rho^1_0$ (the {\\it orientation} coefficient) and $\\rho^2_0$ (the {\\it alignment} coefficient) inform us about the {\\it population differences} among the sublevels. Finally, $\\rho^K_Q$-terms with $Q\\neq0$ account for the coherences between Zeeman sublevels whose magnetic quantum numbers differ by $Q$.  One of the reasons that explain why the density matrix formalism is so suitable for dealing with the generation and transfer of polarized radiation is that, through an emission process, polarization in spectral lines can originate locally either by the splitting of the atomic levels (splitting that can in turn be due either to the Zeeman or the Stark effect) or by the presence of {\\it population differences} and/or {\\it coherences} among the sublevels. Thus, the $\\rho^K_Q$ elements whose self-consistent values are sought at each spatial grid-point, indeed provide the most suitable way of quantifying, at the atomic level, the information that we need to be able to calculate all the ``sources'' and ``sinks'' of polarization within the medium under consideration. The main criticism of this QED theory is that it is based on the approximation of complete frequency redistribution (CRD), i.e. on an assumption that is not adequate for modelling the polarization of several diagnostically important spectral lines. Fortunately, some very recent work has successfully started to incorporate the effects of partial redistribution (PRD) into the framework of the density matrix formalism (see Landi Degl'Innocenti {\\it et al.}, 1997; Bommier, 1997 a,b). These recent efforts to generalize the density matrix theory to PRD are truly important and should be continued because the CRD theory cannot be applied when coherences between non-degenerate levels are present unless the spectrum of the radiation is flat across a frequency interval $\\Delta\\nu$ centred on the line frequency and larger than the frequency separation between the two levels connected by the coherence (see Landi Degl'Innocenti {\\it et al.}, 1997).  The problem of finding the self-consistent values of the irreducible tensor components of the atomic density matrix has been called ``the non-LTE problem of the $2^{nd}$ kind'' (see Landi Degl'Innocenti, 1987). It requires solving jointly the statistical equilibrium (SE) equations for the density matrix elements associated with each level of the assumed atomic model and the Stokes-vector transfer equations for all the radiative transitions involved. This terminology also seems appropriate because, as I shall try to argue below, a better understanding of many of the ``inexplicable'' spectral features of the second solar spectrum can only be achieved by carefully formulating and solving {\\it multilevel} non-LTE problems of the $2^{nd}$ kind, taking fully into account the possibility of {\\it atomic polarization} at {\\it all} the levels of the chosen multilevel atomic model and including the depolarizing role of elastic collisions and magnetic fields.  The word ``inexplicable'' in the preceding paragraphs refers to the impossibility of explaining some of the spectral features of the second solar spectrum by means of theories based on the approximation of neglecting atomic polarization in the {\\it lower} level of the line transition under consideration, i.e. theories based on the assumption that there are no significant differences in the populations of the lower-level sublevels or coherences among them. One example of a mysterious feature of the second solar spectrum that has triggered some recent theoretical work (see Trujillo Bueno and Landi Degl'Innocenti, 1997; Landi Degl'Innocenti, 1998) is the linear polarization pattern observed around the Na I ${\\rm D}_2$ and ${\\rm D}_1$ lines. In fact, in a recent letter in {\\it Nature} that demonstrates the robustness of the density matrix polarization transfer theory including partial frequency redistribution effects, Landi Degl'Innocenti (1998) concluded that the observed Na ${\\rm D}_2$ and ${\\rm D}_1$ linear polarization pattern can be explained by assuming the presence of an amount of {\\it ground}-level atomic polarization as important as (and indeed slightly larger than!) that of the upper level. However, because very small {\\it non-vertical} magnetic fields (and/or elastic collisions) destroy the ground-level atomic polarization of Na, Landi Degl'Innocenti (1998) was forced to rule out in the solar chromosphere both elastic collisions and the existence of turbulent magnetic fields and of horizontal, canopy-like fields stronger than $\\sim0.01$ gauss. This has led to an exciting apparent paradox in solar physics because there are observational evidence for both turbulent fields of the order of 10 gauss and canopy-like horizontal fields (see Jones, 1984; Solanki and Steiner, 1990; Faurobert-Scholl, 1992; Bianda {\\it et. al}, 1998; Stenflo {\\it et. al}, 1998).  The argument in favour of the simplifying approximation of neglecting lower-level polarization is that the lower level of a line transition is generally long-lived, and that it must thus have plenty of time to be depolarized by elastic collisions and/or weak magnetic fields (Stenflo, 1994; 1997). However, this is expected to be the case concerning only the {\\it ground} and {\\it metastable} levels of atomic systems. It is indeed a major simplifying approximation because it implies that the scattered radiation can be expressed linearly in terms of the incident radiation if stimulated emission processes are also neglected (see Landi Degl'Innocenti, 1984). In other words, the approximation of neglecting lower-level atomic polarization allows the study of scattering line polarization problems in terms of phase matrices that are decoupled from the SE equations. However, as we shall show below, the consideration of lower-level atomic polarization leads to a coupled system of {\\it non-linear} equations, even for the simplest case of a two-level model atom.  My motivation for writing this paper is two-fold. Firstly, I would like to provide more arguments concerning the idea (already put forward in the paper by Trujillo Bueno and Landi Degl'Innocenti, 1997) that lower-level atomic polarization is an essential physical ingredient for understanding the second solar spectrum. Secondly, I aim to demonstrate that, contrary to some general beliefs, the density-matrix polarization transfer theory does have a suitable form for practical applications.  To these ends, I will consider here three types of line transitions in the solar atmosphere: ($a$) lines with ${\\rm J}_{l}=0$ and ${\\rm J}_u=1$, ($b$) lines with ${\\rm J}_{l}=1$ and ${\\rm J}_u=0$, and ($c$) lines with ${\\rm J}_{l}=1$ and ${\\rm J}_u=1$. Section 2 summarizes the predictions of the standard theory, which neglects lower-level atomic polarization. Section 3 is dedicated to outlining the formulation of the scattering line polarization problem taking into account atomic polarization in both the upper and the lower levels of such line transitions. Here I will show the self-consistent solution for the density matrix elements and the corresponding emergent fractional linear polarization that results from this more correct treatment. Section 4 discusses model calculations including the depolarizing role of elastic collisions.  As we shall see, my two-level atom scattering line polarization calculations suggest that, if we aim at understanding the second solar spectrum, we need to consider scattering processes in {\\it multilevel} atomic models, taking fully into account the transfer of atomic polarization among all the atomic levels involved. Section 5 argues that, in principle, the above-mentioned ``Na solar paradox'' might admit a {\\it multilevel} solution. However, Section 6 shows that the observed fractional linear polarization in the Mg $b$-lines can only be explained by invoking the presence of atomic polarization in the lower {\\it metastable} levels of the Mg $b_1$ and $b_2$ lines, thus establishing a new paradox: the ``Mg solar paradox''. Finally, Section 7 gives some concluding remarks after pointing out that there is no multilevel solution to this ``Mg solar paradox''. The Appendix is dedicated to a brief description of some iterative methods for the solution of non-LTE problems of the $2^{nd}$ kind, which I consider as ``the road to be taken'' for the development of a general multilevel scattering line polarization code. ",
        "conclusions": "There is a crucial difference between the Na and Mg ``solar paradoxes''. For Na there was still the chance that the multilevel scenario outlined above might help to solve it. However, there is no similar multilevel solution for the ``Mg solar paradox'', which leads to the conclusion that there must indeed exist {\\it ground}-level and {\\it metastable}-level atomic polarization in the solar chromosphere. The three Mg $b$ lines share the same upper level and what cannot be understood by means of the standard transfer theory (which neglects lower-level polarization) is that the observed linear polarization in these three Mg $b$ lines turns out to be similar. Thus, even if one chooses multilevel atomic models for Mg to calculate the self-consistent values of the density-matrix elements, and then solves the Stokes transfer equations to get the emergent fractional linear polarization, but neglecting the dichroism contribution that comes from the atomic alignment of the lower {\\it metastable} levels, one would find again that the ensuing prediction is wrong. The only way I see for increasing the emitted polarization in the Mg $b_1$ and $b_2$ lines, so as to bring it to the same level of that corresponding to the Mg $b_4$ line (that has ${\\rm J}_l=0$), is via the {\\it dichroism} contribution (i.e. the term $-\\eta_{\\rm Q}{\\rm I}$ of Eq. 8). As discussed in Section 3, this {\\it dichroism} can only arise if the magnetic sublevels of the ${\\rm J}_l=2$ and ${\\rm J}_l=1$ {\\it metastable} lower-levels of the $b_1$ and $b_2$ lines are {\\it unequally populated}. The same explanation can be given for other groups of lines belonging to other atoms, and arising from a similar multiplet ($^3{\\rm P}^{\\rm o}\\,-\\,^3{\\rm S}$), with the ensuing appearance of extra ``paradoxes'' for other chemical elements.  As we have seen, the anisotropic illumination of the atoms in a stellar atmosphere can lead to large population imbalances among the lower-level sublevels of many spectral lines. The modelling of the second solar spectrum requires the reliable calculation of the atomic polarization of the lower and upper levels corresponding to the line transitions of interest. To this end, it is crucial to be able to consider {\\it multilevel} atomic models. This goal can presently be achieved by formulating the problems of interest within the framework of the density matrix polarization transfer theory (see Landi Degl'Innocenti, 1983) and by numerically solving the ensuing non-linear and non-local equations with the iterative methods presented in the Appendix. Only when we know the {\\it self-consistent} values of the alignment coefficients of the Na and Mg atomic levels in several solar atmospheric models shall we be able to figure out a possible solution to such ``solar paradoxes''.  Obviously, we are facing a complex problem here, from the observational, theoretical and modelling viewpoints. But it is a highly interesting one, not only because of the fascinating physics that it involves, but mainly because in trying to clarify it we may learn something new about the sun."
    },
    "0710/0710.0668_arXiv.txt": {
        "abstract": "Molecular gas has been searched for and found in unexpectedly large quantities in some collisional debris of interacting galaxies: HI-rich tidal tails, bridges and collisional rings. It was so far observed through  millimeter  observations of the CO line  and detected towards or near regions of star-formation associated to dense condensations of the atomic hydrogen. The discovery of cool H$_{2}$ at distances greater than 50 kpc from the parent (colliding) galaxies, whereas the external disk of spirals is generally considered to be CO--poor,  raised question on its origin and favored the hypothesis of a  local production out of collapsed HI clouds. However recent  observations of a diffuse CO component along tidal debris have challenged this idea. Another recent puzzle is the measurement in the collisional debris of  two interacting systems and four recycled objects of a missing mass, whereas no dark matter is expected there.  One debated interpretation is that this unseen component is cold, ``invisible\"  molecular gas initially present in the disk of spirals. ",
        "introduction": "The collisional debris addressed here  refer to all the material that is expelled into the intergalactic medium during galaxy--galaxy interactions. This is the result of either the tidal forces that shape bridges and tails, of direct high-speed impacts at the origin of rings or head-on collisions forming  systems similar to the so--called ``Taffy\" galaxies (see Struck 1999 for a review). Collisional debris may consist of old and young stars,  gas clouds and dust with a relative proportion that depends very much on the type of interactions and properties of the parent galaxies. Tidal tails with a prominent old stellar population will be common in slow encounters between late-type galaxies -- see the famous Antennae galaxies. Purely stellar tails may exist, when the gaseous counterpart  has for some reason been displaced and is now offset (Mihos 2001). More commonly, the gaseous tail extends further than the stellar one, because it was originally more extended in the disk of its parent galaxies. Collisional rings are particularly gas-rich, being embedded in a faint stellar halo. Finally, long, purely gaseous, tidal tails may form during high-speed collisions, as recently shown by Duc \\& Bournaud (2007). Finally, if the gas in the collisional debris is dense enough, it collapses and locally  form a new generation of stars.  Whereas stellar debris have been known for a long time -- they are the peculiarities  in the optical catalog of Peculiar Galaxies by Arp (1966) --,  their gaseous, HI, counterparts have been systematically studied from less than twenty years (e.g., Hibbard \\& van Gorkom  1996). This was possible thanks to interferometers like the Very Large Array which offer a very large field of view. In situ star formation in collisional debris has been discovered through several ways, including broad-band optical (e.g. Schombert et al. 1990) and UV (e.g., Neff et al. 2005; Boquien et al., 2007) imaging, narrow-band H$\\alpha$ images (e.g. Iglesias-P{\\'a}ramo  \\& V{\\'{\\i}}lchez  2001) as well as mid-infrared photometry.  If ongoing star-formation is present in collisional debris, molecular gas should be present as well. This paper reviews the various attempts that have been done so far to detect it. ",
        "conclusions": "Our detailed, multi-wavelength,  studies of collisional debris, especially its molecular gas content,  raise a number of questions. \\begin{itemize}   \\item Where does the cool H$_{2}$ detected in quantities in tidal tails come from? Locally, out of tidally expelled, dense, HI clouds, as suggested by the kinematical data, and/or from the disk of the parent galaxy, as indicated by the detection of more diffuse CO emission outside the main centers of star-formation? The answer to this question will require to carry-out complete maps of the molecular gas along collisional debris. So far most of the CO detections result from pointed observations, biased towards regions of high HI column density.  \\item How does the star-formation efficiency (the ratio of the star formation rate to the molecular mass) in collisional debris compare with that of individual star-forming regions of spiral disks and classical dwarf galaxies? A first rough comparison is presented in Braine et al. (2001) but should now be updated. The  SFE  should in particular be  determined in an homogeneous way, a difficult task for objects of different sizes and environments.  \\item Is really the missing mass found in collisional debris and hence in spiral disks very cold H$_{2}$? First of all, the  method used to determine the dynamical mass could be challenged. The objects formed in the debris are, spatially, barely resolved and  their inclination is not directly  determined;  however the uncertainties of the method have been well assessed by numerical simulations. Besides, the conversion factor used to derive the mass of luminous  cool  H$_{2}$  from   the CO line intensity is not very constrained; in the relative high-Z environment of TDGs, it should however be not so different than the assumed galactic value. If real, the mismatch between the dynamical and the luminous mass, estimated to a factor of 2--3 in now 4 different objets,  can also be interpreted in a totally different manner. Cold streams of non-baryonic dark matter in disks is a theoretical possibility. Gentile et al. (2007) and Milgrom (2007)  were able to reproduce the observed velocity curves of the recycled objects around NGC~5291 within the MOND framework.   On the other hand, new  simulations including conventional gravity, usual cosmological Dark-Matter halos and the putative cold component of H$_{2}$  in spiral disks (which was not implemented in the numerical modes of Bournaud et al.  2007) seem to also give a good match with the data (Revaz et al., in preparation).  \\end{itemize}  Although being just ``debris\", tidal tails, rings and more in general all the material sent out of galaxies during their merging history, especially the molecular gas,  may tell about some fundamental aspects of astrophysics."
    },
    "0710/0710.1897_arXiv.txt": {
        "abstract": "Inspiral signals from binary compact objects (black holes and neutron stars) are primary targets of the ongoing searches by ground-based gravitational-wave interferometers (LIGO, Virgo, GEO-600 and TAMA-300). We present parameter-estimation simulations for inspirals of black-hole--neutron-star binaries using Markov-chain Monte-Carlo methods. For the first time, we have both estimated the parameters of a binary inspiral source with a spinning component and determined the accuracy of the parameter estimation, for simulated observations with ground-based gravitational-wave detectors. We demonstrate that we can obtain the distance, sky position, and binary orientation at a higher accuracy than previously suggested in the literature. For an observation of an inspiral with sufficient spin and two or three detectors we find an accuracy in the determination of the sky position of typically a few tens of square degrees. ",
        "introduction": "\\label{sec:intro}  Binary systems with compact objects --- neutron stars (NS) and black holes (BH) --- in the mass range $\\sim 1\\,\\Ms - 100\\,\\Ms$ are among the most likely sources of gravitational waves (GWs) for ground-based laser interferometers currently in operation (Cutler \\& Thorne, 2002): LIGO (Barish \\& Weiss 1999), Virgo (Arcese et al.\\ 2004), GEO-600 (Willke et al.\\ 2004) and TAMA-300 (Takahashi et al.\\ 2004). Merger-rate estimates are quite uncertain and for BH-NS binaries current detection-rate estimates reach as high as 0.1\\,yr$^{-1}$ (\\emph{e.g.} O'Shaughnessy et al.\\ 2008) for first-generation instruments. Upgrades to Enhanced LIGO/Virgo (2008--2009) and Advanced LIGO/Virgo (2011--2014) are expected to increase detection rates by factors of about $\\sim 8$ and $10^3$, respectively.  The measurement of astrophysical source properties holds major promise for improving our physical understanding and requires reliable methods for parameter estimation. This is a challenging problem because of the large number of parameters ($> 10$) and the presence of strong correlations among them, leading to a highly-structured parameter space. In the case of high mass ratio binaries (\\emph{e.g.} BH-NS), these issues are amplified for significant spin magnitudes and large spin misalignments (Apostolatos et al.\\ 1994; Grandcl{\\'e}ment et al.\\ 2003; Buonanno et al.\\ 2003). However, the presence of spins benefits parameter estimation through the signal modulations, although still presenting us with a considerable computational challenge. This has been highlighted in the context of LISA observations (see Vecchio 2004; Lang \\& Hughes 2006) but no study has been devoted so far to ground-based observations.   In this \\emph{Letter} we examine for the first time the potential for parameter estimation of spinning binary inspirals with ground-based interferometers, including twelve physical parameters. Earlier studies (\\emph{e.g.} Cutler \\& Flanagan 1994, Poisson \\& Will 1995, Van den Broeck \\& Sengupta 2007) have estimated the theoretical accuracy with which some of these parameters should be measured, without determining the parameters themselves. Also, (R\\\"{o}ver et al.\\ 2006, 2007) have explored parameter estimation for non-spinning binaries. We focus on BH-NS binaries where spin effects are strongest (Apostolatos et al.\\ 1994), while at the same time we are justified to ignore the NS spin. We employ a newly developed Markov-chain Monte-Carlo (MCMC) algorithm (van der Sluys et al.\\ 2008) applied on spinning inspiral signals injected into synthetic ground-based noise and we derive posterior probability-density functions (PDFs) of all twelve signal parameters. We show that although sky position is degenerate when using two detectors, we can still determine the mass and spin parameters to reasonable accuracy.  With three detectors, the sky position and binary orientation can be fully resolved. We show that our accuracies are good enough to associate an inspiral event with an electromagnetic detection, such as a short gamma-ray burst (\\emph{e.g.} Nakar 2007). ",
        "conclusions": "\\label{sec:concl}  We have explored for the first time the parameter estimation of all physical parameters --- including masses, spin, distance, sky location and binary orientation --- on ground-based gravitational-wave observations of binary inspirals with spinning compact objects. We show that for two detectors and sufficient spin ($a_\\mathrm{spin}\\geq0.5$) or for three detectors, the obtained accuracy in sky position, distance and time of coalescence is good enough to allow the identification of electromagnetic counterparts of compact-binary mergers, \\emph{e.g.} short gamma-ray bursts (Nakar 2007). A direct measurement of mass, spin, distance and orientation can be obtained from inspiral GWs, which is notoriously difficult for electromagnetic observations.  The analysis presented here is the first step of a more detailed study that we are currently carrying out, exploring a much larger parameter space, developing techniques to reduce the computational cost of these simulations, and testing the methods with actual LIGO data. The waveform model used here, though adequate for exploratory studies, is not sufficiently accurate for the analysis of real detections, and we are finalising the implementation of a more realistic waveform. Simulations with this improved waveform may also shed light on the degeneracy between mass and spin parameters discussed in Sect.\\,\\ref{sec:results}, and may improve the accuracy of our parameter estimation appreciably (\\emph{e.g.} Van den Broeck \\& Sengupta 2007). Finally, we intend to further develop our Bayesian approach into one of the standard tools that can be included in the analysis pipeline used for the processing of the `science data' collected by ground-based laser interferometers."
    },
    "0710/0710.3776_arXiv.txt": {
        "abstract": "% The opacity of a spiral disk due to dust absorption influences every measurement we make of it in the UV and optical. Two separate techniques directly measure the total absorption by dust in the disk: calibrated distant galaxy counts and overlapping galaxy pairs. The main results from both so far are a semi-transparent disk with more opaque arms, and a relation between surface brightness and disk opacity. In the Spitzer era, SED models of spiral disks add a new perspective on the role of dust in spiral disks. Combined with the overall opacity from galaxy counts, they yield a typical optical depth of the dusty ISM clouds: 0.4 that implies a size of $\\sim$ 60 pc. Work on galaxy counts is currently ongoing on the ACS fields of M51, M101 and M81. Occulting galaxies offer the possibility of probing the history of disk opacity from higher redshift pairs. Evolution in disk opacity could influence distance measurements (SN1a, Tully-Fisher relation). Here, we present first results from spectroscopically selected occulting pairs in the SDSS. The redshift range for this sample is limited, but does offer a first insight into disk opacity evolution as well as a reference for higher redshift measurements. ",
        "introduction": "The number of distant galaxies seen in an {\\em HST} image through the face-on foreground spiral is a direct indication of its opacity, after proper calibration using artificial galaxy counts \\citep[synthetic field method (SFM)][]{Gonzalez98,Holwerda05a}. We have obtained calibrated galaxy counts for a sample of 32 deep {\\em HST/WFPC2} fields. The main results from the disk opacity study are: (1) most of the disks are semitransparent \\citep[][Figure \\ref{f:ra}]{Holwerda05b}, (2) arms are more opaque \\citep{Holwerda05b}, (3) as are brighter sections of the disk \\citep{Holwerda05d}. We did not find a relation between HI profiles and disk opacity but this can be better explored on a single disk \\citep{Holwerda05c}. The optimal distance of the foreground disk is $\\sim$ 10 Mpc, the compromise between crowding effects and solid angle constraints \\citep{Gonzalez03,Holwerda05e}.   \\begin{SCfigure} \\centering \\includegraphics[width=0.5\\textwidth]{holwerda2.eps}% \\caption{The relation between the optical depth from the SED model [$\\tau_m$(SED)], and the observed optical depth from the number of distant galaxies [$\\tau$(SFM)]. Model C from \\protect\\cite{Natta84} is fit through these. Cloud optical depth is 0.4, and  more than a single cloud along the line-of-sight is needed, especially for optically thick disks. Figure from \\protect\\cite{Holwerda07a}.} \\label{f:tc} \\end{SCfigure}    There is an overlap of 12 galaxies with {\\em SINGS}; we compared the disk opacity from number of distant galaxies to the dust surface density obtained from an SED model \\citep{Draine07a}. Expressed as optical depths, the relation between the {\\it observed} optical depth ($\\tau$) and the optical depth inferred from the SED dust mass ($\\tau_m$) is solely a function of the typical cloud optical depth ($\\tau_c$): $\\tau/\\tau_m = (1-e^{-\\tau_c})/\\tau_c$ \\citep{Natta84}. Figure \\ref{f:tc} shows the values for $\\tau$ and $\\tau_m$ and the fit of $\\tau_c$=0.4. This implies that most of the disks are made up of small, optically thin, cold ISM structures with more than one cloud along the line-of-sight, especially in optically thick disks \\citep{Holwerda07a}. From their distribution in the E[I-L] color map, it becomes clear that the number of distant galaxies does not drop exclusively as a result of grand spiral opaque structures but that the unresolved dusty ISM disk is equally important \\citep{Holwerda07b}.    \\begin{figure} \\centering \\includegraphics[width=0.6\\textwidth]{holwerda3.ps}% \\caption{\\label{f:raz}Radial opacity profile for the local (lower), z=0-0.1 (middle) and z=0.1-0.2 (top) spirals from occulting pairs. Open circles are disk sections for the local pairs. Figure from \\protect\\cite{Holwerda07c}.} \\label{f:tt} \\end{figure} ",
        "conclusions": ""
    },
    "0710/0710.4360_arXiv.txt": {
        "abstract": "\\footnote{To appear in proceedings of the Puerto Vallarta Conference on ``New Quests in Stellar Astrophysics II: Ultraviolet Properties of Evolved Stellar Populations'' eds. M. Chavez, E. Bertone, D. Rosa-Gonzalez \\& L. H. Rodriguez-Merino, Springer, ASSP series.  Presented as part of a Ph. D. thesis in the Departamento de F\\'\\i sica, of the Universidad de Guadalajara, M\\'exico.}   Ultracompact (UC) HII regions with Extended Emission (EE) are classical UC~HII regions associated with much larger ($>$1$'$) structures of ionized gas. The efforts to investigate, detect, and understand if the EE is physically related to the UC emission are few.  If they are related, our understanding of UC~HII regions may be affected (e.g., in the estimation of ionizing UV photons). Here we present a brief overview of UC~HII regions with EE (UC~HII+EE) including our most recent effort aimed at searching for UC~HII regions associated with extended emission. ",
        "introduction": "\\label{sec:1}  Ultracompact (UC) HII regions are small (size $\\leq$ 0.1 pc), dense ($\\geq$ 10$^4$ cm$^{-3}$), photoionized hydrogen regions with high emission measure ($\\geq$ 10$^7$ {${\\rm pc\\ cm}^{-6}$}), surrounding recently formed ionizing OB type stars (e.g. Fig.  1a). These characteristics were observationally confirmed by \\cite{WC89} and \\cite{K94}, and more recent reviews are presented by \\cite{C02} and \\cite{LFR06}. The study of UC~HII regions began in 1967 via interferometric observations of compact HII regions (see \\cite{K02} for a summary). Because UC~HII regions are generally surrounded by a natal dust `cocoon', radio--continuum (RC) and infrared (IR) observations are needed to study them. In the RC, the first VLA surveys (e.g., \\cite{WC89}, \\cite{K94}) were made at 2 and 6 cm in configurations A and B, supplying arc--sec resolutions and sensitivities to structures up to 10--20$''$. The IR counterparts were mainly provided by IRAS, with resolutions of $\\sim$30$''$--2$'$.  \\begin{figure} \\centering \\includegraphics[height=9.5cm]{edelafuente_fig1.eps} \\caption{The UC~HII region with EE G60.88--0.13 (IRAS 19442+2427). a) and b) VLA images in configuration B (the UC emission) and configuration C (the EE). c) Combined VLA Multi--Resolution--Clean (MRC) map. This map strongly suggests a direct connection between the UC and the extended emission. d) All IRAC bands image at meso--scales (gray) and superimposed contours from c). Dust is predominant in the region. A spatial location agreement between dust and the HII region is observed.} \\label{fig:1}       % \\end{figure}  Although the presence of large scale structures related to UC~HII regions (Fig.  1b and 8c of \\cite{K99}) had been inferred since 1967, the first interferometric surveys did not detect them.  Nevertheless, their detection is possible with the VLA C and D configurations (e.g., \\cite{K99}), albeit at the expense of resolution towards the UC emission (UCE).  To mitigate these spatial filtering effects, we made multi--configuration VLA observations to provide a multi-scale view of the UC~HII+EE (Fig.1c). Also, by using MSX observations, of higher resolution than IRAS, it is possible both to detect the EE and to resolve the UCE, since the infrared observations are sensitive to the full range of angular sizes. The best IR satellite observations available for this purpose (see Fig. 1d) are from the Spitzer Space Telescope ({\\it http://ssc.spitzer.caltech.edu}). ",
        "conclusions": ""
    },
    "0710/0710.1773_arXiv.txt": {
        "abstract": "We investigate the influence of large-scale meridional circulation on solar p-modes by quasi-degenerate perturbation theory, as proposed by~\\inlinecite{lavely92}. As an input flow we use various models of stationary meridional circulation obeying the continuity equation. This flow perturbs the eigenmodes of an equilibrium model of the Sun. We derive the signatures of the meridional circulation in the frequency multiplets of solar p-modes. In most cases the meridional circulation leads to negative average frequency shifts of the multiplets. Further possible observable effects are briefly discussed. ",
        "introduction": "Meridional circulation is a large-scale flow observed on both hemispheres of the solar surface~\\cite{duvall79,hathaway96,komm93}. Its predominant direction is from the equator to the poles, and its amplitude is of the order of 15\\thinspace m/s. As mass does not accumulate in the polar regions a return flow from the poles to the equator is suspected deeper within the solar interior. The top half of the convection zone contains approximately 0.25\\% of the solar mass, the mass of the bottom half is approximately five times larger. Consequently, a poleward flow of 10\\thinspace m/s in the top half of the convection zone could be compensated by an equatorward flow of 2\\thinspace m/s in the lower half. The transport of magnetic flux from mid to low latitudes by such a flow at the bottom of the convection zone would last approximately 10\\thinspace years, which is close to the period of the solar magnetic cycle.  In addition to magnetic flux, the meridional flow also transports angular momentum. Indeed, the circulation played a key role in early theories of the solar non-uniform rotation as well as of the magnetic cycle~\\cite{bjerknes26,kippenhahn63}. More recently, differential rotation has been explained in mean-field models as a consequence of the Reynolds stresses~\\cite{ruediger80,kueker01,kueker05}, and in three-dimensional numerical models by the influence of the Coriolis force on global convection~\\cite{miesch00,miesch06}. Nevertheless, meridional circulation occurs as well in these models, and in solar-cycle models it has regained popularity, since the mean-field latitude migration along the surfaces of isorotation that occurs in the traditional $\\alpha\\Omega$ dynamo~\\cite{parker55} does not seem to suffice. The effect of the circulation on the butterfly diagram had been demonstrated by ~\\inlinecite{roberts72}; it is considered to be essential in more recent versions of the $\\alpha\\Omega$ dynamo, which therefore have been termed `flux-transport dynamos'~\\cite{choudhuri95,dikpati99,nandy02,rempel06a,rempel06b}.  Local helioseismology has investigated the strength of the meridional flow in the solar interior. By means of ring diagram analysis \\cite{hill88} \\inlinecite{haber02} inverted data for the circulation in a 15\\thinspace Mm deep region below the solar surface. In data from 1998 they found a flow emerging at high northern latitudes with equatorward orientation. This was interpreted as an evolving second cell of circulation. Further studies on the evolution of the flow either by ring-diagrams and by time-distance helioseismology (e.g.~\\opencite{zhao04},~\\opencite{zaatri06}) show predominantly a polward flow with a strong variability in the outer 15\\thinspace Mm of the Sun. The velocity reaches 40\\thinspace m/s. These findings were interpreted as the upper parts of meridional circulation cells in the two hemispheres.  Theoretically, the influence of global-scale stationary flows on solar p-modes was studied in detail by quasi-degenerate perturbation theory~\\cite{lavely92}. These studies were successfully used to solve the forward problem for the influence of differential rotation on the p-modes. The results lead to an improved inversion method for determining the radial dependence of the differential rotation~\\cite{ritzwoller91}. Following~\\inlinecite{lavely92}, \\inlinecite{roth99} studied the influence of large-scale sectoral poloidal flow components that could be related with giant convection cells. They found that such flows yield additional frequency shifts that can only be described with quasi-degenerate perturbation theory, as these frequency shifts are effects of higher order. In a sequel study~\\inlinecite{roth02} were able to show that sectoral poloidal flows could only be found with the current inversion methods of global helioseismology as long as they exceed an amplitude of 10\\thinspace m/s. The meridional flow was found to be not detectable by the current inversion methods as the frequency splittings are fitted by an incomplete set of basis functions, which are tailored to measure only zonal toroidal flows. However, no detailed study on the effect of the meridional circulation on the oscillation frequencies was given, and few other attempts to derive observable signatures of the meridional flow in global helioseismology data exist~\\cite{woodard00}.  In this contribution we concentrate on a theoretical study of the influence of the meridional circulation on the solar p-mode frequencies and describe a possible observable effect. This effect is significantly smaller than the frequency splitting caused by solar differential rotation. But as time series of global oscillation data exist that cover more than 10 years, the necessary frequency resolution might be available to detect it. The advantage of studying the meridional circulation by global helioseismic techniques is a possible inference of information from greater depths. ",
        "conclusions": "In this paper we used simple models of the meridional circulation to investigate their influences on the solar p-mode frequencies. The simplest model consisted of one cell per hemisphere and depth with a maximum horizontal flow velocity of 15\\thinspace m/s on the solar surface. The most complicated model we used had three cells per hemisphere and three revolution cells in depth with a horizontal flow velocity of 15\\thinspace m/s at the surface, too. We performed numerical calculations based on quasi-degenerate perturbation theory to obtain the frequency splittings of the solar p-modes due to this meridional circulation models. We find that the meridional circulation lifts the degeneracies of the multiplets. For the simplest model the shifts are on average only 0.01\\thinspace $\\mu$Hz with a few shifts up to several $\\mu$Hz. Models with more cells per hemisphere affect the p-modes stronger. For the model with three cells per hemisphere we find an average shift of 0.1\\thinspace $\\mu$Hz, with many shifts in the order of 1\\thinspace $\\mu$Hz. In most cases the shifts are negative due to the fact that the next neigboring mode in frequency dominates the shifting; and due to structure of the $l$-$\\nu$ diagram this nearest neighbor has usually a higher frequency causing therefore negative shifts.  Comparing models with the same number of cells in latitude but different number of cells in depth we find only tiny differences in the resulting frequency splittings. In the case of $s=6$ the difference in the frequency splitting is on average 0.03\\thinspace $\\mu$Hz. However a number of differences in the order of 1\\thinspace $\\mu$Hz can occur.  On the Sun the meridional circulation consists probably of a superposition of flow components with different numbers of cells in depth and latitude. In contrast to the models used in our work, the amplitudes of these flow are in reality likely to be very different and highly variable in time. Nevertheless, based on our results we are optimistic that the meridional circulation on the Sun could leave an observable signature in the p-mode frequencies. In order to detect the effect a frequency resolution of at least 0.1\\thinspace $\\mu$Hz must be achieved. This could be done by a several month long time series. In order to be able to distinguish not only the various components of the meridional circulation in the orders $s$ but also in the radial orders $n_c$, several years of data need to be averaged to obtain the required precision in the frequencies. As such long data sets exist it will be worth while to search for this effect in the p-mode frequencies.  Compared to the splitting caused by the differential rotation the effect of meridional circulation is small. But as the effect of the rotational splitting is odd it can be separated from the even meridional splitting. However, e.g. asphericities and the magnetic field might cause symmetrical frequency shifts, too. Therefore, a lot of forward modelling will be necessary before the effect of the meridional circulation can be disentangled from these other effects. In this sense, one possible future extension of our work is the determination of the frequency shifts due to more sophisticated models of the meridional circulation, e.g. from three-dimensional numerical models. This might allow tailoring inversion routines for estimating the meridional circulation in the Sun from global solar oscillation frequencies."
    },
    "0710/0710.3406_arXiv.txt": {
        "abstract": "{Observations of collimated outflows in young stellar objects indicate that several features of the jets can be understood by adopting the picture of a two-component outflow, wherein a central stellar component around the jet axis is surrounded by an extended disk-wind. The precise contribution of each component may depend on the intrinsic physical properties of the YSO-disk system as well as its evolutionary stage. }{ In this context, the present article starts a systematic investigation of two-component jet models via time-dependent simulations of two prototypical and complementary analytical solutions, each closely related to the properties of stellar-outflows and disk-winds. These models describe a meridionally and a radially self-similar exact solution of the steady-state, ideal hydromagnetic equations, respectively. }{ By using the PLUTO code to carry out the simulations, the study focuses on the topological stability of each of the two analytical solutions, which are successfully extended to all space by removing their singularities. In addition, their behavior and robustness over several physical and numerical modifications is extensively examined. Therefore, this work serves as the starting point for the analysis of the two-component jet simulations. }{ It is found that radially self-similar solutions (disk-winds) always reach a final steady-state while maintaining all their well-defined properties. The different ways to replace the singular part of the solution around the symmetry axis, being a first approximation towards a two-component outflow, lead to the appearance of a shock at the super-fast domain corresponding to the fast magnetosonic separatrix surface. These conclusions hold true independently of the numerical modifications and/or evolutionary constraints that the models have been undergone, such as starting with a sub-modified-fast initial solution or different types of heating/cooling assumptions. Furthermore, the final outcome of the simulations remains close enough to the initial analytical configurations showing thus, their topological stability. Conversely, the asymptotic configuration and the stability of meridionally self-similar models (stellar-winds) is related to the heating processes at the base of the wind. If the heating is modified by assuming a polytropic relation between density and pressure, a turbulent evolution is found. On the other hand, adiabatic conditions lead to the replacement of the outflow by an almost static atmosphere.}{} ",
        "introduction": "\\label{sec:intro}  Observations made over the last two decades have shown that one class of the widespread astrophysical phenomenon of collimated plasma outflows (jets) is being launched from the vicinity of most young stellar objects (YSOs) (Burrows et al. \\cite{Bur96}). These supersonic mass outflows are found to be correlated with accretion (Cabrit et al. \\cite{Cab90}; Hartigan et al. \\cite{Har95}), to have narrow opening angles (Ray et al. \\cite{Ray96}) and to propagate for several orders of magnitude of spatial distances ranging from the AU to the pc scales (Dougados et al. \\cite{Dou00}; Hartigan et al. \\cite{Har04}). A central role in the launching, acceleration and collimation of these jets is widely believed to be played by magnetohydrodynamic (MHD) effects, which can also successfully remove the excessive angular momentum, allowing in this way the YSO to accrete and enter the main sequence. Nevertheless, although recent high angular resolution observations put several constraints on the different driving mechanisms proposed, it is not yet clear which is the dominant plasma launching mechanism in YSO jets.  The system of a protostellar object basically contains two dynamical constituents, a central protostar and its surrounding accretion disk. Consequently, in Bogovalov \\& Tsinganos (\\cite{Bog01}) it is argued that jets observed from T Tauri stars most likely consist of two main steady components: (i) an inner pressure driven wind, which is non-collimated if the star is an inefficient magnetic rotator and (ii) an outer magneto-centrifugally driven disk-wind which provides most of the high mass loss rate observed. The relatively faster rotating magnetized disk produces the self-collimated wind which then forces all enclosed outflow from the central source to be collimated as well. This conclusion is confirmed by self-consistent simulations of the MHD equations. More recently, in Ferreira et al. (\\cite{Fer06}) it is argued that for the YSO jets observed in association with T Tauri stars, in addition to the pressure driven stellar outflow and the magneto-centrifugally launched extended ``warm'' disk-wind, a third component may be driven by magnetic processes at the magnetosphere/disk interaction, i.e., a sporadically ejected X-type wind. In addition, in Ferreira et al. (\\cite{Fer00}) a non steady ``two-flow'' scenario was also suggested, regarding a reconnection X- and a disk-wind. Nevertheless, the existence of such sporadic components is not supported by observational data as being one of the major contributors to the steady characteristics of jets, but rather could explain the observed variability in jet emission. On the other hand and within the same framework, recent observations (Edwards et al. \\cite{Edw06}; Kwan et al. \\cite{Kwa07}) show that both, disk and stellar winds are mainly present in T Tauri stars, with the dominant component being determined by the intrinsic physical properties of the particular YSO.  From the analytical studies point of view, the complexity of the launching and collimation mechanisms of jets have forced researchers over the past several years to treat these two components separately. The only available analytical MHD models for jets are those characterized by the symmetries of radial and meridional self-similarity (Vlahakis \\& Tsinganos \\cite{Vla98}). In the former case, the solution is invariant as we look at a constant polar angle and in the latter, as we look at a constant spherical radius. The computational consequence of the respective symmetry is that by employing the separable spherical coordinates ($r$, $\\theta$), the set of coupled MHD equations reduces to a set of ordinary differential equations in $\\theta$, or, in $r$, respectively. The last remaining difficulty is to select solutions which are causally disconnected from the source of the outflow, i.e., those crossing the fast magnetosonic separatrix. In this way, one may construct either radially self-similar solutions closely related to the properties of magneto-centrifugally driven disk-winds (Blandford \\& Payne \\cite{Bla82}; Contopoulos \\& Lovelace \\cite{Con94}; Ferreira \\cite{Fer97}; Vlahakis et al. \\cite{Vla00}, hereafter VTST00), or  meridionally self-similar ones to address pressure driven stellar outflows (Sauty \\& Tsinganos \\cite{Sau94}; Trussoni et al. \\cite{Tru97}; Sauty et al. \\cite{Sau02}, hereafter STT02).  Since each self-similar symmetry corresponds to a particular component, we adopt the following initials: ADO (Analytical Disk Outflow) and ASO (Analytical Stellar Outflow) to refer to radially and meridionally self-similar solutions, respectively.  Apart from the geometry, an intrinsic distinction between these two classes concerns the treatment of the energy equation. Nevertheless, the symmetry difference makes them complementary to each other, since the ADO solution becomes singular at the axis, whereas the ASO is by definition the proper one for modeling the area close to it. In addition, the properties of the launching region of the disk-wind, i.e., at large polar angles, are described more naturally by the ADO model. For a recent review on the analytical work on MHD outflows the reader is referred to Tsinganos (\\cite{Tsi07}).  On the other hand, the increase of computational power along with the development of sophisticated numerical codes have allowed to study the time evolution of the MHD equations, giving us new perspectives of the physics involved. Jet launching and collimation have been mainly investigated with the following two methodologies: (i) by treating the disk as a boundary (Krasnopolsky et al. \\cite{Kra99}; Ouyed et al. \\cite{Ouy03}; Fendt \\cite{Fen06}) and (ii) by including the disk inside the computational box (Casse \\& Keppens \\cite{Cas04}; Meliani et al. \\cite{Mel06}; Zanni et al. \\cite{Zan07}). The former case allows a wider range of physical processes and mechanisms to be studied, whereas the latter, has the advantage of the jet evolution being consistent with that of the disk. However, with the exception of Gracia et al. (\\cite{Gra06}; hereafter GVT06), most numerical studies did not take advantage of the availability of the well studied analytical solutions which also allow a parametric study and therefore a better physical understanding of the problem of jet launching and collimation.  The present work is the first attempt to numerically construct and study a two-component jet, by using as starting point the two well studied classes of analytical self-similar solutions. Towards this goal, in this paper we first address the question of the topological stability of each one of these two classes separately, before we combine them in the following paper.  Concerning the ADO model, we shall use the VTST00 analytical solution, completing and considerably extending the GVT06 analysis. Therein, they found that the disk-wind model may attain a new steady-state configuration close enough to the initial analytical one, provided the assumption of some appropriate approximations around the axis. We further present the first numerical studies of ASO (meridionally self-similar winds), referring to the solutions of STT02, that are essential to model the region around the axis, just where the ADO model fails.  Once the physical properties and stability of these two classes of solutions is clarified, our final aim will be to effectively build up a model that consistently merges the ASO and ADO solutions. Such simulations will be presented in a future work where we will study the launching and propagation of a collimated stellar wind around the system axis surrounded by a disk wind.  Finally, a word on the term topological stability used in this paper. Classical stability theory addresses the question whether a given equilibrium configuration evolves away from (=unstable) or back to (=stable) the initial equilibrium when perturbed. In the present context, topological (or structural) stability refers to the question whether a given configuration preserves its topological properties when subject to various perturbations. Needless to say, that topologically unstable configurations may well be stable from the classical point of view and vice versa.  The paper is structured as follows: In section \\S\\ref{sec:anmod} the formalism of the ADO and ASO solutions is briefly presented. In section \\S\\ref{numst} their implementation is explained and the numerical models to be investigated are presented. Section \\S\\ref{sec:res} reports the results obtained by carrying out the respective simulations. Finally in section \\S\\ref{sec:disc} we discuss our results in the framework of the future matching and we report the conclusions of this work. ",
        "conclusions": "\\label{sec:disc}  In this paper we have studied several physical and numerical aspects concerning two classes of the self-similar models, each associated with a disk- and a stellar-wind, in the framework of the upcoming work that combines them to describe a two-component outflow. These analytical solutions (ADO and ASO) were appropriately modified, implemented as initial conditions and evolved in time. Our main conclusions are the following:  \\begin{itemize} \\item The {\\it Analytical Disk Outflow (radially self-similar) solution} has been successfully validated for its stability and robustness against several physical and numerical issues. This argument holds true even though the analytical solution was in many cases significantly modified. We have constructed numerical models and carried out simulations a) by assuming the extreme cases of an isothermal and an adiabatic evolution,  b) by treating the diverging behavior of the solution at the axis with different kinds of extrapolation schemes, mimicking a stellar wind component and c) by changing the size, resolution and geometry of the computational box.  In all cases, the poloidal critical surfaces, with the exception of the FMSS, were not readjusted, but rather matched perfectly to their initial position. The numerical solution always maintained the property of the successful crossing of all three critical surfaces producing an outflow causally disconnected from the base. This is achieved with the formation of a shock, corresponding to the numerically readjusted FMSS. In particular, this shock acts as a ``wall'' protecting the sub-modified-fast magnetosonic regions (source regions of the disk wind) from any perturbations taking place due to the modification of the models close to the axis (e.g. an effective stellar wind). However, the numerically readjusted FMSS (shock) does not coincide with the analytical one, with this departure being dictated by the respective numerical modifications of the models under consideration. A highly significant result is the fact that such a conclusion holds true even if we initialize the simulation with a sub-modified-fast solution, i.e. a solution with its whole domain causally connected. We found that, during the simulation, such a numerical model self-adapts to produce a shock (corresponding to the FMSS), hence no information of the downstream region can travel back to affect the launching region. This implies that even MHD outflow solutions, that do not successfully cross all three critical points, will probably converge to ``astrophysically correct'' solutions once evolved in time (see also Ferreira \\cite{Fer97}).  On the other hand, the study of GVT06 was successfully extended down to the equator with the help of simulations using spherical coordinates. Furthermore, by adopting different assumptions for the energy source terms, it was shown that the solution is only slightly and accordingly self-modified maintaining all its well defined properties. This is in agreement with the fact that the ADO solution describes essentially a magneto-centrifugally accelerated outflow.  \\item The {\\it Analytical Stellar Outflow (meridionally self-similar) solution}, which was validated in time-dependent simulations for the first time, maintained its well-defined equilibrium as expected. Such conclusion is supported by simulations performed with both the super- and sub-Alfv\\'enic regions included. Quite critical are, contrary to disk winds, the effects of the energetics in such thermally driven models. Although different assumptions of the energy equation in the super-Alfv\\'enic domain did not yield any significant modification of the analytical solution, strong variations of the structure of the axial outflows are found if modifications of the heating/cooling mechanisms occur in the initial accelerating region. In particular a polytropic assumption, mimicking isothermal conditions, would produce a turbulent weak outflow, while an adiabatic evolution asymptotically reaches a static atmosphere. We are tempted to relate the heating intermittency and even a switching off in such an ASO solution with the observed variability of accretion-driven YSO outflows.  \\item All previous statements hold true while being in perfect agreement with physically consistent requirements, such as specifying the correct type of boundary conditions: a) according to the propagation direction of the MHD waves, b) the axisymmetry holding around the axis and c) the constancy of certain physical variables at both a conical surface close to the equatorial plane and a radial one close to the origin, implying the presence of an underlying disk and the stellar surface, respectively. Therefore, the results can safely be trusted since they are not subject to any artificial forcing.  \\item Last, but certainly not least, is the fact that almost all models of both classes reached a steady- or quasi-steady-state. In this context, the upcoming mixing of the two complementary classes of self-similar solutions in order to study a two-component jet is well founded and promising. The final numerical model will incorporate both proposed scenarios of a pressure-driven outflow (ASO) surrounded by an extended, magneto-centrifugally driven disk wind (ADO). This task is undertaken in the second paper of this series. \\end{itemize}"
    },
    "0710/0710.1918_arXiv.txt": {
        "abstract": "We examine production of Li on the surface of a low-mass secondary in a black hole soft X-ray transient (BHSXT) through the spallation of CNO nuclei by neutrons which are ejected from a hot (> 10 MeV) advection-dominated accretion flow (ADAF) around the black hole. Using updated binary parameters, cross sections of neutron-induced spallation reactions, and mass accretion rates in ADAF derived from the spectrum fitting of multi-wavelength observations of quiescent BHSXTs, we obtain the equilibrium abundances of Li by equating the production rate of Li and the mass transfer rate through accretion to the black hole. The resulting abundances are found to be in good agreement with the observed values in seven BHSXTs. We note that the abundances vary in a timescale longer than a few months in our model. Moreover, the isotopic ratio \\nuc{Li}{6}/\\nuc{Li}{7} is calculated to be about 0.7--0.8 on the secondaries, which is much higher than the ratio measured in meteorites. Detection of such a high value is favorable to the production of Li via spallation and the existence of a hot accretion flow, rather than an accretion disk corona system in quiescent BHSXT. ",
        "introduction": "High abundances of Li have been detected in late-type secondaries of black hole soft X-ray transients (BHSXTs) and a neutron star soft X-ray transient (NSSXT) in quiescence~\\citep{martin92,martin94,martin96}, though Li would be destructed in a deep convective envelope of a late-type star. The Li enrichment has not, however, been observed on a late-type secondary in a compact binary with a white dwarf~\\citep{martin95}. These facts strongly suggest that a production mechanism of Li operates in compact binaries~\\citep{yn97,gk99} and that the nature of the primaries is crucial for the mechanism, though rotation might reduce the destruction of Li in the envelope of the secondary~\\citep{mjs05}.   Multi-wavelength spectra of BHSXTs in quiescence are successfully fitted to the radiation from an advection-dominated accretion flow (ADAF) around the black hole~\\citep{nmy96,nbm97}. Density is so low in ADAF, that ions interact inefficiently with electrons. Consequently ions have high temperatures due to viscous heating up to about 30 MeV near the inner edge of ADAF. At such high temperatures,  $\\alpha$-$\\alpha$ reaction proceeds to synthesize Li inside ADAF~\\citep{martin94,yn97}. It is necessary that a fraction $10^{-3}-10^{-4}$ of the accreting gas is transported to the secondary to explain the high abundances of Li observed in BHSXTs. However, such a high fraction is uncertain to be realized due to strong gravity of the black hole and the Coulomb interactions with nuclei inside ADAF~\\citep{gk99}.  Helium breaks via spallation with protons to produce neutrons at the inner region of ADAF. A large fraction of neutrons can be ejected from ADAF, because they do not interact with nuclei through the Coulomb interactions. Neutrons intercepted by the secondary interact with CNO nuclei through spallation to produce Li on the surface~\\citep{gk99}. This scenario is of particular interest, because the Li enrichment is anticipated in secondaries only for BHSXTs and NSSXTs, but for white dwarfs as primaries where ADAF cannot attain enough high temperatures to break helium to nucleons there.  In the present paper, we evaluate the Li abundances on the surface of secondaries in BHSXTs, following the scenario proposed by \\citet{gk99}. To this end, we use updated binary parameters, such as the mass $M$ of a black hole, the mass $M_*$ and radius $R_*$ of a secondary, mass accretion rates derived from the spectrum fitting of multi-wavelength observation of BHSXTs in quiescence, and cross sections of neutron-induced spallation reactions. Then, we compare the resulting abundances with the observed values, and show that the agreement is quite well. Moreover, we predict the isotopic ratio \\nuc{Li}{6}/\\nuc{Li}{7} on the secondaries in BHSXTs. ",
        "conclusions": ""
    },
    "0710/0710.4226_arXiv.txt": {
        "abstract": "We report on simultaneous optical and X-ray observations of the Seyfert galaxy, NGC~3147. The XMM-\\textit{Newton} spectrum shows that the source is unabsorbed in the X-rays ($N_H<5\\times10^{20}$ cm$^{-2}$). On the other hand, no broad lines are present in the optical spectrum. The origin of this optical/X-rays misclassification (with respect to the Unification Model) cannot be attributed to variability, since the observations in the two bands are simultaneous. Moreover, a Compton-thick nature of the object can be rejected on the basis of the low equivalent width of the iron K$\\alpha$ line ($\\simeq130$ eV) and the large ratio between the 2-10 keV and the [O\\textsc{iii}] fluxes. It seems therefore inescapable to conclude that NGC~3147 intrinsically lacks the Broad Line Region (BLR), making it the first ``true'' Seyfert 2. ",
        "introduction": "The basic assumption of Unified Models of Active Galactic Nuclei (AGN) is that type 1 and type 2 objects are intrinsically the same, the apparent difference being solely due to orientation effects \\citep[e.g.][]{antonucci93}. The absorbing medium assumes the fundamental role in this scenario. It is usually envisaged as an optically thick `torus', embedding the nucleus and the Broad Line Region (BLR). If we observe the torus edge-on, all the nuclear radiation, as well as the broad optical lines coming from the BLR, is completely blocked and we classify the source as a type 2. The narrow lines are still visible, because the Narrow Line Region (NLR) is located farther away from the nucleus, beyond the torus.  On the other hand, if the torus does not intercept our line of sight, we observe every component of the spectrum and the object is classified as a type 1.  Simple extrapolations of the optical/UV scenario to the X-ray emission do not, however, always easily fit the observations, leading sometimes to different classifications between the two bands. A larger amount of absorbing material is usually measured from X-ray observations in Seyfert 2s with respect to Seyfert 1s, as expected \\citep[e.g.][]{awaki91,ris99}. However, a number of ``true'' Seyfert 2 candidates have been found, i.e. objects with no absorption in the X-rays and no broad optical lines \\citep[e.g.][]{pappa01,pb02,wol05}, or with a large intrinsic Balmer decrement of the BLR \\citep{bcc03,corr05}. The so-called `naked' AGN may be similar objects, being characterised by the absence of broad lines, but strong variability in the optical band \\citep{hawk04} and, apparently, no absorption in the X-rays \\citep{gliozzi07}. These findings represent a challenge to the Unified Models and may require new classes of objects with intrinsic differences, like the absence of BLR.  Nevertheless, these sources may be highly variable and may change their optical and/or X-ray appearance in different observations. These `changing-look' AGN are not uncommon. In some cases, this behaviour is best explained by a real `switching-off' of the nucleus \\citep[see e.g.][]{mgm03,gua05}, in others by a variable column density of the absorber \\citep[e.g.][]{elvis04,ris05}. If the optical and the X-ray spectrum are taken in two different states of the source, it is clear that the disagreement between the two classifications may be only apparent. Therefore, the key to find genuine `unabsorbed Seyfert 2s' is represented by simultaneous X-ray and optical observations.  NGC~3147 (z=0.00941) belongs to the Palomar optical spectroscopic survey of nearby galaxies \\citep{hfs95}. Its classification as a Seyfert 2 is based on both the relative strength of the low-ionization optical forbidden lines with respect to the hydrogen Balmer lines and the ratio of [O\\textsc{iii}] to H$\\beta$, together with the lack of broad permitted lines \\citep{hfs97}. \\textit{ASCA} provided the first X-ray spectrum, which appeared Seyfert 1-like, without significant absorption and a standard powerlaw index \\citep{ptak96}. The lack of obscuration was later confirmed by \\textit{BeppoSAX} \\citep{dad07} and \\textit{Chandra}, which also showed that no off-nuclear source can significantly contribute to the nuclear emission \\citep{tw03}. In order to reconcile the X-ray data with the optical classification, \\citet{ptak96} suggested that NGC~3147 was a Compton-thick source, given also the relatively large equivalent width (EW) of the iron line. However, this hypothesis was rejected by the use of diagnostic diagrams based on $F_\\mathrm{X}/F_{\\mathrm{[OIII]}}$ and $F_\\mathrm{X}/F_{\\mathrm{IR}}$ ratios \\citep{pb02}. NGC~3147 is, therefore, a genuine candidate to be an unabsorbed Seyfert 2 galaxy. In this paper, we report on simultaneous XMM-\\textit{Newton} and optical observations of NCG~3147, in order to settle the issue.  In the following, errors correspond to the 90\\% confidence level for one interesting parameter ($\\Delta \\chi^2 =2.71$), where not otherwise stated. The adopted cosmological parameters are $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_\\Lambda=0.73$ and $\\Omega_m=0.27$ (i.e. the default ones in \\textsc {Xspec} 12.3.1). ",
        "conclusions": "The best fit model for the XMM-\\textit{Newton} spectra is that of a typical type 1 AGN. The amount of neutral column density measured in excess of the Galactic one is of the same order of the latter, and it is therefore fully consistent with being due to the host galaxy's interstellar medium. The powerlaw index, although somewhat flatter than the average type 1 radio-quiet object, is well within the large dispersion of the hard X-ray $\\Gamma$ distribution \\citep[e.g.][]{pico05}. The EW of the neutral iron K$\\alpha$ line is fully consistent with those typically found in unobscured AGN, especially given its low luminosity \\citep[e.g.][]{bianchi07}. The only peculiar aspect of its X-ray spectrum is represented by the Fe \\textsc{xxv} emission line, whose EW is larger than the one typically arising in a Compton-thin, photoionised gas \\citep[see e.g.][]{bm02,bianchi05}. However, the line is not statistically very significant and the errors on the EW are quite large (see Sect. \\ref{xmmanalysis}).  To test further the consistency of NGC 3147 being a type 1 AGN, we used the QSO template from \\citet{elvis94} and scaled it to the X-ray flux measured by our X-ray observation. The predicted continuum level at optical wavelengths is fully consistent with the measured nuclear spectrum. In addition, we have examined the Optical Monitor \\citep{Mason01} data of our XMM-\\textit{Newton} observation, where the nucleus of NGC~3147 is detected as a point source. Once deconvolved the disk contribution and corrected for Galactic reddening, we get fluxes of $(1.81\\pm 0.14)$ and $(0.94\\pm0.14)\\times 10^{-15}$ erg cm$^{-2}$ s$^{-1}$ \\AA, at 2910 and 2310 \\AA, respectively, leading to a two-point spectral index $\\alpha_{ox}\\simeq-1.33$\\footnote{We adopt the original definition by \\citet{tan79} as the ratio between X-ray and UV flux densities: $\\alpha_{ox}=-0.384 \\log[f(2\\,\\mathrm{keV})/f(2500\\,\\mathrm{\\AA})]$.}, as found for most Seyfert 1s \\citep[e.g.][]{stra05}.  However, the analysis of the \\textit{simultaneous} optical spectrum presented in this paper confirms the lack of broad permitted lines in this source. The very low column density measured in the X-rays ($N_H<5\\times10^{20}$ cm$^{-2}$ at the 99\\% confidence level for two interesting parameters, corresponding to $\\mathrm{A_V}<0.3$) is at odds with the large amount of dust required to obscure the BLR. Generally, it is found that any deviation from the Galactic gas-to-dust ratio in AGN always goes in the opposite direction: the obscuration by dust is \\textit{lower} than what expected from the associated gas column density \\citep{mai01}. It is difficult to find a physical mechanism able to suppress gas without destroying dust.  If the source were Compton-thick, the lack of X-ray absorption could be apparent, because we would not observe the primary emission at all, but only the reflected light, unaffected by any line-of-sight absorption. This solution is untenable for NGC~3147, because of the EW of the neutral iron line, much smaller than the one expected in the case of a reflection-dominated object \\citep[$\\simeq1$ keV: see e.g.][]{mbf96}. However, the (relatively) large EW of the Fe\\textsc{xxv} line may be a signature of reflection from ionised material. In this case, the reprocessed spectrum would be indistinguishable from the primary powerlaw, but an Fe\\textsc{xxv} EW as low as some hundreds of eV would be likely accompanied by detectable emission from Fe\\textsc{xxvi} or lower ionised species \\citep[see e.g.][]{bm02}. This solution would require a rather ad hoc geometry: a Compton-thick neutral absorber blocks the primary emission, while the reprocessed spectrum is instead dominated by an highly ionised material. Moreover, the high ratio between the 2-10 keV and the [O\\textsc{iii}] reddening-corrected flux ($\\simeq20$) strongly rejects the hypothesis that we are not directly observing the primary emission \\citep[in this case, a ratio lower than 1 is expected: see e.g.][and references therein]{pb02}. Therefore, the source is quite unlikely to be Compton-thick.  The only solution left is that the source intrinsically lacks the BLR. In this respect, it is interesting to note that NGC~3147 has no hidden BLR (HBLR) in polarised light (Tran, private communication). It was shown that HBLR sources are characterized, on average, by larger luminosities than non-HBLR ones \\citep[e.g.][]{gh02}. Recently, \\citet{es06} presented a model which depicts the torus as the inner region of a clumpy wind outflowing from the accretion disc. A key prediction of this scenario is the disappearance of the BLR at very low bolometric luminosities. Adopting a constant bolometric correction of 20 \\citep[][but any other choice is irrelevant for the present estimate]{elvis94} to the observed X-ray luminosity, the bolometric luminosity of NGC~3147 is about $5\\times10^{42}$ erg s$^{-1}$, which is somewhat larger than the `threshold' calculated by \\citet{es06} and, in any case, larger than other low-luminosity objects showing broad lines \\citep{ho97}. Moreover, the disappearance of the BLR \\textit{follows} that of the torus, at lower luminosities. Therefore, their model predicts the absence of the torus in NGC~3147, requiring the strong neutral iron line to be produced in the accretion disc. Current data cannot confirm nor reject this hypothesis.  Another possibility is that the BLR is intimately linked to the accretion rate, rather than the luminosity, of the AGN, being formed in accretion disk instabilities occurring around a critical radius \\citep{nic00}. Below a minimum accretion rate, this radius falls below the innermost stable orbit and the BLR cannot form. We found in the recent literature two estimates for the black hole mass of NGC~3147: $6.2\\times10^{8}$ M$_{\\odot}$ \\citep[][inferred from the mass-velocity dispersion correlation]{merl03} and $2.0\\times10^{8}$ M$_{\\odot}$ \\citep[][using a mass-$K_s$ bulge luminosity relation]{dd06}. Combined with the above-derived bolometric luminosity, we get a very low Eddington rate for NGC~3147, roughly ranging between $8\\times10^{-5}$ and $2\\times10^{-4}$. In any case, this is well below the threshold (around $1\\times10^{-3}$) proposed by \\citet{nic00} and \\citet{nic03}, thus supporting this scenario.  In conclusion, it seems inescapable that the key parameter for the observational properties of this source is not orientation, but an intrinsic feature (be it low accretion rate or luminosity), which prevents the BLR to form. NGC~3147 is therefore a ``true'' Seyfert 2 without the BLR, the first unambiguous component of a new class of AGN which requires to be fitted in a more general Unified Model."
    },
    "0710/0710.0319_arXiv.txt": {
        "abstract": "We present new evolutionary models for Type Ia supernova (SN Ia) progenitors, introducing mass-stripping effect on a main-sequence (MS) or slightly evolved companion star by winds from a mass-accreting white dwarf (WD). The mass-stripping attenuates the rate of mass transfer from the companion to the WD.  As a result, quite a massive MS companion can avoid forming a common envelope and increase the WD mass up to the SN Ia explosion.  Including the mass-stripping effect, we follow binary evolutions of various WD + MS systems and obtain the parameter region in the initial donor mass -- orbital period plane where SNe Ia occur. The newly obtained SN Ia region extends to donor masses of $6-7 ~M_\\sun$, although its extension depends on the efficiency of mass-stripping effect.  The stripped matter would mainly be distributed on the orbital plane and form very massive circumstellar matter (CSM) around the SN Ia progenitor. It can explain massive CSM around SNe Ia/IIn(IIa) 2002ic and 2005gj as well as tenuous CSM around normal SN Ia 2006X. Our new model suggests the presence of very young ($\\lesssim 10^8$~yr) populations of SNe Ia, being consistent with recent observational indications of young population SNe Ia. ",
        "introduction": "The nature of Type Ia supernova (SN Ia) progenitors has not been clarified yet \\citep[e.g.,][]{nie04, nom00}, although it has been commonly agreed that the exploding star is a mass-accreting carbon-oxygen white dwarf (C+O WD).  For the exploding WD itself, the observed features of SNe Ia are better explained by the Chandrasekhar mass model than the sub-Chandrasekhar mass model \\citep[e.g.,][]{liv00}. However, there has been no clear observational indication as to how the WD mass gets close enough to the Chandrasekhar mass for carbon ignition; i.e., whether the WD accretes H/He-rich matter from its binary companion [single degenerate (SD) scenario], or two C+O WDs merge [double degenerate (DD) scenario].  Recently, the following two important findings have been reported in relation to the SN Ia progenitors: (1) circumstellar matter (CSM) around the progenitors, and (2) a very young ($\\lesssim 10^8$~yr) population of the progenitors.  {\\bf Circumstellar Matter:} In the SD scenario, H/He-rich CSM is expected to exist around SNe Ia as a result of mass transfer from the companion as well as the WD winds \\citep*[e.g.,][]{nom82, hkn99}. Thus searching for H/He-rich CSM is one of the key observations to identify the progenitors \\citep[e.g.,][]{lun03}. Recently detections of such CSM have been reported for several SNe Ia, i.e., observations of narrow H-emission lines in SNe 2002ic \\citep{hau03} and 2005gj \\citep{ald06,pri07} (Type Ia/IIn or IIa \\citep{den04}), thermal X-rays from 2005ke \\citep{imm06}, and \\ion{Na}{1}~D lines in 2006X \\citep{pat07a}.  The identification of SN 2002ic as an SN Ia has been confirmed by the recent spectral comparison between SN 2005gj and SNe Ia \\citep{pri07}, being against the Type Ic suggestion by \\citet{ben06}.  Several CSM interaction models suggested a $1 - 2 ~M_\\sun$ CSM \\citep{chu04,nom05}. The evolutionary origin of such a massive CSM has been explored by \\citet{liv03} based on a common envelope evolution model, by \\citet{han06} from the delayed dynamical instability model of binary mass transfer, and by \\citet{woo06} based on a recurrent nova model with a red giant companion.  For normal SNe Ia, non-detection of radio has put the upper limit of mass loss rate as $\\dot M / v_{10} \\lesssim 10^{-8} M_\\sun$~yr$^{-1}$, where $v_{10} \\equiv v / 10$~km~s$^{-1}$ \\citep{pan06}. However, the optical observations of SN 2006X have detected variable Na I D lines from CSM, whose expansion velocity and mass have been estimated to be $v_{10} \\sim$ 10 and $\\sim 10^{-4} ~M_\\sun$ \\citep{pat07a}.  Patat et al. have suggested that the CSM in SN 2006X originated from the red-giant companion because of relatively low velocities.  Comparing the SN 2006X light curves with the other normal SNe Ia light curves, \\citet{wan07a} suggested that the obvious deviation, the decline rate is slowing down in a later phase, can be explained by an interaction between ejecta and CSM or a light echo of circumstellar/interstellar matter \\citep[see also][]{wan07b}.  {\\bf Young Population:} According to \\citet{man06}, the present observational data of SNe Ia are best matched by a bimodal population of the progenitors, in which about 50 percent of SNe Ia explode soon after their stellar birth at the {\\sl delay time} of $t_{\\rm delay} \\sim 10^8$ yr, while the remaining 50 percent have a much wider distribution of the {\\sl delay time} of $t_{\\rm delay} \\sim$ 3 Gyr. \\citet{aub07} recently reported evidence for a short (less than 70 Myr) delay time component in the SN Ia population. In this paper, we define the term {\\sl delay time} as {\\sl the age of a binary system at the SN Ia explosion}, in order to compare our results with the earlier results \\citep[e.g.,][]{gre83, gre05, man06}.    \\begin{figure*} \\epsscale{0.6} \\plotone{f1.eps} \\caption{ A schematic configuration of a binary evolution including mass-stripping effect. (a) Here we start a pair of a C+O WD and a more massive main-sequence (MS) star with a separation of several to a few tens of solar radii. (b) When the secondary evolves to fill its Roche lobe, mass transfer onto the WD begins.  The mass transfer rate exceeds a critical rate for optically thick winds.  Strong winds blow from the WD. (c) The hot wind from the WD hits the secondary and strips off its surface. (d) Such stripped-off material forms a massive circumstellar disk or torus and it gradually expands with an outward velocity of $\\sim 10-100$~km~s$^{-1}$.  The interaction between the WD wind and the circumstellar torus forms an hourglass structure.  The WD mass increases up to $M_{\\rm Ia}= 1.38 ~M_\\sun$ and explodes as an SN Ia. When we observe the SN Ia from a high inclination angle such as denoted by ``line of sight,'' circumstellar matter (CSM) can be detected as absorption lines like in SN 2006X. \\label{stripping_evolution}} \\end{figure*}   This kind of short delay times ($t_{\\rm delay} \\lesssim 10^8$~yr) of SNe Ia have been suggested from the distribution of SNe Ia relative to spiral arms \\citep[e.g.,][]{bar94, del94}. Recently, \\citet{dis03} reported, based on the {\\it Chandra} data from four external galaxies: an elliptical galaxy (NGC 4967), two face-on spiral galaxies (M101 and M83), and an interacting galaxy (M51), that in every galaxy there are at least several hundred luminous supersoft X-ray sources (SSXSs) with a luminosity of $\\gtrsim 10^{37}$ erg~s$^{-1}$ and that, in the spiral galaxies M101, M83, and M51, SSXSs appear to be associated with the spiral arms. The latter may indicate that SSXSs are young systems, possibly younger than $10^8$ yr, and has some close relation to the young population of SNe Ia.  The SD scenario has ever not predicted such young populations of $t_{\\rm delay} \\sim 10^8$ yr, corresponding to, at least, the zero-age main-sequence (ZAMS) stars at mass $5-6 ~M_\\sun$ \\citep[see, e.g.,][]{lih97, hknu99, lan00, han04}. In the present paper, we propose a scenario for such a young SN Ia population by introducing {\\sl mass-stripping effect} into binary evolutions. Mass-accreting WDs blow optically thick winds when the mass transfer rate to the WD exceeds the critical rate of ${\\dot M}_{\\rm cr} \\sim 1 \\times 10^{-6} M_\\sun$~yr$^{-1}$ \\citep{hkn96}. The WD wind collides with the secondary's surface and strips off matter. When the mass-stripping effect is efficient enough, the mass transfer rate to the WD is attenuated and the binary can avoid the formation of a common envelope even for a rather massive secondary.  The mass-stripping effect on a MS companion has been first introduced by \\citet{hac03kb, hac03kc}, who analyzed two quasi-periodic transient supersoft X-ray sources, RX~J0513.9$-6951$ and V~Sge: RX~J0513 shows a quasi-periodic oscillation between optical high ($\\sim 100-120$ days) and low ($\\sim 40$ days) states with an amplitude of 1 mag \\citep{alc96}.  RX~J0513 is X-ray bright only during the optical low states \\citep{rei00}. \\citet{hac03kb} proposed a model that the mass transfer is modulated by the WD wind because the WD wind collides with the companion and strips off its surface and attenuates the mass transfer rate.  When the mass transfer rate decreases below the critical rate ${\\dot M}_{\\rm cr}$, the WD wind stops and supersoft X-ray turns on.  This corresponds to an optical low sate. Then the mass-transfer rate recovers because of no attenuation by WD winds and the WD blows winds again. X-ray turns off and an optical high state resumes and the binary starts the next cycle of quasi-periodic oscillation. Such a self-sustained model naturally explains major characteristics of quasi-periodic high and low states and this success encourages us to adopt the same idea in the evolution scenario of supersoft X-ray sources and SN Ia progenitors.  In the present paper, we show that this mass-stripping effect derives (1) formations of circumstellar matter (CSM) around SNe Ia and (2) a very young population of SNe Ia.  We summarize our basic treatments of mass-stripping effect and binary evolutions in \\S 2, and then show our numerical results and their relations to a very young population of SNe Ia in \\S 3.  In \\S 4 we present the origin of CSM around SNe Ia based on our results and show a relation between the very young population of SNe Ia and their massive CSM. Discussion and concluding remarks follow in \\S\\S 5 and 6.    \\begin{figure*} \\epsscale{0.7} \\plotone{f2.epsi} \\caption{ SN Ia evolutions for two typical cases of WIND and CALM. (a) Case WIND: starting from $M_{\\rm WD,0} =1.0 ~M_\\sun$, $M_{2,0} =5.0 ~M_\\sun$, and $P_0 = 2.15$ days with $c_1=3$, the WD reaches the SN Ia explosion in the wind phase at $t=6.57 \\times 10^5$ yr.  The WD mass ($M_{\\rm WD}$), secondary mass ($M_2$), mass loss rate from the secondary (${\\dot M}_2$), WD wind mass loss rate (${\\dot M}_{\\rm wind}$), radius of the secondary ($R_2$), effective radius of the Roche lobe for the secondary ($R_2^*$), and orbital period ($P_{\\rm orb}$) are plotted. Only the orbital period is multiplied by four to easily see its change. (b) Case CALM: starting from $M_{\\rm WD,0} =1.0 ~M_\\sun$, $M_{2,0} =5.0 ~M_\\sun$, and $P_0 = 6.79$ days with $c_1=3$, the WD reaches the SN Ia explosion but in an SSXS phase without winds at $t=  6.93 \\times 10^5$ yr. The WD wind stops at $t=5.5 \\times 10^5$~yr. Even after that, the WD loses its mass due to weak helium shell flashes \\citep{kat99h}.  Here ${\\dot M}_{\\rm wind}$ includes an average mass loss rate by helium shell flashes and thus does not become zero after the optically thick wind of steady hydrogen shell burning stops. Values of the secondary radius ($R_2$) and the Roche lobe radius for the secondary ($R_2^*$) are divided by two to squeeze them into the figure. \\label{evolution_sn2005gj}} \\end{figure*} ",
        "conclusions": "Both Cases WIND and CALM originate from the systems with massive donors, i.e., young population.  It would be important to make some comparisons with the observational data, such as frequency and population.  The red hatched regions in Figures \\ref{zams_evl_con_c3_m110}, \\ref{zams_evl_con_c3_m100}, and \\ref{zams_evl_con_c3_m090} indicate a region in which the progenitor explodes at $t_{\\rm delay} \\le 100$~Myr.  Also the dashed line and the dotted lines correspond to $t_{\\rm delay} =$ 200~Myr and 400~Myr, respectively.  We see in Figure \\ref{birth_rate_population_ms} that Case WIND and thus SNe Ia/IIn (IIa) are realized by the very young system with $t_{\\rm delay} \\lesssim 100-200$~Myr.  If $M_{\\rm WD, 0} \\lesssim 0.9 ~M_\\sun$, we have almost no region of Case WIND, different from the cases of $M_{\\rm WD, 0} \\gtrsim 1.0 ~M_\\sun$.  If all the WD + MS system with $M_{2,0} \\gtrsim 3-6 ~M_\\sun$ ($c_1=3$), $M_{\\rm WD,0} \\gtrsim 1.0 ~M_\\sun$ ($M_{1,0} \\gtrsim 6.5~M_\\sun$), and $P_0 \\sim 0.5 - 2$~days produces SNe Ia/IIn (IIa) events (Table \\ref{condition_sn1a_explosion}), the frequency of these events is estimated to be $\\sim 5$\\% (including both the WD + MS and WD + RG systems with their total number ratio of 4:2).  A group of Type IIn SNe such as SNe 1997cy and 1999E show a very similar spectroscopic and photometric features to SN 2002ic \\citep{wan04, den04, pri07}. If these are in fact all Type Ia/IIn (IIa) SNe, their frequency can be estimated to be $\\sim 5_{-4}^{+7}$~\\% \\citep{pri07}, which is consistent with the above estimate.  Type Ia supernovae play a key role in astrophysics, and thus our progenitor model has important implications. Our model depends essentially on the parameter of stripping effect, $c_1$, which depends on the properties of WD winds, such as asphericity, velocities, and the efficiency of energy conversion. Also we calculate the mass transfer rate using the simple approximate binary models.  In order to improve these parameterization and approximations, we need multi-dimensional hydrodynamical simulations, which are beyond the scope of the present study.  In the present approach, we constrain the $c_1$ parameter observationally, and estimate $c_1 \\sim 7-8$ and $c_1 \\sim 1.5-10$ from the analysis of V~Sge and RX~J$0513.9-6951$, respectively.  With keeping in mind the necessity of further theoretical and observational studies to confirm our new progenitor systems, we summarize the basic results of our new SN Ia scenario:  (1) Mass-accreting WDs blow an optically thick wind when the mass transfer rate to the WD exceeds the critical rate of ${\\dot M}_{\\rm cr} \\sim 1 \\times 10^{-6} M_\\sun$~yr$^{-1}$.  The WD wind collides with the secondary's surface and strips off its surface.  If the mass-stripping effect is efficient enough, the mass transfer rate to the WD is attenuated and the binary can avoid formation of a common envelope even for a rather massive secondary. Including this mass-stripping effect into our binary evolution model of the WD + MS systems, we have found a new evolutionary scenario, in which a companion as massive as 6--$7~M_\\sun$ can produce an SN Ia for a reasonable strength of mass-stripping effect, say $c_1 \\sim 3$.  (2) We have followed simplified binary evolutions and obtained the SN Ia region in the $\\log P_0$--$M_{2,0}$ (initial orbital period -- initial donor mass) plane.   The newly obtained SN Ia region extends to massive donor masses up to $M_{2,0} \\sim 6-7 ~M_\\sun$ for $P_0 \\sim 0.5-10$ days, although its extension depends on the strength of mass-stripping effect, $c_1$, i.e., $M_{2,0} \\sim 7-8 ~M_\\sun$ for $c_1=10$, $M_{2,0} \\sim 5-6 ~M_\\sun$ for $c_1=3$, and $M_{2,0} \\sim 4 ~M_\\sun$ for $c_1=1$.  (3) We have estimated that the SN Ia birth rate in our Galaxy is $\\nu_{\\rm WD+MS} \\sim 0.004$~yr$^{-1}$ (for $c_1 = 3$), which is consistent with the observation.  The rates of young populations, i.e., $t_{\\rm delay} \\le 100$~Myr and $t_{\\rm delay} \\le 200$~Myr, are about 50\\% and 80\\% of the total SN Ia rate of the WD + MS channel.  These short delay times of SN Ia progenitors are consistent with the recent observational suggestions that a half of SNe Ia belong to such a very young population as the delay time of $t_{\\rm delay} \\sim 10^8$~yr.  (4) Another channel of the WD + RG system shows a broad distribution of the delay time over 2--3 Gyr \\citep{hkn99}, thus the two (WD + MS and WD + RG) channels yield a bimodality of the delay time distribution.  (5) The stripped-off material is probably distributed on the orbital plane and forms a massive circumbinary torus (or disk) around SNe Ia. Such circumstellar matter (CSM) may be consistent with the observed CSM feature in SN 2006X. When SN ejecta strongly interact with massive CSM, it can explain the feature of Type Ia/IIn (IIa) SNe 2002ic and 2006gj.  (6) Three different environments of SN Ia explosions can be specified by three different states of WDs just at the SN Ia explosion, i.e., the optically thick WD wind phase (Case WIND), steady hydrogen burning phase without optically thick winds from WDs (Case CALM), and recurrent nova phase (Case RN). In Case WIND, SN Ia ejecta strongly interact with massive CSM like SNe Ia/IIn (IIa) 2002ic and 2005gj because CSM exists near the SN Ia. The estimated rate of Case WIND is $\\sim 5$\\% of the total SN Ia rate, being consistent with the observational estimate. In Cases CALM and RN, SNe show a normal SN Ia feature because the CSM is far from the SN but the ejecta may interact with the CSM in a much later phase.  SN 2006X may be on a border between Case WIND and Case CALM."
    },
    "0710/0710.5220_arXiv.txt": {
        "abstract": "{Although many radio loud quasars and galaxies have been observed in X~rays, systematic studies of well defined samples are rare.} {We investigate the X-ray properties of the most luminous radio sources in the 3CR catalogue, in order to assess if they are similar to the most luminous radio quiet quasars, for instance in the X-ray normalization with respect to the optical luminosity, or in the distribution of the absorption column density.} {We have selected the (optically identified) 3CR radio sources whose 178-MHz monochromatic luminosity lies in the highest factor-of-three bin. The 4 most luminous objects had already been observed in X~rays. Of the remaining 16, we observed with XMM-Newton 8 randomly chosen ones, with the only requirement that half were of type~1 and half of type~2 according to the optical identification.} {All targets have been detected. The optical-to-Xray spectral index, $\\alpha_{ox}$, can be computed only for the type~1s and, in agreement with previous studies, is found to be flatter than in radio quiet quasars of similar luminosity. However, the Compton thin type~2s have an absorption corrected X-ray luminosity syste\\-ma\\-ti\\-cally lower than the type~1s, by a factor which makes them consistent with the radio quiet $\\alpha_{ox}$. Within the limited statistics, the Compton thick objects seem to have a reflected component more luminous than the Compton thin ones.} {The extra X-ray component observed in type~1 radio loud quasars is beamed for intrinsic causes, and is not collimated by the absorbing torus as is the case for the (intrinsically isotropic) disk emission. The extra component can be associated with a relativistic outflow, provided that the flow opening angle and the Doppler beaming factor are $\\sim 1/5$ -- $1/7$ radians.} ",
        "introduction": "A strong X-ray emission is a defining property of Active Galactic Nuclei, and radio loud quasars and galaxies share this property. The most interesting entries of the radio catalogues have been observed repeteadly with all the X-ray satellites launched so far. Although much has been learned, our knowledge is still fragmentary in comparison with the radio quiet AGN. The latter dominate the X-ray sky, and it is relatively straightforward to assemble large samples which can then be studied at other wavelengths. On the contrary, the luminous extragalactic radio sources are rare, and one has to observe them one by one investing large amounts of satellite time.  There are a few points on which a consensus has been established. The type~1 radio loud quasars have an X-ray emission stronger than their radio quiet analogues of similar optical luminosity, which is quantified by a flatter $\\alpha_{ox}$ (e.g. \\cite{brink} versus \\cite{steffen}). There are indications that the X-ray photon index $\\Gamma$ is syste\\-ma\\-ti\\-cally flatter in radio loud quasars (\\cite{shastri}, but \\cite {brink} have a more cautious view). At any rate, a flatter $\\Gamma\\sim 1.5$ -- 1 is well established in flat spectrum radio sources and low frequency peaked blazars (\\cite{grandi}, \\cite{fossati}). The type~2 radio galaxies are commonly found to host absorbed AGN, a paradigmatic case is the near\\-by powerful source Cygnus~A (\\cite{young}). Even if absorbed below energies of several keV, the radio galaxies are frequently detected around 1 keV at a level of a few percent of the main emission (\\cite{diana} and references therein). It is unclear if the residual emission is due to scattering of the main one, or is completely unrelated.  The basic picture is the extension of the so called unified scheme to the case where the accretion disk is complemented with a relativistic outflow (\"jet\"). The radio galaxies and radio quasars are the misaligned and aligned members, respectively, of the same population. The optical and soft X-ray emission of the former is obscured by some intervening medium, perhaps arranged in a toroidal geometry. The jet emission, because of the Doppler boosting, becomes more and more prominent along a sequence where the line of sight gets closer and closer to the jet axis, passing from steep spectrum to flat spectrum radio quasars, and to blazars.  \\begin{table*} \\caption{The sample objects.} \\label{table:1} \\centering \\begin{tabular}{l c c c c c c} \\hline\\hline Name & Type & Redshift & Radio Luminosity & Magnitude & Exposure Time & References \\\\ & & $z$ & L$_{\\rm r}$ & & (ks) & \\\\ \\hline 3C298 & Q & 1.439 & 29.79 & 16.79~(V) &  20~(C) & 1 \\\\ 3C9~~ & Q & 2.012 & 29.72 & 18.21~(V) &  16~(C) & 2 \\\\ 3C257 & G & 2.474 & 29.67 & 18.07~(K) &  30~(X) & 3 \\\\ 3C191 & Q & 1.956 & 29.55 & 18.65~(V) &  17~(C) & 4 \\\\ 3C239 & G & 1.781 & 29.46 & 22.50~(V) &  14~(X) & 5 \\\\ 3C454 & Q & 1.757 & 29.39 & 18.47~(V) &  16~(X) & 5 \\\\ 3C432 & Q & 1.785 & 29.38 & 17.96~(V) &  11~(X) & 5 \\\\ &   &       &       &           &  20~(C) & 4 \\\\ 3C294 & G & 1.786 & 29.35 & 18.00~(K) & 118~(C) & 6 \\\\ 3C249 & G & 1.554 & 29.34 & 18.90~(K) &  35~(X) & 5 \\\\ 3C241 & G? & 1.617 & 29.30 & 23.50~(V) &  25~(X) & 5 \\\\ 3C318 & Q? & 1.574 & 29.30 & 20.30~(V) &  19~(X) & 5 \\\\ 3C205 & Q & 1.534 & 29.28 & 17.62~(V) &   5~(X) & 5 \\\\  \\hline \\end{tabular} \\begin{list}{}{} \\item[] References. (1)~\\cite{aneta}; (2)~\\cite{fabian}; (3)~\\cite{derry}; \\item[] (4)~\\cite{erlund}; (5)~our data; (6)~\\cite{fabian2}. \\end{list} \\end{table*}  In this paper we present and discuss a study of the X-ray properties of the most luminous low frequency radio sources. By selecting at 178 MHz, the frequency of the 3CR catalogue (\\cite{3cr}), we select sources with a luminous extended, isotropic component, which is thought to arise from the accumulation of relativistic plasma over the entire lifetime of the system. The AGN in our sample are then expected to have had a large {\\it time averaged} power in their past life, whereas the X-ray lu\\-mi\\-no\\-si\\-ty depends on the {\\it instantaneous} power at the present time. A correlation between the two holds only in a statistical sense. On the other hand, by selecting with respect to an isotropic component we avoid all the uncertainties connected with the Doppler boosting. We restrict ourselves to the 3CR sources at high Galactic latitude ($|\\ell|>20^{\\circ}$) with optical identification and redshift (\\cite{ident}), and order them with respect to L$_{\\rm r}$~, a radio luminosity parameter equal to the logarithm of the monochromatic luminosity at 178 MHz in W~Hz$^{-1}$~m$^{-2}$~s$^{-1}$. We adopt the concordance cosmology with $\\Omega_{\\rm m}=0.3, \\Omega_{\\Lambda}=0.7$, and $\\rm H_{\\circ}= 70~km~s^{-1}~Mpc^{-1}$. We neglect the $k$-correction because all the objects of interest happen to be steep spectrum radio sources, i.e. they have a spectral slope $\\alpha$ (f$_{\\nu} \\propto \\nu^{-\\alpha}$) close to 1 in the 100-MHz region. In half a decade, from L$_{\\rm r}$~= 29.79 to L$_{\\rm r}$~=29.28, one has 20 objects.  The four most luminous ones had already been observed by previous investigators. We chose at random 8 additional objects among the remaining 16, with the only requirement that 4 were optical type~1 and 4 optical type~2, and observed them with XMM-Newton. One of the 8 (3C294) had previous Chandra observa\\-tions, and we did not reobserve it. One further object (3C432) was also observed with Chandra after our XMM run had been sche\\-du\\-led. Due to the randomness of our choice, we believe that the final sample of 4+8 objects is a fair representation of the highest radio luminosity bin, although a rather slender one. The only infringement to randomness, i.e. the preselection of optical types, implies that we cannot draw inferences about the relative frequency of types from our X-ray data. ",
        "conclusions": "The main point of the present work has been already anticipated in the previous Section, and is the segregation in intrinsic luminosity between the unabsorbed type~1s and the absorption corrected, Compton thin type~2s. This will be shown to have far reaching implications. However our finding is plagued by the small number of objects involved. In order to put the result on firmer grounds, we have collected all the 3CR sources with X-ray observations having: optical identification and redshift; high Galactic latitude ($|\\ell|>20^{\\circ}$); and radio luminosity parameter L$_{\\rm r}$ in the factor-of-three interval immediately below our fiducial sample, i.e. $ 28.67 < \\rm L_{\\rm r} < 29.28$.  The additional sample contains 46 objects, of which only 12 observed in X~rays. With respect to the fiducial sample, the additional one is sparsely observed (26\\% versus 60\\%); furthermore, we have no control on the criteria by which the observed sources were chosen. In order to confirm or disprove the luminosity segregation which we want to investigate, a crucial point is a possible bias of the additional sample with respect to the orientation of the jet axis. One source is the well known blazar 3C454.3; its X-ray emission is completely dominated by the jet, so we do not consider it further in the following. All the remaining 11 sources have steep radio spectra, and none of them is classified as a blazar; for the sake of completeness, we point out that the spectra of three of them (3C380, 3C309.1 and 3C212) become flat above a few GHz, perhaps a hint that their orientation angles are at the lower end of the non-blazar region. Also the distribution of optical properties (4 type 1's, 6 type 2's, and one intermediate type) does not exhibit obvious peculiarities.  Table~\\ref{table:3} gives some basic information about the 11 additional sources; more specifically we list: the 3C name; the optical and X-ray type; the redshift $z$; the radio luminosity parameter L$_{\\rm r}$; the intrinsic and reflected X-ray luminosity, L$_{2-10}$ and L$_{\\rm refl}$ (actually, the logarithm of the luminosity in erg s$^{-1}$); and the reference to the X-ray data. Note that 3C325 is a further case of intermediate optical classification (\\cite{grimes}).  \\begin{figure}[h!] \\resizebox{\\hsize}{!}{\\includegraphics[angle=0]{salvatifig4nuova.eps}} \\caption{The X-ray luminosity of the sources in our sample (fiducial plus additional), plotted against the radio luminosity parameter. Filled squares, quasars; stars, intermediate optical types; empty squares, Compton thin galaxies (the absorption corrected direct component); filled triangles, Compton thick galaxies; filled triangle between parentheses, 3C249 (see text); empty triangles, Compton thin galaxies (the reflected component). The upper and lower solid lines sketch the regions occupied by type 1's and type 2's, respectively; the dotted lines represent the 5\\% fraction of the solid lines.} \\label{figure:4} \\end{figure}  Figure \\ref{figure:4} summarizes the information of Tables \\ref{table:2} and \\ref{table:3}: we plot the logarithm of the X-ray luminosity, Log(L$_{2-10}$) and Log(L$_{\\rm refl}$), against the radio luminosity parameter, L$_{\\rm r}$. The meaning of the symbols is explained in the caption. For the sake of clarity we omit the errorbars: the radio fluxes have typical errors of 1 Jy at the 3C flux limit of 10 Jy, so that the horizontal errorbars are less than 0.04 dex; the X-ray values have typical errors of less than 25\\%, i.e. less than 0.1 dex. Only the reflected components of Compton thin galaxies, the empty triangles in the Figure, have larger errors of the order of 0.3 dex.  The lower right corner is expected to be underpopulated, given the limit of 10~Jy for the 3C catalogue and a typical depth of 10$^{-15}$ erg cm$^{-2}$ s$^{-1}$ for the X-ray exposures.  The two parallel solid lines sketch the regions occupied by the type~1s and the Compton thin type~2s, respectively. The difference of about 0.75 dex is fully consistent with the average values found in the previous Section for the fiducial sample only, i.e. the addition of more sources confirms the luminosity segregation of the two types. Also, the intermediate optical types are confirmed to lie at intermediate X-ray luminosities. The dotted lines indicate the 5\\% fraction of the type~1 and type~2 luminosity, respectively. This particular value is generally regarded as an upper limit for the reflected or scattered components, because of various constraints on geometry, efficiency, and optical depth.  In the case of type~1s the optical emission of the AGN is visible, and one can compute the index $\\alpha_{ox}$. We do that for all the type~1s of the fiducial sample; we derive the rest frame monochromatic fluxes at 2 keV and at 2500 \\AA\\ by means of the observed $\\Gamma$ in the X~rays, and by assuming $\\alpha$ = 0.5 in the optical (f$_{\\nu} \\propto \\nu^{-\\alpha}$). We find $\\alpha_{ox} = 1.35 \\pm 0.11$. At the average optical flux of our sources, 1.9 $10^{31}$ erg cm$^{-2}$ s$^{-1}$ Hz$^{-1}$, the radio quiet AGN have $\\alpha_{ox} = 1.65 \\pm 0.05$ (\\cite{steffen}), thus we recover the well known result (cf. the Introduction) that radio loud quasars are more X-ray luminous than radio quiet ones of the same optical luminosity.  This is not true for the radio galaxies, however. If we take into account the lower intrinsic X-ray luminosity of our type~2s we find $\\alpha_{ox} \\sim 1.64$, fully compatible with the radio quiet one. Here we do not observe directly the optical emission of the AGN, and we have assumed that type~2s and type~1s with the same L$_{\\rm r}$ have the same intrinsic optical luminosity. The additional component characteristic of the jetted sources becomes visible at smaller and smaller viewing angles for longer and longer wavelengths; indeed, even the high luminosity blazars exhibit optical broad lines with equivalent widths within a factor of two of the non blazar AGN (\\cite{pian}, \\cite{wills}). This points to the jet optical emission never being dominant over the disk optical emission, and gives some ground to our assumption.  The X-ray normalization relative to the optical, and the (meager) evidence provided by the X-ray spectral features (slope $\\Gamma$, equivalent width of the Fe line), suggest that at the large viewing angles typical of radio galaxies we see only the X-ray emission due to an accretion disk very similar to the radio quiet AGN. The self-consistent normalization of X-ray, optical and low frequency radio emission does not leave much room for \"fossil\" quasars, where the extended radio lobes survive for a substantial time after the central engine is switched off.  The additional X-ray emission connected with the ``radio loudness'' is not seen in radio galaxies, not because of the intervening absorption, but because this emission is intrinsically beamed. The obvious explanation for the beaming is the relativistic aberration of the jet emission. One can try a more quantitative assessment by adopting the distribution of viewing angles proposed by Barthel (\\cite{bart}), according to whom radio loud AGN appear as type~1s or type~2s if their viewing angle is smaller or larger than $\\sim 44^{\\circ}$; thus the average viewing angle is $\\theta_1\\sim 31^{\\circ}$ and $\\theta_2\\sim 69^{\\circ}$ for the two classes, respectively. Within the type~1 class, one finds the blazar subclass which corresponds to viewing angles smaller than the jet beaming angle, $\\rm \\theta_b < 10^{\\circ}$.  If one chooses the most favorable scenario, i.e. a jet of constant length and a very flat spectral slope ($\\Gamma=1$), the ratio of the brightness in the two directions $\\rm \\theta_b$ and $\\theta_1$ is  $$ \\rm R_X = [(1-\\beta cos\\theta_1)/(1-\\beta cos\\theta_b)]^2,  \\qquad \\beta = \\sqrt{1-\\gamma^{-2}}                       $$  \\noindent where $\\gamma$ is the bulk Lorentz factor of the jet. This ratio is equal to about 80 for $\\gamma = 15$, as suggested by the analysis of the wide band spectral energy distribution of blazars (\\cite{ghisa}). Indeed, by using the maximum viewing angle for a blazar, $\\rm \\theta_b$, instead of the average one, we have further underestimated $\\rm R_X$. We conclude that the typical X-ray luminosity of high L$_{\\rm r}$ blazars should be $ \\rm R_X$ times the typical X-ray luminosity of steep spectrum quasars of comparable L$_{\\rm r}$, i.e. at least 3 10$^{47}$ erg s$^{-1}$. No such blazar is known all over the sky. A possible solution to this inconsistency could be a geometrically wide, stratified jet: here the narrow and fast spine would be responsible for the blazar properties, while the slower (but still relativistic) sheath would produce the wide angle additional component. Similar schemes have been suggested already, both for AGN (\\cite{chiabe}) and for Gamma Ray Bursts (\\cite{grb}).  The above equation for $\\rm R_X$ can in principle be applied to the two directions $\\theta_1$ and $\\theta_2$ to derive the residual jet emission in radio galaxies. This however has nothing to do with the L$_{\\rm refl}$ points in Fig.~\\ref{figure:4}: they refer to a spectral component produced above the absorbing material, at a distance of $\\sim$parsecs from the source, whereas time variability constrains the blazar emission to much smaller dimensions (e.g. \\cite{ghisa}).  The correct comparison for the L$_{\\rm refl}$ points is with the 5-percent-fraction dotted lines. One notes that in Fig.~\\ref{figure:4} the empty triangles lie around or below the lower dotted line, which implies that the reflected component of Compton thin type~2s has the expected normalization with respect to the intrinsic disk component. The filled triangles, instead, are systematically higher, and lie around or below the higher dotted line, as if they were normalized to the wide jet component. In other words, if one insisted in attributing the visible emission of Compton thick objects to reprocessing of the disk emission only, one would have to accept reprocessed fractions as high as 30\\%.  The statistics is certainly not compelling, however a possible scenario is one where the optical depth at large viewing angles is not the same in all objects. Compton thin and Compton thick type~2 objects would be viewed at similar angles, around $70^{\\circ}$, and would have relatively little and relatively large amounts of matter around them, respectively. Thus the Compton thickness would be related with a larger covered solid angle, perhaps so large as to reprocess also the moderately beamed wide jet emission. Under this hypothesis the peculiar position of 3C241 is nicely accounted for. This source exhibits at the same time an intermediate X-ray luminosity and a relatively large N$_{\\rm H}$, as if the line of sight were grazing at the same time the boundaries of the wide jet and the boundaries of a particularly large absorber: indeed, here the reprocessing fraction is as high as 8\\%, and the corresponding empty triangle falls right in the region of Compton thick sources."
    },
    "0710/0710.2400_arXiv.txt": {
        "abstract": "Invisible plasma content in blazar jets such as protons and/or thermal electron-positron ($e^{\\pm}$) pairs is explored through combined arguments of dynamical and radiative processes. By comparing physical quantities required by the internal shock model with those obtained through the observed broadband spectra for Mrk 421, we obtain that the ratio of the Lorentz factors of a pair of cold shells resides in about $2\\sim 20$, which implies that the shocks are  at most mildly relativistic. Using the obtained Lorentz factors, the total mass density $\\rho$ in the shocked shells is investigated. The upper limit of $\\rho$ is obtained from the condition that thermal bremsstrahlung emission should not exceed  the observed $\\gamma$-ray luminosity, whilst the lower limit is constrained from the condition that the energy density of non-thermal electrons is smaller than that of the total plasma. Then we find $\\rho$ is $10^2$-$10^3$ times heavier than that of non-thermal electrons for pure $e^{\\pm}$ pairs, while  $10^2$-$10^6$ times heavier for pure electron-proton ($e/p$) content, implying the existence of a large amount of invisible plasma. The origin of the continuous blazar sequence is shortly discussed and we speculate that the total mass density and/or the blending ratio of $e^{\\pm}$ pairs and $e/p$ plasma could be new key quantities for the origin of the sequence. ",
        "introduction": "The discovery of strong inverse Compton components in $X$ and $\\gamma$-ray emission from jets in active galactic nuclei (hereafter AGN) for a wide range of spatial scales (e.g., Collmar 2001 for review) enables us to probe quantitatively the energetics of relativistic jets. The kinetic power of non-thermal electrons has been estimated by various authors both for inner core jets (i.e., blazars) (e.g., Kino, Takahara and Kusunose 2002, hereafter KTK; Kusunose, Takahara and Kato 2003) and large scale jets (e.g., Tavecchio et al. 2000; Leahy and Gizani 2001, 2002; Kataoka et al. 2003). However, the material content of relativistic jets is not easily constrained by observations since the emission is dominated by that from non-thermal electrons and probably positrons and it is difficult to directly constrain thermal matter content. Hence, the plasma composition in AGN jets, whether normal proton-electron ($e/p$) plasma or electron-positron pairs ($e^{\\pm}$) is a dominant composition, is still a matter of open issue (e.g., Reynolds et al. 1996; Celotti, Kuncic, Rees and Wardle 1998; Wardle et al. 1999; Hirotani et al. 1999; Sikora and Madejski 2000; Ruszkowski and Begelman 2002; Kino and Takahara 2004, hereafter KT04). This problem prevents us from estimating the total mass and energy flux ejected from a central engine.  To constrain invisible matter content such as thermal electron-positron pairs and/or protons co-existing with non-thermal electrons, dynamical considerations are indispensable. In KT04, we proposed a new procedure to constrain the invisible thermal plasma component in classical FR II radio sources. We used the fact that the mass and energy densities of the sum of thermal and non-thermal particles are larger than those of non-thermal electrons which are determined by observations. Here we apply the same technique to the inner core jets of AGNs (i.e., blazars) based on the internal shock model. The internal shock model is believed to be most plausible to explain the production of high energy photons and time variabilities in blazars. It has been widely applied also to the prompt emission of gamma-ray bursts (hereafter GRBs) (e.g., Rees 1978; Rees and Meszaros 1994; Kobayashi, Piran and Sari 1997;, Daigne and Mochkovitch 1998; Ghisellini 1999; Spada, Ghisellini, Lazzati and Celotti 2001). It is worth to note that recently Ghisellini et al. (2005) proposed a structured jet model cosisting of a fast spine surrounded by a slowly moving layer for explaining VLBI scale radio blobs. At the present, however, it is not evident where is the acceleration site of electrons in the structured jet model. This is one of the prime issues which should be answered. Internal shocks are potentially the building blocks of the spine part of the structured jet. Whereas we recognize the importance of the detailed structure of jets, as a first step we focus on the physical condition of the flow based on the simple internal shock model.     The methodology of constraining the invisible plasma content in the emission region is as follows. As mentioned above, a lower limit to the total mass density (sum of non-thermal electrons and invisible plasma) is restricted by the definition that the mass density of total plasma should be smaller than that of the non-thermal electrons. The mass density of non-thermal electrons can be estimated by multi-frequency observations. For this purpose, in \\S \\ref{sec:shock} we review the shock dynamics of two colliding shells. Note that we do not use the simple two point-mass approximation (e.g., Piran 1999; Lazzati et al. 1999; Zhang and  M{\\' e}sz{\\' a}ros 2004 for review) but employ the exact shock dynamics throughout this work. This makes outcomes more accurate. In \\S \\ref{sec:NT} we briefly review the previous results on the amount of non-thermal electrons based on KTK. In \\S \\ref{sec:invisible}, we constrain on the amount of total mass density. As for the upper limit, we use the constraint that bremsstrahlung emission from thermal electron (and positron) component should not exceed the observed $\\gamma$-ray emission. We postulate synchrotron self-Compton (SSC) emission dominance in the $\\gamma$-ray band which is supported by the observed correlations between TeV$\\gamma$-ray and X-ray in TeV blazars (e.g., Takahashi et al. 1996, 2000; Catanese et al. 1997; Maraschi et al. 1999). We can thus bracket the amount of total mass density in the emission region from below and above. In this way we apply this method to the archetypal TeV blazar Mrk 421. In \\S \\ref{sec:dissipation}, we further estimate the shock dissipation rate of the colliding cold shells. The dissipation rate is a widely discussed quantity in literatures concerning gamma-ray bursts (e.g., Lazzati, Ghisellini and Celotti 1999; Piran 1999). The shock dissipation is believed to be the ultimate source of heating and accelerating particles. Summary and discussion are in \\S \\ref{sec:summary}. ",
        "conclusions": ""
    },
    "0710/0710.0775_arXiv.txt": {
        "abstract": "The ultraviolet spectra of all ``weak emission line central stars of planetary nebulae'' (WELS) with available {\\it IUE} data are presented and discussed. We performed line identifications, equivalent width and flux measurements for several features in their spectra. We found that the WELS can be divided in three different groups regarding their UV: i) Strong P-Cygni profiles (mainly in \\ion{C}{4} $\\lambda 1549$); ii) Weak P-Cygni features and iii) Absence of P-Cygni profiles. The last group encompasses stars with a featureless UV spectrum or with intense emission lines and a weak continuum, which are most likely of nebular origin. We have measured wind terminal velocities for all objects presenting P-Cygni profiles in \\ion{N}{5} $\\lambda 1238$  and/or \\ion{C}{4} $\\lambda 1549$. The results obtained were compared to the UV data of the two prototype stars of the [WC]-PG 1159 class, namely, A30 and A78. For WELS presenting P-Cygnis, most of the terminal velocities fall in the range $\\sim 1000-1500$ km s$^{-1}$, while [WC]-PG 1159 stars possess much higher values, of $\\sim$3000 km s$^{-1}$. The [WC]-PG1159 stars are characterized by intense, simultaneous P-Cygni emissions in the $\\sim 1150-2000$\\AA\\, interval of \\ion{N}{5} $\\lambda 1238$, \\ion{O}{5} $\\lambda 1371$ and \\ion{C}{4} $\\lambda 1549$. In contrast, we found that \\ion{O}{5} $\\lambda 1371$ is very weak or absent in the WELS spectra. On the basis of the ultraviolet spectra alone, our findings indicate that [WC]-PG 1159 stars are distinct from the WELS, contrary to previous claims in the literature. ",
        "introduction": "Central stars of planetary nebulae (CSPN) that are hydrogen deficient are generally divided in three main groups: [WR], PG 1159 and [WC]-PG 1159 stars. The first one present spectra with strong and broad emission lines mainly from He, C and O, similar to the Wolf-Rayet stars (WR) of Population I. Depending on the ionization stages of the elements dominating the atmosphere, [WR] stars are subdivided in [WCL], [WCE] and [WO] (Crowther et al. 1998; Acker \\& Neiner 2003). On the other hand, PG 1159 stars are quite distinct objects. They are pre-white dwarfs and show mainly absorption lines of \\ion{He}{2} and \\ion{C}{4} in their spectra (Werner et al. 1997). Only a handful of PG 1159 stars are known to possess wind features (Koesterke et al. 1998; Koesterke \\& Werner 1998). The [WC]-PG 1159 group present strong P-Cygni lines in the ultraviolet (e.g. \\ion{N}{5} $\\lambda 1238$ and \\ion{C}{4} $\\lambda 1549$) and resemble the PG 1159 stars in the optical. Three objects are considered  prototypes of this class: A 30, A 78 and Longmore 4 (at outburst).  The origin and evolution of hydrogen deficient CSPN have been investigated from different point of views during the last decade. Evolutionary models are now able to provide a reasonable match to the observed chemical abundances, although it is still debated the role of binarity and different thermal pulse models (De Marco \\& Soker 2002; Herwig 2001). Nebular analyses as well as sophisticated non LTE atmosphere models have been also very useful to address several questions regarding these central stars: it is generally inferred that the groups above mentioned form the following evolutionary sequence: late type [WR] $\\rightarrow$ early-type [WR] $\\rightarrow$ [WC]-PG 1159 $\\rightarrow$ PG 1159 $\\rightarrow$ non DA white dwarfs (see e.g. Zijlstra et al. 1994; Pen\\~a et al. 2001; Koesterke 2001). However, this scenario has important issues unsolved, such as the C/He mass ratio and the exact position of [WC]-PG 1159 stars in the HR diagram (Hamann 1997; Marcolino et al. 2007). Moreover, the evolutionary status of the so-called ``weak emission line stars'' and their relation to the [WC]-PG 1159 stars is not at all clear. Let us discuss it now.  In an extensive observational study, Tylenda et al. (1993) analyzed spectroscopy data of 77 hydrogen deficient CSPN. In their sample, 39 were classified as [WR] stars. The remaining objects were called ``weak emission line stars'' (WELS or [WELS]). According to these authors, the WELS show emission of \\ion{C}{4} $\\lambda 5805$ (actually \\ion{C}{4} $\\lambda \\lambda 5801, 12$) systematically weaker and narrower than in [WR] stars, and a feature in $\\lambda 4650$, which is possibly a blend of \\ion{N}{3}, \\ion{C}{3} and \\ion{C}{4} emissions. Moreover, \\ion{C}{3} $\\lambda 5696$ is very weak or absent. These optical characteristics were confirmed by Marcolino \\& de Ara\\'ujo (2003) with a homogeneous and higher resolution set of data. Interestingly, by comparing a large sample of WELS, [WC]-PG 1159 and PG 1159 spectra, Parthasarathy et al. (1998) claimed that the WELS and the [WC]-PG 1159 stars actually constitute the same class. However, this claim was based solely on comparisons of optical spectra and was not confirmed by further studies. Indeed, some authors argued that this assertion should be taken with caution until an analysis of a larger sample of objects is performed (e.g. Werner \\& Herwig 2006).  Undoubtedly, the ``weak emission line stars'' (WELS) constitute the least understood class of hydrogen deficient CSPN. An important point is that the WELS might not be descendants of the [WR] stars, as it is explicited in the [WR] $\\rightarrow$ [WC]-PG 1159 (=WELS) $\\rightarrow$ PG 1159 evolution. Pen\\~a et al. (2003) for example, derived lower average nebular expansion velocities for WELS than for [WR] stars, while the contrary would be expected from the long-term action of a stellar wind on a planetary nebula during a [WR] $\\rightarrow$ WELS transition. Furthermore, based on a kinematical study of a large sample of planetary nebulae, Gesicki et al. (2006) raised the interesting possibility that some WELS can be progenitors of the hottest [WR] stars, i.e., the [WO] group (WELS $\\rightarrow$ [WO]). A major issue hinders the determination of the evolutionary status of the WELS: their physical parameters (e.g. $T_{eff}$, $v_\\infty$ and $\\dot{M}$) and chemical abundances remain unknown. The properties of A 30, A 78, and Longmore 4 ([WC]-PG 1159 stars) are known (Koesterke 2001) but again, their identity as WELS is questionable.  So far, most of previous studies involving WELS were done in the optical part of the spectrum. In fact, the very definition of a WELS is based in the optical features in $\\lambda 4650$, \\ion{C}{4} $\\lambda 5805$ and \\ion{C}{3} $\\lambda 5696$ (Tylenda et al. 1993). A few works in the ultraviolet including some WELS can be found in the literature. However, they often include stars of different spectral classes (e.g. Feibelman 2000), and/or are focused in the determination of nebular properties (e.g. Adams \\& Seaton 1982; Pottasch et al. 2005), without special attention to the WELS and their evolutionary state. Motivated by this fact and considering the open questions above described, we investigate in this paper the UV spectra of all WELS with available {\\it IUE} ({\\it International Ultraviolet Explorer}) data. Our main aims are to understand their main UV characteristics; identify and measure the most intense spectral lines; measure wind terminal velocities from the available P-Cygni profiles; and finally, to compare the results to the data of the two prototype [WC]-PG 1159 stars: A30 and A78.  The present paper is organized as follows: in Section 2 we present the observational data retrieved from the Multi-mission Archive at STScI (MAST); in Section 3 we discuss the main characteristics of the UV spectra of the WELS, and present line identifications of the most conspicuous features, as well as equivalent width and line flux measurements. In Section 4 we empirically measure the terminal velocities for all objects presenting P-Cygni profiles in \\ion{N}{5} $\\lambda 1238$ and/or \\ion{C}{4} $\\lambda 1549$, from low and high resolution data. Finally, in Section 5, we discuss and compare the results obtained for the WELS to the data of A 30 and A 78 (the two prototype [WC]-PG 1159 stars), and present the main conclusions of our work. ",
        "conclusions": "\\label{discussion}  From a comparison between several optical spectra of WELS, [WC]-PG 1159, and PG 1159 stars, Parthasarathy et al. (1998) have proposed that the WELS are actually [WC]-PG 1159 stars. This claim was not confirmed by further studies and as we mentioned, some authors have warned about this assertion until a more comprehensive study is achieved (Werner \\& Herwig 2006). After an analysis of the main UV characteristics of the WELS and the two prototype [WC]-PG 1159 stars A30 and A78, our next step was to compare the results obtained for these two class of objects in order to address this important issue.  As we have shown, most of the WELS present a UV spectrum considerably different from the [WC]-PG 1159 stars. While this last class present simultaneously P-Cygni profiles in \\ion{N}{5} $\\lambda 1238$, \\ion{O}{5} $\\lambda 1371$, and \\ion{C}{4} $\\lambda 1549$, the majority of WELS present a very weak or no \\ion{O}{5} $\\lambda 1371$ (see Fig. \\ref{group1}). The same is true at least for some objects regarding the \\ion{N}{5} $\\lambda 1238$ line. The only exceptions are the objects NGC 6543, NGC 6567, and NGC 6572. Their spectra in fact resemble the ones of the [WC]-PG 1159 stars (see Figs. \\ref{wcpg} and \\ref{wcpg2}): their \\ion{O}{5} $\\lambda 1371$ line is clearly visible as well as the other transitions mentioned.  Besides the spectral differences found in the ultraviolet part of the spectrum, we have also found that the terminal velocities of the WELS are considerably lower than in [WC]-PG 1159 stars. Our Table 6 shows that the bulk of the WELS have $v_\\infty$ between $\\sim 1000-1500$ km s$^{-1}$, and A30 and A78 have values about $3000$ km s$^{-1}$. This difference might represent different physical parameters underlying these two class of stars. The theory of radiatively driven stellar winds for example, predicts that the terminal velocity of a star is related with the escape velocity ($v_{esc}$), which in turn depend on other physical parameters (e.g. mass and radius; Abbott 1978; Lamers \\& Cassinelli 1999). Moreover, $v_\\infty$ is also known to correlate with the effective temperature ($T_{eff}$) in several class of stars (see Fig. 8 of Prinja et al. 1990). The $T_{eff}$ tends to be higher for stars with high terminal velocities.  From the considerations above described, we conclude that the [WC]-PG 1159 stars are distinct from the WELS, in contrast with the claim made by Parthasarathy et al. (1998) on the basis of optical spectroscopy. It should be noted however, that the situation for the central stars NGC 6543, NGC 6567 and NGC 6572 is ambiguous. From one side, they have a spectrum compatible with the [WC]-PG 1159 class. On the other hand, they do not present high terminal velocities as it is the case of A30 and A78. From both low and high resolution data it is obtained $v_\\infty$ values less than $2000$ km s$^{-1}$ for these three stars (see Table 6).  If the WELS are not [WC]-PG 1159 stars, what is their role in the evolutionary sequence [WR] $\\rightarrow$ PG 1159 ? Do they form an alternative channel of evolution ? In order to elucidate these and other similar questions, and to further clarify the differences between them and the [WC]-PG 1159 stars, we clearly need to determine their physical parameters and chemical abundances. Non LTE expanding atmosphere models are being computed with this purpose by our group with the CMFGEN code of Hillier \\& Miller (1998). In this way, their position in the HR diagram could be determined, a more efficient comparison to [WR], [WC]-PG 1159 and PG 1159 stars could be made, and their evolutionary status could be better determined."
    },
    "0710/0710.2894_arXiv.txt": {
        "abstract": "In the ongoing HATNet survey we have detected a giant planet, with radius $\\hatcurr\\,\\rjuplong$ and mass $\\hatcurm\\,\\mjuplong$, transiting the bright ($V=10.5$) star \\gschatcur{}. The planet is in a circular orbit with period $\\hatcurP$\\,days and mid-transit epoch \\hatcurT\\,(HJD). The parent star is a late F star with mass $1.29 \\pm 0.06\\,\\msun$, radius $1.46 \\pm 0.06\\,\\rsun$, $\\teff \\sim 6570\\pm80\\,\\mathrm{K}$, $\\mathrm{[Fe/H]} = -0.13 \\pm 0.08$ and age $\\sim 2.3^{+0.5}_{-0.7}\\,\\mathrm{Gy}$. With this radius and mass, \\hatcurb{} has somewhat larger radius than theoretically expected. We describe the observations and their analysis to determine physical properties of the \\hatcur{} system, and briefly discuss some implications of this finding. ",
        "introduction": "\\label{sec:introduction} The detection of transiting exoplanets is very important to exoplanet research because of the information about both planetary radius and mass that comes from photometric transit light curves combined with follow-up radial velocity observations.  The transiting exoplanets known as of this writing span a wide range in the physical parameter space of planetary mass, radius, orbital period, semi-major axis, eccentricity; and parent star parameters, including mass, radius, effective temperature, metallicity, and age. Filling out their distribution in this multidimensional space is certain to give us important information on the origin and evolution of exoplanetary systems. Here we report on the discovery by HATNet of its sixth transiting planet, \\hatcurb{}, an inflated Jupiter-mass gas giant in an essentially circular orbit about an F dwarf star with slightly sub-solar metallicity. ",
        "conclusions": "\\hatcurb{}, with radius $\\hatcurrshort\\,\\rjup$, is similar in size to five low density ``inflated'' planets tabulated by \\citet{kovacs07} (i.e.~WASP-1b, HAT-P-4b, HD~209458b, TrES-4, and HAT-P-1b).  However, its mass of $\\hatcurmshort\\,\\mjup$ is greater than the mass of any of these, and hence it has a larger mean density and surface gravity. For a planet of its mass, age of 2.3 Gy, and stellar flux $F_p$ at the planet given by $F_p = \\lstar / (4 \\pi a^2)$, models of \\cite{Burrows:07} predict a radius of about 1.21 $\\rjup$, assuming that the planet has no heavy-element core. The metallicity of \\hatcur{}, $\\mathrm{[Fe/H]}=-0.13 \\pm 0.08$, is among the smallest of known transiting planet host stars.  If we assume that the bulk composition of \\hatcurb{} tracks the metallicity of its host star, the size of its heavy-element core should be small, but not vanishingly so.  Thus the predicted radius from \\cite{Burrows:07} is comparable to, but perhaps slightly higher, than one might expect for \\hatcurb{}. However, the actual radius found here, $\\rpl = \\hatcurrlong$, lies above the predicted value by about $2\\sigma$, so it appears to be somewhat inflated relative to that model.  \\cite{Hansen:07} proposed that hot jupiters can be placed into two classes based on their equilibrium temperature and Safronov number $\\Theta$, where $\\Theta \\equiv (a/\\rpl) \\times (\\mpl/\\mstar)$ is the ratio of the escape velocity from the surface of the planet to the orbital velocity. When Safronov number is plotted versus equilibrium temperature, transiting hot jupiters seem to fall into two groups, with an absence of objects between. However, the Safronov number for \\hatcurb{} is $0.064 \\pm 0.004$; along with HAT-P-5b \\citep{Bakos:07b} with Safronov number $0.059\\pm0.005$, these two planets appear to fall between the two groups in such a plot.  Hansen \\& Barman also noted a difference in the relation between planet mass and equilibrium temperature for planets of the two classes, but HAT-P-6b and HAT-P-5b appear to fall between the two classes in this respect as well. It would seem that discovery and characterization of a large number of additional transiting exoplanets may be necessary to establish unambiguously whether there is a bi-modal distribution of hot jupiter planets according to their Safronov number."
    },
    "0710/0710.4776_arXiv.txt": {
        "abstract": "The chromosphere of the quiet Sun is a highly intermittent and dynamic phenomenon. Three-dimensional radiation (magneto-)hydrodynamic simulations exhibit a mesh-like pattern of hot shock fronts and cool expanding post-shock regions in the sub-canopy part of the inter-network.  This domain might be called ``fluctosphere''. The pattern is produced by propagating shock waves, which are excited at the top of the convection zone and in the photospheric overshoot layer. New high-resolution observations reveal a ubiquitous small-scale pattern of bright structures and dark regions in-between. Although it qualitatively resembles the picture seen in models, more observations -- e.g. with the future ALMA --  are needed for thorough comparisons with present and future models. Quantitative comparisons demand for synthetic intensity maps and spectra for the three-dimensional (magneto-)hydrodynamic simulations. The necessary radiative transfer calculations, which have to take into account deviations from local thermodynamic equilibrium, are computationally very involved so that no reliable results have been produced so far. Until this task becomes feasible, we have to rely on careful qualitative comparisons of simulations and observations. Here we discuss what effects have to be considered for such a comparison. Nevertheless we are now on the verge of assembling a comprehensive picture of the solar chromosphere in inter-network regions as dynamic interplay of shock waves and structuring and guiding magnetic fields. ",
        "introduction": "The chromosphere of the quiet Sun -- a story full of misunderstandings. Apart from the ongoing controversy concerning the heating mechanism \\cite[(e.g., Fossum \\& Carlsson 2005)]{2005Natur.435..919F)}, many details of the small-scale structure of the chromosphere of inter-network regions are still unknown. Already the term ``chromosphere''\\footnote{\\cite[Rutten (2007, and references therein)]{2007ASPC..368...27R} uses the term ``clapotisphere'' for the shock-dominated subcanopy domain in inter-network regions, whereas his ``chromosphere'' refers to the fibrilar structure  visible in H$\\alpha$ only. As ``clapotisphere'' stands for standing waves, we here introduce the term ``fluctosphere'' instead (fluctus = latin for ``wave'').} is a frequent source of misunderstandings . Certainly the large variety of phenomena observed \\cite[(see, e.g., Judge 2006; Rutten 2006, 2007)]{2006ASPC..354..259J,2006ASPC..354..276R,2007ASPC..368...27R} created a complex puzzle and sometimes apparent contradictions. For instance, the observed UV emission implies high temperatures, whereas the existence of carbon monoxide lines point at much cooler gas \\cite[(Ayres 2002)]{ayres02}. New high-resolution observations -- as reported here -- show a highly dynamic and intermittent pattern that cannot be explained with the classical semi-empirical models by \\cite[Vernazza \\etal\\   (1981, VAL)]{val81} and \\cite[Fontenla \\etal\\   (1993, FAL)]{fal93}. Rather a time-dependent three-dimensional model is mandatory. A self-consistent model that can fulfill all observational constraints would be most valuable for summarising the many faces of the chromosphere, indicating and understanding the most relevant processes. Chromospheric heating is a central issue as it has important implications for the atmospheres of other stellar types.  Here we report on some advances of detailed radiation magnetohydrodynamic simulations in comparison with new high-resolution observations. Some crucial aspects of such comparisons -- which often result in misunderstandings -- are discussed. ",
        "conclusions": "State-of-the-art numerical simulations exhibit a highly dynamic  chromosphere, which is characterised by propagating and interacting shock waves with co-existing hot and cool regions. New observations, as presented here, now have a sufficiently high spatial, temporal {\\em and} spectral resolution to resolve a pattern of bright structures and dark regions. As synthetic intensity images for the cores of the prominent Ca\\,II lines are still missing, these observations can only be compared to the simulations on a qualitative basis at the moment. Nevertheless they clearly support the picture of the chromosphere as highly dynamic and intermittent phenomenon like it was already implied by the pioneering simulations by \\cite[Carlsson \\& Stein (1995, 1998)]{carlsson95, 1998IAUS..185..435C}. Even \\cite[Kalkofen (2004)]{2004IAUS..219..115K} stated that ``[d]etailed observations show the chromosphere to be highly dynamic''. The model atmospheres by VAL and FAL, although certainly very elaborate, suffer from the basic assumption of a one-dimensional stratified static atmosphere. This assumption is clearly questioned by recent high-resolution observations. Strong intensity fluctuations are now observed although one still should need be careful with deriving statements concerning the gas temperature. Frequently it is argued that the amplitudes and minimum values of gas temperature in the model atmospheres are not observed and that they are wrong as they do not agree with the VAL models \\cite[(cf. Kalkofen 2003b)]{2003SPD....34.1101K}. These arguments obviously do not hold any longer in view of new observational results. VAL-type atmospheres should thus be considered as qualitative averages, which could at best be interpreted as variations on large spatial scales. Instead of aiming at an agreement with VAL models, modern 3D radiation (magneto-)hydrodynamical simulations must be directly compared to observations. Certainly the quantitative values of the temperature fluctuations in these models (apart from the low to middle photosphere) still suffer from the simplified treatment of radiative transfer and resulting uncertainties in the energy balance, which are, however, a necessary compromise in order to keep the computations tractable.  The next steps towards realistic chromosphere models requires more work on the modelling of non-equilibrium effects, which are important for the energy balance and thus the temperature amplitudes in the chromosphere. An efficient non-LTE radiative transfer scheme is a major goal.  Observational emphasis should be given to (i)~continued aiming at a combination of high spatial \\textit{and} temporal \\textit{and} spectral resolution, as they are crucial for a meaningful comparison and interpretation, and (ii) the development of new diagnostics as, e.g., the (sub-)millimetre continua. Especially the upcoming ALMA might allow a tomography of the solar atmosphere, finally revealing details of its three-dimensional structure."
    },
    "0710/0710.3683_arXiv.txt": {
        "abstract": "We analyze the rotation curves of 10 spiral galaxies with a newtonian potential corrected with an extra logarithmic term, using a disc modelization for the spiral galaxies. There is a new constant associated with the extra term in the potential. The rotation curve of the chosen sample of spiral galaxies is well reproduced for a given range of the new constant. It is argued that this correction can have its origin from string configurations. The compatibility of this correction with local physics is discussed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.3010_arXiv.txt": {
        "abstract": "Recent observations of the Mira AB binary system have revealed a surrounding arc-like structure and a stream of material stretching 2 degrees away in opposition to the arc. The alignment of the proper motion vector and the arc-like structure shows the structures to be a bow shock and accompanying tail. We have successfully hydrodynamically modelled the bow shock and tail as the interaction between the asymptotic giant branch (AGB) wind launched from Mira A and the surrounding interstellar medium. Our simulations show that the wake behind the bow shock is turbulent: this forms periodic density variations in the tail similar to those observed. We investigate the possiblity of mass-loss variations, but find that these have limited effect on the tail structure. The tail is estimated to be approximately 450\\,000 years old, and is moving with a velocity close to that of Mira itself.  We suggest that the duration of the high mass-loss phase on the AGB may have been underestimated. Finally, both the tail curvature and the rebrightening at large distance can be qualitatively understood if Mira recently entered the Local Bubble. This is estimated to have occured 17\\,pc downstream from its current location. ",
        "introduction": "In the Mira binary system, Mira A is an evolved star which is undergoing a period of enhanced mass-loss as it moves along the asymptotic giant branch (AGB) on route to becoming a white dwarf. The companion star, previously classified as a white dwarf, now has a less clear classification \\citep{karovska05,ireland07} but is less luminous and any stellar outflow is comparably insignificant to Mira A in terms of mass-flux and energetics.  New UV observations of the Mira system \\citep{martin07} have revealed a comet-like tail extending 2 degrees away to the North and an arc-like structure in the South.  \\cite{martin07} postulated that these features are a bow shock and a ram-pressure-stripped tail stretching away in opposition to the bow shock caused by motion through the interstellar medium (ISM). Their postulation is consistent with Mira's proper motion \\citep{turon93} of 225.8 milli-arcseonds per year in the direction 187.1 degrees East of North (corrected for solar motion). Further, they note that at the revised Hipparcos-based distance \\citep{knapp03} of 107 pc, the large space velocity of 130 \\kms, calculated from the proper motion and the radial velocity \\citep{evans67} of 63 \\kms, is further consistent with the bow shock structure.  In this paper, we present hydrodynamical modelling of the Mira system and discuss the implications with respect to postulation of \\cite{martin07}.  We are aiming to fit the position of the bow shock ahead of the central star, the width across the central star, the undulating density profile along the tail, the length of the tail and the ring-like structure one third the way down the tail.  We show the GALEX observation in Figure 1. ",
        "conclusions": "Our models for Mira's evolution suggest that the tail traces half a million years of mass-loss history and ISM interaction. An instability at the bow shock is the cause of the fluctuations seen in the tail. Mass-loss variations of Mira A have much less effect. The curvature in the tail is caused by the differential velocity between Mira and its tail leading to different Galactic orbits. We attribute a gap in the tail to Mira entering the local bubble, 17 pc downstream from its current location. The time to reestablish the bow shock directly leads to the gap in the tail. A ring-like structure is attributed to a vortex shed into the tail just before the bow shock reached its equilibrium position. The tail therefore traces not only the history of Mira itself, but also the structure of the ISM along its path. It's a wonderful prospect."
    },
    "0710/0710.5636_arXiv.txt": {
        "abstract": "We present a parametric strong lensing model of the cluster Abell 2218 based on HST ACS data.  We constrain the lens model using 8 previously known multiply imaged systems, of which 7 have spectroscopically confirmed redshifts.   In addition, we propose five candidate multiply imaged systems and estimate their redshifts using our mass model. The model parameters are optimized in the source plane by a bayesian Monte Carlo Markov Chain as implemented in the the publicly available software Lenstool.  We find rms$_s=0\\farcs12$ for the scatter  of the sources in the source plane, which translates into rms$_i=1\\farcs49$ between the predicted and measured image positions in the image plane. We find that the projected mass distribution of Abell 2218 is bimodal, which is supported by an analysis of the light distribution.  We also find evidence for two structures in velocity space, separated by $\\sim 1000$~km~$s^{-1}$, corresponding to the two large scale dark matter clumps.  We find that the lensing constraints can not be well reproduced using only dark matter halos associated with the cluster galaxies, but that the dark matter is required to be smoothly distributed in large scale halos.   At $100\\arcsec$ ($291$~kpc) the enclosed projected mass is $3.8\\times10^{14}$~M$_\\sun$.  At that radius, the large scale halos contribute $\\sim85\\%$ of the mass, the brightest central galaxy (BCG) contributes  $\\sim9\\%$ while the remaining $\\sim6\\%$ come from the other cluster galaxies.  We find that the model is not very sensitive to the fainter (and therefore by assumption less massive) galaxy sized halos, unless they locally perturb a given multiply imaged system.  Therefore, dark galaxy sized substructure can be reliably constrained only if it locally perturbs one of the systems.  The massive BCG and galaxies which locally perturb a multiply imaged system are reliably detected in the mass reconstruction.  In an appendix we give a self-contained description of the parametric profile we use, the dual pseudo isothermal elliptical mass distribution (dPIE).  This profile is a two component pseudo isothermal mass distribution (PIEMD) with both a core radii and a scale radii. ",
        "introduction": "Dark matter dominates over baryonic matter in the universe, but its nature is not known.  The study of the inner parts of dark matter halos can give insight into the nature of the dark matter, as the steepness of the profile is correlated with the interaction between the dark matter itself and with the baryonic matter.  According to $\\Lambda$CDM simulations, the mass distribution of galaxy clusters should be dominated by their dark matter halos.  Gravitational lensing, which is sensitive to the total matter distribution, visible or dark, is ideal for studying the mass distribution of clusters.  Strong lensing features, consisting of multiply imaged and strongly distorted background sources, provide constraints on the inner parts of the cluster, while weak lensing features, consisting of weakly distorted singly imaged background sources, provide constraints on the outer slope of the surface mass profile \\citep[see e.g.,][]{smail1995a,smail1995b,kneib1996, smail1997,abdelsalam1998,natarajan2002b, bradac2006,gavazzi2007,limousin2007}.  Lensing can therefore provide unique information about the total mass distribution of clusters, from the inner to the outer parts.  In addition, lensing can in principle be used to deduce various cosmological parameters (e.g. $H_0$, $\\Omega_\\Lambda$, $\\Omega_m$).   This has already been extensively applied to lensing on galaxy scales \\citep[see e.g.,][]{schechter1997,koopmans2003} and to a smaller degree for lensing on cluster scales \\citep[see e.g.,][]{soucail2004, meneghetti2005}, where the accuracy of the mass map is a limiting factor.  The accuracy of the mass map is strongly dependent on the number of multiply imaged systems used to constrain it. Therefore, to construct a robust model of the dark matter distribution, accurate enough for cosmography and for using the cluster as a gravitational telescope, it is important to include as many spectroscopically confirmed multiply imaged systems as possible \\citep[see e.g.,][]{ellis2001}.  Abell 2218 is one of the richest clusters in the Abell galaxy cluster catalog \\citep{abell1958, abell1989} and has been successfully exploited as a gravitational lens. A parametric lens model has previously been constructed by \\citet{kneib1995,kneib1996} (using 1 and 2 spectroscopically confirmed systems respectively) and by \\citet{natarajan2002b, natarajan2007} (using 4 and 5 spectroscopically confirmed systems) building on the model of \\citet{kneib1996} and including weak lensing constraints from HST WFPC2 data.  A non-parametric model was constructed by \\citet{abdelsalam1998} using three spectroscopically confirmed multiply imaged systems.  In all of these models, a bimodal mass distribution was required to explain the image configurations (i.e. the models include two large scale dark matter clumps), but the number of constraints were not sufficient to accurately constrain the second large scale dark matter clump.  Abell 2218 has also been used as a gravitational telescope, with  \\citet{ellis2001} discovering a source at $z=5.6$ and \\citet{kneib2004} discovering an even more distant source at $z\\sim6.7$, later confirmed by \\citet{egami2005} using Spitzer data.  \\citet{soucail2004} estimated cosmological parameters based on a lensing model of Abell 2218 using 4 multiply imaged systems.  The latest published lens model of Abell 2218 is by \\citet{smith2005}, who incorporated four multiply imaged systems and weak lensing constraints, using the WFPC2 data.  Although the number of constraints has increased from the initial models, all previous models have assumed that the location of the second dark matter clump coincided with the brightest galaxy in the South East, due to a lack of constraints in its vicinity.  The motivation for revisiting the modeling of Abell 2218 comes from the new ACS data which have not been used before for modeling this cluster and are superior in both resolution, sensitivity and field of view to the previous WFPC2 data set.  These new high quality data allow us to identify several subcomponents in previously known multiple images, thus adding more constraints, and in one case, two more multiple images of a known system. In addition, we have measured a spectroscopic redshift for an arc around the second dark matter clump, which our model predicts to be singly imaged. We also have several new candidate multiply imaged systems, which we add as constraints and estimate their redshift with the model.  In total we have 7 multiply imaged systems with measured spectroscopic redshift and 6 multiply imaged systems without spectroscopic data (of which 5 are new candidate systems).  Finally, the lensing code Lenstool, has undergone significant improvements from previous models, with the incorporation of a Monte Carlo Markov Chain (MCMC), which enables us to not only find the best model in the lowest $\\chi^2$ sense, but the most likely model as measured by its Evidence \\citep{jullo2007}.  The MCMC also allows for a reliable estimate of the uncertainties in the derived model parameters.  The paper is organized as follows:  In Section~\\ref{sec:data}, we give an overview of the data used in this paper.  We compile a list of all currently known and new multiply imaged systems in Abell 2218 and discuss the reliability of the redshift estimate of each one in Section~\\ref{sec:systems}.  The methodology of the strong lensing modeling is described in Section~\\ref{sec:modeling}.    In Section~\\ref{sec:analysis} we present the results of our lensing analysis, and compare them to previous models.  In Section~\\ref{sec:degeneracy} we discuss degeneracies in the modeling.  In Section~\\ref{sec:reliability} we address how reliable our model is, and discuss the smoothness of the dark matter distribution. In Section~\\ref{sec:bimodal} we interpret the bimodality of our model, along with X-ray measurements and an analysis of the distribution of cluster members in velocity space.  We summarize our main conclusions in Section~\\ref{sec:conclusions}.   Throughout the paper, we adopt a flat $\\Lambda$-dominated Universe with $\\Omega_\\Lambda=0.7$, $\\Omega_m=0.3$ and $H_{0}=70\\ \\mathrm{km\\,s}^{-1}\\,\\mathrm{Mpc}^{-1}$.  Following \\citet{smith2005} we place the cluster at $z=0.171$.  At this redshift $1\\arcsec$ corresponds to $2.91$~kpc in the given cosmology. ",
        "conclusions": ""
    },
    "0710/0710.5770_arXiv.txt": {
        "abstract": "We investigate the distribution of massive black holes (MBHs) in the Virgo cluster. Observations suggest that AGN activity is widespread in massive galaxies ($M_*\\simgt 10^{10}\\msun$), while at lower galaxy masses star clusters are more abundant, which might imply a limited presence of central black holes in these galaxy-mass regimes. We explore if this possible threshold in  MBH hosting, is linked to {\\it nature}, {\\it nurture}, or a mixture of both. The {\\it nature} scenario arises naturally in hierarchical cosmologies, as  MBH formation mechanisms typically are efficient in biased systems, which would later evolve into massive galaxies. {\\it Nurture}, in the guise of  MBH ejections following  MBH mergers, provides an additional mechanism that is more effective for low mass, satellite galaxies. The combination of inefficient formation, and lower retention of  MBHs, leads to the natural explanation of the distribution of compact massive objects in Virgo galaxies.  If MBHs arrive to the correlation with the host mass and velocity dispersion during merger-triggered accretion episodes, sustained tidal stripping of the host galaxies creates a population of MBHs which lie above the expected scaling between the holes and their host mass, suggesting a possible environmental dependence. ",
        "introduction": "The nearby Virgo cluster is a perfect laboratory to investigate the evolution of galaxies in a dense environment. Recently, observations of a large sample of galaxies suggested that the properties of nuclei, either quiescent or active, in Virgo galaxies are strongly mass-dependent. This latter finding is in very good agreement with the general trend found also in the SDSS, where very few AGN are found in galaxies with stellar mass $M_*< 10^{10}\\msun$\\citep{Kauffmann2003, Kewley2006}.  Decarli et al. (2007) analyzed nuclear activity in late type galaxies in the Virgo cluster. They conclude, quite remarkably, that at galaxy mass\\footnote{Decarli et al. (2007) measure the dynamical mass of the galaxy within the optical radius, determined at the $25^{th}$ mag arcsec$^{-2}$ isophote in the B band. This is an upper limit to the stellar mass of galaxies, but a lower limit to the total baryonic mass.} $M_{gal}> 10^{10.5}\\msun$ the AGN fraction is unity. As a central black hole is a necessary condition for AGN activity, we conclude that the black hole occupation fraction (BHOF)  must be unity as well. \\cite{Cote2006}, Wehner \\& Harris (2006) and Ferrarese et al. (2006) find that, below $M_*\\sim 10^{10}\\msun$, Virgo galaxies exhibit nuclear star clusters, whose mass scales with $M_*$  in the same fashion as those of the massive black holes detected in brighter galaxies \\citep{Magorrian1998, MarconiHunt2003, Haring2004}. Although the existence of a nuclear star cluster does not rule out a small, hidden  MBH, it is suggestive that Ferrarese et al. (2006) conclude that ``bright galaxies often, and perhaps always, contain supermassive black holes but not stellar nuclei. As one moves to fainter galaxies, nuclei become the dominant feature while  MBHs might become less common and perhaps disappear entirely at the faint end.\"  There are three interlaced sides of the intriguing story which appears to link stellar nuclei and MBHs: 1) understand if (and why) MBHs populate preferentially bright galaxies; 2) understand why stellar nuclei populate preferentially faint galaxies and 3) understand why the ratio of nuclear to galaxy mass is identical to the ratio of MBH to galaxy mass.  In this paper we address the first issue. A possible hint to explain the predominance of star clusters in small galaxies may come from comparing the dynamical time scale and the fragmentation time scale of the infalling gas. Detailed calculations are necessary to test this hypothesis, but it can be argued that in shallow potentials gas could fragment before reaching the center of the galaxy. This is even more suggestive in the case of merger-induced gaseous infall. Regarding the third issue, \\cite{Emsellem2007} for instance show that if galaxies have Sersic profiles, radial compressive forces trigger the collapse of gas in the central regions, and the mass of the nuclear cluster that forms is about 0.1\\% of the mass of the host galaxy. So, nuclear cluster formation and MBH feedback might produce the same scaling relation with galaxy mass, but it is unclear whether this is a coincidence or the result of a single, unexplored process.  Volonteri et al. (2007) investigate the overall distribution of  MBHs in galaxies, as a function of the host velocity dispersion, $\\sigma_*$ , as traced by the host halo (cfr. Ferrarese 2002), and find that the efficiency of  MBH formation decreases with halo mass, and isolated dwarf galaxies are most likely to be devoid of a central  MBH. This is a common feature of  MBH formation models which invoke gas-dynamical processes in the high redshift Universe, either via direct collapse \\citep[e.g.,]{haehnelt1993,LoebRasio1994,Eisenstein1995,BrommLoeb2003,BegelmanVolonteriRees2006,LN2006}, or via an intermediate stage of metal free Population III star \\citep[e.g.,]{CBA84,MadauRees2001,VHM}. The common feature is the need for deep potential wells at cosmic epochs when the protogalaxy population was dominated by minihalos ($M_h<10^7\\msun$). \\\\ \\indent A second physical phenomenon strengthens the likelihood that dwarf galaxies, even if seeded with a  MBH at early times, are now lacking one: the gravitational recoil. That is, the general-relativistic effect which imparts velocity to the center of mass of a merging  MBH binary, when the  MBHs plunge into each other's last stable orbit by emission of gravitational radiation. Gravitational waves carry, in general,  a non-zero net linear momentum, which establishes a preferential direction for the propagation of the waves. As a consequence, the center of mass of the binary recoils in the opposite direction \\citep{redmountrees}, possibly causing the ejection of MBHs from the potential wells of their host galaxies  \\citep[e.g.,][]{Madau2004,MadauQuataert2004, Merrittetal2004, VolonteriRees2006, Haiman2004, SchnittmanBuonanno2007}. The progress of numerical relativity is now leading to a convergence in the estimates of the recoil. Schwarzschild, i.e., non-spinning, black holes \\citep[e.g.,][]{Bakeretal2006} are expected to recoil with velocities below 200 $\\rm{km\\,s^{-1}}$, and a similar range is expected for black holes with low spins, or with spins (anti-)aligned with the orbital axis. However, when the spin vectors have opposite directions and are in the orbital plane, the recoil velocity can be as large as a few thousands $\\rm{km\\,s^{-1}}$ \\citep{Campanelli2007b,Gonzalez2007,Campanelli2007}.  Schnittman (2007) shows that the hierarchical nature of galaxy assembly implies that ejections do not lead to void nuclei, even in the high-recoil limit. This is weaved into the very nature of galaxy assembly, via a series of mergers, in a pattern typically dubbed ``merger tree\". The botanical essence of the tree implies that many branches converge into a central trunk, so that in every generation of the tree, the number of galaxies decreases, while the fraction of  MBHs can increase, even if ejections operate. This is especially true for  ``central galaxies\", which represents the trunk of the merger tree.  The situation is different for satellites in a galaxy cluster: they correspond to loose branches that do not merge into the main trunk. Satellites are indeed galaxies that enter the dark matter halo of the cluster, but do not merge with the central galaxy.  In this paper we develop simple models that describe the dynamical evolution of satellite galaxies in a cluster, focusing on the fate of their nuclei. We estimate the influence of  MBH formation mechanisms, and of  MBH ejections for shaping the BHOF in the Virgo cluster today. ",
        "conclusions": "We explored the influence of formation epoch of  MBHs, bias of their hosts, and  MBHs dynamical evolution on the occupation distribution of  MBHs in galaxy clusters. The possible decrease of the occupation fraction at low galaxy mass proposed by \\cite{Wehner2006} and \\cite{Ferrareseetal2006} is not a surprising result and follows naturally from the evolution of the  MBH population of galaxies in clusters. \\begin{itemize} \\item When the formation mechanisms of  MBH seeds are taken into consideration, implying formation of  MBHs in massive high redshift halos, the BHOF in cluster galaxies in an increasing function of galaxy mass, in line with observations of the AGN fraction in Virgo. \\item The exact mass threshold above which BHOF=1 depends on the details of the formation mechanism (host masses and redshift of formation, ``nature\") and on the dynamical evolution, including  MBH spin magnitude (``nurture\"). \\item The repercussions of ``nurture\" are magnified in cluster galaxies. If a satellite galaxy looses its  MBH due to a dynamical interaction, it has a negligible chance of capturing a new one following a subsequent galaxy merger, except for the central galaxy in the cluster. \\item We also predict that if  MBHs co-evolve with galaxies during galaxy mergers, and satellite galaxies experience tidal stripping during their orbital evolution, then a population of galaxies where the  MBH mass lies above the standard $M_{BH}-M_{gal}$ is expected. Two such examples could be NGC~4486B and NGC~4342. \\end{itemize}  We emphasize that our models do not in any way imply that  MBHs and nuclear clusters are mutually exclusive. We however predict that the occupation fraction decreases with galaxy mass. If, as suggested by Wehner \\& Harris (2006), lacking a MBH is a necessary condition for nuclear cluster formation, radio or hard X-ray observations will inform us of the complementary or mutually exclusive essence of  MBHs and nuclear clusters. \\cite{Gallo2007} recently observed 32 galaxies in Virgo with the Chandra X-ray Observatory, and found, in agreement with \\cite{Decarli2007} that nuclear X-ray activity  increases with the mass of the host galaxy. Intriguingly, at least in one case, VCC1178, a nuclear X-ray source is detected jointly with a central star cluster. \\cite{Seth2006} also report that a small AGN might be hosted within the core of the nuclear star cluster in NGC 4206."
    },
    "0710/0710.2199_arXiv.txt": {
        "abstract": "We report on active region  EUV dynamic events observed  simultaneously at high-cadence with SUMER/SoHO and TRACE.  Although the features appear in the TRACE~Fe~{\\sc ix/x}~171~\\AA\\ images  as  jets seen in projection on the solar disk, the SUMER spectral line profiles suggest that the plasma has been driven along a curved large scale magnetic structure, a pre-existing loop.  The SUMER observations were carried out in spectral lines covering a large temperature range from 10$^4$~K to 10$^6$~K. The spectral analysis revealed that a sudden heating from an energy deposition is followed by a high velocity plasma flow. The Doppler velocities were found to be in the range from 90 to 160 \\kms. The heating process has a duration which is below the SUMER exposure time of 25 {\\rm s} while the lifetime of the events is from 5 to 15 {\\rm min}. The additional check on soft X-ray Yohkoh  images shows that the features most probably reach 3 MK (X-ray) temperatures. The spectroscopic analysis showed no existence of cold material during the events. ",
        "introduction": "A large variety of jet-like phenomena are often observed in the solar atmosphere such as surges, spicules, sprays, Extreme-UltraViolet (EUV) and X-ray jets. X-ray jets (Shibata \\etal\\ 1992) were first identified in data obtained with the Soft X-ray Telescope (SXT) on Yohkoh (Tsuneta \\etal\\ 1991). They represent X-ray enhancements with an apparent collimated motion and were found to have a typical size of 5~$\\times$~10$^3$ -- 4~$\\times$~10$^5$~{\\rm km} and an apparent velocity of 30 to 300 \\kms. Their kinetic energy is estimated  to be 10$^{25}$ -- 10$^{28}$~{\\rm ergs}. Most of the jets were associated with small flares in large X-ray bright points or active regions. Shimojo \\& Shibata (2000) derived the physical parameters of X-ray jets and found temperatures from 3 to 8 MK (determined by using Yohkoh filter ratios) and densities of 0.7 -- 4.0 $\\times$ 10$^9$~{\\rm cm}$^{-3}$. It is strongly believed that they are produced by magnetic reconnection and represent the evaporation flow resulting from the reconnection heating.  EUV jets were studied by Brekke (1999) in off-limb data from the Coronal Diagnostics Spectrometer (CDS) and the Extreme-ultraviolet Imaging Telescope (EIT). From the CDS data it was found that the jet was emitting only at transition region temperatures showing Doppler shifts in the O~{\\sc v} 629.73 \\AA\\ line up to -75 \\kms. The event was also seen in the EIT Fe~{\\sc xii}~195~\\AA\\ passband propagating with an apparent velocity of 180 \\kms. The plasma seemed to be ejected along a large looped magnetic structure. Jets were also analysed in  on-disk data from the  Transition Region And Coronal Explorer (TRACE) taken in the 171~\\AA\\ and 1216~\\AA\\ passbands by Alexander \\& Fletcher (2000). In the 171~\\AA\\ channel the ejected plasma was seen both in emission and absorption which suggests that simultaneously highly collimated hot and cold  material was ejected along the magnetic field lines.  An EUV jet from a new emerging active region (a large Bright Point) was analysed in simultaneous TRACE, EIT and CDS data by Lin \\etal\\ (2006). The authors found the plasma jet to emit in a wide temperature range from 10\\,000~K (He~{\\sc i}) to 2.5~MK (Fe~{\\sc xvi}, the upper temperature limit  of their observations).  H$_\\alpha$ surges are often associated with EUV and X-ray emissions showing the co-existence of cool (H$_\\alpha$) and hot ejections of plasma (Jiang \\etal\\ 2007 and references therein). Only recently, however, have the spatial and temporal relation of these emissions been studied in detail (Jian \\etal\\ 2007) during surge events in a plage area of an active region. The authors first observed the bright structures in TRACE 171 \\AA\\ followed by the cooler  H$_\\alpha$ jet which they interpret as cooling of the hot plasma with a cooling time lasting about 6 -- 15 {\\rm min}. ",
        "conclusions": "This letter presents, to our knowledge for the first time, EUV transient features in an active region identified and analysed in  on-disk SUMER  data and simultaneously obtained TRACE images. These instruments provide data at highest existing (1.5\\arcsec and 1\\arcsec) spatial and 2 \\kms\\ spectral resolution (SUMER). Additionally, the combination of spectrometer and imager data obtained at high cadence (25~{\\rm s} and below) permitted the temporal and spatial evolution, velocities and especially temperatures of EUV active region transients to be derived with the highest possible precision. Three dynamic events were studied in spectral  lines covering a temperature range from 10$^4$ to 10$^6$~K. All three features  showed strong red-shifted emission  in the  O~{\\sc v}~629~\\AA\\  and Mg~{\\sc x}~625~\\AA\\ lines  suggesting high velocity flows which propagate in a direction away from the observer, i.e. towards the solar surface. Considering the magnetic fields structure  of the active region field by the loops seen in TRACE~171 \\AA, we suggest that the features although appearing with a jet-like structure in TRACE~171 \\AA\\ images may rather represent a high velocity flow driven along a curved magnetic field, most probably a pre-existing loop. No signature of the events was found at chromospheric temperatures. Both EIT/Fe~{\\sc xii}~195~\\AA\\ and Yohkoh/SXT showed brightenings in a pixel row indicating the presence of a 1 to 3 MK plasma during the transients.  The lower resolution of these instuments ($\\approx$6\\arcsec) in comparison to TRACE (1\\arcsec) do not permit the events to be identified as jets. The response in the transition region lines is delayed by  around 2~{\\rm min} in respect to the coronal line suggesting cooling of the events. In the future hydrodynamic numerical simulations (see for detail Doyle \\etal\\ 2006b) will be performed  with the results converted into observable quantities to be then compared with the present data. We believe that the capabilities of  the Hinode mission will bring  a better understanding on these features and, more important, the physical mechanism behind them."
    },
    "0710/0710.0964_arXiv.txt": {
        "abstract": "We use a large suite of carefully controlled full hydrodynamic simulations to study the ram pressure stripping of the hot gaseous halos of galaxies as they fall into massive groups and clusters.  The sensitivity of the results to the orbit, total galaxy mass, and galaxy structural properties is explored.  For typical structural and orbital parameters, we find that $\\sim 30\\%$ of the initial hot galactic halo gas can remain in place after 10 Gyr. We propose a physically simple analytic model that describes the stripping seen in the simulations remarkably well.  The model is analogous to the original formulation of Gunn \\& Gott (1972), except that it is appropriate for the case of a spherical (hot) gas distribution (as opposed to a face-on cold disk) and takes into account that stripping is not instantaneous but occurs on a characteristic timescale.  The model reproduces the results of the simulations to within $\\approx 10\\%$ at almost all times for all the orbits, mass ratios, and galaxy structural properties we have explored.  The one exception involves unlikely systems where the orbit of the galaxy is highly non-radial and its mass exceeds about 10\\% of the group or cluster into which it is falling (in which case the model under-predicts the stripping following pericentric passage). The proposed model has several interesting applications, including modelling the ram pressure stripping of both observed and cosmologically-simulated galaxies and as a way to improve current semi-analytic models of galaxy formation.  One immediate consequence is that the colours and morphologies of satellite galaxies in groups and clusters will differ significantly from those predicted with the standard assumption of complete stripping of the hot coronae. ",
        "introduction": "There are marked differences in the observed properties of the field and cluster galaxy populations.  Perhaps the best known difference is the larger fraction of galaxies that are ellipticals or S$0$s (and the correspondingly lower spiral fraction) in clusters relative to the field (e.g., Dressler 1980; Goto et al.\\ 2003).  Not only are the morphologies of cluster galaxies different from those of field galaxies, but so too are a variety of their other observed properties, including colours (e.g., Balogh et al.\\ 2004; Hogg et al.\\ 2004), star forming properties (e.g., Poggianti et al. 1999; Balogh et al.\\ 2000; Gomez et al.\\ 2003), and the distribution and total mass of their gaseous component (e.g., Cayatte et al.\\ 1994; Solanes et al.\\ 2001).  These observed differences indicate that the dense environments of groups and clusters are somehow strongly modifying the properties of galaxies as they fall in.  Uncovering the physical mechanisms that give rise to the observed variation in galaxy properties has been an active topic of research over the past two or three decades (e.g., Dressler 1984; Sarazin 1988). One of the most commonly mentioned processes is ram pressure stripping (Gunn \\& Gott 1972).  Here the gaseous component (which can be composed of both cold atomic/molecular gas and a hot extended component) of the orbiting galaxy is subjected to a wind due to its motion relative to the intracluster medium (ICM).  The gas will be stripped if the wind is sufficiently strong to overcome the gravity of the galaxy. Recently, direct observational evidence for the ram pressure stripping of galaxies in clusters has been provided by long (up to tens of kpc) tails of gas seen to be trailing behind several cluster galaxies (e.g., Sakelliou et al.\\ 2005; Crowl et al.\\ 2005; Vollmer et al.\\ 2005; Sun \\& Vikhlinin 2005; Machacek et al.\\ 2006; Sun et al.\\ 2007a).  Such stripping could at least partially account for the differing properties of cluster and field galaxies.  There have been numerous theoretical studies dedicated to calculating the effects of ram pressure stripping on galaxies using hydrodynamic simulations or semi-analytic models. The vast majority of these studies have focused on the stripping of cold gaseous disks with an emphasis on whether this can account for the observed lower star formation rates (and redder colours) of cluster spirals relative to their field counterparts (e.g., Abadi et al.\\ 1999; Quilis et al.\\ 2000; Vollmer et al.\\ 2001; Okamoto \\& Nagashima 2003; Mayer et al.\\ 2006; Roediger et al.\\ 2006; Hester 2006; Jachym et al.\\ 2007; Roediger \\& Br{\\\"u}ggen 2006; 2007).  However, the stripping of extended {\\it hot} gaseous halos of galaxies is only just beginning to be explored (e.g., Kawata \\& Mulchaey 2007) and has not yet been studied in a detailed and systematic way.  The hot extended component is predicted to exist around most massive galaxies by semi-analytic models and cosmological simulations and is directly observable at X-ray wavelengths in the case of normal ellipticals.  If the hot gaseous halo is completely stripped (as is typically assumed), the only fuel available for star formation is that which resided in the cold component when the galaxy first fell into the cluster. (This process of removing the supply of halo gas is sometimes referred to as ``strangulation'' or ``starvation''.)  However, if the hot halo remains intact for some time it can, via radiative cooling losses, replenish the cold component and potentially significantly prolong star formation.  This, in turn, would affect the colours and morphologies of cluster galaxies (e.g., Larson et al.\\ 1980; Abadi et al.\\ 1999; Benson et al.\\ 2000; Balogh et al.\\ 2000).  Aspects of the stripping of the {\\it hot} gaseous halos of galaxies in clusters have been considered in previous work (e.g., Gisler 1976; Sarazin 1979; Takeda et al.\\ 1984).  Mori \\& Burkert (2000) studied the stripping of dwarf galaxies subject to a constant wind using two-dimensional simulations and found that the relatively shallow potential wells of these systems cannot retain their hot gas component for long. However, these authors did not study more massive systems, such as normal ellipticals and spirals, where stripping of the hot ($\\ga 10^6$ K) halo should be much less efficient due to their higher masses and deeper potential wells.  [Indeed, a recent X-ray survey of massive galaxies in hot clusters by Sun et al.\\ (2007b) has revealed that {\\it most} of the galaxies have detectable hot gaseous halos.]  A few other studies have examined the stripping of more massive systems but not in the context described above.  In particular, they have largely focused on the metal enrichment of the ICM (e.g., Schindler et al.\\ 2005; Kapferer et al.\\ 2007), the X-ray properties of the galaxies (Toniazzo \\& Schindler 2001; Acreman et al.\\ 2003) or the generation of ``cold fronts'' (e.g., Takizawa 2005; Ascasibar \\& Markevitch 2006).  In the present paper, we carry out a detailed study of the ram pressure stripping of the hot gaseous halos of galaxies as they fall into groups and clusters.  This is performed using a large suite of controlled hydrodynamic simulations.  Unlike most previous studies, we use full three-dimensional (3D) simulations in which the galaxies fall into a massive ``live'' group or cluster on realistic orbits.  One important aim is to derive a physically simple and accurate description of the stripping seen in the simulations that can be easily employed in the modelling of observed or cosmologically-simulated galaxies. An additional motivation for deriving such a model is to improve the treatment of ram pressure stripping in semi-analytic models of galaxy formation.  At present, these models typically assume that the hot gaseous halos of galaxies are stripped at the instant they cross the virial radius of the group or cluster.  Clearly, this is not a realistic assumption, especially in the case where the mass of the galaxy is not negligible compared to that of the group or cluster.  Such semi-analytic models tend to predict group and cluster galaxies whose colours are too red compared to observations (e.g., Weinmann et al.\\ 2006; Baldry et al.\\ 2006).  If the ram pressure stripping of the hot gaseous halos of cluster galaxies is not as (maximally) efficient as assumed by these models, the resulting galaxies would be bluer and perhaps in better agreement with observations.  The present paper is structured as follows.  In \\S 2, we present a discussion of our simulation setup and the results of convergence tests that demonstrate the robustness of our findings.  In \\S 3, we first outline a simple analytic model for ram pressure stripping that is based on the original formulation of Gun \\& Gott (1972) but which is appropriate for spherically-symmetrical gas distributions (as opposed to disks). We then compare this model to a wide variety of simulations and demonstrate that it provides an excellent match to the mass loss seen in the simulations.  Finally, in \\S 4, we summarise and discuss our findings. ",
        "conclusions": "Using a suite of carefully controlled 3D hydrodynamic simulations, we have investigated the ram pressure stripping of hot gas in the halos of galaxies as they fall into groups and clusters.  We have proposed a physically simple analytic model that describes the stripping seen in the simulations remarkably well.  This model is analogous to the original formulation of Gunn \\& Gott (1972), except that it is appropriate for the case of a spherical gas distribution (as opposed to a face-on disk) and takes into account that stripping is not instantaneous but occurs on approximately a sound crossing time.  The only pieces of information that the model requires are the initial conditions of the orbiting galaxy (its gas and dark matter profiles), the density profile of the ICM and the orbit [the latter two are needed to calculate $P_{\\rm ram}(t)$].  The model contains two tunable coefficients that are of order unity.  Fixing these coefficients to match the stripping in just one of our idealised uniform medium simulations (see \\S 3.2) leads to excellent agreement with all our other simulations.  With the exception of cases where the mass of the galaxy is greater than about 10\\% of the mass of the group and its orbit is highly non-radial, the analytic model reproduces the mass loss in the simulations to $\\approx 10\\%$ accuracy at all times and for all the orbits, galaxy masses, and galaxy concentrations that we have explored. For cases where the mass of the galaxy exceeds 10\\% of the mass of the group, it will likely be necessary to factor in the effects of tidal stripping and gravitational shock heating, which are neglected by our model.  We re-iterate that the numerical simulations with which our analytic model has been calibrated have been demonstrated to be robust to the adopted resolution and artificial viscosity strength (see \\S 2.2).  Furthermore, as we have demonstrated that KH (and RT) instability stripping is expected to be unimportant, SPH codes should be fully capable of tackling the problem of hot halo gas stripping in galaxies orbiting in groups and clusters.  A direct comparison between the results using the GADGET-2 and FLASH hydrodynamic codes for one of our runs (see Appendix A) confirms this conclusion.  The model we have derived has a number of potentially interesting applications, including modelling observed satellite galaxies and satellite galaxies in cosmological simulations.  One application that we are currently pursuing is the incorporation of our ram pressure stripping model into a semi-analytic model of galaxy formation.  As mentioned in \\S 1, recent observations (Weinmann et al.\\ 2006; Baldry et al.\\ 2006) have revealed that current semi-analytic models predict satellite galaxies whose colours are too red compared to the observed systems.  The implementation of ram pressure stripping in these models is unrealistically efficient since, by assumption, the hot halo of the satellite galaxy is instantly transferred to the more massive system as soon as the satellite galaxy enters the massive system's virial radius.  In reality, the hot gaseous halo of the galaxy will remain intact for a while.  For example, for the most common orbital parameters, we find that between 20\\%-40\\% of the initial hot halo of the galaxy can remain in place even after 10 Gyr of orbiting inside the group or cluster (see Fig.\\ 9; note, however, that the quoted numbers could be sensitive to the adopted hot gas distribution of the galaxy). We note that these predictions are in qualitative agreement with recent {\\it Chandra} X-ray observations of massive galaxies orbiting in hot clusters by Sun et al.\\ (2007b), who find that most of the galaxies have detectable hot gaseous halos. Depending on the efficiency of feedback (e.g., from supernovae winds) in the semi-analytic models, radiative cooling of the remaining hot halo gas will replenish the cold gaseous component at the centre of the galaxy, which in turn will allow star formation to continue for some time.  This will have the effect of making the colour of model satellite galaxies bluer and could resolve the discrepancy between semi-analytic models and observations (Font et al., in prep)."
    },
    "0710/0710.1022.txt": {
        "abstract": "{ We deal with the test of the general relativistic gravitomagnetic Lense-Thirring effect currently being conducted in the Earth's gravitational field with the combined nodes $\\Omega$ of the laser-ranged geodetic satellites LAGEOS and LAGEOS II. One of the most important sources of systematic uncertainty on the orbits of the LAGEOS satellites, with respect to the Lense-Thirring signature, is the bias due to the even zonal harmonic coefficients $J_{\\ell}$ of the multipolar expansion of the Earth's geopotential which account for the departures from sphericity of the terrestrial gravitational potential induced by the centrifugal effects of its diurnal rotation.  The issue addressed here is: are the so far published evaluations of such a systematic error reliable and realistic?  The answer is negative.  Indeed, if the difference $\\Delta J_{\\ell}$ among the even zonals estimated in different global solutions (EIGEN-GRACE02S, EIGEN-CG03C, GGM02S, GGM03S, ITG-Grace02, ITG-Grace03s, JEM01-RL03B, EGM2008, AIUB-GRACE01S) is assumed for the uncertainties $\\delta J_{\\ell}$ instead of using their more-or-less calibrated covariances $\\sigma_{J_\\ell}$, it turns out that the systematic error $\\delta\\mu$ in the Lense-Thirring measurement is about 3 to 4 times larger than in the evaluations so far published based on the use of the covariances of one model at a time separately, amounting up to $37\\%$ for the pair EIGEN-GRACE02S/ITG-Grace03s. The comparison among the other recent GRACE-based models yields bias as large as about $25-30\\%$.  The major discrepancies still occur for $J_4, J_6$ and $J_8$, which are just to which the zonals the combined LAGEOS/LAGOES II nodes are most sensitive. }  %% Keywords should be separated by \\*\\ sign ",
        "introduction": "In the weak-field and slow motion approximation, the Einstein field equations of general relativity get linearized to a form resembling Maxwell's equations of electromagnetism. Thus, a gravitomagnetic field, induced by the off-diagonal components $g_{0i}, i=1,2,3$ of the space-time metric tensor related to the mass-energy currents of the source of the gravitational field, arises \\cite{MashNOVA}. It affects in several ways the motion of, e.g., test particles and electromagnetic waves \\cite{Rug}. Perhaps the most famous gravitomagnetic effects are gyroscope precession \\cite{Pugh,Schi} and the Lense-Thirring\\footnote{According to an interesting historical analysis recently performed in Ref. \\cite{Pfi07}, it would be more correct to speak about an Einstein-Thirring-Lense effect.} precessions \\cite{LT} of the orbit of a test particle, both occurring in the field of a central slowly rotating mass like a planet.  The measurement of gyroscope precession in the Earth's gravitational field has been the goal of the dedicated space-based GP-B mission\\footnote{See \\url{http://einstein.stanford.edu/}} \\cite{Eve, GPB} launched in 2004; its data analysis is still in progress.  In this paper we critically discuss some issues concerning the test of the Lense-Thirring effect performed with the LAGEOS and LAGEOS II terrestrial artificial satellites \\cite{Ciu04} tracked with the Satellite Laser Ranging (SLR) technique \\cite{Pearl02}.  \\cite{vpe76a,vpe76b} proposed measuring the Lense-Thirring nodal precession of a pair of counter-orbiting spacecraft in terrestrial polar orbits and equipped with drag-free apparatus.  A somewhat equivalent, cheaper version of such an idea was put forth in Ref. \\cite{Ciu86} whose author proposed to launch a passive, geodetic satellite in an orbit identical to that of the LAGEOS satellite apart from the orbital planes which should have been displaced by 180 deg apart.\\footnote{ LAGEOS was put into orbit in 1976, followed by its twin LAGEOS II in 1992.}  The measurable quantity was, in this case, the sum of the nodes of LAGEOS and of the new spacecraft, later named LAGEOS III, LARES, WEBER-SAT, in order to cancel the confounding effects of the multipoles of the Newtonian part of the terrestrial gravitational potential (see below). Although extensively studied by various groups \\cite{CSR,LARES,Ioretal02}, such an idea has not been implemented for a long time.  Recently, the Italian Space Agency (ASI) has approved this project and should launch a VEGA rocket for this purpose at the end of 2009-beginning of 2010 (\\url{http://www.asi.it/en/activity/cosmology/lares}).  For recent updates of the LARES/WEBER-SAT mission, including recently added additional goals in fundamental physics and related criticisms, see  Refs. \\cite{LucPao01,Ior02,Ciu04b,Ciu06b,IorJCAP,Ior07c,Ior07d,Ior09}.  Among scenarios involving \\emph{existing} orbiting bodies, the idea of measuring the Lense-Thirring node rate with the just launched LAGEOS satellite, along with the other SLR targets orbiting at that time, was proposed in Ref. \\cite{Cug78}. Tests have been effectively performed using the LAGEOS and LAGEOS II satellites \\cite{tanti}, according to a strategy \\cite{Ciu96} involving a suitable combination of the nodes of both satellites and the perigee $\\omega$ of LAGEOS II. This was done to reduce the impact of the most relevant source of systematic bias, viz. the mismodelling in the even ($\\ell=2,4,6\\ldots$) zonal ($m=0$) harmonics $J_{\\ell}$ of the multipolar expansion of the Newtonian part of the terrestrial gravitational potential:\\footnote{The relation among the even zonals $J_{\\ell}$ and the normalized gravity coefficients $\\overline{C}_{\\ell 0}$ is $J_{\\ell}=-\\sqrt{2\\ell + 1}\\ \\overline{C}_{\\ell 0}$.} they account for non-sphericity of the terrestrial gravitational field induced by centrifugal effects of the Earth's diurnal rotation. The even zonals affect the node and the perigee of a terrestrial satellite with secular precessions which may mimic the Lense-Thirring signature. The three-elements combination used allowed for removing the uncertainties in $J_2$ and $J_4$. In \\cite{Ciu98} a $\\approx 20\\%$ test was reported by using the\\footnote{Contrary to the subsequent CHAMP/GRACE-based models, EGM96 relies upon multidecadal tracking of SLR data of a constellation of geodetic satellites including LAGEOS and LAGEOS II as well; thus the possibility of a sort of $a-priori$ `imprinting' of the Lense-Thirring effect itself, not solved-for in EGM96, cannot be neglected.} EGM96 \\cite{Lem98} Earth gravity model; subsequent analyses showed that such an evaluation of the total error budget was overly optimistic in view of the likely unreliable computation of the total bias due to the even zonals \\cite{Ior03,Ries03a,Ries03b}.  An analogous, huge underestimation turned out to hold also for the effect of non-gravitational perturbations \\cite{Mil87} like direct solar radiation pressure, the Earth's albedo, various subtle thermal effects depending on the the physical properties of the satellites' surfaces and their rotational state \\cite{Inv94,Ves99,Luc01,Luc02,Luc03,Luc04,Lucetal04,Ries03a}, which the perigees of LAGEOS-like satellites are particularly sensitive to. As a result, the realistic total error budget in the test reported in Ref. \\cite{Ciu98} might be as large as $60-90\\%$ or even more (by considering EGM96 only).  The observable used in Ref. \\cite{Ciu04} with the GRACE-only EIGEN-GRACE02S model \\cite{eigengrace02s} and in Ref. \\cite{Ries08} with other global terrestrial gravity solutions was the following linear combination\\footnote{See also \\cite{Pav02,Ries03a,Ries03b}.} of the nodes of LAGEOS and LAGEOS II, explicitly computed in Ref. \\cite{IorMor} following the approach proposed in Ref. \\cite{Ciu96}: \\begin{equation} f=\\dot\\Omega^{\\rm LAGEOS}+ c_1\\dot\\Omega^{\\rm LAGEOS\\ II }, \\lb{combi}\\end{equation} where \\begin{equation} c_1\\equiv-\\rp{\\dot\\Omega^{\\rm LAGEOS}_{.2}}{\\dot\\Omega^{\\rm LAGEOS\\ II }_{.2}}=-\\rp{\\cos i_{\\rm LAGEOS}}{\\cos i_{\\rm LAGEOS\\ II}}\\left(\\rp{1-e^2_{\\rm LAGEOS\\ II}}{1-e^2_{\\rm LAGEOS}}\\right)^2\\left(\\rp{a_{\\rm LAGEOS\\ II}}{a_{\\rm LAGEOS}}\\right)^{7/2}.\\lb{coff}\\end{equation} The coefficients $\\dot\\Omega_{.\\ell}$ of the aliasing classical node precessions \\cite{Kau} $\\dot\\Omega_{\\rm class}=\\sum_{\\ell}\\dot\\Omega_{.\\ell}J_{\\ell}$ induced by even zonals have been analytically worked out in e.g. \\cite{Ior03}; $a,e,i$ are the satellite's semimajor axis, eccentricity and inclination, respectively and yield $c_1=0.544$ for \\rfr{coff}.  The Lense-Thirring signature of \\rfr{combi} amounts to 47.8 milliarcseconds per year (mas yr$^{-1}$). The combination \\rfr{combi} allows, by construction, to remove the aliasing effects due to the static and time-varying parts of the first even zonal $J_2$. The nominal bias (computed with the estimated values of $J_{\\ell}$, $\\ell=4,6...$) due to the remaining higher degree even zonals would amount to about $10^5$ mas~yr$^{-1}$; the need of a careful and reliable modeling of such an important source of systematic bias is, thus, quite apparent. Conversely, the nodes of the LAGEOS-type spacecraft are affected by the non-gravitational accelerations $\\approx 1\\%$ of the Lense-Thirring effect \\cite{Luc01,Luc02,Luc03,Luc04,Lucetal04}. For a comprehensive, up-to-date overview of the numerous and subtle issues concerning the measurement of the Lense-Thirring effect see \\cite{IorNOVA}.  Here, we will address the following questions: \\begin{itemize} \\item Has the systematic error due to the competing secular node precessions induced by the static part of the even zonal harmonics been realistically evaluated so far in literature? (Section \\ref{grav}) % \\item Why has the analysis with the LAGEOS satellites not been %   repeated so far by any other independent team? (\\sref{repeat}) \\item Are other approaches to extract the gravitomagnetic signature from the data feasible? (Section \\ref{approach}) \\end{itemize} ",
        "conclusions": "In this paper we have shown how the so far published evaluations of the total systematic error in the Lense-Thirring measurement with the combined nodes of the LAGEOS satellites due to the classical node precessions induced by the even zonal harmonics of the geopotential are likely optimistic. Indeed, they are all based on the use of elements from the covariance matrix, more or less reliably calibrated, of various Earth gravity model solutions used one at a time separately in such a way that the model X yields an error of $x\\%$, the model Y yields an error $y\\%$, etc. Instead, comparing the estimated values of the even zonals for different pairs of models allows for a much more realistic evaluation of the real uncertainties in our knowledge of the static part of the geopotential. As a consequence, the bias in the Lense-Thirring effect measurement is about three to four times larger than that so far claimed, amounting to tens of parts per cent (37$\\%$ for the pair EIGEN-GRACE02S and ITG-GRACE03s, about 25--30$\\%$ for the other most recent GRACE-based solutions). % % We have also pointed out that, until now, no other tests of the % Lense-Thirring effect have been performed by independent teams, % although it would % be, %at least in principle, a relatively not too demanding task in view of the wide dissemination of SLR stations, for which the LAGEOS satellites are %important targets since long time, and of the freely available softwares to perform the data reduction process. %  Finally, we have pointed out the need of following different strategies in extracting the Lense-Thirring pattern from the data; for instance by explicitly modelling it in fitting the SLR data of LAGEOS and LAGEOS II, and estimating the associated solve-for parameter in a least-square sense along with the other parameters usually determined. %"
    },
    "0710/0710.0141_arXiv.txt": {
        "abstract": "{} {Molecules that trace the high-density regions of the interstellar medium have been observed in (ultra-)luminous (far-)infrared galaxies, in order to initiate multiple-molecule multiple-transition studies to evaluate the physical and chemical environment of the nuclear medium and its response to the ongoing nuclear activity.} {The HCN(1$-$0), HNC(1$-$0), \\HCOP(1$-$0), CN(1$-$0) and CN(2$-$1), CO(2$-$1), and CS(3$-$2) transitions were observed in sources covering three decades of infrared luminosity including sources with known OH megamaser activity. The data for the molecules that trace the high-density regions have been augmented with data available in the literature.} {The integrated emissions of high-density tracer molecules show a strong relation to the far-infrared luminosity. Ratios of integrated line luminosities have been used for a first order diagnosis of the integrated molecular environment of the evolving nuclear starbursts. Diagnostic diagrams display significant differentiation among the sources that relate to initial conditions and the radiative excitation environment. Initial differentiation has been introduced between the FUV radiation field in photon-dominated-regions and the X-ray field in X-ray-dominated-regions. The galaxies displaying OH megamaser activity have line ratios typical of photon-dominated regions. } {} ",
        "introduction": "Large amounts of high-density molecular gas play a crucial role in the physics of (ultra-)luminous (far-)infrared galaxies ((U)LIRGs), giving rise to spectacular starbursts and possibly providing the fuel for the emergence of active galactic nuclei (AGN). A strong linear relation has been found between the far infrared (\\irlums) and HCN luminosity, the latter being an indicator of the high density (\\numd$\\ga$10$^4$\\,\\percc) component of the interstellar medium. This arises predominantly from the nuclear region and indicates a close relation to the nuclear starburst environment and the production of the FIR luminosity \\citep{GaoS2004a}. The molecular gas contributes a significant fraction to the dynamical mass  of the central regions of FIR galaxies \\citep{1991A&ARv...3...47H, YoungScoville1991}.  Emissions of molecules that trace the high-density regions in LIRGs and ULIRGs may serve to study the characteristics of the extreme environments in the nuclear regions of starburst and AGN nuclei. While much of the energy generation of FIR-luminous galaxies has traditionally been attributed to AGN activity using their optical signatures, the NIR and radio signatures suggest significant starburst contributions \\citep{GenzelEA1998, BaanKlockner2006}. Circum-nuclear star formation in the densest regions serves as a power source for the majority of ULIRGs such that they even mimic the presence of an AGN \\citep{BaanKlockner2006}. The ULIRG population is also characterized by OH megamasers (MM) \\citep{Baan1989, HenkelWilson1990} resulting from amplification of radio continuum by FIR-pumped foreground gas. Few H$_2$CO MM are also found among this population \\citep{BaanHU1993, ArayaBH2004}.  The observed molecular emissions from the high-density components could carry the signature of the nuclear region of the ULIRG and its nuclear power generation. Recent studies consider the HCN\\,(1$-$0) versus CO\\,(1$-$0) relation and its implications for the star formation activity \\citep{GaoS2004a, GaoS2004b}. Discussions are ongoing in the literature about whether such molecular emissions are accurate tracers of the high-density medium at the site of star formation \\citep{WuEvansEA2005, GraciaCarpioGPC2006, Papadop2007}. It should be evident that only multi-transition and multi-molecule (multi-dimensional) studies can describe the multi-parameter environment of the nuclear activity and further establish the connection between the molecular signature and the nature of the nuclear energy source. The translation or extrapolation from the characteristics of individual star formation regions in our Galaxy to integrated regions in nearby galaxies at larger distances or even to unresolved nuclear regions in distant galaxies requires detailed multi-species observations and comparisons combined with elaborate theoretical modelling. The interpretation of molecular line data has focused on the nature of the local excitation mechanisms in the form of photon dominated regions (PDR) with far-ultraviolet radiation fields and X-ray dominated regions (XDR) \\citep{1996A&A...306L..21L, MeijerinkS2005}. The molecular properties of these two different regimes can be used as diagnostic tools in interpreting the integrated properties of galaxies.  This paper reflects a multi-dimensional approach to understanding molecular line emissions that started in the early 90's, but is only now coming to fruition. This study originally aimed to establish a better link between molecular megamaser activity and the molecular properties of luminous FIR galaxies but it has changed into a general study of ULIRGs. Here we present data on the molecular characteristics of normal and FIR galaxies across the available range of \\irlums. Data of the CO, HCN, HNC, \\HCOP, CN, and CS line emissions obtained for a representative group of 37 FIR-luminous and OH megamaser galaxies has been joined with the additional data of 80 sources presented in the literature. Early studies of the properties of CS \\citep{MauersbergerHWH1989}, \\HCOP\\ \\citep{NguyenEA1992}, and HNC \\citep{HuttemeisterEA1995} using small numbers of sources studied the relations between line and FIR luminosities and the nature of the central engine. Our molecular emission data show clear tendencies that cover a wide regime of nuclear activity. This database forms the basis for a first evaluation of the emission line properties and for further study and modelling of the line properties. ",
        "conclusions": "Molecular line emissions in FIR-luminous galaxies are tools for multi-dimensional diagnostics of the environmental parameters in the nuclear ISM and the heating processes resulting from the nuclear activity.  A total of 118 line detections made with the 30-m Pico Veleta and the 15-m SEST telescopes for a total of 37 sources have been complemented with partial records published for another 80 sources.  A proper evaluation of integrated emission lines and their diagnostic line ratios can be achieved after synthesis and modelling of a representative description of the integrated nuclear molecular medium under a wide variety of AGN- and starburst-related circumstances. The molecular information in this paper has been diagnosed to first order using modelling results for nuclear emission regions presented in the literature.  The collective behavior of line luminosities and line ratios of the low-density and high-density tracers presents a consistent picture of the molecular medium in the nuclear regions of ULIRGs. The high-density tracers represent the molecular medium in the regions where star formation is taking place, and the low-density tracer represents the relatively unperturbed larger-scale molecular environment. The tracer luminosities increase roughly linearly with FIR luminosity but they show significant scatter in data points due to physical processes in the nuclear region, the chemical history, and the relative age of the nuclear activity. The luminosity of high-density tracers varies linearly with \\irlum except for the CS(3$-$2) luminosity, where a steeper relation is suggestive of a different excitation dependence. The CO(1$-$0) and CO(2$-$1) lines for these high FIR luminosity sources have a slope less than unity and are (almost) consistent with empirical evidence on the relation of SFR versus gas content \\citep{Kennicutt1998}.  First order diagnostics of the nuclear ISM can be based on the collective behavior of high-density tracers HCN, HNC and \\HCOP\\ showing significant differentiation and systematic changes. The line strengths of these three molecules relative to the CO\\,(1$-$0) line show a significant range of 0.03 to 0.23 with a distinct dependence on \\irlums. This dependence has been interpreted in terms of an evolutionary model where the depletion of molecular gas depends on the consumption and destruction of the high-density gas by the ongoing star formation process. Only partial diagnostics could be done using the available CN and CS data, which complements the diagnostics of the three other high-density tracers.  The emission line ratios of the three high-density tracers of the galaxies show a structured distribution that fills selected parts of the parameter space. Interpreting these distributions has been done using modelling column densities that are characteristic of PDR- and XDR-dominated nuclear environments. OH MM and other powerful ULIRGs are mostly characterized by PDR-dominance. The other end of the data distribution is characterized by mixed PDR-XDR and XDR-dominated environments. (U)LIRGs represent the phases of nuclear evolution that generate the highest FIR luminosities and most rapidly deplete the dense molecular component. Using the HCN/CO ratio as an evolutionary indicator, the distributions of data points may reflect evolution from PDR-like to more XDR-like nuclear ISM properties during the course of the outburst.  The \\HCOP/HCN and HNC/HCN ratios serve as indicators of environments with shocks and non-standard heating and with HNC depletion and \\HCOP\\ enhancement in a fraction of the sources. The \\HCOP/HCN ratio also serves as a density indicator and would be higher under PDR circumstances and lower in XDR environments. The observed trends in the \\HCOP/HCN ratio could result from the (indirect) influence of an AGN in the nucleus. More likely the trends in the relative abundance of HCN as compared with other constituents result from evolution of the nuclear environment during the SBN activity.  The detailed interpretations of multi-line multi-molecule emission line behavior and their relation with the local environment require detailed modelling of the physical parameters of the environment together with the excitation, the chemistry, and the radiative transfer for the molecular constituents. Hereby the integration of higher level transitions of the molecules is needed to determine specific excitation temperatures and densities. A comparison with the properties of Galactic emission regions will emphasize the effect of scale size on the line ratios for tracer molecules. Further studies are underway to connect the FIR signature and global heating scenario of ULIRGs \\citep{LoenenBS2006} with modelling of the molecular emissions under varying conditions in a nuclear starburst. The emission scenarios for XDRs and PDRs establish the connection between star formation and other sources of excitation and the integrated emission line parameters observed in extragalactic sources.  \\bigskip"
    },
    "0710/0710.5849_arXiv.txt": {
        "abstract": "{ The Edelweiss programme is dedicated to the direct search for Dark Matter as massive weakly interacting particles (WIMPs) with Germanium cryogenic detectors operated in the Laboratoire Souterrain de Modane in the French Alps at a depth of 4800 mwe. After the initial phase Edelweiss I, which involved a total mass of 1 kg, the second step of the programme, Edelweiss II, currently operates 9 kg of detectors and an active shielding of 100 m$^2$ muon veto detectors and is now in its commissioning phase. The current status and performance of the Edelweiss II set-up in terms of backgrounds will be given, the underground muon flux measured with the muon veto system will be presented. \\PACS{ {95.35.+d}{Dark matter search}   \\and {14.80.Ly}{Supersymmetric partners of known particles} \\and {98.80.Es}{Observational cosmology} \\and {98.70.Vc}{Background radiations, cosmic} } % } % ",
        "introduction": "\\label{intro} Understanding the nature of Dark Matter in the universe is a major challenge for modern cosmology and astrophysics. One of the well-motivated candidates is the generically named Weakly Interacting Massive Particle (WIMP). In the Minimal Supersymetric Standard Model (MSSM) framework, the WIMP could be the Lightest Supersymmetric Particle, which is stable, neutral and massive. The EDELWEISS (Experience pour DEtecter Les Wimps En Site Souterrain) experiment is dedicated to the direct detection of WIMPs trapped in the galactic halo. The detection is performed through the measurement of the recoil energy produced by elastic scattering of a WIMP off target nuclei. The main challenges are the extremely low event rate of $\\le$ 1 evt/kg/year and the relatively small deposited energy of $\\le$ 100 keV. ",
        "conclusions": "\\label{sec:conc}  The EDELWEISS-II setup has been validated with calibration and background runs. Energy resolutions and discrimination capabilities close to those of EDEL\\-WEISS-I have been measured for Ge/NTD detectors. Validation of Ge/NbSi detectors with new aluminium electrodes is in progress and Ge/NTD detectors with interdigitized electrodes scheme have shown promising results in surface laboratory. The muon veto system runs, and study of coincidences between the muon veto and the bolometer systems will be performed in the following months. A new setup for muon-neutron measurement will be soon installed at LSM and will allow to evaluate the muon induced neutron flux. Low background physics runs will be taken with the 28 bolometers setup with the aim to reach sensitivity to WIMP-nucleon cross-section of $\\sim$10$^{-7}$ pb for a WIMP mass of 100 GeV by July 2008. 40 additional detectors (Ge/NbSi and Ge/NTD with interdigitized electrodes) will be added in the two coming years to enhance progressively the sensitivity to $\\sim$10$^{-8}$ pb thanks to active surface rejection, with an expected goal reached by autumn 2009/2010.  \\sloppy \\begin{acknowledgement} {\\bf Acknowledgment:} This work has been partially supported by the EEC Applied Cryodetector network (Contract HPRN-CT-2002-00322), the ILIAS integrating activity (Contract RII3-CT-2003-506222), the HGF initiative and networking fund through the virtual institute VIDMAN and the SFB/Transregio 27 of the Deutsche Forschungsgemeinschaft, DFG. The help of the technical staff of the Laboratoire Souterrain de Modane and of the participating laboratories is gratefully acknowledged. \\end{acknowledgement}"
    },
    "0710/0710.2691_arXiv.txt": {
        "abstract": "We present new optical, near-IR, and mid-IR observations of the young   eruptive variable star V1647~Orionis that went into outburst in late   2004 for approximately two years.  Our observations were taken one year after the star had faded to its pre-outburst optical brightness and show that V1647~Ori is still actively accreting circumstellar material.  We compare and contrast these data with existing observations of the source from both pre-outburst and outburst phases.  From near-IR spectroscopy we identify photospheric absorption features for the first time that allow us to constrain the classification of the young star itself.  Our best fit spectral type is M0$\\pm$2 sub-classes with a visual extinction of 19$\\pm$2 magnitudes and a K-band veiling of r$_K\\sim$1.5$\\pm$0.2.  We estimate that V1647~Ori has a quiescent bolometric luminosity of $\\sim$9.5~L$_{\\odot}$ and a mass accretion rate of $\\sim$1$\\times$10$^{-6}$~M$_\\odot$~yr$^{-1}$.  Our derived mass and age, from comparison with evolutionary models, are  0.8$\\pm$0.2~M$_{\\odot}$ and $\\lesssim$~0.5~Myrs, respectively.  The presence towards the star of shock excited optical [S~II] and [Fe~II] emission as well as near-IR H$_2$ and [Fe~II] emission perhaps suggests that a new Herbig-Haro flow is becoming visible close to the star. ",
        "introduction": "A significant event in the recent history of star formation studies occurred in late 2003 when the young pre-main sequence star V1647~Orionis went into outburst.  This eruption, and the associated one hundred fold increase in optical brightness, resulted in the appearance of a monopolar reflection nebula now known as ''McNeil's Nebula'' after the amateur astronomer, Jay McNeil, who discovered it (McNeil 2004).  The star and nebula remained bright for approximately 18 months before fading rapidly over a six month period.   V1647~Ori was found to be again close to its pre-outburst (Sloan Digital Sky Survey, SDSS) optical brightness in early 2006.  A widely accepted interpretation of such eruptions is one involving a heated disk and a `mass dumping' episode from a circumstellar disk onto a stellar photosphere.  The five magnitude increase in optical brightness seen in V1647~Ori has been attributed to both the addition of significant accretion luminosity and a partial dust clearing due to the high-velocity wind (up to 600~km~s$^{-1}$) emanating from the star/disk engine (Reipurth \\& Aspin 2004; McGehee et al. 2004).  Both star and nebula have been widely studied from the period soon after the outburst occurred through to its fading to its pre-outburst brightness.  In our discussions below, we refer to the period before November 2004 as the {\\it pre-outburst phase} (Brice\\~no et al. 2004), the period from November 2004 to February 2006 as the {\\it outburst phase}, and the period from February 2006 onwards as the {\\it quiescent phase} of V1647~Ori.  Observations spanning all spectral regimes, from x-ray into the optical to the near-IR and on to the far-IR/sub-mm/radio, have been published by many authors in a large number of papers.  Readers are referred to these for detailed background information on V1647~Ori and McNeil's Nebula.  These papers are, in alphabetical order: \\'Abr\\'aham et al. (2004 - the pre-outburst infrared characteristics), Acosta-Pulido et al. (2007 - optical and near-IR imaging photometry, near-IR spectroscopy and polarimetry during the outburst, henceforth AP07), Andrews, Rothberg, \\& Simon (2004 -- mid-IR and sub-mm observations during the outburst), Aspin et al. (2006 -- an historical study of previous outbursts), Brice\\~no et al. (2004 -- the optical outburst history, pre-outburst and during the outburst), Fedele et al. (2007a,b -- optical and near-IR properties during the outburst), Gibb et al. (2006 -- near-IR spectroscopy during the outburst), Grosso et al. (2005 -- x-ray observations during the outburst), Kastner et al. (2004 -- x-ray emission pre-outburst and during the outburst), Kastner et al. (2006 -- x-ray evolution during the outburst), K\\'osp\\'al et al. (2005 -- optical photometry during the outburst), McGehee et al. (2004 -- optical and near-IR photometry pre-outburst and during the outburst), Mosoni et al. (2005 -- mid-IR interferometric observations during the outburst), Muzerolle et al. (2005 -- {\\it Spitzer} observations during the outburst), Ojha et al. (2004 -- near-IR imaging during the outburst), Ojha et al. (2006 -- optical photometry and spectroscopy, and near-IR imaging during the outburst), Reipurth \\& Aspin (2004 -- optical and near-IR imaging photometry and spectroscopy during the outburst), Rettig et al. (2005 -- high spectral resolution near-IR observations during the outburst), Semkov (2004 -- optical photometry during the outburst), Semkov (2006 -- optical light curve during the outburst), Tsukagoshi et al. (2005 -- mm continuum observations during the outburst), Vacca, Cushing, \\& Simon (2004 -- near-IR spectroscopy during the outburst), Vig et al. (2006 -- radio observations during the outburst), and Walter et al. (2004 -- optical photometry and spectroscopy during the outburst).  The significance of eruptive outbursts, similar to the one undergone by V1647~Ori lies in the fact that these periods are considered times when mass accretion rates increase dramatically (Hartmann \\& Kenyon 1996, henceforth HK96).  During periods of outburst, mass accretion rates have been directly observed to increase by up to several orders of magnitude from a pseudo-steady-state value of $\\sim$10$^{-9~to~-7}$~M$_\\odot$~yr$^{-1}$ for classical T~Tauri stars (Hartmann et al. 1998) to $\\sim$10$^{-6~to~-5}$~M$_\\odot$~yr$^{-1}$ for EXor variables, and as high as $\\sim$10$^{-4}$~M$_\\odot$~yr$^{-1}$ for the extremely energetic FUor variables (Hartmann \\& Kenyon 1985).  FUors and EXors are so named after the prototypes of their classes, namely FU~Orionis and EX~Lupi, respectively (Herbig 1989).  Whether all young stars undergo EXor and/or FUor eruptions and, if so, whether these two types of outbursts events are fundamentally the same (on perhaps different scales) or different (in terms of the mechanism involved and the triggering of the eruption) is still a matter of debate.  EXor eruptions are distinct from FUor eruptions in that they are much shorter lived, months to years rather than decades to centuries, and have been empirically determined to be repetitive (Herbig 1989; 2007).  For example, EX~Lupi itself has gone into outburst on several occasions, with the last being in 1993--94 (Lehmann et al. 1995, Herbig et al. 2001).  V1647~Ori is likely another example of an EXor since its recent activity lasted only two years and it was previously observed in outburst from 1966--67 (Aspin et al. 2006) on archival photographic plates.  Mechanisms that have been proposed to initiate an eruptive event, be it EXor-like or FUor-like, include: thermal disk instabilities resulting in a runaway accretion condition (Bell \\& Lin 1994), periodic overloading of the inner regions of the circumstellar disk and subsequent magnetic collapse (HK96), and the close approach of a companion in an eccentric orbit disturbing the stability of the inner regions of the disk (Bonnell \\& Bastien 1992).  In this paper, we present new optical, near-IR, and mid-IR observations of V1647~Ori and McNeil's Nebula taken approximately one year after the source reached its pre-outburst optical brightness.  Below, we compare and contrast our observations with those taken during the pre-outburst and outburst phases of the eruption and attempt to better understand the nature and characteristics of the underlying young star and its circumstellar environment. ",
        "conclusions": "From our optical, near-IR, and mid-IR imaging and spectroscopy of V1647~Ori dating from February to April 2007, approximately one year after its return to its pre-outburst optical brightness, we can conclude that:  \\begin{itemize} \\item The associated nebula, McNeil's Nebula, remained faintly visible suggesting that quasi-static nebula material is still being illuminated by V1647~Ori.  \\item We confirm the findings of Fedele et al. (2007a) in that signposts of shock-excited emission, specifically, in the emission lines of [S~II], [Fe~II], and H$_2$ are present.  We consider that it is a distinct possibility that these emission lines result from a new Herbig-Haro flow.  \\item Our near-IR spectrum shows, for the first time, molecular overtone absorption from CO and atomic absorption from neutral Na and Ca.  We interpret this as evidence that we are now observing the stellar photosphere of V1647~Ori. \\item We have modeled the near-IR spectrum using a template stellar spectrum including both water vapor and water ice absorption and have determined a best-fit parameter set of spectral type M0$\\pm$2 sub-classes, A$_V$=19$\\pm$2 magnitudes, r$_K$=1.5$\\pm$0.2.  \\item We derive values of M$_{bol}$=2.9$\\pm$0.4 magnitudes and  L$_{*}$=5.2$\\pm$2~L$_{\\odot}$ for V1647~Ori. From comparison with theoretical evolutionary tracks, find that, for the adopted T$_{eff}\\sim$3800~K, the star has a mass and age of 0.8$\\pm$0.2~M$_{\\odot}$ and $\\lesssim$0.5~Myrs, respectively.  \\item From near-IR H~I line flux and the A$_V$ values derived above, we estimate the accretion luminosity and mass accretion rate in February 2007 was L$_{acc}$=4.0$\\pm$2~L$_{\\odot}$ and \\.M=1.0$\\times$10$^{-6}$~M$_{\\odot}$~yr$^{-1}$, respectively.  This implies that V1647~Ori is still actively accreting circumstellar material even though it is almost optically invisible and that the accretion rate during the outburst must have been considerably larger that this value.  \\item V1647~Ori is found to have a quiescent phase L$_{bol}$ of  9.25$\\pm$3~L$_{\\odot}$.  \\item For the first time, we see evidence for silicate dust evolution in the mid-IR spectrum of V1647~Ori over the outburst to quiescence period.  We now observe weak silicate absorption at 10~$\\mu$m whereas previously the silicate band was either absent or weakly in emission.  \\item Finally we note that, in February 2007, the spectral energy distribution, SED, of V1647~Ori appears remarkably similar to its pre-outburst SED suggesting that perhaps the derived accretion rate is the normal quiescent phase accretion rate for this object. \\end{itemize}  \\vspace{0.3cm}  {\\bf Acknowledgments}  Based on observations obtained at the Gemini Observatory (under program identification GN-2007A-Q-33 and GS-2005B-Q-13), which is operated by the Association of Universities for Research in Astronomy, Inc., under a cooperative agreement with the NSF on behalf of the Gemini partnership: the National Science Foundation (United States), the Particle Physics and Astronomy Research Council (United Kingdom), the National Research Council (Canada), CONICYT (Chile), the Australian Research Council (Australia), CNPq (Brazil) and CONICET (Argentina).  Colin Aspin and Tracy Beck were visiting astronomers at the Infrared Telescope Facility, which is operated by the University of Hawaii under Cooperative Agreement no. NCC 5-538 with the National Aeronautics and Space Administration, Science Mission Directorate, Planetary Astronomy Program.  This material is based upon work supported by the National Aeronautics and Space Administration through the NASA Astrobiology Institute under Cooperative Agreement No. NNA04CC08A issued through the Office of Space Science."
    },
    "0710/0710.0694_arXiv.txt": {
        "abstract": "We have designed a system to stabilize the gain of a submillimeter heterodyne receiver against thermal fluctuations of the mixing element. In the most sensitive heterodyne receivers, the mixer is usually cooled to 4 K using a closed-cycle cryocooler, which can introduce $\\sim$1\\% fluctuations in the physical temperature of the receiver components. We compensate for the resulting mixer conversion gain fluctuations by monitoring the physical temperature of the mixer and adjusting the gain of the intermediate frequency (IF) amplifier that immediately follows the mixer.  This IF power stabilization scheme, developed for use at the Submillimeter Array (SMA), a submillimeter interferometer telescope on Mauna Kea in Hawaii, routinely achieves a receiver gain stability of 1 part in 6,000 (rms to mean).  This is an order of magnitude improvement over the typical uncorrected stability of 1 part in a few hundred. Our gain stabilization scheme is a useful addition to SIS heterodyne receivers that are cooled using closed-cycle cryocoolers in which the 4 K temperature fluctuations tend to be the leading cause of IF power fluctuations. ",
        "introduction": "\\PARstart{S}{uperconductor} insulator superconductor (SIS) receivers are in use on many millimeter and submillimeter telescopes because of their good spectral line sensitivity. Their continuum sensitivity, however, does not usually reach the theoretical limit because of receiver gain fluctuations.  These arise predominantly from changes in mixer conversion gain, which result from physical temperature changes and local oscillator (LO) power changes. Here we describe a technique to reduce the impact of physical temperature changes on mixer conversion gain variations by changing the intermediate frequency (IF) gain of the receiver in proportion to fluctuations of the physical temperature of the SIS junction used as a mixing element.     A number of radio observatories make use of liquid Helium filled cryostats so that temperature induced receiver gain fluctuations are minimized.  However, due to their lower degree of maintenance and upkeep, closed-cycle Helium cryocooler systems may be preferred. For example multi-receiver systems with high heat loads, such as those in use at the 8-element Submillimeter Array (SMA)\\footnote{The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics, and is funded by the Smithsonian Institution and the Academia Sinica.} \\cite{refs:SMA} or the upcoming 64-element Atacama Large Millimeter Array (ALMA) \\cite{refs:Wootten}, are dependent on the use of cryocoolers. For these applications it is necessary to develop techniques to reconcile the need for highly stable receivers with the practical benefits of cryocoolers.  An obvious approach is to design more stable cryocoolers, another is to compensate for the fluctuations in the mixer's conversion gain.  In the Berkeley-Illinois-Maryland-Association Millimeter Array (BIMA) a heating resistor on the cold stage is used to compensate for cryocooler temperature fluctuations \\cite{refs:plambeck}.  This reduces long time scale temperature drifts, giving 1 mK rms temperature stability over several hours when the data are averaged in 2 second blocks.  The cycling of the cryocooler's displacer also introduces temperature fluctuations on 1-2 second time scales, which are not treated by this technique.  To provide a large thermal capacity that will damp the fluctuations in the temperature of the cryocooler, it is possible to add a liquid helium reservoir.  Using this technique, a ten fold improvement in temperature stability, from 200 mK to 20 mK (peak to peak), has been demonstrated \\cite{refs:sekimoto}.  The ALMA cryocooler also incorporates a helium reservoir and has achieved an rms temperature stability of 1 mK at 4 K \\cite{refs:anna}.   We take a different approach to receiver stabilization.  While monitoring the physical temperature of the mixer we actively vary the IF gain of the receiver to compensate for mixer conversion gain variations.  Such a scheme is straightforward to implement and can be applied to existing systems without redesigning the cryocooler.  We have developed and tested hardware and accompanying software to adjust the third stage gain of the low noise HEMT amplifier immediately following the mixer in proportion to the fluctuations in the physical temperature of the mixer.  We achieve nearly an order of magnitude improvement in receiver stability from 1 part in $\\sim$800 to 1 part in 6,000 (rms to mean ratio) using a 33 millisecond integration time over 10 minutes.  This level of receiver stability equals that which would be obtained if rms fluctuations of the physical temperature of the mixer were held to $<$ 2 mK. ",
        "conclusions": "We have developed a system that stabilizes the gain of a 230 GHz SIS heterodyne receiver to 1 part in 6,000 (rms to mean).  With the use of an open-loop proportional servo system, we monitor the physical temperature of the mixer block and adjust the gain of the third and final stage of a low noise HEMT IF amplifier in order to compensate for the subsequent variations in the conversion gain of the mixer. We have shown that the conversion gain of the mixer varies linearly over the range of typical cryocooler-induced physical temperature fluctuations, and that the gain of the HEMT amplifier varies linearly with its third stage bias current.  Using our gain stabilization system, we routinely achieve a total receiver gain stability of 1 part in 6,000 (rms to mean ratio) for a 0.033 second integration time over 10 minutes, which corresponds to an effective rms temperature fluctuation of 1.7 mK. In comparison, the typical fluctuations in the physical temperature of the mixer at the SMA are on the order of 50-100 mK, and the corresponding power fluctuations are 1 part in several hundred. Therefore, our gain stabilization servo system can provide more than an order of magnitude improvement over a typical unstabilized system.  Our technique introduces a negligible level of instrumental phase error into the source signal. This gain stability system can be implemented on any heterodyne receiver system in which the physical temperature of the mixer can be monitored and the IF amplification can be adjusted in real time."
    },
    "0710/0710.0007_arXiv.txt": {
        "abstract": "The Hubble  Ultra Deep Field  (HUDF) contains a significant  number of \\bb-, \\vv- and  \\ii-band dropout objects, many of  which were recently confirmed to  be young star-forming  galaxies at $z\\!\\simeq\\!4\\!-\\!6$. These galaxies are too  faint individually to accurately measure their radial surface brightness profiles.   Their average light profiles are potentially of  great interest,  since they may  contain clues  to the time since  the onset of  significant galaxy assembly.   We separately co-add \\vv, \\ii- and \\zz-band  HUDF images of sets of $z\\!\\simeq\\!4,5$ and $6$  objects, pre-selected to have nearly  identical compact sizes and the  roundest shapes.  From these  stacked images, we  are able to study the average(d) radial structure  of these objects at much higher signal-to-noise ratio  than possible  for an individual  faint object. Here we explore the reliability and usefulness of a stacking technique of compact  objects at $z\\!\\simeq\\!4\\!-\\!6$ in the  HUDF.  Our results are:  (1) image  stacking provides  reliable and  reproducible average surface  brightness profiles;  (2) the  shape of  the  average surface brightness profile  shows that even  the faintest $z\\!\\simeq\\!4\\!-\\!6$ objects are  \\emph{resolved}; and  (3) if late-type  galaxies dominate the  population  of  galaxies  at  $z\\!\\simeq\\!4\\!-\\!6$,  as  previous \\emph{HST}  studies have  shown for  $z\\!\\lesssim\\!4$, then  limits to dynamical age  estimates for these galaxies from  their profile shapes are comparable with  the SED ages obtained from  the broadband colors. We  also present accurate  measurements of  the sky-background  in the HUDF and its associated 1$\\sigma$ uncertainties. ",
        "introduction": "In  the last  decade,  ground  and space  based  observations of  high redshift  galaxies  have  begun  to  outline  the  process  of  galaxy assembly.   The details of  that process  at high  redshifts, however, remain poorly constrained.  There  is increasing support for the model of  galaxy formation,  in  which the  most  massive galaxies  assemble earlier      than       their      less      massive      counterparts \\citep[e.g.][]{cowi96,guzm97,koda04,mcca04}.   A detailed  analysis of the  `fossil record'  of  the current  stellar  populations in  nearby galaxies   selected   from  the   \\emph{Sloan   Digital  Sky   Survey} \\citep[SDSS;][]{york00} provides  strong evidence for  this downsizing picture  \\citep{heav04,pant07}.   The  increasing number  of  luminous galaxies  spectroscopically   confirmed  to  be   at  $z\\!\\simeq\\!6.5$ \\citep[e.g.][]{hu02,  kodi03,  kurk04,  rhoa04,  ster05,  tani05},  or \\cle0.9 Gyr  after the Big  Bang, also supports this  general picture. In an  alternate hierarchical scenario, arguments have  been made that significant number  of low luminosity  dwarf galaxies were  present at these times,  and were the main  contributor to finish  the process of reionization   of  the  intergalactic   medium  \\citep{yan04a,yan04b}. However,  there  is  presently  little information  on  the  dynamical structure of  these or  other galaxies at  $z\\!\\simeq\\!6$.  It  is not clear  whether  these  objects  represent isolated  disk  systems,  or collapsing spheroids, mergers or other dynamically young objects.  \\citet{ravi06}  used deep, multi-wavelength  images obtained  with the \\emph{Hubble Space Telescope} (\\emph{HST}) Advanced Camera for Surveys (ACS) as part  of the Great Observatories Origins  Deep Survey (GOODS) to analyze 2-D surface brightness distributions of the brightest Lyman Break Galaxies (LBGs)  at $2.5\\!<\\!z\\!<\\!5$.  They distinguish various morphologies  based on the  S\\'{e}rsic index  $n$, which  measures the shape of  the azimuthally  averaged surface brightness  profile (where $n$=1  for exponential  disks and  $n$=4  for a  de Vaucouleurs  law). \\citet{ravi06} find that 40\\% of the LBGs have light profiles close to exponential, as seen for disk  galaxies, and only $\\sim$30\\% have high $n$,  as seen  in  nearby  spheroids.  They  also  find a  significant fraction ($\\sim$30\\%) of galaxies with light profiles \\emph{shallower} than  exponential, which appear  to have  multiple cores  or disturbed morphologies, suggestive  of close  pairs or on-going  galaxy mergers. Distinction  between  these possible  morphologies  and, therefore,  a better estimate of the formation  redshifts of the systems observed at $z\\!\\simeq\\!4\\!-\\!6$  in  particular,  is  important for  testing  the galaxy assembly  picture, and for  the refinement of  galaxy formation models.  One possible technique involves the radial surface brightness profiles of  the most  massive objects  --- those  that will  likely  evolve to become  the massive  elliptical galaxies,  which  we see  in place  at redshifts $z\\!\\lesssim\\!2$  \\citep{driv98,vand03,vand04}.  This can be analytically   understood  in  the   context  of   the  \\citet{lynd67} relaxation formalism and the numerical galaxy formation simulations of \\citet{vana82},  which  describe  collisionless collapse  and  violent relaxation as the formation mechanism for elliptical galaxies.  As the time-scale for  relaxation is shorter in  the inner than  in the outer parts of a galaxy, convergence toward a $r^{1/n}$-profile will proceed from  the  inside  to  progressively  larger  radii  at  later  times. Moreover,  \\citet{korm77} has  shown that  tidal perturbations  due to neighbors can  cause the radial surface brightness  profile to deviate from a  pure de~Vaucouleurs  profile in the  outer parts of  a galaxy. This implies  that the radius where surface  brightness profiles start to deviate significantly from  an $r^{1/n}$ profile \\emph{might} serve as a ``\\emph{virial  clock}'' that traces the time  since the onset of the last major  merger, accretion events or global  starburst in these objects.  Image   stacking  methods   have  been   used  extensively   on  X-ray \\citep{nand02,bran01}  and radio  \\citep{geor03,whit07} data  to study the mean properties (e.g. flux, luminosity) of well-defined samples of sources  that are  otherwise too  faint to  be  detected individually. \\citet{pasc96} applied  such a  stacking method to  a large  number of optically very  faint, compact objects at $z\\!=\\!2.39$  to trace their ``average''   structure.    This   approach   was  also   applied   by \\citet{zibe04} to  detect the presence  of faint stellar  halos around disk  galaxies selected  from  the  SDSS.  An  attempt  to apply  this technique  to  high  redshift   galaxies  in  the  Hubble  Deep  Field \\citep[HDF;][]{will96}  was  not  conclusive  (H.   Ferguson;  private communication) due to the  poorer spatial sampling and shallower depth of   the    HDF   compared   to   the   Hubble    Ultra   Deep   Field \\citep[HUDF;][]{beck06}.  In this paper, we use  the exceptional depth and fine spatial sampling of the HUDF  to study the potential of  this image stacking technique, and will estimate limits to dynamical ages of faint, young galaxies at $z\\!\\simeq\\!4\\!-\\!6$.  The HUDF  reaches $\\sim$1.5~mag deeper than the equivalent  HDF  exposure in  the  \\ii-band,  and  has better  spatial sampling than the HDF.  The  HUDF depth also allows us to characterize the sky background very accurately, which is critical for successfully using  a  stacking  method  to  measure  the  mean  surface-brightness profiles for these faint young galaxies.  This  paper  is  organized  as follows:  In  \\secref{observations}  we summarize the HUDF observations, and in \\secref{sample} we discuss the selection   of    our   $z\\!\\simeq\\!4,5$   and    $6$   samples.    In \\secref{analysis}  we  describe  our  data  analysis,  which  includes accurately  measuring the 1$\\sigma$  sky-subtraction error,  the image stacking method to generate  mean surface-brightness profiles, and our test of  its reliability.  In \\secref{results} we  present and discuss our  results in terms  of the  average surface-brightness  profiles of $z\\!\\simeq\\!4\\!-\\!6$ galaxies, and  in \\secref{conclusion} we conclude with a summary of our results.  Throughout  this paper we  refer to  the \\emph{HST}/ACS  F435W, F606W, F775W,  and F850LP  filters as  the \\bb-,  \\vv-, \\ii-,  and \\zz-bands, respectively.  We assume a \\emph{Wilkinson Microwave Anisotropy Probe} (WMAP)  cosmology   of  $\\Omega_m$=0.24,  $\\Omega_{\\Lambda}$=0.76  and \\Ho=73~km~s$^{-1}$~Mpc$^{-1}$, in  accord with the  most recent 3-year WMAP results  of \\citet{sper07}.  This  implies a current age  for the Universe  of 13.65~Gyr.   All magnitudes  are given  in the  AB system \\citep{oke83}. ",
        "conclusions": "\\label{results}  \\figref{fig8} shows that the  mean surface brightness profiles deviate significantly   from    an   inner   $r^{1/n}$    profile   at   radii $r\\!\\gtrsim$0\\arcspt  27--0\\arcspt  35,   depending  somewhat  on  the redshift bin.   These deviations appear real, with  the break/point of departure located \\cge  1.5--2~mag above the 1$\\sigma$ sky-subtraction error and above  the PSF-wings.  In the following,  we discuss several possible explanations for the observed shapes of our composite surface brightness profiles.  \\subsection{Galaxies with Different Morphologies}  Our test  on nearby galaxies  (\\figref{fig9}) shows that, if  we stack many  galaxies with different  morphologies (early-type,  late-type or spiral galaxies),  it is possible  to get a slope-change  (`break') in the average surface brightness  profile. \\citet{ravi06} find that 40\\% of the brighter LBGs at $2.5\\!<\\!z\\!<\\!5$ have light profiles close to exponential, as seen for disk  galaxies, and only $\\sim$30\\% have high $n$,  as seen  in  nearby  spheroids.  They  also  find a  significant fraction ($\\sim$30\\%) of galaxies with light profiles \\emph{shallower} than  exponential, which appear  to have  multiple cores  or disturbed morphologies, suggestive  of close  pairs or on-going  galaxy mergers. Therefore,  if  galaxies at  $z\\!\\simeq\\!4\\!-\\!6$  have  a variety  of morphological types, then the  shape of the average surface brightness profile that we  see may be due to the stacking  of different types of galaxies.   Therefore, we find  that the  exponential and  the flatter profiles  found by  \\citet{ravi06} for  galaxies  at $2.5\\!<\\!z\\!<\\!5$ also apply to higher redshifts ($z\\!\\ge\\!5$).  Also, we believe that it is  more likely that the high redshift, faint galaxy population consists primarily  of small galaxies with late-type morphologies   and   with  sub-L$^{*}$   luminosities,   as  seen   at $z\\!\\simeq\\!2\\!-\\!3$     \\citep{driv95,driv98}.      So     if     the $z\\!\\simeq\\!4\\!-\\!6$  population consists of  such a  late-type galaxy population, then the slope-change in  the light profiles is likely not the result of co-adding images of objects with disparate morphological types.  \\subsection{Central Star Formation/Starburst}  \\emph{HST} optical  images of galaxies  at $z\\!\\simeq\\!4\\!-\\!6$ sample their rest-frame UV ($\\sim$1200 \\Ang), where the contribution from the actively  star-forming regions (very  young, massive  stars) dominates the   UV-light.    \\citet{hath07}   have   shown  that   galaxies   at $z\\!\\simeq\\!5\\!-\\!6$ are  high redshift starbursts  and these galaxies have similar starburst intensity limit as local starbursting galaxies. Therefore, it  is possible that galaxies  at $z\\!\\simeq\\!4\\!-\\!6$ have centrally concentrated star  formation or starburst.  This possibility is  based  on three  key  assumptions: (1)  most  of  the galaxies  at $z\\!\\simeq\\!4,5,6$     are    intrinsically     later-type    galaxies \\citep{driv98,stei99}; (2)  the Spectral Energy  Distribution (SED) of these galaxies at $z\\!\\simeq\\!4,5,6$ are dominated by early A- to late O-type stars, respectively;  and (3) there are no  old stars with ages at  $z\\!\\simeq\\!4\\!-\\!6$  greater  than  2-1 Gyr  in  WMAP  cosmology, respectively.  \\citet{hunt06}   studied  azimuthally   averaged   surface  photometry profiles for large sample of nearby irregular galaxies. They find some galaxies have double exponentials that  are steeper (and bluer) in the inner  parts compared to  outer parts  of the  galaxy.  \\citet{hunt06} discuss that this  type of behavior is expected  in galaxies where the centrally  concentrated  star  formation  or  starburst  steepens  the surface brightness profiles  in the center. If that  is the case, then one might expect a better correlation between the break in the surface brightness profiles and changes  in color profiles. Unfortunately, for our sample  of galaxies at  $z\\!\\simeq\\!4\\!-\\!6$, we don't  have high- resolution  restframe \\uu\\bb\\vv  color information.   The  objects are generally too  faint for \\emph{Spitzer Space Telescope},  and hence we cannot  confirm  or reject  this  possibility  for  the shape  of  our composite surface brightness profiles.  \\subsection{\\boldmath {Limits to Dynamical Ages for $z\\!\\simeq\\!4,5,6$ Objects}}\\label{ages}  The  average compact $z\\!\\simeq\\!4\\!-\\!6$  galaxy is  clearly extended with respect  to the ACS PSFs  (\\figref{fig8}), and is best  fit by an exponential   profile   ($n\\!<\\!2$)  out   to   a   radius  of   about $r\\!\\simeq$0\\arcspt   35,   0\\arcspt    31,   and   0\\arcspt   27   at $z\\!\\simeq\\!4,5$ and $6$, respectively.  The apparent progression with redshift is  noteworthy.  The  radius at which  the profile  starts to deviate  from  $r^{1/n}$ (in  this  case  at  radius $r$\\cge  0\\arcspt 35--0\\arcspt 27) may be an  important constraint to the dynamical time scale of the system,  as discussed in \\secref{introduction}.  If this argument is valid,  then we can estimate limits  to the dynamical ages of $z\\!\\simeq\\!4,5,6$ galaxies as follows.  In WMAP cosmology,  a radius of $r$\\cge 0\\arcspt  35 at $z\\!\\simeq\\!4$ corresponds   to  $r$\\cge   2.5   kpc.   The   dynamical  time   scale \\citep[e.g.,][]{binn87},  $\\tau_{dyn}$,  goes  as  $\\tau_{dyn}$  =  $C r^{3/2} \\!/\\!\\sqrt{G\\,M}$, where the constant $C=\\pi/2$. For a typical dwarf  galaxy mass  range of  $\\sim10^9\\!-\\!10^8$\\Msun  inside $r$=2.5 kpc,  we  infer  that  the  limits  to  the  dynamical  age  would  be $\\tau_{dyn}\\simeq$ 90--290  Myr, which is the lifespan  expected for a late-type B-star.  This means that the last major merger that affected this  surface brightness  profile  and that  triggered its  associated starburst  may  have occurred  $\\sim$0.20  Gyr before  $z\\!\\simeq\\!4$, ---assuming that the star-formation wasn't spontaneous, but associated with some accretion or a merging event.  Table 2 shows  the break-radius and inferred limits  to dynamical ages for the $z\\!\\simeq\\!4\\!-\\!6$ objects.  At $z\\!\\simeq\\!5$, we find that the  limits   to  dynamical   age  at  the   break  radius   would  be $\\tau_{dyn}\\simeq$ 70--210  Myr, which is the lifespan  expected for a mid B-star,  while at $z\\!\\simeq\\!6$,  $\\tau_{dyn}\\simeq$ 50--150 Myr, which is the lifespan expected for a late O--early B-star.  This means that  the last  major merger  that affected  these  surface brightness profiles at  $z\\!\\simeq\\!5$ and $6$ and that  triggered its associated starburst  may  have occurred  $\\sim$0.14  and  $\\sim$0.10 Gyr  before $z\\!\\simeq\\!5$ and $6$, respectively.  \\begin{deluxetable}{cccc} \\tablewidth{0pt} \\tablecaption{Dynamical Ages for $z\\!\\simeq\\!4-6$ objects in the HUDF\\label{table2}} \\tablenum{2} \\tablehead{\\colhead{Redshift} & \\colhead{``Break'' Radius$^a$} & \\colhead{``Break'' Radius$^b$} & \\colhead{Dynamical Age$^c$} \\\\ \\colhead{$z$} &      \\colhead{(arcsec)} &   \\colhead{(kpc)} &      \\colhead{($\\tau_{dyn}$)} }  \\startdata 4 & 0.35 & 2.5 & 0.09--0.29 Gyr \\\\ 5 & 0.31 & 2.0 & 0.07--0.21 Gyr \\\\ 6 & 0.27 & 1.6 & 0.05--0.15 Gyr \\\\ \\enddata  \\tablenotetext{a}{From composite surface brightness profiles (\\figref{fig6} and \\figref{fig8}).} \\tablenotetext{b}{Radius in kpc corresponding to radius in arcsec at given redshift.} \\tablenotetext{c}{If ``break radius'' interpreted as indicator of dynamical age.} \\end{deluxetable}  The  dynamical time is  a lower  limit to  the actual  time available, since it  assumes matter  starts from rest.   Any angular  momentum at start will increase the available time.  The best-fit SED age from the GOODS \\emph{HST} and \\emph{Spitzer} photometry on some of the brighter of  these objects  --- using  \\citet{bruz03} templates  --- is  in the range  of about  $\\sim$150--650  Myr \\citep{yan05,eyle05,eyle07},  the lower end  of which is consistent  with our limits  to their dynamical age  estimates, while  the  somewhat  larger SED  ages  could also  be affected by the onset of  the AGB in the stellar population increasing the observed  \\emph{Spitzer} fluxes and  hence possibly overestimating ages \\citep{mara05}.   Our age estimates  for $z\\!\\simeq\\!4\\!-\\!6$ are consistent  with the trend  of SED  ages suggested  for $z\\!\\simeq\\!7$ \\citep{labb06}.  It  is noteworthy that, given  the uncertainties, the two  independent  age  estimates  are  consistent. If  our  limits  to dynamical age  estimates for the  image \\emph{stacks} are  thus valid, they are consistent with the SED ages, and point to a consistent young age for these objects.  Furthermore, the  presence of young, massive late  O--early B-stars at $z\\!\\simeq\\!6$ has implications for  the reionization of the universe. From observations of the  appearance of complete Gunn-Peterson troughs in the  spectra of $z\\!\\ga\\!5.8$  quasars \\citep{fan06}, we  know that the epoch of reionization had ended by $z\\!\\simeq\\!6$.  From the steep ($\\alpha$=--1.8)  faint-end  slope   of  the  luminosity  function  of $z\\!\\simeq\\!6$  galaxies, \\citet{yan04a,yan04b}  concluded  that dwarf galaxies,   and   not  quasars,   likely   finished  reionization   by $z\\!\\simeq\\!6$.   Should  the present  interpretation  of their  light profiles  be correct,  then it  would appear  to add  support  to this picture, in the  sense that such objects are  dominated by B-stars and did  not   start  their  most  recent  major   starburst  long  before $z\\!\\simeq\\!6$."
    },
    "0710/0710.4909_arXiv.txt": {
        "abstract": "The origin of galactic cosmic rays is one of the most interesting unsolved problems in astroparticle physics. Experimentally, the problem is attacked by a multi-disciplinary effort, namely by direct measurements of cosmic rays above the atmosphere, by air shower observations, and by the detection of TeV $\\gamma$ rays.  Recent experimental results are presented and their implications on the contemporary understanding of the origin of galactic cosmic rays are discussed. ",
        "introduction": "The Earth is permanently exposed to a vast flux of highly energetic particles, fully ionized atomic nuclei from outer space. The extraterrestrial origin of these particles has been demonstrated by V. Hess in 1912 \\cite{hess} and he named the particles \"H\\\"ohenstrahlung\" or \"Ultrastrahlung\". In 1925 R. Millikan coined the term \"Cosmic Rays\". They have a threefold origin.  Particles with energies below 100~MeV originate from the Sun.  Cosmic rays in narrower sense are particles with energies from the 100~MeV domain up to energies beyond $10^{20}$~eV. Up to several 10~GeV the flux of the particles observed is modulated by the 11-year cycle of the heliospheric magnetic fields.  Particles with energies below $10^{17}$ to $10^{18}$~eV are usually considered to be of galactic origin. A proton with an energy of $10^{18}$~eV has a Larmor radius $r_L=360$~pc in the galactic magnetic field ($B\\approx3$~$\\mu$G). This radius is comparable to the thickness of the galactic disc and illustrates that particles at the highest energies can not be magnetically bound to the Galaxy. Hence, they are considered of extragalactic origin.  The energy density can be inferred from the measured differential energy spectrum $dN/dE$ \\cite{halzenrhoe} \\begin{equation} \\rho_E=\\frac{4\\pi}{c}\\int \\frac{E}{\\beta} \\frac{dN}{dE} dE , \\end{equation} where $\\beta c$ is the velocity of particles with energy $E$.  For galactic cosmic rays the major contribution to the total energy density originates from particles with energies around 1~GeV. Such particles are strongly influenced by the heliospheric magnetic fields. Outside the heliosphere, in the local interstellar environment an energy density $\\rho_E^{LIS}=1.1$~eV/cm$^3$ is obtained (\\rref{gaisserbuch} p.\\,12).  The parameterization of the measured galactic cosmic-ray flux according to the \\modell \\cite{pg} results in a density $\\rho_E^{gal}=0.43$~eV/cm$^3$. This is the measurable energy density at Earth (for an average modulation parameter $M=750$~MeV, see (1) in \\rref{pg}). This implies that less than half of the energy flux can be registered directly at Earth.  In this article we will give an overview on recent experimental results and their implications on the contemporary understanding of the origin of galactic cosmic rays. As space is limited here, the reader may also consider further recent reviews by the author \\cite{pg,origin,cospar06,vulcano,ecrsreview}.  Progress in the understanding of the origin of galactic cosmic rays emerged mainly from observations in three complementary disciplines. The direct measurement of cosmic-ray particles above the atmosphere in outer space and on stratospheric balloons (\\sref{direct}). At energies exceeding $10^{15}$~eV the steeply falling cosmic-ray energy spectrum requires experiments with large detection areas exposed for long times, at present, only realized in ground based installations (\\sref{indirect}). With such detectors extensive air showers are detected, which originate in interactions of high-energy particles in the atmosphere (\\sref{eas}). And, finally, the observation of TeV $\\gamma$-rays (\\sref{gamma}). ",
        "conclusions": "\\label{outlook} In the last decade our understanding of the origin of galactic cosmic rays has been significantly improved by multidisciplinary efforts, combining key observations of the direct and indirect measurements of cosmic rays as well as the detection of $\\gamma$-rays. The observations by the HESS $\\gamma$-ray telescope give clear hints that hadronic particles are accelerated in SNR. The data are compatible with a model of first order Fermi acceleration at strong shock fronts. The particles propagate in a diffusive process through the Galaxy. Parameters of the propagation models have been constraint by direct measurements above the atmosphere. The KASCADE experiment has shown that the energy spectra of the light elements exhibit a cut-off structure, while the spectra of heavier elemental groups follow power laws to higher energies. The observed spectra seem to be compatible with the assumption of power laws and a cut-off energy proportional to the nuclear charge. This implies that the knee in the all-particle energy spectrum is caused by a cut-off of the light elements. The shape of the all-particle spectrum at higher energies is then determined by the subsequent cut-offs of all elemental species in cosmic rays. Most likely, the astrophysical origin of the knee is a combination of the maximum energy reached in the acceleration process and leakage from the Galaxy during propagation \\cite{prop}. More exotic ideas about the cause of the knee are most likely excluded \\cite{cospar06,ecrsreview}.  In conclusion this gives a qualitative 'standard picture' of the origin of galactic cosmic rays. However, several details remain unclear and a precise quantitative description of all aspects of the acceleration and propagation mechanisms is still missing. Among the open questions are: \\\\ It is not clear how to precisely match the spectral indices observed at Earth to the spectra at the sources, being compatible with the TeV $\\gamma$-ray observations and the modification of the spectral slope during propagation.\\\\ A precise astrophysical interpretation of air shower data is at present limited by the understanding of hadronic interactions in the atmosphere, thus the exact shape of the energy spectra at their corresponding knees is unknown.\\\\ Contemporary assumptions on the parameters of cosmic-ray propagation models yield anisotropies in the arrival directions not observed by air shower experiments \\cite{ptuskinaniso,prop}.  In the near future experiments like TRACER aim to reveal details of the cosmic-ray propagation for energies approaching the knee \\cite{muellermerida}. Experiments at the LHC probing the extreme forward direction of phase space will improve the description of high-energy hadronic interactions \\cite{engelpylos}. Also the exploration of the end of the galactic component and the transition to extragalactic particles in the energy range from $10^{17}-10^{18}$~eV will be of importance. Key experiments in this energy region are KASCADE-Grande (a 0.5~km$^2$ extension of the KASCADE experiment) \\cite{grande}, Ice Cube/ Ice Top (a 1~km$^2$ air shower experiment and neutrino telescope at the South Pole) \\cite{icecube,icetop}, and HEAT/AMIGA (a 25~km$^2$ extension of the Auger experiment to lower energies) \\cite{klagesmerida}.  Around $10^{18}$~eV two features appear in the all-particle energy spectrum. The second knee at $E_{2nd}\\approx400$~PeV$\\approx92\\times E_k$, where the spectrum exhibits a steepening to $\\gamma\\approx-3.3$, and the ankle at about 4~EeV, above this energy the spectrum seems to flatten again to $\\gamma\\approx-2.7$. The region around 4~EeV is sometimes also called \"the dip\" in the spectrum.\\\\ A possible cause for the second knee is the end of the galactic component, when all elements successively have reached their cut-off energies, the latter being proportional to their nuclear charge \\cite{pg,prop}. If one assumes that ultra-heavy elements (heavier than iron) play an important role to understand the second knee, the factor of 92 between the energies of the knee and the second knee can be easily understood as the nuclear charge of the heaviest elements in the periodic table.\\\\ The dip is proposed to be caused by electron-positron pair production of cosmic rays on photons of the cosmic microwave background \\cite{berezinskydip}.  The investigation of the transition region and a precise measurement of the galactic all-particle spectrum is also important for an estimate of the energy content of extragalactic cosmic rays.  For example, the extragalactic component needed according to the \\modell to sustain the observed all-particle flux at highest energies has an energy density of $\\rho_E^{exg}=3.7\\cdot10^{-7}$~eV/cm$^3$.  A promising rediscovered technique for the exploration of cosmic rays from the transition region to highest energies is the detection of radio signals from air showers \\cite{allanrev}. Most likely the emission mechanism is coherent synchrotron radiation of electrons with energies around the critical energy (85~MeV) deflected in the magnetic field of the Earth (geosynchrotron radiation) \\cite{huegefalcke}. The LOPES experiment, registering showers in coincidence with the KASCADE-Grande experiment has demonstrated the feasibility of this approach \\cite{radionature}. Radio detection of air showers is also pursued in the CODALEMA experiment \\cite{codalema} and within the LOFAR radio telescope \\cite{lofar}.  Also first radio pulses from air showers have been recorded with antennae set up at the southern site of the Pierre Auger Observatory \\cite{vdbergmerida}."
    },
    "0710/0710.3835_arXiv.txt": {
        "abstract": "We derive a simple consistency relation from the running of the tensor-to-scalar ratio. This new relation is first order in the slow-roll approximation. While for single field models we can obtain what can be found by using other observables, multi-field cases in general give non-trivial contributions dependent on the geometry of the field space and the inflationary dynamics, which can be probed observationally from this relation. The running of the tensor-to-scalar ratio may be detected by direct laser interferometer experiments. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.4412_arXiv.txt": {
        "abstract": "We investigate numerically the long-time behavior of balanced Alfv\\'en wave turbulence forced at intermediate scales. Whereas the usual constant-flux solution is found at the smallest scales, two new scalings are obtained  at the forcing scales and at the largest scales of the system. In the latter case we show, in particular, that the spectrum evolves first to a state determined by  Loitsyansky invariant and later a state close to the thermodynamic equipartition solution predicted by wave turbulence. The astrophysical implications for galactic magnetic field generation are discussed. ",
        "introduction": "Turbulence flows are ubiquitous in astrophysical environments from the solar wind (Goldstein et al., 1999; Galtier, 2006), to interstellar (Elergreen et al., 2004; Scalo et al., 2004), galactic and even intergalactic media (Govoni et al., 2006). At the larger scales, signatures of astrophysical turbulence are found, in particular, in the magnetic field measurements whose origin remains one of the major challenging problems (Pouquet, 1993; Widrow, 2002; Brandenburg, 2005). In this paper, we emphasize a new mechanism for generating large-scale magnetic field. It is based on the resonant interactions of shear-Alfv\\'en waves in a turbulent medium permeated by a strong external magnetic field and influenced by an external forcing at intermediate scales. In this situation, both large-scale kinetic and magnetic energies may be produced by mainly non local interactions. This scenario, although very simple, may be relevant for describing  re-generation (i.e. maintenance) of large-scale galactic fields, e.g. in our galaxy where energy is injected at intermediate scales by stellar winds and supernovae explosions on scales $10-100$pc (Ferriere et al., 2004). ",
        "conclusions": "We have seen that the large-scale magnetic field is produced on a time-scale which is about $10^4$ larger than the time-scale needed to establish the small-scale energy cascade solution. We believe that the final distribution in the large-scale part should correspond to the thermodynamic energy equipartition $k_\\perp^{1}$ even though it would take an extremely long computing time for formation of this spectrum which we were not able to achieve. Instead, we observed formation of a shallower long-term scaling at the large scales, $k_\\perp^{0.6}$, which we believe to be transient.  Because of the energy cascade to the smallest dissipative (resistive and viscous) scales, the small-scale part of the spectrum is important for understanding the total energy dissipation rate. On the other hand, it is this large-scale part of the spectrum that contains most of the wave (magnetic and kinetic) energy itself. One can view it as a sort of powerful energy storage which is charged extremely slowly and in a very inefficient way (because most of the charging energy is wasted via the energy cascade). For example, if the forcing is concentrated in a thin spherical shell with wavenumber lengths between $k$ and $k + \\Delta k$ then the final energy at large scales (assuming the energy equipartition) will be $k/\\Delta k$ times greater than the energy at the forcing scales.  The mechanism of generation of large-scale magnetic fields via nonlinear transfer from an energy source at smaller scales may be relevant, in a very qualitative sense, to maintenance of large-scale galactic fields by small-scale sources provided by supernovae events. To pursue this line further, however, one would have to consider a more realistic thin disk geometry, instead of a simple homogeneous external field considered in the present paper, which could be an interesting future project."
    },
    "0710/0710.4138_arXiv.txt": {
        "abstract": "We present a detailed analysis of week-long simultaneous observations of the blazar \\mrk at 2--60~keV \\xrays (\\rxtenosp) and TeV \\grays (Whipple and HEGRA) in 2001. Accompanying optical monitoring was performed with the Mt. Hopkins 48\" telescope. The unprecedented quality of this dataset enables us to establish the existence of the correlation between the TeV and \\xray luminosities, and also to start unveiling some of its characteristics, in particular its energy dependence, and time variability. The source shows strong variations in both \\xray and \\gray bands, which are highly correlated.  No evidence of a \\xraynosp/\\gray interband lag $\\tau$ is found on the full week dataset, with $\\tau\\lesssim3$\\,ks. A detailed analysis of the March 19 flare, however, reveals that data are \\textit{not} consistent with the peak of the outburst in the 2--4\\,keV \\xray and TeV band being simultaneous.  We estimate a $2.1\\pm0.7$\\,ks TeV lag. The amplitudes of the \\xray and \\gray variations are also highly correlated, and the TeV luminosity increases more than linearly with respect to the \\xray one. The high degree of correlation lends further support to the standard model in which a unique electrons population produces the \\xrays by synchrotron radiation and the \\gray component by inverse Compton scattering. However, the finding that for the individual best observed flares the \\gray flux scales approximately quadratically with respect to the \\xray flux, poses a serious challenge to emission models for TeV blazars, as it requires rather special conditions and/or fine tuning of the temporal evolution of the physical parameters of the emission region. We briefly discuss the astrophysical consequences of these new findings in the context of the competing models for the jet emission in blazars. ",
        "introduction": "\\label{sec:intro}  \\object[Mrk 421]{\\mrk} is the brightest BL Lac object in the \\xray and UV sky and the first extragalactic source detected at TeV energies \\citep{punch92_tev_mkn421}. Like most blazars, its spectral energy distribution shows two smooth broad band components (\\eg \\citealp*{sambruna96_sed,umu97_review}; \\citealp{fossati98_sed}). The first one extends from radio to \\xrays with a peak in the soft to medium \\xray range; the second one extends up to the GeV to TeV energies, with a peak presumed to be around 100~GeV. The emission up to \\xrays is thought to be due to synchrotron radiation from high-energy electrons, while the origin of the luminous \\gray radiation is more uncertain. Possibilities include inverse Compton scattering of synchrotron (synchro self-Compton, SSC) or ambient photons (external Compton, EC) off a single electron population thus accounting for the spectral ``similarity'' of the two components (\\eg \\citealp{macomb95_mkn421,mastichiadis_kirk97,tavecchio_tev_98,% mgc92_3c279,sbr94_ec,dermer_etal92}). Alternative ``hadronic'' models produce \\gray from protons, either directly (proton synchrotron) or indirectly (\\eg synchrotron from a second electron population produced by a cascade induced by the interaction of high-energy protons with ambient photons) \\citep{mucke_etal_2003_SPB_model,bottcher_reimer_2004}. The synchrotron proton scenario may be more favorable for objects like \\mrk \\citep{mucke_etal_2003_SPB_model}, because of the lower density of the diffuse photon fields necessary for processes like pion photoproduction to be effective. Moreover, it is generally true that ``hadronic'' models need a higher level of tuning in order to reproduce the observed highly correlated \\xraynosp/\\gray variability. Hence, in this paper we have not addressed this class of models, and instead focused our limited modeling effort on the pure SSC model.  All of the above models of the \\gray emission from blazars have all had some degree of success in reproducing both single-epoch spectral energy distributions and their relative epoch-to-epoch changes (\\citealp{vonmontigny95_egret}, \\citealp{gg98_sedt}). These favor the SSC model in \\mrk because it is a BL Lac object for which the ratio of thermal (accretion disk and broad line region) and synchrotron photons is $\\sim0.1$, indicating that the EC mechanism is not important. Detailed modeling of ``blue'' BL~Lacs finds that one-component SSC model can generally account for the time-averaged spectral energy distributions \\citep{gg98_sedt}. Some data sets seem to require modifications of the simple model, introducing either multiple SSC components or additional external seed photons (\\eg \\citealp{blazejowski_etal_2005}).  We can, however, further decrease the degeneracy among proposed physical models by taking advantage of blazars' rapid, large-scale time variability with simultaneous \\xraynosp/TeV monitoring (\\eg \\citealp{tavecchio_tev_98,maraschi99_sax_mkn421,henric01_mrk421_x_tev}). Different models produce emission at a given frequency with particles of different energies, cooling times, and cross sections for different processes (\\eg \\citealp{blumenthal_gould70,coppi_blandford90}), and thus are in principle  distinguishable \\citep*{henric02_timedependent}. For example, the SSC model predicts nearly simultaneous variations in both the synchrotron and Compton components (see however \\S\\ref{sec:conclusions}), while other models predict more complicated timing (\\eg \\citealp{umu97_review}).  With the possible exception of the \\xray and \\gray data taken on Mrk\\,501, the multiwavelength observations on which we base our inferences, have often undersampled the intrinsic variability timescales \\citep{buckley96_mrk421,petry00_mrk501,2001ApJ...563..569T,maraschi99_sax_mkn421} and lack a sufficiently long baseline to make a quantitative assertion about the statistical significance of a correlation. Recently, there has even been evidence of an ``orphan'' TeV flare for the Blazar 1ES\\,1959+650 \\citep{krawczynski_etal_2004_pks1959}; a transient $\\gamma$-ray event that was not accompanied by an obvious \\xray flare in simultaneous data.  For what concerns \\mrknosp, there have been regular multiwavelength campaigns in the last several years, planned with an observing strategy focusing on month-long timescales \\citep{blazejowski_etal_2005,rebillot_etal_2006}, and in turn a relatively sparse time sampling (typically one \\rxte snapshot per night, plus binned \\rxtenosp/ASM light curves). These campaigns showed that \\xray and \\gray brightnesses vary ``in step'' and there is certainly a ``loose'' correlation, and also raised some questions about the need of considering additional components to account for the spectra \\citep{blazejowski_etal_2005}, or very high Doppler factors \\citep{rebillot_etal_2006}.  The March 2001 campaign remains the experiment with the highest density coverage at both \\xray and TeV energies, and thus the best dataset to address questions concerning the characteristics of the variability of the two spectral energy distribution (SED) components. Moreover, the brightness state achieved during the week of the observations was unprecedented and it remains unparalleled. Preliminary results were presented in \\citet{jordan01_icrc,fossati03_veritas03}.  In this paper we present the summary of the multiwavelength observations, with a particular focus on the correlated \\xraynosp/TeV variability, also including the TeV data taken by the HEGRA (High Energy Gamma Ray Astronomy) telescope. An account of the HEGRA March 2001 observations was published by the HEGRA collaboration \\citep{aharonian02_mrk421_spectral_var}. \\citet{giebels_etal_2007_mkn421} report on additional TeV observations with the CAT observatory (Cerenkok Array at Themis), simultaneous with the \\rxte data, not included in this paper.  The \\rxte observations, timing and spectral properties are fully presented by Fossati \\etal (2008, in preparation; hereafter F08).  The paper is organized as follows:  the relevant information about the \\xray and TeV observations, and data reduction is given in \\S\\ref{sec:data}. The observational findings are presented in \\S\\ref{sec:analysis}. and discussed in \\S\\ref{sec:conclusions} in the context of the synchrotron self--Compton model.  Section~\\ref{sec:conclusions} summarizes the conclusions. ",
        "conclusions": "\\label{sec:conclusions}  The correlation between the variations in the \\xray and TeV bands is confirmed with unprecedented detail, supporting the idea that the same electron distribution, in the same physical region, is responsible for the emission in both energy bands.  However the details of these findings pose a serious challenge to the emission models. Here we would like to sketch a few selected outstanding issues raised by the correlated variability.  We refer to a forthcoming paper for an in-depth analysis comprising more extensive modeling.  \\begin{itemize} \\item[$\\bullet$] If modeled within the realm of standard values for magnetic field, Doppler factor, source size, the IC scattering responsible for the observed TeV emission occurs in the Klein-Nishina regime. This means that with \\xray and \\gray observations, although we seem to be observing regions of the spectrum that are very similar to each other for what concerns their position with respect to the SED peaks, we are not tracking the evolution of the same electrons (and photons). The extent of the phase and amplitude correlation of the \\xray and \\gray variations is however remarkable, and this sets broad constraints on the characteristics of the processes responsible/governing the variability, \\eg acceleration/injection of particles, dominant cooling cause.  \\item[$\\bullet$] In this context, the observation in K-N regime of a quadratic relationship between synchrotron and IC variations (which would be naturally produced in the Thomson regime because of the effectiveness of self-Compton) constrains the electron spectrum variations to occur over an energy band broad enough to affect also the IC seed photons. Moreover, this variation must be essentially ``achromatic'' (\\ie just a change in normalization), otherwise the extra energy-dependent factor would produce an observable effect.  \\item[$\\bullet$] The observation that the flux--flux path of the better observed flares decay follows closely the bursting path introduces a further complication. If the flare decay is governed by the cooling of the emitting electrons we do not expect the quadratic relationship to hold during the decaying phase. In fact, given the energy-dependent nature of synchrotron (and IC) cooling, with $\\tau_{cool}\\sim E_{ph}^{-1/2}$, the $50$\\,eV seed photons cool on a longer timescale, \\eg $\\sim$10 times longer than the timescale for photons observed at $\\gtrsim3$\\,keV. The possibility that also the electrons contributing to the bulk of the TeV emission have lower energy than those observed in \\xraynosp, compounds the problem. This means that during the flare decay the \\xray and $\\gamma$-ray brightnesses should follow something like a linear relationship, because the IC (TeV) emission will just reflect the evolution of the electron spectrum, scattering a ``steady'' seed photon field.  Plain radiative cooling does not seem to match these observations. A viable mechanisms explaining the flares evolution should allow the concurrent cooling of a broad portion of the electron distribution. On the other hand, brightness variations are accompanied by large spectral changes, and in most cases they are very suggestive of acceleration --or injection-- of the higher energy end of the electron population.  \\item[$\\bullet$] Another recurring discrepancy between data and simple one-zone SSC modeling is that of the TeV spectral shape, which is often harder than model predictions. As we illustrated in the previous section the Klein-Nishina effect plays an important role in this respect. A more careful analysis is warranted, but it is worth noting that one alternative option for addressing this problem is that of considering the effect of additional IC components, off photons external to the blob. \\cite{blazejowski_etal_2005} showed that a multi-component model seems to be required to fit the 2003$-$2004 observations, and it might mitigate the discrepancy in the TeV spectrum. \\cite{gg_ft_mc_2005_structure_jets} discuss the effect of including the effect of the radiation emitted by the putative lower Lorentz factor outer layer of the jet, namely as source of additional seed photons for IC.  \\item[$\\bullet$] Exotic scenarios could address some of these issues, namely by allowing the IC scattering to occur in the Thomson regime, and hence the self-Compton to be effective, thus reinstating the close relationship between the photons and electrons tracked by \\xray and \\gray observations. One such scenario that we discussed briefly would call for very high values of the Doppler beaming factor, that would reduce the intrinsic energies at play, higher magnetic field and very small size of the emission region. In fact recently there has been some interest for less conventional modeling of TeV blazars (\\eg \\citealt{henric01_mrk421_x_tev,rebillot_etal_2006} for \\mrknosp). However, high Doppler factor scenarios raise a series of new issues, or re-open some that have been settled for the more traditional model. First of all, the fact that the beaming cone of the radiation emitted by such a fast blob is going to be much narrower has to be reconciled with the population statistics of blazars and radio-galaxies, their putative parent population (\\citealt{up95_review}). One way of doing it would be to imagine that the jet comprises a very large number of small high Lorentz factor blobs, fanning out filling a wider cone, with aperture consistent with the unification statistics. This would constitute a quite radical change in the jet structure, from the current one where emission is thought to come from internal shocks. There are also implications concerning the statistical properties of the variability, which would likely be due to the combination of the bursting of different blobs, most likely uncorrelated.  \\end{itemize}  The richness and depth of the \\xray and \\gray data of the March 2001 campaign presented in this paper raise the bar for models. The aggregate characteristics illustrated here already challenge the simple traditional SSC model, and the SED-snapshot approach. In order to answer the questions raised by these observations it is of paramount importance to exploit fully the time--axis dimension into the modeling and take a dynamical approach.  The data time density and brightness (and so statistics) are unparalleled, enabling time resolved spectroscopy on timescales of the order of the physically relevant ones, hence allowing to model the phenomenology self-consistently minimizing the need (freedom) to make assumptions as to how to connect spectra taken at different times.  It is likely that this dataset is going constitute the best benchmark for time dependent modeling for some time, despite the great progress made by ground based TeV atmospheric Cherenkov telescopes in the last few years, because of the difficulty of securing long uninterrupted observations with Chandra and XMM-Newton."
    },
    "0710/0710.4248_arXiv.txt": {
        "abstract": "The {\\it INTEGRAL\\/} satellite, which studies the Universe in the hard X-ray and soft Gamma-ray domain, has been operational for 5 years now. The X-ray telescopes, which use the coded mask technique, provide unprecedented spectral and imaging resolution. This led to a number of discoveries, such as the distribution of diffuse emission in the Galaxy, the discovery of highly absorbed sources and fast X-ray transients in the Galactic Plane, localization of $\\sim50$ Gamma-ray bursts, and the resolution of the cosmic X-ray background around its peak at 30 keV. About 300 previously known X-ray sources have been detected and in addition more than 200 new sources have been discovered. {\\it INTEGRAL\\/} provides spectra starting at 3 keV and ranging up to several hundred keV. This article gives a brief overview about the major discoveries of {\\it INTEGRAL}. ",
        "introduction": "ESA's {\\it INTEGRAL\\/} space mission \\cite{INTEGRAL} hosts two major hard X-ray instruments, IBIS and SPI, both coded-mask telescopes. IBIS \\cite{IBIS} provides imaging resolution of 12 arcmin, while SPI \\cite{SPI} is optimized for spectroscopy. Both instruments operate at energies from 15 keV up to several MeV. Co-aligned with these main instruments are the two X-ray monitors JEM-X \\cite{JEMX}, which provides spectra and images in the 3--30 keV band, and the optical camera OMC \\cite{OMC}, which provides photometry in the V filter. {\\it INTEGRAL} was launched on October 17, 2002 from Baikonur into a highly eccentric orbit with a perigee of $9,000$ km and an apogee of $150,000$ km, which avoids as much as possible the Earth's radiation belt and allows for un-interrupted observations of up to 3 days. % ",
        "conclusions": ""
    },
    "0710/0710.2552_arXiv.txt": {
        "abstract": "We expand our Bayesian Monte Carlo method for analyzing the light curves of gravitationally lensed quasars to simultaneously estimate time delays and quasar structure including their mutual uncertainties.  We apply the method to HE1104--1805 and QJ0158--4325, two doubly-imaged quasars with microlensing and intrinsic variability on comparable time scales.  For HE1104--1805 the resulting time delay of $\\Delta t_{AB} = t_A - t_B = 162.2_{-5.9}^{+6.3}$~days and accretion disk size estimate of $\\log(r_s / {\\rm cm}) = 15.7_{-0.5}^{+0.4}$ at $0.2\\mu$m in the rest frame are consistent with earlier estimates but suggest that existing methods for estimating time delays in the presence of microlensing underestimate the uncertainties.  We are unable to measure a time delay for QJ0158--4325, but the accretion disk size is $\\log(r_s / {\\rm cm}) = 14.9\\pm0.3$ at $0.3\\mu$m in the rest frame. ",
        "introduction": "Variability in lensed quasar images comes from two very different sources. Changes in the quasar's intrinsic luminosity are observable as correlated variability between images, while microlensing by the stars in the lens galaxy produces uncorrelated variability.  Measurements of the time delays between the lensed images from the correlated variability can be used to study cosmology (e.g. \\citealt{Refsdal1964} and recently \\citealt{Saha2006,Oguri2007}) or the distribution of dark matter in the lens galaxy \\citep[e.g.][]{Kochanek2006,Poindexteretal.2007,Vuissoz2007}. The microlensing variability can be used to study the structure of the quasar, the masses of the stars in the lens galaxy, and the stellar mass fraction near the lensed images \\citep{Schechter2002,Wambsganss2006}.  It is now possible to use microlensing to measure the correlation of accretion disk size with black hole mass \\citep{Morgan2007}, the wavelength dependence of the size of the accretion disk \\citep{Poindexter2007} or the differing sizes of the thermal and non-thermal X-ray emission regions \\citep{Pooley2007,Dai2007}.  The challenge is that most lensed quasars exhibit both intrinsic and microlensing variability.  To measure a time delay, one must successfully model and remove the microlensing variability such that only intrinsic variability remains. If the microlensing variability has a sufficiently low amplitude or long timescale, it can be ignored \\citep[e.g. PG1115-080,][]{Schechter1997}, but this is a dangerous assumption for many systems.  \\citet{Eigenbrod2005} found that for an $80$~day delay, adding microlensing perturbations with an amplitude of 5\\% (10\\%) to a light curve increased the uncertainty in the time delay by a factor of 2 (6). Existing time delay analyses for lenses with microlensing \\citep[e.g.][]{Paraficz2006,Kochanek2006,Poindexteretal.2007} depend on the intrinsic and microlensing variability having different time scales. These analyses also require that the microlensing variability can be modeled by a simple polynomial function.  This approach will clearly fail if the two sources of variability have similar time scales or if the microlensing variability cannot be easily parameterized.  \\begin{deluxetable}{lcccccc} \\tabletypesize{\\scriptsize} \\tablecaption{HST Astrometry and Photometry of QJ0158--4325 and HE1104--1805} \\tablehead{Lens & \\colhead{Component} &\\multicolumn{2}{c}{Astrometry} &\\multicolumn{3}{c}{Photometry}\\\\ \\colhead{} &\\colhead{} &\\colhead{$\\Delta\\hbox{RA}$} &\\colhead{$\\Delta\\hbox{Dec}$} &\\colhead{H=F160W} &\\colhead{I=F814W} &\\colhead{V=F555W} } \\startdata QJ0158--4325 & A  &$\\equiv 0$ &$\\equiv 0$ &$16.47\\pm0.03$ &$17.81\\pm0.04$ &$18.10\\pm0.13$\\\\ & B  &$-1\\farcs156\\pm0\\farcs003$ &$-0\\farcs398\\pm0\\farcs003$ &$17.27\\pm0.03$ &$18.62\\pm0.11$ &$18.91\\pm0.17$\\\\ & G  &$-0\\farcs780\\pm0\\farcs016$ &$-0\\farcs234\\pm0\\farcs006$ &$16.67\\pm0.13$ &$18.91\\pm0.06$ &$20.36\\pm0.18$\\\\ \\hline HE1104--1805 & A  &$\\equiv 0$ &$\\equiv 0$ &$15.91\\pm0.01$ &$16.40\\pm0.03$ &$16.92\\pm0.06$\\\\ & B  &$+2\\farcs901\\pm0\\farcs003$ &$-1\\farcs332\\pm0\\farcs003$ &$17.35\\pm0.03$ &$17.95\\pm0.04$ &$18.70\\pm0.08$\\\\ & G  &$+0\\farcs965\\pm0\\farcs003$ &$-0\\farcs500\\pm0\\farcs003$ &$17.52\\pm0.09$ &$20.01\\pm0.10$ &$23.26\\pm0.27$\\\\ \\enddata \\label{tab:hst} \\end{deluxetable}  In this paper, we present a new technique for simultaneously estimating the time delay and structure of lensed quasars that exhibit strong microlensing.  In essence, we assume a range of time delays and then determine the likelihood of the implied microlensing variability using the Bayesian Monte Carlo method of Kochanek (2004, see also \\citealt{Kochanek2007}). This allows us to estimate the time delays and the quasar structural parameters simultaneously and include the effects of both phenomena on the parameter uncertainties.   We apply the method to the two doubly-imaged lenses HE1104--1805 \\citep{Wisotzki1993} and QJ0158--4325 \\citep{Morgan1999}.  While HE1104--1805 has a well-measured time delay \\citep{Poindexteretal.2007}, the amplitude of the microlensing ($\\sim 0.05 \\; {\\rm mag \\; yr^{-1}}$ over the past decade) and the fact that it exhibits variability on the 6 month scale of the time delay suggest that it is close to the limit where microlensing polynomial fitting methods \\citep{Burud2001,Kochanek2006} will break down.  QJ0158--4325 clearly shows both correlated and uncorrelated variability, but the polynomial methods cannot reliably produce a time delay estimate. We describe the data and our models in \\S\\ref{sec:data}, our new approach in \\S\\ref{sec:analysis} and the application to the two systems in \\S\\ref{sec:results}.  In \\S~\\ref{sec:discussion}, we discuss the results and their limitations.  We assume a flat $\\Omega_0=0.3$, $\\Lambda_0=0.7$, $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$ cosmology and that the lens redshift of QJ0158--4325 is $z_l=0.5$.  Reasonable changes in this assumed redshift have negligible consequences for the results. ",
        "conclusions": "\\label{sec:discussion}  \\citet{Peng2006} used the width of the \\ion{C}{4}~($\\lambda 1549$\\AA) emission line to estimate the black hole mass $M_{BH}=2.37 \\times 10^9 {\\rm M_\\sun}$ in HE1104--1805 and the width of \\ion{Mg}{2}~($\\lambda 2798$\\AA) emission line to estimate the black hole mass $M_{BH}=1.6 \\times 10^8 {\\rm M_\\sun}$ in QJ0158--4325.  Using these black hole masses, the quasar accretion disk size - black hole mass relation of \\citet{Morgan2007} predicts source sizes at 2500\\AA~of $\\log(r_s/{\\rm cm}) = 15.9 \\pm 0.2$ for HE1104--1805 and $\\log(r_s/{\\rm cm}) = 15.2 \\pm 0.2$ for QJ0158--4325. If we scale our current disk size measurements to 2500\\AA~using the $R_\\lambda \\propto \\lambda^{4/3}$ scaling of thin disk theory and assume an inclination angle $\\cos i = 1/2$, we find $\\log(r_s/{\\rm cm}) = 15.9^{+0.4}_{-0.5}$ for HE1104--1805 and $\\log(r_s/{\\rm cm}) = 14.8 \\pm 0.3$ for QJ0158--4325, fully consistent with the predictions of the \\citet{Morgan2007} accretion disk size - black hole mass relation.  The mixing of intrinsic and microlensing variability in lensed quasar light curves can be a serious problem for estimating time delays \\citep[e.g.][]{Eigenbrod2005} and previous microlensing analyses have been restricted to lenses with known time delays. In HE1104--1805, which must be close to the limits of measuring time delays in the presence of microlensing, we confirm that the approach of fitting polynomial models for the microlensing works reasonably well.  However, the dependence of the delay on the assumed model was a warning sign that the formal errors on the delays were likely to be underestimates, as was recognized by \\citet{Poindexteretal.2007}. In our new, non-parametric microlensing analysis of HE1104--1805 we find a modestly longer delay of $162.2_{-5.9}^{+6.3}$~days that quantifies those concerns.   Estimates of the quasar accretion disk size are little affected by these small shifts in the time delay. In QJ0158--4325, the microlensing amplitude is larger relative to the intrinsic variability, and traditional methods for determining delays fail.  Our new method also fails to measure a delay, but it does allow us to measure the size of the quasar accretion disk despite the uncertainties in the time delay."
    },
    "0710/0710.0282_arXiv.txt": {
        "abstract": "A three fluid system describing the decay of the curvaton is studied by numerical and analytical means. We place constraints on the allowed interaction strengths between the fluids and initial curvaton density by requiring that the curvaton decays before nucleosynthesis while nucleosynthesis, radiation-matter equality and decoupling occur at correct temperatures. We find that with a continuous, time-independent interaction, a small initial curvaton density is naturally preferred along with a low reheating temperature. Allowing for a time-dependent interaction, this constraint can be relaxed. In both cases, a purely adiabatic final state can be generated, but not without fine-tuning. Unlike in the two fluid system, the time-dependent interactions are found to have a small effect on the curvature perturbation itself due to the different nature of the system. The presence of non-gaussianity in the model is discussed. ",
        "introduction": "The problem of determining the evolution of large scale perturbations in a background of a multi-component fluid system is central in modern day cosmology \\cite{Kodama:1985bj, Mukhanov:1990me, Wands:2000dp}. In such a system, the interactions between the different fluids are important in determining the evolution of the curvature perturbation \\cite{Malik:2004tf}. Examples of interacting fluid systems demonstrate the importance of such systems,as  most notably reheating at the end of inflation \\cite{Albrecht:1982mp,Den:1984tn,Kripfganz:1985mn,Bastero-Gil:2002xr,Dvali:2003em,Kofman:2003nx} and the curvaton scenario \\cite{Enqvist:2001zp,Lyth:2001nq,Moroi:2002rd,Moroi:2001ct,Lyth:2002my}.  In contrast to any single fluid system, in a multi-component system the total curvature perturbation, $\\zeta$, generally evolves whenever the non-adiabatic pressure is non-zero, \\ie when interactions between the fluids exist. Evolution of the primordial large scale curvature perturbation can relax the underlying assumptions on the inflationary scenario. Therefore analysis of multi-component fluid systems may affect our view on the physical settings. In addition to the curvature scenario considered in this paper, natural frameworks for such mechanism exist \\eg within a traditional multiple inflationary scenario \\cite{minf} or a string landscape picture \\cite{cliff, multiverse}. Whether a given scenario can effectively modify the primordial spectrum depends on the exact nature of the system.  Recent cosmic microwave surveys have also brought attention into the concept of non-gaussianity \\ie how much the spectrum deviates from gaussian distribution. This is especially important in multifield models of inflation including the curvaton scenario \\cite{Malik:2006pm,Sasaki:2006kq}.  The mechanism how energy is transferred between the fluids can be described by different methods, \\eg by a constant interaction \\cite{Malik:2002jb} or by utilizing the so-called sudden decay approximation \\cite{Lyth:2001nq, Lyth:2002my}. In a recent paper \\cite{Multamaki:2006} we considered relaxing the assumptions behind these approximations by allowing for time dependent interactions while evolving the full large-scale perturbation equations. Such an approach can better model the micro physics behind a particular physical framework by allowing one to choose the strength and the time at which the interaction is turned on. In contrast, if the interaction between is modeled with a constant interaction term, the fluid begins to decay (or interact) when its decay width is of the order of the Hubble rate, $\\Gamma\\sim H$. Physical scenarios relevant to having time (and space) dependent interactions include \\eg phase transitions, multiple inflation scenarios \\cite{minf} and scenarios where locally different decay rates of the inflaton are generated by spatially varying reheating temperature and couplings \\cite{Matarrese:2003tk, Dvali:2003em, Kofman:2003nx}.  In the present paper we consider the curvaton scenario with time dependent interactions between curvaton and other fluids. During inflation the curvaton is a light scalar field that does not contribute to the expansion of the universe; after the inflaton field has decayed into relativistic particles the curvaton begins to oscillate and to decay into radiation and matter. The focus of this article is on this situation: we study how a time-dependent interaction affects the evolution of the curvature perturbation. We study the physically allowed parameter space by utilizing information from known cosmological epochs. Moreover, we calculate the amount isocurvature in terms of curvaton decay widths. Finally, we also discuss how to describe non-gaussianity in the three fluid model and present the non-linearity factor $f_{NL}$ \\cite{Komatsu:2001rj}.  This paper is organized as follows. In sections II and III we present the governing equations of the background and perturbations in the Newtonian gauge, including discussion of non-gaussianity as described by the $f_{NL}$ parameter. In section IV we present our results while the discussion and conclusions are presented in the following section V. ",
        "conclusions": "The curvaton model has gained a lot of attention in the recent years mainly because it can make the inflation potential look more natural \\cite{Lyth:2001nq}.  Since the first model where the curvaton decayed only into radiation \\cite{Lyth:2001nq,Moroi:2001ct}, a number of other possibilities have been explored including a curvaton web model \\cite{Linde:2005yw}, the possibility of multiple curvaton fields \\cite{Assadullahi:2007uw} and different particle models such as axions \\cite{Chun:2004gx,Dimopoulos:2005bp} just to name a few.  In the present paper we have studied a three-fluid model in which the curvaton decays into both radiation and cold dark matter. This has been studied previously also in \\cite{Gupta:2003jc,Lyth:2002my,Gordon:2002gv}. We have assumed that the initial system has no matter content and it is dominated by the radiation which originates from the decay of the inflaton field. Because all of the matter content of the universe comes from the curvaton field we are able to estimate the reheating temperature. This can be additionally used to constraint the parameter space of the model.  We have systematically scanned the parameter space and identified the regions where the model is physically acceptable, \\ie when evolution during and after nucleosynthesis is standard while requiring that the reheating temperature is not unreasonably high. We have identified these regions both when the decay rates are fixed and when the curvaton starts to decay later. These allowed regions are alike once the rescaling $\\Gamma/H_0 \\rightarrow \\Gamma/H(N_*)$ is taken into account and the real difference appears in the initial system temperature.  We find that if the decay rates are comparable to the Hubble rate, a small initial curvaton density is required. Otherwise one needs to fine-tune the decay rates to be much smaller than $H$ at the time of decay. In the continuous interaction case, requiring $\\Gamma_i\\sim H$ leads to a low reheat temperature, but this can be avoided when the matter interaction is delayed.  If the initial curvaton density is large, the final state is naturally adiabatic assuming that the system is otherwise physically acceptable. Note however, that this in turn requires fine-tuning in the decay rates. If $\\Gamma_i\\sim H$, we find that the final state generally contains a large isocurvature component.  We have also we studied non-gaussianity in the framework of the three-fluid models. We find that in the region where the first-order perturbation theory can be applied, the three-fluid model gives no limits on the $f_{NL}$ parameter. This is the result of a conserved curvature perturbation $\\zeta_c$, which carries the initial curvaton perturbation $\\zeta_{\\sigma}$ into the matter perturbation $\\zeta_m$ \\cite{Gupta:2003jc}. Our results differ from the previous results \\cite{Lyth:2002my,Gupta:2003jc} mainly because our $f_{NL}$ is evaluated at the time of last scattering and not at nucleosynthesis. This allows the non-gaussianity to be more easily compared to the observational Sachs-Wolfe effect and we do not have to use a radiation transfer function. In order to calculate the observational non-gaussianity, second-order perturbation theory needs to be applied \\cite{Bartolo:2004if}. This is however beyond the scope of this article and will be the focus of a follow-up paper.  \\subsection*"
    },
    "0710/0710.2622_arXiv.txt": {
        "abstract": "Lyman-alpha (\\lya) is one of the dominant tools used to probe the star-forming galaxy population at high-redshift ($z$). However, astrophysical interpretations of data drawn from \\lya\\ alone hinge on the \\lya\\ escape fraction which, due to the complex radiative transport, may vary greatly. Here we map the \\lya\\ emission from the local luminous blue compact galaxy Haro\\,11, a known emitter of \\lya\\ and the only known candidate for low-$z$ Lyman continuum emission (LyC). To aid in the interpretation we perform a detailed UV and optical multi-wavelength analysis and model the stellar population, dust distribution, ionising photon budget, and star-cluster population. We use archival X-ray observations to further constrain properties of the starburst and estimate the neutral hydrogen column density.  The \\lya\\ morphology is found to be largely symmetric around a single young star forming knot and is strongly decoupled from other wavelengths. From general surface photometry, only very slight correlation is found between \\lya\\ and \\halpha, \\ebv, and the age of the stellar population. Only around the central \\lya-bright cluster do we find the \\lya/\\halpha\\ ratio at values predicted by recombination theory. The total \\lya\\ escape fraction is found to be just 3\\%. We compute that $\\sim 90$\\% of the \\lya\\ photons that escape do so after undergoing multiple resonance scattering events, masking their point of origin. This leads to a largely symmetric distribution and, by increasing the distance that photons must travel to escape, decreases the escape probability significantly. While dust must ultimately be responsible for the destruction of \\lya, it plays little role in governing the observed morphology, which is regulated more by ISM kinematics and geometry. We find tentative evidence for local \\lya\\ equivalent width in the immediate vicinity of star-clusters being a function of cluster age, consistent with hydrodynamic studies. We estimate the intrinsic production of ionising photons and put further constraints of $\\sim 9$\\% on the escaping fraction of photons at 900\\AA. ",
        "introduction": "Four decades ago \\cite{pp67} discussed the prospects of identifying `primeval' galaxies (i.e. galaxies forming their first generation of stars), using both the Lyman decrement and Lyman-alpha (\\lya) emission line as observational tracers. Both methods rely upon the galaxies hosting numerous young, massive stars producing a strong radiation field in the far ultraviolet (UV). The absorption of this radiation field bluewards of the Lyman absorption edge results in the Lyman-break phenomenon, while the reprocessing of the absorbed photons in astrophysical nebulae results in the superimposition of strong hydrogen recombination lines on the galaxy spectra. While it's likely that both features are present in the spectra of high-redshift ($z$) starbursts, from a survey perspective they compete. Lyman-break candidates are identified in multi-broadband imaging surveys where large ranges in redshift (and therefore cosmic volume) can be probed, but are biased towards the detection of galaxies with FUV continua towards the brighter end of the luminosity function (LF) -- less numerous in the hierarchical galaxy formation scenario. On the other hand, because emission lines concentrate a large amount of energy in a very small spectral region, sources with much fainter continuum can be uncovered. Unfortunately, isolation of a line requires either a spectroscopic surveys which typically probe narrow fields of view because of the slit dramatically reducing the size in one dimension, or narrowband imaging which probes only a small range in redshift. The advent of large diameter reflectors, efficient optics, and sensitive photometric devices has resulted in both techniques enjoying much success in recent years.  Young galaxies with low dust contents can be expected to produce \\lya\\ emission with high equivalent widths (\\wlya), $\\sim 240$\\AA\\ for solar-like metallicities ($Z$) and standard initial mass functions (IMF) \\citep{charlotfall93}. This value is predicted to increase significantly as metallicities approach the population {\\sc iii} domain \\citep{schaerer03}. Since the \\lya\\ rest wavelength lies in the FUV, the line is still observable from the ground in the optical domain at $z\\sim 6$ and beyond making it, in principle, the ideal tool by which to identify star-forming galaxies in the early universe. Despite a rocky start \\citep{pritchet93} \\lya\\ emitters (LAEs) are showing up en masse in high-$z$ surveys and perhaps their observed and predicted number-densities are now converging \\citep{ledelliou06}. \\lya\\ is now being used to put constraints on the final stages of cosmic reionisation \\citep{malhotrarhoads04,dijkstra06} and explore high-redshift large scale structure and galaxy clustering \\citep{hamana04,malhotra05,murayama07}. In addition, numerous sources have been detected through narrowband imaging techniques with very high \\wlya \\citep{malhotrarhoads02,shimasaku06}, perhaps indicative of top heavy IMFs or extreme stars, making such targets ideal to search for signatures of the so-far illusive population {\\sc iii} stars. However, so far spectroscopic observations have failed to detect the strong He{\\sc ii}$\\lambda 1640$\\AA\\ feature expected from pop {\\sc iii} objects (eg. \\citealt{dawson04,nagao07}). Finally, \\lya\\ is also being used to estimate the cosmic star formation rate density at the highest redshifts \\citep{fujita03,yamada05}. Such star formation rates (SFR) are typically estimated assuming case B recombination \\citep{brocklehurst71} and the \\halpha\\ SFR calibration of \\cite{kennicutt98}, although the large spread in SFR(\\lya) {\\em vs.} SFR(FUV) and consistent underestimates from SFR(\\lya) (eg. \\citealt{murayama07}) suggest that such a calibration should be used with caution. This caution must be extended to all cosmological studies in light of the fact that only a fraction of UV-selected high-$z$ targets show a \\lya\\ feature in emission \\citep[e.g.][]{shapley03}.  The complexities of using \\lya\\ as a cosmological tool result from the fact that it is a resonant line, and its potential cosmological importance has motivated a number of studies of star-forming galaxies at low-$z$ and theoretical models of resonant line radiative transfer. Early studies with the {\\em International Ultraviolet Explorer (IUE)} initially mirrored the first results at high-$z$: \\lya\\ was typically weak or absent in local starbursts. This systematic weakening of \\lya\\ was first attributed to dust absorption (eg. \\citealt{charlotfall91}) and an anticorrelation between \\wlya\\ and metallicity was found in the {\\em IUE} sample \\citep{charlotfall93}. However, the damped \\lya\\ absorption seen in some local starbursts, and the failure of dust corrections to reconcile \\lya\\ with the fluxes predicted by recombination theory \\citep{giavalisco96}, is indicative of selective attenuation of \\lya. This is clearly exemplified by the fact that the most metal-poor galaxies known at low-$z$ ({\\sc i}Zw\\,18 and SBS\\,03350-52) are shown to be damped \\lya\\ absorbers \\citep{kunth94,thuanizotov97} in spectroscopic observations from the {\\em Goddard High Resolution Spectrograph (GHRS)} and {\\em Space Telescope Imaging Spectrograph (STIS)}. This is to be expected if the starburst is enshrouded by a static layer of neutral hydrogen that is able to resonantly trap \\lya, thereby greatly increasing the path-lengths of the photons in order to escape the host, and exponentially increasing the chance of their destruction by dust grains \\citep{neufeld90}. Further {\\em GHRS} observations \\citep{kunth98} showed that when \\lya\\ is seen in emission, it almost systematically shows a P\\,Cygni profile and systematic velocity offset from low ionisation state metal absorption features in the neutral ISM, suggesting that an outflowing medium is an essential ingredient in the formation of the line profile and \\lya\\ escape physics. Similar results have also been found at high-$z$ \\citep{shapley03,tapken07}. Hydrodynamic models of expanding bubbles \\citep{tt99} predicted \\lya\\ line profiles to follow an evolutionary sequence starting with pure absorption at the earliest times, developing through a pure emission phase into P\\,Cygni profiles as the ISM is driven out by mechanical feedback from the starburst, and fading back into absorption at late times. This allowed \\cite{mashesse03} to reconcile the variety of \\lya\\ profiles observed at low-$z$ in such an evolutionary sequence. Further theoretical studies \\citep[e.g.][]{ahn03,verhamme06} have shown how a wide variety of P\\,Cygni-like and asymmetric profiles can develop, depending on the ISM properties and geometry. The line formation becomes more complex still if the ISM is multiphase \\citep{neufeld91,hansen06} and certain physical and kinematic configurations may lead to increases in the fraction of escaping \\lya\\ photons compared to non-resonant radiation, thereby effectively boosting \\wlya. Recent advances have been made in the understanding of observed line profiles with the implementation of full 3-D codes with arbitrary distributions of gas, ionisation, temperature, dust, and kinematics \\citep{verhamme06}.  Since nebular ionisation does not occur in situ with ionising sources, recombination line imaging (e.g. \\halpha) may reveal a morphology different from that found by imaging of the stellar continuum. This phenomenon may be much more significant for \\lya\\ photons which resonantly scatter. While \\lya\\ and \\halpha\\ have the same points of origin, it is likely that internal \\lya\\ scattering events may cause \\lya\\ photons to be emitted far from their production sites, and not be spatially correlated with \\halpha\\ or other non-resonant recombination lines. Targeted spectroscopic observations are therefore liable to miss a fraction of the emission and, while containing neither frequency nor kinematic information, \\lya\\ imaging becomes an invaluable complement to the spectroscopic studies. This was the motivation for our imaging survey of local starbursts using the {\\em Advanced Camera for Surveys (ACS)} (\\citealt{kunth03,hayes05,ostlin07} in preparation). This is a truly unique dataset for a number of reasons. Firstly the angular resolution of the {\\em ACS} allows us to map the \\lya\\ morphology on scales of 5--15~pc; 2--3 orders of magnitude better than typical studies at high-$z$. The addition of \\halpha\\ (absent in almost all high-$z$ studies due to an inconvenient rest wavelength) allows us to quantitatively examine the decoupling of \\lya\\ from non-resonant lines and estimate the global escape fractions. Multiband UV and optical data allow us to map dust reddening, stellar ages and masses, ionising photon production, and other properties of the host, all at the same resolution as \\lya. The first detailed \\lya\\ imaging study of a local starburst, ESO\\,338-IG04 \\citep{hayes05} found emission and absorption varying on very small scales in the central starburst regions,  and little or no correlation with the FUV morphology. The starburst is surrounded with a large, diffuse, low surface brightness \\lya\\ halo that contributes $\\sim 70$\\% to the global \\lya\\ luminosity, resulting from the resonant decoupling and diffusion of \\lya. The total escape fraction  was found to be just $\\sim 5$\\%, implying any global values (eg. SFR) that would be estimated from \\lya\\ alone would be seriously at fault.  Feedback processes from star-formation are capable of driving galaxy-scale `superwinds' that shock heat and accelerate the ambient medium and circumnuclear gas, resulting in large-scale, diffuse X-ray nebulae \\citep{heckman90}. Typically, hot X-ray emitting regions exhibit a tight morphological correlation with with warm gas as probed by optical emission lines \\citep{grimes05}. Indeed, such X-ray nebulae are near-ubiquitous in starbursts \\citep{strickland04} and have been calibrated as tracers of star-formation rates \\citep{ranalli03,colbert04}. \\cite{grimes05} have also shown that observed X-ray spectra are well reproduced by a single thermal plasma  over a range of galaxy classifications spanning dwarfs, discs, and ultra-luminous infrared galaxies (ULIRG), although the central regions of some ULIRGs require an additional power-law component. Since the diffuse X-ray component can be so valuable for an understanding of the wind properties, and feedback is an essential ingredient in the escape of \\lya\\ \\citep{kunth98}, X-ray information provides a valuable supporting dataset for a detailed investigation of \\lya.  In this article we turn our attention to another target in our sample, the well-known, luminous ($M_B$ = -20.5), low metallicity ($\\log(O/H)+12=7.9$; \\citealt{bergvall02}) blue compact galaxy (BCG). It exhibits a complex morphology consisting of three main star-forming condensations and an unrelaxed kinematic structure \\citep{ostlin99,ostlin01}, suggestive of a dwarf galaxy merger. It is actively star-forming, exhibits a large number of bright young star clusters \\citep{ostlin00}, is a known emitter of \\lya\\ \\citep{kunth98}, and the only known local candidate emitter of Lyman continuum (LyC) \\citep{bergvall06,grimes07}. The FUV continuum luminosity puts Haro\\,11 at the brighter end of the distribution of \\lya\\ emitters at $z\\sim 3.1$ \\citep{gronwall07}. We utilise {\\em HST/ACS} images in the FUV ({\\em Solar Blind Channel; SBC}), to examine \\lya\\ and nearby continuum, and broadband images in the UV and optical, and narrowband \\halpha\\ ({\\em High Resolution Camera; HRC, and Wide Field Camera; WFC}) to examine dust, ages in the stellar population, the star cluster population as a whole, and estimate the LyC production. We use deep ground-based narrowband images in \\halpha\\ and \\hbeta\\ in order to estimate extinction in the gas phase. X-ray observations of Haro\\,11 have been obtained using the {\\em Chandra} satellite \\citep{grimes07} but currently their status is still proprietary. We hence also exploit serendipitous off-axis observations from the {\\em Chandra} and {\\em XMM-Newton} X-ray observatories to study the wind properties and internal photoelectric absorption.  The article is arranged in the following manner: in Section~\\ref{sect:data} we describe the observations and data reductions; in Section~\\ref{sect:res} we present the results; in Section~\\ref{sect:andis} we analyse discuss the results; and in Section~\\ref{sect:conc} we present our concluding remarks. We assume a cosmology of $H_0=72$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_\\mathrm{M}=0.3$ and $\\Omega_\\Lambda=0.7$ throughout. The redshift of Haro\\,11 is taken to be 0.020598 from NED, corresponding to a luminosity distance of 87.1~Mpc.     \\section[]{Observations, reductions, and data processing}\\label{sect:data}   \\subsection{{\\em HST} Ultraviolet and optical}   \\subsubsection{Observations}  This study makes use of images from all three channels of the {\\em ACS} onboard {\\em HST}: the {\\em SBC} for the FUV; the {\\em HRC} for the 2200\\AA\\ and $U-$band; and the more sensitive {\\em WFC} for images in the optical domain. The details of the observations and post-reduction processing of the {\\em HST} UV and optical dataset are described in a companion paper, \\cite{ostlin07}. Observations were performed in the bandpasses listed in Table~\\ref{tab:exptime}. Briefly, {\\em F122M} and {\\em F140LP} correspond to \\lya\\ on-line and continuum, respectively, while {\\em FR656N} covers \\halpha\\ on-line and {\\em F550M} measures line-free continuum bluewards of \\halpha. The remainder are broadband continuum observations, selected in order to avoid the strongest emission lines. \\begin{table} \\caption{Observations and exposure times} \\centering \\begin{tabular}{@{}lllll@{}} \\hline \\hline Bandpass & &Channel & ExpTime [ s ] & \\# Split \\\\ \\hline {\\em F122M}   & \\lya           & {\\em SBC} & 9095 & 5 \\\\ {\\em F140LP}  &\\lya\\ cont      & {\\em SBC} & 2700 & 5 \\\\ {\\em F220W}   &$\\sim 2200$\\AA  & {\\em HRC} & 1513 & 3 \\\\ {\\em F330W}   &$U-$band        & {\\em HRC} & 800  & 2 \\\\ {\\em F435W}   &$B-$band        & {\\em WFC} & 680  & 2 \\\\ {\\em F550M}   &medium $\\sim V$ & {\\em WFC} & 471  & 2 \\\\ {\\em FR656N}  &\\halpha         & {\\em WFC} & 680  & 2 \\\\ {\\em F814W}   &$I-$band        & {\\em HRC} & 100  & 1 \\\\ \\hline \\end{tabular} \\label{tab:exptime} \\end{table}  All images are `drizzled' using the {\\tt MULTIDRIZZLE} task in {\\tt IRAF/STSDAS} onto the same pixel sampling scale (0.025\\arcsec /pixel) and position angle. The inverse variance weight maps were saved from the drizzle process since they provide an estimate of the error in each pixel. Remaining cosmic rays, charge transfer tracks, and blemishes are removed from the CCD observations using the {\\tt CREDIT} task. Remaining band-to-band discrepancies in the astrometric alignment were rectified using the {\\tt GEOMAP} and {\\tt GEOTRAN} tasks. Since our dataset comprises observations between 1200 and 9000\\AA\\ and utilises three different {\\em ACS} channels, we address the issue of variations in the point spread function (PSF) of the images. PSF models of each band were generated using the {\\tt PSF} task in {\\tt DIGIPHOT/DAOPHOT}, with the resulting models being used to convolve all images to the PSF of {\\em F550M} (the broadest emission-line free PSF in our dataset). PSF models were re-computed for the convolved images and compared to that of {\\em F550M}; all frames showed PSF full width at half maximum consistent with {\\em F550M} at below the 5\\% level.   \\subsubsection{\\lya\\ continuum subtraction and SED modeling}\\label{sect:contsub}  As described in \\cite{hayes05}, and in more detail in Hayes et al. (2007, in preparation), continuum subtraction of \\lya\\ using these filters requires more sophisticated techniques than most emission lines. Because the offline filter is removed from the online filter by $\\Delta \\lambda / \\lambda = 0.22$ and the FUV continuum evolves rapidly as a function of $\\lambda$, age, and \\ebv, it is imperative to understand the behaviour of the continuum between {\\em F140LP} and {\\em F122M}. Special care is also required because of the broad nature of the {\\em F122M} filter which has a rectangular width around 10\\% of the central wavelength, and transmits both the geocoronal \\lya\\ line and Milky Way \\lya\\ absorption profile. In addition, P\\,Cygni profiles may result in a reduction of the net emission due to the blue-side absorption cancelling some or all of the emission.  We defined in \\cite{hayes05} the {\\em Continuum Throughput Normalisation} (\\ctn) factor as the factor that, for a given spectrum, scales the continuum flux sampled in {\\em F140LP} to the flux that would be expected from continuum processes alone in {\\em F122M}. The procedure of estimating \\ctn\\ in each pixel utilises each of the images between {\\em F140LP} and {\\em F814W}, the filter throughput profiles, and fitting {\\em Starburst99} spectral evolutionary models \\citep{leitherer99,vazquez05}. In \\cite{hayes05} we demonstrated that we could find non-degenerate best-fitting spectra if SED datapoints sample the UV continuum slope and Balmer/4000\\AA\\ break. That way, for each pixel we fit burst age and \\ebv\\ using  standard $\\chi^2$ minimisation. The method has been substantially developed and is described in detail in Hayes et al. (2007, in preparation). We now use continuum-subtracted \\halpha\\ to first map the contribution to the overall SED from nebular gas alone. Thanks to observations at $V$ and $I$ we are also able to constrain the contribution from any stellar population that underlies the current starburst. Essentially we measure and subtract the gas spectrum and fit age and mass in two stellar components, applying the same reddening for all three SEDs. For each pixel, we are able to find the age of the stellar populations and \\ebv, treating the nebular gas and two stellar components independently. With the best-fitting spectral reconstruction (composite {\\em Starburst99} spectra) at each {\\em HST} pixel we then map the \\ctn\\ factor and are able to reliably continuum subtract \\lya.  In Hayes et al. (2007, in preparation) we present extensive simulations, designed to test the reliability of this methodology for an array of input spectra. We determine that the method must account for the presence of an underlying stellar population and nebular continuum emission. Both these components may affect the $U-B$ colour and result in bad fits and poor recovery of \\ctn\\ in cases where they dominate the optical luminosity. Our method of reconstructing the nebular continuum and fitting multiple stellar components allows robust recovery of the FUV continuum features and \\ctn. An independent measure of the metallicity is also found to be a requirement although the metallicity of Haro\\,11 is well known. Details of the IMF are found not to have a detrimental impact when multiple stellar component fitting  (i.e. a more complex star-formation history) is used. Overall we find that the software employed here is able to always recover input \\lya\\ equivalent widths to within 30\\% for `weak' \\lya\\ emission (\\wlya=10\\AA) and to within 10\\% when the \\lya\\ line is stronger (\\wlya=100\\AA). With regard to the integrated fluxes and visual morphologies, we found in \\cite{hayes05} that varying the parameters of the model spectra had a minimal effect on our results. Morphologies were always indistinguishable by eye and integrated \\lya\\ fluxes self-consistent to within $\\sim 25$\\%, even when pushing the parameters outside the regimes deemed reasonable for the galaxy in question.  This procedure also provides numerous other outputs against which we can compare \\lya. The full list of output maps is: continuum subtracted \\lya\\ and \\halpha, \\ctn-factor at \\lya\\ ({\\em F122M/F140LP}), continuum flux densities at 900\\AA\\ (\\flyc), 1500\\AA\\ (\\ffuv), and 2200\\AA\\ (\\fnuv), \\ebv, the age of the two stellar components, the mass of the two stellar components, the stellar equivalent width in absorption of \\halpha\\ and \\hbeta, and $\\chi^2$. These allow us to compare \\lya\\ fluxes and equivalent widths with various local stellar ages, \\ebv, local star-formation rates, and estimate the ionising photon production.  Haro\\,11 metallicity is known to be around 20\\% solar \\citep{bergvall02} so for this study we adopt the evolutionary tracks generated with metallicity $Z=0.004$. Since we are concerned with OB-dominated, young starbursts we use the tracks of the Geneva group. In the absence of a quantitative estimate, the IMF was assumed to follow that of Salpeter ($\\alpha=-2.35$) in the range 0.1~M$_\\odot$ to 120~M$_\\odot$. Star-formation history was taken to be that of an instantaneous starburst (single stellar population) with a mass normalisation of $10^6$~M$_\\odot$.     \\subsubsection{Binning}\\label{sect:reductionbinning}  As described in the introduction, we can expect to see \\lya\\ features (either emission or absorption) around clusters where variations can be expected on small scales, and/or large-scale diffuse emission at low surface brightness. Typically $S/N$ per resolution element is high in the vicinity of strong continuum sources but significantly lower $(< 1)$ in the diffuse regions. Ordinarily, such problems are overcome by smoothing. While fixed kernel smoothing may improve $S/N$ in diffuse regions it smears out spatial details in regions where $S/N$ is good and smoothing would not be necessary. Adaptive smoothing may maintain some spatial resolution but does not necessarily conserve surface brightness.  Adaptive binning provides an alternative to the smoothing approach and for this study we make use of the Voronoi tessellation code of \\cite{die06}, a generalisation of the Voronoi binning algorithm of \\cite{cappel03}. Voronoi tessellation overcomes the drawbacks of smoothing by binning together groups of pixels to conglomerate resolution elements (frequently referred to as ``spaxels\"), recomputing the signal and noise in each new bin. Pixels are continuously accreted until a threshold $S/N$ has been met in each spaxel. The advantage of this adaptive binning method is that in high-$S/N$ regions, the spatial sampling remains high because spaxels are typically small, whereas in diffuse regions $S/N$ is improved greatly. Diffuse emission regions, by definition, show variations over much larger spatial scales. Since each individual pixel is used exactly once, surface brightness is always conserved.  The {\\em F140LP} FUV science observation is assigned as the ``reference\" image, and used to generate the binning pattern, using the inverse variance map output by the {\\tt STSDAS/MULTIDRIZZLE} task to compute $S/N$. Pixels are accreted into spaxels until $S/N=5$ has been met, with a maximum size of $40^2$pixels $ = 1\\square$\\arcsec\\ in order to speed up the process and prevent spaxels from growing arbitrarily large.   \\subsubsection{Super star clusters}  Point-like sources are identified in the deepest of the optical bandpasses, {\\em F435W} and {\\em F550M} using the {\\tt DAOFIND} task in {\\tt IRAF}. To be considered a detection, a cluster must be present in both of these bands. Each object was then inspected by eye and any detections that were clearly spurious were removed from the catalogue. For crowded fields PSF-photometry is preferred to standard aperture methods in order to eliminate cross-contamination from neighbouring clusters. However, some clusters may be extended enough to be resolved by our observations, thus precluding the use of PSF-fitting methods and limiting us to aperture photometry. Aperture photometry was performed using the {\\tt PHOT} task in {\\tt IRAF} in all bands using the {\\em F435W+F550M} detected catalogue. An aperture of 0.10\\arcsec\\ was used, and sky was sampled in a circular annulus of radius between 0.1 and 0.15\\arcsec. Aperture corrections were then applied in accordance with those given in \\cite{siri05} for the {\\em HRC} and {\\em WFC} bandpasses. While {\\em SBC/F140LP} aperture corrections have been computed in the past, \\citep[e.g.][]{dieball05}, they are not available in the published literature and our own aperture corrections, computed with the {\\tt TinyTim} software \\citep{tinytim}, are included here in Table~\\ref{tab:apcorr}, Appendix~\\ref{sect:apcorr}. Since we have no a priori knowledge of the continuum slope, we adopt the {\\em F140LP} aperture correction for $\\beta=0$, using 0.1\\arcsec\\ as in the other bandpasses. Since {\\em SBC/F122M} contains the \\lya\\ line, no point-source photometry was performed in this bandpass itself. Instead, we perform aperture photometry on the continuum subtracted \\lya\\ and \\halpha\\ images in the same 0.1\\arcsec\\ apertures at the position of each cluster with no re-centering applied. This gives a direct measure of the \\lya\\ flux and equivalent width in the vicinity of each SSC.  In order to estimate the properties of the SSCs, the same SED-fitting software described in Section~\\ref{sect:contsub} was applied to the aperture extracted fluxes. Age, \\ebv, photometric mass, etc. were computed, allowing us to compare these properties with \\lya\\ in the immediate vicinity of each SSC.       \\subsection{ESO -- New Technology Telescope}  Haro\\,11 was observed during the nights of 18, 19, and 20 September 2004, using the {\\em New Technology Telescope} at ESO La Silla, as part of an observing run to obtain \\halpha, \\hbeta, and [O{\\sc iii}] narrowband images for all southern targets in our {\\em HST} \\lya\\ sample (Atek et al., in preparation). On the night beginning 18 Sept seeing was good, not exceeding 1.2\\arcsec for the duration, although thin cirrus prevents a direct calibration of these data. On the night of 19 Sept, observational conditions were ideal: photometric and with sub-arcsecond seeing throughout. The final night was still photometric but seeing deteriorated to $>2$\\arcsec\\ and images were unusable for science purposes. The good-seeing data from the night of 18 was calibrated through secondary standard stars in the field, using data from the photometric nights of 19 and 20 Sept. Narrowband imaging was performed at \\halpha, \\hbeta, and nearby continuum. Spectrophotometric standard stars LDS749B, Feige 110, and GD50 were selected (on the grounds of their high spectral resolution: $1-2$\\AA, in order to resolve the stellar absorption features around \\halpha\\ and \\hbeta) from the catalog of \\cite{oke90}, and were observed at regular intervals for the duration of each night in all filters. Imaging observations were performed using both the {\\em ESO Multi-Mode Instrument (EMMI)} \\citep{dek86} and {\\em Super Seeing Imager 2 (SuSI2)} \\citep{odor98}, interchangeably. The instruments were used in $2\\times 2$ pixel binning mode to reduce the readout noise, providing a plate-scale of 0.1665\\arcsec~pix$^{-1}$ and 0.130\\arcsec~pix$^{-1}$, and a field-of-view (FoV) of $9.1 \\times 9.9$\\arcsec\\ and $2.2 \\times 2.2$\\arcsec, for {\\em EMMI-R} and {\\em SuSI2}, respectively. Table \\ref{tab:ntt} summarises the imaging observations included in this article, lists the ESO filters used, and the total exposure times in each band.  \\begin{table} \\caption{NTT observations} \\centering \\begin{tabular}{@{}llcc@{}} \\hline \\hline Observation & Camera & ESO filter \\# & ExpTime [ s ] \\\\ \\hline \\halpha        & {\\em EMMI-R} & 598 &  900 \\\\ \\halpha\\ cont. & {\\em EMMI-R} & 597 &  1200 \\\\ \\hbeta         & {\\em SuSI2}  & 549 &  2866 \\\\ \\hbeta\\ cont.  & {\\em EMMI-R} & 770 &  1800 \\\\ \\hline \\end{tabular} \\label{tab:ntt} \\end{table}  Data were first reduced by the standard routines in {\\tt NOAO/IRAF}: bias subtraction and flat-field correction using well exposed sky- and dome-flats. Images were aligned using the {\\tt GEOMAP/GEOTRAN} tasks and smoothed to the seeing of the worst seeing image using the {\\tt GAUSS} task. All {\\em HST} bandpasses are aligned with the {\\em NTT} frames and the PSF is degraded by convolution with a Gaussian kernel. \\halpha\\ and \\hbeta\\ were continuum subtracted to obtain the nebular emission fluxes, accounting for contamination by [\\nii] (\\halpha\\ only) and underlying stellar absorption. [\\nii] contamination is estimated using [\\nii]$\\lambda 6583$\\AA/\\halpha$=0.189$ \\citep{bergvall02}. Stellar absorption is estimated from the best-fitting {\\em Starburst99} spectrum at each pixel by modifying the spectral fitting code described in Sect.~\\ref{sect:contsub}. The {\\em Starburst99} stellar libraries \\citep{mart05} were built upon the latest model atmospheres and include full line-blanketing for all stars and non-LTE effects for hot stars, thereby providing the best estimate of the stellar absorption features available. With the age of the stellar population we measure the equivalent width of \\hbeta\\ using the same line and continuum wavelength windows as the Lick index \\citep{lick}. The windows used for \\halpha\\ are of the same size as those for \\hbeta, but scaled to \\halpha\\ (i.e. the same windows $\\times 6563/4861=1.35$). \\ebv\\ is generated from \\halpha/\\hbeta, assuming a temperature of 10,000~K and intrinsic line ratio of 2.86, using the extinction law of the SMC following \\cite{Gordon03,Fitzpatrick_Massa88}.    \\subsection{X-Ray: {\\em Chandra} and {\\em XMM-Newton}}\\label{sect:resxray}  Haro\\,11 happens to be located 14\\arcmin\\ away from the Cartwheel galaxy, of which X-ray observations have been obtained with both the {\\em Chandra} and {\\em  XMM-Newton} telescopes \\citep{wolt04}. Respective total integration times in {\\em Chandra} and {\\em XMM-Newton} were 80 and 70 ksec. Regarding {\\em XMM-Newton}, Haro\\,11 only falls within the field--of--view of the {\\em MOS2} detector, having unfortunately been missed by both the {\\em MOS1} and {\\em PN} chips. Both the {\\em Chandra} and {\\em XMM-Newton} datasets have been obtained from their respective archives. Since both telescopes operate with curved focal-plane configuration, the point spread function degrades severely with angular distance from the optical axis. For the {\\em XMM-Newton} detection, the PSF is so distorted that all spatial information has been lost, but the object is bright enough that a one-dimensional spectrum can be extracted from the {\\em MOS2} data. In the {\\em Chandra} observation, Haro\\,11 lies on the S1 chip: the object is clearly detected and we were able to perform a spatial analysis of the source (see Sect.~\\ref{sect:res_xray}).  The {\\em Chandra} data were processed following the X-Ray Data Centre pipeline software. The level 1 events were reprocessed using {\\tt Ciao~3.4} task {\\tt acis\\_process\\_events}. The whole band (0.2-10 keV) image was smoothed applying the {\\tt csmooth} task. The image was smoothed with a minimal signal--to--noise of 3 using a circular Gaussian kernel. The source spectrum was extracted from a circular region of radius 30\\arcsec. The background was obtained from the combination of four circular regions located close to the source and avoiding other X-ray sources in the field. Source counts were grouped to have at least 20 counts per bin to allow modified $\\chi^2$ minimisation technique \\citep{kendall73} in the spectral analysis. Redistribution matrix and auxiliary response matrix files were generated.  The {\\em XMM-Newton} data were processed using the standard {\\tt Science Analysis System, SAS, v.7.0.0} \\citep{gabriel04}. The most up-to-date files available as of January 2007 were used for the reduction process. The time intervals corresponding to high background events were removed using the method followed in \\cite{piconcelli04}. The resulting exposure for the {\\em MOS2} data is only of 44 ks. No sign of pile-up was found in the {\\em MOS2} image, according to the {\\tt epaplot  SAS} task. A visual inspection of the 0.2-10~keV image shows a distorted shape due to location of the source, close to the edge of the CCD. Therefore, no further analysis of the {\\em XMM-Newton} image has been performed. The source spectrum was obtained, as for the {\\em  Chandra} observation, from a circular region with a radius of 30\\arcsec. The background was extracted from a circular region close to the source and avoiding other X-ray sources in the  field. Source counts were grouped to have at least 20 counts per bin. {\\tt SAS} appropriate tasks were used to generate the distribution matrix and auxiliary response matrix files.  On-axis observations of Haro\\,11 have also been obtained with the {\\em Chandra} telescope \\citep{grimes07} although their archive status is still proprietary. Quantities derived from that dataset are, however, compatible with those presented here using off-axis archival observations (see Section~\\ref{sect:res_xray}). ",
        "conclusions": "\\label{sect:conc}  Using {\\em HST/ACS} we have mapped, calibrated, and analysed the \\lya\\ emission from nearby luminous blue compact galaxy Haro\\,11. The \\lya\\ emission has been compared to: \\halpha; \\ebv\\ as determined through both \\halpha/\\hbeta\\ and UV continuum; the kinematic structure of the galaxy from previous studies; and the properties of the super star clusters. We have used archival X-ray data to map the hot outflowing gas and to put constraints on the evolutionary state of the burst. From SED fitting we have estimated and mapped the predicted Lyman-continuum production at 900\\AA, and compared this direct measurements. Most notably: \\begin{itemize}  \\item{Our photometry reproduces spectroscopic fluxes as determined from the {\\em IUE} satellite. The total \\lya\\ flux is found to be $79.6 \\times 10^{-14}$~erg~s$^{-1}$~cm$^{-2}$, corresponding to a luminosity of $7.22 \\times 10^{41}$~erg~s$^{-1}$.  }  \\item{The escaping fraction of \\lya\\ photons is found to be $\\sim 3$\\%.}  \\item{\\lya\\ shows almost no spatial correlation with \\halpha. The \\lya\\ morphology shows a central bright emission region where \\lya\\ photons escape the galaxy directly, surrounded by a low surface brightness diffuse halo that results from multiple resonant scatterings. 90\\% of the \\lya\\ output is in the halo component. \\lya/\\halpha\\ at the values predicted by recombination theory are seen only in the brightest central regions.}  \\item{Little correlation is seen between \\lya\\ and dust. The main central \\lya\\ emitting region shows bright emission from a region where \\ebv\\ is significantly greater than regions that show only \\lya\\ absorption. In the regions of halo emission, \\ebv\\ extends beyond 1 although \\lya\\ is resonantly scattered and cannot feel the effect of such high extinction. This again indicates that dust is not the major regulatory factor governing the \\lya\\ morphology, which appears to be driven more by the H{\\sc i} distribution and its kinematic structure. }  \\item{X-ray observations reveal a diffuse, soft component centred on the central star forming knots, and showing hot wind-driven gas pushed outside the central starburst region as probed by UV bands. The brightest X-ray regions are found to be spatially coincident with the regions of highest surface brightness in \\lya. This indicates that \\lya\\ emission from central regions may be the result of a perturbed and outflowing ISM, accelerated by the release of mechanical energy from the SSCs located around knot \\C. }   \\item{From the super star clusters, we find that peak \\lya\\ equivalent widths are small at the youngest ages (1~Myr), reach a maximum from clusters in the range 2.5--4~Myr, and decline again to zero at ages $\\gtrsim8$~Myr. This is qualitatively consistent with models of \\lya\\ emission resulting from outflows in the neutral ISM.  }  \\item{Haro\\,11 is the only galaxy known to emit Lyman-continuum radiation ($\\lambda < 912$\\AA) in the nearby universe. From fitting spectral synthesis models we estimate the continuum flux at 900\\AA\\ to be $f_{900,0}=12.3\\times10^{-14}$~erg~s$^{-1}$~cm$^{-2}$~\\AA$^{-1}$. This corresponds to an escape fraction of 9\\% at 900\\AA.  }   \\end{itemize}"
    },
    "0710/0710.1767_arXiv.txt": {
        "abstract": "{The shell-type supernova remnant RX J1713.7--3946 was observed during three years with the H.E.S.S. Cherenkov telescope system. The first observation campaign in 2003 yielded the first-ever resolved TeV gamma-ray image. Follow-up observations in 2004 and 2005 revealed the very-high-energy gamma-ray morphology with unprecedented precision and enabled spatially resolved spectral studies. Combining the data of three years, we obtain significantly increased statistics and energy coverage of the gamma-ray spectrum as compared to earlier H.E.S.S. results. We present the analysis of the data of different years separately for comparison and demonstrate that the telescope system operates stably over the course of three years. When combining the data sets, a gamma-ray image is obtained with a superb angular resolution of 0.06 degrees. The combined spectrum extends over three orders of magnitude, with significant gamma-ray emission approaching 100 TeV. For realistic scenarios of very-high-energy gamma-ray production, the measured gamma-ray energies imply efficient particle acceleration of primary particles, electrons or protons, to energies exceeding 100 TeV in the shell of RX J1713.7--3946.}  \\begin{document} ",
        "introduction": "\\vspace{-0.4cm}  \\begin{figure*} \\begin{center} \\includegraphics [width=0.78\\textwidth]{icrc0524_fig1.eps} \\end{center} \\caption{\\underline{\\textbf{Upper panel:}} H.E.S.S.\\ gamma-ray excess images from the region around RX~J1713.7$-$3946\\ are shown for three years. \\underline{\\textbf{Lower panel:}} 1D distributions generated from the non-smoothed, acceptance-corrected gamma-ray excess images.  } \\label{fig1} \\end{figure*}       The energy spectrum of cosmic rays measured at Earth exhibits a power-law dependence over a broad energy range. Starting at a few GeV $(1~\\mathrm{GeV} = 10^9~\\mathrm{eV})$ it continues to energies of at least $10^{20}~\\mathrm{eV}$. The power-law index of the spectrum changes at two characteristics energies: in the region around $3 \\times 10^{15}~\\mathrm{eV}$ -- the \\emph{knee} region -- the spectrum steepens, and at energies beyond $10^{18}~\\mathrm{eV}$ it hardens again. Up to the knee, cosmic rays are believed to be of Galactic origin, accelerated in shell-type supernova remnants (SNRs) -- expanding shock waves initiated by supernova explosions~\\cite{HillasReview}. However, the experimental confirmation of an SNR origin of Galactic cosmic rays is difficult due to the propagation effects of charged particles in the interstellar medium. The most promising way of proving the existence of high-energy particles in SNR shells is the detection of very-high-energy (VHE) gamma rays ($E > 100~\\mathrm{GeV}$), produced in interactions of cosmic rays close to their acceleration site~\\cite{DAV}.  Recently the VHE gamma-ray instrument H.E.S.S. has detected two shell-type SNRs, RX~J1713.7$-$3946~\\cite{Hess1713a,Hess1713b} and RX~J0852.0--4622~\\cite{HessVelaJr_a,HessVelaJr_b}. The two objects show an extended morphology and exhibit a shell structure, as expected from the notion of particle acceleration in the expanding shock fronts. While it is difficult to attribute the measured VHE gamma rays unambiguously to nucleonic cosmic rays (rather than to cosmic electrons), the measured spectral shapes favour indeed in both cases a nucleonic cosmic-ray origin of the gamma rays~\\cite{Hess1713b,HessVelaJr_b}.  Apart from the first unambiguous proof of multi-TeV particle acceleration in SNRs, the question of the highest observed energies remains an important one. Only the detection of gamma rays with energies of 100~TeV provides experimental proof of acceleration of primary particles, protons or electrons, to the \\emph{knee} region (1~PeV). Here we present a combined analysis of H.E.S.S.\\ data of RX~J1713.7$-$3946\\ of three years, from 2003 to 2005. A comparison of the three data sets demonstrates the expected steady emission of the source as well as the stability of the system. Special emphasis is then devoted to the high-energy end of the combined spectrum. ",
        "conclusions": "\\vspace{-0.4cm} The complete H.E.S.S.\\ data set of the SNR RX~J1713.7$-$3946\\ recorded from 2003 to 2005 is presented here. Very good agreement is found for both the gamma-ray morphology and the differential energy spectra over the years. The combined analysis confirms the earlier findings nicely: the gamma-ray image reveals a thick, almost circular shell with significant brightness variations. The spectrum follows a hard power law with significant deviations at higher energies (beyond a few TeV).  In the combined image using $\\sim63$~hours of H.E.S.S.\\ observations an unprecedented angular resolution of $0.06\\degr$ is achieved. The high-energy end of the combined spectrum approaches 100~TeV with significant emission $(4.8\\sigma)$ beyond 30~TeV. Given the systematic uncertainties in the spectral determination at these highest energies and comparable statistical uncertainties despite the long exposure time, this measurement is presumably close to what can be studied with the current generation of imaging atmospheric Cherenkov telescopes.  From the largest measured gamma-ray energies one can estimate the corresponding energy of the primary particles. In case of $\\pi^0$-decay gamma rays, energies of 30~TeV imply that primary protons are accelerated to $30~\\mathrm{TeV} / 0.15 = 200~\\mathrm{TeV}$ in the shell of RX~J1713.7$-$3946. On the other hand, if the gamma rays are due to Inverse Compton scattering of VHE electrons, the electron energies at the current epoch can be estimated in the Thompson regime as $E_\\mathrm{e}~\\approx~20~\\sqrt{E_\\gamma}~\\mathrm{TeV} \\approx 110~\\mathrm{TeV}$. At these large energies Klein--Nishina effects start to be important and reduce the maximum energy slightly such that $\\sim100~\\mathrm{TeV}$ is a realistic estimate.  RX~J1713.7$-$3946\\ remains an exceptional SNR in respect of its VHE gamma-ray observability, being at present the remnant with the widest possible coverage along the electromagnetic spectrum. The H.E.S.S.\\ measurement of significant gamma-ray emission beyond 30~TeV without indication of a termination of the high-energy spectrum provides proof of particle acceleration in the shell of RX~J1713.7$-$3946\\ beyond $10^{14}$~eV, up to energies which start to approach the region of the cosmic-ray \\emph{knee}."
    },
    "0710/0710.4468_arXiv.txt": {
        "abstract": "To study the effect of metallicity on the mass-loss rate of asymptotic giant branch (AGB) stars, we have conducted mid-infrared photometric measurements of such stars in the Sagittarius (Sgr dSph) and Fornax dwarf spheroidal galaxies with the 10-$\\mu$m camera VISIR at the VLT.  We derive mass-loss rates for 29 AGB stars in Sgr dSph and 2 in Fornax.  The dust mass-loss rates are estimated from the $K-[9]$ and $K-[11]$ colours. Radiative transfer models are used to check the consistency of the method. Published IRAS and Spitzer data confirm that the same tight correlation between $K-[12]$ colour and dust mass-loss rates is observed for AGB stars from galaxies with different metallicities, i.e. the Galaxy, the LMC and the SMC.  The derived dust mass-loss rates are in the range 5$\\times10^{-10}$ to 3$\\times10^{-8}$ M$_{\\odot}$yr$^{-1}$ for the observed AGB stars in Sgr dSph and around 5$\\times10^{-9}$ M$_{\\odot}$yr$^{-1}$ for those in Fornax; while values obtained with the two different methods are of the same order of magnitude.  The mass-loss rates for these stars are higher than the nuclear burning rates, so they will terminate their AGB phase by the depletion of their stellar mantles before their core can grow significantly. Some observed stars have lower mass-loss rates than the minimum value predicted by theoretical models. ",
        "introduction": "Stars with initial mass in the range 0.8--8 M$_{\\odot}$ go through an asymptotic giant branch (AGB) phase towards the end of their evolution.  This evolutionary phase is dominated by strong mass loss.  The star expels material at rates up to $10^{-4}M_{\\odot}$yr$^{-1}$, eventually ejecting between 20 and 80 per cent of its initial main-sequence mass. This process is of great importance for the chemical evolution of our Galaxy. Mass loss from AGB stars contributes to around half of the gas recycled by stars (Maeder 1992), and is a dominant source of Galactic dust.  The mass-loss mechanism of AGB stars is a two-step process. First, shocks due to pulsations from the star produce gas over-densities in an extended atmosphere (e.g. H\\\"ofner et al. 1998). This triggers the formation of dust. Secondly, radiation pressure accelerates the dust to the escape velocity. The gas is carried along through friction with the dust particles. Pulsations alone can explain mass-loss rates up to about $10^{-7}M_{\\odot}$yr$^{-1}$, but the much higher rates observed require dust-driven winds (Bowen \\&\\ Wilson 1991).  The importance of metallicity on the mass-loss rates of AGB stars is not well understood. In low metallicity environments less dust is expected to form, leading to lower predicted mass-loss rates.  Theoretical work by Bowen \\& Willson (1991) predicts that for metallicities below [Fe/H]$=-1$ dust-driven winds fail, and the wind becomes pulsation-driven. This would affect the evolution of a low metallicity host galaxy in two obvious ways. First the stellar dust production would be much reduced and therefore the composition of the material out of which new stars and planets were forming would be significantly different. Secondly the much weaker dust-driven winds allow the degenerate core of the AGB star to grow for longer, resulting in much higher masses for the remnant white dwarfs. In extreme cases, the core could reach the Chandrasekhar limit before mass loss terminates its evolution. The Bowen \\&\\ Willson mass-loss rates predict the occurrence of AGB supernovae at very low metallicities (Zijlstra 2004). It is therefore important to test whether dust-driven winds exist at low metallicity.  Observations in the Magellanic Clouds and the Galaxy have shown that the dust mass-loss rates are smaller in the Magellanic Clouds.  Assuming that the dust-to-gas ratio is a linear function of metallicity ([Fe/H]=$-0.6$ (Venn 1999) and [Fe/H]=$-0.3$ for the Small and Large Magellanic Clouds respectively) yields the conclusion that the total mass-loss rate (dust+gas) may be similar in the three galaxies (van Loon 2000, 2006), although this assumption may not be strictly correct, as the dust in the carbon stars comes from primary carbon.  In order to obtain further constraints on the effect of metallicity on the mass-loss rates from AGB stars, and to know if dust-driven mass loss can occur at very low metallicities, we need to study the mass-loss rates from more metal-poor AGB stars.  The dusty circumstellar envelopes surrounding these AGB stars absorb the radiation from the central star and re-emit it in the thermal infrared. Observations at infrared wavelengths have enabled the study of mass-loss from Galactic and Magellanic Clouds AGB stars. Thanks to the sensitivity achievable with mid-infrared instruments on 8m class ground-based telescopes, it is now possible to perform such studies for AGB stars in more distant Local Group galaxies.  We thus observed AGB stars in the Fornax dSph (Fornax) and Sagittarius (Sgr dSph) dwarf spheroidal galaxies using the mid-infrared camera VISIR on the VLT (ESO, Chile). Fornax is a satellite galaxy of the Milky Way, at a distance of $\\sim$ 138 kpc. Sgr dSph is located behind the Galactic disc and bulge at a distance of $\\sim$24 kpc.  It is currently being torn apart by tidal forces (Majewski et al. 2003). Those galaxies have low metallicities (see section 2). Furthermore the distances of both galaxies are quite well known, making an estimation of the mass-loss rates easier than it is for stars within the Galaxy. ",
        "conclusions": "We have presented a study of mass-loss from AGB stars in the Sagittarius and Fornax dwarf spheroidal galaxies using mid-infrared photometry obtained with VISIR (VLT, ESO).  We have shown that the well known relation between $K-[12]$ colours and mass-loss rates among Galactic stars is also valid for AGB stars in lower metallicity galaxies. We used this relation and radiative transfer models to estimate dust mass-loss rates for 15 stars in Sgr dSph and 2 stars in Fornax.  Our study showed that some AGB stars in the Sgr dSph  and Fornax galaxies are losing mass. The estimated dust mass-loss rates are in the range 5$\\times10^{-10}$ to $3\\times10^{-8}$ M$_{\\odot}$yr$^{-1}$ for the stars in Sgr dSph and around 5$\\times10^{-9}$ M$_{\\odot}$yr$^{-1}$ for those in Fornax. The values obtained with the two different methods are of the same order of magnitude. These mass-loss rates are higher than the nuclear burning rates, so these stars will terminate their AGB evolution by the depletion of their stellar mantles, before their cores can grow significantly.  Using previous work to get an estimation of the dust-to-gas mass ratio for the stars observed we found that most of them had very low mass-loss rates for their luminosity. This is in contradiction with theoretical predictions (see e.g Gail \\& Sedlmayr 1987)."
    },
    "0710/0710.4142_arXiv.txt": {
        "abstract": "We infer the large-scale source parameters of dusty galaxies from their observed spectral energy distributions (SEDs) using the analytic radiative transfer methodology presented in Chakrabarti \\& McKee (2005). For local ultra-luminous infrared galaxies (ULIRGs), we show that the millimeter to far-infrared (FIR) SEDs can be well fit using the standard dust opacity index of 2 when self-consistent radiative transfer solutions are employed, indicating that the cold dust in local ULIRGs can be described by a single grain model. We develop a method for determining photometric redshifts of ULIRGs and sub-mm galaxies from the millimeter-FIR SED; the resulting value of $1+z$ is typically accurate to about $10\\%$.  As such, it is comparable to the accuracy of near-IR photometric redshifts  and provides a complementary means of deriving redshifts from far-IR data, such as that from the upcoming $\\it{Herschel~Space~Observatory}$.  Since our analytic radiative transfer solution is developed for homogeneous, spherically symmetric, centrally heated, dusty sources, it is relevant for infrared bright galaxies that are primarily powered by compact sources of luminosity that are embedded in a dusty envelope.  We discuss how deviations from spherical symmetry may affect the applicability of our solution, and we contrast our self-consistent analytic solution with standard approximations to demonstrate the main differences. ",
        "introduction": "The far-infrared (FIR) spectral energy distribution (SED) is a vital implement in understanding the physical conditions of dusty sources. Chakrabarti \\& McKee (2005, henceforth CM05), presented self-consistent analytic radiative transfer solutions for the spectra of unresolved, homogeneous, spherically symmetric, centrally heated, dusty sources.  We showed that from two colors in the millimeter (mm) and FIR portion of the spectrum one can approximately infer the mm to FIR, and that this in turn determines the luminosity to mass ratio, $L/M$, and surface density, $\\Sigma$, which (at low redshift) are distance-independent parameters. With a distance measurement, one can further infer the size, mass, and luminosity of the source.  We extensively compared our analytic solutions against a well-tested numerical scheme, DUSTY (Ivezic \\& Elitzur 1997), to find excellent agreement with the numerical results.  Here, we apply this methodology to dusty galaxies to derive their large-scale source parameters.  We discuss applications to protostars and radiative transfer methodology of clumpy envelopes in a separate forthcoming paper.  From observations of their SEDs, the IRAS all-sky survey characterized ULIRGs as a class of extremely luminous ($L_{8-1000~\\micron}>10^{12}L_{\\odot}$) galaxies that emit most of their energy in the FIR (Soifer et al. 1984; Aaronson \\& Olszewski 1984; Soifer et al. 1987; Sanders \\& Mirabel 1996).  These galaxies were then understood to be a new class of objects, quite distinct from those studied by optical surveys as little correlation was found between their optical and infrared luminosities (Sanders \\& Mirabel 1996).  That the demographics of these galaxies is not simply an extrapolation of normal galaxies can be seen from the fact that the luminosity function on the bright end ($L_{\\rm IR} \\ga 1 \\times 10^{11} L_{\\odot}$) is significantly in excess of the Schechter function (Rieke \\& Lebofsky 1986).  Theoretical models have suggested that heavily starbursting systems like ULIRGs can be produced via mergers of roughly equal mass galaxies (Toomre \\& Toomre 1972; Mihos \\& Hernquist 1996), and recent observations corroborate the idea that local ULIRGs are products of major mergers (Dasyra et al. 2006).  While the limited sensitivity of IRAS did not allow for a characterization of the redshift evolution of this dusty, luminous population of galaxies, observations by the $\\it{Spitzer~Space~Telescope}$ indicate that there is a strong evolution in this population (on the basis of near and mid-IR observations converted to total IR luminosities using observational templates) out to $z \\sim 1$ (e.g. Le Floch et al. 2005), with the contribution from ULIRGs to the comoving IR energy density increasing by an order of magnitude from local systems to $z \\sim 1$.   Understanding the FIR SEDs of dusty galaxies has renewed importance today.  Submillimeter galaxies (SMGs; $F_{850~\\micron} \\ga 1~\\rm mJy$) (Ivison et al. 2000; Blain et al. 1998) are luminous ($L \\ga 10^{12} L_{\\odot}$), dusty galaxies at moderate redshifts (the median redshift  in the Chapman et al. 2003, 2005 samples is $z \\sim 2$). They are faint at optical wavelengths and were discovered in the first deep extragalactic surveys in the sub-mm wavebands (the SCUBA Cluster Lens Survey; Smail et al 1997, 2002). SMGs produce a significant fraction of the energy output of the high redshift early universe, and hence represent a cosmologically significant population (Smail et al 1997, Blain et al 2002, Blain et al 1999).  Chapman et al. (2005) find that SMGs and Lyman break galaxies contribute equally to the star formation density at $z \\sim 2-3$ and that, when extrapolated to lower fluxes, SMGs may be the dominant site of massive star formation at this epoch. Upcoming instruments, such as the $\\it{Herschel~Space~Observatory}$ and SCUBA-2, will be able to perform routine observations of SMGs at rest-frame FIR wavelengths, which is critical for observationally determining the bolometric luminosities of this high redshift, cosmologically significant population.  In contrast to local ULIRGs, the suggested formation mechanisms and evolutionary scenarios for SMGs remain varied in nature, ranging from primeval, heavily accreting galaxies undergoing a starburst (Rowan-Robinson 2000; Efstathiou \\& Rowan-Robinson 2003) to products of gas-rich major mergers undergoing intense feedback (Chakrabarti et al. 2007b), with recent observations favoring the latter scenario (Nesvadba et al. 2007; Bouche et al. 2007).  In particular, Bouche et al. (2007) suggest that dissipative major mergers may have produced the SMG population on the basis of their finding that the SMG population has lower angular momenta and higher matter densities compared to the UV/optically selected population.   A self-consistent analytic method of inferring source parameters from the observed SEDs may be an useful alternative to SED templates in analyzing upcoming FIR data sets. We give the general relations for observed quantities in terms of the redshift, and graphically depict the variation of these quantities with redshift, which is significant even at $z\\sim 1$. We show that this implies that one can estimate the value of $1+z$ for ULIRGs and SMGs from the mm and FIR SED with an accuracy $\\sim 10\\%$ (this is comparable to typical accuracies of photometric redshift codes, which estimate $z$ to typical accuracies of $\\sim 10\\%$, e.g. the IMPZ code of Babbedge et al. 2004, or the widely used HYPERZ code of Bolzonella et al. 2000),  given that our assumptions are satisfied: (1) the mm-FIR spectrum is due to reprocessing of emission from a central, dust-enshrouded source; (2) the dust can be approximated as being homogeneous and spherically distributed, with a density that varies as a power of the radius; (3) the source is sufficiently opaque that emission from the dust destruction front is negligible; and (4) the luminosity-to-mass ratio, $L/M$, of high-redshift ULIRGs and SMGs is similar to that of low-redshift ULIRGs (this last assumption is verified for the small set of SMGs for which adequate data currently exist). Chapman et al. (2005) point out that the dust temperatures inferred for SMGs are significantly lower than those of local ULIRGs and conclude that FIR photometric redshifts have an uncertainty $\\Delta z \\simeq 1$; as we shall see, our method is significantly more accurate.  We primarily focus our discussion of SMGs on sources observed recently by Kovacs et al. (2006) using the $350~\\micron$ band of SHARC-2, which is currently the most direct observational probe of the rest-frame FIR of high redshift SMGs.  The organization of the paper is as follows: in \\S 2, we review the basics of the analytic radiative transfer methodology presented in CM05 and collect the main expressions in Appendix A; in \\S 2.1, we explain the general procedure of applying our results, and contrast our solution with standard approximations in Appendix B. \\S 3 is devoted to a treatment of ULIRGs, where we infer the large scale parameters of a dozen local ULIRGs by fitting to the FIR SED.  In \\S 4 we present SED fits for a sample of SMGs. In \\S 5 we present the principal result of this paper, a method of inferring redshifts from FIR SEDs, and we demonstrate its applicability both with a simulated test case and with data of SMGs.  We conclude in \\S 6. In Appendix B we discuss standard approximations, such as the Hildebrand (1983) prescription for the mass in terms of the mm flux, and modified blackbody single temperature models in fitting the SEDs of ULIRGs and protostars (Yun \\& Carrilli 2002, henceforth YC02).  For purposes of illustration, we graphically contrast our solution and the standard approximations against the numerical results from DUSTY over the astrophysical parameter space. ",
        "conclusions": "We have applied the shell methodology for radiative transfer developed in CM05 to a range of extragalactic sources, from local ULIRGs to high redshift SMGs.  The main results are:  1. Using the general expressions for the SEDs of dusty sources given in CM05, we have shown how to derive the light-to-mass ratio, $L/M\\delta$, and the mean surface density, $\\Sigma\\delta$, of dusty galaxies at cosmological distances.  Here $M$ and $\\Sigma$ refer to the gas mass, and $L$ is the FIR luminosity as determined from observations at rest wavelengths $\\ga 60$~\\micron. The effective dust-to-gas ratio, $\\delta$, is the ratio of the actual FIR opacity to the one we have adopted, which is twice the Weingartner \\& Draine (2001) opacity; the factor 2 allows for ice mantles. Approximate expressions are given for the case in which the radius of the dusty envelope, $R_c$, is much larger than the characteristic photospheric radius, $\\rch$ (i.e., for $\\rct\\equiv R_c/\\rch\\gg 1$).  2.  The long-wavelength slope of ULIRGs can be fit with a standard dust opacity curve ($\\kappa_\\nu\\propto \\nu^2$ in the FIR) when a self-consistent radiative transfer solution is employed.  This confirms the conclusion of Dunne \\& Eales (2001), who used two-temperature fits to the SED. Comparison with the three sources in DS98 for which there are resolved measurements for the galaxies in our sample (Arp 220, Mrk 231 and IRAS 10565+2448) shows that on average our mass estimates are about 1.3 times greater than theirs, which is excellent agreement in view of the approximations in each method and the possible variations in the value of $\\delta$ from galaxy to galaxy.  3. From an analysis of the FIR data of 10 local ULIRGs, we find that the mm to FIR SED can be well-described by our simple construct of a spherically symmetric dust envelope surrounding a central source of luminosity.  We find that a density profile $\\rho \\propto r^{-2}$ provides the best to the FIR data and is consistent with CO masses.  We report our findings for the luminosities, masses, and sizes of these 10 ULIRGs.  4.  We find that local ULIRGs and high redshift SMGs ($z \\sim 2$) have similar $L/M\\delta$ ratios, with the SMGs in our sample having values $\\sim 1.45$ times smaller. The SMGs in our sample have luminosities about 5.5 times larger, and masses about 8 times larger, than the ULIRGs in our sample.   5.  We have developed a method of inferring FIR photometric redshifts. The accuracy of the method depends on whether the galaxy has a light-to-mass ratio comparable to the template, but since $1+z$ scales as only the 1/6 power of $L/M\\delta$, significant variations are allowed. Using Arp 220 as an example, we showed that our method should work for redshifts $z\\la 10$. We tested our method on a sample of five SMGs for which good photometric data are available. Under the assumption that the SMGs had the same light-to-mass ratio as a sample of local ULIRGs, we were able to infer values of $1+z$ for the SMGs with an average accuracy of 10\\% and a worst accuracy of 15\\%. If we used the average light-to-mass ratio for the SMGs, the average accuracy improved to about 5\\%. As the samples of high-redshift galaxies grow, our knowledge of both the typical light-to-mass ratio and the accuracy of FIR photometric redshifts will improve. Our method should be applied to galaxies that radiate most of their energy in the FIR, such as ULIRGs and SMGs.  Whether this condition is met can be determined by confirming that the near-IR (or optical/UV) luminosity is significantly less than that emerging in the FIR. Our method works even if AGN provide a significant fraction of the luminosity, although our sample is not large enough to determine how large the AGN fraction can become before our method breaks down. As our method is analytic, it can be employed to quickly obtain photometric redshifts of large samples of SMGs, as are expected to be detected in the FIR by the upcoming $\\it{Herschel}$ mission.  This FIR photometric redshift method provides a complementary means of inferring the redshift when near-IR methods are not available or are not viable.  6. We discuss how our approximation for the radiative transfer compares with the standard single-temperature SED in Appendix B. The accuracy of the single-temperature blackbody approximation degrades for extended envelopes, $\\rct \\ga 100$, but the approximation is typically accurate to within a factor 2 for compact envelopes, for which the source function does not probe a large range of temperatures.  The effective dust temperature in the single-temperature approximation is close to the outer core temperature.  \\bigskip \\bigskip"
    },
    "0710/0710.1677_arXiv.txt": {
        "abstract": "We report the negative results from a search for 6.7 GHz methanol masers in the nearby spiral galaxy M33. We observed 14 GMCs in the central 4 kpc of the Galaxy, and found 3$\\sigma$ upper limits to the flux density of $\\sim$9 mJy in spectral channels having a velocity width of 0.069 \\kms. By velocity shifting and combining the spectra from the positions observed, we obtain an effective 3$\\sigma$ upper limit on the average emission of $\\sim$1mJy in a 0.25 \\kms\\ channel. These limits lie significantly below what we would expect based on our estimates of the methanol maser luminosity function in the Milky Way. The most likely explanation for the absence of detectable methanol masers appears to be the metallicity of M33, which is modestly less than that of the Milky Way. ",
        "introduction": "6.7 GHz methanol masers are the second brightest maser transition ever observed, and are typically much brighter than OH masers. The properties of methanol masers in the Magellanic clouds \\citep{sinclair1992,ellingsen1994b,beasley1996} are consistent with those of our Galaxy, given appropriate consideration for different galactic properties such as metallicity. The SMC has an oxygen abundance 12 + log(O/H) = 7.96 \\citep{vermeij2002}, which is a factor $\\simeq$ 6 smaller than that of GMCs in the Milky Way \\citep{peimbert1993}.  However, 6.7 GHz methanol masers have not been discovered in galaxies beyond the Magellanic clouds. In particular, surveys have shown that there is no analog to OH megamasers at 6.7 GHz \\citep{ellingsen1994a,phillips1998,darling2003}. It would be of interest to detect methanol masers in a Milky Way like galaxy. If the number of sources detected were large, one could derive the methanol maser luminosity function since all sources would be at the same distance. Further, one could look for correlations of methanol masers with giant molecular cloud masses, other types of masers, etc. The Arecibo radio telescope offers unequaled sensitivity for targeted surveys for methanol masers, and the nearest spiral galaxy which can be observed with this instrument is M33.  It is difficult to develop an optimal strategy to search for extragalactic methanol masers since the luminosity function of methanol masers in our Galaxy is unknown. This is mostly due to difficulties in determining distances to methanol masers, which is compounded by the kinematic distance ambiguity. We have adopted the following approach. We selected sources from the general catalog \\citep{pestalozzi2005}, eliminating sources wihtin 10 degrees of longitudes 0 and 180 degrees since these suffer from large uncertainties in kinematic distance. Assuming that all masers are at their near kinematic distance, we then scaled the source flux density to the distance of M33, which we take to be 840 kpc, averaging the results of \\cite{lee2002} and \\cite{freedman1991}, which gives a result consistent with the distance determined by \\cite{kim2002}. The resulting histogram of ``expected'' methanol maser flux densities is showin in Figure \\ref{expected_hist}  \\begin{figure} \\begin{center} \\includegraphics[angle=0]{f1.eps} \\caption{Distribution of expected flux densities of methanol masers in M33 based on observed masers in the Milky Way.} \\label{expected_hist} \\end{center} \\end{figure}  We thereby obtained an estimate for the distribution of flux density of methanol masers in M33, which of course must be regarded with considerable skepticism, as the original Galactic sample has significant biases, not to mention the unknown differences in the properties of massive star formation in M33 and the Milky Way. It is also inherently pessimisitic in the sense that some of the Galactic masers are at their far kinematic distance, and hence are more luminous than we have assumed. With these caveats, we found a ``high flux density tail'' for our hypothetical methanol maser population in M33 that extends from 100 mJy down to 1 mJy. Approximately 16\\% of the total number of masers would be expected to have flux density greater than 1 mJy, with about 6\\% having flux density greater than 3 mJy. ",
        "conclusions": "It is difficult to be highly quantiative about the lack of detection of methanol maser emission in M33 beyond the statements given above. Given that the 40\\arcsec\\ Arecibo beam subtends a region 160 pc in size at the distance of M33, we should be sensitive to a methanol maser located anywhere within the individual giant molecular clouds being observed. Based on our estimate of the Milky Way methanol maser luminosity function, we would expect a significant fraction of masers to have a flux density $\\geq$ 10 mJy, significantly higher than the individual limits we have obtained, and an order of magnitude greater than the averaged limit derived by combining the results from the 14 GMCs observed.  Explanations of the rarity of methanol masers in external galaxies have focused on (1) insufficient methanol density over path where amplification could take place and (2) insufficient pumping to invert the methanol transition in question \\citep{phillips1998}. Both of these can result from low metallicity. A reduced abundance of oxygen, for example, will likely reduce both the abuance of methanol (CH$_3$OH) and of dust, which is required to convert the short wavelength radiation from massive young stars to the infrared wavelengths required for maser pumping. The rarity of masers in the Magellanic clouds has been discussed in similar terms by \\cite{beasley1996}. The O/H ratio in the Magellanic clouds is dramatically lower than that of the Milky Way (see discussion in Beasley et al. 1996 and the more recent measurement by Vermeij \\& van der Hulst 2002). For M33, the O/H ratio is only slightly less than that of the Milky Way. \\cite{vilchez1988} determined that 12 + log(O/H) = 9.0 at the center of M33, falling to $\\sim$8.5 at distances between 2 and 5 kpc from the center of the galaxy. The best fit line of \\cite{crockett2006} to their new data plus previous data has a very small slope of -0.12 dex kpc$^{-1}$ and 12 + log(O/H) = 8.3 at 5 kpc from the center of M33. These values may be compared to Galactic values of 12 + log(O/H) = 8.51 for Orion and 8.78 for M17 \\citep{peimbert1993}. The GMCs we have observed are located between 0.5 and 4 kpc from the center of M33, with a mean distance of 1.9 kpc. We conclude that the relevant O/H ratio in M33 is between 0.1 and 0.4 dex below that of the Milky Way. This difference is much smaller than that between the SMC and the Milky Way, for which log(O/H) differs by $\\simeq$ 0.8 dex. If the relatively small difference between the metallicity in the Milky wand M33 is responsible for the lack of methanol masers in the latter, it suggests that the maser luminosity must be a very senstive function of the galactic metallicity.  \\cite{henkel1987} have detected thermal methanol emission from two galaxies, IC342 and NGC253, finding that the fractional abundance of methanol is $\\simeq$10$^{-7.5}$. This is similar to that found in GMCs in the Milky Way, but is a factor of at least 100 lower than that required for high brightness methanol maser luminosity as discussed by \\cite{sobolev1997}. However, this difference is also found in comparing hot cores in the Milky Way, presumed to be the sites of methanol masers, with more extended molecular cloud regions. The methanol abundance may be greatly increased in regions near hot stars by thermal desorption of molecules frozen onto grain surfaces. We have no direct evidence regarding the methanol abundance in M33, so it is possible that the lower O/H ratio does yield an insufficient methanol abundance to produce highly luminous masers.  There are individual regions within M33, presumably powered by massive young stars \\citep{hinz2004}, which are powerful far--infrared sources. The infrared flux is a critical requirement for maser pumping as elaborated in models of \\cite{cragg1992}, \\cite{sobolev1994}, and \\cite{sobolev1997}. The transition to the second torsional excited state occurs at a wavelength of $\\simeq$ 30 \\mic, so that dust temperatures of at least 150 K are required to achieve high maser brightness \\citep{sobolev1997}. Among our targets was number 8 of \\cite{engargiola2003}, which lies within 15\\arcsec of the optical nebula NGC 604. It has an IRAS luminosity of 6.8$\\times$10$^7$ \\Ls \\citep{rice1990} and the associated GMC has a mass derived from CO of 4.4$\\times$10$^5$ \\Ms \\citep{engargiola2003}. By Galactic standards this region would seem likely to harbour a high luminosity methanol maser, but no such emisson was detected.   \\cite{fix1985} were unsuccessful in a search using the VLA for highly luminous OH masers in M33. Their limit of $\\simeq$ 25 mJy (5$\\sigma$) was sufficient to elminate the presence of any type I maser as luminous as the brightest type I masers in the Milky Way The present work thus reemphasizes the mystery of the lack of luminous masers in M33."
    },
    "0710/0710.5923_arXiv.txt": {
        "abstract": "The primary optical caustic surface behind a Kerr black hole is a four-cusped tube displaced from the line of sight. We derive the caustic surface in the nearly asymptotic region far from the black hole through a Taylor expansion of the lightlike geodesics up to and including fourth-order terms in $m/b$ and $a/b$, where $m$ is the black hole mass, $a$ the spin and $b$ the impact parameter. The corresponding critical locus in the observer's sky is elliptical and a point-like source inside the caustics will be imaged as an Einstein cross. With regard to lensing near critical points,  a Kerr lens is analogous to a circular lens perturbed by a dipole and a quadrupole potential. The caustic structure of the supermassive black hole in the Galactic center could be probed by lensing of low mass X-ray binaries in the Galactic inner regions or by hot spots in the accretion disk. ",
        "introduction": "Black holes (BHs) are among the most fascinating predictions of the general theory of relativity. Recent progresses in mass measurements of compact objects have supplied compelling evidence for their existence. Nearly every galaxy is supposed to contain a very massive BH at its center in the mass range from $10^{6}$ to $10^{9}~M_\\odot$ \\citep{na+qu05}. Despite strong hints that massive BHs mostly rotate rapidly, their spins are still unmeasured so that a full characterization of their properties is not possible. Analyses of luminosities and spectra of accretion disks around spinning BH are very promising tool, but the interpretation can be sometimes unclear due to uncertainties in the physics of the inflowing gas. The classical test of gravitational deflection of light rays passing near compact bodies can provide an alternative probe. In fact, the theoretical bases of lensing  are  well understood and, on the observational side, the supermassive BH supposed to be hosted in the radio source Sgr~A* in the Galactic center, with a mass of $\\sim 3.6\\times 10^6 M_\\odot$ and at a distance of $7.6~\\mathrm{Kpc}$ from the Earth \\citep{eis+al05}, offers a very appealing target for future space- and ground-based experiments.  Lensing by rotating BHs has been studied from the very beginning of the Kerr spacetime \\citep{bo+li67}. Angular momentum first appears in gravitational lensing through terms $\\sim  m a /b^2$ in the deflection angle. Up to this order, light deflection is well understood. Very different approaches  can be undertaken \\citep{pi+ro77,ep+sh80,iba83,ba+ho04,bra86,dym86,gli99,kop97,ko+sc99,ko+ma02,as+ka00,ser03b} to show the degeneracy between a Kerr BH and a suitably displaced Schwarzschild lens \\citep{asa+al03,we+pe07}.  Such  analyses can be  easily extended to general spinning mass distributions \\citep{ser02,se+ca02,ser03,ser05,ser07} and show many interesting features in the magnification pattern and in the image properties \\citep{ser03,ser05,we+pe07} .  Investigations up to higher orders require the full consideration of the lightlike geodesics \\citep{car68,cha83}.  The optical structure of the primary caustic surface \\cite{ra+bl94}  as well as the appearance of both stars \\citep{cu+ba73} and accretion disk \\citep{vie93,be+do05} orbiting a Kerr BH have been detailed through numerical investigations. A clear analytical picture of the relativistic caustics in the strong deflection limit has also emerged \\citep{boz+al05,bo+sc07}. The missing part, which we are going to provide here, is an analytical treatment of map singularities and caustics in the weak deflection limit. This is the prerequisite for the study of critical points and lensing map inversion near them.   The paper is organized as follows. In Section~\\ref{sec:geod}, we recall basic properties of lightlike geodesics in Kerr spacetime. Sections~\\ref{sec:caus} discusses the singularities of lens mapping and lensing near caustics. Section~\\ref{sec:disc} is devoted to some final considerations. In this paper, we will use units $G=c=1$, with $c$ the light speed in the vacuum, so that the constant $m (\\equiv G M /c^2)$ is the gravitational radius. ",
        "conclusions": "\\label{sec:disc}  The analytical derivation performed in this paper corroborates and extends knowledge about the primary caustics which, in a earlier work \\citep{ra+bl94}, have been investigated only numerically and provides the first study, to our knowledge, of critical points and lens mapping near them. Our approach is the natural complement to qualitative methods which give some information on lensing properties without actually solving the equation for lightlike geodesics, such as the Morse theory \\citep{ha+pe06}.  Together with the theoretical motivation, an equally compelling reason for investigating gravitational lensing in a Kerr spacetime comes from lensing observations towards Sgr~A* which are coming into the reach of observability and for which the  weak-deflection limit approximation at the lowest order is not applicable. The stars surrounding the Galactic center have been considered as suitable targets for detection of lensing effects. Sgr~A* is expected to be lensing nearly ten sources at any given time for observations down to $K \\sim 21$ \\citep{jar98,al+st99,cha+al01} . Considering the only few stars whose orbital parameters have already been accurately determined, detectable lensing events are expected to occur in a temporal span of $\\sim 30$~years \\citep{bo+ma05}. The typical  radius of these sources ($\\sim 10^{8}-10^{9}~\\mathrm{m}$) is much larger than the caustic width, so that the effects of the finite astroid size are washed out. Nevertheless, the astrometric shift of the caustic center, together with the precession orbital effects related to the spinning central body, must be accounted for when considering future experiments.  Other appealing sources for lensing by Sgr~A* are low mass X-ray binaries (LMXBs), whose emission mostly originates in a region a few tens of kilometers across, consisting of the inner accretion disk around the BH accreting from the companion. Tens of thousands of stellar-masses BHs and neutron stars are likely to have settled dynamically into the central parsec of the Galaxy  and perhaps few hundreds of them might have stellar companions \\cite{mun+al05}. Several tens of LMXBs have been already detected in the very inner regions down to a minimum projected distance of only $0.1~\\mathrm{pc}$ \\citep{mun+al05b}, with a detected overabundance of transients X-ray binaries within $1~\\mathrm{pc}$ \\citep{mun+al05}.  These sources have been considered for detection of relativistic images \\citep{boz+al05}  and could as well, since their small radius, probe the primary caustic surface.  Compact emission regions in the clumpy and unsteady accretion flow near Sgr~A* are other interesting lensing sources \\citep{br+lo05,br+lo06}.  Relativistic images of orbiting bright-spots could be detected in the near future with observations achieving submilliarcsecond resolution at infrared and submillimetre wavelengths. However, we remark that in our approximation the source distance must be large and close inner orbits can not be considered.  The finite size of the primary caustics behind Sgr~A* could then be probed by lensing of either LMXBs or compact hot spots in the accretion flow whereas the effects of angular momentum, such as the shift in the caustic position, affect even lensing of main sequence stars in the Galactic center cluster."
    },
    "0710/0710.5112_arXiv.txt": {
        "abstract": "Physical sizes of extended radio galaxies can be employed as a cosmological ``standard ruler'', using a previously developed method. Eleven new radio galaxies are added to our previous sample of nineteen sources, forming a sample of thirty objects with redshifts between 0 and 1.8.  This sample of radio galaxies are used to obtain the best fit cosmological parameters in a quintessence model in a spatially flat universe, a cosmological constant model that allows for non-zero space curvature, and a rolling scalar field model in a spatially flat universe. Results obtained with radio galaxies are compared with those obtained with different supernova samples, and with combined radio galaxy and supernova samples.  Results obtained with different samples are consistent, suggesting that neither method is seriously affected by systematic errors. Best fit radio galaxy and supernovae model parameters determined in the different cosmological models are nearly identical, and are used to determine dimensionless coordinate distances to supernovae and radio galaxies, and distance moduli to the radio galaxies.  The distance moduli to the radio galaxies can be combined with supernovae samples to increase the number of sources, particularly high-redshift sources, in the samples.  The constraints obtained here with the combined radio galaxy plus supernovae data set in the rolling scalar field model are quite strong.  The best fit parameter values suggest that $\\Omega_m$ is less than about 0.35, and the model parameter $\\alpha$ is close to zero; that is, a cosmological constant provides a good description of the data.  We also obtain new constraints on the physics of engines that power the large-scale radio emission. The equation that describe the predicted size of each radio source is controlled by one model parameter, $\\beta$, which parameterizes the extraction of energy from the black hole. Joint fits of radio galaxy and supernova samples indicate a best fit value of $\\beta$ that is very close to a special value for which the relationship between the braking magnetic field strength and the properties of the spinning black hole is greatly simplified, and the braking magnetic field strength depends only upon the spin angular momentum per unit mass and the gravitational radius of the black hole.  The best fit value of $\\beta$ of 1.5 indicates that the beam power $L_j$ and the initial spin energy of the black hole $E$ are related by $L_j \\propto E^2$, and that the relationship that might naively be expected for an Eddington limited system, $L_j \\propto E$, is quite clearly ruled out for the jets in these systems. ",
        "introduction": "Recent cosmological studies using CMBR, supernovae, and other types of astronomical sources and phenomena have greatly improved our understanding of the recent expansion and acceleration history of the universe.  Whereas a consistent picture has emerged (the ``concordance cosmology''), in which the dynamics of the universe is currently dominated by a mysterious dark energy, its physical nature remains one of the key outstanding problems of physical science today. For a summary of developments in this field see Ratra \\& Vogeley (2007).  Aside from the CMBR, most experimental methods to study the expansion history of the universe, and thus its matter-energy contents, require samples of standardisable sources whose distances can be determined in a consistent way, e.g., the supernovae of type Ia. It is clear that both low and high redshift sources play an important role in these studies. Low to moderate redshift sources allow us to define and probe the acceleration of the universe and the properties of the dark energy at the current epoch, whereas higher redshift sources allow us to probe the properties of the dark energy at earlier epochs and possible changes in its properties, which is perhaps the best path towards understanding its physical nature. Confidence in luminosity and coordinate distance determinations, which are the foundation of the studies discussed here, is bolstered when more than one method yields the same results.  Like supernovae, powerful radio galaxies are observed out to redshifts greater than one, and their observable properties can be used to determine the coordinate distances to them, or equivalently the luminosity distances or distance moduli (Daly 1994). It is thus interesting to study these sources and to determine whether cosmological results obtained with radio galaxies agree with those obtained with supernovae and other methods. In addition, the radio galaxy and supernova samples may be analyzed jointly to improve the determinations of the radio galaxy model parameters.  In this paper, eleven new radio galaxies are combined with a previously studied sample of nineteen sources, to yield a sample of thirty radio galaxies suitable for cosmological studies. This sample is analyzed here in three cosmological models. The first model allows for quintessence and non-relativistic matter in a spatially flat universe; the second model allows for non-relativistic matter, a cosmological constant, and space curvature; and the third model allows for a rolling scalar field in a spatially flat universe (Peebles \\& Ratra 1988). All of these models are based on the equations of General Relativity, and rely upon this as the correct theory of gravity.  The primary objectives are to obtain and compare constraints on model and cosmological parameters using different samples of type Ia supernovae, extended radio galaxies, and a combined sample of supernovae and radio galaxies. Similar results obtained with the radio galaxy and supernovae methods, which determine similar quantities over similar redshifts, will bolster our confidence in each method. This is because the methods rely upon measurements of entirely different quantities that are then applied using completely different astrophysical arguments, and thus provide independent measures of coordinate distances and cosmological parameters. Results obtained from the combined sample allow strong constraints to be placed on the radio galaxy model parameter, which provides a direct link to and diagnostic of the physics of the energy extraction that produces the large-scale jets in these systems. Finally, if very similar best fit model parameters are obtained in different cosmological scenarios, tben these parameters can be used to determine the dimensionless coordinate distance to each source.  The dimensionless coordinate distances thus obtained are then used for a separate study.  In addition, the dimensionless coordinate distances to the radio galaxies can be used to define the distance modulus to each radio galaxy, which can be combined with those of supernovae to increase the sample sizes.  The use of powerful extended radio galaxies as a cosmological tool is reviewed in section 2.  Results obtained with radio galaxies, supernovae, and combined samples of radio galaxies and supernovae are presented in section 3.  Determinations of distance moduli to radio galaxies are presented in section 4.  The main results and conclusions of the paper are summarized in section 5. ",
        "conclusions": "A sample of thiry radio galaxies was used to determine cosmological parameters and the radio galaxy model parameters $\\kappa_{RG}$ and $\\beta$ in three standard cosmological models. Nearly identical values of $\\kappa_{RG}$ and $\\beta$ are obtained in each cosmological model indicating that they are not strongly affected by the context in which they are determined (see Tables 1, 2, and 3), thus they can be used to determine the dimensionless coordinate distance to each source.  Three supernovae samples are considered both separately and jointly with the radio galaxy sample and are analyzed in the context of three standard cosmological models (see Tables 1, 2, and 3) to determine cosmological parameters and the model parameter $\\kappa_{SN}$.  Nearly identical values of $\\kappa_{SN}$ are obtained for each supernovae sample in the context of each cosmological model, thus they can be used to determine the dimensionless coordinate distance to each source.  They can also be used to determine the effective distance modulus to each of the radio galaxies, which can then be combined with those of the supernovae to increase the sample sizes, particularly at high redshift.  Constraints on cosmological parameters obtained with radio galaxies are consistent with, though weaker than, those obtained with supernovae alone.  There are no inconsistencies between results obtained with radio galaxies and supernovae.  For example, in the context of a standard quintessence model, radio galaxies alone indicate that the universe is accelerating today with about 90 \\% confidence (see Figure 6). The consistency between results obtained with radio galaxies alone, supernovae alone, and the combined supernovae and radio galaxy samples suggests that neither method is plagued by unknown systematic errors; both methods seem to be working well. The radio galaxy and supernovae methods are completely independent, based on a completely different physics and observations, and provide independent measures of distances to sources at similar redshift. The facts that the cosmological parameters obtained in specific models are consistent, and that the coordinate distances to sources at similar redshift are consistent suggests that systematic errors are not playing a major role in either method.   Since nearly identical values of $\\kappa_{SN}$, $\\kappa_{RG}$, and $\\beta$ are obtained in each cosmological model, they can be used to solve for the dimensionless coordinate distance to each source, $y$.  Values of $y$ to each radio galaxy, obtained using the best fit values of $\\kappa_{RG}$ and $\\beta$ indicated by fits to the joint sample of 192 supernovae and 30 radio galaxies, are listed in Table 4. The best fit values of $y$ listed in Table 4 are combined with the best fit value of $\\kappa_{SN}$ obtained for the Davis et al. (2007) sample to obtain the distance modulus to each of the radio galaxies, $\\mu_D$, which are listed in Table 4 and which can be added to those already listed in Davis et al. (2007).  Similarly, the values of $y$ to the radio galaxies obtained using the best fit values of $\\kappa_{RG}$ and $\\beta$ indicated by fits to the joint sample of 182 supernovae and 30 radio galaxies were obtained (but are not listed), and were combined with the best fit value of $\\kappa_{SN}$ obtained for the Riess et al. (2007) sample to obtain values of $\\mu_R$ for the radio galaxies that can be combined with those already listed by Riess et al. (2007). Since nearly identical values of $\\kappa_{SN}$, $\\kappa_{RG}$, and $\\beta$ and their uncertainties are obtained in each cosmological model, the average of the values obtained in the different cosmological models and their uncertainties were used.  New constraints were obtained on the model parameter $\\alpha$ in the rolling scalar field model of Peebles \\& Ratra (1988). These constraints are rather strong and indicate that $\\alpha$ is close to zero for reasonable values of $\\Omega_m$. Thus, a cosmological constant provides a good description of the data, and there is no indication that the energy density of the dark energy is changing with redshift.  New constraints were obtained on the radio galaxy model parameter $\\beta$, suggesting values of $\\beta$ close to 1.5. This is interpreted in the context of a standard magnetic braking model of energy extraction from a rotating black hole (e.g. Blandford 1990).  This is a very special value of $\\beta$ for which the braking magnetic field strength depends only upon the spin angular momentum per unit mass and the gravitational radius of the black hole.  This suggests that when the magnetic field strength reaches this maximum or limiting value, the relativistic outflow is triggered. The fact that the  magnetic field strength does not depend explicitly on the black hole mass for this special value of $\\beta$ may explain why it is that this paricular type of radio source is able to provide a modified standard yardstick for cosmological studies.  We have provided further evidence that radio galaxies can be used to determine coordinate distances to sources, and thus cosmological parameters through a standard angular diameter test. Our comparative study shows that values of cosmological parameters obtained with radio galaxies are consistent with those obtained with supernovae. Supernovae alone and radio galaxies alone both indicate that the universe is accelerating at the current epoch when these data are analyzed in specific models such as a a quintessence model in spatially flat universe, a lambda model in universe that allows for non-zero space curvature, and a rolling scalar field model in a spatially flat universe.  All of these models rely upon the equations of General Relativity, and the results obtained in these models would not be correct if General Relativity is not the correct theory of gravity. These data are analyzed in a model-independent way that does not rely upon General Relativity by Daly et al. (2008), who show that the supernovae data alone and the radio galaxy data alone indicate that the universe is accelerating at the current epoch independent of whether General Relativity is the correct theory of gravity. When expressed as coordinate distances, both radio galaxy and supernova samples can be combined, and used in a joint cosmological analysis, as it was done, e.g., by Daly \\& Djorgovski (2003, 2004).  We use these expanded and combined samples in a separate paper (Daly et al. 2008)."
    },
    "0710/0710.2397_arXiv.txt": {
        "abstract": "{The initial-final mass relationship (IFMR) for stars is important in many astrophysical fields, such as the evolution of galaxies, the properties of type Ia supernovae (SNe Ia) and the components of dark matter in the Galaxy.} {The purpose of this paper is to obtain the dependence of the IFMR on metallicity.} {Following Paczy\\'{n}ski \\& Zi\\'{o}lkowski (1968) and Han et al. (1994), we assume that the envelope of an asymptotic giant branch (AGB) or a first giant branch (FGB) star is lost when the binding energy of the envelope is equal to zero ($\\Delta W=0$) and the core mass of the AGB star or the FGB star at the point ($\\Delta W=0$) is taken as the final mass. Using this assumption, we calculate the IFMRs for stars of different metallicities.} {We find that the IFMR depends strongly on the metallicity, i.e. $Z=0.0001, 0.0003, 0.001, 0.004, 0.01, 0.02, 0.03, 0.04, 0.05, 0.06, 0.08$ and $0.1$. From $Z=0.04$, the final mass of the stars with a given initial mass increases with increasing or decreasing metallicity. The difference of the final mass due to the metallicity may be up to 0.4 $M_{\\odot}$. A linear fit of the initial-final mass relationship in NGC 2099 (M37) shows a potential evidence of the effect of metallicity on the IFMR. The IFMR for stars of $Z=0.02$ obtained in the paper matches well with those inferred observationally in the Galaxy. For $Z\\geq 0.02$, helium WDs are obtained from the stars of $M_{\\rm i}\\leq 1.0 M_{\\odot}$ and this result is upheld by the discovery of numerous low-mass WDs in NGC 6791 which is a metal-rich old open cluster. Using the IFMR for stars of $Z=0.02$ obtained in the paper, we have reproduced the mass distribution of DA WDs in Sloan DR4 except for some ultra-massive white dwarfs.} {The trend that the mean mass of WDs decreases with effective temperature may originate from the increase of the initial metallicities of stars. We briefly discuss the potential effects of the IFMR on SNe Ia and at the same time, predict that metal-rich low-mass stars may become under-massive white dwarfs.} ",
        "introduction": "\\label{sect:1} White dwarfs (WDs) are the endpoint of the evolution of stars with initial masses ranging from about 0.1 $M_{\\odot}$ to about 8 $M_{\\odot}$. The vast majority of stars in the Galaxy belong to the mass range and over 97\\% of the stars in the Galaxy will eventually end up as WDs (Fontaine et al. \\cite{FON01}). As a result, WDs may give direct information about star formation during the Galaxy's earliest epochs (Gates et al. \\cite{GAT04}). WDs show their importance in many fields. It has been about 50 yrs since Schmidt (\\cite{SCH59}) recognized the usefulness of WDs as cosmochronometers. WDs may also be the dominating component of dark matter in the Galaxy (Alcock et al. \\cite{ALC99}; Chabrier \\cite{CHA99}). As a matter of fact, WDs are very important for type Ia supernovae (SNe Ia) since it is believed that SNe Ia are from the thermonuclear runaway of carbon-oxygen white dwarfs (CO WD) (see the reviews by Hillebrandt \\& Niemeyer (\\cite{HN00}) and Leibundgut (\\cite{LEI00})). It is well known that the masses of WDs are typically of the order of half that of the Sun, while their radii are similar to that of a planet. However, a detailed knowledge of WDs is still unclear in observation and theory (Fontaine et al. \\cite{FON01}; Moroni \\& Straniero \\cite{MOR02}, \\cite{MOR07}). Among all the uncertainties of WDs, the correlation between initial-final mass relationship (IFMR) and metallicitiy is a very important one. It is well known that low metallicity leads to a larger CO WD for a given initial mass with $Z\\leq 0.02$ (Umeda et al. \\cite{UME99a}). However, it is necessary to check the cases of $Z>0.02$ since $Z\\in[0.06-0.1]$ is possible for some ultra-luminous galaxies (Roberts \\& Hynes \\cite{RH94}; Ruiz-Lapuente et al. \\cite{RUI95}; Terlevich \\& Forbes \\cite{TER02}).  The IFMR for stars over a large mass range (e.g. 0.8-8 $M_{\\odot}$) is a powerful input to chemical evolution models of galaxies (including enrichment in the interstellar medium) and therefore can enhance our understanding about star formation efficiencies in these systems (Ferrario et al. \\cite{Fer05}; Kalirai et al. \\cite{KAL07b}). The IFMR is also an important input for modelling the luminosity functions of Galactic disk WDs and the cooling sequences of halo clusters, which may directly yield the age of the Galactic disk and halo components (Kalirai et al. \\cite{KAL07b}). Meanwhile, the IFMR represents the mass loss of a star over its entire evolution and it is possible to get some indications about the origin and evolution of hot gas in elliptical galaxies from the IFMR (Mathews \\cite{MAT90}). Since Weidemann (\\cite{WEI77}) showed the first comparison between observations and theoretical predications of the IFMR, many observations gave constraints on the relationship by studying the properties of white dwarfs in open clusters or field white dwarfs, especially during the last decade (Herwig \\cite{HER95}; Reid \\cite{REI96}; Koester \\& Reimers \\cite{KR96}; Finley \\& Koester \\cite{FK97}; Claver et al. \\cite{CLA01}; Williams et al. \\cite{WIL04}; Ferrario et al. \\cite{Fer05}; Williams et al. \\cite{WIL06}). Some empirical relations have also been suggested based on different observations (Weidemann \\cite{WK83}, \\cite{WEI00}; Ferrario et al. \\cite{Fer05}; Williams et al. \\cite{WIL06}, Dobbie et al. \\cite{DOB06b}). At the same time, numerous sets of stellar evolution models have been calculated to study the relationship in theory and to give the IFMRs for stars of different metallicity (Han et al. \\cite{HAN94}; Girardi et al. \\cite{GIR00}). Although great progress has been made in observation and theory, there still exist many uncertainties, i.e. what is the origin of the intrinsic scatter of WD mass or whether there is any dependence of the IFMR on the metallicity (Kalirai et al. \\cite{KAL05}; Williams \\cite{WIL06}). The correlation between the IFMR and the metallicity has been established in theory for many years (Han et al. \\cite{HAN94}; Girardi et al. \\cite{GIR00}), but the evidence of the dependence of the IFMR on the metallicity was not found until 2005 (Kalirai et al. \\cite{KAL05}). Many theoretical calculations showed that the core mass at the first thermal pulse (TP) in the asymptotic giant branch (AGB) may be taken as the final mass (see the review by Weidemann \\cite{WEI00}). However, these theoretical studies only focused on some special metallicities and it is also difficult for some stellar evolution codes to determine which is the first thermal pulse. In the paper, we use a simple but robust method to systemically study the IFMR over a wide metallicity range, i.e. $0.0001\\leq Z\\leq 0.1$. We also briefly discuss the potential effect of the IFMR on SNe Ia.  In section \\ref{sect:2}, we describe our model and physical inputs. We give the results in section \\ref{sect:3} and show discussions and conclusions in section \\ref{sect:4}.  \\begin{figure*} \\centering \\includegraphics[width=150mm,height=160mm,angle=270.0]{mimf2.ps} \\caption{Initial-final mass relationship (IFMR) for stars of different metallicities. The points of $M_{\\rm f}< 0.5 M_{\\odot}$ represent helium WDs. The thick solid lines are the mass threshold for the second dredge-up (see the text for details) and the thick dotted ones show the transition between carbon-oxygen white dwarfs and oxygen-neon-magnesium ones in our models. The dashed line in the lower panel shows the IFMR for stars of $Z=0.04$ as a comparison. } \\label{Fig1}% \\end{figure*} ",
        "conclusions": "\\label{sect:4} \\subsection{the uncertainty of the initial-final mass relationship }\\label{subs:4.1} The major uncertainty of the IFMR is from the assumption that the final mass is equal to the core mass when the binding energy of the envelope $\\Delta W=0$ for an AGB/FGB star. Although the consistency between the mass distribution of WDs in the paper and that from Sloan DR4 upholds the assumption, there is no direct observational evidence to verify it. Another widely accepted choice of the final mass is the core mass of a star at the beginning of the thermal-pulse AGB (TPAGB, the first TP or the end of early AGB) (see the review by Weidemann \\cite{WEI00}). From this choice, taking $Z=0.02$ as an example, the general trend of the IFMR is divided into three segments as follows: (i), the core mass is almost constant from 1 $M_{\\odot}$ to 2.5 $M_{\\odot}$, around 0.52 $M_{\\odot}$, which is smaller than that in this paper by about 0.03-0.1 $M_{\\odot}$. (ii), the core mass strongly increases with the initial mass up to 4.0 $M_{\\odot}$. This trend is similar to that in this paper, but the core mass is smaller than that in this paper by 0.03-0.05 $M_{\\odot}$. (iii), the core mass slowly increases with the initial mass up to 6-7 $M_{\\odot}$. The final mass is smaller than that in this paper by 0.1-0.2 $M_{\\odot}$. According to above discussion , the maximum variation of the final mass derived from different assumptions may be as large as 0.2 $M_{\\odot}$.  In this paper, we skipped the detailed evolution of TP and simply treated it as an average evolution, which may lose some information on the structural change of the envelope caused by thermal pulse at the early phase of thermally pulsing AGB. We therefore chose a model with $M_{\\rm i}=3.00 M_{\\odot}$ and $Z=0.02$ to examine the influence of thermal pulses on $\\Delta W$ and the result is shown in Fig. \\ref{Fig8}. From the figure, we see that there is a time that $\\Delta W$ changes its sign but returns back immediately during a full amplitude of TP. The change of the sign of $\\Delta W$ is from the decrease of stellar radius which results from the decrease of total luminosity (see panels 3 and 4). This implies that superwind might be periodically interrupted, but the timescale of the interruption is very short ($\\sim100$yr), and then its influence is small. There are many studies focusing on the TPs previously and we briefly summarize them in the following. Generally, the final mass after several TPs is larger than the core mass at the first TP. Forestini \\& Charbonnel (\\cite{FC97}) calculated a model of $M_{\\rm i}=3 M_{\\odot}$ assuming a Reimers mass-loss law with $\\eta$ increasing from 2.5 to 5.0 in the final AGB stage. They obtained $M_{\\rm f}=0.60M_{\\odot}$ after 19 TPs, but the core mass is 0.54$M_{\\odot}$ at the first TP. Dominguez (\\cite{DOM99}) showed that the thermal pulse starts when the core mass is 0.571$M_{\\odot}$ for a star of $M_{\\rm i}=3 M_{\\odot}$, while the final mass increases to $M_{\\rm f}=0.71 M_{\\odot}$ after 26 TPs if the mass loss prescription of Groenewegen \\& de Jong (\\cite{GRO94}) was adopted (Straniero et al. \\cite{STRA97}). If the fairly strong mass-loss law of Bl$\\ddot{\\rm o}$cker (\\cite{BLOC95}) was adopted, the final mass becomes 0.63$M_{\\odot}$ after 20 TPs from 0.53$M_{\\odot}$ at the first TP (Bl$\\ddot{\\rm o}$cker \\cite{BLOC95}). However, the final mass of 0.68 $M_{\\odot}$ is obtained from the model with $M_{\\rm i}=3.0 M_{\\odot}$ and $Z=0.02$ in this paper.   {In Fig. \\ref{Fig8}, we see that $\\Delta W=0$ for a star with $M_{\\rm i}=3.00 M_{\\odot}$ and $Z=0.02$ is achieved before TPAGB phase, which might be very important for galactic chemical evolution since it is widely believed that s-process elements are created in TPAGB phase of a low-mass star (Han et al. \\cite{HAN95}; Busso et al. \\cite{BGW99}). Actually, $\\Delta W=0$ for all the models in this paper is achieved at EAGB phase. These results seem to indicate that the s-process elements could not be created. In fact, as mentioned in subsection \\ref{subs:2.1}, it is more likely that $\\Delta W=0$ is only a lower time limit for superwind, i.e. a superwind starts at or soon after $\\Delta W=0$. Therefore, the stars achieving $\\Delta W=0$ at EAGB phase still have opportunities to enter into TPAGB phase, and contribute to s-process elements.} So, the final mass shown in this paper might be different from a real one. As mentioned above, the core mass in 3 $M_{\\rm \\odot}$ models grows by about 0.1 $M_{\\rm \\odot}$ during TPAGB evolution. For a 1.5 $M_{\\rm \\odot}$ star with $Z=0.02$, the final core mass is about 0.03 $M_{\\rm \\odot}$ larger than that in this paper if the mass loss prescription of Groenewegen \\& de Jong (\\cite{GRO94}) was adopted (Dominguez \\cite{DOM99}). The main result in this paper then still holds.   In Fig. \\ref{Fig4}, we see that when $M_{\\rm i}<3.5M_{\\odot}$, the IFMR in this paper is well consistent with the empirical relations. However, for the higher initial mass, the IFMR in this paper is larger than the empirical relations given by Weidemann (\\cite{WEI00}) and Ferrario et al. (\\cite{Fer05}). This discrepancy is mainly from our assumption. As mentioned in subsection \\ref{subs:2.1} and above paragraph, $\\Delta W\\geq0$ is a necessary condition to eject the envelope of an AGB star. Since $\\Delta W=0$ is achieved at the second dredge-up for the star with $M_{\\rm i}>3.5M_{\\odot}$ during EAGB phase and the core is decreasing at this phase, the necessary condition means that final mass may be overestimated. The uncertainty of $M_{\\rm f}$ attributed to the definition of core in this paper may also reduce the discrepancy. Meanwhile, as mentioned in subsection \\ref{subs:3.1}, the final mass might be overestimated if the helium envelope of a star reaches $\\Delta W=0$ again after the ejection of hydrogen-rich envelope. This overestimate is more likely for high initial mass, since the helium envelope is thicker, and then reaches $\\Delta W=0$ more easily (see subsection \\ref{subs:3.1}). However, since the errors from the observations are also very large (see the error bar in Fig. \\ref{Fig4}), the IFMR obtained in this paper is well located in the error range and therefore, is still consistent with observations.      \\subsection{low-mass white dwarf}\\label{subs:4.2} Kalirai et al. (\\cite{KAL07a}) suggested that helium white dwarfs may be obtained in metal-rich old clusters, i.e. NGC 6791. It is a reasonable assumption that the same mechanism in NGC 6791 would work for metal-rich field stars. Recently, Kilic, Stanek \\& Pinsonneault (\\cite{KIL07c}) rehandled the data in Valenti \\& Fischer (\\cite{VF05}) and found that there have been metal-rich stars with [Fe/H]$>$ 0 at all times in the local Galactic disk, although the metallicity distribution of disk stars peaks below solar metallicity for stars with ages greater than about 5 Gyrs. Considering that only 5\\% of all WDs in the local disk are single low-mass white dwarfs ($<$ 0.45$M_{\\odot}$) (Liebert et al. \\cite{LIB05}; Kepler et al. \\cite{KEP06}) and the fraction of metal-rich stars with ages greater than 9 Gyrs is 21\\%, Kilic, Stanek \\& Pinsonneault (\\cite{KIL07c}) argued that only the stars of [Fe/H]$>$+0.3 can lose their hydrogen-rich envelope to produce helium WDs on the FGB. From our study, however, all stars with $M_{\\rm i}\\leq1.0 M_{\\odot}$ and $Z\\geq0.02$ will lose their hydrogen-rich envelope and finally become helium white dwarfs. This inconsistency may result from the overestimate of metal-rich stars in the sample of Valenti-Fischer (Reid et al. \\cite{REI07}) adopted by Kilic, Stanek \\& Pinsonneault (\\cite{KIL07c})  As shown in our study, there should be numerous He WDs in metal-rich old clusters, which has indeed been observed in NGC 6791 (Kalirai et al. \\cite{KAL07a}). Here, we emphasize that we do not assume any mass-loss mechanism while Kalirai et al. (\\cite{KAL07a}) assumed a metal-enhanced stellar wind on the red giant branch (RGB). Our study shows that $\\Delta W = 0$ in several FGB stars with high metallicities and small initial masses can be achieved before these stars get to the tip of FGB, indicting that the relative number of the tip RGB stars in metal-rich old clusters is smaller than that in metal-poor clusters. This fact is directly observed in the RGB luminosity functions of two open clusters, e.g. NGC 188 and NGC 6791 (see Fig 8 in Kalirai et al \\cite{KAL07a}). Meanwhile, a similar effect should be seen in the local population of RGB stars. Luck \\& Heiter (see Fig 9 in their paper, \\cite{LH07}) make a comparison of the metallicity histograms between dwarfs and giants within 15 pc of the Sun and found that the dwarf population has a high metallicity tail extending up to [Fe/H]$\\sim$ 0.6, while the giants show a sharp drop in numbers after [Fe/H]=0.2 and no giants with [Fe/H]$>$ 0.45 are observed in the field.  The fate that low-mass metal-rich stars will become He WDs before helium is ignited on FGB might give some constrain on planetary nebulae (PNe). It is widely accepted that PNe originate from AGB stars or are related to binary evolution. For the case of binary evolution, one component in the binary system fills its Roche lobe at AGB phase and the mass transfer is dynamically unstable which leads to a common envelope phase. After the ejection of the common envelope, a PN may form. However, since some metal-rich stars may not experience AGB phase, we might speculate that the number of PNe in metal-rich galaxies is relatively lower than that in metal-poor galaxies. Interestingly, Buzzoni et al (\\cite{BUZ06}, see their Fig 11) really observed a relatively low number of PNe per unit galaxy luminosity in more metal-rich elliptical galaxies. Gesicki \\& Zijlstra (\\cite{GZ07}) compared the mass distribution of the central stars of planetary nebulae (CSPN) with those of WDs. These two distribution are very different, i.e. the CSPN mass distribution is sharply peaked at 0.61 $M_{\\odot}$ ranging from 0.55 $M_{\\odot}$ to 0.66 $M_{\\odot}$, while the WD distribution peaks at a slightly lower mass and shows a much broader range of masses. Gesicki \\& Zijlstra (\\cite{GZ07}) suggested that this difference may imply that only some WDs have gone through the PN phase. Our models provide a channel for WDs to avoid the PN phase.   Based on our results, WDs from single stars are always larger than 0.4 $M_{\\odot}$, whatever the composition of the WD is. Then, extremely low-mass WDs ($~0.2 M_{\\odot}$) are possible in binary systems or once in binary systems since no stellar population is old enough to produce such extremely low-mass WDs through single star evolution. Only recently, several extremely low-mass ($~0.2 M_{\\odot}$) WDs were discovered in the field (Liebert et al. \\cite{LIB04}; Kawka et al. \\cite{KAW06}; Eisenstein et al. \\cite{EIS06}; Kilic et al. \\cite{KIL07a}), and some of them are companions of pulsars (van Kerkwijk et al. \\cite{VANK05}). Kilic et al. (\\cite{KIL07b}) also found a low-mass WD in a binary system. If an extremely low-mass WD were a single star, it is very likely that it could have a very high space velocity since it may come from a close double-degenerate binary, where its companion has gone through a supernova event that disrupted the binary (Hansen \\cite{HANS03}). LP 400-22 might be a case from this channel (Kawka et al. \\cite{KAW06}). In any case, hard work is need to systematically search for the companions of WDs with mass less than 0.4 $M_{\\odot}$.   \\subsection{the potential effect of the initial-final mass relationship on type Ia supernovae}\\label{subs:4.3} As the best cosmological distance indicators, Type Ia supernovae (SNe Ia) have been successfully applied to determine the cosmological parameters ,e.g. $\\Omega_{\\rm M}$ and $\\Omega_{\\rm \\Lambda}$ (Riess et al. \\cite{REI98}; Perlmutter et al. \\cite{PER99}). Phillips relation (Philips \\cite{PHI93}) is used when taking SNe Ia as the distance indicators. It is assumed that Phillips relation is correct at high redshift, although the relation was obtained from a low-redshift sample. This assumption is precarious since the exact nature of SNe Ia is still unclear, especially the progenitor model and explosion mechanism (Hillebrandt \\& Niemeyer \\cite{HN00}; Leibundgut \\cite{LEI00}). If the properties of SNe Ia evolve with redshift, the results for cosmology might be different. Since metallicity decreases with redshift, a good way to study the correlation between the properties of SN Ia and redshift is to study the correlation between the properties of SN Ia and metallicity. It is widely believed that SNe Ia are from thermonuclear runaway of a carbon-oxygen white dwarf (CO WD) in a binary system. The CO WD accretes material from its companion to increase its mass. When its mass reaches its maximum stable mass, it explodes as a thermonuclear runaway and almost half of the WD mass is converted into radioactive nickel-56 (Branch \\cite{BRA04}). The mass of nickel-56 determines the maximum luminosity of SNe Ia. The higher the mass of nickel-56 is, the higher the maximum luminosity is (Arnett \\cite{ARN82}). Some numerical and synthetical results showed that metallcity may affect the final amount of nickel-56, and thus the maximum luminosity of SNe Ia (Timmes et al. \\cite{TIM03}; Travaglio et al. \\cite{TRA05}; Podsiadlowski et al. \\cite{POD06}). There is also much evidence about the correlation between the properties of SNe Ia and metallicity in observations (Branch \\& Bergh \\cite{BB93}; Hamuy et al. \\cite{HAM96}; Wang et al. \\cite{WAN97}; Cappellaro et al. \\cite{CAP97}), e.g. the maximum luminosity of SNE Ia is proportional to the metallicity (Shanks et al. \\cite{SHA02}).  Two progenitor models of SNe Ia have competed for about three decades. One is a single degenerate model, which is widely accepted (Whelan \\& Iben \\cite{WI73}). In this model, a CO WD increases its mass by accreting hydrogen- or helium-rich matter from its companion, and explodes when its mass approaches the Chandrasekhar mass limit. The companion may be a main-sequence star (WD+MS) or a red-giant star (WD+RG) (Yungelson et al. \\cite{YUN95}; Li et al. \\cite{LI97}; Hachisu et al. \\cite{HAC99a}, \\cite{HAC99b}; Nomoto et al. \\cite{NOM99}; Langer et al. \\cite{LAN00}). Hachisu \\& Kato (\\cite{HK03a}, \\cite{HK03b}) suggested that supersoft X-ray sources, which belong to WD+MS channel, may be good candidates for the progenitors of SNe Ia. The discovery of the companion of Tycho's supernova also verified the reliability of the model (Ruiz-Lapuente et al. \\cite{RUI04}; Ihara et al. \\cite{IHA07}). Another progenitor model of SNe Ia is a double degenerate model (Iben \\& Tutukov \\cite{IBE84}; Webbink \\cite{WEB84}), in which a system consisting of two CO WDs loses orbital angular momentum by gravitational wave radiation and merges. The merger may explode if the total mass of the system exceeds the Chandrasekhar mass limit (see the reviews by Hillebrandt \\& Niemeyer \\cite{HN00} and Leibundgut \\cite{LEI00}). In both of the progenitor models, the CO WD which finally explodes as a SN Ia should approach or exceed the Chandrasekhar mass limit. Obviously, a higher mass CO WD may fulfill this situation more easily than a low mass one. According to our results, the mass of a CO WD with a given initial mass will be higher in the circumstance with extremely high metallicity or extremely low metallicity than that in the middle-metallicity circumstance. Metallicity is therefore very relevant to SNe Ia.        \\subsection{Conclusions}\\label{subs:4.3} We use the method of Han et al. (\\cite{HAN94}) to calculate the IFMR for stars of different metallicities. The conclusions are as follows. \\begin{enumerate} \\item There is an obvious dependence of the IFMRs on the metallicity and the dependence is not monotone. When $Z=0.04$, the final mass of the CO WD for a given initial mass is smallest. For higher or lower $Z$, the mass of CO WD will be higher. The difference of the final mass derived from different metallicties is up to $0.4 M_{\\odot}$.   \\item The initial-final mass relationship for stars of $Z=0.02$ is consistent with observations.   \\item For $Z\\geq0.02$, a helium white dwarf is formed from a star of $M_{\\rm i}\\leq 1.0M_{\\odot}$. The final masses of helium WDs for a given initial mass slightly decreases with metallicity.   \\item Incorporating the IFMR for stars of $Z=0.02$ in the paper into Hurley's single stellar population synthesis code, we reproduce the mass distribution of DA WDs in Sloan DR4 except for some extra-massive WDs.  \\item We reconfirm the discovery of Kalirai et al. (\\cite{KAL05}) that the initial-final mass relationship derived from observation in NGC 2099 might be evidence of the dependence of the IFMR on metallicities and that the metallicity of NGC 2099 may be about $0.01$, although the observational error of white dwarfs in NGC 2099 is large. It should be encouraged to program more accurate observations to find a larger sample of white dwarfs in globular clusters. Such programs may help to confirm the dependence of the initial-final mass relationship on the metallicity.     \\item We bring up again Willson's suggestion (Willson \\cite{WIL00}) that the effect of metallicity may be the origin of the phenomenon that the mean mass of WDs decrease with effective temperature.  \\end{enumerate}"
    },
    "0710/0710.5262_arXiv.txt": {
        "abstract": "I share a few reminiscences and observations of 40 years of Pulsars. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.0392_arXiv.txt": {
        "abstract": "We report 349 radial velocities for 45 metal-poor field red giant and red horizontal branch stars, with time coverage ranging from 1 to 21 years. We have identified one new spectroscopic binary, HD~4306, and one possible such system, HD~184711. We also report 57 radial velocities for 11 of the 91 stars reported on previously by Carney et al.\\ (2003). All but one of the 11 stars had been found to have variable radial velocities. New velocities for the long-period spectroscopic binaries BD$-1$~2582 and HD~108317 have extended the time coverage to 21.7 and 12.5 years, respectively, but in neither case have we yet completed a full orbital period. As was found in the previous study, radial velocity ``jitter\" is present in many of the most luminous stars. Excluding stars showing spectroscopic binary orbital motion, all 7 of the red giants with estimated $M_{\\rm V}$ values more luminous than $-2.0$ display jitter, as well as 3 of the 14 stars with $-2.0 < M_{\\rm V} \\leq\\ -1.4$. We have also measured the line broadening in all the new spectra, using synthetic spectra as templates. Comparison with results from high-resolution and higher signal-to-noise (S/N) spectra employed by other workers shows good agreement down to line broadening levels of 3 \\kms, well below our instrumental resolution of 8.5 \\kms. As the previous work demonstrated, most of the most luminous red giants show significant line broadening, as do many of the red horizontal branch stars, and we discuss briefly possible causes. The line broadening appears related to velocity jitter, in that both appear primarily among the highest luminosity red giants. ",
        "introduction": "In a previous paper (Carney et al.\\ 2003 hereafter C2003), we discussed results from over two thousand high-resolution (R = 32,000), low signal-to-noise (10 to 50 per resolution element) spectra of 91 field metal-poor red giant branch (RGB) and red horizontal branch (RHB) stars. Radial velocities were obtained by cross-correlating each observed spectrum with a synthetic spectrum that closely matched the adopted temperature, gravity, and metallicity for the star. Sixteen stars were found to be single-lined spectroscopic binaries, and orbital solutions were presented for 14 of them. Excluding those 14 stars, observations of the stars covered spans from 2956 days to 6670 days, roughly from 8 to 18 years, with an average of 13.7 years.  The use of synthetic spectra enabled us to measure line broadening as well as radial velocities. We found some anticipated results as well as some surprises. For example, studies of RGB stars in globular clusters (Gunn \\& Griffin 1979; Mayor \\& Mermilliod 1984; Lupton et al.\\ 1987; Pryor et al.\\ 1988; C\\^{o}t\\'{e} et al.\\ 1996, 2002; Brown \\& Wallerstein 1992; Kraft et al.\\ 1997; Mayor et al.\\ 1997) have shown that the most luminous stars, generally with $M_{\\rm V} \\leq\\ -1.4$, show velocity variability that is unlikely to be associated with orbital motion. About half of the stars studied by C2003 with estimated $M_{\\rm V}$ values more luminous than $-1.4$ also showed such velocity variability, which is often referred to as ``jitter\". The term's ambiguity describes the absence of a well-understood cause of the velocity variability.  As expected, the spectroscopic binary frequency, for periods less than 6000 days, was very similar to that found for metal-poor dwarfs and subgiants (Latham et al.\\ 2002; Goldberg et al.\\ 2002). There was some evidence for a dearth of short-period binaries among the RGB and RHB stars, which is not too surprising. Most short period systems are expected to undergo mass transfer when the more massive star begins to expand and fills its Roche lobe. The period does shorten initially, until the donor star's mass becomes less than that of the original secondary star, following which the orbital separation widens and the period increases. Among the fourteen spectroscopic binaries in C2003 with orbital solutions, two systems with small ratios of orbital semi-major axis to estimated primary stellar radius have circular orbits, presumably the result of tidal interactions. These two stars also had higher line broadenings than other RGB stars, presumably due to higher rotational velocities, a consequence of tidally-induced locking of the rotational and orbital periods\\footnote{C2003 referred to line broadening as rotational velocities since the values were determined using rotationally-broadened synthetic spectra. Macroturbulence also contributes to line broadening, so throughout this paper we refer to the more general term, line broadening, and refer to $V_{\\rm broad}$ rather than $V_{\\rm rot}$~sin~$i$.}.  The most unexpected result involved the dependence of the measured line broadening on the stars' evolutionary state. For most of the RGB stars, the line broadening was generally smaller than our instrumental resolution (about 8.5 \\kms). But at the highest luminosities, the mean line broadening rose to values as high as about 12 \\kms\\ at $M_{\\rm V} \\approx\\ -2.6$, and perhaps as high as 15 \\kms\\ if our interpretation of the apparent periodicity in the velocity jitter of two stars was interpreted correctly as due to the rotational period. These very high levels of line broadening seemed to rule out macroturbulence. We also discounted transfer of core rotational angular momentum to the surface of the RGB stars because we should have seen significant increases in line broadening where internal mixing manifests itself by changes in atmospheric abundances. Having ruled out alternative explanations, we speculated that a ``spin up\" in outer envelope rotational velocities could have been caused by the absorption of giant planets. This would imply that metal-poor stars might have such planets {\\em and} that they exist at larger orbital distances from the host stars than have been found to date in many metal-rich disk stars.  We also observed significant line broadening in the RHB stars. To first order, this makes sense. If the much larger stars near the tip of the RGB stars rotate, then so should their descendents. But perhaps their descendents should be stars whose enhanced rotation helps them shed their outer envelopes, which should lead to core helium-burning stars blueward of the instability strip, known as blue horizontal branch (BHB) stars. Soker (1998) and Siess \\& Livio (1999) argued that absorption of a planet by a luminous red giant could explain the high rotation velocities often found among field BHB stars (Kinman et al.\\ 2000; Behr 2003) and cluster BHB stars (Peterson et al.\\ 1983; Peterson 1985a; Peterson et al.\\ 1995; Cohen \\& McCarthy 1997; Behr et al.\\ 2000a,b; Recio-Blanco et al.\\ 2002). More generally, theorists have struggled to explain the presence of both RHB and BHB stars in clusters whose stars share the same metallicity. Some ``second parameter\" must be at work, presumably leading to either larger (RHB) envelope masses or smaller ones (BHB stars). If rotation is the second parameter, then why do RHB stars in the field appear to have rotational velocities comparable to BHB stars, when allowance is made for their different radii?  C2003 suggested four follow-up studies. First, expand the sample to ascertain if our original sample was unusual in some manner. This paper reports on the results of a ``hasty\" study of 45 additional metal-poor field stars\\footnote{By ``hasty\", we mean that some stars were observed for less than two years.}. We also sought to obtain line broadening measures for giants in globular clusters, and initial data for four clusters are in hand. Results will be published later. If planets do exist around metal-poor stars, this would suggest that disk instability is a viable mechanism for planet formation, and a high-precision radial velocity survey of roughly two hundred metal-poor field stars was begun (see Sozzetti et al.\\ 2006). Finally, we have to ask if the line broadening is due to rotation or to macroturbulence. This can be determined with very high-resolution, high-S/N spectra that enable Fourier transform studies of line profiles, following Gray (1982; 1984) and Gray \\& Toner (1986; 1987). We have completed acquisition and analysis of such spectra and will report on the results in a future paper (Carney et al.\\ 2007). ",
        "conclusions": "We have obtained 349 new radial velocities and line broadening measures for 45 metal-poor RGB and RHB stars, as well as 57 such measures for 11 of the stars we studied previously (C2003). A comparison of our derived values for line broadening with results from Behr (2003) and de~Medeiros et al.\\ (2006) shows that our lower-resolution, lower-S/N, and limited wavelength coverage  spectra yield excellent results. We believe the good agreement testifies to the power of high-resolution synthetic spectra as templates. We have identified one new spectroscopic binary, HD~4306, and a possible second one, HD~184711, although we note that the latter's radial velocity variability may be due to velocity jitter rather than orbital motion.  We draw attention to the observed correlation between variable radial velocity (``jitter\") and line broadening. The significant line broadening seen in the metal-poor field RHB stars is hard to explain: it may be a combination of ``spin up\" when a slowly rotating luminous RGB star settles on to the horizontal branch, or temperature-dependent macroturbulence may be involved. Very high-resolution and very high-S/N spectroscopy should reveal the relative balance of rotation and macroturbulence in RGB and RHB stars."
    },
    "0710/0710.1866_arXiv.txt": {
        "abstract": "We study singularities which can form in a spherically symmetric gravitational collapse of a general matter field obeying weak energy condition. We show that no energy can reach an outside observer from a null naked singularity. That means they will not be a serious threat to the Cosmic Censorship Conjecture (CCC). For the timelike naked singularities, where only the central shell gets singular, the redshift is always finite and they can in principle, carry energy to a faraway observer. Hence for proving or disproving  CCC  the study of timelike naked singularities will be more important. Our results are very general and are independent of initial data and the form of the matter. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.1327_arXiv.txt": {
        "abstract": "{Despite the enormous progress occurred in the last 10 years, the Gamma-Ray Bursts (GRB) phenomenon is still far to be fully understood. One of the most important open issues that have still to be settled is the afterglow emission above 10 keV, which is almost completely unexplored. This is due to the lack of sensitive enough detectors operating in this energy band. The only detection, by the \\sax/PDS instrument (15-200 keV), of hard X-ray emission from a GRB (the very bright GRB\\,990123), combined with optical and radio observations, seriously challenged the standard scenario in which the dominant mechanism is synchrotron radiation produced in the shock of a ultra-relativistic fireball with the ISM, showing the need of a substantial revision of present models. In this respect, thanks to its unprecedented sensitivity in the 10--80 keV energy band, \\simbolx{}, through follow--up observations of bright GRBs detected and localized by GRB dedicated experiments that will fly in the $>$2010 time frame, will provide an important breakthrough in the GRB field. ",
        "introduction": "Gamma--Ray Bursts (GRBs) are short and intense flashes of low--energy gamma--rays coming from random directions in the sky at unpredictable times and with a rate of $\\sim$ 300/year as measured by all--sky detectors in low Earth orbit. In the last 10 years, observations allowed huge steps forward in the comprehension of these phenomena, such as their cosmological distance scale, their huge luminosities, their host galaxies, the likely association of \"long\" ($\\sim$2--1000 s) GRBs with the collapse of peculiar massive stars and of \"short\" ($<$$\\sim$2 s) GRBs with the merging of compact objects (NS---NS, NS--BH). See, e.g., \\citet{Meszaros06} for a recent review. However, there are still several open issues, one of the most important being the emission mechanisms in play and their relative contribution to the total radiation. Follow--up observations at longer wavelengths (X--ray, optical, radio) of GRB fields generally lead to the detection of delayed, fading emission (the afterglow). According to the general interpretation, the afterglow emission is described reasonably well, in the framework of the fireball model \\citep{Cavallo78,Meszaros97}, as synchrotron emission from accelerated electrons when a relativistic shell collides with an external medium, the interstellar medium in our case. In this scenario the afterglow spectrum at any given time consists generally of a four segments power law. The spectral and temporal indices are linked together by relationships that depend on the geometry of the fireball expansion \\citep{Sari98} and the properties of the circum--burst environment (density, distribution). The average temporal decaying index,$\\sim$1.3, and spectral photon index, $\\sim$2.2, obtained from observations \\citep{Depasquale06} give an electron spectral index p$\\sim$2.2$-$2.5, which is indeed typical of shock acceleration.  \\begin{figure*}[t!] \\centerline{\\includegraphics[clip=true,width=10cm]{amati_f1.eps} } \\caption{\\footnotesize Simulated \\simbolx{} 15--60 keV image of a bright afterglow (like GRB\\,990123) observed for 100 ks starting from 48 hrs since the GRB onset. } \\label{image} \\end{figure*} ",
        "conclusions": "Despite the enormous observational progress occurred in the last 10 years, the GRB phenomenon is still far to be fully understood. One of the main open issues is the understanding  of physical mechanisms at the basis of prompt and afterglow emission; the case of GRB\\,990123 shows that measurements of the nearly unexplored GRB hard ($>$ 15 keV) X-ray afterglow emission can provide very stringent test to emission models. Thanks to its unprecedented sensitivity in the 15--60 keV energy band, \\simbolx{} can provide a significant step forward in this field. Simulations based on observed distribution of X--ray afterglow fluxes and spectral and decay indices show that even a 100 ks TOO observation of a bright GRB starting 2 days after the event can provide sensitive spectral measurements and allow to discriminate different emission components for a significant fraction of events. Moreover, it is likely that significantly lower TOO stat times (12/24 hr) will be possible for a few event/year, thus allowing sensitive hard X--rays measurements also for medium intensity GRB afterglows. The needed GRB detection and few arcmin localizations will be provided by space missions presently planned to be in flight in the $>$2012 time frame and by optical telescopes."
    },
    "0710/0710.3114_arXiv.txt": {
        "abstract": "The Aarhus code is the result of a long development, starting in 1974, and still ongoing. A novel feature is the integration of the computation of adiabatic oscillations for specified models as part of the code. It offers substantial flexibility in terms of microphysics and has been carefully tested for the computation of solar models. However, considerable development is still required in the treatment of nuclear reactions, diffusion and convective mixing. ",
        "introduction": "What has become ASTEC started its development in Cambridge around 1974. The initial goal was to provide an improved equilibrium model for investigations of solar stability, following earlier work by \\citet{Christ1974}. However, with the initial evidence for solar oscillations and the prospects for helioseismology \\citep{Christ1976} the goals were soon extended to provide models for comparison with the observed frequencies. Given the expected accuracy of these frequencies, and the need to use them to uncover subtle features of the solar interior, more emphasis was placed on numerical accuracy than was perhaps common at the time. The code drew some inspiration from the Eggleton stellar evolution code \\citep{Egglet1971}, which had been used previously, but the development was fully independent of that code. An early description of the code was given by % \\citet{Christ1982}, with further extensive details provided by \\citet{Christ1978}; many aspects of this still stand and will only be summarized here. With the increasing quality and extent of the helioseismic data the code was further developed, to allow for more realistic physics; a major improvement was the inclusion of diffusion and settling \\citep{Christ1993}. This led to the so-called Model~S of the Sun \\citep{Christ1996} which has found extensive use as reference for helioseismic investigations and which at the time provided reasonably up-to-date representations of the physics of the solar interior. In parallel, extensions have been made to the code to consider the evolution of stars other than the Sun; these include the treatment of convective cores and core overshoot, attempts to model red-giant evolution and the inclusion of helium burning. This development is still very much under way.  For use in asteroseismic fitting a version of the code has also been developed in the form of a subroutine with a reasonably simple calling structure which also includes the computation of adiabatic oscillation frequencies as part of the computation. The combined package is available as a single tar file, making the installation relatively straightforward, and the code has been successfully implemented on a variety of platforms. Nevertheless, it is sufficiently complex that a general release is probably not advisable. ",
        "conclusions": "No stellar evolution code is probably ever finished or fully tested. ASTEC has certainly proved useful in a number of applications, and the results for the Sun, as applied to helioseismology, are perhaps reasonably reliable, at least within the framework of `standard solar modelling'. It is obvious, however, that application to the increasingly accurate and detailed asteroseismic data will require further development. The tests provided through the ESTA collaboration and extended through the HELAS Coordination Action are certainly most valuable in this regard."
    },
    "0710/0710.5732_arXiv.txt": {
        "abstract": "{} {We performed an integral field spectroscopic study for the HII galaxy IIZw70 to investigate the interplay between its ionized interstellar medium (ISM) and the massive star formation.} {Observations were taken in the optical spectral range from $\\lambda$3700~\\AA~-~6800~\\AA~with the Potsdam Multi-Aperture Spectrophotometer (PMAS) attached to the 3.5 m telescope at Calar Alto Observatory. We created and analyzed maps of spatially distributed emission-lines (at different stages of excitation), continuum emission, and properties of the ionized ISM (ionization structure indicators, physical-chemical conditions, dust extinction, kinematics). We investigated the relation of these properties to the spatial distribution and evolutionary stage of the massive stars.} {For the first time we have detected Wolf-Rayet (WR) stars in this galaxy. The peak of the ionized gas emission coincides with both the location of the maximum of the stellar continuum emission and the WR bump. The region of the galaxy with lower dust extinction corresponds to the region that shows the lowest values of velocity dispersion and radial velocity. The overall picture suggests that the ISM of this region is being disrupted via photoionization and stellar winds, leading to a spatial decoupling between gas+stars and dust clouds. The bulk of dust appears to be located at the boundaries of the region occupied by the probable ionizing cluster. We also found that this region is associated to the nebular emission in HeII$\\lambda$4686 and to the intensity maximum of most emission lines. This indicates that the hard ionizing radiation responsible for the HeII$\\lambda$4686 nebular emission can be related to the youngest stars. Within $\\sim$ 0.4 $\\times$ 0.3 kpc$^{2}$ in the central burst, we derived oxygen abundances using direct determinations of $T_{\\rm e}$[OIII]. We found abundances in the range 12+log(O/H)= 7.65 - 8.05, yielding an error-weighted mean of 12+log(O/H)=7.86 $\\pm$ 0.05 that has been taken as the representative oxygen abundance for IIZw70.} {} ",
        "introduction": "\\label{intro}  The HII galaxies are gas-rich, dwarf systems that have experienced intense recent or ongoing violent star formation (SF). These objects were identified for the first time by Haro (1956) and Zwicky (1964), and are characterized by the presence of giant HII regions that dominate their observable properties at optical wavelengths. In particular, their integrated spectra are very similar to those of giant extragalactic HII regions in spiral galaxies (Sargent \\& Searle 1970).  The analysis of their emission line spectrum indicates that HII galaxies are metal poor objects (1/40 Z$_{\\odot}$ $\\lesssim$ Z $\\lesssim$ 1/3 Z$_{\\odot}$; e.g. Kunth \\& Sargent 1983).   HII galaxies are ideal laboratories for probing the interplay between massive SF and the ISM in low metallicity environments.  The massive SF process in dwarf galaxies has effects on the properties of the surrounding ISM. A large number of massive stars (between 10$^{4}$ - 10$^{6}$ solar masses of gas are transformed into stars in HII galaxies) are formed almost simultaneously within relatively small volumes (e.g. Melnick 1992). Millions of years after the onset of the burst, the most massive stars begin to explode as supernovae, creating violent, short-lived injections of kinetic energy and metallic elements into the ISM. The disruption of the ISM may significantly affect the spatial distribution of gas and dust particles in the regions close to the massive star cluster, and also determine the way in which new metals ejected by massive stars are mixed with the original gas from which the stars formed (Kunth \\& \\\"Ostlin 2000).  High spatial resolution imaging has revealed that in many HII galaxies the ionized material presents a complex structure: star clusters in the main body (with a non-uniform distribution of SF knots, ensembles of star clusters, or individual super stellar clusters) of the galaxy and external gas components outside the young stellar clusters (e.g. Martin \\& Kennicutt 1995, V\\'{\\i}lchez \\& Iglesias-P\\'aramo 1998, Papaderos et al. 2002). Thus, a two-dimensional analysis of the ionized material in HII galaxies yields spatially resolved information on properties of the ionized gas.  The study of the distribution of these properties is an important issue for our understanding of the interplay between the massive stellar population and the ISM.  However, up to now only a few two-dimensional spectroscopic studies of HII galaxies have been performed. Recent work by Izotov et al. (2006) presents two-dimensional spectroscopy of the extremely metal-deficient HII galaxy SBS 0335-052E. They found a small gradient of the electron temperature $T_e$ and oxygen abundance, and the presence of an ionized gas outflow in the perpendicular direction to the galaxy disk. Cair\\'os et al. (2002), using two-dimensional spectroscopy, obtained the ionized gas velocity field in the central part of Mrk 370.   In this paper, we present a two-dimensional spectroscopic study for the HII galaxy IIZw70. This galaxy is number 70 in Zwicky's second list of ``Compact Galaxies and Compact Parts of Galaxies, Eruptive and Post-eruptive Galaxies''(Zwicky \\& Zwicky 1971).  IIZw70 has been classified as an HII galaxy (Sargent 1970). The basic data of IIZw70 are shown in table~\\ref{tab1}. This object is interacting with its nearby companion IIZw71 at a projected distance of 23 kpc (assuming a distance to the IIZw70-71 system of 18.1 Mpc; see table 1). The ongoing interaction is indicated by interferometric \\ion{H}{I} studies (Balkowski et al. 1978, Cox et al. 2001) which revealed a gaseous streamer connecting IIZw70 with the polar ring galaxy candidate IIZw71 (Whitmore et al. 1990, Reshetnikov \\& Combes 1994). Cair\\'os et al. (2001) show that IIZw70 presents very elongated outer isophotes and a very blue (U-B = -0.89) nuclear starburst in what seems to be an edge-on disk. They also find that the starburst activity is concentrated in the optical center of the galaxy and is surrounded by faint gaseous emission with quite a distorted morphology.  \\begin{figure} \\center{ \\includegraphics[width=9cm,clip]{7987fig1.ps}} \\caption{Image of IIZW70 in the {\\it r (left)} and {\\it u (right)} bands from Sloan Digital Sky Survey. The blue box  mosaic represents the field used in this work ({\\it left}), which is located approximately in the central parts of the galaxy; the green box marks the likely ionizing clusters location, as indicated with a box on the {\\it u}  image ({\\it right}). North is up, east is left. The spatial scale is shown in each panel, where 1 kpc $\\sim$ 11\\hbox{$^{\\prime\\prime }$} assuming a distance to the galaxy of 18.1 Mpc (see table 1).} \\label{SDSS} \\end{figure}  In figure~\\ref{SDSS} we show the images of IIZw70 in the {\\it r} (left image) and {\\it u} (right image) bands extracted from SDSS\\footnote{The Sloan Digital Sky Survey website is http://www.sdss.org}. Over the {\\it r} filter image, we marked the field of view used in this work (blue box mosaic).  In this work, we investigate the relation between massive SF process and the ionized ISM in IIZw70 using integral field spectroscopy (IFS). In order to do this we created and analyzed maps of spatially distributed emission-lines, continuum emission and physical-chemical properties of the ionized gas (electron temperature and density, gaseous metal abundances, dust extinction, excitation). We present a discussion on the significance of the chemical abundance variation found here.  We also present a brief study of the kinematic properties of the ionized gas.  The relationship between the physical-chemical and kinematic properties spatial distribution is discussed.   \\begin{table}% \\caption{Basic data of IIZw70  \\label{tab1}} \\renewcommand{\\footnoterule}{} \\begin{minipage}{\\textwidth} \\begin{tabular}{lc} \\hline Parameter      &Value \\\\ \\hline Name           &IIZw70  \\\\ Other designations                &UGC 9560, Mrk 829    \\\\ R.A. (J2000.0)                 &14h 50m 56.5s         \\\\ DEC. (J2000.0)                &+35d 34' 18''         \\\\ redshift   &0.004         \\\\ M$_{B}$\\footnote{Absolute magnitude in B from Cair\\'os (2001)}(mag)    &-16.3         \\\\ m$_{B}$\\footnote{Apparent magnitude in B from Deeg et al. (1997)}(mag)  &15.1         \\\\ D\\footnote{Distance to the galaxy from Cox et al. (2001)}(Mpc)       &18.1         \\\\ M$_{HI}$\\footnote{Neutral hydrogen mass from Thuan \\& Martin (1981)}(\\msun) &0.34 x 10$^{9}$          \\\\ Z\\footnote{Metallicity from Kobulnicky \\& Skillman (1996)}/Z$_{\\odot}$                          & 1/5  \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table}  In the following section we describe the observations and data reduction. In Sect.3 the results are presented. We discuss our results in Sect.4. ",
        "conclusions": "\\subsection{Ionized gas and massive stars}  We analyzed integral field spectroscopic observations of the HII galaxy IIZw70 in order to investigate the interplay between the massive stars and the ionized gas. This galaxy is interacting with its nearby companion IIZw71. The contiguous streamer of gas between IIZw70 and IIZw71, as well as the SF activity seen in both, indicate an ongoing interaction and possible gas interchange between these two galaxies (Cox et al. 2001). We studied the spatial distribution of properties of the ionized gas (electron temperature and density, chemical abundances, dust extinction, excitation, kinematics) from different emission-line maps.  These maps reveal a central starburst (Cair\\'os et al. 2001, 2001) surrounded by an extended lower excitation, low surface brightness ionized component. From our integrated flux in H$\\alpha$, corrected by extinction, we derived a star formation rate of SFR(H$\\alpha$) = 0.3 M$_{\\odot}$/yr for IIZw70 following Kennicutt (1998); this value is in agreement with the value we derived using the flux reported in Gil de Paz et al. (2003) from CCD H$\\alpha$ imaging. Kewley et al. (2002), using long-slit spectroscopic mapping, derived a SFR(H$\\alpha$) = 0.19 M$_{\\odot}$/yr\\footnote{We calculated this value converting the SFR(H$\\alpha$) given in Kewley et al. (2002) from the distance they adopted to the distance adopted in this paper (see table 1).}. Different integration areas and C(H$\\beta$) values could be possible reasons for the discrepancy between our SFR(H$\\alpha$) measurement and the value found by Kewley et al. (2002).  For the first time we detected the presence of WR stars in IIZw70. For the blue bump, our measurements of the equivalent width, EW(WRbump) = -2.2 $\\pm$ 0.6 \\AA, and of the ratio 100 x I(WRbump)/I(H$\\beta$) = 4.6 $\\pm$ 1.1 (section 3.3) are in good agreement with the predictions of the evolutionary synthesis models Starburst99 (assuming a Salpeter IMF, M$_{upper}$ = 120 M$_{\\odot}$ and Z=0.004 with high mass loss; Leitherer et al. 1999). The models predict EW(WRbump) values around -2 \\AA~ at $\\sim$ 3.2 Myr after an instantaneous burst. Corresponding predictions for 100 x I(WRbump)/I(H$\\beta$) give values around 3.6 Myr. This fact indicates that the central burst in IIZw70 is very young.  Assuming an age for the burst of 3.2 Myr (as deduced from our measurements), we obtained a ratio of WR/O = 0.056 from Starburst99. The presence of a bump at $\\sim$ 5800 \\AA~ could be an indicator of the existence of WC stars. A broad feature at this wavelength has been marginally detected in the same spectra, and though the measured EW is consistent with the same model predictions, its statistical significance is poor due to the low S/N. Overall, a total number of $\\sim$ 17 WN+WC stars are predicted by the models from our findings.  The nebular emission of HeII$\\lambda$4686 detected in this work is shown to be spatially extended, though it is not coincident with the location of the WR bumps by a few arcsec\\footnote{We should bear in mind, however, that a small contribution of nebular HeII$\\lambda$4686 could have remained unnoticed in the broad feature}. This nebular HeII emission is seen to the southwest of the box encompassing the probable ionizing clusters (see figure 4), suggesting the existence of very hard ionizing radiation which, in principle, could not be necessarily related to the WR stars located towards the northeast corner.  In the overall picture, IIZw70 is dominated by the presence of a young central starburst with a significant SFR. The peak of the H$\\alpha$ emission is located close to the probable youngest stellar clusters (see figure~\\ref{cumulo.WR.HeII}).  These young starburst could be producing strong winds which have been able to develop an ionized gas outflow from which we see the blueshifted component as a feature in the radial velocity map of the galaxy (see figure 10). This outflow is thought to have been responsible for the minimum in optical extinction (slightly displaced from H$\\alpha$ emission peak by $\\sim$2$^{\\prime\\prime }$ to the west) apparent in the H$\\alpha$/H$\\beta$ map, where dust could have been swept up by the wind. The characteristic size ($\\sim$ 150 pc) and velocity ($\\sim$ 60 km/s) of the outflow zone are consistent with previous findings in the literature (see e.g. Roy et al. 1991) for the giant HII complex NGC~2363). We can also see that the area with lower dispersion velocity partially overlaps with the zone of the gas outflow, between the probable ionizing clusters location and the zone of minimum extinction. From the kinematical point of view, this picture appears to be consistent with the velocity dispersion map. The narrowest line profiles are dominated by the bulk of the emission from the central HII regions, as expected (Telles et al. 2001), whereas the rest of the ionized gas, with a lower surface brightness, seems to be suffering the effects of the winds and its kinematics should be associated to the structures generated by the likely interaction of the winds with the ambient medium. This situation is reminiscent of what has been found for other extragalactic HII regions and HII galaxies (e.g. Mu\\~noz-Tu\\~n\\'on et al. 1996, Telles et al. 2001) which host strong starbursts.   \\subsection{Chemical enrichment and abundance variations}  In this section, based on our chemical abundance measurements, we present an analysis on the chemical enrichment of the ionized gas and chemical abundance variations in IIZw70.  The oxygen abundance derived here for the integrated spectrum of IIZw70 is 12 + log(O/H) = 7.83 $\\pm$ 0.04. Previous works reported direct values of oxygen abundance for IIZw70 using long-slit spectra positioned in different places across the brightest regions of IIZw70. Lequeux et al. (1979) quoted a value of 8.07; Kobulnicky \\& Skillmann (1996) derived an oxygen abundance of 8.07 $\\pm$ 0.08 whereas Shi et al. (2005) found 12+log (O/H) = 7.69. In principle, the differences between our O/H measurement and the values found in the literature could be explained by different integration areas (see figure 9, bottom panel; Kewley et al. 2005) and atomic parameters used.  Thus far, in most works on chemical abundances in dwarf galaxies there is no clear evidence of abundance variations. Though selected examples of nitrogen enhancement in dwarf galaxies have been reported in the literature, notably the case of NGC~5253 (e.g. L\\'opez-S\\'anchez et al. 2006 and references therein), this is not the case of O/H. Many works in the literature have searched for a possible oxygen abundance variation within dwarf galaxies. Kobulnicky \\& Skillman (1996) presented optical spectroscopic chemical measurements of the ISM in the irregular galaxy NGC 1569 that reveal no substantial localized chemical self-enrichment (i.e., ``local pollution''; see e.g. Kunth \\& Sargent 1986; Pagel, Terlevich \\& Melnick 1986; Pilyugin 1992). V\\'{\\i}lchez \\& Iglesias-P\\'aramo (1998), using bi-dimensional, long-slit spectroscopy, showed that IZw18 presents a substantially homogeneous chemical composition over the whole galaxy. From long-slit spectra of 16 HII regions in the dwarf irregular NGC 1705, Lee \\& Skillman (2004) reported that there is no significant spatial variation of oxygen abundances for this galaxy. Izotov et al. (2006), from two-dimensional spectroscopy of the HII galaxy SBS0335-052E, found a variation of 0.4 dex in the oxygen abundance. They concluded that these variations may not be statistically significant due to unaccounted error estimates.  Lee et al. (2006) obtained a slightly significant (3.2$\\sigma$) oxygen abundance gradient of -0.16 $\\pm$ 0.05 dex kpc$^{-1}$ within the Local Group dwarf irregular galaxy NGC~6822. However they claim further deep high-quality spectra of nebulae and stars are needed to distinguish clearly between either a zero or a nonzero slope.  In this work we calculated the O/H abundance in 16 positions across IIZw70 directly from the measurement of [OIII]$\\lambda$4363. From these values we found a maximum variation in the derived oxygen abundances of $\\sim$ 0.4 dex, similar to the range reported by Izotov et al. (2006) for SBS0335-052E. In order to study whether this variance is statistically significant for an effective abundance variation we must include all other unaccounted uncertainties (variable seeing, flux calibration, etc), which surely contribute to the total oxygen abundance error but are difficult to estimate. In any case, our observations indicate that, within $\\sim$ $\\pm$ 0.2 dex, the ionized gas in the central burst of IIZw70 appears to be chemically homogeneous in O/H over spatial scales of hundreds of parsecs. In the case of the neon-to-oxygen ratio, Ne/O, the range of values as derived from our point by point abundance analysis, does not show inhomogeneity within the errors; as is also the case of the nitrogen-to-oxygen ratio, N/O (see section 3.4).  All these findings give us the upper limits to any chemical inhomogeneity in the ionized gas in the brightest part of IIZw70. In accordance with these results we can say that we are probably not detecting a ``local pollution'' case in IIZw70, consistent with most results obtained by other works cited previously.  A favourable explanation to this apparent degree of homogeneity in dwarf galaxies is given by the scenario presented in Tenorio-Tagle (1996). According to this scenario, in those places in which one expects that newly produced metals could be returned to the ISM, the ``local pollution'' is not detected in the warm phase of the ISM. Rather the newly synthesized metals are returned into the hot phase of the ISM, and there is a significant time delay before these newly produced metals can be detected in the warm phase of the ISM. Variations of the oxygen abundance in the ISM of HII galaxies could give constraints to the physical mechanisms involved in the recycling processes (e.g. Recchi et al. 2001).  Analysing the predictions of chemical evolution models, we verified that continuous SF models (Moll\\'a \\& D\\'\\i az 2005) predict an enhancement in O/H of up to 0.4 dex, to be produced in spatial scales of $\\sim$ 1 kpc along their disk, over timescales of 8 Gyr. The important point here is that these models predict O/H enhancements that could be achieved keeping the N/O ratio approximately constant throughout the process, and at an absolute value consistent with our observations (see section 3.4) to within 0.1 dex. Further observations with IFUs in 10 m class telescopes and additional measurements of chemical abundances in IIZw70, covering a larger area, should be required in order to draw a firmer conclusion on the statistical contribution of local variations to the chemical homogeneity.  Given the interest in the problem of chemical homogeneity of dwarfs galaxies and its role in their chemical evolution (e.g. Legrand et al. 2001, Recchi et al. 2001) a clearer view of the observational situation is probably warranted."
    },
    "0710/0710.0784_arXiv.txt": {
        "abstract": "Spectroscopic modeling of Type II supernovae (SNe) generally assumes steady-state. Following the recent suggestion of Utrobin \\& Chugai, but using the 1D non-LTE line-blanketed model atmosphere code CMFGEN, we investigate the effects of including time-dependent terms, which are generally neglected, that appear in the statistical and radiative equilibrium equations. We base our discussion on the ejecta properties and the spectroscopic signatures obtained from time-dependent simulations, investigating different ejecta configurations (slow, standard, and fast), and covering their evolution from one day to six weeks after shock breakout. Compared to equivalent steady-state models, our time-dependent models produce SN ejecta that are systematically over-ionized, affecting helium at one week after explosion, but ultimately affecting all ions after a few weeks. While the continuum remains essentially unchanged, time-dependence effects on observed spectral lines are large. At the recombination epoch, H{\\sc i} lines and Na{\\sc i}\\,D are considerably stronger and broader than in equivalent steady-state models, while Ca{\\sc ii}\\,8500\\AA\\ is weakened. If time dependence is allowed for, the He{\\sc i} lines at 5875\\AA\\ and 10830\\AA\\ appear $\\sim$3 times stronger at one week, and He{\\sc i}\\,10830\\AA\\ persists as a blue-shifted absorption feature even at 6 weeks after explosion. Time dependence operates through the energy gain from changes in ionization and excitation, and, perhaps more universally across SN types, from the competition between recombination and expansion, which in-turn, can be affected by optical-depth effects. Our time-dependent models compare well with observations of the low-luminosity low-velocity SN\\,1999br and the more standard SN\\,1999em, reproducing the H$\\alpha$ line strength at the recombination epoch, and without the need for setting unphysical requirements on the magnitude of nickel mixing. ",
        "introduction": "Because of radiative cooling and the fast expansion of the exploding mantle of the progenitor massive star, photospheric-phase Type II SN spectra evolve rapidly. At shock breakout the spectral energy distribution (SED) peaks in the far-UV and X-rays. Subsequently the SED centers in the UV, then in the optical, and later in the infrared until the object fades from view. From a radiative transfer perspective, the cooling induces a recombination to lower ionization stages that impose a strong blanketing effect on the energy distribution. A few days after explosion, the ejecta radiate a UV-dominated SED, with He{\\sc ii}\\,4686\\AA\\ (only observed at a few days after shock breakout; Dessart et al. 2007), Balmer/Paschen lines of Hydrogen, He{\\sc i}\\,5875\\AA, He{\\sc i}\\,10830\\AA, some multiplets of O{\\sc ii} and N{\\sc ii} multiplets at \\,4600\\AA\\ (see Dessart \\& Hillier 2006a; Baron et al. 2007), N{\\sc ii}\\,5400\\AA, and a few isolated resonance lines such as Mg{\\sc ii}\\,2802\\AA\\ and Al{\\sc iii}\\,1859\\AA. As the ejecta continue to cool, hydrogen eventually recombines, with a contemporaneous enhancement in line-blanketing due to the switch from Fe{\\sc iii} to Fe{\\sc ii}. In Dessart \\& Hillier (2006) and Dessart et al. (2007), we successfully modeled these epochs for SNe 1999em, 2005cs, and 2006bp, using steady-state non-LTE CMFGEN models and a power-law density distribution with an exponent of ten (the immediate post-breakout phase requires higher values for this exponent). These spectroscopic analyses deliberately focused on the first $\\sim$45 days after shock breakout, since, beyond that time, we encountered severe difficulties in reproducing the hydrogen lines. Specifically, at late times, synthetic line profiles were sizably narrower than observed, suggesting that line formation regions predicted by CMFGEN were confined to velocities/radii that were too small. Despite a continued success with reproducing the overall continuum energy distribution, CMFGEN failed to reproduce important line profiles, whenever hydrogen started recombining at and above the photosphere.  The H$\\alpha$ problem has not been clearly emphasized in the literature, where one can see a great disparity in model atmosphere assumptions and agreement between theoretical predictions and observations, but no direct link to physical/numerical issues. In the eighties and early nineties, the recognized importance of non-LTE effects confronted the strong limitations of computer technology, so that only a few species were treated in non-LTE (i.e., usually hydrogen and helium), while the metals responsible for line blanketing were treated in LTE. \\citet{Eastman_Kirshner_1989} followed this approach to model the first ten days of SN 1987A, and did not encounter obvious difficulties with hydrogen lines. \\citet{Schmutz_etal_1990}, using an approximate non-LTE technique, had, on the contrary, great difficulty reproducing any of the Balmer lines, suggesting clumping as the culprit. \\citet{Hoeflich_1988} reproduced the SN 1987A spectral evolution and the hydrogen lines, over many months, using a large ``turbulent'' velocity. In the more sophisticated non-LTE CMFGEN \\citep{HM_98,DH_05a} models presented in \\citet{DH_06a}, we found that decreasing the turbulent velocity weakened line-blanketing effects and increased, although only modestly, the strength of hydrogen lines. However, beyond 40 days after explosion, this tuning had no longer any important influence on the hydrogen lines.  The non-LTE model atmosphere code PHOENIX predicts strong Balmer lines at the hydrogen recombination epoch \\citep{Mitchell_etal_2001,Baron_etal_2003}, but this may stem from their adoption of non-thermal ionization/excitation due to $^{56}$Ni at the photosphere, sometimes just a few days after explosion and in mass shells moving at $\\ge$10000\\,\\kms in SN 1987A \\citep{Mitchell_etal_2001} or $\\sim$9000\\,\\kms in SN 1993W \\citep{Baron_etal_2003}. Hydrodynamical simulations of core-collapse SNe predict that $^{56}$Ni has velocities of at most $\\sim$4000\\,\\kms, the nickel fingers being strongly decelerated at the H/He interface \\citep{Fryxell_etal_1991,Kifonidis_etal_2000, Kifonidis_etal_2003}. The magnitude of this disagreement extends far beyond the uncertainties of explosion models and suggests a genuine incompatibility.  Interestingly, H{\\sc i} lines remain strong for months in all Type II SNe, in objects as diverse as the ``peculiar'' SN 1987A, the ``plateau'' SN 1999em, and the ``low-luminosity'' SN 1999br. CMFGEN models computed for these objects were unable to reproduce Balmer lines after $\\sim$4 days in SN 1987A, $\\sim$40 days in SN 1999em, and $\\sim$20 days in SN 1999br, all coincident with hydrogen recombination in the ejecta. These three SNe have very different inferred ejecta properties and observed light curves. SN 1999br even synthesized an order of magnitude less $^{56}$Ni than average \\citep{Pastorello_etal_2004} for the Type II class. The only common property between these objects, which is connected to the H$\\alpha$ problem, is the recombination of the ejecta to a lower ionization state at the corresponding epoch.  Recently, \\citet[UC05]{UC_05} proposed that the effect of time-dependence, and the energy associated with changes in ionization/excitation, lead to a strong H$\\alpha$ line profile in SN 1987A during the recombination epoch. They also found that barium lines were affected, and that, with their more consistent approach, Ba{\\sc ii}\\,6142\\AA\\ could be fitted using the LMC metallicity value. The steady-state models of \\citet{Mazzali_etal_1992} supported instead an abundance enhancement of five.  Time dependence has been invoked in the past by \\citet{Fransson_Kozma_1993} to explain the late-time light curve of SN 1987A, and the theoretical study of \\citet{Pinto_Eastman_2000a, Pinto_Eastman_2000b} showed that time dependence {\\it in the radiation field} had a critical impact on the radiative transfer in Type Ias. \\citet{Pinto_Eastman_2000a, Pinto_Eastman_2000b}, however, treat the material in LTE, i.e., do not solve the rate equations, focusing instead on the time-dependent {\\it diffusion} of photons through an optically thick Type Ia SN ejecta. Similarly, \\citet{Kasen_etal_2006} neglect explicit time dependence in the rate equations, computing the ionization and excitation state of the medium in LTE. Thus, although there is at present growing interest in accounting for time dependence in the radiation field, the often-used expedient of LTE, to maintain low CPU costs, has forced the neglect of both non-LTE and time-dependence in the level populations. One exception to these time-dependent LTE approaches is the work of \\citet{Hoeflich_2003}, who treats time-dependence in the rate equations but, to our knowledge, has not discussed the associated effects on Type II SN spectra. In the present study, we investigate thoroughly the effects of time-dependence in the rate equations and their impact on inferred ejecta properties. Note that time dependence in the radiation field is accounted for by adjusting the base luminosity so that the emergent synthetic flux matches the observed flux, using the bolometric-light evolution of SN 1999em as a guide \\citep{DH_06a}.  Here, we report the salient features of several {\\it time-dependent} non-LTE CMFGEN simulations, covering the early evolution for a range of Type II SN ejecta. We confirm the results of UC05 that time-dependence induces an over-ionization of the recombining ejecta and that it solves the H$\\alpha$ problem. However unlike UC05, we self-consistently solve the radiation transfer equation --- the coupling between the level populations and the radiation field is calculated and fully allowed for. Our more comprehensive study of the full spectrum further predicts that all lines, not just those of hydrogen, are significantly affected. The ionization of the ejecta and its evolution are so strongly modified that we predict even He{\\sc i}\\,10830\\AA\\ many weeks after explosion.  In the next section, we present out treatment of time-dependence in CMFGEN, which focuses here on the terms appearing in the statistical and radiative equilibrium equations (an appendix also provides further details). In \\S2.5, we present the various model calculations performed to illustrate our discussion. In \\S3, we present our results, discussing the effects of time dependence on the ejecta properties as well as the associated spectroscopic signatures. We then present in \\S4 a comparison with a few representative observations, focusing on the well observed Type II Plateau SN 1999em and the low luminosity Type II SN 1999br. In \\S5, we discuss the implications of such time-dependent effects on our understanding and on our modeling of photospheric-phase Type II SN spectra, before giving our conclusions in \\S6. Note that the salient features and key results presented here are also available in a concise, ``letter'', format in \\citet{DH_06b}, to which we refer the hurried reader. ",
        "conclusions": "We have presented an extension to the non-LTE model atmosphere code CMFGEN to treat time-dependent terms in the statistical and radiative equilibrium equations. We use implicit first order differencing, and the resulting equations are solved by a partial linearization technique in a manner virtually identical to that used to solve the steady state equations \\citep{HM_98}. We confirm the findings of Utrobin \\& Chugai (2005), which were related to H$\\alpha$ and Ba{\\sc ii}\\,6142\\AA\\  at a few days after explosion in the spectrum of SN 1987A, that the time dependent terms are important for the analysis of Type II SN spectra.   We find that the inclusion of time-dependence terms in the statistical equilibrium equations produces ejecta, in Type II SN, that are systematically over-ionized relative to steady-state models. Spectroscopically, the associated changes in level populations and optical depths alter the line profiles of all species throughout the spectral range. Qualitatively, lines appear in general stronger and broader, as is seen for the Balmer, the Paschen, and the Brackett series of hydrogen or for Na{\\sc i}\\,D. Exceptions (e.g. Ca{\\sc ii}\\,8500\\AA) arise when the altered ionization inhibits the necessary recombination of a given ion (e.g., Ca$^{++}$). Because of optical depth effects, the inclusion of time-dependent terms in the statistical equilibrium equations can even influence the ionization state of the gas in regions where the classic recombination times-scale is much shorter than the flow time-scale. Surprisingly, we predict the presence of He{\\sc i}\\,10830\\AA\\ 6 weeks after the explosion, naturally caused by the frozen, and full, ionization of the outer ejecta. The resulting detached blueshifted absorption and flat-topped emission profile is supported by observations, although the magnitude of the absorption component may require the additional contribution from the cold dense shell that can form at the interface between the SN ejecta and the pre-SN wind \\citep{Chugai_etal_2007}.  The effects of time dependence observed stem from the energy gain that follows changes in ionization and excitation, and from the similarity in recombination and expansion time-scales (crucial at large distances above the photosphere). The recombination time-scale can be lengthened by optical-depth effects, and in some cases this may allow time-dependence effects to be important even at the photosphere. The present study indicates that neglecting time dependence may compromise the results of quantitative analyses of Type II SN spectra, affecting the inferred ejecta ionization (and its origin) or chemical abundances. In particular, time dependence offers a natural means to sustain strong and broad lines at late times, as observed in {\\it all} Type II SN spectra, without invoking radioactive contributions. Such non-thermal excitation/ionization is expected to be relevant at the photosphere after a few months, but not as soon as a few weeks, and is subject to strong variations accross the SN class (SNe 1999br or 2005cs, which boast a strong H$\\alpha$ at the recombination epoch, have symptomatically low $^{56}$Ni yields compared to the standard Type II-Plateau SN 1999em).  The time-dependence effects we observe just one week after explosion in the outer, fully ionized, SN ejecta also support the idea that the energy gain drawn out of recombining ions is not a fundamental driver. Time-dependence effects can lead to over-ionization for He, C, N, O, and metals, under fully-ionized hydrogen conditions. We thus surmise that time dependence should operate, with a magnitude to be determined, in SN ejecta of all types, primarily because they all combine the properties of fast expansion and low density.  We have investigated the potential impact of time dependence on the correction factors used in the Expanding Photosphere Method, or, essentially, whether the influence on the electron density generates a systematic shift in the magnitude of flux dilution. We find no sizable and systematic deviation from the correction factors obtained, at the same color temperature, with steady-state CMFGEN models. This supports our use of steady-state CMFGEN models for distance determinations based on early-time Type II SN observations (Dessart \\& Hillier 2006a; Dessart et al. 2007). However, the good agreement at late times between time-dependent CMFGEN models and late time photospheric-phase observations motivates the extension of the time baseline to include the recombination epoch, thus allowing for multi-epoch observations that cover the entire Plateau phase, thereby reducing the errors on the inferred distance.  For the modeling of the longer evolution of Type II SN ejecta and their radiation, a time-dependent approach is warranted. To achieve a higher level of consistency, CMFGEN is under developement to follow as well the time-dependent evolution of the radiation field, together with options for chemical stratification and energy deposition from radioactive isotopes. This versatility will allow us to study SN ejecta of any type and over months after explosion."
    },
    "0710/0710.2262_arXiv.txt": {
        "abstract": "We have developed a method for measuring higher-order weak lensing distortions of faint background galaxies, namely the weak gravitational flexion, by fully extending the Kaiser, Squires \\& Broadhurst method to include higher-order lensing image characteristics (HOLICs) introduced by Okura, Umetsu, \\& Futamase. We take into account explicitly the weight function in calculations of noisy shape moments and the effect of higher-order PSF anisotropy, as well as isotropic PSF smearing. Our HOLICs formalism allows accurate measurements of flexion from practical observational data in the presence of non-circular, anisotropic PSF. We test our method using mock observations of simulated galaxy images and actual, ground-based Subaru observations of the massive galaxy cluster  A1689 ($z=0.183$).  From the high-precision measurements of spin-1 first flexion, we obtain a high-resolution mass map in the central region of A1689. The reconstructed mass map shows a bimodal feature in the central $4'\\times 4'$ region of the cluster. The major, pronounced peak is associated with the brightest cluster galaxy and central cluster members, while the secondary mass peak is associated with a local concentration of bright galaxies. The refined, high-resolution mass map of A1689 demonstrates the power of the generalized weak lensing analysis techniques for quantitative and accurate measurements of the weak gravitational lensing signal. ",
        "introduction": "Propagation of light rays from a distant source to the observer is governed by the gravitational field of intervening mass fluctuations as well as by the global geometry of the universe. The images of background sources hence carry the imprint of the gravitational potential of intervening cosmic structures, and their statistical properties can be used to test the background cosmological models.   Weak gravitational lensing is responsible for the weak shape-distortion and magnification of the images of background sources due to the gravitational field of intervening matter (e.g., Bartelmann \\& Schneider 2001; Umetsu, Futamase, \\& Tada 1999). To the first order, weak lensing gives rise to a few -- $10\\%$ levels of elliptical distortions in images of background sources, responsible for the second-order derivatives of the gravitational lensing potential. Thus, the weak lensing signal, measured from tiny but coherent quadrupole distortions in galaxy shapes, can provide a direct measure of the projected mass distribution of cosmic structures. However, practical weak lensing observations subject to the effects of atmospheric seeing, isotropic/anisotropic PSF, and (residual) camera distortion across the field of view, which must be examined from the stellar shape measurements and corrected for in the weak lensing analysis. Practical methods for PSF corrections and shear measurements/calibrations have been studied and developed by many authors, such as pioneering work of  Kaiser, Squires, \\& Broadhurst (1995, hereafter KSB), the Shapelets technique which describes PSF and object images in terms of Gaussian-Hermite expansions (Refregier 2003), and recent systematic, collaborative efforts by The Shear TEsting Programme  (Heymans et al. 2006; Massey et al. 2007).  Thanks to these successful developments in weak lensing techniques as well as in instrument technology, the quadrupole weak lensing has become one of the most important tools in observational cosmology to map the mass distribution in individual clusters of galaxies (e.g., Kaiser \\& Squires 1993; Broadhurst et al. 2005a; Okabe \\& Umetsu 2007; Umetsu \\& Broadhurst 2007), measure ensemble-averaged mass profiles of galaxy-group sized halos from the galaxy-galaxy lensing signal (e.g., Hoekstra et al. 2001; Hoekstra et al. 2004; Parker et al. 2005; Mandelbaum et al. 2006), study the statistical properties of the large scale structure of the universe from the cosmic shear statistics (e.g., Bacon et al. 2000; van Waerbeke et al. 2001; Hamana et al. 2003), and search for galaxy clusters by their mass properties (e.g., Schneider 1996; Erben et al. 2000; Umetsu \\& Futamase 2000; Wittman et al. 2001).   In recent years, there have been theoretical efforts to include the next higher order distortion effects as well as the usual quadrupole distortion effect in the weak lensing analysis (Goldberg \\& Natarajan 2002; Goldberg \\& Bacon 2005; Bacon et al. 2006; Irwin \\& Shmakova 2006; Goldberg \\& Leonard 2007; Okura, Umetsu, \\& Futamase 2007). We have proposed in Okura, Umetsu, \\& Futamase (2007, hereafter OUF) to use certain convenient combinations of octopole/higher multipole moments of background images which we call the Higher Order Lensing Image's Characteristics (HOLICs), and have shown that HOLICs serve as a direct measure for the next higher-order weak lensing effect, or the gravitational flexion (Goldberg \\& Bacon 2005) and that the use of HOLICs in addition to the quadrupole shape distortions can improve the accuracy and resolution of weak lensing mass reconstructions based on simulated observations.   Recently, Goldberg \\& Leonard (2007) extended the HOLICs approach for flexion measurements to include observational effects, namely the Gaussian weighting in shape-moment calculations (see the appendix therein) and the isotropic PSF effect, under the assumption that PSF is nearly circular, and tested their extended HOLICs approach with simulated and HST/ACS observations. Leonard et al. (2007) have applied the extended HOLICs method to reconstruct the projected mass distribution in the central region of the massive galaxy cluster A1689 at $z=0.183$, and revealed substructures associated with small clumps of galaxies. Further Leonard et al. (2007) found that in dense systems such as galaxy clusters the HOLICs technique is robust and less sensitive than the Shapelet technique to contamination by light from the extended wings of lens/foreground galaxies.   In the present paper we develop a method for measuring flexion by the HOLICs approach by fully extending the KSB formalism; We take into account explicitly the effects of Gaussian weighting in calculation of noisy shape moments and higher-order PSF anisotropy as well as isotropic PSF smearing. We then apply our method to actual, ground-based Subaru observations of A1689, and perform a mass reconstruction in the central region of A1689.  The paper is organized as follows. We first summarize in \\S 2 the basis of weak gravitational lensing and the flexion formalism. In \\S 3, we derive the relationship between HOLICs and flexion by incorporating Gaussian smoothing in shape measurements in the presence of isotropic and anisotropic PSF. The practical method to correct the isotropic/anisotropic PSF effects will be presented in \\S 4. In Section 5 we use simulations to test our flexion analysis method based on the HOLICs moment approach. We then perform a weak lensing flexion analysis of A1689 by our fully-extended HOLICs approach, and perform a mass reconstruction of A1689 from the HOLICs estimates of flexion. Finally summary and discussions are given in \\S 6. We refer interested readers to a complete appendix\\footnote {Full appendix is available in electronic form at http://www.asiaa.sinica.edu.tw/keiichi/OUF2/appendix.pdf.} for details of the derivation of flexion-observable relationships in practical observations. ",
        "conclusions": "In the present paper, we have developed a method for weak lensing flexion analysis by fully extending the KSB method to include the measurement of HOLICs (OUF). In particular, we take into account explicitly the weight function in calculations of noisy shape moments and the effects of spin-1 and spin-3 PSF anisotropies, as well as isotropic PSF smearing, in the limit of weak lensing and small  PSF anisotropy ($q$). The higher order weak lensing effect induces a centroid shift in the observed image of the background (Goldberg \\& Bacon 2005; OUF; Goldberg \\& Leonard 2007). In weighted moment calculations, this will yield in the flexion measurement additional correction terms (relevant to $W'(x), W''(x), W'''(x)$) that must be taken into account by properly expanding the weight function $W(x)$. It is found that neglecting these additional terms originated from the Taylor expansion of $W(x)$ yields the same result as obtained by Goldberg \\& Leonard (2007; see Appendix therein). We extended the KSB formalism to include the higher-order isotropic and anisotropic PSF effects relevant to spin-1 and spin-3 HOLICs by following the prescription given by KSB and Bartelmann \\& Schneider (2001), which provides direct relations between the observable HOLICs and underlying flexion in the weak lensing limit.  We have implemented in our analysis pipeline our flexion analysis algorithm based on the HOLICs moment approach, and tested the reliability and limitation of our PSF correction scheme using numerical simulations. Our simulation results show that (i) after applying our PSF correction method the PSF-induced anisotropies in HOLICs of mock galaxy images can be considerably reduced by a factor of $10$--$100$, depending on the strength of PSF anisotropy, (ii) those small galaxies whose angular size is smaller than or comparable to the size of PSF suffer from severe anisotropic PSF effects, and that (iii) there is an overall trend that the fractional correction factor is larger for larger galaxy images. Therefore, our simulation results support the reliability of our PSF-correction scheme and its practical implementation.   Based on the simulation results, we have applied our flexion analysis pipeline to ground-based $i'$ imaging data of the rich cluster A1689 ($z=0.183$) taken with Subaru/Suprime-Cam. Our flexion analysis of Subaru A1689 data revealed a non-negligible, significant effect of higher-order PSF anisotropy induced in stellar images (Figures \\ref{fig:zetaq_field} and \\ref{fig:zetaq}). It is therefore important in practical flexion measurements to quantify and correct for the higher-order anisotropic PSF effects.  Our mass reconstruction from the first-flexion measurements shows two significant ($>4\\sigma$) mass structures associated with concentrations of bright galaxies in the central cluster region: the first peak ($5.2\\sigma$) associated with the central concentration of bright galaxies including the cD galaxy, and the second peak ($4.4\\sigma$) associated with a clump of bright galaxies located $\\sim 1'$ northeast of the cluster center. This significant detection of the second peak confirms earlier ACS results from the strong lensing analysis (Broadhurst et al. 2005b; Halkola et al. 2006; Leonard et al. 2007) and the combined strong lensing, weak shear, and flexion analysis by Leonard et al. (2007). The central mass peak, however, was not recovered in the earlier flexion analysis by Leonard et al. (2007) based on HST/ACS data. Leonard et al. (2007) attributed this to their relatively large reconstruction error at the cluster center, although they have a very large number density of background galaxies, $\\bar{n}_g\\approx 75$ arcmin$^{-2}$. On the other hand, owing to our conservative selection criteria for the background sample, the mean number density of background galaxies used for the present analysis is $\\bar{n}_g=7.75$ arcmin$^{-2}$, which is almost one order of magnitude smaller than that of the ACS data, and is about $20\\%-30\\%$ of a typical number density of magnitude/size-selected background galaxies usable for the quadrupole shape measurements in ground-based Subaru observations ($\\bar n_g\\sim 30-40 {\\rm arcmin}^{-2}$). However, we found that it is rather important to remove small/faint galaxy images and noisy outliers in flexion measurements since they are likely to be affected by the residual PSF anisotropy and/or observational noise in the shape measurement (\\S \\ref{subsec:sim}). Besides, the smaller the object, the larger the amplitude of intrinsic flexion contributions. Recall that flexion and HOLICs have a dimension of length inverse: The response to flexion is size-dependent, and the amplitude of intrinsic flexion is inversely proportional to the object size. Indeed, we find that inclusion of smaller objects results in a noisy reconstruction. Similar values of the background number density, $\\bar n_g\\simlt 10 {\\rm arcmin}^{-2}$,  have been used in recent quadrupole weak lensing analyses based on Subaru observations (e.g., Broadhurst et al. 2005a; Umetsu \\& Broadhurst 2007; weak lensing cluster mass measurements of Okabe \\& Umetsu 2008). In their studies only objects redder than the cluster sequence are selected in color-magnitude space for their weak lensing analysis, because such a red population is expected to comprise only background galaxies ($\\bar z_s\\sim 0.9$; see, e.g., Medezinski et al. 2007), made redder by relatively large $k$-corrections and with negligible contamination by cluster galaxies (Broadhurst et al. 2005a; Medezinski et al. 2007). However, the smaller number of objects implies a coarser angular resolution in the map-making for achieving a proper signal-to-noise ratio (e.g., per-pixel ${\\rm S/N}\\simgt 1$). With $\\bar n_g\\sim 10 {\\rm arcmin}^{-2}$ for cluster quadrupole weak lensing, typical angular resolutions are about $1-2$ arcmin (e.g., Gaussian FWHM, or boxcar width). On the other hand, flexion measures essentially the gradient of the tidal gravitational shear field (i.e., $F,G\\propto \\phi(r)/r^3$), and hence is relatively sensitive to small-scale structures. Therefore, our successful reconstruction of the mass substructures with a small background density, $\\bar n_g\\sim 8 {\\rm arcmin}^{-2}$, could be attributed to the superior sensitivity of flexion to small scale structures (see OUF for detailed discussions) and the here-adopted selection criteria for a background galaxy sample for weak lensing flexion analysis.   Finally, we emphasize that our HOLICs formalism here is different from the earlier work by Goldberg \\& Leonard (2007) in that (1) additional correction terms for the centroid shift, relevant to the derivatives of the weight function, have been included and (2) the spin-1 and spin-3 PSF anisotropies, as well as the isotropic PSF smearing, have been taken into account under the assumption of small PSF anisotropy ($q[\\btheta]$), as done in the KSB formalism. Our flexion-based mass reconstruction of A1689 demonstrates the power of the generalized flexion analysis techniques for quantitative and accurate measurements of the weak gravitational lensing effects."
    },
    "0710/0710.5383_arXiv.txt": {
        "abstract": "We present further \\sirtf\\ Space Telescope observations of the recurrent nova {\\rs}iuchi, obtained over the period 208--430~days after the 2006 eruption. The later \\sirtf\\ IRS data show that the line emission and free-free continuum emission reported earlier is declining, revealing incontrovertible evidence for the presence of silicate emission features at 9.7 and 18\\mic. We conclude that the silicate dust survives the hard radiation impulse and shock blast wave from the eruption. The existence of the extant dust may have significant implications for understanding the propagation of shocks through the red giant wind and likely wind geometry. ",
        "introduction": "\\rs\\ is a recurrent nova (RN) that erupted in 1898, 1933, 1958, 1967, 1985 \\citep{wallerstein}. It consists of a semi-detached binary (orbital period 455.7~days) comprising a roche-lobe-filling red giant (M2III) and a massive ($\\gtsimeq1.2\\Msun$) white dwarf \\citep[WD;][]{fekel}. The eruption follows a thermonuclear runaway on the surface of the WD \\citep{tnr}. In the case of the \\rs\\ class of RN, however, the ejected material runs into, and shocks, a dense red giant (RG) wind \\citep{bk85,obrien92}.  The 1985 eruption was, for the first time, the subject of a multi-wavelength observational campaign, from the radio to the X-ray \\citep{vnu}. Its most recent eruption, on 2006 February 12.83 \\citep[][ we take this to define the origin of time post-outburst]{hirosawa}, was the subject of an even more intensive observational campaign. Infrared (IR) observations \\citep{das06,evans07a,evans07b}  showed evidence for the shock, seen also at radio \\citep{eyres,obrien} and X-ray \\citep{bode06,sokoloski,ness,osborne} wavelengths, as the ejecta interacted with the RG wind.  Observations of \\rs\\ obtained with the Spitzer Space Telescope \\citep[{\\it Spitzer};][]{werner, gehrz07} during the period from 2006 April 16 through 26~UT ($\\simeq 67$ days after outburst) were described by \\cite{evans07b}. We present here further IR observations conducted as part of a long duration synoptic campaign, obtained later in the outburst with the \\sirtf\\ Infrared Spectrometer \\citep[IRS;][]{houck}. ",
        "conclusions": "We have presented further \\sirtf\\ IRS observations of \\rs\\ in the aftermath of its 2006 eruption. Unlike our earlier observation, which showed emission by the hot shocked gas, more recent IR spectra reveal the presence of silicate dust.  The dust we see could not have been formed in the 2006 eruption. It most probably survived the 2006 eruption, because a portion of the RG wind (including the dust) did not see the eruption at all, and because the passing shock failed to destroy any dust that survived the UV blast at the eruption. This conclusion may have major implications for the evolution of the shock, and of the long-term survival of the RG wind between eruptions. We predict that the dust we  see around \\rs\\ will persist until the next eruption."
    },
    "0710/0710.0790_arXiv.txt": {
        "abstract": "{Water vapor emission at 22 GHz from masers associated with star-forming regions is highly variable.}{We present  a database of up to 20  years of monitoring of a sample of 43 masers within star-forming regions.  The sample covers a large range of luminosities of the associated IRAS source and is representative of the entire population of \\hdo\\ masers of this type.  The database forms a good starting point for any further study of \\hdo\\ maser variability.}{The observations were obtained with the Medicina 32--m radiotelescope, at  a rate of 4--5 observations per year.} {To provide a database that can be easily  accessed through the web,  we give for each source: plots of the calibrated spectra, the velocity--time--flux density plot, the light curve of the integrated flux, the lower and upper envelopes of the maser emission, the mean spectrum, and the rate of the maser occurrence as a function of velocity. Figures for just  one source  are given in the text for representative purposes. Figures for all the sources are given in electronic form in the on-line appendix. A discussion of the main properties of the \\hdo\\ variability in our sample will be presented in a forthcoming paper.} {} ",
        "introduction": "Since the discovery of 22 GHz \\hdo\\ maser emission associated with young stellar objects (YSOs) within star-forming regions (SFRs), variability of maser emission is well-known.  Changes as large as several orders of magnitude in the maser emission have been observed (e.g., Little et al.~\\cite{lit77}; Liljestr\\\"om et al.~\\cite{lil89}; Claussen et al.~\\cite{cla96}; Comoretto et al.~\\cite{com90}; Wouterloot et al.~\\cite{wou95}). At the same time,  velocity drifts of individual components of up to a few \\kms\\ per year  have also been reported (e.g.,  Brand et al.~\\cite{BCCFPPV2003}). The variability can be slow or burst-like and covers all ranges of timescales, from hours-days to months-years. In the present study we are mainly concerned with the latter and, for the first time, deal with a large sample of sources (43) and with a time-span of up to 20 years, with 4--5 spectra per year.  Due to the large amount of  telescope time required to follow the evolution of \\hdo\\ maser  emission, inevitably the variability is more easily  monitored through single dish observations. In fact, besides our campaign only the Pushchino 22--m single-dish maser patrol covers a comparably long period (e.g., Rudnitskij et al.~\\cite{russi}). Furuya et al.\\ (\\cite{furuya}) carried out a multiepoch 22 GHz H$_{2}$O maser survey towards 173 low-mass YSOs (Class 0 to Class III sources) using the Nobeyama 45 m telescope. This was the first complete water maser survey towards Class 0 sources in the northern sky. However, their observations extend non-uniformly over a period of only three years.   It is well-known that the variability depends strongly  on the luminosity of the  SFR (usually derived from the associated IRAS source).  For instance, there is a minimum luminosity ($\\sim$25 L$_\\odot$)  below which \\hdo\\ maser emission may be present only for about one third of the entire duration of the maser activity (Persi et al.~\\cite{per94}; Claussen et al.~\\cite{cla96}). The results obtained so far suggest that the lower the SFR luminosity, the higher the observed degree of variability of the \\hdo\\ maser emission. Higher luminosity sources may show more steady components. Consequently, a large  sample of sources is  needed to better understand the dependence of the variability on other parameters of the SFRs, in particular on  their luminosities which cover a range of several orders of magnitude, and on their (molecular) environment.  The use of a single dish instrument is of course a limitation because interferometric observations clearly show that \\hdo\\  masers within a SFR often consist of  many spatially separated, unresolved components, generally clustered in groups (usually with different velocities), which in most cases cannot be separated in single dish observations (e.g., Forster \\& Caswell~\\cite{for89}, ~\\cite{for99}; Tofani et al.~\\cite{tof95}). However, Felli et al.~(\\cite{FMRC2006}) reported a case in which single dish observations were able to separate the evolution of spatially distinct components. Nevertheless, frequent single dish observations have the advantage of being able to follow potential  velocity drifts of individual {\\it velocity} components, which can be easily  identified in the spectra, and thus give a better view of the dynamics (e.g., accelerations) occurring in the circumstellar environment of the YSO.  In two earlier papers (Valdettaro et al.~\\cite{VPBCCFP2002}; Brand et al.~\\cite{BCCFPPV2003}), we set the basis for a systematic study of the \\hdo\\ variability in SFRs with the Medicina 32--m radiotelescope  using a sample of 14 sources which  covered a wide range of luminosities and had been observed  for more than 10 years. Here we report on a larger sample of 43 SFRs (including the former 14 sources)  and increase  the time coverage up to 20 years. The database is presented in the form of plots of the calibrated spectra and additional plots of derived quantities, in a form that can be easily accessed through  the web. In a forthcoming paper, we will analyze  the properties of \\hdo\\ variability of our  sample. ",
        "conclusions": "We present the observational results of a systematic study extending for almost 20 years of the \\hdo\\ maser variability in 43 SFRs with luminosities of the associated IR sources between 20 and 1.5$\\times 10^6$ L$_\\odot$. This database provides the backbone for the discussion of the main long-term properties of maser emission that will be presented in a forthcoming paper. We have identified several ways to describe graphically the main aspects of the \\hdo\\ maser emission. These include:  \\begin{itemize} \\item [(a)] the spectra in a compressed form, with an autoscaled  flux density scales (the same sets of plots, but with fixed linear and logarithmic flux density scales, are also available from our WEB pages); \\item [(b)] the velocity--time--flux density  plots, which  conveniently describe the morphology of the variability of the maser emission. These diagrams are particularly useful for recognizing the presence of possible velocity drifts and separating steady components from bursts of short duration; \\item [(c)] the velocity--integrated flux density as a function of time. This light curve   describes the overall emission  of the maser components associated with a SFR; \\item [(d)] the mean spectrum; \\item [(e)] the upper and lower  envelopes of the maser emission. The upper envelope represents the maximum emission that the SFR could produce if {\\it all} the velocity components were present {\\it simultaneously} and emitting at their {\\it maximum} rate. Similarly, the lower envelope pinpoints the steady components and their lowest level  (but $> 5 \\sigma$) of emission; \\item [(f)] the number of times that  the maser emission  is well above the noise ($> 5 \\sigma$) as a function of velocity is  directly gauged by the histogram of the rate--of--occurrence. \\end{itemize}"
    },
    "0710/0710.0759_arXiv.txt": {
        "abstract": "{  We suggest  to  use the   observationally measured  and  theoretically justified correlation between size and rotational velocity of galactic discs as a  viable method to select a   set of high redshift  standard rods which may  be  used to explore the   dark energy  content of  the universe via the classical angular-diameter  test.  Here we explore  a new strategy for  an optimal implementation  of this test.  We propose to use the rotation speed of high redshift galaxies as a standard size indicator and  show how high  resolution multi-object spectroscopy and ACS/HST high  quality spatial images,  may be combined  to measure the amplitude of the  dark  energy density  parameter  $\\Omega_{Q}$, or to constrain  the cosmic equation  of state  parameter  for a smooth dark energy   component  ($w=p/\\rho, \\;\\; -1\\le  w  <  -1/3$).  Nearly 1300 standard rods with high velocity rotation in the bin $V=200\\pm 20$km/s are expected in a field of 1 sq. degree and over the redshift baseline $0<z<1.4$. This sample is sufficient  to constrain the cosmic equation of state parameter  $w$ at a level   of $20\\%$ (without  priors in the [$\\Omega_m,\\Omega_Q$]  plane)  even  when  the [OII]$\\lambda 3727$\\AA~ linewidth-diameter relationship is  calibrated with a scatter of $\\sim 40\\%$. We evaluate how systematics may  affect the proposed tests, and find that a  linear standard rod  evolution, causing galaxy dimensions to be up to $30\\%$ smaller at $z=1.5$,  can be uniquely diagnosed, and will   minimally   bias  the  confidence     level contours   in   the [$\\Omega_{Q}$, $w$] plane.  Finally, we   show how to derive,  without {\\it a priori} knowing the specific functional form of disc evolution, a cosmology-evolution diagram with which it is possible to establish a mapping between different cosmological models and the amount of galaxy disc/luminosity evolution expected at a given redshift. ",
        "introduction": "Several  and    remarkable progresses  in  the  understanding   of the dynamical status of the universe,  encourage us to believe that, after roaming from paradigm to paradigm, we are finally converging towards a well-founded, internally consistent standard model of the universe.  The  picture emerging from  independent  observations and  analysis is sufficiently coherent  to   be referred to  as   the {\\it concordance} model \\citep[e.g.][]{teg06}. Within this  framework, the universe is  flat  ($\\Omega_K = -0.003^{+0.0095}_{-0.0102}$) composed   of  $\\sim 1/5$ cold dark matter ($\\Omega_{cdm} \\sim 0.197^{+0.016}_{-0.015}$) and $\\sim 3/4$ dark energy  ($\\Omega_{\\Lambda} = 0.761^{+0.017}_{-0.018}$),  with large negative pressure ($w =-0.941^{+0.017}_{-0.018}$), and with a very low baryon content ($\\Omega_b = 0.0416^{+0.0019}_{-0.0018}$). Mounting and compelling evidence for accelerated expansion of the universe, driven by a  dark energy component, presently relies  on our  comprehension of the mechanisms with   which Supernovae Ia (SNIa) emit  radiation (see  \\citet{per99,  rie01})   and  of the    physical processes that  produced   temperature fluctuations  in   the primeval plasma (see \\citet{lee01, deb02, hal02, spe06}.)  Even  if the ambitious task of   determining geometry and evolution of the  universe as a whole, which  commenced in the 1930s, now-day shows that the relativistic Friedman-Lema\\^{\\i}tre model passes impressively demanding  checks, we are faced with  the  challenge of developing and adding new  lines  of evidence  supporting   (or falsifying)  the {\\it concordance} model. Moreover, even if   we parameterize our   ignorance about dark energy describing its nature only via  a simple equation of state $w=p/\\rho$, we only have  loose constraints on the precise value of the w parameter or on its functional behavior.  In this  spirit we  focus  this  analysis on  possible   complementary approaches    to determining   fundamental cosmological    parameters, specifically on geometrical tests.   A whole arsenal of classical geometrical methods has been developed to measure global properties of the universe. The  central feature of all these  tests  is the attempt to  directly  probe the various operative definitions    of  relativistic   distances   by  means   of   scaling relationships in which an  observable  is expressed  as a function  of redshift ($z$) and of the  fraction of critical density contributed by all forms of matter and energy ($\\Omega$).  The most remarkable among these classical  methods are the {\\it Hubble diagram} (or  magnitude-redshift relation $m=m(M,z,\\Omega$)), the {\\it Angular          diameter     test}    (or   angle-redshift   relation $\\theta=\\theta(L,z,\\Omega$)), the {\\it Hubble test} (or count-redshift relation N=N(n,z,$\\Omega$))   or the  {\\it Alcock-Paczinsky  test} (or deformation-redshift relation $\\Delta z/ z \\Delta \\theta\\equiv k= k(z, \\Omega$)). The common key idea is to constrain cosmological parameters by measuring,  at various cosmic epochs, the   scaling of the apparent values $m$, $\\theta$, N, $k$ of some  reference standard in luminosity (M),  size  (L),  density (n)   or   sphericity and  compare   them to corresponding model predictions.  The  observational viability of these  theoretical strategies has been remarkably proved by the Supernova Cosmology Project \\citep{per99} and the  High-z Supernova Team  \\citep{rie01} in  the case  of  the Hubble diagram. With a parallel strategy,  \\citet{new02} recently showed that a variant of the Hubble test ($N(z)$ test) can be in principle applied to distant  optical clusters selected in  deep redshift survey such as VVDS \\citep{lef05} and DEEP2  \\citep{dav00},  in order to measure  the cosmic equation-of-state parameter $w$.  Unfortunately,  the   conceptually simple pure   geometrical tests of world models,  devised     to anchor relativistic cosmology  to    an observational basis,    have so  far   proved   to   be difficult   to implement. This   is because the  most effective  way to constrain the evolution of the cosmological metric  consists in probing deep regions of   the universe      with a  primordial    class   of   cosmological objects. Besides   the complex instrumental    technology this kind of experiments  requires,   it becomes  difficult  at  high   redshift to disentangle the effects  of   object evolution from the  signature  of geometric evolution.   Since geometrical tests are by definition independent from predictions of   theoretical models or  simulations,  as well  as from assumptions about content, quality  and   distribution of matter in   the universe (mass  fluctuations  statistics (e.g. Haiman,    Mohr \\& Holder  2001, Newman et al.  2001), galactic  biasing  (e.g. Marinoni et al.   1998, Lahav et al.  2001,  Marinoni et  al.  2006), Halo   occupation models (e.g. Berlind et al. 2001,  Marinoni \\& Hudson 2002,  van den Bosch et al.   2006)) it  is  of  paramount  importance  to  try  to  devise an observational way to implement them. The technical maturity of the new generation of   large  telescopes, multi  object  spectrographs, large imaging   detectors and space   based  astronomical observatories will allow these  tests to be  more effectively applied  in the near future \\citep{hut00}. In  this paper, we describe a  method to select a class of  homologous  galaxies  that   are at  the  same  time  standard  in luminosity and size,  that can be in  principle applied to data coming from  the zCOSMOS   spectro-photometric   survey (Lilly et  al.  2006, of the deep universe.   An observable relationship exists between the speed of rotation $V$ of a spiral galaxy and  its metric radial  dimension $D$  as well as  its total   luminosity    L   \\citep{tul77,bot80}.    From   a theoretical perspective, this set of scaling relations are expected and explicitly predicted in the context of CDM models of galaxy formation \\citep{MO}. The  Tully-Fisher  relations for   diameter  and luminosity  have been extensively  used in the local universe  to determine the distances to galaxies and the value of  the Hubble constant.   We here suggest that they may be used in  a cosmological context to  select in a physically justified way, high redshift  standard rods since galaxies  having the same    rotational  speed will  statistically     have the same narrow distribution in physical sizes.   The picture gets complicated  by the fact  that the standard model  of the universe implies some sort of evolution  in its constituents. In a non  static, expanding universe, where  the scale  factor changes with time,  we  expect various   galaxy properties, such  as  galaxy metric dimensions, to be an explicit function of  redshift. In principle, one may break this circular argument  between model and evolution with two strategies: either by understanding the  effects of different standard rod evolutionary patterns  on  cosmological parameters, or  by looking for cosmological predictions  that  are independent from  the specific form of the disc evolution function.  In this  paper, following the first approach,  we  study how different disc  evolution functions  may   bias the  angular-diameter  test.  We simulate the diameter-redshift experiment using the amount of data and the realistic errors expected in   the context of the zCOSMOS  survey. We  then   evaluate how  different disc   evolution functions   may be unambiguously recognized from the data and  to what extent they affect the estimated values of the various  cosmological parameters.  We also explore the second approach and  show how cosmological information may be  extracted, without any   knowledge about the particular functional form  of the standard rod  evolution, only by requiring  as a prior an estimate  of the upper limit value  for the relative disc evolution at some reference redshift.  This paper is  set out as follow: in  \\S  2 we review the  theoretical basis of the angular diameter  test. In \\S 3  we describe the proposed strategy to select high redshift standard rods.  In \\S 4 we digress on how implement in practice the '$\\theta-z$' test with zCOSMOS data, and in  \\S 5  we present  the zCOSMOS  expected statistical constraints on cosmological  parameters. In   \\S  6 we   discuss different   possible approaches   with which   to  address the    problem of  standard  rod evolution. Conclusions are drawn in \\S 7. ",
        "conclusions": "The scaling of the apparent angular diameter of galaxies with redshift $\\theta(z)$ is  a powerful discriminator  of  cosmological models. The goal of  this  paper  is  to    explore the   potentiality of a    new observational  implementation  of the  classical angular-diameter test and to study its performances and limitations.  We propose to use the velocity-diameter relationship, calibrated using the [OII]$\\lambda 3727$\\AA~ line-widths,  as a tool to select standard rods and probe world  models. As for  other purely geometrical test of cosmology,   a  fair  sampling  of   the   galaxy  population  is  not required.  It is however imperative  to have high quality measurements of the structural parameters of high redshift galaxies (disc sizes and rotational   velocity).   Surveys with HST    imaging  and high enough spectral resolution will thus  provide the fundamental ingredients for the practical realization of the recipe we have presented.  In  order to avoid any luminosity  dependent selection effect (such as for exemple Malmquist bias) it is necessary to apply the proposed test to  high velocity rotators.   We show  that nearly  1300 standard rods with rotational velocity  in the bin $V  \\sim 200\\pm  20$ km s$^{-1}$) are  expected in a field  of size  1 deg$^2$  over  the redshift range $0<z<1.4$. Interestingly this large sample can be quickly assembled by the currently underway zCOSMOS   deep redshift survey, which  uses the VIMOS   multi-object spectrograph  at   the  VLT  to   target galaxies photometrically selected using high-resolution ACS images.  Even allowing a  scatter   of $40\\%$ for  the  [OII]$\\lambda 3727$\\AA~ linewidth-diameter  relationship for disc galaxies,   we show that the angular-diameter diagram constructed using this  sample is affected by a scatter  of  only $ \\sim  5\\%$ per  redshift  bin of amplitude $dz=0.1$. This   scatter translates into  a   $20\\%$ precision in  the \"geometric\" measurement of the dark energy  constant equation of state parameter $w$, through a test  performed without  {\\it priors} in  the [$\\Omega_{m}, \\Omega_{Q}$] space.  Current  theoretical  models suggest   that  large  discs (i.e.   fast rotators)  evolve  weakly with  cosmic time  from  $z=1.5$ down to the present epoch. Anyway, we  have explored how  an eventual evolution of the velocity-selected standard rods might affect the implementation of the test.   We have shown that any  possible evolution in the standard rods may be   unanbiguously revealed by  the  fact that even   a small decrement with redshift of the disc sizes shifts the inferred value of the matter density parameter   into {\"\\em a-priori}  excluded\" regions ($\\Omega_m<0.2$).  We  have shown that  a linear  (as expected on  the  basis  of various theoretical models)  and   substantial (up to   $40\\%$  over the range $0<z<1.5$)   disc evolution minimally biases   the inferred value of a dark energy    component that  behaves  like   Einstein's cosmological constant $\\Lambda$. Moreover  we have shown  that assuming that  discs evolve in a linear-like way as a function  of redshift, and that their sizes were  not more than  $30\\%$ smaller  at $z=1.5$ with  respect to their present  epoch dimension, then  the angular diameter test can be used  to  place  interesting  constraints  in  the  [$\\Omega_{Q},  w$] plane. In particular, assuming a scatter  of $5\\%$ per redshift bin in the  angular  diameter-redshift  diagram (nearly  corresponding to the scatter expected  for   a sample  of   1300  rotators with  $0<z<1.4$, $dz=0.1$, whose diameter   is    locally calibrated with  a     $40\\%$ precision), we have shown that the input fiducial $[\\Omega_Q,w]$ point is still within the 1$\\sigma$ error contours  obtained by applying the angular diameter test to the evolved data.   Finally, we have outlined the strategy to derive a cosmology-evolution diagram with which it is possible to  establish an interesting mapping between different cosmological   models  and  the amount  of    galaxy disc/luminosity evolution     expected  at  a  given    redshift.  The construction of this  diagram   does not  require an  {\\it   a-priori} knowledge of  the     particular  functional  form  of   the    galaxy size/luminosity evolution. By   reading  this diagram, one   can infer cosmological  information  once  a    theoretical  prior  on   disc or luminosity evolution at a given redshift  is assumed. In particular if the amplitude of the relative disc evolution at $\\bar{z}=1.5$ is known to  better than   $\\sim 30\\%$, then   an  Einstein-de Sitter  universe ($\\Omega_m=1$) may be geometrically  discriminated from a flat, vacuum dominated one ($\\Omega_m=0.3$,  $\\Omega_Q=0.7$).  {\\it Viceversa}, one can use  the cosmology-evolution diagram to  place constraints  on the amplitude of the  galaxy disc/luminosity  evolution, once a  preferred cosmology is chosen.   In conclusion, given    the simple ingredients entering   the proposed implementation strategy, nothing, besides evolution of discs, could in principle bias  the test. Even so,  evolution can be  easily diagnosed and, under some general conditions,  it can be  shown that it does not compromise the  possibility of detecting the  presence of  dark energy and constraining the value of its equation of state.  In the following papers  of this series \\citep[]{paperII,paperIII}, we implement the    proposed  strategy   to  a  preliminary    sample  of velocity-selected high redshift rotators."
    },
    "0710/0710.0874_arXiv.txt": {
        "abstract": "The identification of point sources poses a great challenge for the high energy community.  We present a new approach to evaluate the likelihood of a set of sources being a Galactic population based on the simple assumption that galaxies similar to the Milky Way host comparable populations of gamma-ray emitters.  We propose a luminosity constraint on Galactic source populations which complements existing approaches by constraining the abundance and spatial distribution of any objects of Galactic origin, rather than focusing on the properties of a specific candidate emitter.  We use M31 as a proxy for the Milky Way, and demonstrate this technique by applying it to the unidentified {EGRET} sources.  We find that it is highly improbable that the majority of the unidentified {EGRET} sources are members of a Galactic halo population (e.g., dark matter subhalos), but that current observations do not provide any constraints on all of these sources being Galactic objects if they reside entirely in the disk and bulge.  Applying this method to upcoming observations by the Fermi Gamma-ray Space Telescope has the potential to exclude association of an even larger number of unidentified sources with any Galactic source class. ",
        "introduction": "\\label{intro}   The Energetic Gamma Ray Experiment Telescope ({EGRET}) measured gamma-ray emission at energies greater than 100 MeV across the entire sky. More than half of the sources in the third {EGRET} catalog \\citep{hartman_etal_99} were unidentified at the time of publication, and since then only a handful of those sources have been associated with known low-energy counterparts. Theoretical work has produced many candidate sources for these detections, including known Galactic source classes not previously confirmed as gamma-ray emitters such as microquasars \\citep{paredes_etal_00}, and newly proposed populations of high energy Galactic sources such as dark matter annihilation in subhalos~\\citep[e.g.,][]{berezinsky_etal_03,bergstrom_etal_99,blasi_olinto_tyler_03,calcaneoroldan_moore_00,tasitsiomi_olinto_02,taylor_silk_03,ullio_etal_02} and in mini-spikes around intermediate mass black holes \\citep{bertone_etal_05}.  The most direct approach to the confident association of gamma-ray sources with suitable counterparts is multiwavelength follow-up of individual objects \\citep[e.g.,][]{bloom_97, zook_97, gehrels_michelson_99,crb_99,tornik_02,fegan_05, tdr_05, lapalombara_etal_06, paredes_etal_08, reimer_funk_07}. However, the sources in the third {EGRET} catalog typically have large error boxes, making association with known sources by positional coincidence difficult, and the large number of unidentified sources makes multiwavelength follow-up for every object impractical.  Spectral and variability information can help by strengthening or excluding possible associations~\\citep[e.g.,][]{merck_96,roberts_etal_05}.  However, large uncertainties in {EGRET's} spectral and variability data\\footnote{Note that here we refer to general variability properties, such as variability indices, rather than correlated multiwavelength variability which, if present, can confirm the identification of a gamma-ray source with a lower-energy counterpart.}, particularly for the faintest sources, hinder the use of this information to extract sub-populations within the unidentified sources \\citep[e.g.,][]{grenier_00, reimerASSL01, grenier03, zhang_etal_04}. Figure~\\ref{fig:variability} shows the distribution of the initially unidentified {EGRET} sources in the parameters $\\delta$ (variability index from \\citep{nolan_etal_03}) and $\\alpha$ (spectral index from the 3EG catalog) with associated uncertainties. Although in principle association of sources with members of some source classes may be possible using correlations in these parameters, from Figure~\\ref{fig:variability} it is clear that any clustering in the $\\alpha-\\delta$ plane would be significantly smeared by the measurement uncertainties. For this reason, using such clustering to pick out distinct classes of gamma-ray emitters is very difficult with the {EGRET} data, although future experiments that can probe smaller variability timescales and better determine source spectra may be able to use this information to extract source populations.  \\begin{figure}[t] \\centering \\includegraphics[width=0.6\\textwidth]{f1.eps} \\caption{Spectral and variability indices and associated uncertainties for sources originally unidentified in the 3EG catalog.  Spectral data is as published in the 3EG catalog; variability data is taken from \\citet{nolan_etal_03}.\\label{fig:variability}} \\end{figure}  The situation is less favorable for associating sources with members of theoretically-motivated classes of gamma-ray emitters without previously established detections.  The uncertainties inherent in the emission features of unconfirmed source classes make association of individual sources tentative at best.  Association of unidentified sources with members of new source classes by positional correlation with low-energy counterparts has been proposed \\citep{stu_der_95, colaf_02}, but these techniques can produce spurious results \\citep{torres_reimer_05}. Furthermore, approaches to associating sources with new populations which rely on matching unidentified gamma-ray sources with objects detected in other wavelengths cannot be used if the candidate counterparts are not necessarily detectable in other energy ranges (e.g., the case of annihilating dark matter clumps). While emission from dark matter annihilation may have unique signatures such as a ``smoking gun'' spectral line, for many scenarios it is unlikely that even next-generation experiments could detect such a signal.  One way to make progress in making source associations that does not require any knowledge of candidate counterparts is the $\\log N-\\log S$ diagram, the distribution of the number of sources with observed flux \\citep[e.g.,][]{reimer_thompson_01, reimerASSL01, grenier_01, grenier03, bhatta_03, reimer_05}.  However, variability, large uncertainties in measured fluxes, and the presence of multiple populations within the unidentified sources \\citep{grenier_01} pose problems for this method. Still, one advantage of using $\\log N-\\log S$ distributions is that they enable comparison between Galactic and extragalactic populations without requiring a large number of detailed assumptions.  While this is a population-level approach, rather than one which focuses on individual objects, it can help distinguish between Galactic and extragalactic populations, and thus provides information not only about the source populations themselves but also about possible contributions from unresolved members of the same source class to diffuse backgrounds \\citep{pavlidou_siegal-gaskins_fields_etal_08}.  The source distance is perhaps the most obvious difference between extragalactic and Galactic populations, and is an essential ingredient in connecting the source luminosity, which is often motivated by theoretical arguments but not directly observable, with the measured flux.  Previous work has considered the expected distances and luminosities implied by associating unidentified sources with specific Galactic populations \\citep{mukh_95,grenier03}.  However, a challenge for this approach is finding an appropriate ``standard candle'' with which to compare the derived luminosities. In this work we revisit this type of approach, and introduce a novel standard candle with which we can compare the total luminosity of a proposed Galactic population: the total luminosity of M31\\@. By assuming that the nearby galaxy M31 hosts gamma-ray emitting populations similar to those of the Milky Way, we determine whether the bulk of the unidentified sources in a given catalog can be of Galactic origin and place constraints on the spatial distribution of any proposed population without assuming a specific candidate emitter.  Our approach uses only angular position and flux information to evaluate the plausibility of Galactic distributions of these sources. In anticipation of the large source catalogs expected from the recently launched Fermi Gamma-ray Space Telescope\\footnote{http://fermi.gsfc.nasa.gov} and other experiments in the near future, we demonstrate the potential of this technique by applying it to the existing {EGRET} data set.  Confident association of high-energy sources with counterparts is a difficult problem involving many degeneracies.  As a result, individual constraints on the nature of detected sources are often weak. For this reason, as many {\\it independent} constraints as possible are needed to robustly determine the nature of the unidentified sources. The approach presented here, which uses M31 as a standard candle for the gamma-ray luminosity of normal galaxies, involves its own uncertainties which we discuss in detail below.  But, these uncertainties are independent of those of other methods, so in this way our proposed constraint adds a new piece to the puzzle. In \\S\\ref{constraints} we discuss the set of unidentified sources from the third {EGRET} catalog, outline our approach for placing constraints on unidentified source populations, and present our results.  We discuss the potential of this approach in the {\\it Fermi} era in \\S\\ref{fermi} and conclude in \\S\\ref{conc}. ",
        "conclusions": "\\label{conc}  We outlined a new approach to constraining the origin of unidentified gamma-ray sources based on a luminosity constraint.  This approach provides a constraint on candidate populations which is complementary to, and can be used in conjunction with, constraints from existing approaches. We demonstrated this method by applying it to the unidentified sources in the third {EGRET} catalog.  Using the observational upper limit for the gamma-ray luminosity of M31 along with the assumption that M31 and the Milky Way host similar gamma-ray emitting populations, we constrained the allowed spatial distribution of these sources.  We found that it is highly unlikely that a substantial fraction of the {EGRET} unidentified sources are members of a Galactic halo population.  For the case of source populations expected to be associated with the disk and bulge, we find that within reasonable uncertainties all of the {EGRET} unidentified sources can be of Galactic origin without exceeding our assumed luminosity upper limit, although isotropy studies could likely further constrain this scenario.  The new information {\\it Fermi} brings, with respect to both known gamma-ray sources and possible detections of proposed emitters, will help to determine the nature of the unidentified sources.  Statistical techniques such as those employed in this study will be particularly useful in making meaningful statements about gamma-ray emitting populations when large numbers of individual identifications are not feasible.   \\ack We thank J. Beacom, A. Strong, O. Reimer, and I. Grenier for insightful comments and suggestions relating to this work, and D. Thompson for his advice with regard to the observational upper limits from the {EGRET} data.  We are extremely grateful to S. Digel for detailed discussions about the {EGRET} upper limits and diffuse backgrounds.  This work was supported by the Kavli Institute for Cosmological Physics at the University of Chicago through grants NSF PHY-0114422 and NSF PHY-0551142 and an endowment from the Kavli Foundation and its founder Fred Kavli. Support for VP was provided by NASA through the GLAST Fellowship Program, NASA Cooperative Agreement: NNG06DO90A."
    },
    "0710/0710.0335_arXiv.txt": {
        "abstract": "The microquasar \\source~ is one of the most appealing source of the Galactic Centre region. The high energy feature detected once by SIGMA has been searched in the last years by \\integral, but never confirmed. Classified as a persistent source, on 2004 it showed a quiescent-like state. In fact for few month \\source~ was below the detector sensitivity level. We present the long term temporal behaviour of \\source~ observed by \\integral~ and \\rxte~ in 2004 and 2005, as well as preliminary results on possible spectral transitions. ",
        "introduction": "\\source~ was classified as a Black Hole Candidate and reported as the hardest persistent source in the Galactic Centre region \\citep{suny91}. Because of the two-sided radio jets associated to the source, it was classified as a microquasar \\citep{mira92}.  The nature of the normal star companion still remains a mystery. This fact is likely due to the dense surrounding medium, with high concentrations of dust and a large hydrogen column density  (N$_{H} \\sim$ 1.05 $\\times 10^{23}$ cm$^{-2}$ \\citep{gallo02}), which depletes photons from the infrared up to soft X-ray energies.  \\rxte~ long term observations provided an orbital and a super-orbital modulations of 12.5 dy and 600 days, respectively \\citep{smith02a}. The upper-limits on the magnitude obtained in the infrared band \\citep{marti00}, as well as the modulation period measured by \\rxte~suggest that the stellar companion in \\source~ might be a low-mass star.  The electron-positron annihilation radiation detected by SIGMA in 1992 \\citep{bouchet91} has been searched longly by numerous experiments. Reference \\citep{smith96} used simultaneous CGRO/BATSE data to search for the 1992 \\source~ transient, and a $3 \\sigma$ upper limit of 1.8$\\times 10^{-3}$ photons s$^{-1}$ cm$^{-2}$, which contradicted the SIGMA flux, was given. OSSE was also used to observe that feature: no evidence of 511 keV line was found \\citep{jung95}. The history of these and other gamma-ray line transients is reviewed by \\citep{harris98}. Any 511 keV transient event from a compact object in the Galactic Bulge has been searching by \\integral; up to now only upper limits have been provided \\citep{decesare04}.  Thanks to simultaneous \\integral~ and \\rxte~ observations performed in 2003, \\citep{delsanto05} reported on the first broad-band spectral study of \\source. These authors showed a spectral transition from the canonical low/hard state to a peculiar intermediate/soft state, which occurred when the source flux was decreasing. After few days the source quenched in the soft gamma-ray and it was at the detector sensitivity level when observed with X-ray instruments, as also firstly reported by \\citep{greb04} and \\citep{mark04}. However, a similar behaviour was found some years before with SIGMA observations \\citep{kuz97}.  Recently a model of the source emission, from radio to gamma-rays, with a cold-matter dominated jet has been presented \\citep{bosch06}. These authors find out that jet emission cannot explain the high fluxes observed at hard X-rays without violating at the same time the constraints from the radio data, favoring the corona origin of the hard X-rays.  \\begin{figure}[t!] \\centering \\includegraphics[height=8cm,angle=90]{pca_6_12.ps}  \\caption{\\source~ \\rxte/PCA ligh curve based on short pointings in the energy range 6--12 keV.\\label{pca}} \\end{figure}    \\begin{figure*}[t!] \\centering \\includegraphics[height=8cm,angle=90]{1E1740_spring05.ps} \\includegraphics[height=8cm,angle=90]{1E1740_aut05.ps} \\caption{IBIS/ISGRI count rate vs MJD binned by pointing (1800-2200 s) is shown. Observations were performed in Spring 2005 (left) and Autumn 2005 (right). \\label{ibis}} \\end{figure*} ",
        "conclusions": "Black holes in low mass systems (LMXB) are known to spend most of the time in the canonical Low/Hard state \\citep{gallo02}. During this state, a radio activity due to jets emission can be observed; it is strongly suppressed in the high/soft state \\citep{gallo03}.  \\source~ appears to show related radio and X-ray variability \\citep{mirabel93}, although it is radio underluminous in the context of the radio-X-ray luminosity correlation found for black hole candidate X-ray binaries and associated with the accretion/ejection activity \\citep{gallo03}. This might be due to a particularly radiatively efficient corona.  There is a wide debate about the possible origin of the X-rays in microquasars since both the corona and the jet scenarios seem to be roughly consistent in some cases with observations \\citep{marc05}. In \\source~ the hard X-ray emission has been interpreted with a corona emission \\citep{delsanto05}; later, \\citep{bosch06} showed that the hard X-rays cannot be explained if coming from a jet since it would imply an energetic efficiency in the jets significantly larger than for the corona emission.  As observed with \\rxte, for \\source~ and GRS 1758--258 the softest spectra observed are related to dropping photon fluxes \\citep{smith02b}. In 2003, the first broad-band study of \\source~ showed a spectral transition to the intermediate/soft state before the source quenching \\citep{delsanto05}. Our 2005 simultaneous \\integral~ and \\rxte~ observations were performed in high flux levels; as expected \\source~ was in its Low/Hard state. Presumably, a spectral transition would have occurred after few months (see Fig. \\ref{pca}). Temporal and spectral analysis of 2006 are in progress.  The so-called ``dynamical'' soft state \\citep{smith02b} occurring during flux decreasing is explained in the context of two simultaneous and independent accretion flows (i. e., thin and thick). It would occur when the inner geometrically thin disc has yet to respond to a drop in accretion rate that has already depleted the external halo. The delay between the two accretion flows depleting is expected in the LMXB because of the longer viscous timescale of the optically thick disc compared with the halo \\citep{smith02b}.    \\begin{figure}[t] \\includegraphics[height=8cm,angle=90]{hr_1E1740_bin.ps} \\caption{IBIS/ISGRI hardness ratio with orbit bin size.\\label{hr}} \\end{figure}"
    },
    "0710/0710.2937_arXiv.txt": {
        "abstract": "We report null results on a two year photometric search for outburst predictors in SS Cyg. Observations in Johnson V and Cousins I were obtained almost daily for multiple hours per night for two observing seasons. The accumulated data are put through various statistical and visual analysis techniques but fails to detect any outburst predictors. However, analysis of 102 years of AAVSO archival visual data led to the detection of a correlation between a long term quasi-periodic feature at around 1,000-2,000 days in length and an increase in outburst rate. ",
        "introduction": "SS Cyg is one of the most popular and studied variable stars. It is a cataclysmic variable (CV-a.k.a dwarf novae) of the U Gem (UG) class, and prototype of the SS Cyg (UGSS) subclass \\citep{kho85}.  UG class CVs are a binary system consisting of a late main sequence secondary star orbiting a white dwarf primary. The secondary star continuously fills its Roche lobe and transfers mass to the primary, which forms an accretion disc around the primary as angular momentum is conserved and spread within the disc. Occasionally the accretion disc flares up into a bright state referred to as an \"outburst\". The UGSS subclass have outbursts that are usually separated on the order of months, but are largely unpredictable. Outbursts typically have an amplitude of 2-6 magnitudes and last a few days to a few weeks.  The most popular fundamental model for the outburst mechanism in UG stars is the disk instability model first proposed by Osaki (1976). In it, the region where the in-falling matter hits the accretion disc, known as the hot spot, reaches a critical mass. This is a result of a higher rate of mass transfer from the secondary to the disc than from the disc to the white dwarf (due to viscosity).  When the leftover material reaches a critical mass, the accretion disc becomes unstable. Viscosity increases, along with the mass transport rate, through the disc. Both increases are driven by angular momentum. Eventually the system blazes in additional light as the mass built up emits energy due to enhanced ionization of Hydrogen. Only about 3-10\\% of the material from the disc actually falls onto the white dwarf (see panel 3 of figure 13 in Cannizzo [1993]).   \\subsection{Outburst Types}  SS Cyg's quiescence magnitude is around v=12. Outbursts are usually expected every 4-10 weeks, with a duration of 7-18 days each \\citep{jev03}. During outburst, the star reaches maximum light near magnitude v=8. Cannizzo and Mattei (1992) describe the outbursts as bimodal with long outbursts typically lasting $>$12 days and short outbursts lasting $<$12 days.  They assigned the letter \"L\" to the long outbursts and \"S\" to the short and anomalous outbursts. The anomalous outbursts tend to have a linear, lower rate of change and smaller amplitude than long and short outbursts (Figure 1). They found the most common sequence as LS (with 134 occurrences), LLS (69), LSSS (14), and LLSS (8). About 89\\% of all outbursts can fall into one of these four sequences.  \\begin{figure} \\epsscale{1} \\plotone{f1.eps} \\caption{A sample of the SS Cyg visual light curve showing the three category of outbursts: long, short and anomalous (left to right).} \\end{figure}  \\subsection{Outburst Cycles \\& Periodicity}  Honey et al. (1989) photometrically observed SS Cyg over a single long and a single short outburst cycle. They reported substantial flickering on the order of minutes but were unable to find any modulation (including near the orbital period) or periodicity to an amplitude limit of 0.05 magnitude. Giovannelli, Martinez-Pais and Graziati (1992) review the flickering behavior of SS Cyg in both outburst and quiescence, but find no correlation with outburst type, although they do find an inverse correlation between flickering amplitude and system brightness.  Cannizzo and Mattei (1992) analyzed the historical light curve of the American Association of Variable Star Observer's (AAVSO) International Database (hereafter: AID), which at the time included 29,387 individual daily means from 1896 September 27 to 1992 April 7.  They found no correlation between outburst duration time and cycle time but did confirm a correlation between quiescent magnitude and cycle time, caused by a variation in the mass transfer rate of up to a factor of 2 on yearly time scales and a variation of 20\\%-30\\% on decadal time scales. They did not find any periodicity in cycle time.  In a later study \\citep{can98}, using a smaller subset of the data (1963 - 1997), they discovered a variable decay rate in some outburst declines. In a minority of the outbursts, the decay rate slows down about two-thirds of the way into the decline for about a day and a half, then the decay returns to its previous rate.  The significance of the break is proportional to the length of the total decline. The break manifests itself as a 20\\%-300\\% decrease in the decay rate for around one day.  Cannizzo and Mattei (1998) refers to this phenomenon as a \"glitch\", we refer to it hereafter as the \"Cannizzo Glitch\".  Ak et al. (2001) analyzed the quiescence magnitude and outburst cycles of 23 dwarf novae mostly using Fourier analysis. They found no correlation between outburst cycle period and masses of component stars, mean outburst interval, mean outburst duration, mean decline and rise rates of outbursts, absolute quiescent magnitudes and outburst states.  In a poster, Hill and Waagen (2005) suggested the existence of a pre-outburst brightening of about 2\\% over the course of around 12 hours followed by a plateau which lasts another 12 hours leading up to the onset of the outburst. The analyzed data set included about half of the AID visual data and excluded CCD observations. This was the only search for predictors found in published literature.  \\subsection{Long Term Variation}  Kiplinger, et al. (1988) report a 0.15 magnitude amplitude, 7.2 year period in the quiescent behavior of SS Cyg, determined through Fourier analysis of one day means in the AID data from 1896 - 1984. Also, Bianchini (1988) discovered a similar 6.9 year period through Fourier analysis of the intervals between outbursts and attributes it to solar type variation in the secondary \\citep{Bia92}.  However, Richman, Applegate and Patterson's (1994) own analysis of the individual data points plus various long term moving averages find that the power of the reported $\\sim$7 year period is not significantly greater than that of other low frequency signals in the power spectrum and caution that the reported significance is due to the eye being, \"...very prone to spot one to three cycles of periodic behavior in any randomly varying time series.\" However, their analysis excluded all quiescent data points within 12 days on either side of the outburst, thus if the source for the reported $\\sim$7 year period was found in activity near the outburst then it would be hidden from their analysis.  Hemplemann and Kurths (1990) O-C analysis found secular variation within 100 years as well as deviations from the sample mean over intervals of a few tens of cycles. Jevtic, Mattei and Schweitzer (2003) performed a nonlinear analysis using Poincar\\'{e} section and found a period of around 52 years. It is unclear whether the two periods are related as a fundamental and a harmonic pair. ",
        "conclusions": "We have analysed intensive {\\it V} and {\\it I$_{c}$} band observations of SS Cyg over two observing seasons. We were unable to detect any predictors or other activity which could be used to predict an oncoming outburst. However, the combined uncertainty of our photometric data was high enough to warrant further investigation with more precise observations. In particular, some of the data suggest a rise in the ({\\it V-I$_{c}$}) color beginning five days prior to outburst. The rise level was barely within our uncertainty so we cannot report it. A similar dataset in ({\\it B-$I_{c}$}) may increase the photometric sensitivity to such a feature, if it exists.  We also analysed 102 years of AAVSO visual observations. No periodicity was detected which had not already been reported. However, long term quasiperiodic features were detected on the order of 1,000-2,000 days. Their existence is moderately correlated with shortened intervals between outbursts. If the quasiperiodic features are due to solar type variation, as previously reported, then it is possible that the enhanced activity drives additional mass transfer thus decreasing the mean time between outbursts."
    },
    "0710/0710.1872_arXiv.txt": {
        "abstract": "We present a new \\sch orbit-superposition code that is designed to model discrete datasets composed of velocity measurements of individual kinematic tracers in a dynamical system. This constitutes an extension of previous implementations that can only address continuous data in the form of (the moments of) velocity distributions, thus avoiding potentially important losses of information due to data binning. Furthermore, the code can handle any combination of available velocity components, i.e., only line-of-sight velocities, only proper motions, or a combination of both. It can also handle a combination of discrete and continuous data. The code determines the combination of orbital mass weights (representing the distribution function) as a function of the three integrals of motion $E,L_z,$ and $I_3$ that best reproduces, in a maximum-likelihood sense, the available kinematic and photometric observations in a given axisymmetric gravitational potential. The overall best fit is the one that maximizes the likelihood over a parameterized set of trial potentials. The fully numerical approach ensures considerable freedom on the form of the distribution function $f(E,L_z,I_3)$. This allows a very general modeling of the orbital structure, thus avoiding restrictive assumptions about the degree of (an)isotropy of the orbits. We describe the implementation of the discrete code and present a series of tests of its performance based on the modeling of simulated (i.e., artificial) datasets generated from a known distribution function. We explore pseudo-datasets with varying degrees of overall rotation and different inclinations on the plane of the sky, and study the results as a function of relevant observational variables such as the size of the dataset and the type of velocity information available. We find that the discrete \\sch code recovers the original orbital structure, mass-to-light ratio, and inclination of the input datasets to satisfactory accuracy, as quantified by various statistics. The code will be valuable, e.g., for modeling stellar motions in Galactic globular clusters, and modeling the motions of individual stars, planetary nebulae, or globular clusters in nearby galaxies. This can shed new light on the total mass distributions of these systems, with central black holes and dark matter halos being of particular interest. ",
        "introduction": "\\label{sec.intro}  The study of the internal dynamics of stellar systems plays an essential role in astronomy.  From the observed positions and velocities of the stars in galaxies and globular clusters it is possible to infer their total (dark+luminous) mass distribution, which, in particular, provides information on the presence and properties of dark halos and massive black holes. In turn, this structural knowledge constrains theories for the formation and evolution of these systems.  The dynamical state of a stellar system is determined by its phase space distribution function, $f({\\vec r}, {\\vec v})$, which counts the stars as a function of position ${\\vec r}$ and velocity ${\\vec v}$. Typically, however, only three of the six phase-space coordinates are available observationally: the projected sky position $(x',y')$, and the velocity $v_{z'}$ along the line of sight (LOS). Proper motion observations can provide the additional velocities $(v_{x'},v_{y'})$, but such data are generally not available (with the notable exception of some Galactic globular clusters). To make progress with the limited information available, the dynamical theorist is often forced to make simplifying assumptions about geometry (e.g., that the system is spherical) or about the velocity distribution (e.g., that it is isotropic). Such assumptions can have strong effects on the inferred mass distribution (\\citealt{bin82}). To obtain the most accurate results it is therefore important to make models that are as general as possible. Of particular importance for collisionless, unrelaxed systems such as galaxies is to constrain the velocity anisotropy using available data, rather than to assume it a priori.  In a collisionless system the distribution function satisfies the collisionless Boltzmann equation. Analytical methods to find solutions of this equation usually rely on the Jeans Theorem, which states that the distribution function must depend on the phase-space coordinates through integrals of motion (quantities that are conserved along a stellar orbit). In a spherical system all integrals are known analytically, namely, the energy $E$ and the components of the angular momentum vector ${\\vec L}$. Analytical models for spherical systems are therefore fairly easily constructed. In an axisymmetric system things are more complicated (e.g., \\citealt{bt87,mer99}). Only two integrals are known analytically, $E$ and the vertical component $L_{\\rm z}$ of the angular momentum vector\\footnote{We adopt the notation in which $(x,y,z)$ denote the coordinates intrinsic to the axisymmetric stellar system, with the plane $x-y$ being the equatorial plane, and $z$ the symmetry axis. These relate via the inclination $i$ to the observable coordinates $(x',y')$ on the plane of the sky (aligned, respectively, along the projected major and minor axes of the stellar system), and $z'$ the line-of-sight direction, positive in the direction away from us.}, but there is generally a third integral for which no analytical expression exists. Therefore, it is not generally possible to construct an axisymmetric model analytically. The special class of so-called `two-integral' ($f=f(E,L_z)$) models (e.g., \\citealt{bat93,deh94,ver02}) has its uses (e.g., \\citealt{mag98,vdm06}), but these have an isotropic velocity distribution in their meridional plane, which need not be a good fit to real dynamical systems.  The most practical way to model a general axisymmetric system is to do it numerically. While a few methods exist to do this (e.g., \\citealt{m2m,nmagic}), the most common approach uses Schwarzschild's (1979) method. One starts with a trial guess for the gravitational potential $\\Psi$ and then numerically calculates an orbit library that samples integral space in some complete and uniform way. The orbits are integrated for several hundred orbital periods, and the time-averaged intrinsic and projected properties (density, LOS velocity, etc.) are stored as the integration progresses. The construction of a model consists of finding a weighted superposition of the orbits that: (1) reproduces the observed stellar or surface brightness distribution on the sky; and (2) reproduces all available kinematical data to within the observational error bars. Additional constraints can be added to enforce that the distribution function in phase space be smooth and reasonably well behaved, e.g., through regularization or by requiring maximum entropy.  Several axisymmetric Schwarzschild codes have been developed in the last decade (e.g., \\citealt{vdm98,cre99,geb00,val04,tho04}). These codes deal with the situation in which information on the line-of-sight velocity distribution (LOSVD) is available for a set of positions on the projected plane of the sky. This is the case, e.g., when the kinematical data are from long-slit or integral-field spectroscopic observations of unresolved galaxies. The optimization problem for such data can be reduced to a linear matrix equation for which one needs to find the least-squares solution with non-negative weights \\citep{rix97}. One dimension of the matrix corresponds to the number of orbits in the library, while the other corresponds to the number of (luminosity, kinematical and regularization) constraints that must be reproduced. Both dimensions are typically in the range $10^3$--$10^4$. Nonetheless, efficient numerical algorithms exist to find the solution, which yield the orbital and the velocity distribution of the model, as well as the $\\chi^2$ of the fit to the kinematical data. The procedure must then be iterated with different gravitational potentials, to determine the potential that provides the overall best $\\chi^2$. The existing codes have been used and tested extensively (e.g., \\citealt{cre00,cap02,cap06,geb03,ben05,dav06}). Some questions remain, e.g., about the importance of smoothing in phase space, the exact meaning of the confidence regions determined using $\\Delta \\chi^2$ contours, and, in some situations, valid concerns have been raised regarding whether the available data contain enough information so as to warrant the conclusions of the \\sch modeling \\citep{val04,cre04,kra05}. Nevertheless, on the whole Schwarzschild codes have now been established as an accurate and versatile tool to study a wide range of dynamical problems.  A disadvantage of the existing codes is that they cannot be easily applied to the large class of problems in which the kinematical observations come in the form of discrete velocity measurements, rather than as LOSVDs. This is encountered, e.g., when modeling the dynamics of galaxies at large radii, where the low-surface brightness prevents integrated-light spectroscopy. The only available data are then often of a discrete nature, e.g., via the LOS velocities of individual stars in galaxies of the Local Group (e.g., \\citealt{bill00,jan01,jan02,lok02,wil04,lok05,wal06,geh06}), or via planetary nebulae (e.g., \\citealt{dou02,rom03,teo05}) and globular clusters (e.g., \\citealt{cote01,tom04}) surrounding giant ellipticals. The kinematical data available for clusters of galaxies, consisting of redshifts for individual galaxies, are of a similarly discrete nature (e.g., \\citealt{lok03}). The typical datasets in all these cases consist of tens to hundreds of LOS velocities. Galactic globular clusters constitute another class of object for which kinematical data is often available only as discrete measurements, rather than in the form of LOSVDs. From ground-based observations, data sets of individual LOS velocities can be available for up to thousands of stars in these systems (e.g., \\citealt{sun96,may97,rei06}), and for $\\omega$ Cen it has been possible to assemble large samples of proper motions as well \\citep{vleu00}. With the capabilities of {\\it HST}, accurate proper motion data sets with up to $\\sim 10^4$ stars are now becoming available for several more Galactic globular clusters (e.g., \\citealt{mcn03,mcl06}).  Note that discrete datasets do not necessarily provide better or worse information than datasets obtained from integrated-light measurements. Both types of data have their advantages and disadvantages. For discrete datasets, for example, interloper contamination can be a problem (see also the end of Section~\\ref{sec:logL} below). By contrast, for integrated-light measurements, it is often difficult to constrain the wings of the LOSVD due to uncertainties associated with continuum subtraction. Which type of data is most appropriate and most easily obtained depends on the specific object under study. This is therefore not a question that we address in this paper.  Instead, we focus on the issue of how to best analyze discrete data, if that happens to be what is available.  Analyses of discrete datasets have often been more simplified than the analyses that are now common for integrated-light data. For example, the observations are analyzed using the Jeans equations (e.g., \\citealt{ger02,lok03,cote03,dou07}), often with the help of data binning to calculate rotation velocity and velocity dispersion profiles (see, however, the ``spherical'' Schwarzschild models of M87 of \\citealt{rom01}). The disadvantage of such an approach is that not all the information content of the data is used, including information on deviations of the velocity histograms from a Gaussian. Such deviations are important because they constrain the velocity dispersion anisotropy of the system (e.g., \\citealt{vdm93,ger93,ger98}). This anisotropy is an important ingredient in some existing controversies, e.g. regarding the presence of dark halos around elliptical galaxies \\citep{rom03,dek05}. Loss of information can be avoided when large numbers of datapoints are available, as is often the case for globular clusters. It is then possible to create velocity histograms for binned areas on the projected plane of the sky, after which analysis can be done with existing Schwarzschild codes (e.g., \\citealt{bos06}). While this is possible for large datasets, such an approach is not viable for the more typical, smaller datasets that are often available. The availability of Schwarzschild codes that can fully exploit the information content of such smaller datasets would therefore be valuable to advance this subject.  Motivated by these considerations we set out to adapt our existing Schwarzschild code \\citep{vdm98} to deal with discrete datasets. This does not constitute a trivial change, since it changes the constrained superposition procedure from a linear matrix problem to a more complicated maximum likelihood one. For each observed velocity of a particle in the system the question becomes: what is the probability that this velocity would have been observed if the model is correct? The overall likelihood of the data, given a trial model, is the product of these probabilities for all observations. Such likelihood problems have previously been solved for spherical systems \\citep{mer93,vdm00,wu06} and the special class of axisymmetric $f(E,L_z)$ systems \\citep{mer97,wu07}. However, for the axisymmetric Schwarzschild modeling approach the problem corresponds to finding the minimum of a function in a space with a dimension of $10^3$--$10^4$. We show in this work, via the \\sch modeling of simulated datasets, that this problem can indeed be solved successfully and efficiently. Moreover, we follow \\cite{glenn06} and implement in our new code the ability to calculate and fit proper motions in addition to LOS velocities. Applications of the code to real datasets will be presented in forthcoming papers.  The structure of the paper is as follows. In Section \\ref{sec:logL} we phrase the new problem of fitting a \\sch model to a dataset of discrete velocities (of one, two, or three dimensions) of individual kinematic tracers in terms of a likelihood formalism. Section \\ref{sec:code} describes the implementation of the discrete fitting procedure into our existing \\sch code. At the same time, we summarize here the major steps involved in the construction of the probability matrix that describes the likelihood of a given kinematic data point belonging to some particular orbit of the library. We then present in Section \\ref{sec:tools} sets of simulated data that we use for the purpose of testing the performance of the discrete \\sch code. We also describe the known input distribution functions from which these data were drawn. The application of the code to the simulated datasets is presented in Section \\ref{sec:tests}. We present a thorough analysis of the accuracy with which our discrete \\sch code recovers the known distribution function, mass-to-light ratio and inclination used to generate the simulated data. Finally, in Section \\ref{sec:end} we summarize our findings and present our conclusions. ",
        "conclusions": "\\label{sec:end}  Discrete kinematic datasets, composed of velocities of individual tracers (e.g., red giants, planetary nebulae, globular clusters, galaxies, etc.), are routinely being assembled for a variety of stellar systems of all scales (\\S\\,\\ref{sec.intro}). These include not only LOS-velocity surveys. High-quality proper-motion databases already exist for Galactic globular clusters, and future facilities hold the promise of providing the same for stars in the nearest galaxies. However, the most sophisticated tools typically being used in the modeling of these observations were actually developed for the analysis of kinematic data in the form of LOSVDs, a rather different type of velocity information than the case of the velocities of kinematic tracers on a one-by-one basis. As a consequence, the information content of any particular dataset of a discrete nature is likely not being fully exploited. We thus have developed a specific tool for the modeling of discrete datasets, which we have presented in this paper along with detailed tests of its performance based on the modeling of simulated data.  The new tool consists of a \\sch orbit-superposition code that, adapted from the implementation of \\citet{vdm98}, can handle any number of (one-, two-, or three-dimensional) velocities of individual kinematic tracers without relying on any binning of the data. Under the only assumptions that the system is in steady-state equilibrium (i.e., the gravitational potential is not changing in time) and may be well approximated as axisymmetric, the code finds the distribution function (a function of the three integrals of motion $E$, $L_z$, and $I_3$) that best reproduces the observations (the velocities of the tracers as well as the overall light distribution) in a given potential. The fact that the distribution function is free to have any dependence on the three integrals of motion allows for a very general description of the orbital structure, thus avoiding common restrictive assumptions about the degree of (an)isotropy of the orbits.  Unlike previous implementations of the \\sch technique, we cast the problem of finding the best superposition of orbits using a probabilistic approach, i.e., by building a likelihood function representing the probability that the entire set of measurements would have been observed assuming a particular form for the gravitational potential (\\S\\,\\ref{sec:logL}). In this case, and in contrast with the old continuous versions, the dependence of the likelihood function on the orbital weights is non-linear, and the optimization problem can not be reduced to a linear matrix equation. Instead, it becomes a problem of the maximization of a likelihood with respect to the set of weights associated to all possible combinations of the integrals $(E,L_z,I_3)$ that comprise the orbit library (\\S\\,\\ref{sec:logL}), and which accounts for the observed positions and (any-dimensional) velocities of all particles in the dataset, including their uncertainties (\\S\\,\\ref{sec:pij}). After extensive testing, a conjugate gradient algorithm was found to converge satisfactorily to the correct solution and was adopted for the remaining tests of the code's overall performance (\\S\\,\\ref{sec.mkfitin}).  In order to assess the reliability of our discrete \\sch code, we applied it to several sets of simulated data, i.e., artificially generated kinematic observations obtained from a model of an axisymmetric galaxy of which the orbital structure, mass distribution, and inclination are known in advance. Pseudo-datasets were generated from a two-integral phase-space distribution function with varying degrees of overall rotation, types of velocity information (only-LOS, only proper motions, and both), total number of particles, and for two different inclinations on the plane of the sky (\\S\\,\\ref{sec:data}).  Using the various simulated datasets, we studied the recovery of the input orbital structure or DF, mass-to-light ratio, and inclination. For the purposes of these tests, we assumed complete knowledge of the radial profile of the underlying mass distribution and a mass-to-light ratio that remains constant as a function of radius. These restrictions are easily (and must be) lifted when modeling data on real systems, in which case one needs to explore a range of plausible underlying potentials and allow for variations of the mass-to-light ratio to properly account for the possibility of central black holes and dark halos.  Inside the region constrained by data, we find that the distribution function (represented by the corresponding distributions of orbital mass weights) and streaming characteristics of the input datasets are satisfactorily recovered by the \\sch fits when the correct inclination and mass-to-light ratio are known (Figs. \\ref{fig:1Dplots} to \\ref{fig:Ebins55is}). As measured by the mean absolute deviations between the integrated weight distributions, the agreement between the fitted and the input orbital weight distributions as a function of $E$, $L_z$, and $I_3$ is typically of the order of 3\\%, 10\\%, and 20\\%, respectively (the numbers for our worst case being 5\\%, 16\\%, and 25\\%). When eliminating the dependence on $I_3$, the agreement between the fitted and input $E-L_z$ distributions is of the order of 15\\%, with the net rotational behavior of the input datasets cleanly recovered (Figs.  \\ref{fig:2Dplot55ns} and \\ref{fig:2Dplot55is}). Thus, we conclude that the discrete \\sch code can successfully recover the orbital structure of the system under study.  Assuming that the inclination of the system on the plane of the sky is known, we quantified the recovery of the input mass-to-light ratio as a function of the size of the input dataset (Fig. \\ref{fig:MLparabN}) and of the type of kinematic information available (Fig. \\ref{fig:MLparab2}). We studied both the best-fit value as well as the uncertainty in its determination (Fig. \\ref{fig:errors_ML}). The statistical expectation of better results when the amount of observational information is larger (either regarding the number of datapoints or the number of velocity components) is clearly reproduced by our discrete \\sch models. For the smallest datasets used in our testing ($N=100$), and regardless of whether using only-LOS velocities, only proper motions, or both, the best-fit mass-to-light ratio is within 5-10\\% of the input value, with formal $1\\sigma$ uncertainties of the order of 15\\%. When increasing either the number of available measurements or the number of measured velocity components, the mass-to-light ratio is always recovered to better than $\\sim 10\\%$ accuracy, with the corresponding random ($1\\sigma$) uncertainties in the range of 5-10\\%. The discrete \\sch code, therefore, recovers the mass-to-light ratio of the input datasets to satisfactory levels of accuracy.  The recovery of both the mass-to-light ratio and inclination when neither of these quantities are known in advance (as is usually the case with real observations) was studied using a grid of discrete \\sch models, exploring also the dependence on the type of velocity components available (Fig. \\ref{fig:55is_LOS_MU}). We find that the mass-to-light ratio was again successfully recovered, but the best-fit inclination was not identified correctly using small orbit libraries. We found that this was remedied by better sampling the available $(E,L_z,I_3)$ integral space using a larger orbit library (Fig. \\ref{fig:grid_55_90}). For our input datasets with $i=55\\grad$, the best-fit inclination obtained by our models with a large orbit library is $57\\grad$, while for input datasets with $i=90\\grad$ we obtain a best-fit model with $i=80\\grad$. Given the known difficulty of \\sch models in general for determining the inclination of stellar systems, and considering the low relative importance of this parameter compared to other properties such as the orbital structure and the mass-to-light ratio, we regard this small disagreement for the high inclination datasets as acceptable.  In summary, we have shown that our new \\sch code, designed to adequately handle modern datasets composed of discrete measurements of kinematic tracers, doing this without any loss of information due to data binning or restrictive assumptions on the distribution function, is able to constrain satisfactorily the orbital structure, mass-to-light ratio, and inclination of the system under study. Applications to data for Galactic globular clusters and nearby dE galaxies will be presented in future papers. These are only two examples of a large range of dynamical problems in astronomy to which a discrete \\sch code like ours can be applied, so we expect this new tool will contribute to the better understanding of stellar systems in general."
    },
    "0710/0710.1043_arXiv.txt": {
        "abstract": " ",
        "introduction": "The detection of broad polarised lines in Seyfert~2 nuclei stands that this type  of AGNs are  intrinsically the same objects  than Seyfert~1 galaxies and the differences  in their observational properties can be understood as orientation  effects.  Within the Unified Models, the central continuum  source of AGNs and the  broad line region (BLR) are surrounded  by an optically  and geometrically thick  structure of dust  and molecular gas,  likely following  a toroidal  geometry.  The orientation with respect to our line of sight of Seyfert~2 galaxies is such  that  emission  from  the  central engine  is  shielded  by  the molecular torus  and therefore its continuum and  broad emission lines are only observed in polarised light as a result of scattering outside the torus.  Firstly  confirmed  in NGC~1068  (Miller  et  al.  1991), the  central emission  in Seyfert~2  is scattered  by free  electrons in  a conical structure placed along  the axis of the torus,  the so-call ionisation cones, generating  a polar scattering  spectrum with a  position angle perpendicular  to  the scattering  cone.  In  contrast, the  polarised spectrum  of Seyfert~1  galaxies does  not conform  with  simple polar scattering,  exhibiting a wide  diversity of  polarisation properties. These results  state that  the geometry of  a single  polar scattering region is  incomplete to explain  simultaneously all types  of Seyfert polarised spectra.  Based on a study of 36 Seyfert~1 galaxies Smith et al.   (2002) proposed  a model  in  which the  broad-line emission  is originated in a rotating disk  and scattered in two different regions: the  classical  {\\it  ionisation   cones}  responsible  of  the  polar scattering  and an  equatorial  scattering region  located within  the torus  and  co-planar with  the  rotating  disc.   Therefore, the  two extreme cases are  the face--on Seyfert~1 which exhibit  a null or weak polarisation  as a  result  of cancellation  of  polar and  equatorial scattering  and the Seyfert~2,  dominated by  polar scattering  as the equatorial scattering region is completely obscured. In Seyfert~1 with an  intermediate line  of  sight angle,  both  scattering regions  are visible but in general equatorial polarisation dominates.  Interestingly,  a   peculiar  type  of   Seyfert~1  galaxies  exhibits polarised spectra  similar to those  of Seyfert~2, i.e.   dominated by polar scattering.   According to  the model proposed  by Smith  et al. (2002),  these {\\it  polar--scattered}  Seyfert 1  galaxies should  be observed at  an inclination comparable  with the torus  opening angle, and suffer therefore only a moderate extinction through the torus rim. Smith et al.  (2004) estimate  that between 10\\% and 30\\% of Seyfert~1 galaxies are dominated by polar scattering.  Their spectropolarimetric observations identified twelve of  this type of objects.  These twelve objects constitute  the complete sample of  all known polar--scattered Seyfert~1 galaxies.   X--ray studies  of these distinctive  {\\it polar--scattered} Seyfert~1 galaxies are a powerful tool to  prove the basis of the model proposed by Smith  et al.  (2004) and  therefore to further test  the scheme of Unified Models for  AGNs. We have analysed the  eight objects included in  the Smith  et al.  (2004) sample  with available  X-ray  data.  We present for  the first  time, X-ray analysis  performed with  \\xmm\\ of four them  (Fairall 51, Mrk  704, ESO 323-G077, and  IRAS 15091-2107). The results  of the four remaining  ones (Mrk 231, NGC  3227, Mrk 766, and NGC 4593) have been obtained from the literature. In the following section, we  present the results on  the analysis of  the four objects observed by  the first time by  \\xmm.  In Section~\\ref{sect:results}, we  combine   our  results  with  the  four   already  published  {\\it polar--scattered} Seyfert  1 galaxies  and address the  conclusions of this work. ",
        "conclusions": ""
    },
    "0710/0710.5330_arXiv.txt": {
        "abstract": "% {The observed abundance peculiarities of many chemical species relative to the expected cluster metallicity in blue horizontal-branch (BHB) stars presumably appear as a result of atomic diffusion in the photosphere.  The slow rotation (typically $v\\sin{i}<$ 10 km s$^{-1}$) of BHB stars with effective temperatures $T_{\\rm eff}>$ 11,500~K supports this idea since the diffusion mechanism is only effective in a stable stellar atmosphere.} {In this work we search for observational evidence of vertical chemical stratification in the atmospheres of six hot BHB stars: B84, B267 and B279 in M15 and WF2-2541, WF4-3085 and WF4-3485 in M13.} {We undertake an abundance stratification analysis of the stellar atmospheres of the aforementioned stars, based on acquired Keck HIRES spectra.} {We have found from our numerical simulations that three stars (B267, B279 and WF2-2541) show clear signatures of the vertical stratification of iron whose abundance increases toward the lower atmosphere, while the other two stars (B84 and WF4-3485) do not. For WF4-3085 the iron stratification results are inconclusive. B267 also shows a signature of titanium stratification. Our estimates for radial velocity, $v\\sin{i}$ and overall iron, titanium and phosphorus abundances agree with previously published data for these stars after taking the measurement errors into account. The results support the hypothesis regarding the efficiency of atomic diffusion in the stellar atmospheres of BHB stars with $T_{\\rm eff}>$ 11,500~K. } {} ",
        "introduction": "According to the current understanding of stellar evolution, the horizontal-branch (HB) stars are post-main sequence stars that burn helium in their core and hydrogen in a shell (e.g. Moehler \\cite{Moehler04}). In this paper we consider the HB stars that are located in the blue part of the HB, to the left of the RR Lyrae instability strip. Most researchers call them blue horizontal-branch (BHB) stars to distinguish them from the red horizontal-branch (RHB) stars, which exhibit different observational properties. Sandage \\& Wallerstein (\\cite{S+W60}) have found from analysis of the colour-magnitude diagrams of globular clusters\\footnote{Most of the known BHB stars are found in globular clusters.} that the HB % generally becomes bluer with decreasing metallicity. Derived masses of the cool ($T_{\\rm eff}<$11,500~K) BHB stars in the globular cluster NGC~6388 (Moehler \\& Sweigart \\cite{Moehler+S06}) are in a good agreement with the predictions of canonical HB evolution, except for the hot BHB stars with $T_{\\rm eff}>$11,500~K, where the estimated stellar masses seem to be lower than the canonical values.  \\begin{table*}[th] \\parbox[t]{\\textwidth}{ \\centering \\caption[]{Journal of Keck+HIRES spectroscopic observations of the selected hot BHB stars from Behr (\\cite{Behr03b}).} \\begin{tabular}{lcccccc} \\hline \\hline Cluster/Star&     HJD    &Exposure&  S/N & Coverage & Seeing & Slit' width\\\\ &2450000+&   Time (s)   & &  (\\AA)  & (arcsec) & (arcsec) \\\\ \\hline M13/WF2-2541& 1046.7902 &3$\\times$1500& 44 & 3885-6292 & 0.90 & 0.86 \\\\ M13/WF4-3085& 1052.7338 &3$\\times$1200& 37 & 3888-5356 & 1.10 & 0.86\\\\ M13/WF4-3485& 1053.7793 &3$\\times$1200& 34 & 3888-5356 & 0.90 & 0.86\\\\ M15/B84     & 1053.9024 &4$\\times$1400& 34 & 3888-5356 & 0.80 & 0.86\\\\ M15/B267    & 1053.9731 &4$\\times$1400& 26 & 3888-5356 & 0.80 & 0.86\\\\ M15/B279    & 1053.8312 &4$\\times$1400& 34 & 3888-5356 & 0.90 & 0.86\\\\ \\hline \\end{tabular} } \\label{tab1} \\end{table*}   The Hertzsprung-Russell diagrams of some globular clusters show long blue tails (an extension of the HB), populated by very hot BHB stars and extreme horizontal branch (EHB) stars. Published data on the BHB stars argue that the hot BHB stars show remarkable differences in physical properties when compared to the cool BHB stars.  Using high-precision photometry of stars in M13, Ferraro et al. (\\cite{Ferraro+98}) have found gaps in the distribution of stars along the blue tail. One of these gaps, labeled as G1, is located at $T_{\\rm eff}\\sim$11,000-12,000~K. Grundahl et al. (\\cite{Grundahl+98}) used the results of Str\\\"{o}mgren $uvby\\beta$-photometry finding good agreement between the theoretical prediction of stellar evolution models and the observed location of BHB stars, except for the hot BHB stars, whose $u$-magnitudes are brighter than predicted. It appears that this $u$-jump is observed for the hot BHB stars and coincides with the temperature range of the G1 gap in M13.  Similar $u$-jumps have also been found for other globular clusters (Grundahl et al. \\cite{Grundahl+99}). For the hot BHB with effective temperatures up to 20,000 K, the surface gravities derived from the fits of Balmer and helium line profiles appear to be lower than the predictions of stellar evolution models, while the gravities derived for the stars outside this temperature range are in good agreement with theoretical predictions (Moehler et al. \\cite{Moehler+95, Moehler+97a, Moehler+97b, Moehler+03}). The stellar rotation velocity distribution of BHB stars also appears to have a discontinuity at $T_{\\rm eff} \\simeq $ 11,500~K (Peterson et al.~\\cite{Peterson+95}; Behr et al. \\cite{Behr+00a}; Recio-Blanco et al.~\\cite{RB+04}), indicating that the hotter stars show modest rotation with $v\\sin{i}<$ 10 km s$^{-1}$, while the cooler stars rotate more rapidly on average. Comprehensive surveys of abundances also show that the hot BHB stars have abundance anomalies when compared to the cool BHB stars in the same cluster (Glaspey et al. \\cite{Glaspey+89}; Grundahl et al. \\cite{Grundahl+99}; Behr et al. \\cite{Behr+99, Behr+00b}, Behr \\cite{Behr03a}; Fabbian et al. \\cite{Fabbian+05}; Pace et al. \\cite{Pace+06}).   The observed phenomena such as the low gravity, photometric jumps and gaps, abundance anomalies and slow rotation suggest that atomic diffusion could be important in the stellar atmospheres of hot BHB stars. Atomic diffusion arises from the competition between radiative acceleration and gravitational settling. This can produce a net acceleration on atoms and ions, which results in their diffusion through the atmosphere (Michaud~\\cite{Michaud70}). In order for atomic diffusion to produce a vertical stratification of the abundances of particular elements, the stellar atmosphere must be hydrodynamically stable. According to Landstreet (\\cite{Landstreet98}), photospheric convection should be very weak at the effective temperatures of BHB stars. Theoretical atmospheric models of Hui-Bon-Hoa, LeBlanc \\& Hauschildt (\\cite{Hui-Bon-Hoa+00}) showed that the observed photometric jumps and gaps for hot BHB stars can be explained by elemental diffusion in their atmosphere. Behr (\\cite{Behr03b}) has shown that adoption of a microturbulent velocity of 0 or 1 km s$^{-1}$ provides the best fit to line strengths in the spectra of hot BHB stars. This fact supports the proposal that strong velocity fields are not present in the atmospheres of hot BHB stars.  While synthesizing spectral line profiles, Khalack et al. (\\cite{Khalack+07}) have recently found vertical abundance stratification of sulfur in the atmosphere of the field BHB star HD~135485. In this paper we also attempt to detect signatures of vertical abundance stratification of elements from line profile analyses of several other BHB stars for which we have appropriate data. Together with the data on stratification of the sulfur abundance in HD~135485, new positive results would provide a convincing argument in favour of efficient atomic diffusion in the atmospheres of hot BHB stars. In Sec.~\\ref{obs} we discuss the properties of the acquired spectra, while in Sec.~\\ref{mod} we describe details concerning the simulation routine and adopted atmospheric parameters for the program stars. The evidence for vertical stratification of some chemical species is given in Sec.~\\ref{vert}, while the estimation of mean abundances and velocities is described in Sec.~\\ref{mean}. A discussion follows in Sec.~\\ref{discuss}. ",
        "conclusions": "\\label{discuss}  In this paper we continue our attempts to detect vertical abundance stratification in the atmospheres of BHB stars. After the report by Bonifacio et al. (\\cite{Bonifacio+95}) of the vertical stratification of helium in the atmosphere of Feige 86, it became clear that the abundances of other chemical species may also be stratified.   We devoted special interest to iron because hot BHB stars usually have an enhanced iron abundance (e.g. Behr~\\cite{Behr03b}), suggesting that this element may be strongly affected by diffusion. Analysing the spectra of another hot BHB star HD~135485 (Khalack et al. \\cite{Khalack+07}) we did not find direct evidence of iron stratification, but revealed strong signatures of sulfur depletion in the deeper atmospheric layers.  However, HD~135485 is different from the other BHB stars in that its spectrum shows evidence of helium enrichment (in comparison with the solar abundance), while in the atmospheres of the other BHB stars helium is depleted.   Therefore, we directed our attention to BHB stars where the iron abundance is near the solar abundance or enhanced, and helium is depleted. The results obtained argue that at least three stars (B267 and B279 in M15 and % WF2-2541 in M13) show clear signatures of vertical stratification of their iron abundance, while for WF4-3085 the results are suggestive, but not conclusive. The other two stars studied here (B84 in M15 and WF4-3485 in M13) do not show stratification of iron and their averaged iron abundance is close to solar (but is enhanced in comparison with its cluster value). B267 shows also a signature of vertical stratification of titanium (see Fig.~\\ref{B267}b).  Since our simulations show that the turnup feature observed in the iron stratification profile is strongly dependent on microturbulent velocity, the value of the abundance at low optical depths is uncertain. Of course, if the theoretical framework supposes that these abundance gradients are due to atomic diffusion, microturbulence should be weak since a stable atmosphere is needed for diffusion to be dominant. It should be noted that for corresponding optical depths the abundance profile of iron is similar in the three BHB stars that exhibit stratification. The reason that the other two stars in our study do not show clear signs of iron stratification (B84 and WF4-3485) might be related to evolutionary effects or the presence of other competing hydrodynamical processes. The absence of Ti\\,{\\sc ii} and P\\,{\\sc ii} lines in their spectra might be evidence of this.  In conclusion, the results shown here add to the mounting evidence of the existence of vertical abundance stratification, and hence atomic diffusion, in the atmospheres of BHB stars."
    },
    "0710/0710.5106_arXiv.txt": {
        "abstract": "{ The excess of diffuse galactic gamma rays above 1 GeV, as observed by the EGRET telescope on the NASA Compton Gamma Ray Observatory, shows all the key features from Dark Matter (DM) annihilation: (i) the energy spectrum of the excess is the same in all sky directions and is consistent with the gamma rays expected for the  annihilation of  WIMPs with a  mass   between 50-100 GeV; (ii) the intensity distribution of the excess in the sky is used to determine the halo profile, which was found to correspond to the usual profile from N-body simulations with additional substructure in the form of two doughnut-shaped  structures  at radii of 4 and 13 kpc; (iii)  recent N-body simulations of the tidal disruption of the Canis Major dwarf galaxy show that it is a perfect progenitor of the ringlike Monoceros tidal stream of stars at 13 kpc with ring parameters in agreement with the EGRET data; (iiii)  the mass of the outer ring is so large, that its gravitational effects influence both the gas flaring and the rotation curve of the Milky Way. Both effects are clearly observed in agreement with the DMA interpretation of the EGRET excess. } % ",
        "introduction": "If dark matter  (DM) is created thermally during the Big Bang the present relic density  is inversely proportional to $\\langle\\sigma v\\rangle$, the annihilation cross section $\\sigma$ of DM particles, usually called WIMPS (Weakly Interacting Massive Particles), times their relative velocity. The average is taken over these velocities. This inverse proportionality is obvious, if one considers that a higher annihilation rate, given by $\\langle\\sigma v\\rangle n_\\chi$, would have reduced the relic density before freeze-out, i.e. the time, when the expansion rate of the Universe, given by the Hubble constant, became equal to or larger than the annihilation rate.  For the present value of $\\Omega h^2=0.105 \\pm 0.008$, as measured by WMAP \\cite{wmap}, the thermally averaged total cross section at the freeze-out temperature of $m_\\chi/22$ must have been around $3\\cdot 10^{-26} ~{\\rm cm^3s^{-1}}$ \\cite{jungman}. If the s-wave annihilation is dominant, as expected  in many supersymmetric models, then the annihilation cross section is energy independent, i.e. the cross section given above is also valid for the cold temperatures of the present universe \\cite{susy}. Such a large cross section will lead to a production rate of mono-energetic quarks\\footnote{The quarks are mono-energetic, since the kinetic energy of the cold dark matter particles is expected to be negligible with the mass of the particles, so the energy of the quarks equals the mass of the WIMP.} in our Galaxy, which is 40 orders of magnitude above the rate produced at any accelerator. The fragmentation of these  quarks will lead to a large flux of gamma rays with a characteristic energy spectrum quite different from the background of cosmic ray interactions with the interstellar material of the Galaxy. In addition, gamma rays have the advantage that they point back to the source and do not suffer energy losses, so they are the ideal candidates to trace the dark matter density. The charged components interact with Galactic matter and are deflected by the Galactic magnetic fields, so they do not point back to the source. Therefore the charged particle fluxes have large uncertainties from the pro\\-pagation models, which determine how many of the produced particles arrive at the detector. For gamma rays the propagation is straightforward: only the ones pointing towards the detector will be observed.  An excess of diffuse gamma rays compatible with dark matter annihilation (DMA) has indeed been observed by the EGRET telescope on board of NASA's CGRO (Compton Gamma Ray Observatory)\\cite{us}. The excess was observed in all sky directions, which would imply that DM is not dark anymore, but shining in gamma rays.  Of course, such an important observation needs to be scrutinized heavily. Before discussing the criticism the evidence and new confirmation from N-body simulations and the gas flaring is presented in the next section. \\begin{figure} \\begin{center} \\vspace*{-0.65cm} \\includegraphics [width=0.42\\textwidth,clip]{Plots/fig2.eps} \\caption[]{Fit of the shapes of background and DMA signal to the EGRET data in the direction of the Galactic centre The light shaded (yellow) area indicates the background using the shapes known from accelerator experiments, while the dark shaded (red) area corresponds to the signal contribution from DMA for a 60 GeV WIMP mass. The intermediate (blue) shaded area corresponds to a variation of the WIMP mass between 50 and 70 GeV.} \\label{fig2} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0710/0710.1808_arXiv.txt": {
        "abstract": "Observations of a large solar flare of December 13, 2006, using Solar Optical Telescope (SOT) on Hinode spacecraft revealed high-frequency oscillations excited by the flare in the sunspot chromosphere. These oscillations are observed in the region of strong magnetic field of the sunspot umbra, and may provide a new diagnostic tool for probing the structure of sunspots and understanding physical processes in solar flares. ",
        "introduction": "Solar flares represent a process of conversion of magnetic energy into heat, kinetic energy of plasma eruptions and high-energy particles. Solar flares may excite various types of oscillations and waves in various layers of the Sun, from \"sunquakes\" in the interior \\citep{Kosovichev1998} to coronal Morton waves \\citep{Moreton1960} and coronal loop oscillations \\citep{Aschwanden1999}. The mechanisms of these oscillations and waves are not fully understood yet, but obviously related to the energy release and transport properties. For instance, the seismic response (\"sunquakes\") are believed to be related to the hydrodynamic impact on the lower atmosphere and photosphere by shocks generated in the area heated by high-energy electrons. The spatial-temporal properties of the seismic wave source are closely related to the properties of hard X-ray source \\citep{Kosovichev2006a,Kosovichev2006b}. The sunquakes represent packets of high-frequency acoustic waves traveling through the Sun's interior. The waves propagate through the active regions and sunspots with strong magnetic field without significant distortion of the wave front and large changes in the travel times. The magnetic field effects in the sunquake waves, which to some extent are obviously present, have not been detected.  Here, we report on observations from Hinode spacecraft of a different type of oscillations excited by a solar flare. These oscillations are observed in the chromosphere of the sunspot umbra and inner penumbra, and, thus, probably represent some kind of magnetohydrodynamic oscillatory modes. Unfortunately, the relatively low cadence of the observations did not allow us to identify the specific mode of the oscillations. Nevertheless, these oscillations carry potentially interesting information about the flare energy release and transport and properties of sunspots, and deserve further observational and theoretical studies. ",
        "conclusions": "The observations of the solar flare of December 13, 2006, from Hinode reveal a new type of flare-excited oscillations. The oscillations observed in Ca~II~H images appeared in the chromospheric layers of the sunspot umbra immediately after the impulsive phase. They had the amplitude 2--4 times larger than the pre-flare oscillations in the umbra. Also, their frequency seemed to be higher. There is a weak evidence that during the first 30-40 min the oscillations represent waves traveling through the umbra in the direction away from the flare ribbon with a speed of 50--100 km/s. Then, the oscillation become more irregular with some occasional wave packets. The lifetime of these oscillations is probable about 8 hours. The uncertainties in the data analysis and interpretation are caused by the low cadence of the Hinode observing program during the flare. The images were taken every 2 min while the characteristic oscillation period is about 3 min. Most of the oscillation power is probably even at the shorter periods. Thus, for detailed studies of these oscillations it will be important to increase the image cadence.  The image cadence should be sufficiently high to capture the initial waves traveling with a speed of $\\sim 100$ km/s according to our preliminary estimates. This speed indicates that the waves are of an MHD type, and if their speed is of the order of magnitude of the Alfven speed then they should propagate rather low in the sunspot chromosphere. Thus, simultaneous observations in photospheric lines would be interesting. The oscillation amplitude was several times higher than the amplitude of preflare umbral oscillations, which can be as high as  5--6 km/s \\citep{Yoon1995}, and thus may reach supersonic velocities of 10--20 km/s.  Sunspot oscillations have been studied intensively for many years (for a review see \\citet{Staude1999}) but these Hinode observations seem to be first that show enhanced oscillations in the umbra, associated with a solar flare. Further investigations of these oscillations are of great interest for understanding the processes in solar flares and sunspots."
    },
    "0710/0710.1039_arXiv.txt": {
        "abstract": "{Turbulent motions in stellar convection zones generate acoustic energy, part of which is then supplied to normal modes of the star. Their amplitudes result from a balance between the efficiencies of excitation and damping processes in the convection zones. } {We develop a formalism that provides the excitation rates of non-radial global modes excited by turbulent convection. As a first application, we  estimate the impact of non-radial effects on excitation rates and amplitudes of high-angular-degree modes which are observed on the Sun. } {A model of stochastic excitation by turbulent convection has been developed to compute the excitation rates, and it has been successfully applied to solar radial modes (Samadi \\& Goupil 2001, Belkacem et al. 2006b). We generalize this approach to the case of non-radial global modes. This enables us to estimate the energy supplied to  high-($\\ell$) acoustic modes. Qualitative  arguments as well as numerical calculations are used to illustrate the results. } {We find  that non-radial effects  for $p$~modes are non-negligible: \\\\ - for high-$n$ modes (i.e. typically $n > 3$) and for high values of $\\ell$; the power supplied to the oscillations depends on the mode inertia.\\\\ - for low-$n$~modes, independent of the value of $\\ell$, the excitation is dominated by the non-diagonal components of the Reynolds stress term. } {We carried out a numerical investigation of high-$\\ell$ $p$~modes and we find that the validity of the present formalism is limited to $\\ell < 500$ due to the spatial separation of scale assumption. Thus, a model for very high-$\\ell$ $p$-mode excitation rates calls for further theoretical developments, however the formalism is valid for solar $g$~modes, which will be investigated in a paper in preparation.  } ",
        "introduction": "\\label{intro} Amplitudes of solar-like oscillations result from a balance between stochastic excitation and damping in the outermost layers of the convection zone, which extends to near the surface of the star. Accurate measurements of the rate at which acoustic energy is supplied to the solar $p$~modes are available from ground-based observations (GONG, BiSON) as well as from spacecraft (SOHO/GOLF and MDI).  From those measurements and a comparison with theoretical models, it has been possible to demonstrate that excitation is due to eddy motions in the uppermost part of the convection zone and by advection of entropy fluctuations.  Stochastic excitation of \\emph{radial} modes by turbulent convection has been investigated by means of  several semi-analytical approaches \\citep{GK77,GK94,B92,Samadi00I}; they differ from each other in the nature of the assumed  excitation sources, the assumed simplifications and approximations, and also by the way the turbulent convection is described \\citep[see reviews by][]{Stein04,Houdek06}.  Two major mechanisms have nevertheless been identified to drive the resonant $p$~modes of the stellar cavity: the first is related to the Reynolds stress tensor and as such represents a mechanical source of excitation; the second is caused by the advection of turbulent fluctuations of entropy by turbulent motions (the so-called ``entropy source term'') and as it such represents a thermal source of excitation \\citep{GK94,Stein01B}. Samadi \\& Goupil (2001, hereafter Paper~I) proposed a generalized formalism, taking the Reynolds and entropy fluctuation source terms into account.  In this model, the source terms are written as functions of the turbulent kinetic energy spectrum and the temporal-correlation function. This allows us to investigate several possible models of turbulence \\citep{Samadi02II,Samadi02I}.  The results have been compared with GOLF data for radial modes, and the theoretical values are found to be in good agreement with the observations \\citep{Samadi02I}.  Part of the remaining discrepancies has been recently removed by taking into account the asymmetry introduced by turbulent plumes \\citep{Belkacem06a,Belkacem06b}.  In this paper we take an additional step in extending the \\cite{Samadi00I} formalism to the case of non-radial global modes. This will enable us to estimate the excitation rates for a wide variety of $p$ and $g$ modes excited in different types of stars.  The present model provides the energy supplied to the modes by turbulence in inner, as well as outer, stellar convective regions, provided the turbulent model appropriate for the relevant region is used.  Studies of stochastic excitation of solar radial modes \\citep{Samadi02II,Samadi02I} have given us access mainly to the radial properties of turbulence. The present generalized formalism enables us to take into account the horizontal properties of turbulence (through the non-radial components of the Reynolds stress) in the outermost part of the convective zone.  In the Sun, high-angular-degree $p$~modes (as high as one thousand) have been detected \\citep[e.g.,][]{Rabello04}. From an observational point of view, \\cite{Woodard01} found that the energy supplied to the mode increases with $\\ell$, but that above some high-$\\ell$ value, which depend on the radial order $n$ \\citep[see][Fig.~2]{Woodard01}, the energy decreases with increasing $\\ell$. They mentioned the possibility of an unmodelled mechanism of damping. Hence one of the motivations of this work is to investigate such an issue. As a first step,  we develop here a theoretical model of the stochastic excitation taking into account the $\\ell$-dependence of the source terms to seek a physical meaning for such a behaviour of the amplitudes.  Modelling of the mechanisms responsible for the excitation of non-radial modes is not only useful for high-$\\ell$ acoustic modes but also for gravity modes, which are intrinsically non-radial. As for $p$ modes, $g$ modes are stochastically excited by turbulent convection; the main difference is that the dominant restoring force for $g$ modes is buoyancy. We however stress that convective penetration is another possible excitation mechanism for $g$~modes \\citep[e.g.][]{Dintrans05}. Such modes are trapped in the radiative interior of the Sun, so their detection promises to give a much better knowledge of the deep solar interior.  However, such modes are evanescent in the convection zone; their amplitudes at the surface are very small and their detection remains controversial.  A theoretical prediction of their amplitudes is thus an important issue. It requires an estimation of the excitation rates but also  of the damping rates. Unlike $p$~modes, the damping rates cannot be inferred from observations and this introduces considerable uncertainties; e.g. theoretical estimates of the $g$-mode amplitudes \\citep{Gough85,Kumar96} differ from each other by orders of magnitudes, as pointed out by \\cite{CD2002}.  We thus stress that the present  work focuses on the excitation rates -- damping rates are not investigated.  A specific study of gravity modes will be considered in a forthcoming paper.  The paper is organized as follows: Sect.~2 introduces the general formalism, and a detailed derivation of the Reynolds and entropy source terms is provided.  In Sect.~3, we demonstrate that the formalism of \\cite{Samadi00I} is a special case and an asymptotic limit  of the present model.  In Sect.~4, we determine, using qualitative arguments, the different contributions to the excitation rates and identify the dominant terms involving  the angular degree (\\,$\\ell$\\,).  Sect.~5 presents the numerical results: excitation rates are presented and discussed. Sect.~6 discusses the limitations of the model and some conclusions are formulated in Sect.~7. ",
        "conclusions": "\\label{discussion}  \\subsection{The separation of scales} \\label{sep_ech}  The main assumption in this general formalism appears in \\eq{mean_square_amplitude_trapping}, where it has been assumed that the spatial variation of the eigenfunctions is large compared to the typical length scale of turbulence, leading to what we call  \\emph{the separation of scales}. In order to test this assumption, one must  compare the oscillation wavelength to the turbulent one  or equivalently the wavenumbers. To this end, we use the dispersion relation \\cite[see][]{Unno89} \\begin{equation} \\label{dispersion} k_r^2 = \\frac{\\omega^2}{c_s^2}(1 - \\frac{S_\\ell^2}{\\omega^2})(1 - \\frac{N^2}{\\omega^2}) \\quad \\mbox{ and} \\quad k_h^2 = \\frac{L^2}{r^2} \\end {equation} where $N$ is the buoyancy frequency, $S_\\ell$ the Lamb frequency and $k_r,k_h$ the radial and horizontal oscillation wavenumbers, respectively, and $L^2 = \\ell(\\ell+1)$.\\\\ For the turbulent wavenumber,  we choose to use, as a lower limit, the convective wavenumber $k_{conv} = 2 \\pi / L_c$ where $L_c$ is the typical convective length scale. Thus, the assumption of separation of scales is fulfilled provided \\begin{equation} k_{r,h} / k_{conv} \\ll 1 \\end{equation} In Figs.~\\ref{sep_radial} are plotted the ratios $k_r/k_{conv}$ and $k_h/k_{conv}$ respectively.  Those plots  focus on the uppermost part of the solar convection zone where most of the excitation takes place. The assumption of separation of scale  is valid for the horizontal component of the oscillation since  one has $k_h / k_{conv} \\ll 1$ (for $\\ell \\leq 500$) in the region where excitation is dominant. However, we must recall that our criterion  is based on  the mixing length to compute $k_{conv}$.  As shown by \\cite{Samadi02II} using 3D numerical simulations, the convective length scale (computed using the CESAM code, see Sect.~\\ref{calculation_method}) must be multiplied by a factor around five in order to reproduce the injection scale ($L_c$) in the superadiabatic layers.  Hence, for a more conservative criterion, we must then  multiply  the ratio $k_h / k_{conv}$ by a factor of five, this leads to a ratio near unity for  $\\ell \\approx 500$ (see Fig.~\\ref{sep_radial}).  Thus, for higher values of the angular degree, the separation-of-scale hypothesis becomes doubtful.  \\begin{figure}[t] \\begin{center} \\includegraphics[height=6cm,width=9cm]{sep_horizontal.eps}\\\\ \\includegraphics[height=6cm,width=9cm]{sep_radial.eps} \\caption{{\\bf Top:} Ratio of the horizontal oscillation wavenumber to the convective wavenumber ($k_h / k_{conv}$), versus the normalized radius ($r/R$). $k_{conv}$ is computed using the mixing length theory such that $k_{conv} = 2 \\pi / L_c$ ($L_c$ is the mixing length) and $k_{r}$ is computed using the dispersion relation \\eq{dispersion}. Note that the ratio $k_h / k_{conv}$ is computed for a frequency of around $\\nu= 3$ mHz, depending of the angular degree (\\,$\\ell$\\,). {\\bf Bottom:} the same as in the top but for the ratio $k_r / k_{conv}$.} \\label{sep_radial} \\end{center} \\end{figure}  Concerning the radial component of the oscillation wavenumber, the limiting value of $\\ell$ seems to be the same (i.e. $\\ell=500$). Thus, we conclude that for modes of angular degree lower than 500 one can use the separation of scales assumption. For $\\ell >$~500 the characteristic length of the mode becomes smaller than the characteristic length $L_c$ of the energy bearing eddies. Those modes will then be excited by turbulent eddies with length-scale smaller than $L_c$, i.e. lying in the turbulent cascade.  These eddies inject less energy into the mode than the energy bearing eddies do, since they have  less kinetic energy.  We can then expect that -- at fixed frequency -- they received less energy from the turbulent eddies than the low-degree  modes.  A theoretical development is currently underway to properly treat the case of very high $\\ell$ modes.  \\subsection{The closure model}  A second approximation  in the present formalism is the use of a closure model.  The uppermost part of the convection zone is a turbulent convective system composed of two flows (upward and downward) and the probability distribution function of the fluctuations of the vertical velocity and temperature does not obey a Gaussian law \\citep{Lesieur97}.  Thus, the use of the quasi-normal approximation (QNA, \\citet{Million41}) which is exact for a normal distribution, is no longer rigorously correct.  A more realistic  closure model has been developed in \\cite{Belkacem06a} and can be easily be adapted for high-$\\ell$ modes.  This alternative approach  takes into account  the existence of two flows (the up- and downdrafts) within  the convection zone.  However, the QNA is nevertheless often used for the sake of simplicity as is the case here.  Note that, when using the closure model with plumes, it is no longer  consistent to assume that the third-order velocity moments strictly vanish, however as shown by \\cite{Belkacem06a,Belkacem06b}, their contribution is negligible in the sense that their effect is smaller than the accuracy of the presently available observational data.  \\subsection{Mode inertia}  We have shown that the excitation rates for high-$\\ell$ and $n$ modes are sensitive to the variation of the mode inertia ($I$).  $I$ depends on the structure of the stellar model and the properties of the eigenfunctions in  these external regions. \\cite{Samadi06} have shown that different local formulations of convection can change the mode inertia by a small amount. This sensitivity then affects the computed excitation rates ($P$).  However, the changes induced in $P$  are found to be smaller than the accuracy to which the mode excitation rates  are derived from the current observations \\citep[see][]{Baudin05,Belkacem06b}.  Furthermore, concerning the way the modes are obtained, we have computed non-adiabatic eigenfunctions using the time-dependent formalism of Gabriel for convection \\citep[see][]{grigahcene05}.  The mode inertia obtained with these non-adiabatic eigenfunctions exhibits a $\\nu$ dependency different from those obtained using adiabatic eigenfunctions (the  approximation adopted in the present paper).  On the other hand, the mode inertia using non-adiabatic eigenfunctions \\citep[see][for details]{Houdek1999} obtained according to Gough's time-dependent formalism of convection \\citep{Gough77} shows smaller differences with the adiabatic mode inertia. Accordingly, the way the interaction of oscillation and time-dependent convection is modelled affects the eigenfunctions differently.  As explained in Sect.~\\ref{surf}, the formalism developed in this paper can be an efficient tool to derive constraints on the mode inertia to discriminate between the different treatments of convection.  Further work is thus needed on that issue."
    },
    "0710/0710.3499_arXiv.txt": {
        "abstract": "High-resolution ultraviolet observations of the black hole X-ray  binary Cygnus X-1 were obtained using the Space Telescope Imaging Spectrograph on the Hubble Space Telescope. Observations were taken at two epochs roughly one year apart;  orbital phase ranges around $\\phi_{orb}$ = 0 and 0.5 were covered at each epoch. We detect P Cygni line features from high (N~V, C~IV, Si~IV) and absorption lines from low (Si~II, C~II) ionization state material. We analyze the characteristics of a selection of P Cygni profiles and note, in particular,  a strong dependence on orbital phase for the high ionization material:  the profiles show strong, broad absorption components when the X-ray source is behind the companion star and noticeably weaker absorption when the X-ray source is between us and the companion star.  We fit the P~Cygni profiles using the Sobolev with Exact Integration method applied to a spherically symmetric stellar wind subject to X-ray photoionization from the black hole. Of the wind-formed lines, the Si\\,IV doublet provides the most reliable estimates of the parameters of the wind and X-ray illumination. The velocity $v$ increases with radius $r$ (normalized to the stellar radius) according to $v=v_\\infty(1-r_\\star/r)^\\beta$, with $\\beta\\approx0.75$ and $v_\\infty\\approx1420$~km~s$^{-1}$. The microturbulent velocity was $\\approx160$~km~s$^{-1}$. Our fit implies a ratio of X-ray luminosity (in units of 10$^{38}$ erg s$^{-1}$) to wind mass-loss rate (in units of 10$^{-6} \\msun~yr^{-1}$) of L$_{X,38}/\\dot M_{-6} \\approx 0.33$, measured at $\\dot M_{-6}$ = 4.8. The lines from the lower ionization species and the He~II~$\\lambda$1640 absorption are consistent with formation in the photosphere of the normal companion.  Our models determine parameters that may be used to estimate the accretion rate onto the black hole and independently predict the X-ray luminosity. Our predicted L$_x$ matches that determined by contemporaneous RXTE ASM remarkably well, but is a factor of 3 lower than the rate according to Bondi-Hoyle-Littleton spherical wind accretion.  We suggest that some of the energy of accretion may go into powering a jet. ",
        "introduction": "Cygnus X-1, the only Galactic X-ray binary with a high mass companion where existing observations require a black hole for the compact object, was first discovered in 1962 (Cowley 1992, and references therein). In addition to the well-established high mass function found with optical observations, the X-ray data of Cyg X-1 display transitions from a high flux state in the 2-10 keV band (where a strong soft component dominates) to a low flux state (where the soft component largely disappears) that has been interpreted as characteristic of black hole systems. It also displays highly broadened Fe K$\\alpha$ emission (Miller~\\etal~2002) that is consistent with models for X-ray reflection in Galactic black holes and AGNs. The broad-line shape of Fe K$\\alpha$ may be caused by Doppler shifts and the gravitational field of the black hole.  The Cyg X-1 system consists of a supergiant star and a compact object. The mass of the compact object is in the range 7-20 $\\msun$ (Shaposhnikov \\& Titarchuk 2007; Ziolkowski 2005) the mass of the visible star is in the range 18-40$\\msun$ (Ziolkowski 2005; Tarasov~\\etal~2003; Brocksopp~\\etal~1999). The binary orbital period is 5.6 days (Bolton 1972). Miller~\\etal~(2005) using Chandra/HETG observations find that the X-ray spectrum of Cygnus X-1 at phase 0.76 is dominated by absorption lines, in strong contrast to spectra of other HMXBs such as Vela X-1 and Cen X-3. Schulz~\\etal~(2002) report marginal evidence for ionized Fe transitions with P-Cygni type profiles at orbital phase 0.93 whereas Marshall~\\etal~(2001) find no evidence for such line profiles at phase 0.84. Miller~\\etal~(2005) suggest that, while the spectra of Cen X-3 can be modelled by a spherically-symmetric wind (Wojdowski~\\etal~2003), the X-ray absorption spectrum of Cyg X-1 requires dense material preferentially along the line of sight;  considered together, the Chandra spectra provide evidence in X-rays for a focused wind in Cygnus X-1.  Our Space Telescope Imaging Spectrometer (STIS) observations of Cygnus X-1 were obtained when Cyg X-1 was in its soft/high X-ray state and show line profiles that change significantly between orbital phases 0.0 and 0.5. We interpret these changes in terms of models that include the effects of X-ray photoionization on the stellar wind of the normal companion or of a focused wind. We test our model predictions with contemporaneous X-ray observations. The observations and analysis are described in Section 2, our models of the line profiles are presented in Section 3, and our interpretation and conclusions are discussed in Section 4. \\\\ ",
        "conclusions": "The Space Telescope Imaging Spectrograph on Hubble provides the highest resolution ultraviolet spectra taken of Cyg X-1 to date.  Observations were taken at two epochs roughly a year apart:  at each epoch orbital phases when the compact object is behind the stellar companion and when the compact object is in front of the companion star were covered. We find P~Cygni profiles from high ionization (N\\,V, C\\,IV, Si\\,IV) gas. For both epochs the P~Cygni profiles show significantly less absorption at phases when the compact object is in the line of sight. RXTE observations indicate that the X-ray flux of the system was at a similar level at each epoch.  The observed changes can be attributed entirely to orbital effects.  We interpret this to mean that X-rays from the compact object photoionize the wind from the massive companion resulting in reduced absorption by the wind material.  P Cygni profiles of selected species are consistent with the Hatchett-McCray effect, in which X-rays from the compact object photoionize the stellar wind from the companion star, thereby reducing absorption. This effect also appears in UV observations of LMC X-4 and SMC X-1 (Boroson et al. 1999; Vrtilek et al. 1997; Treves et al. 1980). SEI models can fit the observed P Cygni profiles and provide measurements of the stellar wind parameters. The Si\\,IV fits are the most reliable and we use them to determine L$_x/\\dot{M} =4.8 \\pm0.3 \\times 10^{42}$ ergs s$^{-1}$ M$_{\\odot}^{-1}$ yr, where L$_{X}$ is the X-ray luminosity and $\\dot{M}$ is the mass-loss rate of the star. The results from the C\\,IV and N\\,V lines are less reliable because they are saturated and the C\\,IV fit does not match the data well.  For these fits we fixed the terminal velocity $\\nu_{\\infty}$ and the microturbulent velocity to those given by our fits to Si\\,IV.  The best fit values for the optical depth in the ambient wind are high ($\\ge 10$).  Once saturated, the OB star wind lines hardly change with large optical depth, but when much of the wind is ionized by the black hole, the ion fraction in the remaining regions can have a significant effect on the line profile.  The reduced absorption when the compact object is in the line of sight is inconsistent with focusing of the wind toward the compact object, as has been suggested by several authors (e.g, Sowers~\\etal~1998; Tarasov~\\etal~2003; Miller~\\etal~2005), as then we would expect more absorbers in the line-of-sight and hence increased P Cygni absorption. Also, IUE observations taken at 8 orbital phases show a continuous variation in the P Cygni profiles with maximum absorption at phase 0.0 and minimum at 0.5 (Treves~\\etal~1980; van Loon~\\etal~2001). Further high spectral resolution ultraviolet observations of Cyg X-1 will be necessary to study the behavior of the P Cygni lines during different X-ray states.  In an analysis of 2 years of RXTE/ASM data Wen~\\etal~(1999) found the 5.6 day orbital period of Cyg X-1 during the X-ray low/hard state, but no evidence of the orbital period during the high/soft state. Wen \\etal~suggest that absorption of X-rays by a stellar wind from the companion star can reproduce the observed X-ray orbital modulations in the hard state: The lack of modulation in the soft state could be due either to a reduction of the wind density during the soft state or to partial covering of a central hard X-ray emitting region by an accretion stream. Gies~\\etal~(2003) used the results of a four year spectroscopic monitoring program of the H$\\alpha$ emission strength of HDE226868, the normal companion to Cyg X-1, to argue that the low/hard X-ray state occurs when there is a strong, fast wind and accretion rate is low, while in the high/soft state a weaker, highly ionized wind attains only a moderate velocity and the accretion rate increases. The interpretations of both Wen~\\etal~(1999) and Gies~\\etal~(2003) are inconsistent with the fact that the {\\it total} X-ray luminosity from 1.3-200 keV remains constant during both the X-ray soft and hard states (Wen~\\etal~1999): the designation of X-ray high or X-ray low during these states is only applicable for the 1.5-12 keV ASM band. Since the 1.3-200 keV X-ray luminosity is unchanged from the hard to soft state, fluctuations in the narrow ASM band cannot be due to reduction in accretion, or obscuration of the X-ray source; rather it is a physical change that causes the dominant emission mechnism to switch from thermal to power-law.  We note that the orbital modulation observed by Wen~\\etal~in the hard state ($\\pm 1.6$ ASM cts/sec around the average) is less than the errors on the counts during the soft state ($\\pm 3$ ASM cts/sec); we suggest that the quality of the RXTE ASM is not sufficient to detect this low level orbital modulation during the soft state. Ultraviolet observations clearly show orbital modulation during both X-ray hard and soft states (Gies~\\etal~2007). Our wind models explain both the X-ray and ultraviolet flux during the soft/high state.  We need high spectral resolution ultraviolet observations during the hard/low state to determine if there is a change in wind density between states.  Our determination of the mass-accretion rate can be considered as a positive check on the Hatchett-McCray models, as the Si\\,IV line gave $L_x=1.6\\times10^{37}$~erg~s$^{-1}$, and the model is subject to systematic uncertainties. However, this value of $L_x$ is a factor of 3 lower than the best estimate of the accretion rate according to Bondi-Hoyle-Littleton spherical wind accretion. We suggest that some of the energy of accretion may go into powering the jet.  Our test of the dependence of our results on X-ray luminosity confirms the utility of our models and demonstrates that the wind outside the shadow zone is still sensitive to X-ray illumination, an arrangement which allows us to fit $L_{\\rm x}$ as a free parameter in our model.  In future observations of the time-variability of the wind lines, the light travel-time effects may be used to advantage, as the wind may act as a ``low-pass filter\" to the X-ray observations, with the filter cutoff indicating the size of the ionized region (Kallman, McCray, \\&\\ Voit 1987)."
    },
    "0710/0710.3450_arXiv.txt": {
        "abstract": "We performed a timing analysis of the 2003 outburst of the accreting X-ray millisecond pulsar \\xtejb\\ observed by RXTE. Using recently refined orbital parameters we report for the first time a precise estimate of the spin frequency and of the spin frequency derivative. The phase delays of the pulse profile show a strong erratic behavior superposed to what appears as a global spin-up trend.  The erratic behavior of the pulse phases is strongly related to rapid variations of the light curve, making it very difficult to fit these phase delays with a simple law.  As in previous cases, we have therefore analyzed separately the phase delays of the first harmonic and of the second harmonic of the spin frequency, finding that the phases of the second harmonic are far less affected by the erratic behavior. In the hypothesis that the second harmonic pulse phase delays are a good tracer of the spin frequency evolution we give for the first time a estimation of the spin frequency derivative in this source. The source shows a clear spin-up of $\\dot \\nu = 2.5(7) \\times 10^{-14}$ Hz sec$^{-1}$ (1 $\\sigma$ confidence level).  The largest source of uncertainty in the value of the spin-up rate is given by the uncertainties on the source position in the sky. We discuss this systematics on the spin frequency and its derivative. ",
        "introduction": "Binary systems in which one of the two stars is a neutron star (NS hereafter) are among the most powerful X-ray sources of our Galaxy. The emission of X-rays is due to the matter transferred from the companion star and accreted onto the NS, and to the release of the immense gravitational energy during the fall or in the impact with the NS surface. A sub-category of such systems is called Low Mass X-ray Binaries (LMXB). LMXBs are characterized by low NS superficial magnetic fields ($< 10^9~ \\mathrm{Gauss}$) and by the low-mass ($< 1 \\, M_\\odot$) of the companion star. The so-called recycling scenario \\citep[see for a review][]{Bhatta_91} sees in the millisecond radio pulsars the last evolutionary step of LMXBs, where the torques due to the accretion of matter and angular momentum, together with the relatively weak magnetic fields, are able to spin-up the NS to periods of the order of one millisecond. When the accretion phase terminates and the companion star stops transferring matter, the NS can switch on as a millisecond radio pulsar, although no example has been reported yet.  The recycling scenario received the long awaited confirmation only in 1998 with the discovery of the first millisecond X-ray pulsar in a transient LMXB; the first LMXB observed to show coherent pulsations at a frequency of $\\sim 400$ Hz was \\saxj\\ \\citep{Wijnands_98}, in which the NS is orbiting its companion star with a period of $\\sim 2.5$ hr \\citep{Chakra_98}.  Why millisecond X-ray pulsars were so elusive is an argument still debated in literature. A possible reason can be due to the relatively low magnetic fields of these sources which has therefore less capability to channel the accreting matter onto the polar caps, then the chance to see a pulsed emission from a LMXB is quite low \\citep[see e.g.][]{Vaughan_94}, especially at high accretion rates. However, to date 10 LMXBs have been discovered to be accreting millisecond pulsars (see \\citealp{Wijnands_05}~for a review on the first 6 discovered, for the last four see \\citealp{Kaaret_06, Krimm_07, Casella_07, Altamirano_07}), and all of them are in transient systems. They spend most of the time in a quiescent state, with very low luminosities (of the order of $10^{31} - 10^{32}$ ergs/s) and rarely they go into an X-ray outburst with luminosities in the range $10^{36} - 10^{37}$ ergs/s. Although the recycling scenario seems to be confirmed by these discoveries, from timing analysis of accreting millisecond pulsars we now know that some of these sources show spin-down while accreting~ \\citep{Galloway_02, Papitto_07}. This means that it is of fundamental importance to study the far from being understood details of the mechanisms regulating the exchange of angular momentum between the NS and the accreting matter, and chiefly the role of the magnetic field in this exchange. The main way to do this is the study of the pulse phase shifts and their relations with other physical observable parameters of the NS.  The pulse phase shifts are frequently affected by intrinsic long-term variations and/or fluctuations (with which we mean an erratic behavior of the phase delays possibly caused by variations in the instantaneous accretion torques or movements of the accretion footprints on the NS surface, see \\citealp{DiSalvo_07}\\ for a review). Examples of this complex behavior of the pulse phase shifts in accreting millisecond pulsars were already reported in literature.  \\cite{Burderi_06}, who analyzed the 2002 outburst of the accreting millisecond pulsar \\saxj\\ found a jump of 0.2 in the pulse phases of the first harmonic which is not present in the second harmonic phases, which show a much more regular behavior. This change is in correspondence of a change in the slope in the exponential decay of the X-ray light curve (see also \\cite{Hartman_07} for a discussion of the complex phase behavior in other outbursts of \\saxj).  \\cite{Papitto_07} found that the second harmonic of XTE~J1814--338 follows the first harmonic giving approximately the same spin frequency derivative.  A clear model which can explain this behavior is still missing, but these observational evidences seem to suggest that perhaps the second harmonic has a more fundamental physical meaning. For instance it may be related to the emission of both the polar caps while the first harmonic may be dominated by the most intense but less stable polar cap. Another possible explanation comes from possible shape and/or size variations of the accretion footprints related to variations of the accretion rate. \\cite{Romanova_03} found a such behavior in their numerical simulations.  In this paper we report the results of a timing analysis performed on \\xtejb, making use of an improved orbital solution \\citep{Riggio_07}. As in the cases mentioned above, \\xtejb\\ shows erratic fluctuations of the phase delays of the first harmonic and a much more regular behavior of the phase delays derived from the second harmonic. In the hypothesis that the second harmonic pulse phase delays are a good tracer of the spin frequency evolution we can derive a spin-up rate of $2.5(7) \\times 10^{-14}$ Hz/s (1 $\\sigma$ confidence level). ",
        "conclusions": "We have analyzed a long RXTE observation of the accreting millisecond pulsar \\xtejb\\ and reported the results of an accurate timing analysis performed on a time span of about $120$ days, the longest outburst of an accreting millisecond pulsar for which a timing analysis has been performed to date.  We find that the phase delays derived from the first harmonic show an erratic behavior around a global parabolic spin-up trend.  This behavior is similar to that previously shown by two accreting millisecond pulsar, \\saxj\\ \\citep{Burderi_06} and \\xtejc\\ \\citep{Papitto_07}. In the case of the 2002 outburst of \\saxj, the phase delays of the first harmonic show a shift by about 0.2 in phase at day 14 from the beginning of the outburst, when the X-ray flux abruptly changed the slope of the exponential decay.  On the other hand, the phase delays of the second harmonic in \\saxj\\ showed no sign of the phase shift of the first harmonic, and could be fitted by a spin-up in the first part of the outburst plus a barely significant spin-down at the end of the outburst.  In the case of \\xtejc, the fluctuations in the phase delays were visible both in the first harmonic and in the second harmonic, superposed to a global parabolic spin-down trend. \\cite{Papitto_07} have shown that the post-fit phase residuals were strongly anti-correlated to variations of the X-ray light curve. These fluctuations were interpreted as due to movements of the accretion footprints (or accretion column) induced by variations of the X-ray flux.  In the case of \\xtejb, the fluctuations in the phase delays affect mostly the first harmonic, which shows a trend that is very difficult to reproduce with a simple model.  As in the case of \\xtejc, the post-fit phase residuals are clearly anti-correlated with variations observed in the X-ray light curve; from Figure~\\ref{fig1} we see that the phases decrease when the X-ray flux shows rapid increases. It is important to note that the anti-correlation visible between the post-fit phase delays and the X-ray flux is independent of the spin-down or spin-up behavior of the source, since it is observed in \\xtejc, which shows spin-down, and in \\xtejb, which shows spin-up. The correlation between the phase delays and the X-ray flux affects the second harmonic only marginally.  Indeed, there are a few points in the phase delays of the second harmonic that are significantly below the global trend observed in the phase delays, and all of them correspond to flares in the X-ray light curve. Excluding these points marginally affects the values we obtain for the spin frequency and its derivative, but gives a significant improvement of the $\\chi^2$ of the fit.  We find that the phase delays of the harmonic are fitted by a parabolic spin-up model. We have also showed that the quality of the fit is much improved if we use a more physical model in which the spin-up rate decreases exponentially with time following the decrease of the X-ray flux (and hence of the inferred mass accretion rate).  In fact, if the spin-up of the source is related to the mass accretion rate, then it should not be constant with time, but, in first approximation, should decrease proportionally with the mass accretion rate onto the NS. For instance, assuming that accretion of matter and angular momentum occurs at the corotation radius $R_{co}$, the relation between the spin frequency derivative and the mass accretion rate is, from the angular momentum conservation law, $\\nudot = \\Mdot \\, \\sqrt{G M R_{co}}/ 2 \\pi I$, where G is the gravitational constant, M the NS mass and I is the NS moment of inertia; this gives a lower limit on the mass accretion rate since the specific angular momentum at the corotation radius is the maximum that can be transferred to the NS. In the case of \\xtejb, the duration of the outburst is particularly long (about 120 days), and the effect of the global decrease of the mass accretion rate during the outburst should be particularly relevant for this source. Indeed in this case the fit we obtain using an exponentially decreasing spin-up rate is significantly better than using a constant spin-up rate.  From the fit of the phase delays of the second harmonic of \\xtejb\\ with the model discussed above we find a mass accretion rate at the beginning of the outburst of $4(1) \\times 10^{-10}$ M$_\\odot$ yr$^{-1}$.\\footnote{For this estimation we adopted the value of $I=10^{45}$ g cm$^2$, M = 1.4 M$_\\odot$ and NS radius $R_{NS} = 10^6$ cm.} We can compare this mass accretion rate with the X-ray flux of the source at the beginning of the outburst that was $2 \\times 10^{-9}$ ergs cm$^{-2}$ s$^{-1}$ \\citep{Falanga_05} from which we derive an X-ray luminosity of $4.7 \\times 10^{36}$ ergs s$^{-1}$ and a distance to the source of 4.4(6) kpc. Clearly this is only a crude estimation of the distance on the basis of our timing results and future independent estimates are needed to confirm or disprove our hypothesis."
    },
    "0710/0710.1663_arXiv.txt": {
        "abstract": "Recent observations by \\xmm\\ detected rotational pulsations in the total brightness and spectrum of several neutron stars.  To properly interpret the data, accurate modeling of neutron star emission is necessary. Detailed analysis of the shape and strength of the rotational variations allows a measurement of the surface composition and magnetic field, as well as constrains the nuclear equation of state. We discuss our models of the spectra and light curves of two of the most observed neutron stars, \\rxj\\ and \\onee, and discuss some implications of our results and the direction of future work. ",
        "introduction": "}  Thermal radiation from the surface of neutron stars (NSs) can provide invaluable information on the physical properties and evolution of NSs. NS properties, such as the mass $M$ and radius $R$, in turn depend on the poorly constrained physics of the stellar interior, such as the nuclear equation of state (EOS) and quark and superfluid/superconducting properties at supra-nuclear densities. Many NSs are also known to possess strong magnetic fields ($B \\sim 10^{12}-10^{13}$~G), with some well above the quantum critical value ($B\\gg B_{\\rm Q}\\equiv 4.4\\times 10^{13}$~G).  The observed thermal radiation originates in a thin atmospheric layer (with scale height $\\sim 1$~cm) that covers the stellar surface. To properly interpret the observations of NS surface emission and to provide accurate constraints on their physical properties, it is important to understand in detail the radiative behavior of NS atmospheres in the presence of strong magnetic fields (see \\citep{pavlovetal95,holai01,zavlinpavlov02,holai03,zavlin07}, for more detailed references on observations and on previous works in NS atmosphere modeling). The properties of the atmosphere, such as the chemical composition, EOS, and radiative opacities, directly determine the characteristics of the observed spectrum. While the surface composition of the NS is unknown, a great simplification arises due to the efficient gravitational separation of light and heavy elements \\citep{alcockillarionov80}. A pure hydrogen atmosphere is expected even if a small amount of accretion/fallback occurs after NS formation; the total mass of hydrogen needed to form an optically thick atmosphere can be less than $\\sim 10^{16}$~g. On the other hand, a heavy element atmosphere may be possible if no accretion takes place.  The strong magnetic fields present in NS atmospheres significantly increase the binding energies of atoms, molecules, and other bound states (see \\citep{lai01}, for a review).  Abundances of these bound states can be appreciable in the atmospheres of cold NSs (i.e., those with surface temperature $T\\lesssim 10^6$~K; \\citep{laisalpeter97,potekhinetal99}). In addition, the presence of a magnetic field causes emission to be anisotropic and polarized; this must be taken into account when developing radiative transfer codes. The most comprehensive early studies of magnetic NS atmospheres focused on a fully ionized hydrogen plasma and moderate field strengths ($B\\sim 10^{12}-10^{13}$~G; \\citep{miller92,shibanovetal92,pavlovetal94,zaneetal00}). These models are expected to be valid only for relatively high temperatures. More recently, atmosphere models in the ultra-strong field ($B\\gtrsim 10^{14}$~G) and relevant temperature regimes have been presented (\\citep{ozel01,zaneetal01,holai03,lloyd03,vanadelsberglai06}; see also \\citep{bezchastnovetal96,bulikmiller97}, for early work), and all of these rely on the assumption of a fully ionized hydrogen composition. Magnetized non-hydrogen atmospheres have been studied by \\citep{miller92,rajagopaletal97}, but because of the complexity of the atomic physics, the models were necessarily crude (see \\citep{moriho07}, for more details). Only recently has self-consistent atmosphere models \\citep{hoetal03,potekhinetal04,moriho07} using the latest EOS and opacities for partially ionized hydrogen \\citep{potekhinchabrier03,potekhinchabrier04} and mid-$Z$ elements \\citep{morihailey02,morihailey06} been constructed.  The atmosphere models discussed above only describe emission from a local patch of the stellar surface. By taking into account surface magnetic field $\\vecB$ and temperature $T$ distributions, we can construct more physically correct models of emission from NSs. However, these spectra from the whole NS surface are necessarily model-dependent, as the $\\vecB$ and $T$ distributions are unknown. Nevertheless, detailed comparisons of the models with rotation phase-resolved observations is a powerful tool to study NSs, e.g., spectral features that vary with phase are essential to disentangling magnetic field effects from other parameters and to probe the magnetic field geometry on the surface of the star. Indeed there have been recent works attempting to fit magnetic atmosphere spectra to observations of NSs (see \\citep{ho07}, and references therein). Here we describe some of the details and observational applications of our work. ",
        "conclusions": ""
    },
    "0710/0710.2899_arXiv.txt": {
        "abstract": "We report the detection of \\hi~21~cm absorption from the $z=2.289$ damped Lyman-$\\alpha$ system (DLA) towards TXS~0311+430, with the Green Bank Telescope.  The 21~cm absorption has a velocity spread (between nulls) of $\\sim 110$~km~s$^{-1}$ and an integrated optical depth of $\\int \\tau {\\rm d}V = (0.818 \\pm 0.085)$~km~s$^{-1}$.  We also present new Giant Metrewave Radio Telescope 602~MHz imaging of the radio continuum.  TXS~0311+430 is unresolved at this frequency, indicating that the covering factor of the DLA is likely to be high. Combining the integrated optical depth with the DLA \\hi~column density of \\nhi\\ = $(2 \\pm 0.5) \\times 10^{20}$~\\cm, yields a spin temperature of $\\ts = (138 \\pm 36)$~K, assuming a covering factor of unity.  This is the first case of a low spin temperature ($< 350$~K) in a $z > 1$ DLA and is among the lowest ever measured in any DLA. Indeed, the $\\ts$ measured for this DLA is similar to values measured in the Milky Way and local disk galaxies. We also determine a lower limit (Si/H)~$\\gtrsim 1/3$ solar for the DLA metallicity, amongst the highest abundances measured in DLAs at any redshift.  Based on low redshift correlations, the low $\\ts$, large 21~cm absorption width and high metallicity all suggest that the $z \\sim 2.289$ DLA is likely to arise in a massive, luminous disk galaxy. ",
        "introduction": "\\label{sec:intro}  Damped Lyman-$\\alpha$ systems (DLAs) are the highest column density absorbers seen along QSO lines-of-sight, with neutral hydrogen column densities \\nhi\\ $\\ge 2.0\\times 10^{20}$~\\cm. They have long been identified as the precursors of today's galaxies and the primary gas reservoir for star formation at high redshifts.  Despite the recognised importance of the absorbers, their typical size, structure and internal physical conditions remain issues of controversy (e.g. \\citealt{wolfe05}). For example, DLA metallicities, now measured in over 100 absorbers, show very little evolution between $z \\sim 3$ and $z \\sim 0$, with low-metallicity ($-2<$~[Z/H]~$< -1$, where Z~$\\equiv$~Zn, S or Si\\footnote{In the usual notation, [Z/H]~$\\equiv$~log(N(Z)/N(\\hi))$-$log(N(Z)/N (\\hi))$_{\\odot}$.  We use solar values from \\citet{lodders03}.}) absorbers the norm at all redshifts (e.g. \\citealt{kulkarni05}).  This lack of metallicity evolution runs contrary to expectations that the interstellar metallicity should rise towards lower redshifts if DLAs trace the bulk of gas in galaxies.  Moreover, only a few tens of DLAs have their galactic counterparts identified, of which only a small fraction have been spectroscopically confirmed (e.g. \\citealt{chen03}). Our understanding of the basic physical properties of the absorbing galaxies, such as their mass, size, star formation rates and luminosity, remains very limited.  \\hi~21~cm absorption studies of DLAs towards radio-loud quasars provide an independent probe of physical conditions in the absorbers. They can be combined with optical measurements of the \\hi\\ column density of the DLA (from the Lyman-$\\alpha$ line) to obtain the column-density-weighted harmonic mean spin temperature $\\ts$ of the absorbing gas, allowing one to determine the temperature distribution of the \\hi\\ along the line of sight.  Measurements of $\\ts$ in a large number of DLAs may ultimately be used to infer other galactic properties of high redshift systems. For example, in the local Universe, large spiral disks like the Milky Way and M31 typically have low $\\ts$ values ($\\lesssim 300$~K; \\citealt{braun92}) while high $\\ts$ values ($\\sim 1000$~K; \\citealt{young97}) are more common in dwarf galaxies. Tentative evidence for a similar trend has been found in low~$z$ DLAs, out to $z \\sim 0.7$ \\citep{kanekar03}. If such correlations hold for DLAs at all redshifts, measurements of $\\ts$, and its evolution, would provide interesting insights into galaxy evolution.  Unfortunately, despite a number of searches over the past two and a half decades (e.g. \\citealt{briggs83,carilli96,chengalur00,kanekar03}), spin temperature estimates are available for only $15$~DLAs (of which $10$ are lower limits) at $z \\gtrsim 1.7$ (Kanekar et al, in preparation). All of these previous high $z$ measurements yield high spin temperatures of $\\ts \\gg 300$ K.  There are a number of reasons why the crop of $\\ts$ estimates has increased slowly over the last 25 years, including observational issues such as the frequency coverage of radio telecopes and widespread radio frequency interference (RFI) at the low frequencies ($\\lesssim 1$~GHz) of the redshifted 21~cm line. However, an additional important reason for the relatively-small present 21~cm absorption sample is simply the dearth of known DLAs towards radio-loud QSOs suitable for 21~cm absorption follow-up. We have hence been conducting an optical survey of low-frequency-selected radio-loud quasars, specifically designed to increase the number of $\\ts$ estimates in the redshift range $2<z<4$.  While the survey is still in progress, we report here its first results, the detection of 21~cm absorption from the $z \\sim 2.289$ DLA towards TXS~0311+430, the first case of a low spin temperature in a high-$z$ DLA. ",
        "conclusions": "\\label{sec:discuss}  Spin temperature estimates have so far been obtained in more than thirty DLAs at all redshifts, including fifteen detections of 21~cm absorption (Kanekar et al, in preparation).  Only four absorbers, all at $z < 0.6$, show low spin temperatures, $\\ts < 350$~K, typical of lines of sight through the Milky Way and local disks \\citep{kanekar03}\\footnote{Note that this does not include the $z \\sim 0.656$ DLA towards 3C336 \\citep{curran07}, where the background radio source is strongly lobe-dominated, with a very small core fraction. The very extended radio structure (corresponding to a spatial extent of $\\sim 195 \\: h_{71}^{-1}$~kpc at $z \\sim 0.656$) implies that the radio and optical absorption almost certainly arise from very different lines of sight (a possibility discussed by Curran et al.).}  Another eleven systems, all with $\\ts > 500$~K, have been detected in 21~cm absorption (of which only four are at $z \\gtrsim 2$; \\citealt{wolfe79,wolfe81,kanekar06,kanekar07}) while the remaining DLAs without detectable 21~cm absorption all have $3\\sigma$ lower limits of $> 700$~K on the spin temperature.  The preponderance of high spin temperature estimates in high-$z$ DLAs has usually been attributed to a relatively low fraction of the cold neutral phase of \\hi\\ (the CNM), with most of the gas in the warm phase (the WNM) (e.g. \\citealt{carilli96,chengalur00}). This assumes that the absorbers have a similar two-phase structure to that seen in the ISM of the Milky Way (e.g. \\citealt{wolfe03b}).  The low spin temperature of the $z \\sim 2.289$ DLA towards TXS~0311+430 would then imply that a higher fraction of the \\hi\\ along the line of sight is in the CNM compared with the majority of DLAs at $z \\gtrsim 2$.  \\citet{chengalur00} argued that the high derived $\\ts$ values of high-$z$ DLAs could be due to the low metallicities of typical high-$z$ DLAs, with the paucity of metals resulting in fewer radiation pathways for gas cooling (see also \\citealt{young97}). If this is correct, one would expect DLAs with low spin temperatures (such as the absorber towards TXS~0311+430) to have significantly higher metallicities than those of the general DLA population \\citep{kanekar01a}. As expected, and assuming that the Si\\,{\\sc ii}~$\\lambda 1808$ equivalent width has not been over-estimated, the $z \\sim 2.289$ DLA has [Si/H]~$\\ge -0.48$, one of most metal-rich DLAs yet discovered.  It has also been suggested that the observed low 21~cm optical depth in high-$z$ DLAs arises due to covering factor effects (e.g. \\citealt{curran05}).  In the present case, any reduction in the covering factor below the assumed value of unity can only strengthen the case for a low spin temperature (since $\\ts = (138 \\pm 36) \\times f$~K). The only way to alter the conclusion that a sizeable fraction of \\hi\\ along the line of sight is in the cold phase is if there are large spatial differences in the \\hi\\ column densities along the optical and radio lines of sight (e.g. \\citealt{wolfe03b}). For example, if the optical QSO lies behind a ``hole'' in the \\hi\\ column density distribution, the average \\hi\\ column density against the radio QSO could be larger than that measured from the Lyman-$\\alpha$ line, implying a higher spin temperature from Eqn.~\\ref{eqn:tspin}. Unfortunately, it is very difficult to directly test this possibility.  While Galactic \\hi\\ column densities derived from Lyman-$\\alpha$ absorption studies are in excellent agreement with those obtained from \\hi\\ 21~cm emission observations in the same directions, despite the very different spatial resolutions in the two methods, such comparisons have only been carried out for a fairly small number of high latitude lines of sight \\citep{dickey90}.  A similar comparison between \\nhi\\ values derived from Lyman-$\\alpha$ absorption and 21~cm emission has only been possible in one DLA, the $z \\sim 0.009$ absorber towards SBS~1543+593.  Here, the \\hi\\ column densities agree to within a factor of $\\sim 2$, despite the extremely poor spatial resolution ($\\sim 5.3 \\: h_{71}^{-1}$~kpc) of the radio observations \\citep{chengalur02}. While both of these studies suggest that the \\nhi\\ values along the radio and optical lines of sight are likely to be comparable, we cannot formally rule out the possibility of differences in an individual absorber. However, the fact that the expected high metallicity is indeed seen in the $z \\sim 2.289$ DLA (assuming that our Si~{\\sc ii}~$\\lambda$1808 measurement is accurate) is consistent with the interpretation of a high CNM fraction.  Finally, \\citet{kanekar03} noted a relationship between spin temperature and absorber morphology, in that, at low redshifts, low $\\ts$ values are only found in DLAs identified with luminous disk galaxies, while low-$z$, low-luminosity DLAs, associated with dwarf or LSB galaxies, are all found to have high spin temperatures ($\\ts \\gtrsim 700$~K).  It has not been possible to test this empirical relationship at $z > 1$, as the host galaxies of high-$z$ DLAs are rarely detectable. However, if the relationship does extend out to high redshifts, we would expect the $z \\sim 2.289$ DLA to be a massive, luminous disk galaxy.  We note that both the high metallicity and the large velocity spreads seen in the optical low-ionization metal lines and the 21cm absorption are consistent with the absorption arising in a massive galaxy. If so, it should be possible to detect the absorber host with deep imaging; this would be the first direct test of the $\\ts$-morphology relationship at high redshifts.  In summary, we have detected damped Lyman-$\\alpha$ and \\hi~21~cm absorption at $z = 2.289$ towards the quasar TXS~0311+430. We obtain a DLA spin temperature of $\\ts = (138 \\pm 36) \\times f$~K, the first case of a low spin temperature estimate in a high redshift DLA. The low spin temperature, high metallicity and large velocity spread of the 21~cm and metal lines all suggest that the absorber is likely to be a massive disk galaxy."
    },
    "0710/0710.0252_arXiv.txt": {
        "abstract": "{The ANTARES telescope is being built in the Mediterranean Sea. The detector consists of a 3D array of photomultipliers (PMTs) that detects the Cherenkov light induced by the muons produced in neutrino interactions. Other signatures can also be detected. Since the neutrino fluxes from point-like sources are expected to be small, it is of the utmost importance to take advantage of the ANTARES pointing accuracy (angular resolution better than 0.3 degrees for muon events above 10 TeV) to disentangle a possible signal from the unavoidable atmospheric neutrino background. In order to distinguish an excess of neutrino events from the background, several searching algorithms have been developed within the ANTARES collaboration. In this contribution, the discovery potential and sensitivity to point-like sources of the ANTARES neutrino telescope are presented.}  \\begin{document} ",
        "introduction": "The ANTARES collaboration~\\cite{ANTARES} has started the construction of an underwater neutrino telescope in the Mediterranean Sea at a depth of 2475~m. Seven of the twelve lines of the detector have been deployed so far (May 2007) and the whole detector will be commissioned in early 2008. Each line is equipped with 75 10'' photomultiplier tubes (PMTs) joined in triplets making 25 floors along the line. The lines are kept vertical by means of a buoy located at their upper end. The mean distance among lines is about 65 m and the instrumented length starts at 100 m over the seabed and covers about 350 m. The angular resolution for the ANTARES telescope at high energies (above 10~TeV) is better than 0.3$^{\\circ}$. At lower energies, the angular resolution is dominated by the angle between the muon track and the original neutrino direction.  The angular resolution together with the effective area determine the performance of a neutrino telescope. The good pointing accuracy and large effective area at high energies (0.06 km$^2$ at 100 TeV) enables us to search for point-like neutrino sources or, in case no hint of neutrinos source is found, to set restrictive upper limits to neutrino fluxes. ",
        "conclusions": "In this contribution we have reviewed the expected performance of the ANTARES neutrino telescope regarding the search for point-like neutrino sources. Several searching algorithms have been devised in ANTARES. Among them, unbinned techniques turn out to be more efficient than the standardised binning approach. In addition, the method based on the EM algorithm presents very good results without the trade-off of higher dependence on the estimated detector performances. % Discovery potentials in terms of number of events and required neutrino flux to claim the existence of a source at 50\\% of probability has been presented. The comparison with other experiments in term of the sensitivity has also been presented. \\\\ \\\\ {\\it This work is supported by Spanish MEC grants FPA2003-00531 and FPA2006-04277.}"
    },
    "0710/0710.2127_arXiv.txt": {
        "abstract": "We have used high precision differential astrometry from the Palomar High-precision Astrometric Search for Exoplanet Systems (PHASES) project and radial velocity measurements covering a time-span of 20 years to determine the orbital parameters of the 88 Tau A system. 88 Tau is a complex hierarchical multiple system comprising a total of six stars; we have studied the brightest 4, consisting of two short-period pairs orbiting each other with an $\\sim$18-year period.  We present the first orbital solution for one of the short-period pairs, and determine the masses of the components and distance to the system to the level of a few percent. In addition, our astrometric measurements allow us to make the first determination of  the mutual inclinations of the orbits. We find that the sub-systems are not coplanar. ",
        "introduction": "88 Tau (HD 29140, HR 1458, HIP 21402) is a bright ($m_V=4.25, m_K=3.69\\pm 0.25$; Skrutskie et al. 2006\\nocite{2mass}), nearby ($\\sim 50$pc) hierarchical sextuple stellar system \\citep{msc}. The A component contains a pair of systems (designated Aa and Ab) in an $\\sim18$-year \\citep{balega99} orbit that has been resolved by speckle interferometry \\citep{mac87}. The Aa component is a known spectroscopic binary system ($P\\sim 3.57$-day), with a composite spectral type of A5m \\citep{c69}. In previous work it had been noted \\citep{bc88} that the A system is likely complex, with possibly as many as 5 components. \\citet{balega99} noted a discrepancy between the total estimated mass of this system based on photometry and spectral types, and the total mass derived from the visual orbit and {\\em Hipparcos} parallax.  In this work we have determined that, like the Aa component, the Ab component is a double-lined binary; this newly-resolved binary has a period of 7.89 days. Finally, there is a common-proper-motion companion, labeled B, located $\\sim$69 arcseconds away from the A system;  it, too, is known to be a binary \\citep{tg01}.  For clarity we provide a schematic of this complex system in Figure \\ref{fig1}.  \\begin{figure} \\figurenum{1} \\plotone{f1.eps} \\figcaption{A schematic diagram of the 88 Tau system. \\label{fig1}}  \\end{figure} There are several reasons why multiple stellar systems such as 88 Tau merit attention: first, binary orbits make it possible to measure accurate stellar masses and distances, while the larger number of presumably co-eval stars allows one to impose the additional constraint that any given model must accurately match all of the stars. This approach has proven particularly fruitful when applied to another famous hierarchical sextuple system: Castor\\citep{tr02}. Second, as outlined in \\citet{st02}, the relative orientations of the orbital angular momenta allow one to constrain the properties of the cloud from which the stars are thought to have formed, as well as the subsequent dynamical decay process. Despite their value, observational problems have limited the number of triple or higher-order systems with accurately measured orbits to fewer than 10. Given their hierarchical nature it is often the case that either the close system is unresolvable or the outer system has an impractically long orbital period.  With the advent of long-baseline stellar interferometry, and more recently phase-referenced long-baseline interferometric astrometry \\citep{lm04} capable of 10--20 $\\mu$-arcsecond astrometric precision between pairs of stars with separations in the range 0.05--1 arcsecond, it has become possible to resolve the orbital motion of several interesting multiple systems \\citep{kapPeg,v819Her}. Here we report on astrometric and radial velocity measurements of the 88 Tau A system, which allow us to constrain the orbits of the 3.57-day, 7.89-day and 18-year components with improved precision, and for the first time provide a relative orientation of the orbits as well as component masses.  Astrometric measurements were made with the Palomar Testbed Interferometer \\citep{colavita99} as part of the Palomar High-precision Astrometric Search for Exoplanet Systems (PHASES) program \\citep{limits}.  The Palomar Testbed Interferometer 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 the Space Interferometry Mission (SIM). It operates in the J (1.2 $\\mu$m), H (1.6 $\\mu$m), and K (2.2 $\\mu$m) bands and combines starlight from two out of three available 40 cm apertures. The apertures form a triangle with 86 and 110 m baselines. ",
        "conclusions": "PHASES interferometric astrometry has been used together with radial velocity data to measure the orbital parameters of the quadruple star system 88 Tau A, and in particular to resolve the apparent orbital motion of the close Aa1-Aa2 and Ab1-Ab2 pairs. We have made the first determination of the period of the Ab binary system and found it to consist of a pair of nearly equal-mass G stars. The amplitude of the Ab1-Ab2 Center-of-Light motion is only $\\sim 65 \\mu$as, indicating the level of astrometric precision attainable with interferometric astrometry. We are able to resolve the orbital motion of all of the components, and hence determine the orbital inclinations and component masses with a precision of a few percent. Finally, we are able to determine the mutual inclinations of the various orbits."
    },
    "0710/0710.5531_arXiv.txt": {
        "abstract": "{} {Procyon A, a bright  F5 IV-V Sun-like star, is  justifiably regarded as a prime asteroseismological target. This star was repeatedly observed by MOST, a specialized microsatellite providing long-term, non-interrupted broadband photometry of bright targets. So far, the widely anticipated p modes eluded direct photometric detection, though  numerous independent approaches hinted for the presence of signals in the f$\\sim0.5-1.5$\\,mHz range.} {Implementation of an alternative approach in data processing, as well as combination of the MOST data from 2004 and 2005 (264\\,189 measurements in total) helps to reduce the instrumental noise affecting previous reductions, bringing the $3\\sigma$ detection limit down to $\\sim$5.5\\,part-per-million in the $f=0.8-1.2$\\,mHz range.} {This enables to cross-identifiy 16 p-mode frequencies (though not their degrees) which were previously detected via high-precision radial velocity measurements, and provides an estimate of the large spacing,  $\\delta\\nu =0.0540$\\,mHz at $f\\sim1$\\,mHz. The relatively low average amplitude of the detected modes, $a=5.8\\pm0.6$\\,ppm, closely matches the amplitudes inferred from the ground-based spectroscopy and upper limits projected from WIRE photometry. This also explains why such low-amplitude signals eluded the direct-detection approach which exclusively relied  on the MOST 2004 (or 2005) data processed by a standard pipeline.} {} ",
        "introduction": "Procyon A (F5 IV-V) has served as a prime target in many asteroseismological campaigns (see the reviews by Bedding \\& Kjeldsen, 2006,2007, and references therein). While high-precision measurements of radial velocities revealed the  presence of variability caused by solar-like oscillations (\\cite{ma99}) and provided a preliminary identification of p-mode frequencies (\\cite{egge}; \\cite{mart}; \\cite{lecc}), the photometric studies met either limited success (\\cite{brun}) or outright null-detection (\\cite{matt}; \\cite{gun}). The null-detection result was hotly debated ever since (\\cite{bedd};  \\cite{regu}). Attempting to resolve the controversy, we applied an alternative data-processing approach to the abundant photometric data acquired by MOST in 2004-2005. ",
        "conclusions": ""
    },
    "0710/0710.5707_arXiv.txt": {
        "abstract": "In February 1997, the Japanese radio astronomy satellite HALCA was launched to provide the space-bourne element for the VLBI Space Observatory Programme (VSOP) mission. Approximately twenty-five percent of the mission time was dedicated to the VSOP Survey of bright compact Active Galactic Nuclei (AGN) at 5~GHz. This paper, the fifth in the series, presents images and models for the remaining 140 sources not included in Paper III, which contained 102 sources.  For most sources, the plots of the \\uv\\ coverage, the visibility amplitude versus \\uv\\ distance, and the high resolution image are presented.  Model fit parameters to the major radio components are determined, and the brightness temperature of the core component for each source is calculated.  The brightness temperature distributions for all of the sources in the VSOP AGN survey are discussed. ",
        "introduction": "The radio astronomy satellite HALCA (Highly Advanced Laboratory for Communications and Astronomy) was launched by the Institute of Space and Astronautical Science in February 1997 to participate in Very Long Baseline Interferometry (VLBI) observations with arrays of ground radio telescopes.  HALCA provides the longest baselines of the VSOP, an international endeavor that has involved over 28 ground radio telescopes, five tracking stations and three correlators \\citep{hir98,hir00a}.  HALCA was placed in an orbit with an apogee height above the Earth's surface of 21,400\\,km, a perigee height of 560\\,km, and an orbital period of 6.3~hours.  During the seven years of HALCA's mission lifetime, about 75\\% of observing time was used for projects selected by international peer-review from open proposals submitted by the astronomical community in response to Announcements of Opportunity.  This part of the mission's scientific programme constituted the General Observing Time (GOT). The remaining observing time was devoted to a mission-led survey of active galactic nuclei at 5\\,GHz: the VSOP Survey Program. The major goal of the Survey was to determine the statistical properties of the sub-milliarcsecond structure of a complete sample of AGNs.  \\citep{hir00b, fom00a}. % Following the end of the formal international mission period in February 2002, the Japanese-dominated effort continued survey observations until October 2003, when HALCA lost its attitude control capability. This occurred well after the end of the original planned mission lifetime.  This paper is the fifth in the series of VSOP Survey related papers. \\citet{sco04} (henceforth P-III) contains the results for 102 sources which were observed and reduced before 2001 October. \\citet{hor04} (henceforth P-IV) analyzed the cumulative visibilities of those sources to obtain the `typical source structure'. This paper contains the additional 140 survey sources which were successfully observed by VSOP and completes the survey programme observing results.  The brightness temperature properties of the entire sample of sources are discussed. ",
        "conclusions": "We have presented images, models and comments of the 140 sources which were observed as part of the VSOP Survey project that were not covered in P-III. We have combined the brightness temperature measurements and limits found for the entire sample to produce the T$_b$ distribution for the VSS.  We find that about half of the AGN sample of sources reported upon in this paper have significant radio emission in the core component, with $T_b \\ge 10^{12}$~K in the source frame.  Since the maximum brightness temperature one is able to determine using only ground-based arrays is of the order of $10^{12}$~K, our results confirm the necessity of using space VLBI to explore the extremely high brightness temperature regime.  In addition, our Survey results clearly show that by using space VLBI with higher sensitivity, and somewhat higher resolution, the radio cores of many AGN can be successfully imaged.  Because of the variability of many of the sources in the Survey sample, detailed spectral indices of the core components are difficult to determine. However, many of the sources were observed with the VLBA at 15 GHz as part of the VLBA2cm2 survey, and the spectral properties of the cores will be reported elsewhere \\citep{VLBA2cm2}.  It was not possible to slew the HALCA satellite during the observing runs, therefore HALCA was not able to participate in scans of fringe finders, or flux calibrators. It is the absence of these which forces us to label a number of experiments with no space fringes as failures, when it could be the effects of the source structure.  The design of VSOP-2 will allow fringe checks, and also phase referencing experiments, to be performed \\citep{hir00c}.  {% The completion of this survey has been a Quixotic endeavor, and possibly: ``vino a dar en el m\\'as estra\\~no pensamiento que jam\\'as dio loco en el mundo; y fue que le pareci\\'o convenible y necesario, as\\'i para el aumento de su honra como para el servicio de su rep\\'ublica, hacerse caballero andante'' \\citep{quixote}. }"
    },
    "0710/0710.3121_arXiv.txt": {
        "abstract": "We have examined images from the Precision Solar Photometric Telescope (PSPT) at the Mauna Loa Solar Observatory (MLSO) in search of latitudinal variation in the solar photospheric intensity. Along with the expected brightening of the solar activity belts, we have found a weak  enhancement of the mean continuum intensity at polar latitudes (continuum intensity enhancement $\\sim0.1 - 0.2\\%$ corresponding to a brightness temperature enhancement of $\\sim2.5{\\rm K}$). This appears to be thermal in origin and not due to a polar accumulation of weak magnetic elements, with both the continuum and CaIIK intensity distributions shifted towards higher values with little change in shape from their mid-latitude distributions. Since the enhancement is of low spatial frequency and of very small amplitude  it is difficult to separate from systematic instrumental and processing errors. We provide a  thorough discussion of these and conclude that the measurement captures real solar latitudinal intensity variations. ",
        "introduction": "Latitudinal variations in the thermal structure of the solar convective envelope have been implied or required by theories ranging from non-general-relativistic gravity \\citep[]{dic64} to those explaining the solar differential rotation \\citep[e.g.][]{kit95,dur99,rem05,mie06}, meridional circulation, and torsional oscillations \\citep[]{spr03}. Models of the solar differential rotation and meridional circulation reproduce the observed helioseismic rotation profiles \\citep[]{thom96,sch98} best when baroclinicity is introduced via a latitudinal temperature gradient in region of the solar tachocline which spreads into the convection zone proper by turbulent convection \\citep[]{rem05,mie06}. While it is estimated from such models that in the solar photosphere the pole may be as much as a few degrees warmer and the mid-latitudes a couple of degrees cooler than the equatorial region \\citep[e.g][]{bru02,mie06}, these estimates are model dependent and measurements have proven difficult. The difficulties arise because of the intrinsically low amplitude thermal signal expected and the presence of magnetic structures which introduce small scale intensity fluctuations of comparable or greater amplitude.  Thus, previous observations have yielded wide ranging results (Table~\\ref{table1}).  Here we examine full disk images from the Precision Solar Photometric Telescope (PSPT)  at the Mauna Loa Solar Observatory (MLSO). The telescope should by design readily allow measurement of latitudinal temperature variations if they are a few degrees in magnitude. In \\S2 below we review the telescope capabilities and data processing techniques necessary to realize those.  In \\S3 we present the results obtained and in \\S4 discuss the uncertainties in these due to random and systematic errors. Finally, we conclude in \\S5 by putting our methods and results in the context of previous measurements. ",
        "conclusions": "We have measured an enhancement of the photospheric continuum intensity in the solar polar regions using PSPT images. We have carefully examined the properties of the signal and believe it to be of solar origin.  Moreover, the properties of the intensity distributions suggest that the signal is thermal not magnetic in origin, although contributions from very weak network elements associated with the pole-ward extending branch of the solar cycle \\citep[e.g.][]{how81,mul84,she06} cannot be ruled out without processing significantly more data and employing masking of greater severity. This is planned, but limitations of the current ground based solar photometric observations may make a decisive measurement difficult. Space based imaging photometry would be ideal for this and other intriguing problems.  This is particularly true if solar cycle dependencies are to be uncovered.  Such measurements would help constrain global solar models by adding a thermodynamic constraint to the dynamical constraints currently provided by helioseismology \\citep[]{kld88}. Current numerical simulations seem to show little indication of a local temperature minimum at the equator, as our surface measurement implies, but rather display a mid-latitude minimum with a weak local maximum at the equator and stronger one at the pole.  This is true even when a monotonically decreasing pole to equator profile is imposed in the solar tachocline \\citep[]{rem05,mie06}. While we find some hint of equatorial brightening after application of our most restrictive activity masks (bottom row Figure~\\ref{fig3}), it has extremely low amplitude and the measurement is extremely uncertain.  Our efforts can be distinquished from previous work in several ways. Early work (references $1-7$ in Table~1) struggled with instrumental precision, systematic error, and detector gain correction. We too have taken pains to address these. More importantly, no distinction was made in that early work between magnetic and thermal contributions to the signal, other than perhaps qualitative selection of nonactive regions. We have emphasized this distinction by applying a series of masks based on CaIIK intensity to the images before analysis. Previous efforts by Kuhn et al. (references $8-10$ in Table~$10$) have, for the most part, also attempted to separate quiet Sun and facular contributions to the latitudinal signal. The approach in those studies differed from ours in that the facular contribution to the latitudinal variation was estimated statistically based on models of the quiet Sun and facular color and intensity distributions. Interestingly, the pole/mid-latitiude temperature contrast measured in those studies is in good agreement with our measurement. The previous authors, however, reported an additional low latitude brightness enhancement, only seeen very marginally, if at all, in our study at the most severe masking levels. We note, that the one measurement of solar latitudinal intensity variation available using spacecraft data \\citep[]{kuh98} did not remove the facular contribution during analysis. The latitudinal temperature profile obtained in that work looks very similar to our Figure~$6a$, also unmasked.  Finally, we caution that our analysis does not completely eliminate possible contributions to the solar photospheric intensity from an abundance of unresolved magnetic elements of very low magnetic flux density. Correlations between magnetic flux density and CaIIK intensity extend to very low values, making masking based on these images possible, but also suggesting continued contamination of the continuum measurements at very low contrasts (Rast 2003). Unambiguous measurement of a strictly thermal signal in the solar photosphere may require space-based full-disk photometric imaging at high precision and resolution over many wavelengths, combined with detailed modeling of the radiation field. Even that effort may be stymied if in fact there is no quiet Sun."
    },
    "0710/0710.4537_arXiv.txt": {
        "abstract": "Due to projection effects,  coronagraphic observations cannot uniquely  determine parameters relevant to the geoeffectiveness of CMEs, such  as the true propagation speed, width, or source location. The Cone Model for Coronal Mass Ejections (CMEs) has been studied in this respect  and it could be used to obtain these parameters. There are evidences that some CMEs initiate from a flux-rope topology. It seems that these CMEs should be elongated along the flux-rope axis and the cross section of the cone base should be rather elliptical than circular. In the present paper we applied an asymmetric cone model to get the real space parameters of frontsided halo CMEs (HCMEs) recorded by SOHO/LASCO coronagraphs in 2002. The cone model parameters are generated through a fitting procedure to the projected speeds measured at different position angles on the plane of the sky. We consider models with the apex of the cone located at the center and  surface of the Sun. The results are compared to the standard symmetric cone model. ",
        "introduction": "A halo coronal mass ejection (HCME) was first recorded by Howard in 1982 (Howard \\emph{et al.}, 1982). Since then, HCME are routinely recorded in white light by coronagraphs placed in  space. In coronagraphic observations, HCMEs appear as an enhancement surrounding the entire occulting disk.  HCMEs originating close to the disk center are often responsible for the severest geomagnetic storms (Gosling, 1993; Kahler, 1992; Webb \\emph{et al.}, 2000). For space weather forecast it is very important to determine the kinetic and geometric parameters describing HCMEs. Unfortunately coronagraphic observations are subjected to projection effects. Viewing in the plane of the sky does not allow to determine the true 3-D space velocity, width and source location of a given CME. It is widely accepted that the geometrical structure  of CMEs may be described by the cone model (e.g., Howard \\emph{et al.}, 1982; Fisher and Munro, 1984; St.Cyr \\emph{et al.}, 2000; Webb \\emph{et al}, 2000; Zhao \\emph{et al.}, 2002; Michalek \\emph{et al.}, 2003; Xie \\emph{et al.}, 2004; Xue \\emph{et al.}, 2005). Assuming that the shape of HCMEs is a cone and they propagate with constant angular widths and speeds, at least in their early phase of propagation, a technique was developed (Michalek \\emph{et al.}, 2003) which can determine the following parameters: the linear distance $r$ of the source location measured from the solar disk center,  the angular distance $\\gamma$ of the source location measured from the plane of the sky, the angular width $\\alpha$ (cone angle =$0.5\\alpha$) and the 3-D space velocity $V_{S}$ of a given HCME. This technique required measurements of the sky-plane speed and the moment of the first appearance of the halo CME above opposite limbs. If we determine spatial parameters using only two measurements, large random errors would occur. What is more, this technique was limited to asymmetric events not originating from close to the center of the Sun. A similar cone model was used recently by Xie \\emph{et~al.}~(2004) to determine the angular width and orientation of HCMEs. To improve accuracy, in the present attempt the space parameters of HCMEs are determined by fitting the cone model to projected speeds ($V_{P}$) obtained from height-time plots at different position angles. Although  many thousands of CMEs were recorded by LASCO coronagraphs the 3D structure of  CMEs is still open question. Many authors believe that CMEs originate from a flux-rope geometry, (e.g., Chen \\emph{et al.}, 1997; Dere \\emph{et al.}, 1999; Chen \\emph{et al.}, 2000; Plunket \\emph{et al.}, 2000; Forbes 2000; Krall \\emph{et al.}, 2001; Chen and Krall 2003). If CMEs have a flux-rope geometry,  they should be elongated along the flux-rope axis and the cross section of the cone base should be rather elliptical than circular. In the present approach we consider the asymmetric cone model where an eccentricity and orientation of the cone base are new free parameters. We try  to identify where to locate the apex of the cone, either at   the center of the Sun (Zhao \\emph{et al.}, 2002; Xie \\emph{et al.}, 2004) or on the solar surface (Michalek \\emph{et al.}, 2003). It is important to note that the real elliptical cone model was first developed by Cremades and Bothmer (2004). The model was introduced based on observations of cylindrical shaped CMEs (Cremades and Bothmer, 2004, 2005). They applied it to 32 halo CMEs. In their approach, the best parameter values describing the ellipse are determined from a LASCO image sequence that showed a sharp leading edge. In our method we derive the best-fit parameter values for  halo CMEs by working in the velocity space.  The paper is organized as follows: in Section~2 the asymmetric cone model is presented, numerical simulations and fitting procedure are explained in Section~3 and in Section~4 the results are described. Final conclusions are given in Section 5. ",
        "conclusions": "In this paper we have presented the new asymmetric cone model  to determine the geometrical  parameters characterizing  HCMEs. We applied this model to all  frontsided HCMEs originating close to the disk center and listed in the SOHO/LASCO catalog in 2002. We estimated, with very good accuracy (the average r.m.s error is equal to 60~km~s$^{-1}$), the cone parameters for $15$ events. It was shown (from the polar plots and the r.m.s errors) that the ACM can reproduce, around the entire occulting disk, the measured projected speed much better than the SCM. We have to note that sometimes the individual projected speeds may significantly differ from the fitting model even for the ACM. In these cases errors are probably caused by inaccurate measurements. Unfortunately, most HCMEs are not sufficiently bright around the entire occulting disk to generate the height-time plots with the same good accuracy for  all considered position angles. The biggest errors may arise when the faintest part of a given event is considered or when measurements are disturbed by another event appearing in the LASCO field of view. We also confirmed that most of the events have better fits when the cone apex is located at the center of the Sun.  We used this model for the frontsided full HCMEs originating close to the disk center but it can also be applied to limb and even partial HCMEs.  When considering limb or partial HCMEs the accuracy will be slightly worse. The determined  parameters for HCMEs are similar to that derived by the other cone models, but for some events differences could amount to even 20\\%.  We have to remember that this approach has two shortcomings: (1) CMEs may be accelerating, moving with constant speed or decelerating at the beginning phase of propagation. This means that the constant velocity assumption may be invalid.(2) CMEs may expand in addition to radial motion. Then the measured sky-plane speed is a sum of the expansion speed and the projected radial speed. This would also imply that  CMEs may not be a rigid cone as we had assumed. Unfortunately,  having observation from only one spacecraft there is no possibility to overcome these assumptions. There are two attempts  to get the real space parameters of CMEs. Some scientist (e.g. Xie \\emph{et al.}, 2004) try to get the real parameters of CMEs by considering the elliptical shapes of observed CMEs.  In our approach we consider velocities of CMEs as a function of the position angle. We take into account the projected velocities (many points) from an entire halo CME, not only from a few chosen  points. We do not introduce any assumption about the apparent (geometrical) shapes of CMEs. Therefore, we expect that our determination of parameters may have  better accuracy. The determined real velocities of halo CMEs could be potentially useful for space weather forecast. But we have to note that although it is true that faster interplanetary CMEs are more geoeffective, the strength and topology of magnetic field is still crucial to define geoeffectiveness of CMEs (e.g., Bothmer, 2003)."
    },
    "0710/0710.1062_arXiv.txt": {
        "abstract": "In this paper, we report on the gas-phase abundance of singly-ionized iron (Fe II) for 51 lines of sight, using data from the {\\it Far Ultraviolet Spectroscopic Explorer} ({\\it FUSE}).  Fe II column densities are derived by measuring the equivalent widths of several ultraviolet absorption lines and subsequently fitting those to a curve of growth.  Our derivation of Fe II column densities and abundances creates the largest sample of iron abundances in moderately- to highly-reddened lines of sight explored with {\\it FUSE}, lines of sight that are on average more reddened than lines of sight in previous {\\it Copernicus} studies.  We present three major results.  First, we observe the well-established correlation between iron depletion and $\\nHavg$ and also find trends between iron depletion and other line of sight parameters (e.g.~$\\fHmol$, $\\ebv$, and $\\av$), and examine the significance of these trends.  Of note, a few of our lines of sight probe larger densities than previously explored and we do not see significantly enhanced depletion effects.  Second, we present two detections of an extremely weak Fe II line at 1901.773 \\AA{} in the archival STIS spectra of two lines of sight (HD 24534 and HD 93222).  We compare these detections to the column densities derived through {\\it FUSE} spectra and comment on the line's $f$-value and utility for future studies of Fe II.  Lastly, we present strong anecdotal evidence that the Fe II $f$-values derived empirically in through {\\it FUSE} data are more accurate than previous values that have been theoretically calculated, with the probable exception of $f_{1112}$. ",
        "introduction": "\\label{s:FeII_intro} Iron is both a relatively abundant element in the Galaxy and a major constituent in most grain models.  Fe I has an ionization potential of only 7.87 eV, while Fe II has an ionization potential of 16.18 eV.  Therefore, the dominant form of gas-phase interstellar iron found in H I regions should be Fe II.  In H II regions, iron should be found as a mix of Fe II, Fe III, and Fe IV.  In $\\Hmol$ regions, some Fe I may exist.  The average gas-phase abundance of iron that has been established implies that the vast majority of interstellar iron is tied up in dust.  However, observed variations in the iron abundance, while having only marginal implications for the dominant ingredients of dust, may still shed light on the physical conditions of interstellar clouds.  Three major studies of interstellar gas-phase iron have been carried out: \\citet{SavageBohlin} and \\citet*{Jenkins1986} with {\\it Copernicus} data; and \\citet*{SRF2002} with {\\it FUSE} data.  (For the rest of this paper, these three papers will be referred to as SB1979, JSS1986, and SRF2002, respectively.)  \\citet{Howk} also studied Fe II with {\\it FUSE} but for the purpose of empirically determining oscillator strengths ($f$-values) in the {\\it FUSE} wavelength region rather than studying abundances and depletions.  In summarizing previous studies and discussing our own results, it will be useful to define a few terms.  By abundance, we mean the ratio of an element relative to hydrogen, given either logarithmically or linearly.  We will define the iron abundance as Fe II/$\\Htot=\\NFeII/[\\NHI + 2\\NHmol]$, and neglect other ionization states of iron and hydrogen for the reasons discussed above.  The quantity $\\delta$ is the ratio of the gas-phase abundance seen in a particular line of sight to an assumed cosmic abundance standard.  Despite difference in terminology in the literature (cf.~SB1979 and SRF2002), we will refer to the quantity $D=-\\log{\\delta}$ as the depletion.  It is the logarithm of the ratio of the total amount of an element to the amount in the gas-phase.  Therefore, an increase in depletion means an increase in the amount of that element tied up in dust grains and/or molecules.  This definition of depletion is useful because errors in the assumed cosmic standard produce a constant shift in the quantity $D$ and trends can still be analyzed.  SB1979 found the interstellar ratio of iron to hydrogen to be equal to $4.5\\times10^{-7}$.  They observed depletions in the range of $D=0.11-2.49$.  SB1979 found positive correlations between $\\NFeII$ and the quantities $\\NHtot$ and $\\ebv$.  These results do not consider depletion, and could be consistent with constant depletion.  However, they did find iron depletion to be correlated, to varying degrees, with several parameters.  The first very rough correlation they cite is between depletion and $E_{1330-V}/\\ebv$, a measure of dust grain size.  SB1979 suggested that this implies that small grains are removed through grain growth and coalescence as opposed to destruction.  However, other dust-related parameters such as visual polarization from grains and the gas-to-dust ratio (as measured by $\\NHtot / \\ebv$) were not found to correlate with iron depletion.  Clearer correlations did exist for $\\nHavg$ ($\\equiv \\NHtot/r$, where $r$ is the line-of-sight pathlength) and $\\ebv/r$, two quantities that should show similar correlations because of the correlation between $\\NHtot$ and $\\ebv$ \\citep{Bohlin1978}.  In these correlations, iron depletion increases with increased average density in the line of sight.  Two other measures of cloud conditions that SB1979 found to be correlated to iron depletion were the molecular fraction of hydrogen, $\\fHmol$---positively correlated with depletion---and kinetic temperatue, $T_K$---negatively correlated with depletion.  These correlations were somewhat rough, however.  All of these correlations ($\\nHavg$, $\\ebv/r$, $\\fHmol$, and $T_K$) imply that iron depletion increases in denser environments, possibly due to density-dependent formation or destruction processes or shielding from $\\Hmol$.  Finally, SB1979 also found a correlation between the depletions of Fe and Ti, a result that, when combined with other results from the literature \\citep{deBoerLamers, Stokes}, suggests similarities between grains that contain Fe, Ca, Mn, and Ti.  JSS1986 measured the abundances of several refractory elements (Mg II, P II, Cl I, Cl II, Mn II, Fe II, Cu II, and Ni II) with {\\it Copernicus}.  The major conclusion of the JSS1986 study was that depletions of all the elements are strongly correlated with the average hydrogen density in the line of sight, $\\nHavg$, with increasing depletion related to increased density.  Depletions of various elements are also correlated with each other.  JSS1986 interpreted these results largely in the context of the model of \\citet{Spitzer1985}, where the ISM is considered to be comprised of two types of clouds, warm and cold.  According to this model, individual lines of sight should exhibit properties that are a result of sampling both types of clouds, and the best parameter for characterizing whether warm or cold clouds are sampled is $\\nHavg$.  JSS1986 found $\\logFeIIH=-6.46\\pm0.06$ for cold clouds and $\\logFeIIH=-5.84\\pm0.05$ for warm clouds.  The observed random deviations from a smooth trend in the relationship with $\\nHavg$ were larger than those expected from just the errors in measuring the iron and hydrogen column densities; JSS1986 used this fact to estimate that the true dispersion in the iron abundance is between 0.03 and 0.11 dex.  It is worth noting that JSS1986 expressed two observational concerns that may have affected their data:  (1) the possibility of ``hidden,'' highly-saturated components that, though undetected, might contribute significantly to the measured column densities and (2) significant contributions to the column densities from H II regions that violate the assumption that the true iron abundance, $\\NFe/\\NHtot$, is accurately represented by $\\NFeII / [\\NHI + \\NHmol]$.  SRF2002 examined Fe II abundances with {\\it FUSE}, and explored 18 reddened lines of sight with greater average extinction, reddening, and molecular fraction of hydrogen than the SB1979 and JSS1986 surveys, though the SRF2002 sample did not probe lines of sight with values of $\\nHavg$ that were significantly larger than the densest lines of sight in SB1979 and JSS1986.  SRF2002 found that while Fe II depletion increases with increased $\\ebv$ up to about $\\ebv \\approx 0.35\\magnitude$, primarily in the {\\it Copernicus} targets, the trend levels off at larger values of $\\ebv$ (see Fig. 3 of SRF2002).  No correlation was found between depletion and $\\av$ (Fig. 4 of SRF2002), for a wide range of extinction ($\\av \\approx 1-3.5\\magnitude$); in fact, the depletions appear to be very constant.  SRF2002 similarly concluded that there is little correlation between iron depletion and the molecular fraction of hydrogen (Fig. 5 of SRF2002, including the results of SB1979).  Lastly, Fig. 6 of SRF2002 shows the correlation of iron depletion with average line of sight density, $\\nHavg$.  That figure shows that the SRF2002 results are consistent with the correlation already observed by SB1979 and JSS1986, though, as stated above, SRF2002 did not probe larger $\\nHavg$ than those surveys.  SRF2002 discussed in detail potential interpretations of the observed correlations (or lack thereof) and the fact that significantly enhanced depletions are $not$ observed in clouds with large $\\ebv$, $\\av$, and $\\fHmol$, contrary to what was expected prior to the survey.  That discussion involved the creation and destruction of grain mantles, the warm and cold cloud mixing model of the ISM suggested by \\citet{Spitzer1985}, and the definition of {\\it translucent lines of sight} as opposed to {\\it translucent clouds}.  This latter idea is also discussed in \\citet{SnowMcCall}, with the principle being that diffuse material can make a significant contribution to the reddening and extinction in a line of sight that might otherwise be assumed to match the assumed characteristics of a ``translucent cloud.''  We will return to some of these ideas in \\S \\ref{s:FeII_results}.  A more recent but smaller survey by \\citet{Miller2007} examined the iron and silicon abundances in six lines of sight characterized as ``translucent.''  \\citeauthor{Miller2007}~found that both iron and silicon depletion increases in lines of sight with increased values of the extinction parameter $c_4$ \\citep[in the scheme of][]{FM1988}.  Larger values of $c_4$ should correspond to an increase in the population of small grains.  Therefore, \\citet{Miller2007} suggest that when additional depletion of iron and silicon is observed, those elements have been incorporated into smaller grains.  It is important to note that \\citet{Miller2007} used weak lines in the STIS wavelength region to measure the Fe II abundance.  Four of the six lines of sight from \\citet{Miller2007} have also been observed with {\\it FUSE} and have measurements of the iron abundance found in either SRF2002 (HD 27778 and HD 207198) or this paper (HD 147888 and HD 152590).  We have carried out this study for the following reasons (roughly in decreasing order of importance):  \\begin{enumerate} \\item To increase the sample size of reddened lines of sight with measurements of iron abundances and depletions.  All of our lines of sight have $\\logHtot \\gtrsim 21$.  Additionally, $\\sim \\frac{2}{3}$ of our lines of sight are characterized by $\\av \\gtrsim 1$,  $\\sim \\frac{2}{3}$ by $\\ebv \\gtrsim 0.3$, and  $\\sim \\frac{2}{3}$ by $\\fHmol \\gtrsim 0.1$.  In conjunction with the 16 unique lines of sight from SRF2002---we have reanalyzed HD 24534 and HD 73882 in this paper---our data nearly quadruples the size of the overall sample of reddened lines of sight examined with {\\it FUSE}.  Additionally, while most of the average line of sight densities in our sample are not significantly denser than the densities in previous studies, two lines of sight, HD 147888 and HD 179406, do probe larger $\\nHavg$ than any line of sight in SRF2002. \\item To probe lines of sight with less extinction and/or reddening than SRF2002.  This is useful because certain extinction parameters (specifically $\\av$ and $\\rv$) were not available in many cases for {\\it Copernicus} targets (cf.~Figure 4 of SRF2002).  In addition, many {\\it Copernicus} targets were short-pathlength lines of sight toward bright stars; using {\\it FUSE}, we are able to probe fainter targets and longer pathlengths (though there is some overlap between the two samples). \\item To present two detections of a very weak line of Fe II ($\\lambda_{\\rm rest}$=1901.773 \\AA{}) in archival STIS spectra (lines of sight toward HD 24534 and HD 93222).  We place limits on the $f$-value of this line and discuss its utility for future determinations of Fe II. \\item To evaluate the $f$-values of the Fe II used in this study, comparing the empirical values from \\citet{Howk} and the theoretical calculations of \\citet{RU1998}. \\item To make simple models of Fe II in several lines of sight where we are also studying the 1260 \\AA{} line (along with other absorption lines) of Si II (A. G. Jensen \\& T. P. Snow, in preparation). \\end{enumerate}  In \\S \\ref{s:FeII_obsdata} we discuss our observations and data reduction, including comments on the absorption lines used and our measurement methods.  In \\S \\ref{s:methods} we discuss our methods of deriving column densities, abundances, and depletions.  In \\S \\ref{s:FeII_results} we discuss our results and in \\S \\ref{s:FeII_summary} we summarize. ",
        "conclusions": "\\label{s:FeII_results} As stated above, our derived column densities are given in Table \\ref{coltable}, while curves of growth can be seen in Figures \\ref{fig:cogs1-15}-\\ref{fig:cogs46-51}.  For our sample (not including the unique lines of sight from SRF2002) we find that Fe II/$\\Htot=3.2\\pm0.1\\times10^{-7}$, an average weighted by the inverse squares of the errors in the abundances.  This is somewhat smaller than the results of SB1979, Fe II/$\\Htot=4.5\\times10^{-7}$; however, the two numbers are calculated differently.  Our calculation is a weighted average of the abundances in each line of sight; the SB1979 abundance is a ``total'' abundance, calculated from the ratio of total observed Fe II to total observed hydrogen.  If we calculate a total abundance in this manner, we obtain a total abundance that is larger, Fe II/$\\Htot=5.6\\times10^{-7}$.  The weighted average tends to favor the smaller abundances which also have smaller errors, whereas calculating a total abundance slightly favors the tail of larger abundances, which are up to a few times larger than the median.  Whether we consider an average or the total iron abundance, our results are distinctly different from the results of SRF2002; the 16 lines of sight unique to that study have a weighted average of $1.9\\pm0.2\\times10^{-7}$ and a total abundance of $3.5\\times10^{-7}$.  The fact that our abundances are larger in both cases has a simple interpretation---our study covers some of the same highly-reddened ground as SRF2002, but with an additional sample of much less-reddened lines of sight.  We note here that the averages of JSS1986, discussed in \\S \\ref{s:FeII_intro}, correspond to $1.4^{+0.2}_{-0.1}\\times10^{-6}$ for the ``warm'' ISM and $3.5\\pm0.5\\times10^{-7}$ for the ``cold'' ISM; our weighted average is similar to the latter sample, but our total abundance falls squarely between the two (as it should if our sample covers both ranges of the two JSS1986 samples).  Much more insightful than the averages, however, are that we also see several statistically significant correlations between iron depletion and various line of sight parameters; these correlations and their interpretations are discussed in \\S \\ref{ss:FeII_correlations}.  In these plots and analyses, we include the results of SRF2002 but not SB1979 or JSS1986.  As discussed in \\S \\ref{s:FeII_intro}, our sample, compared to these previous samples, is either larger or covers more parameter space.  There are potential systematic errors between the {\\it FUSE} studies (this paper and SRF2002) and SB1979 and JSS1986.  SB1979 used the 1122.0, 1133.7, and 1144.9 \\AA{} lines of Fe II, while JSS1986 used the 1096.9, 1122.0, and 1133.7 \\AA{} lines.  In addition to the differences in absorption line sets between those studies and ours, SB1979 used $f_{1133.7}=6.3\\times10^{-3}$ and $f_{1144.9}=0.15$, while JSS1986 adopted $f_{1133.7}=4.8\\times10^{-3}$.  As noted in Table \\ref{linetable}, we used the $f$-values adopted by \\citet{Howk} of $f_{1133.7}=5.5\\times10^{-3}$ and $f_{1144.9}=0.106$.  However, in comparing their results with SB1979, SRF2002 concluded that the SB1979 abundances should be increased (due to the downward revision of $f$-values) by an average of 20\\%, much smaller than the inherent variation in abundances.  Therefore, there is not a clear reason to suspect that $f$-value differences lead to significant systematic errors.  On the other hand, the fact that this study and SRF2002 are based on more absorption lines (7 in most cases) covering a broader range in $f$-value than the studies by SB1979 and JSS1986 (3 lines at most) is one reason to prefer the {\\it FUSE} results; at the least, it justifies including only the SRF2002 results in our analysis.  Our search for the weak 1901.773 \\AA{} line of Fe II and the two probable detections that resulted are discussed in \\S \\ref{ss:1902line}.  A brief comparison of the previous theoretical \\citep{RU1998} and empirical \\citep{Howk} $f$-values is given in \\S \\ref{ss:f-values}.  Finally, we discuss the overall Galactic abundance of iron (and, briefly, the implications for dust models) in \\S \\ref{ss:cosmic_iron}.   \\subsection{Correlations} \\label{ss:FeII_correlations} In this subsection we discuss the correlations and anticorrelations found between iron abundances and depletions and various line of sight parameters (measures of hydrogen content, extinction, reddening, and line of sight pathlength).  As discussed, above, we have included the lines of sight from SRF2002.  We use the Fe II column densities determined in that paper, but in order to calculate abundances and depletions and analyze correlations we find our own values from the literature for the line of sight parameters just mentioned.  This is the reason for any observed discrepancies between our plots and the plots in SRF2002.  We have not included the data points from either SB1979 or JSS1986 for the reasons discussed above.  Correlations between iron depletion and measures of hydrogen (total hydrogen column density, average hydrogen volume density, and the molecular fraction of hydrogen) are discussed in \\S \\ref{sss:FeII_hydro_corr}; correlations between iron depletion and extinction and reddening parameters are discussed in \\S \\ref{sss:FeII_ext_corr}; an anticorrelation between iron depletion and distance (i.e.~line of sight pathlength) is discussed in \\S \\ref{sss:FeII_dist_anticorr}; and lastly, we provide a few comments on outlying data points in \\S \\ref{sss:outliers}.%  \\subsubsection{Correlations with Hydrogen} \\label{sss:FeII_hydro_corr} Iron depletion shows a clear correlation with increased average hydrogen volume density; see Figure \\ref{fig:logFeIIHlognh}.  This is not surprising; this trend was already known for Fe II (JSS1986, SRF2002), and known or at least suggested for several other elements (JSS1986; \\citeauthor{Cartledge2001}~\\citeyear{Cartledge2001}, \\citeyear{Cartledge2004}).  The physical explanation usually invoked for this phenomenon is the model by \\citet{Spitzer1985}, discussed in \\S \\ref{s:FeII_intro}, that measurements of individual lines of sight are probing two fundamental cloud types, warm and cold, and that the relative mix of the two cloud types along the line of sight determines the observed depletions.  The average hydrogen volume density, $\\nHavg$, is usually taken to be the best diagnostic for the determining the mix of the two cloud types being probed.  In general, we have not probed significantly larger $\\nHavg$ than previous studies, with one exception:  HD 147888 has an average hydrogen volume density of 13.6 cm$^{-3}$, more than twice as dense than any of the lines of sight in this sample, SRF2002, or SB1979.  (Note that HD 179406 is also more dense than any line of sight in SRF2002, though only half as dense as HD 147888.)  While 147888 is among the most highly depleted lines of sight in this sample ($\\logFeIIH=-6.87$), it does not show depletion that is particularly extreme; seven other lines of sight have $\\logFeIIH<-6.8$ within their errors.  Although we do not want to overstate the importance of one data point, this line of sight could potentially support the \\citet{Spitzer1985} model at least up to this larger $\\nHavg$.  It would seem that even in this dense line of sight we are not probing clouds with the physical conditions expected of translucent clouds---or, perhaps, models of translucent clouds need to be adjusted to explain the observed lack of extreme depletions.  We also find that iron depletion is correlated with the total hydrogen column density.  \\citet{WakkerMathis} have suggested that $\\NHI$ is in and of itself a good indicator of depletions.  While the correlations of depletion with respect to both $\\NHI$ and $\\nHavg$ are similar, we examined the partial correlation coefficient of the iron abundance and hydrogen volume density while hydrogen column density held constant (an example of using partial correlation coefficients in this manner is outlined in JSS1986).  When we examine the correlation of the logarithmic iron abundance with the logarithm of $\\nHavg$ and $\\logHtot$, the partial correlation coefficient of the iron abundance and $\\nHavg$ is -0.447.  Using a $t$-distribution for the appropriate number of degrees of freedom (in the manner of JSS1986), we find that the probability that the observed correlation coefficient has an absolute magnitude this large if the true correlation coefficient is 0 (we will refer to this as the probability of the null hypothesis) is 0.03\\%.  If we examine the converse, the partial correlation between iron depletion and $\\logHtot$ with $\\lognHavg$ held fixed is still significant: the coefficient is -0.267; the two-sided probability of the null hypothesis is 3.4\\%.  These results imply that while observationally speaking, both $\\lognHavg$ and $\\logHtot$ are very good predictors of abundances and depletions, $\\nHavg$ appears to have more fundamental physical significance.  We note here that $\\logFeIIH$, $\\lognHavg$, and $\\logHtot$ were chosen as opposed to their linear counterparts because these variables show the strongest correlations.  We also note that using these correlation coefficients makes several implicit assumptions, including a linear correlation between the variables being examined, and Gaussian distributions of the variables for the true overall populations.  The validity of these assumptions potentially influences the validity of the above analysis, but we do conclude that $\\nHavg$ is more significant.   SB1979 noted a possible rough correlation between iron depletion and the molecular fraction of hydrogen, $\\fHmol$.  However, SRF2002, including the SB1979 points, did not conclude that there was an overall correlation.  Examining Fig. 5 of SRF2002 reveals that in general, the lines of sight with a large molecular fraction (primarily the {\\it FUSE} lines of sight) have relatively constant depletion, while lines of sight with a smaller molecular fraction (primarily the {\\it Copernicus} lines of sight) show a great deal of scatter in iron depletion (over an order of magnitude in the iron abundance).  Analyzing only our data points and those from SRF2002, we {\\it do} observe an overall correlation between iron depletion and $\\fHmol$; see Figure \\ref{fig:logFeIIHHf}.  With the exception of a few lines of sight, this is in fact one of the strongest correlations seen in our data.  We have briefly explored the possibility that in regions with a large $\\fHmol$ this additional iron is found in the form of gas-phase Fe I, but conclude that the gas-phase Fe I abundance is, in general, at least an order of magnitude too small to account for the difference.   Despite the strength of the overall correlation, there are several outlying points.  Lines of sight with a large $\\fHmol$ but little depletion include HD 210121 (from our sample) and HD 27778 and HD 62542 (from SRF2002), while lines of sight with a small $\\fHmol$ but larger depletions include HD 147888, HD 152236, and HD 164740.  We will return to these outlying data points in \\S \\ref{sss:outliers}.  Ultimately we conclude that while $\\fHmol$ is strongly correlated with iron depletion, it is not a perfect indicator of the physical conditions required for large depletions, and environments that are merely very dense without necessarily forming significant $\\Hmol$ can still show large depletions.  As above, to examine the independent significance of the $\\fHmol$ correlation, we examined the partial correlation coefficients between iron depletion and $\\fHmol$ with $\\nHavg$ held fixed, we conclude that the correlation with $\\fHmol$ is largely secondary to the correlation with $\\nHavg$.  If the three variables considered are $\\logFeIIH$, $\\fHmol$, and $\\lognHavg$, the probability of the null hypothesis is well over 50\\% whether $\\fHmol$ is considered linearly or logarithmically.  If the linear values of the iron abundance or $\\nHavg$ are used, the probability of the null hypothesis $\\nHavg$ decreases (i.e.~that the correlation is more real is likely); however, as before, $\\logFeIIH$ and $\\lognHavg$ are examined because these two variables exhibit a much stronger correlation than the linear variables.  Again, we note the implicit assumption of linear correlations between whatever variables (linear, logarithmic, or otherwise) are under examination.  In the sense that the correlation with $\\fHmol$ is not necessarily independently significant, we have not shown anything new.  However, we have shown that in our sample a stronger correlation exists than in previously studied samples (e.g.~SB1979, JSS1986, SRF2002).  Why the SB1979 sample shows such a large spread in depletion at lower values of $\\fHmol$ that is not seen in our sample is unclear.  We have discussed earlier the potential for systematic errors between {\\it Copernicus} and {\\it FUSE} studies, but even if those errors are larger than we estimate, they do not account for the spread within the self-consistent {\\it Copernicus} data.  Scattered light, which was worse for {\\it Copernicus} than {\\it FUSE}, may be a small factor in causing some variations, but is almost certainly not responsible for the majority of the scatter.  Inherent variation over the short pathlengths of some of the SB1979 targets may also be factor.  It is worth noting that while our results show a few points of scatter at both large and small $\\fHmol$, the SB1979 scatter is exclusively at small $\\fHmol$.  At $\\fHmol \\gtrsim 0.3-0.4$, there are no SB1979 targets with small depletions, but rather the few SB1979 targets with $\\fHmol$ this large more or less conform to the trend seen in our sample.  This compares well with \\citet{Cartledge2006}, where trends of depletion for elements such as magnesium are also seen with respect to $\\fHmol$---though not as strong as between depletion and $\\nHavg$---and with some scatter at small $\\fHmol$.    \\subsubsection{Correlations with Extinction and Reddening Parameters} \\label{sss:FeII_ext_corr} Figure \\ref{fig:logFeIIHred} shows $\\logFeIIH$ as a function of four different reddening or extinction parameters---the relative color excess between the $B$ and $V$ bands, $\\ebv$; $\\ebv$ scaled by line-of-sight pathlength, $\\ebvdist$ the total visual extinction, $\\av$; and $\\av$ scaled by line-of-sight pathlength, $\\avdist$.  Both $\\ebv$ and $\\av$ are measured in magnitudes, and we calculate $\\ebvdist$ and $\\avdist$ in magnitudes per parsec.  The parameter of $\\ebv$ is known to correlate well with $\\NHtot$, and is thought to be strongly correlated with the total dust column density in the line of sight.  Given the correlation we see between iron depletion and $\\NHtot$, it is not surprising that we also see a correlation between iron depletion and $\\ebv$---denser environments are correlated with both iron depletion and total dust content.  SRF2002 noted that the correlation between iron depletion and $\\ebv$ seems to break off at $\\ebv \\approx 0.35 \\magnitude$.  We do not observe this trend.  In our combined sample, we see a very strong correlation overall.  The slope we see for the entire sample is about half as steep as the slope seen in SRF2002 for lines of sight with $\\ebv < 0.35$, but a few times steeper than the slope of the lines of sight with $\\ebv > 0.35$ in that paper.  If we take the same approach as SRF2002 and look for the point where the correlation disappears for larger $\\ebv$, we must restrict the sample to the 23 lines of sight with $\\ebv \\geq 0.5\\magnitude$.  In this sense, we still see the same effect noted in SRF2002, which they interpreted as a threshold effect, wherein the conditions that increase $\\ebv$ do not necessarily require additional iron per hydrogen atom.  The parameter $\\av$ also correlates well with the total hydrogen column density and should also be correlated with the total dust column density.  As discussed in \\S \\ref{s:FeII_intro}, SRF2002 did not find any correlation between iron depletion and $\\av$ for the 18 lines of sight in that paper.  They were not able to compare their 18 lines of sight with the SB1979 or JSS1986 results because those papers did not report on $\\av$ for those lines of sight.  However, taking our sample and the SRF2002 sample, we do find a correlation between iron depletion and $\\av$.  The reason for this difference seems to be the wider range of $\\av$ examined by our sample.  When we restrict the sample to lines of sight with $\\av \\geq 1.5\\magnitude$, the correlation weakens, and it disappears for lines of sight with $\\av \\geq 2\\magnitude$ (at this point, however, it should be noted that the restricted sample contains only 11 lines of sight).  If $\\ebv$ and $\\av$ are correlated with the total dust column density, then dividing by the pathlength of the line of sight results in quantities ($\\ebvdist$ and $\\avdist$) that should be correlated with the total dust volume density.  Not surprisingly, iron depletion is correlated with both of these quantities, at least as strongly as it is correlated with the corresponding extinction parameters integrated over the entire line of sight.  Calculating partial correlation coefficients indicates that the probability of the null hypothesis for a correlation between iron depletion and $\\ebvdist$ with $\\nHavg$ held constant is between 10\\% and 20\\%, for all possible combinations of linear and logarithmic quantities.  The relationship between the iron abundance and $\\avdist$ with $\\nHavg$ held constant is less clear; the probability of the null hypothesis is 42\\% if $\\avdist$ and $\\nHavg$ are logarithmic, but less than 1\\% if they are linear.  Reversing the correlations, we see that iron depletion is always very well-correlated with $\\nHavg$ even when other variables are held constant.  We therefore conclude that the most important variable of interest for determining iron depletion is still $\\nHavg$, but that these measures of dust density ($\\ebvdist$ and $\\avdist$) are also useful quantities that independently correlate with the iron depletion to a small degree.  That the measures of dust density should correlate with iron depletion is also physically motivated in that the depleted material should be found in the dust.  That $\\nHavg$ is still the variable most strongly correlated with iron depletion potentially indicates that the way iron is incorporated into the dust is nonuniform, and/or that the dust grains that contain some iron have variable extinction characteristics.   Dividing these extinction parameters by line of sight pathlength also somewhat resolves the ``threshold'' effects seen for the integrated extinction parameters alone.  Correlations, though both weaker and less significant, remain---rather than disappear---even when only the densest lines of sight are considered.  However, the fact that the correlations do weaken and we do not see extreme depletions indicates that we are not observing lines of sight that are dominated by translucent clouds, though perhaps some lines of sight do probe these clouds to a limited extent.   The quantity $\\rv$ is the ratio of total visual extinction to selective extinction ($\\rv$ is defined as $\\av / \\ebv$) and is correlated with grain size (because larger grains contribute significantly to $\\av$ but not to $\\ebv$).  We find that while there is the hint of a correlation between iron depletion and with $\\rv$ (e.g.~a linear regression of the linear abundance and $\\rv$ has a statistically significant slope), upon further examination the correlation is not particularly significant.  For example, using a Pearson correlation coefficient, the probability of the null hypothesis for logarithmic iron depletion and $\\rv$ is 35\\%.  Rank correlation coefficients (Spearman's $\\rho$ and Kendall $\\tau$, using IDL's R\\_CORRELATE function) that do not carry any implicit assumptions about the functional form of the correlation also show that any correlation is slight at best.  As discussed below in \\S \\ref{sss:outliers}, we also find anecdotal cases where lines of sight that do not conform to some of the observed trends (e.g.~$\\fHmol$) tend to be correlated with $\\rv$ values different from the interstellar average of 3.1; however, we cannot draw the conclusion that there is an overall correlation between depletion and $\\rv$.       Lastly, we will point out a few interesting anecdotal cases regarding cases of extreme depletion and large reddening and extinction.  Given that the overall trends hold, we could point out many cases, but we will restrict ourselves to the three cases with the largest depletions:  from most to least depletion, HD 164740, HD 110432, and HD 147888.  When considering the combined sample, HD 164740 has the largest $\\av$ and second largest $\\ebv$, HD 110432 has the second largest $\\avdist$, and HD 147888 has the largest $\\avdist$ and $\\ebvdist$.  And though we do not claim a conclusive trend between depletion and $\\rv$ overall, we note that these three lines of sight all have large $\\rv$ of $\\gtrsim4$, including HD 164740 with the second largest $\\rv$ in the sample of 5.36.  \\subsubsection{Anticorrelation with Distance} \\label{sss:FeII_dist_anticorr} We find that iron depletion is anticorrelated with distance, i.e.~line of sight pathlength.  With the exception of a few outlying points (HD 210121 from our sample and HD 27778 and HD 62542 from SRF2002), iron depletion decreases out to line of sight pathlengths of about 2 kpc.  For lines of sight with pathlengths from about 2--6 kpc, iron depletion appears constant to within $\\sim0.4\\dex$, significantly less variation than the decrease of an order of magnitude found in lines of sight with pathlengths less than 2 kpc (again, ignoring the three outlying points).  The fact that we see constant depletion for long pathlengths is sensible for two reasons.  First, lines of sight with pathlengths of at least a few kiloparsecs should be sampling the conditions of several clouds, and therefore the total iron abundances and depletions should remain relatively constant.  Secondly, these longer lines of sight are, on average, less dense and less reddened per unit distance than the shorter lines of sight.  However, it cannot be assumed that the lines of sight are truly uniform; slight variations, therefore, may nevertheless contribute to our observed correlations with other line of sight parameters (such as reddening and extinction).    \\subsubsection{Outlying Data Points} \\label{sss:outliers} We have already mentioned six outlying data points.  HD 27778, HD 62542, and HD 210121 have smaller iron depletions but larger values of $\\nHavg$, $\\fHmol$, $\\ebvdist$, and $\\avdist$, breaking the trends that we have just discussed.  Conversely, HD 147888, HD 152236, and HD 164740 have large iron depletions despite small values of $\\fHmol$.  We will first approach the somewhat easier task of interpreting the second set of outliers.  HD 147888, HD 152236, and HD 164740 all have very large hydrogen column and/or volume densities.  Additionally, these lines of sight all have relatively large values of extinction and reddening---HD 147888 has the largest $\\avdist$ and $\\ebvdist$ in the sample, while HD 152236 and HD 164740 both have large pathlength-integrated values of $\\av$ and $\\ebv$ (when scaled by pathlength, HD 152236 is in the upper half of the sample for both parameters and HD 164740 is in the upper third).  Therefore, given the trends we see with respect to all of these parameters, it is not surprising that we see large iron depletions.  To explain the low fractions of $\\Hmol$, we note that all three of these lines of sight have $\\rv$ larger, and in two cases much larger, than the interstellar average of 3.1:  4.06 for HD 147888, 3.29 for HD 152236, and 5.36 for HD 164740.  Large values of $\\rv$ imply a large average grain size, but a large average grain size also implies that the $\\Hmol$ formation rate may be small because of the reduced surface area per unit volume \\citep[see][]{Snow1983}.  Comparing the three lines of sight, HD 152236 has the smallest iron depletion, smallest $\\rv$, and largest $\\fHmol$, while HD 164740 has the greatest iron depletion, largest $\\rv$, and smallest $\\fHmol$, with HD 147888 in the middle in all cases.  These trends fit our interpretation that these lines of sight have smaller molecular fractions of hydrogen due to increased grain size, but otherwise follow the trends between iron depletion and overall gas and dust density.  Additionally, we should note that the increased grain size is unlikely to cause the additional iron depletion, because the reduced average grain size that suppresses $\\Hmol$ formation should also reduce the sticking rate of atoms and ions to grains.  Rather, the increased grain size is likely a result of grain coagulation in environments where iron is already highly depleted.  It is more difficult to interpret the other three lines of sight with relatively low iron depletions despite having large hydrogen volume densities, large molecular fractions of hydrogen, and larger values of reddening and extinction (both per unit length and pathlength integrated).  We first note that all three of these lines of sight have relatively large errors in the iron abundance, so it is possible that these lines of sight are merely statistical fluctuations in the overall correlation.  However, assuming this is not the case, there are a few comments that we can make about these lines of sight.  The atomic hydrogen column density that we report for HD 210121 is based on a 21 cm measurement from \\citet{WF1992}.  If this measurement is saturated, then both the iron abundance and the molecular fraction of hydrogen are somewhat smaller and can be partially reconciled with the correlation.  We note that while the total hydrogen column density derived from adding the $\\Hmol$ measurement of \\citet{Rachford2002} and the $\\NHI$ measurement of \\citet{WF1992} implies $\\logHtot=21.00$, the relationship of \\citet{Bohlin1978}, which correlates $\\ebv$ with $\\Htot$, implies that $\\logHtot=21.37$.  Also worth noting is that if the HD 210121 curve of growth is fit with the equivalent width measurements being artificially weighted equally (i.e.~an unweighted curve-of-growth fit), the result for the iron column density is substantially smaller (by $0.31\\dex$).  Therefore, either an underestimated value of $\\NHtot$ or an overestimate iron column density could be responsible for an iron abundance that seemingly deviates from certain trends.  It is worth noting that the difference between the weighted and unweighted fits is a phenomenon that is generally not observed in other lines of sight.  We have used the Fe II column density for HD 27778 from SRF2002, but we note that \\citet{Miller2007} derived a much smaller Fe II column density (and resulting iron abundance) for this line of sight.  \\citet{Miller2007} used the 2260 \\AA{} line to derive the Fe II column density in HD 27778; we do not have an adequate basis for comparing the use of this line (primarily in terms of its $f$-value) to our curve-of-growth results.  However, we note than an adoption of the \\citet{Miller2007} column density would eliminate this line of sight from consideration as an ``outlier.''  It is worth noting that the abundances of molecules such C$_2$ and CN are relatively large in HD 27778 \\citep{Federman1994}, HD 62542 \\citep{Gredel1993}, and HD 210121 \\citep{Gredel1992}.  The properties of the diffuse interstellar band (DIB) absorption features are somewhat unique in the HD 62542 \\citep{Snow2002, ABM2005} and HD 210121 \\citet{Thorburn2003} lines of sight, with unusually strong ``C$_2$ DIBs'' (DIBs correlated with C$_2$) and unusually weak ``classical'' DIBs (other DIBs not correlated with C$_2$).  All three lines of sight also have $\\rv$ at least slightly less than the interstellar average of 3.1: 2.73 for HD 27778, 2.83 for HD 62542, and 2.07 for HD 210121.  The last point, about the small values of $\\rv$, implies that the average grain size is smaller in these lines of sight.  With a population of small grains, the surface area per unit volume is increased and the rate of $\\Hmol$ formation may increase.  This may be the case for these lines of sight---supported by the fact that of the three lines of sight, HD 210121 has the smallest value of $\\rv$, and presumably the smallest average grain size, but the largest fraction of $\\Hmol$.  Also potentially supporting this interpretation are the large abundances of other molecules in these lines of sight.  However, an important caveat is that the increased surface area per unit volume should also increase sticking of gas-phase atoms to grains.  Therefore, it is still unclear why the iron depletions are so low.  If the smaller average grain size is the result of destruction processes (e.g.~shocks), then the answer may be that the destruction processes may somehow preferentially destroy iron-bearing grains, releasing iron back to the gas phase.  Again, however, it is very unclear why this would be the case, as shocks are thought to destroy grain mantles rather than grain cores, and the major source of depletion for iron should be grain cores, not mantles.  While it is beyond the scope of this paper to explore this issue in detail, we have also examined correlations between $\\rv$, $\\nHavg$, and $\\fHmol$ in our entire sample.  We see, for instance, that $\\fHmol$ increases with $\\nHavg$ up to about $\\nHavg \\approx 2$ cm$^{-3}$, where the correlation begins to break down and, in fact, becomes an anticorrelation (though not as statistically significant as the correlation at low densities).  Therefore, we are possibly seeing evidence of the low $\\Hmol$ formation rate in dense environments.  There is a significant amount of scatter in our attempt to correlate $\\rv$ and $\\fHmol$, with most lines of sight centered near $\\rv=3.1$ with a wide range of $\\fHmol$.  Nevertheless, there is an anticorrelation between $\\rv$ and $\\fHmol$, significant at the 1-$\\sigma$ level, that can be found using several different correlation methods---a simple linear regression, in addition to Pearson, Spearman's $\\rho$, and Kendall $\\tau$ correlation coefficients.  This anticorrelation exists for both the entire range in $\\rv$ and for a narrower range in $\\rv$ near the ISM average of 3.1.  Again, given that $\\rv$ is correlated with grain size, this could be interpreted as limited evidence that the $\\Hmol$ formation rate is lower in environments with larger grains.  However, radiation may also be an important factor; this is briefly discussed below.  While less relevant to the issue of $\\Hmol$ formation, we also note that there is not a statistically significant correlation between $\\nHavg$ and $\\rv$.  However, in spite of the explanation that variations in the correlation between the iron abundance and $\\fHmol$ might be explained by unusual values of $\\rv$, we should also note that unusual values of $\\rv$ do not guarantee a break from this correlation.  For example, HD 38087, with the largest $\\rv$ in our sample, does not break from the trend.  This may be in part due to the fact that while $\\rv$ gives a rough measure of average grain size, it does not give precise information about the details a grain population, nor does it measure total grain surface area (as opposed to surface area per unit grain volume).  For example, while grain coagulation reduces the surface area per unit volume, an increased average grain size due to mantle growth will increase the total grain surface area \\citep{Snow1983}.  An additional consideration is $\\Hmol$-dissociating radiation.  The HD 147888 line of sight passes through the $\\rho$ Oph cloud complex (HD 147888 is $\\rho$ Oph D), which is known to have a high internal ultraviolet radiation field.  This may be generally true regarding lines of sight with large values of $\\rv$---for example, Figure 1 of \\citet{Draine2003}, calculated using the data of \\citet{Fitzpatrick1999}, shows that for the same value of $\\av$, a typical line of sight with $\\rv=4$ has significantly less extinction (about 3 magnitudes) at 1000 \\AA{} than a typical line of sight with $\\rv=3.1$.  Thus, a higher far-UV radiation field allowed by the reduced far-UV extinction almost certainly plays some role in the small values of $\\Hmol$ in some lines of sight with large values of $\\rv$.  In some cases, the far-UV radiation may be the dominant cause.  However, it should be noted that \\citet{Snow1983} nevertheless argues, based on the large number of lines of sight that are both dense and have small values of $\\fHmol$ in the survey of \\citet{ChaffeeWhite}, that a high radiation field is unlikely to be exclusively responsible for these cases.  In conclusion regarding the issue of grain size affecting $\\Hmol$ formation rates, we have shown a few interesting anecdotal cases wherein applying this hypothesis to these outlying points seems to have some explanatory power.  However, this cannot necessarily be considered a uniform effect or the dominant cause of these outlying points.  Making definitive predictions would require further information about the grain population (namely, total grain mass) and the local radiation field.  \\subsection{The 1901.773 \\AA{} Line of Fe II} \\label{ss:1902line} As mentioned in \\S \\ref{ss:FeII_HSTdata}, Fe II has a very weak transition at 1901.773 \\AA{}.  This wavelength region is within the STIS wavelength coverage for 17 lines of sight in the combined sample of this paper and SRF2002 (near the edge of coverage in many lines of sight).  Prior to this paper, there have been no published detections of this line in interstellar absorption due to its small $f$-value \\citep[calculated to be $7.00\\times10^{-5}$ by][]{RU1998}, although weaker lines of Fe II at 2267 \\AA{} \\citep[$f=2.16\\times10^{-5}$;][]{CardelliSavage1995} and 2234 \\AA{} \\citep[$f=1.29\\times10^{-5}$;][]{Miller2007} have been detected.  We examined the available STIS data (see \\S \\ref{ss:FeII_HSTdata} for our methods) in an attempt to detect this feature, and our search resulted in two detections (HD 24534 and HD 93222) and 15 upper limits.  In this section, the detections will be discussed.  Below (\\S \\ref{ss:f-values}) we will discuss the upper limits and the derivation of an $f$-value from these data.  The results are recorded in Table \\ref{eqwidths1902}.      In the spectrum of HD 24534, we find a small feature at 1901.85 \\AA{} with a central depth of $\\sim5\\%$ of the continuum that we identify as the 1901.773 \\AA{} line.  We cite the following pieces of evidence in favor of this identification:  (1) the velocity offset of $\\approx13\\kmpers$  precisely matches that of the dominant velocity component of other lines of Fe II and other elements seen in this line of sight, in particular the 1355.5977 \\AA{} O I and the 1608.4511 line of Fe II (the latter of which is too saturated and too similar in $f$-value to the 1144.9 \\AA{} line to be of additional use in this study); (2) the equivalent width is roughly consistent with the column density measurement, as discussed below; and (3) no other significant ground-state transitions of any abundant elements exist within 1 \\AA{} of the Fe II line.  HD 24534 \\AA{} was studied in SRF2002, who found a Fe II column density of $\\logFeII=14.42^{+0.14}_{-0.13}$ and a $b$-value of $7.5^{+3.3}_{-2.0}\\kmpers$.  Because we detected a feature that appears to be the 1901.773 \\AA{} line of Fe II, we have reanalyzed the {\\it FUSE} spectra to independently determine our own column density and $b$-value, $\\logFeII=14.63\\pm0.06$ and $b=6.4\\pm0.5\\kmpers$.  Though the difference in the column density of $0.21\\dex$ is nontrivial, the column densities are nearly consistent within the 1-$\\sigma$ errors.  The difference likely arises due to the fact that, because of our coadding procedures, we were able to detect the 1127 \\AA{} line of Fe II (the weakest of the {\\it FUSE} lines), which SRF2002 did not include in their analysis of HD 24534, and our equivalent widths are better constrained.  Based on the revised column density that we have derived and the \\citet{RU1998} $f$-value of $7.00\\times10^{-5}$, the expected $\\eqw$ of the 1901.773 \\AA{} line for HD 24534 is 0.96 m\\AA{}.  This is in excellent agreement with our measured equivalent width of $0.94\\pm0.17$ m\\AA{}.  This fit, however, is subject to a possible systematic error.  The area where we identify the feature is slightly asymmetric, and the fit of the line depends on whether we choose a narrow fit that focuses on the portion of the feature where the depth is the greatest or a broader fit that covers the entire feature, including the redward asymmetry.  If we choose the latter fit, we obtain a larger $\\eqw$ of $1.73\\pm0.26$ (consistent with the other fit to within 2-$\\sigma$).  Both fits have a similar reduced $\\chi^2$ of $\\approx 0.73$.  We have selected the narrower fit because of its consistency with the column density, but note that this choice ignores this possible systematic error.  We observe a similar feature in the spectrum of HD 93222 that we identify as this line.  In this spectrum (with a lower S/N than HD 24534), the feature is at a central wavelength of 1901.62 \\AA{} and has a depth of $\\approx10\\%$ of the continuum.  For the same three reasons as before (matching of velocity offset with other lines, rough consistency with derived column density, and lack of other potential identifications), we identify this feature as the 1901.773 \\AA{} line.  The expected $\\eqw$ for HD 93222, based on the column density derived through the {\\it FUSE} lines, is 5.76 m\\AA{}.  However, in HD 93222 we identify at least two marginally resolved cloud components in the FUV absorption lines; the velocity offset of the feature we detect matches the velocity offset of the weaker, narrower component.  However, we have also analyzed the doublet of Mg II found at 1240 \\AA{}, and in that case find three major components---at -23, -6, and 7$\\kmpers$.  The feature we detect and identify as the 1901.773 \\AA{} line is at -23$\\kmpers$ with respect to its rest wavelength; the -23$\\kmpers$ component is the strongest (i.e.~largest column density) component for Mg II.  If we assume that the velocity structure of Mg II and Fe II is similar, then the stronger Fe II component (at $\\sim0\\kmpers$) in the FUV lines is not a single stronger component but a blend of the two weaker components seen for Mg II (-6 and 7$\\kmpers$), and the weaker Fe II component is identified with the strongest Mg II component (both at -23$\\kmpers$). We determined that the -23$\\kmpers$ component contains approximately 43\\% of the total column Mg II density.  Assuming the same distribution for Fe II implies an expected equivalent width of 2.45 m\\AA{} for the -23$\\kmpers$ component, which agrees within the errors with our measured value of $1.95\\pm0.76$ m\\AA{}.  Our simple simulations show that the smaller components (at -6 and 7$\\kmpers$) of the 1901.773 \\AA{} line may be simply lost in the noise.  Both of these comparisons are based on the assumption that the \\citet{RU1998} $f$-value is correct.  Working the above problem in reverse, we have taken our equivalent widths and column densities and calculated an empirical $f$-value in \\S \\ref{ss:f-values}.  Although the value we calculate is consistent with the \\citet{RU1998} value, this line is worth more study to further constrain its $f$-value.  In any case, these detections at the very least roughly confirm the line's $f$-value, which in turn is extremely important for future work with COS, now scheduled to be installed on {\\it HST} in 2008.  COS will have greatly increased sensitivity (depending on wavelength) but somewhat lower resolution than the highest dispersion modes of STIS.  It is possible that this line will be able to reveal Fe II column densities down to at least $\\logFeII\\approx14$ and possibly smaller, without fear of saturation effects.  Detecting such small column densities with COS may not be necessary, however, as many COS targets will have much larger total hydrogen column densities; thus, iron column densities will also be larger unless there are enhanced depletion effects.  If enhanced depletion effects are observed, however, this line may be of utility as our rough measure of its $f$-value is larger than the $f$-value of the weak Fe II lines used by \\citet{Miller2007}.  Therefore, this line may be more easily detected in lines of sight with extreme depletion and/or poor S/N, while still being weak enough to guarantee reasonable freedom from saturation.  \\subsection{Fe II $f$-values} \\label{ss:f-values} \\citet{Morton2003}, which is a definitive compilation of $f$-values and other atomic data for UV lines of astrophysical interest, quotes the $f$-values theoretically calculated by \\citet{RU1998} rather than the more recent empirical $f$-values derived by \\citet{Howk} using {\\it FUSE} data.  However, SRF2002 used the \\citeauthor{Howk}~$f$-values in their study.  Figures \\ref{fig:cogs1-15}-\\ref{fig:cogs46-51} show significant anecdotal evidence that the cases with the smallest errors in the equivalent widths produce extremely self-consistent curves of growth using the \\citet{Howk} $f$-values.  The same holds true, in large part, for the curves of growth in SRF2002.  We have carried through our curve-of-growth method using both the \\citet{Howk} and \\citet{RU1998}.  In 39 of 51 cases, using the \\citet{RU1998} $f$-values produces values of $\\chi^2$ that are larger than the $\\chi^2$ values that result from using the \\citet{Howk} values.  When comparing the 102 cases (51 lines of sight using each set of $f$-values), the 10 poorest fits use the \\citet{RU1998} $f$-values.  It is also worth noting that in several cases the \\citet{RU1998} $f$-values produce a column density of Fe II that is unreasonably large---more than an order of magnitude larger than any of the column densities that are produced using the \\citet{Howk} $f$-values, and producing gas-phase abundances on the order of $10^{-5}$, much larger than has been historically observed.  Though some of the initial calculations using the \\citet{Howk} $f$-values similarly produce very large column densities, an alternate solution (i.e., values in the $\\chi^2$ array that are separated in parameter space from the best solution, but where $\\chi^2$ is below the 1-$\\sigma$ cutoff) always presents itself.  This is not true of the calculations using the \\citet{RU1998} $f$-values, where 15 of 51 lines of sight (including 4 of the 12 cases where the \\citeauthor{RU1998} $f$-values improve $\\chi^2$ compared to using the \\citeauthor{Howk} $f$-values) fail to have a reasonable solution.  In another four of the 12 cases where use of the \\citet{RU1998} $f$-values improves the $\\chi^2$ compared to the \\citet{Howk} $f$-values the improvement is a factor of two or more.  In these cases, however, the reduced $\\chi^2$ values using the \\citet{Howk} $f$-values are $\\lesssim1$, with a maximum of 1.37 for HD 147888.  The other cases (HD 41117, HD 168941, and HD 179406) generally have much larger errors in the equivalent widths (contributing to the small values of $\\chi^2$) and less than the full complement of lines measured (e.g.,~only three lines are measured for HD 41117).  Figure \\ref{fig:cogsfval} shows a sample comparison of the best-fit curves of growth for HD 195965 using the \\citet{Howk} $f$-values and the \\citet{RU1998} $f$-values.  We have selected HD 195965 because the error in the column density using the \\citet{Howk} $f$-values is among the smallest in our sample ($^{+0.03}_{-0.02}\\dex$); in addition, this line of sight shows very little evidence of a multiple-component velocity structure.  Examining Figures \\ref{fig:cogs1-15}-\\ref{fig:cogs46-51}, shows that there are many, many cases where the \\citet{Howk} $f$-values produce self-consistent curves of growth, particularly when the equivalent width errors are small.  It is very unlikely that the curves of growth would be this consistent across a wide range of lines of sight if the \\citet{Howk} $f$-values were not roughly correct.  It is beyond the scope of this paper to present a more formal analysis of the $f$-values than what has been presented above (though we further discuss one specific exception below).  We should note that the \\citet{Howk} and \\citet{RU1998} $f$-values are consistent with each other within their errors.  However, the difference is nontrivial for constructing a curve of growth, as can be seen in Figure \\ref{fig:cogsfval}.  The discrepancy in the two sets of $f$-values for the seven absorption lines in this study ranges from $\\approx5\\%$ to a factor of 2.5; on average, the difference is 30\\%, or $0.11\\dex$.  The consistency of our best curves of growth, however, argues that the \\citet{Howk} $f$-values are much more accurate than this, perhaps to within 0.02 or $0.03\\dex$ (5-7\\%).  However, this is more of an argument for the self-consistency of the \\citet{Howk} $f$-values, rather than a strict constraint on the accuracy of their collective magnitude.  A uniform error in the magnitude of the $f$-values would affect our derived column densities but would not significantly affect the relative correlations discussed in \\S \\ref{ss:FeII_correlations}.  The main exception to our argument for the self-consistency of the \\citet{Howk} $f$-values is the 1112 \\AA{} line.  Using these $f$-values implies that the 1112 \\AA{} line should have a larger $\\eqw$ than the 1133 \\AA{} line by 9\\% in the linear case (where $\\eqw \\propto f\\lambda^2$).  In the 15 lines of sight where SRF2002 measured both lines, $W_{1133} > W_{1112}$ in 11 cases, though the 1-$\\sigma$ errors in $\\eqw$ do overlap in 14 of the 15 cases ($W_{1133} > W_{1112}$ for HD 197512 even when errors are taken into account).  However, the weighted averages tell a different story:  $W_{1133}=33\\pm1$ m\\AA{}, while $W_{1112}=27\\pm1$ m\\AA{}.  We find essentially the same result in our data.  Figures \\ref{fig:cogs1-15}-\\ref{fig:cogs46-51} show that in the cases with the smallest errors in the equivalent widths and the most self-consistent curve of growth, the 1112 \\AA{} line nearly always deviates from the curve of growth.  Our weighted averages are $W_{1133}=49\\pm1$ m\\AA{} and $W_{1112}=46\\pm1$ m\\AA{} (the smaller difference between the two in our sample may be due to the fact that with increased $\\eqw$, the 1133 \\AA{} line is slightly saturated in more cases).  Interestingly, \\citet{RU1998} did find that $f_{1112} < f_{1133}$.  Again, we chose HD 195965 to examine this effect.  If we fit a curve of growth while omitting the 1112 \\AA{} line, we find a slightly revised column density ($\\logFeII=14.87\\pm0.03$ as opposed to $\\logFeII=14.85^{+0.03}_{-0.02}$).  If we then fit the 1112 \\AA{} line to this curve, we find that $f_{1112}=4.61\\times10^{-3}$, in much better agreement with \\citeauthor{RU1998} than \\citeauthor{Howk}~(see Table \\ref{linetable}).  In test cases, revising the $f$-value of the 1112 \\AA{} line does not significantly alter our curve-of-growth results; therefore, we have not reanalyzed our column densities in light of this potential $f$-value revision.  We are also presenting the first published detections in interstellar absorption of the 1901.773 \\AA{} line of Fe II.  Therefore, we can use our results to constrain the $f$-value of this line.  If $\\eqw$ and $\\lambda$ are in \\AA{} and $N$ is in cm$^{-2}$, then the following equation applies: \\begin{equation}\\label{eq:weakeqw} \\eqw=8.85\\times10^{-21}Nf\\lambda^2 \\end{equation} Solving for the $f$-value and substituting in the wavelength of the 1901.773 \\AA{} line, we find that $f=\\eqw/[3.20\\times10^{-14}N]$.  Considering both the errors in $\\eqw$ and $\\logFeII$, our derived $f$-value is $6.9\\pm1.6\\times10^{-5}$ for HD 24534 and $5.5\\pm2.4\\times10^{-5}$ for HD 93222.  The weighted average of these results is $6.5\\pm1.3\\times10^{-5}$.  This agrees, within the errors, with the \\citet{RU1998} theoretical value of $7.00\\times10^{-5}$.  Note, however, the assumptions that went into both equivalent width measurements (the choice between fits for the line in the HD 24534 spectrum, and the distinct cloud components assumed for HD 93222).  At the suggestion of the anonymous referee, we have also taken a different approach to calculate the $f$-value of the 1901.773 \\AA{} line.  In this approach, we attempt to measure $\\eqw$ for the lines even when no line is visibly apparent and there is no feature that will be fit by a program utilizing $\\chi^2$ minimization.  This is done by using other lines as a proxy---the 1239.9523 \\AA{} line of Mg II \\citep[from our work in][]{JensenMgII} in most cases, and the 1608.4511 \\AA{} line of Fe II for 24534.  The fits to these lines are examined for the range over which the profile model is less than 99\\% of the continuum.  Given the saturation level in these lines, this cutoff value should include nearly all of the equivalent width.  This wavelength range is then translated to the same velocity range near the 1901.773 \\AA{} line, and a measurement of $\\eqw$ is made by summing the depth of the data points relative to the assumed continuum.  Fluctuations from the noise are assumed to cancel out to zero, and residual equivalent widths can be potentially observed.  To calculate an error on $\\eqw$, we adapt the formula presented by \\citet{Jenkins1973} for the maximum equivalent width of an undetected absorption line: \\begin{equation}\\label{eq:upperlimits} W_{\\lambda,max}=\\frac{N_{\\sigma} d\\lambda \\sqrt{M}}{\\rm S/N} \\end{equation} In this equation, $W_{\\lambda,max}$ is the upper limit on $\\eqw$, $N_{\\sigma}$ is the number of $\\sigma$ confidence desired, $d\\lambda$ is the wavelength spacing of the pixels, $M$ is the number of consecutive pixels required for detection, and S/N is the signal-to-noise of the local continuum.  We use translate this upper limit on the equivalent width into an error on the equivalent widths that we measure.  In other words, we assume $W_{\\lambda,max}=\\sigma(W_{\\lambda})$ where the $W_{\\lambda,max}$ term is evaluated with $N_{\\sigma}=1$.  Even though the resulting values of $\\eqw$ do not correspond to visually apparent features and are not statistically significant in many cases, the goal of this procedure is to use these measurements to derive an $f$-value for the line.  The $f$-value is calculated in each case using Equation \\ref{eq:weakeqw}.  The error in the $f$-value is carried through using standard error propagation, considering errors in both $\\eqw$ and column density.  The number of pixels used in determining the error from Equation \\ref{eq:upperlimits} is the number of pixels that meet the condition mentioned above, that the range where the model fit is less than 99\\% of the continuum.  Within this alternate method, we approach the problem in two different ways.  The first way is to perform the summation over the entire range where all components of the proxy fit are less than 99\\%, calculating the $f$-value based on the column density derived through the {\\it FUSE} lines.  In this case, the weighted average is $3.1^{+1.2}_{-1.1}\\times10^{-5}$.  The other way is to perform the summation over only the velocity range where the dominant component is less than 99\\% of the continuum.  We then derive the $f$-value using an adjusted column density, using the Fe II column densities of this paper (or SRF2002), but scaled by the relative fraction of the column density found in the dominant component of the proxy fit.  Using this method, the weighted average of the fits is $6.5\\pm1.5\\times10^{-5}$.  The fact that summing over the larger range results in a smaller weighted average shows that either (1) the noise fluctuations over such small ranges do not cancel on average or (2) there is some error in fitting the continuum.  We prefer the method of directly analyzing the lines that are actually observed and fit with statistical significance.  However, the alternate method of summing over the data points is also in rough agreement with the previously calculated value of the lines $f$-value, particularly when we only analyze the dominant component.  Nevertheless, this method depends on calculating a statistically significant average from statistically insignificant individual measurements, and is subject to systematic error of making an assumption about what range over which the summation should be performed.  In either case, as stated in \\S \\ref{ss:1902line}, we have made an important step in roughly confirming the $f$-value of this line.  However, further study to improve these constraints is important so that this line can be used in future studies with COS.  \\subsection{The Cosmic Value of Fe/H} \\label{ss:cosmic_iron} In the past several years, many papers have attempted to summarize the current knowledge of stellar and solar abundances, with implications for the cosmic abundance ``standards'' in the ISM, if consistent standards in fact exist.  Three papers of particular importance for this discussion are \\citet{SnowWitt}, \\citet{SofiaMeyer}, and \\citet{Lodders}.  \\citet{SnowWitt} argued for a cosmic iron abundance (including all iron in both gas and dust) of $\\logFeH=-4.57$, based on an average of field B stars ($\\logFeH=-4.51$), cluster B stars ($\\logFeH=-4.49$), and disk F and G stars ($\\logFeH=-4.74$).  For iron (along with many other elements), \\citeauthor{SnowWitt} concluded that the Sun was substantially overabundant ($\\logFeH=-4.33$), an anomalous data point relative to the true Galactic abundances.  \\citeauthor{SofiaMeyer}, examined B stars (without differentiation between cluster or field stars) and disk F and G stars, and found a weighted average of $\\logFeH=-4.55$ for each sample.  The B star abundances of both \\citeauthor{SnowWitt} and \\citeauthor{SofiaMeyer} are reasonably consistent, but with a significant difference between the F and G star abundances; \\citeauthor{SofiaMeyer} explain this difference as a result of restricting their sample to stars with ages $\\leq 2$ Gyr.  In any case, the final adopted cosmic iron abundance of both papers is consistent to with 0.02$\\dex$, well within the errors.  Contrary to \\citeauthor{SnowWitt}, however, \\citet{SofiaMeyer} argued that the solar abundances may in fact reasonably represent cosmic abundances for many elements, in part due to downward revisions to solar abundances \\citep[mainly][]{Holweger}.   \\citet{Lodders} summarized CI chondritic abundances and solar photospheric abundances, and used both to derive protosolar abundances.  \\citeauthor{Lodders} adopts a protosolar iron abundance of $\\logFeH=-4.46$.  The overall agreement (0.10$\\dex$---25\\%---not including the errors) between all types of measurements (stellar and solar) is the among the best for any of the elements discussed in \\citeauthor{SnowWitt} and \\citeauthor{SofiaMeyer}.  Given this agreement, we can infer the amount of iron in dust fairly confidently.  We adopt a cosmic abundance of iron of $3.1\\pm^{+2}_{-1}\\times10^{-5}$ (a weighted average of both \\citeauthor{SofiaMeyer} measurements and the \\citeauthor{Holweger} solar abundance).   Our weighted average implies that only $\\sim$1\\% of the cosmic abundance in the gas-phase, and our maximum gas-phase abundances are at most approximately $2\\times10^{-6}$.  This places a very tight constraint on iron that can be used in dust models \\citep*[e.g.][]{ZDA2004}.  However, there are reasons to question any cosmic abundance standard, as we have discussed previously \\citep{Snow2000, JensenOI, JensenNI}---many processes, particularly those in stellar formation, could cause abundances in stars to deviate from the abundances of the ISM.  SB1979 found some significant spatial variations in the Fe/H abundance---namely that Scorpius-Ophiuchus lines of sight were substantially more depleted than two lines of sight in Cygnus.  In our sample, we do not note any particularly strong spatial effects.  Most of our lines of sight are in or very near the disk ($|b|<10^{\\circ}$), and within that range there is no hint of a correlation between iron depletion and Galactic $b$.  There are slight trends with Galactic longitude $l$ and overall location; in particular, many of the stars in the Crux/Musca region are somewhat less depleted ($\\logFeH \\approx -6$) than the average of our sample, though this is not true of the more reddened (and nearby) HD 110432 in the SRF2002 sample.  However, this deviation is not substantial compared to the overall scatter in our sample.    Iron abundances and depletions were explored in the past by SB1979 and JSS1986 with {\\it Copernicus} and more recently by SRF2002 with {\\it FUSE}.  We have undertaken a survey that covers a wider range of reddening and extinction than covered in any of these studies.  We find evidence that iron depletion correlates with many line of sight parameters (such as $\\nHavg$, $\\ebv$, $\\ebvdist$, $\\av$, $\\avdist$, and $\\fHmol$).  Some of these correlations have been noted previously while others have not.  The fact that we observe trends that have not been previously observed may be explained by the following:  (1) this is the largest survey of interstellar Fe II yet performed with {\\it FUSE} and, in conjunction with data points from SRF2002, probes the widest range in line of sight parameters in a self-consistent manner; (2) {\\it Copernicus} suffered from scattered light to a greater degree than {\\it FUSE} (of order 10\\% for {\\it Copernicus} compared to only a few percent for {\\it FUSE}), which may be responsible, to a limited degree, for some scatter in the measured iron abundances; (3) many {\\it Copernicus} lines of sight had very short pathlengths, over which intrinsic scatter may be significant; and (4) quite simply, the {\\it Copernicus} lines of sight have not been analyzed in the light of some of these parameters, especially $\\av$.  Correlations between iron depletion and $\\ebv$, $\\ebvdist$, $\\av$, $\\avdist$, and $\\fHmol$ can all be interpreted as related to the correlation between iron depletion and the average line of sight density $\\nHavg$.  While a few of our lines of sight probe slightly larger density and/or extinction than previous studies, we do not see depletions that are particularly extreme relative to previously observed depletions, suggesting that at best our lines of sight only partially probe true translucent clouds, and are instead dominated by lines of sight with integrated ``translucent'' levels of reddening and extinction.  Also of note are two detections of the 1901.773 \\AA{} line of Fe II in the spectra of HD 24534 and HD 93222.  This very weak line is potentially very important for determining iron abundances in a straightforward fashion, without fear of saturation, with the Cosmic Origins Spectrograph.  Further detections are needed to better constrain the $f$-value, but our detections make clear that the previously calculated theoretical value \\citep{RU1998} is at least roughly correct, and the line should be detectable at typical Fe II column densities, even in highly-reddened lines of sight with potentially extreme depletions, with the high S/N expected for COS.  Finally, we briefly discuss the $f$-values of the FUV absorption lines of Fe II used in this study.  Though the theoretical $f$-values of \\citet{RU1998} and the empirical $f$-values of \\citet{Howk} are consistent within the errors, the differences are nontrivial for constructing a curve of growth.  We find that the \\citet{Howk} empirical $f$-values produce self-consistent curves of growth in far more cases than the \\citet{RU1998} $f$-values.  We find an exception in the case of the 1112 \\AA{} line--to obtain self-consistent curves of growth; an $f$-value closer to the \\citet{RU1998} $f$-value is preferred."
    },
    "0710/0710.3067_arXiv.txt": {
        "abstract": "The Parity doublet model containing the SU(2) multiplets including the baryons identified as the chiral partners of the nucleons is applied for neutron star matter. The chiral restoration is analyzed and the maximum mass of the star is calculated. ",
        "introduction": "The doublet parity model assumes, besides the nucleons, the presence of particles with opposite chirality named chiral partners. At high density environment as the one in neutron stars interior, there is the possibility of creating such heavy particles. The main difference between this model and the usual chiral model (Ref.~\\refcite{chi}) is that in this one there is a term of bare mass in the lagrangian density. This mass, called $M_O$ appears with fields that are a mixture of the fields of the particles and their chiral partners in such a way that it does not break chirality (Ref.\\refcite{par}).  It is assumed that the star is in chemical equilibrium and the baryons interact through the mesons $\\sigma$, $\\omega$ and $\\rho$ (included in order to reproduce the high asymmetry between neutrons and protons). Electrons are included to insure charge neutrality. The lagrangian density of the system contains besides the kinetic and the interaction terms an explicit symmetry breaking term in order to reproduce the masses of the pseudo-scalar mesons. The coupling constants of the baryons are adjusted to reproduce the baryonic vacuum masses. In the high-density limit the nucleon and its chiral partner have degenerate masses ($M_0=790 MeV$) as the sigma field goes to zero and chiral symmetry is restored:  \\begin{eqnarray} M^*_\\pm=\\sqrt{\\left[\\frac{(M_{N_+}+M_{N_-})^2}{4}-M_0^2\\right]\\frac{\\sigma^2}{\\sigma_0^2}+M_0^2}\\pm\\frac{M_{N_+}-M_{N_-}}{2}\\frac{\\sigma}{\\sigma_0}. \\end{eqnarray}  A good candidate for the nucleon chiral parter is the N'(1535), but since the identification of chiral partners is still on its first steps we study the case with N'(1200) and N'(1500) for comparison. This variation has drastic consequences on the results. Four different cases first studdied in Ref.~\\refcite{par} are applied to neutron stars. ",
        "conclusions": "With increasing density .i.e. toward the center of the star,  chiral partners begin to appear, reaching a point where they exist at the same  rate as they corresponding particles. The decrease in the scalar condensates signals the restoration of the chiral phase. Depending on the parameters, the phase transition turns out to be a continuous cross-over or of first order for P4.  The maximum mass of the star is higher when the coupling constant $g_4$ is set to zero (P1 and P3) so in order to reproduce the most massive star observed, that has $M=2.1^{+0.4}_{-0.5}M_{\\odot}$ (Ref.~\\refcite{21}), the best option would be the P1 description, because although the P3 predicts a higher maximum mass for the star, its phase transition occurs at a chemical potential bigger than 1700 MeV, too hight compared to predictions."
    },
    "0710/0710.0702_arXiv.txt": {
        "abstract": "We present constraints on violations of Lorentz Invariance based on Lunar Laser Ranging (LLR) data.  LLR measures the Earth-Moon separation by timing the round-trip travel of light between the two bodies, and is currently accurate to a few centimeters (parts in $10^{11}$ of the total distance).  By analyzing archival LLR data under the Standard-Model Extension (SME) framework, we derived six observational constraints on dimensionless SME parameters that describe potential Lorentz-violation.  We found no evidence for Lorentz violation at the $10^{-6}$ to $10^{-11}$ level in these parameters. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.1018_arXiv.txt": {
        "abstract": "We show that near-infrared observations of the red side of the Ly$\\alpha$ line from a single gamma ray burst (GRB) afterglow cannot be used to constrain the global neutral fraction of the intergalactic medium (IGM), $\\bar{x}_H$, at the GRB's redshift to better than $\\delta \\bar{x}_H \\sim 0.3$.  Some GRB sight-lines will encounter more neutral hydrogen than others at fixed $\\bar{x}_H$ owing to the patchiness of reionisation.  GRBs during the epoch of reionisation will often bear no discernible signature of a neutral IGM in their afterglow spectra.  We discuss the constraints on $\\bar{x}_H$ from the $z = 6.3$ burst, GRB050904, and quantify the probability of detecting a neutral IGM using future spectroscopic observations of high-redshift, near-infrared GRB afterglows.  Assuming an observation with signal-to-noise similar to the Subaru FOCAS spectrum of GRB050904 and that the column density distribution of damped Ly$\\alpha$ absorbers is the same as measured at lower redshifts, a GRB from an epoch when $\\bar{x}_H = 0.5$ can be used to detect a partly neutral IGM at $97\\%$ confidence level $\\approx 10$\\% of the time (and, for an observation with three times the sensitivity, $\\approx 30$\\% of the time). ",
        "introduction": "Tomorrow, a gamma ray burst (GRB) may be observed that originates from the death of one of the first stars, during the epoch of reionisation. Despite the great distance to this burst, it will be the brightest gamma ray source on the sky for several tens of seconds, one of the brightest cosmological X-ray sources for hours, and its afterglow will be observable for weeks in the near-infrared (and, for the first few hours, brighter than any $z \\sim 6$ QSO).  Much of the optical and near-infrared light will be obscured by the Ly$\\alpha$ forest, and this obscuration will enable the strongest constraint to date on the neutral hydrogen fraction of the intergalactic medium (IGM) at the burst's redshift.  In fact, such an occurrence may already have been realised. \\citet{haislip06}, \\citet{kawai05}, \\citet{tagliaferri05}, and \\citet{totani06} observed and analysed the optical/near-infrared afterglow of GRB050904, identified to be at $z =6.3$ -- possibly during the reionisation epoch and the GRB with the highest confirmed redshift.  \\citet{totani06} derived the constraint on the global neutral fraction $\\bar{x}_{H} < 0.6$ at $z = 6.3$.  In this paper, we discuss the assumptions that went into their analysis, and we investigate how realistic modelling of reionisation can affect constraints on $\\bar{x}_H$ from GRB050904 and from future $z>6$ GRBs.  The Swift satellite has greatly increased the sample of GRBs with known redshifts in the last two years \\citep{gehrels04}. Future missions such as EXIST \\citep{grindlay06} and JWST \\citep{gardner06} will further enhance our ability to detect high-redshift GRBs and will enable more detailed follow-up studies of their near-infrared afterglows.  Interestingly, approximately one-half of Swift bursts are ``dark bursts'' -- bursts that have detected X-ray afterglows, but that have no measurable optical emission (e.g., \\citealt{filliatre06}).  While it is probable that most dark bursts originate from low-redshift, dust-rich galaxies, a fraction of dark bursts may originate from $z>6$ and are ``dark'' because Ly$\\alpha$ absorption from the high-redshift IGM absorbs the optical emission (e.g., \\citealt{malesani05}).  In addition to their extreme luminosity, there are several other advantages to studying reionisation with GRBs compared to other probes of this epoch.  First, the afterglows of high-redshift GRBs are observed at earlier (brighter) times in the source frame than those at lower redshifts, so the dimming owing to increased luminosity distance is nearly cancelled, and the observed flux is almost independent of redshift \\citep{lamb01, ciardi99}.  Second, unlike the spectra of galaxies and quasars, the intrinsic afterglow spectrum of a GRB is a featureless power-law at the relevant wavelengths, allowing a more precise measurement of absorption owing to a neutral IGM \\citep{barkana04b}.  Finally, since the theoretical expectation is that most of the star formation at $z \\gtrsim 6$ occurs in halos with $m \\sim 10^9 \\; \\Msun$ and because observations at $z \\gtrsim 6$ currently probe only the most massive galaxies and QSOs ($m \\gtrsim 10^{11} ~\\Msun$), high-redshift GRB host galaxies should be less massive than galaxies selected in another manner.  Consequently, GRB host galaxies will sit in smaller HII regions during reionisation (on average) than galaxies selected by different means.  Therefore, GRBs will suffer a larger Ly$\\alpha$ IGM absorption feature.   In this work, we do not concentrate on wavelengths blueward of source-frame Ly$\\alpha$ (in the Ly$\\alpha$ forest) to derive constraints from GRBs.  Any blueward flux indicates the presence of ionised gas at the redshift of the transmission.  However, at high redshifts there is little or no flux in the Ly$\\alpha$ forest, even in ionised regions, owing to the increase in density with increasing redshift, the decrease in the size and in the number of voids, and the decrease in the amplitude of the ionising background (e.g., \\citealt{becker06} and \\citealt{lidz07}).  As a result, it is difficult to distinguish a partly ionised IGM from a fully ionised one with the $z > 6$ forest \\citep{becker06, lidz07}.  Future observations of the $z >6$ Ly$\\alpha$ forest from additional QSOs and GRBs will aid reionisation studies, but is unclear whether such studies will ever provide definitive evidence for neutral pockets in the IGM. In contrast, the shape of the line profile redward of Ly$\\alpha$ \\emph{is} sensitive to a substantially neutral IGM and, therefore, can be used to unambiguously detect reionisation \\citep{miralda98}.  Little is known about the rate of GRBs at $z > 6$.  We assume that the rate of long GRBs traces the massive star formation rate (SFR) for most calculations in this work.\\footnote{We do not consider the other class of GRBs, the ``short'' GRBs, in this study.  These bursts, while still cosmological, are more local than long GRBs and typically do not have a detected afterglow.}  The assumption that the GRB rate traces the massive SFR is supported by observations of lower redshift GRB host galaxies \\citep{bloom02, djorgovski01}.  However, \\citet{kistler07} found that the GRB rate is four times higher at $z \\approx 4$ than if the GRB rate exactly traces the SFR.  Other properties of a galaxy apart from its massive SFR might be correlated with its rate of GRBs.  For example, \\citet{stanek06} found that $z < 0.25$ GRBs -- GRBs that are typically under-luminous -- are preferentially in metal-poor galaxies.  Making the assumption that the GRB rate traces the observed SFR, \\citet{salvaterra08} predicted that SWIFT will be triggered by $1-4$ bursts a year above $z = 6$\\footnote{This rate is larger than the SWIFT rate of $1$ identified $z >6$ bursts in three years, but it is possible that these numbers can be reconciled in light of these dark bursts.} and that the EXIST mission would observe $10-60$ bursts a year.  Other studies have predicted even larger rates \\citep{bromm02, daigne06}.   At $z > 6$, POPIII stars with average masses of $\\sim 100~\\Msun$ may exist.  It is unclear whether the death of a POPIII star can result in a GRB.  \\citet{fryer01} identified a mechanism that might produce GRBs from POPIII stars.  However, once the interstellar metallicity reaches a critical value of $\\sim 10^{-3.5}$ solar in a high-redshift galaxy, POPIII star formation quenches and the normal mode of POPII star formation begins (e.g., \\citealt{mackey03, yoshida04}), and this mode {\\it is} known to produce GRBs.  Most if not all of reionisation likely owes to photons from POPII-like stars (e.g., \\citealt{sokasian04, trac06}).  In Section \\ref{redwing}, we discuss the absorption profiles of a neutral IGM as well as of a damped Ly$\\alpha$ absorber (DLA), and, in Section \\ref{reionisation}, we discuss our simulations of reionisation and their implications for the amount of IGM absorption in a GRB afterglow spectrum.  Section \\ref{fits} quantifies the detectability of a neutral IGM from GRB afterglow spectra, and Section \\ref{GRB050904} describes the constraints on $\\bar{x}_H$ from the $z = 6.3$ burst, GRB050904.  For our calculations, we adopt a cosmology with $\\Omega_m = 0.27$, $\\Omega_\\Lambda = 0.73$, $\\Omega_b = 0.46$, $\\sigma_8 = 0.8$, $n = 1$, and $h = 0.7$, which is consistent with the most recent cosmic microwave background and large scale structure data \\citep{spergel06}.  We express all distances in comoving units unless otherwise noted.  When this project was nearing completion, we learned of a similar effort by \\citet{mesinger07b} and refer the reader there for a complementary discussion. ",
        "conclusions": "GRBs are the most luminous sources at high redshifts, and their smooth power-law spectra are ideal for isolating the effects of absorption owing to a neutral IGM.  However, observations must separate the impact of IGM absorption from that of a DLA to detect a neutral IGM. If no damping wing feature from IGM absorption is detected in the spectrum of a high-redshift GRB, one must be careful to conclude that reionisation is complete at the redshift of interest.  We have shown that there is a wide probability distribution of HII region sizes, and that there is large variation in the distribution of $\\bar{x}_H$ along different sight-lines.  A non-detection of neutral hydrogen from a GRB afterglow spectrum might arise because the GRB host galaxy sits within a large HII region.  If absorption owing to a neutral IGM is detected, it will be impossible to infer $\\bar{x}_H$ from a single GRB to better than $\\delta \\bar{x}_H \\sim 0.3$ because of the patchiness of reionisation.  Assuming an observation with similar sensitivity to the Subaru FOCAS spectrum of GRB050904, that the distribution of DLAs is the same as found at lower redshifts, and that the redshift of the GRB is known, a GRB from a redshift at which $\\bar{x}_H \\approx 0.5$ can be used to detect a partly neutral IGM at $98\\%$ C.L.  $\\approx 10\\%$ of the time (and, for an observation with $3$ times the sensitivity, $\\approx 30\\%$ of the time).  If $\\bar{x}_H < 0.5$, these percentiles for detection are even smaller.  Weaker DLAs enhance the probability of detecting a neutral IGM, but too weak of a DLA may prevent a precise redshift determination, which is essential for tight constraints on $\\bar{x}_H$.  Since the $z = 6.3$ burst GRB050904 has a DLA with $N_{\\rm HI} = 10^{21.6}$ cm$^{-2}$, the absorption on the red side of the line is dominated by the DLA \\citep{totani06}.  While this burst may favour a model with an ionised universe over a neutral universe \\citep{totani06}, a weaker DLA is necessary to be able to constrain $\\bar{x}_H$ in the spectrum of a high-redshift GRB.  This is particularly true if the GRB occurs within a large HII region.  GRB050904 was observed spectroscopically by the Subaru telescope $3.4$ days after the prompt gamma ray emission.  If this afterglow had been observed hours after the burst, the flux would have been more than an order of magnitude larger.  To detect a neutral IGM, it is crucial for optical and near-infrared spectrographs to observe candidate high-redshift GRBs as soon as possible after the prompt gamma ray emission.  A high signal-to-noise ratio is critical to distinguish IGM absorption from that arising from a DLA.  Such an observing programme is worthwhile given the promise that GRBs have as probes of the epoch of reionisation."
    },
    "0710/0710.4950.txt": {
        "abstract": "{} {\\par We study the forced rotation of Titan seen as a rigid body at the equilibrium Cassini state, involving the spin-orbit synchronization.} {\\par We used both the analytical and the numerical ways. We analytically determined the equilibrium positions and the frequencies of the 3 free librations around it, while a numerical integration associated to frequency analysis gave us a more synthetic, complete theory, where the free solution split from the forced one. } {\\par We find a mean obliquity of 2.2 arcmin and the fundamental frequencies of the free librations of about 2.0977, 167.4883, and 306.3360 years. Moreover, we bring out the main role played by Titan's inclination on its rotation, and we suspect a likely resonance involving Titan's wobble.} {} ",
        "introduction": "\\par Since the terrestrial observations of Lemmon et al. (\\cite{Lemmon93}), the rotation of Titan, Saturn's main satellite, has been assumed to be synchronous or nearly synchronous. This has been confirmed by Lemmon et al. (\\cite{Lemmon95}) and by Richardson et al. (\\cite{Richardson04}) with the help of Voyager I images. In this last work, Titan's rotation period is estimated at $15.9458 \\pm 0.0016$ days, whereas its orbital period is  $15.945421 \\pm 0.000005$ days.  \\par The spin-orbit synchronization of a natural satellite is very common in the solar system (such as for the Moon and the Galilean satellites of Jupiter) and is known as a Cassini state. This is an equilibrium state that has probably been reached after a deceleration of the spin of the involved body under dissipative effects, like tides.  \\par Recently, Henrard and Schwanen (\\cite{Henrard04}) have given a 3-dimensional elaborated analytical model of the forced rotation of synchronous triaxial bodies, after studying the librations around the Cassini state. This model has been successfully applied by Henrard on the Galilean satellites Io (\\cite{Henrard05i}) and Europa (\\cite{Henrard05c}), seen as rigid bodies. Such studies require knowing some parameters of the gravitational field of the involved bodies, which cannot be considered as spheres. Another analytical study has been performed for Mercury by D'Hoedt and Lema\u00eetre (\\cite{DHoedt04}), for the case of a $3:2$ spin-orbit resonance.  \\par Since the first fly-bys of Titan by the Cassini spacecraft, we have a first estimation of the useful parameters, more particularly Titan's $J_2$ and $C_{22}$ (Tortora et al. \\cite{Tortora06}), so a similar study of Titan's rotation can be made. In this paper, we propose a study of Titan's forced rotation, where Titan is seen as a rigid body. The originality of this study over Henrard's previous studies is that we use both the analytical and the numerical tools and compare our results. ",
        "conclusions": "\\par This paper offers a first study of Titan's rotation, where Titan is seen as a rigid body. We obtain a quasiperiodic decomposition of the forced solution, which can be split from the free solution in which Titan's obliquity plays an overwhelming role. Moreover, we find good matching between the frequencies of the free librations around the equilibrium, analytically and numerically evaluated, despite a model of circular orbit in the analytical study. However, we find a slight difference in the equilibrium obliquity. Finally, we cannot exclude a resonance between the proper mode $\\Phi_6$ and Titan's wobble.  \\par The next fly-bys of Cassini spacecraft should give us more information on Titan's gravitational field, so we should be able to make a more accurate study on its rotation, that could include direct perturbations on the other Saturnian satellites. These perturbations are supposed to be small (see for instance Henrard \\cite{Henrard04}) and should be negligeable compared to the uncertainties we have on Titan's gravitational parameters. After that, the next step is to consider Titan as a multilayer non-rigid body and to study the consequences of its internal dissipation on the rotation."
    },
    "0710/0710.3192_arXiv.txt": {
        "abstract": "M85 optical transient 2006-1 (M85\\,OT\\,2006-1) is the most luminous member of the small family of V838~Mon-like objects, whose nature is still a mystery. This event took place in the Virgo cluster of galaxies and peaked at an absolute magnitude of $M_{I}\\approx-13$. Here we present Hubble Space Telescope images of M85\\,OT\\,2006-1 and its environment, taken before and after the eruption, along with a spectrum of the host galaxy at the transient location. We find that the progenitor of M85\\,OT\\,2006-1 was not associated with any star forming region. The $g$ and $z$-band absolute magnitudes of the progenitor were fainter than about $-4$ and $-6$~mag, respectively. Therefore, we can set a lower limit of $\\sim50$\\,Myr on the age of the youngest stars at the location of the progenitor that corresponds to a mass of $<7$~M$_{\\odot}$. Previously published line indices suggest that M85 has a mean stellar age of $1.6\\pm0.3$~Gyr. If this mean age is representative of the progenitor of M85\\,OT\\,2006-1, then we can further constrain its mass to be less than $2$~M$_{\\odot}$. We compare the energetics and mass limit derived for the M85\\,OT\\,2006-1 progenitor with those expected from a simple model of violent stellar mergers. Combined with further modeling, these new clues may ultimately reveal the true nature of these puzzling events. ",
        "introduction": "\\label{Introduction}   M85 Optical Transient 2006-1 (M85\\,OT\\,2006-1; J122523.82$+$181056.2) was discovered on 2006 Jan 6 by the Lick observatory supernova search team (Filippenko et al. 2001\\footnote{http://astro.berkeley.edu/$\\sim$bait/kait.html}) as a faint, $V\\sim19.3$~mag transient in the galaxy M85 (NGC\\,4382), which is at a distance of $17.8$~Mpc (Mei et al. 2007). Subsequent spectroscopy, as well as visible light and infra-red (IR) photometry, presented in Kulkarni et al. (2007), showed that M85\\,OT\\,2006-1 has a recession velocity of $880\\pm130$~km~s$^{-1}$, and is therefore associated with M85. Moreover, we showed that the temporal and spectral properties of this object are unlike those of supernovae, novae, or luminous blue variables.  M85\\,OT\\,2006-1 peaked at absolute $I$-band magnitude of about $-13$. The light curve settled into a $\\sim60$~day plateau, followed by a decrease in bolometric luminosity during which the black-body emission peak shifted toward near-IR wavelengths. The early spectrum of M85\\,OT\\,2006-1, obtained six weeks after discovery, resembles that of a $\\sim4600$~K black body, with H$\\alpha$ and H$\\beta$ narrow emission lines (full width at half maximum of $\\sim350\\,$km\\,s$^{-1}$), along with several other unidentified emission lines. Spitzer IR observations obtained about six months after the discovery revealed a $\\sim1000\\,$K black body spectral energy distribution (Rau et al. 2007).  The spectral and temporal properties of this object resemble those of M31-RV (discovered by Rich et al. 1989; e.g., Mould et al. 1990; Bryan \\& Royer 1992), V838~Mon (discovered by Brown 2002; e.g., Kimeswenger et al. 2002; Bond et al. 2003; Corradi \\& Munari 2007), and possibly the less studied object V4332~Sgr (Martini et al. 1999). However, the M85 transient is the most luminous member of the V838~Mon class. The favored model for this emerging class of V838~Mon-like objects (also known as luminous red novae\\footnote{this term was introduced by Kulkarni et al. 2007.}) is that they are the result of stellar mergers (e.g., Soker \\& Tylenda 2006). However, other models have been suggested to explain these objects (e.g., Retter \\& Marom 2003; Lawlor 2005). The nature  of these events, with their energetics lying between the realms of supernovae and novae, remains uncertain.  In this paper, we present Hubble Space Telescope (HST)- Advanced Camera for Surveys (ACS)/Wide Field Camera (WFC) and Near Infrared Camera and Multi-Object Spectrometer (NICMOS) observations, as well as Palomar 5\\,m spectroscopy, of the environment of M85\\,OT\\,2006-1. The observations are used to characterize the environment of the transient and to set a limit on the mass of the progenitor. ",
        "conclusions": "\\label{Disc}   Although several models exist for V838~Mon-like objects (e.g., Soker \\& Tylenda 2003; Lawlor 2005), in the absence of detailed simulations, the nature of these objects remain elusive. A clue to their origin can be derived from their environment, luminosity function and rate. Given that only a small number of these objects are known, and they were found serendipitously in various searches, the luminosity function and rate are not well constrained. However, the fact that at least two events were observed in our Galaxy (i.e., V838~Mon and V4332~Sgr) in the last $\\sim13$ years suggests that they have a higher rate than SNe. We can set a lower limit on their rate, of $0.019\\,$yr$^{-1}\\,$L$_{MW}^{-1}$, at the $95\\%$ confidence level, where L$_{MW}$ is the Milky Way luminosity.    Now we discuss the implications of our observations for a specific model for V838~Mon-like objects. Soker \\& Tylenda (2006) presented a model for violent stellar mergers in which, prior to the merger, the spins and orbital frequencies of the binary star are losing synchronization due to the Darwin instability (e.g., Eggleton \\& Kiseleva-Eggleton 2001). They found that for a given primary mass, the maximal energy production obtained for a binary mass ratio of $\\sim1/50$, is $\\sim2.5\\times10^{-3}GM_{1}^{2}/R_{1}$, where $G$ is the gravitational constant, and $M_{1}$ and $R_{1}$ are the mass and radius of the primary star. Given the upper limit on the progenitor mass, based on the mean stellar age in M85, $<2$~M$_{\\odot}$, and assuming a main-sequence mass-radius relation, $R\\propto M^{0.7}$, the maximum available energy in their model is short by a factor of three in the total energy production, as compared to the radiated energy of M85\\,OT\\,2006-1 in the first two months, $\\sim8\\times10^{46}$~ergs (assuming a distance of $17.8$~Mpc to M85; Mei et al. 2007). Moreover, it is expected that a large fraction of the energy will go into lifting the outer region of the star rather than radiated away. Furthermore, if the primary is an evolved star, then its radius will be larger than the radius of a main sequence star with the same mass, and the extracted energy will be even smaller. This suggests that either more detailed modeling of violent stellar mergers is required, or that this event is not the result of a violent stellar merger. Another possible solution is that the mass of the progenitor is somewhat larger. A larger progenitor mass will still be consistent with our upper limit of $7$~M$_{\\odot}$ which is based on the absence of stars brighter than $I\\sim-6$~mag. For example, according to Soker \\& Tylenda (2006) model, a $7$~M$_{\\odot}$ progenitor can yield $\\sim4$ times more energy than a $2$~M$_{\\odot}$ progenitor and may explain the discrepancy. We note, however, that other kinds of instabilities can lead to stellar mergers (e.g., in triple systems) and that the above comparison is valid only for the specific case discussed by Soker \\& Tylenda (2006).     Existing hydrodynamical simulations of the common envelope phase in stellar mergers (and also star $+$ neutron star mergers) predict that the total dissipated energy is of the order of that observed in V838\\,Mon and M85\\,OT\\,2006-1 (e.g., Taam \\& Bodenheimer 1989; Terman et al. 1995; Terman \\& Taam 1996). Moreover, simulations of the common envelope phase predicts that most of the envelope will be ejected in the equatorial plane (e.g., Taam \\& Ricker 2006). Indeed, in Rau et al. (2007) we reported evidence suggesting that the expansion of M85\\,OT\\,2006-1 is asymmetric. However, more detailed hydrodynamical simulations of the vast parameter space  available for stellar mergers are needed in order to understand these processes and to test if V838\\,Mon-like objects are indeed the results of stellar mergers.  To summarize, we show that, in contrast to V838\\,Mon, but similarly to M31\\,RV, M85\\,OT\\,2006-1 was probably produced by members of an old stellar population ($>1\\,$Gyr), and that its progenitor/s mass was probably $\\ltorder2\\,$M$_{\\odot}$. These constraints narrow down the allowed venue of stellar models for the nature of this event."
    },
    "0710/0710.4584_arXiv.txt": {
        "abstract": "We construct a physically motivated model for predicting the properties of the remnants of gaseous galaxy mergers, given the properties of the progenitors and the orbit.  The model is calibrated using a large suite of SPH merger simulations.  It implements generalized energy conservation while accounting for dissipative energy losses and star formation. The dissipative effects are evaluated from the initial gas fractions and from the orbital parameters via an ``impulse\" parameter, which characterizes the strength of the encounter.  Given the progenitor properties, the model predicts the remnant stellar mass, half-mass radius, and velocity dispersion to an accuracy of 25\\%. The model is valid for both major and minor mergers. We provide an explicit recipe for semi-analytic models of galaxy formation. ",
        "introduction": "\\label{sec:intro}  Major mergers between galaxies are central to the formation and evolution of elliptical galaxies \\citep{TT72,T77,MH94dsc}. The hierarchical buildup of galaxies in the $\\Lambda$CDM cosmology consists of a sequence of mergers, of which a significant fraction are ``major,\" involving progenitors with a mass ratio larger than 1:3. The gravitational interactions in such mergers have a dramatic effect on the dynamics and morphology of the galaxies, in particular turning rotating disks into pressure-supported spheroids. If the progenitors also contain gas, the mergers induce starbursts followed by gas consumption, which leads to aging stellar populations. The modeling of major mergers is therefore a key element in the attempts to confront the broad picture of galaxy formation with detailed observations. This is commonly performed via simulations incorporating Semi-Analytic Models (SAMs), where the complex physical processes are modeled using simplified parametric recipes.  Advanced SAMs are currently attempting the non-trivial task of following the sizes and internal velocities of galaxies. For disk galaxies, sizes are evaluated using the halo virial radii $R_{\\rm vir}$ and spin parameters $\\lambda$ via $R_{\\rm disk} \\simeq \\lambda R_{\\rm vir}$, with some modifications due to the halo density profile (Fall \\& Efstathiou 1980; Mo, Mao \\& White; Bullock, Dekel et al. 2001; Dutton et al. 2006). The sizes of the remnants of gas-poor (``dry\") mergers, where the dominant interaction is gravitational, can be extracted from the properties of the progenitors and the orbital energy by assuming conservation of energy and relaxation to virial equilibrium \\citep{Cole00}. These considerations work well in simulations of relatively dry mergers, quite independently of the details of the orbit.  While dry mergers may dominate in the formation of the most massive galaxies \\citep{Naab06}, the most common mergers are ``wet\" mergers of gaseous galaxies. It is thought that gas processes play an important role in the formation of ellipticals \\citep{RobertsonFP, Dekel06, Ciotti07}.  As demonstrated below, the sizes predicted by dissipationless energy conservation can be off by a factor of a few for gas-rich mergers. Our goal is to construct a more accurate recipe to predict the size and velocity dispersion of the remnant of a wet merger given the properties of the progenitors and the orbital parameters. Cosmological mergers involve a complex mixture of variables (such as orbital parameters, gas fractions, mass ratio and bulge fraction) and physical processes (such as star formation and feedback), all of which can influence the properties of the merger remnants. Given a rich suite of high-resolution SPH merger simulations \\citep{thesis, Cox05}, that span the available parameter space and physical processes, albeit in a rather sparse and nonuniform manner, we seek a model that will properly represent the simulation results and enable an interpolation between them as well as an extrapolation to outside the simulated regime.  For such a recipe to be successful, it should be based on a toy model that grasps the essence of the main physical processes involved in wet mergers. Our intuition is guided by the finding from the simulations that the remnants are more compact when the initial gas fraction is higher and when the first passage involves a stronger tidal impulse, namely when a larger fraction of the orbital energy turns into internal kinetic energy. At a first glance, this may seem surprising, as a system that gains energy is not expected to become more tightly bound, and indeed, the remnants of dry mergers are not very sensitive to the strength of the impulse. This dependence on gas fraction and on the impulse implies that a higher gas fraction and a stronger impulse are associated with a higher degree of dissipation, via shocks, collisions of gas clouds, and induced gas flows toward the centres of the merging systems.  The resultant higher gas densities enhance the energy losses to radiation, leaving behind a more tightly bound remnant.  In parallel, the higher degree of dissipation yields a stronger burst of star formation, which tends to be focused in the central region of the remnant. This understanding is the basis for our proposed recipe, which characterizes the merger by the impulse at first passage, evaluates the associated degree of dissipation and the resultant radiative energy losses and star formation, and accounts for these energy losses in the energy balance. A few free parameters with values of order unity can hopefully compensate for the crude approximations made. These approximations include, for example, an assumption of structural homology between the progenitors and remnant. The physically motivated recipe is then calibrated using the merger simulations, and its success is to be judged by its accuracy in matching the simulated remnant properties.  In \\S \\ref{sec:methods} we describe the simulations used for this study. In \\S \\ref{sec:model} we present the details of our model for predicting remnant properties. In \\S \\ref{sec:Mass} we generalize the model to unequal mass mergers. In \\S \\ref{sec:Caveats} we discuss certain limitations of our model, and in \\S \\ref{sec:Conclusions} we summarize our conclusions. Appendix A discusses the details of our impulse approximation. Appendix B presents an explicit recipe for SAMs. ",
        "conclusions": "\\label{sec:Conclusions}  We have developed a simple toy model for the physical processes involved in wet mergers of galaxies, and have calibrated it using a suite of hydrodynamical merger simulations. This modeling helped us to gain a better understanding of these processes, and provides a practical semi-analytic recipe for predicting post-merger galaxy properties in SAMs.  Crude models of this sort have been used by \\citet{Cole00}, \\citet{Galics03}, and \\citet{Shen03}, but these models did not account for energy losses through dissipative processes, and they have not been calibrated against realistic merger simulations. Using a suite of merger simulations, we have demonstrated the key role of dissipative energy losses in determining the final radii and velocity dispersions of  merger remnants. We found that the dissipative effects depend on the initial orbits of the progenitors.  More violent, lower angular momentum orbits create greater disturbances in the gas disks, which in turn radiate more energy and produce more stars.  This orbital ``violence'' can be parameterized through an impulse approximation for energy exchange between the orbital and internal components during the first close pass of the encounter.  We present a physically-motivated, simulation-calibrated model that is capable of predicting star formation,  central dark matter fraction, remnant radius and remnant velocity dispersion, given the properties of the progenitors and the initial orbital parameters of a merger.  The non-dissipative energy conservation model often predicts radii that are off by a factor of $\\sim 2-3$, and it does not reproduce the spread due to orbital variations. Our model, which accounts for the dissipative energy losses, results in only $\\sim 25\\%$ errors in the predicted radius and velocity dispersion when a wide variety of progenitor types is considered. For a given progenitor type, the error in remnant properties is reduced to $\\sim 10\\%$, indicating that our model correctly captures the variation of remnant properties due to merger orbit.  Since we used the whole available simulation suite to calibrate our model, via a few proportionality constants of order unity, a proper evaluation of the model performance is yet to be pursued using an independent suite of simulations."
    },
    "0710/0710.1197_arXiv.txt": {
        "abstract": "% ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.4590_arXiv.txt": {
        "abstract": "The long, bright gamma-ray burst GRB~070125 was localized by the Interplanetary Network. We present light curves of the prompt gamma-ray emission as observed by Konus-WIND, RHESSI, Suzaku-WAM, and \\textit{Swift}-BAT. We detail the results of joint spectral fits with Konus and RHESSI data. The burst shows moderate hard-to-soft evolution in its multi-peaked emission over a period of about one minute.  The total burst fluence as observed by Konus is $1.79 \\times 10^{-4}$~erg/cm$^2$ (20~keV--10~MeV). Using the spectroscopic redshift $z=1.548$, we find that the burst is consistent with the ``Amati'' $E_{peak,i}-E_{iso}$ correlation.  Assuming a jet opening angle derived from broadband modeling of the burst afterglow, GRB~070125 is a significant outlier to the ``Ghirlanda'' $E_{peak,i}-E_\\gamma$ correlation.  Its collimation-corrected energy release $E_\\gamma = 2.5 \\times 10^{52}$~ergs is the largest yet observed. ",
        "introduction": "The prompt gamma-ray emission of gamma-ray bursts (GRBs) is the most extensively studied aspect of these energetic explosions.  Indeed, for twenty-five years after the discovery of GRBs \\citep{kleb73}, the prompt emission was the only GRB observable available.  With the first afterglow observations at longer wavelengths \\citep{cost97,vanp97}, detailed analysis of burst models became possible. Presently, the \\textit{Swift} satellite is detecting $\\sim$ 100 bursts per year, most with rapid localization and followup.  The exact mechanism which produces the prompt gamma-ray emission, with its characteristic smoothly broken power-law spectrum, has not been definitively established.  Recent efforts to correlate burst observables with the intrinsic burst energetics have increased the importance of detailed spectral fitting for localized bursts \\citep[for a review, see][]{zhan07c}. Some correlations involve the peak spectral energy \\Ep, which is often above the $\\sim$150 keV cutoff of the \\textit{Swift} Burst Alert Telescope (BAT) passband.  Several current observatories are capable of detailed spectral analysis of GRBs over the full range of \\Ep.  Konus-W \\citep{apte95} is a double scintillator instrument on the WIND spacecraft.  The Ramaty High Energy Solar Spectroscopic Imager (RHESSI) is a solar observatory which uses nine germanium detectors to image the Sun at X-ray to gamma-ray energies \\citep{lin02}.  RHESSI's detectors are unshielded and receive emission from astrophysical sources like GRBs.  The Wide-Band All-Sky Monitor (WAM) \\citep{yama05} aboard Suzaku is the large BGO anticoincidence shield for the Suzaku Hard X-Ray Detector.  AGILE \\citep{trav06} and GLAST \\citep{ritz07} will give additional coverage at the energy range of \\Ep\\ and extend spectral coverage for GRBs up to tens of GeV.  In this paper, we present Konus, RHESSI, and Suzaku observations of the bright GRB~070125.  In Section \\ref{sec-obs}, we discuss the observations and the localization of the burst by the IPN.  Section \\ref{sec-lc} contains the burst light curves, and in Section \\ref{sec-fits} we conduct joint spectral fits to the Konus and RHESSI data. ",
        "conclusions": "While GRB~070125 had a large measured prompt gamma-ray fluence, its spectral properties are unremarkable.  The values of the best-fit spectral parameters are similar to those observed for other bright bursts \\citep[e.g.,][]{kane06}, and the spectral evolution observed is similarly common.  The environment of GRB~070125 is unique, however \\citep{cenk08, chan08, updi08}, requiring a broad jet opening angle in broadband afterglow models \\citep{chan08}. After collimation correction, GRB~070125 has the most energetic prompt emission yet observed and is a significant outlier to the correlation between peak energy and $E_{\\gamma}$.  GRB~070125 appears to weaken the claim that the Ghirlanda correlation has low dispersion. GRB~070125 is not a ``recognizable'' outlier to the Ghirlanda relation in the sense of \\citet{ghir07}, as it is highly consistent with the Amati relation. Its jet parameters have been derived from a rich and well-sampled afterglow dataset.  While the circumburst environment of this GRB is unusually dense, this only highlights the assumption of a fairly narrow range of efficiency and density parameters for the majority of GRBs where broadband modeling of the afterglow has not been possible.  The true dispersion of the correlation may in fact be larger.  The physical significance of GRB spectrum--energy correlations has been questioned \\citep[e.g.][]{butl07,butl08}.  In particular, detector trigger thresholds affect burst detection, and more complex selection effects govern the measurement of peak energies, redshifts, and afterglow breaks.  These effects can influence the sample of GRBs with known redshift, $E_{peak,i}$, and $E_{\\gamma}$.  \\citet{ghir08} examined the effect of trigger and spectral analysis thresholds in the \\Ep--fluence plane, finding that the \\textit{Swift}-detected burst sample was truncated by the spectral analysis threshold.  Neither threshold truncated the pre-\\textit{Swift} burst sample.  We were unable to confirm the source of the systematic shift in \\Ep\\ and fluence between the two instruments for this burst. Minor radiation damage was becoming noticeable in RHESSI detector 8 near the time of this work, mostly below the 65 keV cut utilized here. It is also possible that the Monte Carlo simulation of the RHESSI response is less accurate for such extreme off-axis angles, where a greater number of interactions with the cryostat may be expected.  Our previous work had found excellent agreement in all fit parameters for independent RHESSI and Konus spectral fits for GRB 051103 and GRB 050717.  For the short GRB 051103, Konus found $\\Ep = 1920 \\pm 400$~keV and a 20~keV--10~MeV fluence of $4.4 \\pm 0.5 \\times 10^{-5}$~ergs/cm$^2$ \\citep{kgcn051103, fred07}.  A RHESSI fit yielded $\\Ep = 1930 \\pm 340$~keV and 20~keV--10~MeV fluence of $4.5 \\times 10^{-5}$~erg/cm$^2$ \\citep{bellhead06}. \\citet{krim06} found for a cutoff power-law fit to Konus data for GRB 050717 a best fit value of \\Ep $=2101 ^{+1934}_{-830}$~keV.  A RHESSI fit to the same burst found \\Ep $=1550 ^{+510}_{-370}$~keV \\citep{wigg06}. Those bursts had RHESSI off-axis angles of 97 and 110 degrees, respectively.  Joint spectral fits to \\textit{Swift}-BAT and RHESSI data for 25 bursts co-observed by the two instruments between December 2004 and December 2006 indicated that no offset in response normalization was needed for the two instruments \\citep{bellsf08}. However, for two of three bursts occurring during or after December 2006, the RHESSI data showed a significant deficit relative to \\textit{Swift}-BAT. The RHESSI polar angles for all three late bursts were between 90 and 110 degrees. These fits were conducted using only detectors 1 and 7, which do not appear to have radiation damage in background spectra during this interval. Nonetheless, these results suggest that the observed offset in the RHESSI and Konus fit parameters found here is more likely a consequence of increased radiation damage in the RHESSI detectors than a geometric effect or a generic offset in the RHESSI simulations.  Future analysis of archival bursts may help identify the source of any systematic effects present here.  It is clear, however, that joint fits between instruments capable of constraining the full range of \\Ep\\ are valuable in providing the most accurate and precise determination of the fit parameters."
    },
    "0710/0710.4559_arXiv.txt": {
        "abstract": "We present BIMA observations of a 2$\\arcmin$ field in the northeastern spiral arm of M31.  In this region we find six giant molecular clouds that have a mean diameter of 57$\\pm$13 pc, a mean velocity width of 6.5$\\pm$1.2 \\kms, and a mean molecular mass of 3.0 $\\pm$ 1.6 $\\times$ 10$^5$\\Msun.  The peak brightness temperature of these clouds ranges from 1.6--4.2 K.  We compare these clouds to clouds in M33 observed by \\citet{wilson90} using the OVRO millimeter array, and some cloud complexes in the Milky Way observed by \\cite{dame01} using the CfA 1.2m telescope.  In order to properly compare the single dish data to the spatially filtered interferometric data, we project several well-known Milky Way complexes to the distance of Andromeda and simulate their observation with the BIMA interferometer.  We compare the simulated Milky Way clouds with the M31 and M33 data using the same cloud identification and analysis technique and find no significant differences in the cloud properties in all three galaxies.  Thus we conclude that previous claims of differences in the molecular cloud properties between these galaxies may have been due to differences in the choice of cloud identification techniques.  With the upcoming CARMA array, individual molecular clouds may be studied in a variety of nearby galaxies.  With ALMA, comprehensive GMC studies will be feasible at least as far as the Virgo cluster.  With these data, comparative studies of molecular clouds across galactic disks of all types and between different galaxy disks will be possible. Our results emphasize that interferometric observations combined with the use of a consistent cloud identification and analysis technique will be essential for such forthcoming studies that will compare GMCs in the Local Group galaxies to galaxies in the Virgo cluster. ",
        "introduction": "In the inner disk of the Milky Way, the dominant component of the interstellar medium is molecular gas which is distributed in discrete cloud complexes \\citep{scoville87, scoville90}.  Nearly all star formation in our Galaxy is associated with these clouds \\citep{scoville87, blitz93}.  Therefore, understanding the distribution and properties of molecular clouds is a prerequisite for understanding galaxy evolution.  Studies of the molecular gas emission in the Milky Way show that while the cloud mass spectrum extends over several orders of magnitude, a majority of the molecular gas mass is contained in large, massive cloud complexes, which are typically $\\sim$40 pc in diameter and have masses on the order of 1 $\\times$ 10$^5$ \\Msun; these complexes are usually referred to as giant molecular clouds (GMCs)\\footnote{In this paper, we will use the words giant molecular clouds, GMCs and cloud complexes interchangeably.}  and are considered to be the basic organizational unit of the molecular interstellar medium \\citep{scoville90, combes91, young91}.  Compared to the Milky Way, the properties of GMCs in external galaxies are not well known.  The primary reason for this is a lack of high resolution, high sensitivity data which are necessary for spatially resolving individual complexes.  Only in the Magellanic Clouds can single dish telescopes resolve extragalactic GMCs \\citep{israel93,rubio93, fukui01}.  Even in the nearest spirals (e.g., M31 or M33), the angular size of a typical 40 pc GMC is $\\sim$12$\\arcsec$, approximately one half the resolution of the IRAM 30m single dish telescope at 2.6mm, the wavelength of the CO (J=1--0) line emission.  For these galaxies, interferometric observations are necessary to resolve GMCs.  GMCs have been detected using millimeter interferometers in the two nearest spiral galaxies, M33 \\citep{wilson90, engargiola03,rosolowsky07b} and M31 \\citep{vogel87,wilson93,loinard98,rosolowsky07a}. In M33, \\citet{wilson90} observed 19 different fields and identified over 30 GMCs but only 9 GMCs were suitable for a study of cloud properties. The number of GMCs detected in most spiral galaxies has been less than a dozen, presumably due to the small coverage area. This is beginning to change.  Recently a survey of M33 was undertaken by the BIMA array. The angular resolution was coarse ($\\sim$13$\\arcsec$) but adequate to identify nearly 150 GMCs (\\citealt{engargiola03}, see also \\citealt{rosolowsky07b}).  With the upcoming CARMA array, ground-breaking studies like the M33 survey are likely to become routine for external galaxies.  It is important to note that previous studies of GMCs in external galaxies have used different instruments with varying angular and linear resolutions and varying sensitivity.  Moreover, these studies have used different cloud-identification methods to define GMCs and determine their properties.  Yet, in almost all of these studies, the observations have detected discrete molecular features similar to Galactic GMCs. Particularly notable is the consistent linewidth-size relationship in most of these observations.  However, several studies have noted differences between Galactic GMCs and those found in some irregular and dwarf galaxies.  In the irregular galaxies NGC 6822 \\citep{wilson94a} and the SMC \\citep{rubio93}, the clouds are slightly less massive, and perhaps smaller than clouds seen in the Milky Way.  In the LMC, \\citet{cohen88} find that for a given size and line width, clouds are about six times fainter in CO than comparable clouds in the Milky Way. The same is true of the dwarf NGC 1569 where the CO-to-H$_2$ conversion factor is reported to be six times larger than the Galactic value \\citep{taylor99}.  In IC10, the cloud sizes and masses are similar to the Milky Way clouds, but the CO-to-H$_2$ conversion factor may be twice as high as that seen in the Galaxy if the clouds are assumed to be self-gravitating \\citep{wilson91}.  These studies suggest that there are noticeable differences between the Galactic GMCs and those found in irregular or dwarf galaxies, that may be attributed to the different host environments (e.g., metallicity) of these galaxies.  GMCs in the three nearest spiral galaxies, M31, M33 and the Milky Way can be observed at sufficiently high resolution ($\\sim$20 pc) to remove ambiguities from coarse resolution.  Are the GMC properties in these spiral galaxies fundamentally different?  \\citet{wilson90} claim that no GMC larger than 4$\\times$10$^5$ \\Msun is seen in M33, and unlike the Milky Way, 50\\% of the molecular mass in M33 resides in small clouds with masses less than 8$\\times$10$^4$ \\Msun, based on a comparison of interferometric and single dish flux from the CO(J=1--0) line.  On the other hand, analysis of the $^{13}$CO/$^{12}$CO line ratio for a number of M33 clouds suggests that the total mass measured from the $^{12}$CO (J=1-0) line may be overestimated and the actual conversion factor from the CO to H$_2$ mass may be lower in the diffuse gas; in that case, the amount of mass in GMCs would constitute most of the molecular mass in M33 \\citep{wilson94b}.  With coarser resolution data, \\citet{engargiola03} find a steep GMC mass spectrum in M33 and report a characteristic mass of 7 $\\times$ 10$^4$\\Msun. They conclude that molecular clouds in M33 are formed rapidly from the atomic gas and are short lived entities.  On the other hand, in M31 \\citet{vogel87} and \\citet{wilson93} note that the four GMCs they detect are similar to Galactic GMCs. \\citet{loinard99} smooth a section of the Carina arm in the Milky Way to the resolution of single dish M31 maps and conclude that there are general similarities between the distribution of molecular gas in M31 and the Milky Way.  Similar results were also found by \\citet{heyer00} who compared the molecular gas in M31 and the Milky Way using single dish observations; in their study,\\citet{heyer00} find a number of similarities in the molecular gas proeprties between M31 and the Milky Way such as the amount of molecular gas at radii larger than 8 kpc, large arm to interarm contrasts, and similar surface brightness, line widths and spacings of GMCs in spiral arms.  \\citet{rosolowsky07a} has recently completed a large survey of clouds in M31.  He also finds that the GMCs in M31 are similar to those found in the Milky Way. These results depend on comparison of properties derived from datasets which are obtained with a different instrument and analyzed with different methods of cloud identification.  The Milky Way data, for example, come from single dish observations with a signal to noise that is at least 2-4$\\times$ higher than in the M33 and M31 datasets. The datasets also differ in that single-dish telescopes are inherently different than the aperture synthesis maps: the interferometer will miss smooth, extended emission whereas the single-dish data may suffer from beam dilution.  The Milky Way survey also suffers from varying linear resolution.  Finally, the methods for identifying clouds in the different Milky Way studies vary.  These studies (e.g., \\citealt{dame86,sanders85}) have usually projected a typical (l,b,v) cube on one of its axes and used a integrated intensity contour to identify discrete features they call clouds or complexes.  The sizes were measured from the total area of the cloud, assuming that the cloud was spherical \\citep{dame86}, or by measuring the chord across the velocity centroid in a position-velocity diagram \\citep{sanders85}.  In M33, \\citep{wilson90} used a two-tiered selection criterion which required that a cloud be 3$\\sigma$ or more in 2 adjacent channels and that its summed flux be at least 3$\\sigma$ above the noise.  The lack of consistency in these methods and differences in observing techniques are worrisome. A proper comparative study of GMCs in these galaxies requires that the differences described above be eliminated. This is particularly important in view of future telescopes such as CARMA and ALMA which will have the ability to mosaic large regions of galaxies like M31 and M33.  The large scale surveys will have the detail necessary for in-depth comparisons of the molecular ISM in Local Group and more distant galaxies.  These future studies will address the fundamental question of how the molecular ISM varies from galaxy to galaxy.  Knowing this is critical to understanding star formation and galaxy evolution.  In this paper, we make an attempt to do a fair comparison of the molecular emission in the three nearest spiral galaxies, i.e., M31, Milky Way and M33.  We compare the M31 complexes to the simulated Milky Way complexes and previous surveys, using a cloud-identification technique in which we identify complexes using a single integrated intensity contour (\\S \\ref{newres}), and discuss how previous conclusions about M33 would change if the same technique of cloud identification were used (\\S \\ref{just}).  An important addition to this analysis would be the addition of single dish data for the M31 and M33 clouds which would allow for an even more comprehensive comparison.  Of course, one would still need to project the Milky Way clouds to the distance of Andromeda to compare the clouds at the same spatial resolution.  Our method is intended to lay the groundwork for future studies that will compare molecular clouds across galactic environments and across different galaxy types with telescopes such as CARMA, SMA and ALMA.  In \\S \\ref{methods}, we describe two automated cloud identification algorithms (Gaussclump developed by \\citet{stutzki90}, and Clumpfind developed by \\citet{williams94}), and their advantages and disadvantages.  While we do not advocate using these algorithms to describe the characteristics of GMC populations in a galaxy, these algorithms may be used to compare the M31, M33 and the simulated Milky Way data.  We describe the results of this experiment in \\S \\ref{auto} and discuss our overall conclusions in \\S \\ref{conc}. Preliminary results from this work have been published in \\citet{sheth00}. ",
        "conclusions": "\\label{conc} In the northeastern spiral arm of M31, we have identified six distinct, large complexes of molecular gas.  All of these complexes lie along the spiral arm dust lanes, or in one case (Complex E), along a dust spur in between the dust lanes.  These complexes have a mean diameter of 57$\\pm$13 pc, a mean velocity width of 6.5$\\pm$1.2 \\kms, and a mean molecular mass of 3.0$\\pm$1.6 $\\times$ 10$^5$ \\Msun and are indistinguishable from those found in the Milky Way e.g., \\citet{dame86}.  Meaningful comparison of GMCs in different galaxies requires consistent analysis of data taken with different instruments and different cloud identification techniques.  This paper represents an attempt at eliminating such differences and comparing GMCs in M31, M33 and the Milky Way in as consistent a manner as possible. We have simulated three Milky Way complexes at the distance of M31 and observed them with the BIMA array.  The simulations show that interferometers are excellent at recovering molecular complexes in the Local Group galaxies.  We compared the simulated Milky Way complexes to the M31 data using an integrated intensity contour method and found that the complexes thus identified and analyzed fell on the same linewidth size relationship as found by \\citep{dame86}.  The M33 clouds if analyzed in this manner also yield larger clouds than previously stated.  Larger interferometric surveys of these galaxies are necessary to compare the GMC distribution and mass function to that in the Milky Way (e.g., \\citealt{engargiola03, rosolowsky07a}).  Finally, we compared two automated algorithms, Clumpfind and Gaussclump, and found that Gaussclump was better at recovering cloud properties than Clumpfind for low dynamic range data.  Our experiments revealed that such algorithms are resolution dependent and the inherent clumpy nature of molecular emission prevents these methods from identifying clouds larger than one or two resolution elements.  Hence we caution against using these methods to characterize GMC populations in galaxies.  Still, these may be used to quickly compare properties of clouds in data with similar noise and resolution characteristics. Using Gaussclump, we find that the cloud properties in M31, M33 and simulated Milky Way data are indistinguishable."
    },
    "0710/0710.4135_arXiv.txt": {
        "abstract": "We present an analysis of the optical colors of 413 Virgo cluster early-type dwarf galaxies (dEs), based on Sloan Digital Sky Survey imaging data. Our study comprises (1) a comparison of the color-magnitude relation (CMR) of the different dE subclasses that we identified in Paper III of this series, (2) a comparison of the shape of the CMR in low and high-density regions, (3) an analysis of the scatter of the CMR, and (4) an interpretation of the observed colors with ages and metallicities from population synthesis models. We find that the CMRs of nucleated (\\,dE(N)\\,) and non-nucleated dEs (\\,dE(nN)\\,) are significantly different from each other, with similar colors at fainter magnitudes ($\\mr\\gtrsim 17$ mag), but increasingly redder colors of the dE(N)s at brighter magnitudes. We interpret this with older ages and/or higher metallicities of the brighter dE(N)s. The dEs with disk features have similar colors as the dE(N)s and seem to be only slightly younger and/or less metal-rich on average. Furthermore, we find a small but significant dependence of the CMR on local projected galaxy number density, consistently seen in all of $u-r$, $g-r$, and $g-i$, and weakly $i-z$. We deduce that a significant intrinsic color scatter of the CMR is present, even when allowing for a distance spread of our galaxies. No increase of the CMR scatter at fainter magnitudes is observed down to $\\mr \\approx 17$ mag ($\\Mr \\approx -14$ mag). The color residuals, i.e., the offsets of the data points from the linear fit to the CMR, are clearly correlated with each other in all colors for the dE(N)s and for the full dE sample, implying that, at a given magnitude, a galaxy with an older stellar population than average typically also exhibits a higher metallicity than average. Given the observational data for Virgo dEs presented here and in the previous papers of this series, we conclude that there must be at least two different formation channels for early-type dwarfs in order to explain the heterogeneity of this class of galaxy. ",
        "introduction": "\\label{sec:int}  It has long been known that a close correlation exists between the colors and luminosities of early-type galaxies, in the sense that more luminous objects have redder colors \\citep[e.g.,][]{bau59,deV61,fab73,sanvis78a,cal83}. A striking observation is the universality of the color-magnitude relation (CMR): it was found to be equal, within the measurement errors, for E and S0 galaxies within clusters, groups, or the field \\citep{fab73,sanvis78a,sanvis78b,bow92}, leading \\citet{fab73} to state that the colors of elliptical galaxies ``are independent of all physical properties studied other than luminosity''. \\citeauthor{fab73} also showed that a similar relation exists between the strength of spectral absorption features and luminosity, which basically is the spectroscopic analogue to the CMR. She interpreted the CMR as a trend of increasing metallicity with luminosity, which today is still considered to be the primary determinant of the CMR \\citep[e.g.,][]{kod97,cha06}. This can be understood with a higher binding energy per unit mass of gas in more massive galaxies, leading to stronger enrichment of the stellar populations. Recently, \\citet{ber03IV} showed that color seems to correlate even more strongly with velocity dispersion than it does with luminosity. This implies that the CMR itself is most likely just a combination of the relation of luminosity and velocity dispersion \\citep[``Faber-Jackson relation'',][]{fab76} and that of color and velocity dispersion. As \\citet{mat05} point out, the latter might hint at a more fundamental relation, namely between galaxy metallicity and mass.  While this trend includes the dEs, there has been disagreement about whether or not they follow the same CMR as the giant ellipticals. \\citet{deV61} found the dwarfs to be ``systematically bluer'' than the giants (but nevertheless following a CMR), whereas \\citet{cal83} reported a linear CMR over a range of $-15 \\le \\MV \\le -23$ mag. However, his Figure~3 actually suggests that the slope of the CMR might indeed be slightly different for the dEs, in the same way as the results of \\citet{deV61} suggested (i.e., with decreasing magnitude, dwarfs become bluer more rapidly than giants do). \\cite{mat05} found a change in the slope of the Faber-Jackson relation for ``faint early-type'' galaxies --- here, ``faint'' means $-17.3 \\le \\Mb \\le -20.5$ mag, thus reaching only slightly into the dwarf regime. \\citet{der05} found an even larger difference in slope for a sample of 15 dEs, but argued that this is consistent with theoretical models, due to the dynamical response to starburst-induced mass loss, which is stronger for objects of lower mass. Significant constraints to such models could only be provided with a better understanding of dE formation. In contrast to giant ellipticals, formation mechanisms proposed for dEs in clusters are typically not based on an early formation epoch, but rather, on infall and subsequent transformation of late-type galaxies through gas-stripping and tidally enhanced star formation \\citep[e.g.,][]{dav88,moo96,vZe04a,sab05,del07}.  After having established a subdivision scheme of Virgo cluster early-type dwarf (dE) galaxies into subclasses with different shapes and distributions \\citep[Paper III of this series]{p3}, we can now proceed to the next logical step, namely to exploiting the wealth of data provided by the Sloan Digital Sky Survey (SDSS) Data Release 5 \\citep[DR5,][]{sdssdr5} by a multicolor analysis of our sample of 413 dEs. However, the colors of different dEs, or of dEs of different subclasses, can not straightforwardly be compared with each other: the existence of the CMR requires such a comparison to be done either at fixed magnitude or with a correction for magnitude differences. For this reason, we shall explore the  dE colors mainly through an analysis of their CMR.  One would naively expect that, if two given dE subclasses formed through different mechanisms, their resulting CMRs should display differences as well, since the relation of galaxy mass to the properties of its stellar population should depend to some extent on how and when the latter was formed. On the other hand, the apparent universality of the CMR -- mainly defined for giant ellipticals -- would seem to argue against such differences. \\citet{conIII} found considerable scatter of the CMR of dEs in the Perseus cluster at magnitudes $\\Mb \\ge -15$ mag, apparently caused by two different sequences of dwarfs in color-magnitude space. They argue that the early-type dwarfs must have multiple origins --- similar to the conclusion of \\citet{pog01} about an apparent bimodal metallicity distribution of Coma cluster dwarfs that might hint at more than one formation channel. We should be able to test the presence of such a bimodality in more detail, given our ``preparatory work'', namely the separation of dE subclasses that have different shapes and distributions (Paper III). Moreover, the SDSS multicolor data enable us to construct CMRs in more than one color, and also to analyze our galaxies in color-color space, thereby translating colors into ages and metallicities. Since \\citet{rak04} found the dE(nN)s in the Coma and Fornax clusters to be younger and to have a higher metallicity than the dE(N)s, we can perform a similar analysis for our Virgo cluster galaxies, allowing us to test the similarity of dE populations of different clusters. ",
        "conclusions": "\\label{sec:discussion}  We have analyzed the colors of 413 Virgo cluster dEs by constructing color-magnitude relations (CMRs) for different dE subclasses and different local densities, as well as by comparing them to theoretical colors from population synthesis models of \\citet{bc03}. We found significant differences between the CMRs of dE(N)s and dE(nN)s, as well as between the CMRs at low and high local projected densities. The models imply that the brighter dE(nN)s are younger than the dE(N)s and/or have  lower metallicities. The dE(di)s are more similar to the dE(N)s, yet still seem to be slightly younger and/or less metal-rich on average. A significant intrinsic color scatter of the CMR is present. The color residuals about the CMR are correlated between different colors for the dE(N)s, and partly also for the dE(nN)s, such that a galaxy falling on the blue side of the CMR in one color also does so in the other color. We find no increase in the color scatter at fainter magnitudes down to $\\mr \\approx 17$ mag ($\\Mr \\approx -14$ mag).  The dE(N)s are consistent with having a nearly ``perfect'' intrinsic correlation of colors, i.e., if the intrinsic $u-r$ color of a dE(N) lies on the red side of the respective CMR, the same is true in almost all cases for the intrinsic $i-z$ color. This is particularly interesting, since $u-r$ is more sensitive to the age of the stellar population, while $i-z$ is sensitive to metallicity. A simple, straightforward interpretation is that when the stars in a dE are on average older than the typical value at that dE's luminosity, then they are also more metal rich, and vice versa. Assuming a direct correlation between luminosity and galaxy mass, we can speculate that the intrinsic scatter of the CMR could, for a given initial mass, reflect a spread in star formation rate or in the efficiency with which gas is turned into stars, perhaps caused by environmental effects. Neglecting other possible effects, a higher SFR at a given initial (gas) mass would lead to stronger enrichment, i.e.\\ to a higher metallicity, than a lower SFR. The gas would be consumed more rapidly, thus reaching the end of star formation earlier than with a lower SFR, and consequently yielding older stars on average.   \\subsection{Subclass colors and stellar populations} \\label{sec:sub_discuss1}  While we found in Section~\\ref{sec:subclasses} that dE(nN)s and dE(N)s follow different CMRs, we then discovered in Section~\\ref{sec:density} that the CMR depends on environmental density. Since  dE(nN)s and dE(N)s populate different density regimes (Paper III), we should compare the CMR of the dE(nN)s to that of the \\emph{low-density} subsample of dE(N)s: the median density of the latter is 1.18 (in units of the logarithm of the number of galaxies per square degree), and it is 1.20 for the dE(nN)s. In contrast, the median density of the full sample of dE(N)s is 1.37. However, the CMR of the low-density dE(N)s is still significantly different from that of the dE(nN)s in $g-r$ and $g-i$; the probability for a common distribution is $\\le$0.1\\% for the half-light aperture. Likewise, the color difference between dE(N)s and dE(nN)s at the bright reference magnitude is $0.09$ mag in $u-r$ and $0.05$ mag in both $g-r$ and $g-i$, using the half-light aperture. This difference only changes by $0.01$ mag in $u-r$ and $g-i$ and even less in $g-r$ when considering only the low-density dE(N)s. Thus, the colors of the (bright) dE(nN)s do differ from those of the dE(N)s even if we allow for the different sampling in density.  The same test can be done for the dE(di)s: their densities (median value 1.18) are also much more comparable to those of the low-density dE(N)s than to those of the full sample. While the color difference between the dE(di)s and the full dE(N) sample at the bright reference magnitude is only $0.02$ mag in $u-r$, $0.03$ mag in $g-r$, and $0.04$ mag in $g-i$, the statistical comparison of their CMRs yields significant differences (Section~\\ref{sec:subclasses}). This changes when we compare the dE(di)s to only the low-density dE(N) sample. Although the colors of the dE(di)s are still slightly bluer than those of the low-density dE(N)s, none of these differences is statistically significant. Whether or not this could indicate a close relation between dE(di)s and dE(N)s despite their very different shapes will be discussed in Section~\\ref{sec:sub_discuss2}.  \\citet{rak04} determined ages and metallicities for 91 dEs in the Coma and Fornax clusters, based on narrowband photometry. They derived ages above 8 Gyr for the dE(N)s, which they found to be about 5 Gyr older than the dE(nN)s. At least qualitatively, and in a relative sense, this would be consistent with our results. As for the metallicities, \\citeauthor{rak04} found the dE(N)s to have \\emph{lower} metallicities than the dE(nN)s, conjecturing that ``globular clusters and dEN galaxies are primordial and have metallicities set by external constraints such as the enrichment of their formation clouds.'' This is not in agreement with our results for the brighter magnitudes: there, we find the metallicities of the dE(N)s to be either similar (in the inner part of the galaxies, see Figure~\\ref{fig:modelsplusdata}) or higher (Figure~\\ref{fig:modelsplusdata_metal}) than those of the dE(nN)s. Due to the fact that the Virgo cluster is a dynamically less relaxed structure than the Coma and Fornax clusters, it would be particularly interesting to find systematic differences in the stellar content of their galaxy populations. However, it would be premature to conclude that the dE(nN)s and dE(N)s in Virgo behave inversely to those in the other clusters --- the sample of \\citet{rak04} only comprises 10 dE(nN)s of the Coma cluster and 9 dE(nN)s of the Fornax cluster, and moreover, the dE(nN)s show a considerable color scatter in the metallicity-sensitive color $i-z$. Whether or not these issues can account for the differences needs to be analyzed in future work.  The question of whether our fitted CMRs would be consistent with being mainly a luminosity-metallicity relation can be qualitatively addressed  by comparing the respective CMR values at the bright and faint reference magnitude in Figure~\\ref{fig:modelsplusdata_metal}. Overall consistency with  this interpretation is present for all subclasses, but we cannot exclude the additional presence of at least some age differences with magnitude. A general problem for the interpretation of the colors also is that the color values for some data points lie at ``too large'' an age, i.e., they fall above the age of the Universe for our model tracks. We thus need to emphasize again that many simplifications entered the calculation of these tracks, like the fact that the models are calculated at a fixed metallicity, or the rather simple star formation histories that we consider. All our interpretations of observed dE colors are always done within this simplified framework of stellar population models.  \\subsection{The minimum number of formation scenarios} \\label{sec:sub_discuss2}  On the basis of our observational analyses presented here and in the previous papers of this series, we now attempt to answer the question of how many different dE formation mechanisms there must be \\emph{at least} in order to explain the diversity of dE subclasses. To start with, how confident can we be that the (flat) dE(di)s do not belong to the same intrinsic (sub-)class as the (round) dE(N)s? The distribution of intrinsic shapes of the brighter dE(N)s, as deduced in Paper III, is rather broad and includes a significant number of flat objects, even down to axial ratios of 0.3. Most of the range of intrinsic axial ratios of the dE(di)s is thus covered by the range of values of the dE(N)s. Moreover, 73\\% of the dE(di)s are nucleated. While their colors are somewhat different from the full sample of dE(N)s, the difference is not anymore significant when compared only to the low-density dE(N)s, which have the same median density as the dE(di)s (Section~\\ref{sec:sub_discuss1}).  If disk substructure, like spiral arms or bars, could only occur in the flattest dEs, due to, e.g., the kinematical configuration of these objects, we would have automatically selected only intrinsically flat galaxies in our search for disk features (Paper I), and would obviously have found their flattening distribution to be consistent with disk galaxies. The fact that these show no central clustering could then be explained, for example, by the much stronger tidal heating that a galaxy experiences in denser regions of the cluster, leading to an earlier destruction of disk features \\citep[cf.][]{mas05}. Similarly, if dE(N)s and dE(di)s originated from a morphological transformation of infalling late-type spirals through galaxy harassment \\citep{moo96}, one could imagine that the amount of transformation depended on how close the encounters with massive galaxies were that led to it --- and the probability for close(r) encounters is obviously higher in the cluster center, leading to rounder objects without disk features.  Why, then, are there almost no fainter dE(di)s? While we concluded in Paper I that we most likely missed a significant amount of dE(di)s at fainter magnitudes due to our detection limits, we also argued that the true number fraction \\emph{does} decrease when going to fainter objects. As a possible explanation, we can speculate that disk substructure might be more likely to occur in more massive galaxies, possibly connected to the presence of a certain amount of rotational velocity. These qualitative considerations demonstrate that dE(di)s and dE(N)s could, in principle, have formed through the same formation process, keeping our counter of necessary dE formation mechanisms at 1 for the moment.  Could the dE(nN)s be also related to the dE(N)s and be formed by the same process? At least for the bright subsamples, the colors of dE(nN)s and dE(N)s differ significantly, indicating younger average stellar ages of the dE(nN)s, or lower metallicities, or both (Section~\\ref{sec:stelpop}). This still holds true even when the different density distributions of dE(nN)s and dE(N)s are accounted for (Section~\\ref{sec:sub_discuss1}). One might thus conjecture that the dE(nN)s simply formed more recently. The significantly flatter shapes of the bright dE(nN)s could then possibly be explained in the sense that they still need to experience several (further) encounters with massive galaxies, leading to further morphological transformation. One would, though, need to invoke another assumption, namely that the faint dE(nN)s, which are already significantly rounder than the bright dE(nN)s, were much more affected by the first tidal encounters, while the bright dE(nN)s were able to partly preserve their initial shape\\footnote[7]{~However, note that the bright dE(nN)s might not even have experienced any tidal encounters yet, but could have been born as thick, puffy systems \\citep{kau07}.} --- this could possibly be explained by the difference in mass.  Another requirement of this scenario would be that the distribution of dE(nN)s within the cluster, which is not centrally concentrated at all (Paper III), would need to shift towards significantly larger local densities within the next few Gigayears, in order to be similar to the distribution of today's dE(N)s. However, \\citet{conI} derived a two-body relaxation time for the Virgo dEs of much more than a Hubble time. Even violent relaxation, which probably only applies for the case of infalling or merging groups, would take at least a few cluster crossing times $t_{\\rm cr}$, with $t_{\\rm cr}\\approx 1.7$ Gyr for Virgo \\citep{bos06}. It thus seems difficult to reconcile these numbers with the  required dynamical process.  A further, obvious point is that nuclei would need to form soon in the dE(nN)s. Perhaps nuclei are currently being formed in the centers of the dE(bc)s where we are witnessing ongoing star formation (Paper II). Such a scenario would lead to nuclei whose stars were clearly younger than the vast majority of their host galaxies' stars. This would, however, be hardly consistent with the results of \\citet{acsvcs8}, who found \"old to intermediate-age populations\" of dE nuclei, and of \\citet{lotz04}, who measured similar colors of nuclei and dE globular clusters. An alternative would be that most nuclei form through coalescence of globular clusters \\citep[e.g.][]{oh00}. Yet in neither case do we find an explanation for why the ratio of nucleated and non-nucleated dEs should increase strongly with luminosity \\citep{san85b} if the latter were the immediate progenitors of the former.  Taken together, our observational results do not allow to explain dE formation  with less than two different processes. We thus conclude by repeating the statement of \\citet{vZe04a}, ``we caution against single-channel evolutionary scenarios.'' Early-type dwarfs are not a homogeneous class of objects, and we strongly recommend to separately analyze the properties of dEs belonging to different subclasses in any future study of dEs."
    },
    "0710/0710.0699_arXiv.txt": {
        "abstract": "The recent discovery of super-Earths (masses $\\leq$ 10 $M_{\\oplus}$) has initiated a discussion about conditions for habitable worlds. Among these is the mode of convection, which influences a planet's thermal evolution and surface conditions. On Earth, plate tectonics has been proposed as a necessary condition for life. Here we show, that super-Earths will also have plate tectonics. We demonstrate that as planetary mass increases, the shear stress available to overcome resistance to plate motion increases while the plate thickness decreases, thereby enhancing plate weakness. These effects contribute favorably to the subduction of the lithosphere, an essential component of plate tectonics. Moreover, uncertainties in achieving plate tectonics in the one earth-mass regime disappear as mass increases: super-Earths, even if dry, will exhibit plate tectonic behaviour. ",
        "introduction": "Until recently, Earth was the largest terrestrial object known to exist. However, five super-Earth planets (a class defined as having a mass between 1-10 $M_{\\oplus}$- earth-masses) have been detected in the last few years \\citep{Rivera_et_al:2005,OGLE-5.5:2006,Lovis_et_al:2006,Udry_et_al:2007}. The five planets have masses in the 5-10 $M_{\\oplus}$ range, but we do not have information on their sizes and cannot be sure if these are really rocky terrestrial planets. However, their discovery provides some evidence that super-Earths might be common and it is only a matter of chance that our Solar System has none. Some of these planets might be in the 'habitable zone', where the radiation from the star allows for the presence of liquid water, but only their thermal and chemical evolution will determine if they are, in fact, habitable. In turn, their thermal evolution and surface conditions depend on and affect their tectonic regime. Currently, Earth is the only planet where plate tectonics is active. Furthermore, this mode of convection has dominated our planet's geological history, is associated with geochemical cycles and thus, has been proposed as a required mechanism for life on Earth \\citep{Walker_et_al:1981}. Here we address whether or not super-Earths are likely to have plate tectonics or be in a stagnant lid convection like Mercury and Mars. ",
        "conclusions": "In summary, convection is more vigourous in massive terrestrial planets, making their lithospheres thinner and therefore reducing lithospheric strength. Furthermore, they achieve larger stresses owing primarily to larger velocities and therefore can more easily overcome the lithospheric resistance to deformation. Plates may reach negative buoyancy on super-Earths despite their relative younger ages. This scenario is suitable for the failure of the plate and subsequent subduction, which is a necessary step for plate tectonics. Given that Earth's convective state leads to plate tectonics, the more favorable conditions experienced by super-Earths will inevitably lead to plate tectonics. Furthermore, planets of similar mass should have the same potential to exhibit plate tectonics. Conversely, this physics can help explain why small planets like Mars, Mercury and the Moon do not exhibit plate tectonics.  \\subsection{Role of Water} Venus is only slightly smaller than Earth and does not exhibit plate tectonics, although some authors \\citep{Turcotte:1993,Jellenik_water:2005J} have suggested it may have in the past. This observation indicates that the $\\sim$1 earth-mass case falls within a zone of transition between `hard' stagnant lid and mobile plate regimes. In this case, characteristics other than M may be important to the dynamics of the lithosphere. For example, the high surface temperature of Venus might lead to a weak, highly deformable boundary layer that would not support the coherent plate-like behaviour that characterizes oceanic plates on Earth. Moreover, plate strength is relatively large compared to the mantle driving force in the one earth-mass case; yield stresses are on the order of 1-5 GPa for olivine (the representative upper mantle mineral) \\citep{Chen_et_al:1998}, whereas our calculations, in agreement with more detailed models \\citep{Becker_Oconnell:2001}, suggest an underlying driving force of only 10 MPa. Since slip can occur on pre-existing faults at stress values of a few MPa, the existence of plate tectonics in the one earth-mass regime may thus depend crucially on the conditions required to initiate subduction. The presence of water is one possible mechanism to reduce the yield strength of a plate and friction on faults. Experiments show that water reduces the yield strength of olivine by 62\\% when raising the temperature from 25 to 400$^{\\circ}$C at 10 GPa, compared to a drop of 39\\% in dry olivine \\citep{Chen_et_al:1998}. Hence, the hydration level of Venus' mantle, which is 1-2 orders of magnitude lower than on Earth \\citep{Zolotov_et_al:1997}, may make it very difficult for convective forces within this planet to overcome plate resistance.  For larger planets, super-Earths, these issues become less relevant. A wet super-Earth will clearly have enough driving force to sustain subduction.  But, more importantly, the consequences for initiating subduction associated with the hydration of a one earth-mass planet (i.e., a reduction of the yield strength by half) would be similar to a doubling of the mass of the planet (Fig. 2). That is to say, both scenarios would be as likely to initiate and maintain subduction.   \\subsection{Atmospheric Observables} The difference between a Super-Earth with active plate tectonics and one with stagnant lid is in the access of upper mantle material and gasses to the atmosphere. The first case allows several global geochemical cycles to operate, like the CO$_2$ and SO$_2$ ones. For example, cases in our Solar System comprise:  Earth with a CO$_2$ cycle and possibly early Mars with a SO$_2$ cycle \\citep{Halevy_et_al:2007}. Earth has had stable modest levels of atmospheric CO$_2$ (between 160-7000 ppm -- \\citet{Royer_et_al:2001}) in the last 0.5 Gy whereas Venus\u2019 levels stand today at 96\\%. A planet with plate tectonism and a carbonate rock reservoir has an efficient built-in cycle that stabilizes climate at temperatures within the liquid water regime \\citep{Kasting:1996}. A super-Earth that has plate tectonics and weathering capabilities can be expected to have CO$_2$ atmospheric concentrations that would yield temperatures around liquid water. Therefore, evidence against the presence of plate tectonics on an exoplanet would be the detection of high values of CO$_2$ for the age of the star, type of star and orbital distance. An SO$_2$ based atmosphere is also possible and the same reasoning would apply, since the sulfur cycle operates analogously to the carbon cycle. But obviously, more theoretical research is necessary to model the details and predict the right observable signatures.  In conclusion, we show here that as mass increases, the process of subduction, and hence plate tectonics, becomes easier. Therefore, massive super-Earths will very likely exhibit plate tectonics. In the future with TPF by NASA and Darwin by ESA it might be possible to use spectroscopy to identify atmospheric signatures suggesting plate tectonism on these objects. This class of planets offers the possibility of finding Earth analogs and, in particular, make attractive targets in the search for habitable planets."
    },
    "0710/0710.2655_arXiv.txt": {
        "abstract": "The detector material Cadmium Zinc Telluride (CZT) achieves excellent spatial resolution and good energy resolution over a  broad energy range, several keV up to some MeV. Presently, there are two main methods to grow CZT crystals, the Modified High-Pressure Bridgman (MHB) and the High-Pressure Bridgman (HPB) process. The study presented in this paper is based on MHB CZT substrates from the company Orbotech Medical Solutions Ltd. \\cite{Orbotech}. Former studies have shown that high-work-function materials on the cathode side reduce the leakage current and therefore improve the energy resolution at lower energies. None of the studies have emphasized on the anode contact material. Therefore, we present in this paper the result of a detailed study in which for the first time the cathode material was kept constant and the anode material was varied. We used four different anode materials : Indium, Titanium, Chromium and Gold, metals with work-functions between 4.1~eV and 5.1~eV. The detector size was  2.0$\\times$2.0$\\times$0.5~cm$^3$ with 8$\\times$8 pixels and a pitch of 2.46~mm. The best performance was achieved with the low work-function materials Indium and Titanium with energy resolutions of 2.0~keV (at 59~keV) and 1.9~keV (at 122~keV) for Titanium  and 2.1~keV (at 59~keV) and 2.9~keV (at 122~keV) for Indium.  Taking into account the large pixel pitch of 2.46~mm, these resolutions are very competitive in comparison to those achieved with detectors made of material produced with the more expensive conventional HPB method. We present a detailed comparison of our detector response with 3-D simulations. The latter comparisons allow us to determine the  mobility-lifetime-products ($\\mu\\tau$-products) for electrons and holes. Finally, we evaluated the temperature dependency of the detector performance and $\\mu\\tau$-products. For many applications temperature dependence is important, therefore, we extended the scope of our study to temperatures as low as -30$^{\\circ}$C. There are two important results. The breakdown voltage increases with decreasing temperature, and electron mobility-life-time-product decreases by about 30\\% over a range from 20$^{\\circ}$C to -30$^{\\circ}$C. The latter effect causes the energy resolution to deteriorate, but  the concomitantly increasing breakdown voltage makes it possible to increase the applied bias voltage and restore  the full performance. ",
        "introduction": "\\label{Introduction} Cadmium Zinc Telluride (CZT) has emerged as the 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. The high density of CZT ($\\rho\\,\\simeq$~5.76 g/cm$^3$) and high average atomic number ($\\simeq$~50) result in high stopping power and in a large cross section for photoelectric interactions.  The main application of CZT detectors is the detection of photons in the 10~keV to $\\sim$1~MeV energy range. CZT has a major impact in various fields including medical imaging, homeland security applications, and space-borne X-ray and gamma-ray astronomy. To give a few examples of the latter, the Swift mission launched in 2004 carries a wide field of view X-ray telescope for the discovery of gamma-ray bursts (GRBs) in the energy range from 15~keV to 150~keV \\cite{swift}. This Burst Alert Telescope makes use of the coded mask technique to localize GRBs with an accuracy of 1-2~arcmin. It is built out of an array of 32,768 co-planar 2$\\times$4$\\times$4~mm$^3$ CZT detectors covering a total area of 0.5~m$^2$. The proposed EXIST (Energetic X-ray Imaging Survey Telescope) mission \\cite{Grindlay} also uses the coded mask approach. With 15,000 pixelated CZT detectors, each with a volume of 2.0$\\times$2.0$\\times$0.5~cm$^3$ and with 16$\\times$16 pixels, an angular resolution of 5 arcmin will be achieved and the energy range between 10~keV and 600~keV will be covered.\\\\ In this paper, we present a detailed study of CZT grown with the Modified Horizontal Bridgman (MHB) process by the company Orbotech Medical Solutions Ltd. \\cite{Orbotech} (previously Imarad). This growth technique gives uniform substrates at high yields and thus modest costs. One of the disadvantages of the MHB technique is a somewhat lower bulk resistivity of $10^9\\,\\Omega\\,$cm compared to $10^{10}\\,\\Omega\\,$cm obtained with the more conventional High Pressure Bridgman (HPB) process. This results in higher leakage currents and therefore degrades the energy resolution at lower energies ($<$ 100~keV).   Several authors recognized that it is possible to reduce dark currents in MHB detectors by using {\\sl p}-type-intrinsic-{\\sl n}-type (PIN) and metal-semiconductor-metal (MSM) contacts \\cite{Nemirovski:01,Vadawale:04,Nari:98,Nari:00,Nari:02}. Both schemes, PIN and MSM, can indeed reduce the dark currents by factors $>$10 and improve on the energy resolution of the detectors at low ($<$ 100 keV) energies. Encouraged by the result of these authors, we experimented with different cathode and anode materials on several MHB substrates, optimizing for the first time the cathode and anode contacts separately. We discussed a comparison of different  cathode materials in an earlier paper \\cite{Jung:05}.  Gold (Au) was among the metals yielding the best performance in terms of dark current suppression and energy resolution as cathode contact material. With the objective to finalize our studies on contact materials, we present in this paper the results on different anode contact materials with an Au cathode and added a study of the performance at lower temperatures which is especially important for space borne applications. We derived $\\mu\\tau$-products of electrons and holes needed for detector simulations. These reproduce the measured data well, and can be used to further improve detector performance and design.  The paper is structured as follows: \\\\ In section \\ref{sec:detectorfabrication}, origin and some general properties of our substrate are presented. We explain how we used this substrate to fabricate detectors suited for our measurements and give a short description of our experimental setup. Then in section \\ref{sec:char}, we discuss crystal inhomogeneities and present the results of photoluminescence mapping of the Zn content and of infrared transmission microscopy. Subsequently, in section \\ref{sec:contact}, the choice of the contact materials is explained, and the design of the detectors we used for comparison measurements is described in detail. Section \\ref{sec:performance} presents the results of our comparison experiments and a discussion of the influence of anode contact materials on the performance.  In section \\ref{sec:Simulation}, we show that we can reproduce experimental data with computer simulations. For this, $\\mu\\tau$-products of electrons and of holes need to be known. We explain how we measured the $\\mu\\tau$-products of electrons and used simulations to derive the $\\mu\\tau$-products of holes. In section \\ref{sec:T}, we present studies of the temperature dependence on the detector response and the electronic properties of CZT. ",
        "conclusions": "\\label{sec:disc} In this paper, we present the results of a detailed study of the performance of a MHB CZT detector from the company Orbotech Medical Solutions Ltd. \\cite{Orbotech}. \\\\ Using photoluminescence, we mapped the spatial distribution of the zinc content, which showed a variation of  $\\sim$3\\%. This variation does not affect the energy resolution of the detector if information about the location of the interaction is available (e.g. pixelated detectors), because in this case the effect can be corrected for. With infrared transmission microscopy, we could detect some crystal defects and could correlate poor pixel performance with such defects. \\\\ We tested various anode contact materials and found that Ti showed the best performance of 2.0~keV, 1.9~keV and 7.3~keV at 59~keV, 122~keV and 662~keV, respectively. We are using this metal now to contact all our MHB detectors. We showed detailed comparisons of experimentally measured and simulated detector response. After adjusting the hole mobility and lifetime, our simulation gives a good description of the measured data and shows that we have a deep understanding of our data. Therefore, the simulation can be used to optimize the detector design as  the pixel width and pitch.\\\\ We studied the detector properties as a function of temperature. The electron $\\mu\\,\\tau$-product drops by about 30\\%, when cooling the substrate from 20$^{\\circ}$C to -30$^{\\circ}$C with constant bias voltage, and the energy resolution deteriorates accordingly, a behavior that has been noted already for MHB substrates \\cite{Nari:00}. For HPB substrates, a corresponding decrease of the electron $\\mu\\,\\tau$-product has been observed \\cite{Stur:05}. We were able to recover almost all of the room temperature performance by increasing the bias voltage at low temperatures, which is possible, because at lower temperatures the breakdown voltage is increased. Taking everything into account, we conclude, that the overall performance of CZT detectors, produced with the cost-effective MHB process, is comparable to the performance of CZT detectors which are produced with the expensive HPB process. \\\\[2ex] {\\it Acknowledgments:} We thank Uri El Hanany from Orbotech Inc. for several free CZT detectors. We acknowledge S. Komarov, L.~Sobotka, D.~Leopold, and J.~Buckley for helpful discussions. Thanks to electrical engineer P.~Dowkontt, and electrical technician G.~Simburger for their support. This work is supported by NASA under contracts NNG04WC176 and NNG04GD70G, and the NSF/HRD grant no.\\ 0420516 (CREST). The authors at Fisk University gratefully acknowledge financial support from the National Science Foundation through the Fisk University Center for Physics and Chemistry of Materials (CPCoM), Cooperative Agreement CA: HRD-0420516 (CREST program), and from US DOE through the National Nuclear Security Administration (NNSA), Office of Nonproliferation Research and Engineering (NA-22), grant no. DE-FG52-05NA27035."
    },
    "0710/0710.4186_arXiv.txt": {
        "abstract": "{ Existence of GZK neutrinos (ultra high energy neutrinos) have been justified although the flux is very low. % A new method is desired to use a huge mass of a detector medium to detect them. % A fundamental study of radar method was carried out to measure % microwave reflection from electromagnetic energy deposit by X-ray irradiation % in a small rock salt sample. % The reflection rate of $1 \\times 10^{-6}$ was found at the energy deposit of $ 1\\times 10^{19}$ eV % which was proportional to square of the X-ray intensity suggesting the effect to be coherent scattering. % The decay time of the reflection was several seconds. This effect implies a large scale natural rock salt formation could be utilized like % a bubble chamber irradiated by radio wave instead of visible light to detect GZK neutrinos. \\PACS{ {61.80.Cb}{X-ray effects}   \\and {95.55Vj}{Neutrino, pion, and other elementary particle detectors: cosmic ray detectors} } % } % ",
        "introduction": "\\label{intro} GZK (Greisen, Zatsepin and Kuzmin) neutrinos are generated by collision between % ultra high energy (UHE) cosmic rays ($\\geq 4 \\times 10^{19}$ eV) and  cosmic microwave background % Ref.~\\cite{Greisen}. % Both have been observed, then GZK neutrinos ($\\geq 10^{16}$ eV) % have been justified to exist although the flux is very low ($\\sim$ 1 km$^{-2}$ day$^{-1}$). % The flux can be estimated to a certain extent % since flux of UHE cosmic rays and density of cosmic microwave background are known. % GZK neutrinos could become a standard candle of UHE neutrinos for finding out other % conceivable UHE neutrino sources such as active galactic nuclei, topological defects and $\\gamma$-ray bursts etc. % in the universe. % In order to detect GZK neutrinos, a huge mass of a detection medium as large as % 50 Gt (3 $\\times$ 3 $\\times$ 3 km$^3$ for a rock salt case) is needed due to their low flux. % Detecting radio wave from Askryan effect (coherent Cherenkov effect) Ref.~\\cite{Askaryan} is a % promising way to utilize a large mass of natural rock salt Ref.~\\cite{Stanley} or % ice bed in Antarctica % but it needs a lot of bore holes to be installed for radio wave detection antennas.  We had measured attenuation length of natural rock salt samples as well as % synthetic rock salt samples of single crystal by a perturbed cavity resonator method. % We found long attenuation lengths for electric field % more than 200 m at 0.3 - 1 GHz in the samples of natural rock salt % Ref.~\\cite{Chiba} as shown in fig. \\ref{fig:attenuation}. % Data noted as \"Hockley\\_in\\_situ\" in the legend are % in-situ measurements at Hockley salt mine Ref.~\\cite{Hockley}. The long attenuation length allows us to % utilize a rock salt formation as a UHE neutrino detector Ref.~\\cite{Gorham}.  \\begin{figure*} \\includegraphics[width=1.\\textwidth,height=0.68\\textwidth,angle=0]{attenuation.eps} \\caption{Attenuation lengths for electric field are plotted with respect to frequency for % synthetic and natural rock salt. % The upper and the lower straight lines are % fitted to synthetic and natural rock salt samples of Asse salt mine, respectively, % assuming loss tangent is constant with respect to the frequency.} \\label{fig:attenuation}       % \\end{figure*} A basic study of a new detection method has been carried out to measure % radio wave reflection from electromagnetic energy deposit by X-ray irradiation % in a small rock salt sample. % If we could detect the radio wave reflection by an enough reflection rate, a large scale natural rock salt formation would work as % a radio bubble chamber using radio wave instead of visible light to get % 3 dimensional information of the energy deposit of the shower generated  by the UHE neutrino. ",
        "conclusions": "\\label{sec:3} Microwave was reflected at the rate of $10^{-6}$ from X ray irradiated rock salt % in the energy deposit of $1\\times 10^{19}$ eV. % Life time of particles which was responsible to the decay time of % the microwave reflection was several seconds. % It is long enough to employ periodic transmission of radar pulses without triggered by % reception of Askaryan radio wave. % Microwave scattering from the irradiated rock salt was coherent % but the particle species working as the scattering targets is not known. %  Power of radio wave emitted by Askaryan effect   % increases proportionally to the frequency. % At the higher frequency the attenuation length becomes short % as shown in fig. \\ref{fig:attenuation}. % On the contrary radar method is not imposed such a restriction on the frequency. % Strong artificial radar pulses are available for the radar method. % Consequently, we could get long range of detection. %  Times and amplitude from several receiving antennas could give us 3 dimensional % information of the UHE shower with its energy. % In order to utilize a salt dome with a diameter of 3 km and a depth of 3 km (50 Gt), % several bore holes are needed in which the transmitting and receiving antennas are installed. Radar method would have a potential to realize Salt Neutrino Detector to act like a Radio % Bubble Chamber to detect GZK neutrinos. %  If we could get the peak power of the radar considerably larger than 1 GW, % the range becomes long enough to know whether GZK neutrinos exist or not without expensive bore holes. % The antennas would be installed slightly under a floor in an excavated space of a rock salt dome. %   \\begin{acknowledgement} Work is partially supported by a Grant in Aid for Scientific Research for Ministry of Education, Science, % Technology and Sports and Culture of Japan, and Funds of Tokubetsu Kenkyuhi, at Seikei University. % We express deep appreciation to KEK-PF staffs, especially Dr. Kazuyuki Hyodo who extended us his hospitality % throughout this experiment, without him this measurement could not be carried out smoothly. \\end{acknowledgement}"
    },
    "0710/0710.4465_arXiv.txt": {
        "abstract": "We study the effect of the neutron star spin -- kick velocity alignment observed in young radio pulsars on the coalescence rate of binary neutron stars. Two scenarios of the neutron star formation are considered: when the kick is always present and when it is small or absent if a neutron star is formed in a binary system due to electron-capture degenerate core collapse. The effect is shown to be especially strong for large kick amplitudes and tight alignments, reducing the expected galactic rate of binary neutron star coalescences compared to calculations with randomly directed kicks. The spin-kick correlation also leads to a much narrower NS spin-orbit misalignment. ",
        "introduction": "There is an increasing interest of a broad astrophysical community to coalescing binary compact stars as primary sources of gravitational waves for the ground-based gravitational wave observatories. Double neutron stars (DNS), observed as binary pulsars, remain to be the most reliable objects for gravitational wave searches. On-going LIGO science runs \\citep{Abbot07} have already set first experimental upper limits on their galactic rates of a few per year. Astrophysical estimates of the DNS coalescence rate, which are based on the binary pulsar statistics or can be obtained from population synthesis simulations, are model-dependent and vary within more than an order of magnitude around the value $10^{-5}$ per year (see recent reviews \\citealt{PYu,Kalogera&07} and references therein).  The kick velocity imparted to a newborn neutron star is an important phenomenological parameter of the core collapse supernovae and represents one of the major uncertainties in the theory of binary star evolution. The origin of the kicks remains unclear and a number of physical models have been suggested (see, for example, \\citealt{Lai} and references therein). For post-supernova evolution of a binary, both the amplitude of the kick and its space direction are important. The distribution of the kick amplitudes is usually obtained from the analysis of radio pulsar proper motions \\citep{Hobbs}. The direction of kicks (for example, with respect to the spin axis of the neutron star) is more difficult to infer from observations. Recently, several observational clues appeared indicating possible NS spin-kick alignment. A noticeable spin-kick alignment has been inferred from polarization measurements of radio emission of pulsars \\citep{Johnston_ea, Rankin07, Johnston_ea07}, as well as from X-ray observations of pulsar wind nebulae around young pulsars \\citep{Helfand,Kargaltsev}. Implications of these findings to the formation of double pulsars were discussed by \\cite{Wang_ea06}. The possibility and conditions for such an alignment in the model of the kick origin by multiple random kicks during NS formation (proposed by \\citep{Spruit}) were studied by \\cite{Wang_ea07}. The implication of NS kick-spin correlation to the plausible birth-kick scenarios was also discussed by \\cite{Ng&Romani}.  Here we explore the effect of NS spin -- kick correlation on the formation and galactic coalescence rate of double neutron stars (DNS) which are primary targets for modern gravitational wave detectors. We show that the tighter alignment, the smaller is the DNS merging rate with respect to models with random kick orientation. The effect is especially important for large kick amplitudes ($\\sim 400$ km/s). We calculate the spin-orbit misalignment of the components of DNS which can be important for GW data analysis. We also considered a scenario in which no (or insignificant) kick accompanies the formation of a neutron star in binary systems from the main-sequence progenitors in a restricted mass range (8-11 $M_\\odot$ or so) due to electron-capture collapse of O-Ne-Ng degenerate stellar core proposed by \\cite{Podsiadlowski_ea04} and further elaborated by \\cite{vdH04,vdH07}. This hypothesis is phenomenologically based on the existence of long-period Be X-ray binaries with low eccentricities \\citep{Pfahl_ea02}. It is consistent with the evolutionary analysis of double neutron star formation \\citep{vdH07} and has been used in some population synthesis studies of DNS, see for example \\cite{Dewi_ea05, Dewi_ea06}. ",
        "conclusions": "We have shown that the spin-velocity correlation observed in radio pulsars, suggesting the NS spin-kick velocity alignment, may have important implications to GW studies. First, the tight alignment reduces the galactic rate of double neutron star coalescences (especially for large kicks 300-400 km/s) relative to models with random kicks. Second, the spin-kick correlation results in a specific distribution of NS spin -- orbit misalignments. In turn, analysis of the NS spin-orbit misalignments inferred from GW signals during DNS mergings can be potentially used to put independent bounds on the still elusive nature of NS kicks."
    },
    "0710/0710.3590_arXiv.txt": {
        "abstract": "We review recent results of SPH simulations of gravitational instability in gaseous protoplanetary disks, emphasizing the role of thermodynamics in both isolated and binary systems. Contradictory results appeared in the literature regarding disk fragmentation at tens of AU from the central star are likely due to the different treatment of radiation physics as well as reflecting different initial conditions. Further progress on the subject requires extensive comparisons between different codes with the requirement that the same initial conditions are adopted. It is discussed how the local conditions of the disks undergoing fragmentation at $R < 25$ AU in recent SPH simulations are in rough agreement with the prediction of analytical models, with small differences being likely related to the inability of analytical models to account for the dynamics and thermodynamics of three-dimensional spiral shocks. We report that radically different adaptive hydrodynamical codes, SPH and adaptive mesh refinement (AMR), yield very similar results on disk fragmentation at comparable resolution in the simple case of an isothermal equation of state. A high number of refinements in AMR codes is necessary but not sufficient to correctly follow fragmentation, rather an initial resolution of the grid high enough to capture the wavelength of the strongest spiral modes when they are still barely nonlinear is essential. These tests represent a useful benchmark and a starting point for a forthcoming code comparison with realistic radiation physics. ",
        "introduction": "Physical fragmentation in astrophysical systems such as gravitationally unstable protoplanetary disks depends on the competition between gravity and thermal pressure, at least in the limit in which the contribution of magnetic fields is neglected. Knowing whether existing numerical simulations are modeling correctly both gravity and pressure is then of paramount importance in this context. Thermal pressure is affected by the details of cooling and heating, either by radiation, convection or viscosity. The increasingly more sophisticated simulations designed in the last few years have explored the effect of  thermodynamics on disk fragmentation. finding in general that this is less likely than in older models in which the gas was evolved using a locally isothermal equation of state (Boss 2002a,b; Mayer et al. 2002; Rice et al. 2003). Currently, it is debated whether fragmentation into Jupiter-sized clumps is possible, especially at distances less than $50$ AU from the central star (Durisen et al. 2007; Stamatellos \\& Whitworth 2007). It is well understood that the disk must cool on a timescale comparable to the orbital time in order for fragmentation to occur (Rice et al. 2003, Gammie 2001; Johnson \\& Gammi 2003; Clarke et al. 2007). At a few tens of AU the optically thick disk midplane  cools via radiation on a timescale longer than the local orbital time, hence the only chance for disks to cool efficiently is via a non-radiative mechanism. This could be either convection or turbulent diffusion associated with shock bores (Boss 2004; Mayer et al. 2007; Boley \\& Durisen 2006). Finally, analytic calculations that include convection predict no fragmentation at radii $< 50$ AU for disks with masses significantly smaller than the mass of the central star, as it is the case in T Tauri disks (Rafikov 2005, 2007).  Four groups  have implemented a scheme for radiative transfer in three dimensional simulations. Two of them, using, respectively, an SPH and a finite-difference polar grid code, and similar initial conditions, find that fragmentation can happen at  $R < 20$ AU (Boss 2004; Mayer et al. 2006), while the  two other groups use , respectively, an SPH (Stamatellos \\& Whitworth 2007) and a cylindrical grid code (Boley et al. 2006) with nearly identical initial conditions and find that disk fragmentation does not happen at $R < 50$ AU (Stamatellos \\& Whitworth 2007 find that fragmentation is possible at larger radii, $R \\sim 100$ AU). These contradictory results were obtained with four different codes. The two sets of initial conditions were also significantly different, the disks in Mayer et al. (2007) and Boss (2006) having much higher surface densities than those in Boley et al. (2006) and Stramatellos \\& Whitworth (2007), and thus being more prone to fragmentation.  Numerical codes differ not only in the way they implement radiative transfer but also in more basic aspects of the algorithm such as how they compute gravity and how they solve the energy equation. Further progress in elucidating whether disk instability is a  possible formation mechanism for giant planets clearly require that the different codes adopted in this area are compared on identical conditions, first in simple models with a fixed equation of state and then on progressively more sophisticated models with radiative transfer. Here we summarize the results recently obtained with the SPH code GASOLINE (Wadsley, Stadel \\& Quinn 2004) with a scheme for radiative transfer for the case of both isolated and binary systems. Then we discuss the results of the first stage of a code comparison project in which the isothermal disks are evolved with GASOLINE and with one of the most advanced grid codes currently available, the adaptive mesh refinement (AMR) code FLASH (Fryxell et al. 2000). ",
        "conclusions": ""
    },
    "0710/0710.0300_arXiv.txt": {
        "abstract": "{In the last couple of years a population of very massive ($M_\\star>10^{11}$ M$_\\odot$), high-redshift ($z\\ge2$) galaxies has been identified, but its role in galaxy evolution has not yet been fully understood.} {It is necessary to perform a systematic study of high-redshift massive galaxies, in order to determine the shape of the very massive tail of the stellar mass function and determine the epoch of their assembly.} {We selected high-$z$ massive galaxies at 5.8$\\mu$m, in the SWIRE ELAIS-S1 field (1 deg$^2$). Galaxies with the 1.6$\\mu$m stellar peak redshifted into the IRAC bands ($z\\simeq1-3$, called ``IR-peakers'') were identified. Stellar masses were derived by means of spectro-photometric fitting and used to compute the stellar mass function (MF) at $z=1-2$ and $2-3$. A parametric fit to the MF was performed, based on a Bayesian formalism, and the stellar mass density of massive galaxies above $z=2$ determined.} {We present the first systematic study of the very-massive tail of the galaxy stellar mass function at high redshift. A total of 326 sources were selected. The majority of these galaxies have stellar masses in excess of $10^{11}$ M$_\\odot$ and lie at $z>1.5$. The availability of mid-IR data turned out to be a valuable tool to constrain the contribution of young stars to galaxy SEDs, and thus their $M_\\star/L$ ratio. The influence of near-IR data and of the chosen stellar library on the SED fitting are also discussed. The $z=2-3$ stellar mass function between $10^{11}$ and $\\sim10^{12}$ M$_\\odot$ is probed with unprecedented detail. A significant evolution is found not only for galaxies with $M\\sim10^{11}$ M$_\\odot$, but also in the highest mass bins considered. The comoving number density of these galaxies was lower by more than a factor of 10 at $z=2-3$, with respect to the local estimate. SWIRE 5.8$\\mu$m peakers more massive than $1.6 \\times10^{11}$ M$_\\odot$ provide 30$-$50\\% of the total stellar mass density in galaxies at $z=2-3$. } {} ",
        "introduction": "\\label{sect:intro}    Tracing the formation of galaxies and understanding the epoch {\\em when} the bulk of their baryonic mass was assembled represents one of the major problems of modern cosmology, particularly controversial when dealing with massive ($M_{\\textrm{stars}} > 10^{11}$ M$_\\odot$) objects.  The assembly of massive galaxies is one of the critical questions in the cosmic evolutionary scenario. The uniform properties of local early-type galaxies and of the fundamental plane have inspired the so called ``monolithic collapse'' scenario \\citep{eggen1962,chiosi2002}, in which galaxies formed in the remote past through huge events of star formation and subsequently evolved passively across cosmic time. On the other hand, in the more recent ``hierarchical'' scenario \\citep{white1978,kauffmann1996,kauffmann1998,somerville1999}, massive galaxies assemble by mergers of lower-mass units, with the most massive objects being born in the latest stages of evolution, at $z\\le1$.   The availability of several powerful tracers of star formation (e.g. UV continuum, optical recombination lines, far-IR emission, sub-mm light) has favored the popularity of studies of the comoving star formation density \\citep{madau1996,lilly1996} in galaxies at various redshifts. It is now well determined that the Universe experienced an epoch of enhanced star formation in the past, peaking at $z\\simeq1-2$, with a subsequent decline of at least one order of magnitude to the present time \\citep[e.g.][ among others]{hopkins2004,rudnick2003,flores1999,madau1996}.  An alternative approach consists of studying the mass already assembled in galaxies, instead of the amount of stars being formed. The integral of the past star formation density provides the stellar mass density at a given epoch: a complementary constraint on cosmic galaxy evolution.  Thus, the build up of the stellar mass across cosmic time has become one of the major topics in observational cosmology, and has overtaken the classic Madau-Lilly diagram as the central tool for studying galaxy evolution. The large observational effort dedicated to this subject has shown that the global stellar mass density increases from early epochs to the low-redshift Universe \\citep[e.g.][]{brinchmann2000,dickinson2003,fontana2003, fontana2004,fontana2006,rudnick2003,rudnick2006,drory2005}. Very deep surveys have been exploited to describe the shape of the stellar mass function at high redshift \\citep{fontana2006,drory2005,gwyn2005}, but a clear picture on the role of very massive galaxies has not yet emerged.   Several pieces of evidence exist that fully formed massive galaxies were already in place at redshift $z\\sim 2-3$.  A substantial population of luminous red galaxies at redshifts $z>2$ (known as ``distant red galaxies'', DRGs) was found in the Faint InfraRed Extragalactic Survey \\citep[FIRES,][]{franx2003}. Based on near-IR  spectroscopy \\citep{foerster2004,vandokkum2004}, these galaxies turned out to be massive ($M=1-5\\times10^{11}$ M$_\\odot$), evolved (ages of $1-2.5$ Gyr) systems, probably descendants of galaxies which started forming at redshift $z>4$. Based on FIRES data, \\citet{rudnick2003} inferred that DRGs contribute $\\sim50$\\% of the global stellar mass density at $z=2-3$.   \\citet{dickinson2003} exploited Hubble Deep Field North NICMOS data to derive the stellar masses of galaxies up to $z=3$. The study of the global stellar mass density highlighted that $50-75$\\% of the present-day stellar mass was already in place at $z\\simeq1$, while only $3-14$\\% had been already assembled at $z\\simeq2.7$.  Direct determinations of the galaxy mass function based on near-IR deep imaging by the Spitzer Space Telescope indicate that the number of massive galaxies does not significantly evolve up to at least $z\\simeq1$ \\citep{franceschini2006,fontana2006,bundy2005}.  Using data from the Gemini Deep Deep Survey \\citep[GDDS,][]{abraham2004}, \\citet{glazebrook2004} identified a population of red ($[I-K]>4$ Vega mag.) galaxies at $z\\simeq2$ with stellar masses in excess of $10^{11}$ M$_\\odot$. These objects contribute roughly 30\\% of the total stellar mass density of the Universe at that epoch. \\citet{mccarthy2004} estimated the age of red galaxies at $z=1.3-2.2$ in the GDDS, deriving a median age of 1$-$3 Gyr, and a star formation history dominated by very powerful bursts (300$-$500 M$_\\odot$ yr$^{-1}$). These massive galaxies must have undergone a rapid formation process at $z>1$.  Exploiting data from the K20 survey \\citep{cimatti2002c,cimatti2002b,cimatti2002a}, \\citet{daddi2004} identified few luminous $K$-band selected galaxies at $1.7 < z < 2.3$ with stellar masses  $M \\simeq 10^{11}-5\\times 10^{11}\\ [M_\\odot]$. Combining deep K20 spectroscopy and HST-ACS imaging, \\citet{cimatti2004} discovered four old, fully assembled spheroidal galaxies at $1.6 < z < 1.9$: the most distant such objects currently known. The stellar mass of these galaxies turned out to be in the range $1-3\\times10^{11}$ M$_\\odot$. \\citet{fontana2004} studied K20 galaxies as well, showing that massive ($M>10^{11}$ M$_\\odot$) galaxies are easily found up to $z\\simeq2$. These authors also report on the stellar mass function: only mild evolution ($\\sim$2$-$30\\%) is detected to $z=1$, but only $\\sim35$\\% of the $z=0$ stellar mass locked up in massive objects was assembled by $z=2$.  At even higher redshifts, \\citet{rigopoulou2006} have recently studied a population of $z\\sim3$ Lyman-break galaxies (LBGs) with stellar masses in excess of $10^{11}$ M$_\\odot$. \\citet{mclure2006} identified nine LBGs at $z\\ge5$ in the UKIDSS \\citep{lawrence2006} survey, over an area of 0.6 deg$^2$. A stacking analysis suggests that the typical stellar mass of these sources is $>5 \\times 10^{10}$ M$_\\odot$. \\citet{mobasher2005} analyzed the properties of J-dropouts in the Hubble Ultra Deep Field \\citep[HUDF,][]{beckwith2006}, exploiting Spitzer photometry. They identified a $z\\sim6.5$ candidate that was interpreted as a post-starburst galaxy with a surprisingly high stellar mass of $5.7 \\times 10^{11}$ M$_\\odot$ (but see, for example, Yan et al. \\citeyear{yan2004} for a different interpretation).  \\citet{drory2005} studied the stellar mass function of galaxies in the FORS Deep Field \\citep[FDF,][]{heidt2003} and GOODS/CDFS \\citep{giavalisco2004} field, over a total area of 90 arcmin$^2$ and found that the total stellar mass density at $z=1$ is 50\\% of the local value. At $z=2$, 25\\% of the local mass density was already assembled, and at $z=3$ and $z=5$, at least 15\\% and 5\\% of stellar mass, respectively, was already in place. Massive ($M>10^{11}$ M$_\\odot$) galaxies existed over the whole redshift range probed, up to $z=5$. The number density of these massive galaxies evolves very similarly to galaxies with $M>10^{10}$  M$_\\odot$, decreasing by 0.4 dex to $z=1$, 0.6 dex to $z=2$, and 1 dex to $z=4$.  By analyzing the properties of $K$-selected galaxies in the $\\sim131$ arcmin$^2$ of the GOODS-CDFS survey, \\citet{caputi2006} found that the vast majority ($85-90$\\%) of local $M > 2.5 \\times 10^{11}$ M$_\\odot$ galaxies appears to be already in place at $z\\sim1$. These authors also infer that roughly $65-70$\\% of these galaxies assembled at $z=1-3$ by means of obscured, intense bursts of star formation, while the remaining could be in place at even higher redshifts ($z= 3-4$).  The observational challenge is that large volumes are needed to find representative samples of such rare very massive galaxies at high redshift. The Spitzer Wide-area InfraRed Extragalactic survey \\citep[SWIRE,][]{lonsdale2003,lonsdale2004} observed $\\sim$49 deg$^2$ in the seven Spitzer channels, and is therefore ideal to find rare objects. Its volume is large enough to detect $\\sim$85 DM haloes of mass $> 10^{14} M_\\odot$ in the $2<z<3$ redshift range \\citep{jenkins2001,mo2002}. These haloes are predicted to host the most luminous ($L_{bol}>10^{12}\\ L_\\odot$) and most massive ($M_\\star>\\textrm{several}\\ 10^{11}\\ M_\\odot$) galaxies ever to exist.  The Infrared Array Camera \\citep[IRAC,][]{fazio2004}, onboard Spitzer \\citep{werner2004}, samples the restframe near-IR light of distant galaxies. A near-IR selection not only directly probes the low-mass stars dominating the baryonic mass of a galaxy, but also is minimally affected by dust extinction. Therefore Spitzer is best suited to the study of the stellar content of galaxies up to $z=3$. Moreover Spitzer allows detection of galaxies that would be missed by restframe UV selection, for example distant red galaxies \\citep[e.g.][]{daddi2004}.  We take advantage of the shape of near-IR spectral energy distribution (SEDs) of galaxies to identify high-redshift objects on the basis of IRAC colors. Our selection is based on the detection of the 1.6$\\mu$m stellar peak in galaxies \\citep{sawicki2002,simpson1999}, redshifted to the IRAC domain. It is also worth noting that galaxies selected in this way benefit from a negative k-correction in the IRAC bands, because the slope of a galaxy's SED is negative redward of the 1.6$\\mu$m peak. In this way we performed a systematic search for $M\\gtrsim10^{11}$ M$_\\odot$ galaxies at $z>1$.  The analysis is carried out in the central square degree of the ELAIS-S1 SWIRE field, where optical, near-IR ($J$, $K_s$) photometry and optical spectroscopy are available (Berta et al., \\citeyear{berta2006}, Dias et al., in prep., La Franca et al., in prep.). This area, and the sampled volume, are bigger than any other previously explored for studying very massive galaxies at $z=1-3$, which were limited to very deep, pencil-beam surveys \\citep[e.g.][]{dickinson2003,drory2005,gwyn2005,fontana2006}.  This paper is structured as follows. Section {\\bf 2} presents the data available in the ELAIS-S1 field; in Sect. {\\bf 3} we present our selection criterion; then Sect. {\\bf 4} discusses the photometric estimate of redshifts. Section {\\bf 5} deals with the estimate of the stellar mass in galaxies, and presents a very detailed analysis of the influence of mid-IR and near-IR constraints on it. Experiments with different stellar libraries and IMFs are also discussed. Section {\\bf 6} presents the stellar mass function of our galaxies, including completeness correction and a parametric fit based on a Bayesian formalism. Finally, Sects. {\\bf 7} and {\\bf 8} discuss results and draw our conclusions.  Throughout this work, we adopt a standard $H_0=71$ $[$km s$^{-1}$ Mpc$^{-1}]$, $\\Omega_m=0.27$, $\\Omega_\\Lambda=0.73$ cosmology, unless otherwise stated. ",
        "conclusions": "Sampling the very massive tail of the stellar mass function at high redshift and estimating its contribution to the global stellar mass density is a critical task in modern cosmology, motivated by the recent evidence that a ``downsizing'' effect exists in the evolution of stellar mass across cosmic time.  We have exploited SWIRE/Spitzer and ancillary data in the ELAIS-S1 area to perform a systematic search for high redshift ($z\\gtrsim1.0$) massive ($M>10^{11}$ M$_\\odot$) galaxies. High redshift systems have been isolated by identifying the 1.6$\\mu$m restframe stellar peak shifted to IRAC wavelengths (3.6$-$8.0 $\\mu$m). The availability of near-IR ($J$ and $K_s$ band) data allowed us to avoid low-redshift interlopers and selection aliasing due to bright 3.3 $\\mu$m PAH features at $z\\simeq0.4$. A total of 203 5.8$\\mu$m-peakers and 123 4.5$\\mu$m-peakers have been identified over one square degree in the ELAIS-S1 field.  We have performed an extensive SED analysis, based on mixed stellar population synthesis, focused on deriving the stellar masses of the selected sample. The advantage of near-IR and mid-IR constraints, as well as the dependence of results on the choice of the IMF and SSP library have been explored in detail. The main results are: \\begin{itemize} \\item because of the shallow 5.8$\\mu$m flux cut adopted, the SWIRE IR-peaker sample consists of very massive galaxies, the majority of sources having $M_\\star>10^{11}$ M$_\\odot$. Objects in the range $z=1-2$ peak at $M_\\star\\simeq10^{11}$ M$_\\odot$, while the distribution of 5.8$\\mu$m-peakers ($z=2-3$) is centered at $M_\\star\\simeq2\\times 10^{11}$ M$_\\odot$. Typical uncertainties in stellar mass estimate (due to degeneracies in the SFH space) range between 0.1 and 0.3 dex, depending on multiwavelength coverage. The emission of $\\sim$30\\% of the sources turns out to be dominated by stars older than 1 Gyr, and also in the majority of the remaining cases old stellar populations do contribute to the observed SEDs. \\item the availability of mid-IR data provides a valuable constraint on the recent star formation history of individual galaxies. If no 24$\\mu$m flux (nor upper limit) were available, the resulting stellar masses could be underestimated and the spread in mass would be wider. \\item despite being very useful in the selection process, near-IR ($J$ and $K_s$) data turned out to be effective in constraining the D4000 break only in $\\sim$15\\% of cases, and preferentially for 4.5$\\mu$m-peakers. \\item the choice of a \\citet{chabrier2003} IMF, instead of a \\citet{salpeter1955} one, leads to systematically lower stellar masses, resulting in a rigid shift of the stellar mass distribution by $\\sim0.3$ dex. \\item when including thermally-pulsing AGB stars in the SSP library \\citep{maraston2005}, the fraction of objects dominated by old stellar populations increases, but the overall stellar mass distribution does not change significantly, because of the different $M_\\star/L$ ratios of SSPs. \\end{itemize} The stellar mass estimates have been used to compute the comoving number density of galaxies as a function of stellar mass (i.e. the observed stellar mass function), adopting the accessible volume formalism. Following \\citet{fontana2004}, we have corrected the samples for mass incompleteness and recovered the mass function down to $1.25\\times10^{11}$ and $1.6\\times10^{11}$ M$_\\odot$ for 4.5$\\mu$m- and 5.8$\\mu$m-peakers respectively.  Unfortunately, the selection, based on a 5.8$\\mu$m flux cut turned out to be only partially sensitive to 4.5$\\mu$m peakers, therefore only a lower limit on the mass function of $z=1-2$ sub-sample could be set.  The observed stellar mass function of 5.8$\\mu$m-peakers was reproduced with a parametric function, using the STY \\citep{sandage1979} approach. The uncertainties in the stellar masses of individual sources were automatically included in the analysis by using a Bayesian formalism, and the best fit was obtained with a MCMC sampling of the parameter space. The stellar mass function was finally integrated to derive the stellar mass density locked in 5.8$\\mu$m-peakers with $M>1.6\\times10^{11}$ M$_\\odot$, at $z=2-3$. The results of the mass function analysis are: \\begin{itemize} \\item the wide area surveyed by SWIRE allows the very massive tail of the stellar mass function to be probed with unprecedented detail at $z=2-3$, extending previous analyses up to $M_\\star=7\\times10^{11}$ $[$h$_{70}^{-2}$ M$_\\odot]$. \\item at $M<5\\times10^{11}$ M$_\\odot$, a significant intrinsic evolution has been detected across the redshift range $z=2-3$, strongly dependent on mass. The dependence or the 5.8$\\mu$m-peakers number density on $\\left(1+z\\right)$ has powers of $\\sim-0.4$ and $\\sim-0.65$ for $M\\sim2\\times10^{11}$ and $4\\times10^{11}$ M$_\\odot$ respectively. In the highest mass bins ($M\\ge5\\times10^{11}$ M$_\\odot$) the number density of IR-peakers keeps nearly constant. \\item comparison to literature data for the $z=2-3$ mass function shows an overall agreement of the 5.8$\\mu$m-peaker comoving number density to that of the K20, MUNICS, and GOODS-MUSIC surveys, in the higher mass regime. At lower masses a significant evolutionary correction should be applied, when the mass function is averaged over a wide redshift range. \\item a significant evolution of the stellar mass function of $M\\gtrsim10^{11}$ M$_\\odot$ galaxies with respect to the local estimate was detected: $\\Phi(z=2-3)\\le0.1\\times\\Phi(z=0)$. Combining 5.8$\\mu$m-peakers and literature data \\citep[e.g.][]{fontana2006}, this implies that the bulk of massive galaxies was not yet in place by the time the Universe was $\\sim3$ Gyr old, but must have been assembled in the following $\\sim1.5$ Gyr of evolution. \\item current hydro-dynamical models significantly overestimate the number density of massive galaxies, while the semi-analytic approach underestimates it. \\item since SWIRE 5.8$\\mu$m-peakers sample only the very massive tail of the mass function, the Schechter slope $\\alpha$ cannot be constrained. Despite its very low significance, the best fit value $\\alpha=-0.30$ is similar to that found in the literature. The other best fit parameters are: $M^\\ast = 1.66 \\times10^{11}$ $[$h$_{70}^{-2}$ M$_\\odot]$ $\\pm$1 dex; $\\Phi^\\ast=0.00022^{+0.00004}_{-0.00009}$ $[$ h$_{70}^3$ Mpc$^{-3}]$. \\item the integrated stellar mass density of 5.8$\\mu$m peakers is $\\rho_\\star = 1.18 \\times 10^7$ $[$h$_{70}$ M$_\\odot$ Mpc$^{-3}]$, with a 3$\\sigma$ range of $\\pm 0.3$ dex. Only a lower limit could be set for 4.5$\\mu$m-peak galaxies: $\\rho_\\star\\ge6.55\\times10^6$ $[$h$_{70}$ M$_\\odot$ Mpc$^{-3}]$. \\item on average SWIRE massive 5.8$\\mu$m-peakers provide 30$-$50\\% of the total stellar mass density in galaxies at $z=2-3$. \\item 5.8$\\mu$m-peakers provide less than $\\sim10$\\% of the stellar mass locked in massive galaxies ($M=10^{11}-10^{12}$ M$_\\odot$) in the local Universe. \\end{itemize} The analysis carried out on SWIRE massive galaxies at $z>1.5$ over one square degree, highlighted the complementarity of wide-shallow and deep pencil-beam surveys.  On one hand, sampling the faint end of the luminosity function, i.e. the low-mass end of the mass function, is needed in order to constrain the shape of the mass function and the total stellar mass density in galaxies at high redshift.  On the other hand, very massive galaxies are rare objects on the sky, with a number density $1.2\\times10^{-4}$ $[$h$_{70}^3$ Mpc$^{-3}]$ and it is necessary to explore large volumes of Universe in order to fully characterize them. The natural extension of this analysis is to build the stellar mass function of high-$z$ galaxies over the whole SWIRE 49 deg$^2$ area, taking advantage of what we have learned thanks to the full multi-wavelength coverage in ELAIS-S1.  At the time being, not much information is known about the environment hosting these massive high-$z$ objects. Further analyses of this population should be carried out probing also the environmental frame they belong to, in order to correctly interpret their role in the ``downsizing'' scenario."
    },
    "0710/0710.2902_arXiv.txt": {
        "abstract": "We report on our initial analysis of a deep 510 ks observation of the Galactic oxygen-rich supernova remnant (SNR) G292.0+1.8 with the {\\it Chandra X-ray Observatory}. Our new {\\it Chandra} ACIS-I observation has a larger field of view and an order of magnitude deeper exposure than the previous {\\it Chandra} observation, which allows us to cover the entire SNR and to detect new metal-rich ejecta features. We find a highly non-uniform distribution of thermodynamic conditions of the X-ray emitting hot gas that correlates well with the optical [O {\\small III}] emission, suggesting the possibility that the originating supernova explosion of G292.0+1.8 was itself asymmetric. We also reveal spectacular substructures of a torus, a jet, and an extended central compact nebula all associated with the embedded pulsar J1124$-$5916. ",
        "introduction": "Supernova Remnant}  Figure~\\ref{fig:fig1} shows an X-ray color image of G292.0+1.8 created from our deep {\\it Chandra} data. The outer boundary of the radio SNR is overlaid with a white contour. The energy bands displayed in each color were chosen to emphasize major atomic line emission that illustrates the distribution of electron temperatures and ionization states across the SNR (Table~\\ref{tbl:tab1}). We note that line and underlying continuum emission is included in each subband; e.g., the bright blue emission near the SNR's center is dominated by synchrotron emission from the PWN (see the inset in Figure~\\ref{fig:fig1}; Hughes et al.\\ 2001). Our deep ACIS-I exposure comprehensively images the entire SNR to its faint outermost edge, which matches well the extent of the radio remnant (Figure~\\ref{fig:fig1}) and traces the location of the blast wave. Ejecta knots (Park et al.\\ 2004, Gonzalez \\& Safi-Harb 2003) appear brightly colored -- yellow, green or blue -- in Figure~\\ref{fig:fig1}.  They extend closest to the rim in the west and south and are furthest away in the SE quadrant. Due to their red-orange color in Figure~\\ref{fig:fig1} and positional coincidence with [O {\\small III}] ejecta \\citep{wink06}, the southernmost X-ray knots are likely to be O-rich as well. The SNR's full diameter is $\\sim$9$\\farcm$6 (N-S) and $\\sim$8$\\farcm$4 (E-W), corresponding to physical sizes of $\\sim$14.7$-$16.7 $d_6$ pc, where $d_6$ is the distance to G292.0+1.8 in units of 6 kpc, the distance we assume throughout \\citep{gaen03}.  Our new image of G292.0+1.8 shows little evidence for the featureless, spectrally-hard, and geometrically-thin X-ray filaments that trace sites of efficient particle acceleration in other young SNRs, such as Cas A, Tycho, and SN 1006. Instead the SNR's rim is defined by spectrally-soft emission that is faint and diffuse in places (e.g., the SE rim) and filamentary elsewhere. Particle acceleration is evidently occurring under rather different conditions in G292.0+1.8 than other young Galactic SNRs.  The SNR interior contains a complex network of knots and filaments with a variety of colors and morphologies.  The overall color distribution of these features is highly asymmetric; red-orange emission is dominant in the S-SE, while the W-NW regions are bright in emission appearing green-blue in color. We are confident that these variations largely reflect differences in the underlying distributions of gas temperature and ionization in the metal-rich ejecta, rather than variable foreground absorption.  This is supported by the relative spatial distributions of the Ne He$\\alpha$ and Ly$\\alpha$ lines.  Because these lines are separated in energy by only $\\sim$0.1 keV, their flux ratio is insensitive to absorption variations. The equivalent width maps show that Ne He$\\alpha$ emission is relatively enhanced in the S-SE, while Ne Ly$\\alpha$ emission is enhanced in the W-NW \\citep{park02}. The metal abundance variation is also unlikely a significant contributor for the observed large-scale color variation based on our hardness ratio (HR) analysis (see below).  We constructed a HR map (Figure~\\ref{fig:fig2}a) to investigate further the variation in the thermodynamic state of the ejecta across G292.0+1.8. (For technical reasons we added the continuum component [$E$ $<$ 1.25 keV] to the soft band line emission to enhance the hard band lines and suppress the strong hard continuum emission from the PWN.)  There are prominent enhancements in Figure~\\ref{fig:fig2}a near the W and NW boundary of the SNR that generally trace ``fingers'' of hot ejecta protruding out to the very boundary of the SNR (green-blue filaments in Figure~\\ref{fig:fig1}). There is also an overall large-scale variation from HR values $\\sim$2.4 in the hard ``NW'' region to $\\sim$1.1 in the soft ``SE'' region (Figure~\\ref{fig:fig2}a).  Example spectra extracted from small regions within the larger ``NW'' and ``SE'' regions clearly show the distinctive line ratios responsible for the observed HR variation (Figure~\\ref{fig:fig2}b). We can draw on our previous work fitting the spectra of individual knots to relate HR values to plasma temperatures.  Spectral region 3 from Park et al. (2004) lies in the hard ``NW'' region, and spectral analysis yields a best-fit $kT \\sim 5$ keV, corresponding to HR $\\sim$ 2.5. We find that this HR value strongly traces the electron temperature rather than on individual elemental abundances. Although we did not study knots in the soft ``SE'' region previously, our preliminary spectral modeling of the SE regions favors significantly lower temperatures ($kT$ $\\sim$ 0.7 keV), corresponding to HR values $\\sim$ 1.0. Thus, the observed HR distribution reveals a large-scale spatial variation of thermal condition of metal-rich ejecta, which could not be detected by previous low-resolution data \\citep{hughes94}; i.e., a significantly higher temperature of the hot gas ($kT$ $\\sim$ 5 keV) is implied in the N-W regions, while a relatively lower-temperature plasma ($kT$ $\\la$ 1 keV) prevails in the S-E regions of the SNR.  There is also a very similar and highly significant gradient in the optical properties of G292.0+1.8.  The region with the lowest HR values is coincident with the bulk of the optical [O {\\small III}] emission indicated by ``SE'' region in Figure~\\ref{fig:fig2}a \\citep{ghav05,wink06}. Across the projected middle of the SNR (within an $\\sim$3$'$ wide region aligned roughly N-S), there are dozens of isolated optical knots (generally uncorrelated with X-ray knots), while on the western edge no optical emission appears at all. This high degree of correlation between the optical and X-ray properties suggests the possibility that radiative cooling in the ejecta is responsible for the SE-NW gradient in observed properties. In this picture, the SE ejecta would be undergoing significant, perhaps catastrophic cooling. Across the projected middle of the SNR, cooling is probably just beginning to occur in isolated knots that happen to have the appropriate thermodynamic conditions, while the emission toward the NW remains fully nonradiative.  Variation in the ambient density surrounding G292.0+1.8 could provide an explanation for the observed asymmetry in the ejecta properties. However, there is no evidence for a higher density in the SE regions \\citep{braun86} as would be expected. In fact, 60 $\\mu$m images\\footnote{The apparent differences in the detailed structure of the two IR images in Figure~\\ref{fig:fig1} are unlikely to be real because of the highly processed nature of both images. The {\\it ISO} image was constructed from a rastor scan of many pointings using a $3\\times 3$ pixel ($\\sim$45$^{\\prime\\prime}$ per pixel) IR array that, unfortunately, did not fully cover the SE rim of the SNR. The {\\it IRAS} image, which did cover the entire SNR, was the result of an image reconstruction process intended to improve the angular resolution from the native value of several arcminutes to the level of 1$^{\\prime}$$-$2$^{\\prime}$. Given these significant differences, the overall general agreement between the images, specifically, the brightness of the SW rim and the faintness toward the SE, seems reasonable to us.} (see upper right and upper left insets to Figure~\\ref{fig:fig1}) show that around the rim of the SNR the SE region is a minimum in flux, while the SW is a maximum. Thus, albeit somewhat speculative at the current stage of the analysis, we raise the possibility that the ejecta asymmetry has its origin in some intrinsic asymmetry of the SN explosion itself, such as a variation in the density or velocity distribution from SE to NW. Further work including detailed X-ray spectral analysis of ejecta and hydrodynamic studies appropriate for G292.0+1.8 are required to test this asymmetric explosion scenario. ",
        "conclusions": ""
    },
    "0710/0710.5525_arXiv.txt": {
        "abstract": "Temperature anisotropies in the Cosmic Microwave Background (CMB) are affected by the late Integrated Sachs-Wolfe (lISW) effect caused by any time-variation of the gravitational potential on linear scales. Dark energy is not the only source of lISW, since massive neutrinos induce a small decay of the potential on small scales during both matter and dark energy domination. In this work, we study the prospect of using the cross-correlation between CMB and galaxy density maps as a tool for constraining the neutrino mass. On the one hand massive neutrinos reduce the cross-correlation spectrum because free-streaming slows down structure formation; on the other hand, they enhance it through their change in the effective linear growth. We show that in the observable range of scales and redshifts, the first effect dominates, but the second one is not negligible. We carry out an error forecast analysis by fitting some mock data inspired by the Planck satellite, Dark Energy Survey (DES) and Large Synoptic Survey Telescope (LSST).  The inclusion of the cross-correlation data from Planck and LSST increases the sensitivity to the neutrino mass $m_{\\nu}$ by 38\\% (and to the dark energy equation of state $w$ by 83\\%) with respect to Planck alone. The correlation between Planck and DES brings a far less significant improvement. This method is not potentially as good for detecting $m_{\\nu}$ as the measurement of galaxy, cluster or cosmic shear power spectra, but since it is independent and affected by different systematics, it remains potentially interesting if the total neutrino mass is  of the order of 0.2~eV; if instead it is close to the lower bound from atmospheric oscillations, $m_{\\nu} \\sim 0.05$~eV, we do not expect the ISW-galaxy correlation to be ever sensitive to $m_{\\nu}$. ",
        "introduction": "As photons pass through a changing gravitational potential well, they experience a redshift or a blueshift, depending on whether the well grows or decays respectively. Cosmic microwave background (CMB) photons can experience such variations between the time of last scattering and their detection now. This effect was first described by Sachs and Wolfe in 1967~\\cite{Sachs:1967er}, and hence is dubbed the integrated Sachs-Wolfe effect (ISW). During a Cold Dark Matter (CDM) and/or baryon dominated era, the gravitational potential distribution remains frozen, and the ISW effect has no net effect on the blackbody temperature of CMB photons. This property is crucially related to the fact that non-relativistic matter (like CDM and baryons) has a vanishing sound speed, and experiences gravitational clustering on all sub-Hubble scales after photon decoupling, as described by the Poisson equation. In such a situation, the universal expansion and the gravitational contraction compensate each other in such a way as to maintain a static gravitational potential. However, when the expansion rate is affected by any type of matter with a non-vanishing sound speed, e.g. during Dark Energy (DE) domination, the gravitational perturbations decay and the cosmic photon fluid experiences a blue shift, acquiring extra temperature perturbations related to the intervening pattern of matter perturbations. It was first proposed by Crittenden and Turok in 1995~\\cite{Crittenden:1995ak} to cross correlate maps of temperature perturbations in the CMB with those of matter overdensities in large scale structures (LSS), in order to measure a possible acceleration of the universe's expansion. However, the CMB and LSS data available at that time were not good enough for such an ambitious goal, and the first strong indication of a positive acceleration came in 1998 from the side of type-Ia supernovae~\\cite{Perlmutter:1998np,Riess:1998cb}.  Analyses of the first (2003) and second (2006) data releases of the Wilkinson Microwave Anisotropy Probe (WMAP)~\\cite{Spergel:2003cb,Spergel:2006hy} were the first to indicate the existence of Dark Energy independent of acceleration, by means of the location of the second peak in the CMB power spectrum.  Simultaneously, a number of interesting papers presented the first detections of the ISW effect by cross-correlating WMAP anisotropy maps with various LSS data sets ~\\cite{Fosalba:2003ix,Boughn:2003yz,Fosalba:2003iy,Afshordi:2003xu, Padmanabhan:2004,Cabre:2006qm,Cabre:2006uj,Giannantonio:2006du,McEwen:2007}, now able to give an independent measure for the acceleration of the expansion of the universe.  The domination of Dark Energy is not the only source of gravitational potential evolution and of a net ISW effect. On small cosmological scales, as soon as matter perturbations exceed the linear regime, gravitational perturbations start to grow and to redshift CMB photons. This effect, called the Rees-Sciama effect, has not been significantly detected until now~\\cite{Puchades:2006gs}. CMB photons can also be scattered by gravitational lensing \\cite{SeljakZaldarriaga:2000} and by the Sunyaev-Zeldovich (SZ)  effect \\cite{SZ:1969} (see \\cite{Fosalba:2003iy,Afshordi:2003xu,Afshordi:2007} for detections in CMB-LSS cross-correlation analysis). An other party expected to affect the evolution of gravitational perturbations --at least by a small amount-- is the background of massive neutrinos. Over thirty years ago massive neutrinos were proposed as a Hot Dark Matter (HDM) candidate, and later ruled out as the dominant dark component, since HDM tends to wash out small scale overdensities during structure formation~\\cite{Primack:2001ib}.  Observed neutrino oscillations however constrain neutrinos to have a mass~\\cite{Maltoni:2004ei,Fogli:2005cq}. In addition, the presence of a Cosmic Neutrino Background (CNB) is strongly suggested on the one hand by the abundance of light elements produced during primordial nucleosynthesis~\\cite{Cuoco:2003cu,Steigman:2005uz,Mangano:2006ur}, and on the other hand by CMB anisotropies~\\cite{Crotty:2003th,Pierpaoli:2003kw,Barger:2003zg,Trotta:2004ty,Hannestad:2005jj,Hannestad:2006mi,Ichikawa:2006vm,Ichikawa:2006vm,deBernardis:2007bu,Hamann:2007pi}. Therefore, a small fraction of HDM is expected to coexist with the dominant CDM component. On small cosmological scales (for instance, cluster scale), the free-streaming of massive neutrinos should induce a slow decay of gravitational and matter perturbations~\\cite{Bond:1980ha}, acting during both matter and Dark Energy domination. This effect depends on the total neutrino mass summed over all neutrino families, $m_{\\nu}=\\sum_i m_i$, unlike laboratory experiments based on tritium decay or neutrinoless double-beta decay, which probe different combinations: hence, a cosmological determination of the total neutrino mass would bring complementary information to the scheduled particle physics experiments~\\cite{Hannestad:2006zg,Lesgourgues:2006nd}. % The free streaming of massive neutrinos has not yet been detected~\\cite{Hannestad:2007tu}, but there are good prospects to do so in the future, since the smallest total neutrino mass allowed by data on atmospheric neutrino oscillations ($m_{\\nu} \\geq \\sqrt{\\Delta m^2_{\\rm atm}} \\sim 0.05$~eV) implies at least a 5\\% suppression in the matter/gravitational small-scale power spectrum~\\cite{Hannestad:2006zg,Lesgourgues:2006nd}. A positive detection --even in the case of minimal mass-- could follow from the analysis of future galaxy/cluster redshift surveys~\\cite{Lesgourgues:2004ps,Wang:2005vr,Hannestad:2007cp}, weak lensing surveys~\\cite{Song:2003gg,Hannestad:2006as}, Lyman-$\\alpha$ forest analysis, cluster counts~\\cite{Wang:2005vr}, etc. The goal of measuring the neutrino mass from cosmology is very ambitious since each of these methods suffers from its own source of systematics (bias issues, modeling of non-linear clustering, ...). Therefore, a robust detection could only be achieved by comparing the results from various types of experiments.  The goal of this work is to describe a possible cosmological determination of the absolute neutrino mass scale through the ISW effect induced by neutrino free-streaming on CMB temperature maps, using as an observable the cross-correlation function of galaxy-temperature maps. This possibility was investigated previously by Ichikawa and Takahashi \\cite{Ichikawa:2005hi} (and suggested again recently in \\cite{Kiakotou:2007pz}). As neutrinos slow down the growth of structure, we expect the blueshift caused by an accelerated expansion to be more pronounced if neutrinos have a larger mass. On the other hand, the distribution of matter inducing the late ISW effect is smoother in case of free-streaming by massive neutrinos. These two antagonist effects should in principle induce some mass-dependent variations in the galaxy-temperature cross-correlation function.  In section~\\ref{sec:theory} of this paper we give an outline of the theory of the ISW-effect in the presence of a neutrino mass. In section~\\ref{sec:mock}, we use some mock data with properties inspired from the Planck satellite, Dark Energy Survey (DES) and Large Synoptic Survey Telescope (LSST) in order to show the potential impact of this method in the future. ",
        "conclusions": "We have studied here the possibility to use the cross-correlation between CMB and galaxy density maps as a tool for constraining the neutrino mass.  On one hand massive neutrinos reduce the cross-correlation spectrum because their free-streaming slows down structure formation; on the other hand, they enhance it because of the behavior of the linear growth in presence of massive neutrinos. Using both analytic approximations and numerical computations, we showed that in the observable range of scales and redshifts, the first effect dominates, but the second one is not negligible. Hence the cross-correlation between CMB and LSS maps could bring some independent information on neutrino masses.  We performed an error forecast analysis by fitting some mock data inspired from the Planck satellite, Dark Energy Survey (DES) and Large Synoptic Survey Telescope (LSST).  For Planck and LSST, the inclusion of the cross-correlation data increases the sensitivity to $m_{\\nu}$ by 38\\%, $w$ by 83\\% and $\\Omega_{dm} h^2$ by 30\\% with respect to the CMB data alone. With the fiducial model employed in this analysis (based on eight free parameters) the standard deviation for the neutrino mass is equal to 0.38~eV for Planck alone and 0.27~eV for Planck plus cross-correlation data. This is far from being as spectacular as the sensitivity expected from the measurement of the auto-correlation power spectrum of future galaxy/cluster redshift surveys or cosmic shear experiments, for which the predicted standard deviation is closer to the level of 0.02~eV, leading to a 2$\\sigma$ detection even in the case of the minimal mass scenario allowed by current data on neutrino oscillations (see \\cite{Lesgourgues:2006nd} for a review). However, the method proposed here is independent and affected by different systematics. So, it remains potentially interesting, but only if the neutrino mass is not much smaller than $m_{\\nu} \\sim 0.2$~eV."
    },
    "0710/0710.0907_arXiv.txt": {
        "abstract": "This paper presents the results of a {\\em Spitzer} IRAC $3-8$ $\\mu$m photometric search for warm dust orbiting 17 nearby, metal-rich white dwarfs, 15 of which apparently have hydrogen dominated atmospheres (type DAZ).  G166-58, G29-38, and GD 362 manifest excess emission in their IRAC fluxes and the latter two are known to harbor dust grains warm enough to radiate detectable emission at near-infrared wavelengths as short as 2 $\\mu$m.  Their IRAC fluxes display differences compatible with a relatively larger amount of cooler dust at GD 362.  G166-58 is presently unique in that it appears to exhibit excess flux only at wavelengths longer than about 5 $\\mu$m.  Evidence is presented that this mid-infrared emission is most likely associated with the white dwarf, indicating that G166-58 bears circumstellar dust no warmer than $T\\sim400$ K.  The remaining 14 targets reveal no reliable mid-infrared excess, indicating the majority of DAZ stars do not have warm debris disks sufficiently opaque to be detected  by IRAC. ",
        "introduction": "The {\\em Spitzer Space Telescope} opened a new phase space to white dwarf researchers interested in the infrared properties of degenerate stars and their environments.  The majority of white dwarfs are inaccessible from the ground beyond 2.4 $\\mu$m due to their intrinsic faintness combined with the ever-increasing sky brightness towards longer wavelengths \\citep*{gla99}. This limits any white dwarf science which aims to study matter radiating at $T<1500$ K.  Prior to the launch of {\\em Spitzer}, only one previously published, directed mid-infrared study of white dwarfs exists ; an {\\em Infrared Space Observatory} search for dust emission around 11 nearby white dwarfs, 6 of which have metal-rich photospheres \\citep*{cha99}.  Owing to the superb sensitivity of {\\em Spitzer} \\citep*{wer04}, a Cycle 1 IRAC program was undertaken to search for warm dust emission associated with cool, hydrogen atmosphere white dwarfs with photospheric metals, the DAZ stars.  This paper presents a synopsis of the IRAC results, including the detection of $5-8$ $\\mu$m flux excess at G166-58, $3-8$ $\\mu$m data on G29-38 and GD 362, and also includes Gemini $3-4$ $\\mu$m spectroscopy of G29-38. ",
        "conclusions": "Although a survey of only 17 stars does not allow robust statistics, it is clear that the majority of DAZ white dwarfs do not harbor warm dust of sufficient emitting surface area to be detected with IRAC. If most or all of these stars do host circumstellar material, the fractional luminosities must be relatively low compared to currently known dusty white dwarfs.  This could result from a modest amount of dust (as in the zodiacal cloud), large particle sizes, or cooler material further from the star.  Another possibility is that any warm dust produced within the Roche limit of a white dwarf is swiftly destroyed through mutual collisions, not unlike ice in a blender, as it orbits with Keplerian velocities near $0.003c$.  In optically thin disks, particles with orbital period $p$ will collide on a timescale given by $t_{\\rm coll}=p/\\tau$.  The tidal rings at G29-38 and GD 362 have been modeled to extend from approximately $0.2-0.4$ $R_{\\odot}$, where a typical orbital period is only $p=0.6$ hr and the resulting collision timescale for $\\tau\\sim0.01$ is $t_{\\rm coll}=2.5$ dy. If a sizeable fraction of dust produced in a tidal disruption event is initially optically thin, then both collisions and Poynting-Robertson drag will compete to quickly annihilate this material.  The ratio of these two timescales for dust particles orbiting a distance $D$ from a star of mass $M$ can be written as  \\begin{equation} \\gamma  = \\frac{t_{\\rm pr}}{t_{\\rm coll}} = 693 \\ \\frac{\\sqrt{MD} \\rho a\\tau}{LQ_{\\rm pr}} \\end{equation}  \\noindent where $M$ and $L$ are in solar units, $D$ is in AU, $\\rho$ is in g ${\\rm cm}^{-3}$ , and $a$ is in microns.  This fraction reaches a minimum for 0.1 $\\mu$m grains with 1 g ${\\rm cm}^{-3}$ at the inner disk edge, and yields $\\gamma_{\\rm min}>10$ for all possible white dwarf disk parameters. Table \\ref{tbl6} gives minimum values for $\\gamma$ at the inner edges of the disks at G29-38, GD 362, and G166-58.  Therefore collisions will erode dust grains faster than they can be removed by angular momentum loss. This is also true in the case of optically thick disks where: 1) the bulk of material is shielded from starlight, and hence the Poynting-Robertson effect is diminished, and 2) the collision timescale is less than half the orbital period \\citep*{esp93}.  Hence, for a wide range of disk densities, it is plausible that mutual collisions within an evolving dust ring at a typical white dwarf will result in the relatively rapid self-annihilation of the micron size grains required to radiate efficiently at $3-30$ $\\mu$m.  Following the tidal disruption of an asteroid, if one models the dust produced as a collisional cascade, the expected particle size distribution behaves classically as $n(a)\\propto a^{-3.5}$ \\citep*{doh69}.  For dust at main sequence stars, this distribution is reshaped on short timescales as sub-micron size grains are removed by radiation pressure which, as shown above, does not apply in the case of white dwarfs.  In the absence of radiation pressure, the average particle in will have a size $\\bar{a}= 5/3$ $a_{\\rm min}$.  For practical purposes, at white dwarfs one can assume $a_{\\rm min}=0.01$ $\\mu$m, where particles are already inefficient absorbers and emitters of infrared radiation, and anything smaller approaches the size of gas molecules and atoms.  With such a distribution of extant dust, 99.7\\% of the particles will have sizes $a\\leq0.1$ $\\mu$m, leaving a paltry fraction of larger particles which could effectively support infrared emission from the disk.  Clearly, such small particles and gas might be present at most or all white dwarfs which show signs of circumstellar accretion such as the DAZ stars.  On the other hand, the persistence of warm dust disks at several white dwarfs \\citep*{jur07b} must be explained despite the fact that in some or most cases where it is produced, it may also be efficiently destroyed. One possibility is that the disk density (which could contain gas) becomes sufficiently high as to damp out collisions in the disk and also make it optically thick, thus somewhat protecting it from self-erosion and drag forces simultaneously.  The evolution of such a dense, fluid-like ring is then dominated by viscous forces (differential rotation and random motions) which cause it to spread, losing energy in the process \\citep*{esp93}.  The maximum lifetime of such a ring occurs at minimum viscosity  \\begin{equation} t_{\\rm ring} \\approx \\frac{ {{\\rho}^2} {w^2} p} {2 \\pi {{\\sigma^2}} } \\end{equation}  \\noindent where $w$ is the radial extent of the ring and $\\sigma$ is the surface mass density \\citep*{esp93}.  If the mass of a large solar system asteroid, $10^{24}$ g, were spread into a tidal ring of negligible height ($h<10$ m), a radial extent $0.2-0.4$ $R_{\\odot}$, consisting of micron size particles orbiting a typical white dwarf, the resulting volume mass density (0.55 g ${\\rm cm}^{-3}$) would be sufficiently high that the mean free path of particles is on the same order as their size.  This could effectively damp out collisions, thus minimizing viscosity, and with a resulting surface mass density of $\\sigma\\approx550$ g ${\\rm cm}^{-2}$, permit a potential disk lifetime -- in the absence of competing forces -- longer than the Gyr white dwarf cooling timescales.  However, for sustained accretion rates as low as $\\dot{M}=10^{10}$ g ${\\rm s}^{-1}$, a $10^{24}$ g disk would become fully consumed within several Myr.  Large rocks and colder material orbiting at $D\\ga100$ $R_{\\odot}$ will be unaffected by any of the aforementioned processes, and such a reservoir of material is strictly necessary to supply some fraction of DAZ white dwarfs with photospheric metals, regardless of circumstellar dust production (collisional versus tidal) and evolution (persistence versus destruction).  The overall number of white dwarfs with remnant planetesimal belts may be rather high based on a growing number of detections.  If one takes 12\\% (\\S4.1) as the fraction of DAZ stars with circumstellar dust as observed by {\\em Spitzer} to date, 20\\% as the fraction of DAZ stars among cool DA white dwarfs \\citep*{zuc03}, and 80\\% as the number of cool DA stars among all white dwarfs in the field \\citep*{eis06}, then a lower limit to the number of white dwarfs with asteroid-type belts is at least 2\\%.  This fraction could be as high as 20\\% if the majority of metal-rich white dwarfs harbor circumstellar matter, which raises important questions about the implied frequency of planetesimal belts around main-sequence stars and the current detection rate (see the Appendix of \\citealt*{jur06}).  Owing to their low luminosities, white dwarfs which may have been polluted by heavy elements in winds or transferred material from substellar companions \\citep*{deb06,dob05,zuc03,sio84} are easily identified with IRAC observations \\citep*{mul07,han06,far05a,far05b} down to T dwarf temperatures.  There is no evidence of such companions in the data presented here, ruling out all but the coldest brown dwarfs, active planets and moons as close orbiting, companion-like polluters (Farihi et al. 2008, in preparation)."
    },
    "0710/0710.5180_arXiv.txt": {
        "abstract": "Mass bounds on dark matter (DM) candidates are obtained for particles that decouple in or out of equilibrium  while ultrarelativistic with {\\bf arbitrary} isotropic and homogeneous distribution functions. A coarse grained Liouville invariant primordial phase space density $ \\mathcal D $ is introduced which depends solely on the distribution function at decoupling. The density  $ \\mathcal D $ is explicitly computed and combined with recent photometric and kinematic data on dwarf spheroidal satellite galaxies in the Milky Way  (dShps) and the observed DM density today yielding upper and lower bounds on the mass, primordial phase space densities and velocity dispersion of the DM candidates. Combining these constraints with recent results from $N$-body simulations yield estimates for the mass of the DM particles in the range of a few keV. We establish in this way a direct connection between the microphysics of decoupling \\emph{in or out} of equilibrium and the constraints that the particles must fulfill to be suitable DM candidates. If chemical freeze out occurs before thermal decoupling, light bosonic particles can Bose-condense. We study  such Bose-Einstein {\\it condensate} (BEC) as a dark matter candidate. It is shown that depending on the relation between the critical ($ T_c $) and decoupling ($ T_d $) temperatures,  a BEC light relic could act as CDM but the decoupling scale must be {\\it higher} than the electroweak scale. The condensate hastens the onset of the non-relativistic regime and tightens the upper bound on the particle's mass. A non-equilibrium scenario which describes particle production and partial thermalization, sterile neutrinos produced out of equilibrium and other DM models is analyzed in detail and the respective bounds on mass, primordial phase space density and velocity dispersion are obtained. Thermal relics with $ m \\sim \\mathrm{few}\\,\\mathrm{keV} $  that decouple when ultrarelativistic and sterile neutrinos produced resonantly or non-resonantly lead to a primordial phase space density compatible with {\\bf cored} dShps and disfavor cusped satellites. Light Bose-condensed DM candidates yield phase space densities consistent with {\\bf cores} and if $ T_c\\gg T_d $ also with cusps. Phase space density bounds on particles that decoupled non-relativistically  combined with recent results from N-body simulations suggest a potential tension for WIMPs with $ m \\sim 100\\,\\mathrm{GeV},T_d \\sim \\,10\\,\\mathrm{MeV} $. ",
        "introduction": "Although the existence of dark matter (DM) was inferred several decades ago \\cite{zwoo}, its nature still remains elusive. Candidate dark matter particles are broadly characterized as cold, hot or warm depending on their velocity dispersions. The clustering properties of collisionless DM candidates in the linear regime depend on the free streaming length, which roughly corresponds to the Jeans length with the particle's velocity dispersion replacing the speed of sound in the gas. Cold DM (CDM) candidates feature a small free streaming length favoring a bottom-up hierarchical approach to structure formation, smaller structures form first and mergers lead to clustering on the larger scales.  Among the CDM  candidates are weakly interacting massive particles (WIMPs) with $m \\sim 10-10^{2}\\,\\mathrm{GeV}$. Hot DM (HDM) candidates feature large free streaming lengths and favor top down structure formation, where larger structures form first and fragment. HDM particle candidates are deemed to have masses in the few $\\mathrm{eV}$ range, and warm DM (WDM) candidates are intermediate with  a typical mass range $ m \\sim 1-10 \\,\\mathrm{keV} $.  \\medskip  The \\emph{concordance} $ \\Lambda\\mathrm{CDM} $ standard cosmological model emerging from CMB, large scale structure observations and simulations favors the hypothesis that DM is composed of primordial particles which are cold and collisionless \\cite{primack}. However, recent observations hint at possible   discrepancies with the predictions of the $ \\Lambda\\textrm{CDM} $ concordance model: the satellite and cuspy halo problems. \\\\  The satellite problem, stems from the fact that CDM favors the presence of substructure: much of the CDM is not smoothly distributed but is concentrated in small lumps, in particular in dwarf galaxies for which there is scant observational evidence so far. A low number of satellites have been observed in Milky-Way sized galaxies \\cite{kauff,moore,moore2,klyp}. This substructure is a consequence   of the CDM power spectrum which favors small scales becoming non-linear first, collapsing in the bottom-up hierarchical manner and surviving the mergers as dense clumps \\cite{moore,klyp}. \\\\  The cuspy halo problem arises from the result of large scale $N$-body simulations of CDM clustering which predict a monotonic increase of the density towards the center of the halos \\cite{dubi,frenk,moore2,bullock,cusps}, for example the universal Navarro-Frenk-White (NFW) profile $ \\rho(r)\\sim r^{-1}(r+r_0)^{-2} $ \\cite{frenk} which describes accurately clusters of galaxies, but indicates a divergent cusp at the center of the halo. Recent observations seem to indicate central cores in dwarf galaxies \\cite{dalcanton1,van,swat,gilmore}, leading to the 'cusps vs cores' controversy. \\\\  A recent compilation of observations of dwarf spheroidal  galaxies dSphs \\cite{gilmore}, which are considered to be prime candidates for DM subtructure \\cite{spergel}, seem to \\emph{favor} a core with a smoother central density and a low mean mass density $ \\sim 0.1\\,M_{\\odot}/\\mathrm{pc}^3 $ rather than a cusp \\cite{gilmore}. The data  cannot yet rule out cuspy density profiles which allow a maximum density $ \\lesssim 60 \\,M_{\\odot}/\\mathrm{pc}^3 $ and the interpretation and analysis of the observations is not yet conclusive \\cite{dalcanton1,van2}.  These \\emph{possible} discrepancies have rekindled an interest in WDM particles,  which feature a velocity dispersion larger than CDM particles, and consequently larger free-streaming lengths which smooth-out the inner cores and would be prime candidates to relieve the cuspy halo and satellite  problems \\cite{turok}.  \\medskip  A possible WDM candidate is a sterile neutrino \\cite{dw,este,kusenko} with a mass in the $\\mathrm{keV}$ range and produced via their mixing and oscillation with an active neutrino species either non-resonantly \\cite{dw}, or through MSW (Mikheiev-Smirnov-Wolfenstein) resonances in the medium \\cite{este}. Sterile neutrinos can decay into a photon and an active neutrino (more precisely the largest mass eigenstate decays into the lowest one and a photon) \\cite{pal} yielding the possibility of direct constraints on the mass and mixing angle from the diffuse X-ray background \\cite{Xray}.  \\medskip  Observations of cosmological structure formation  via the Lyman-$\\alpha$ forest provide a complementary probe of primordial density fluctuations on small scales which yield an indirect constraint on the masses of WDM candidates. While constraints from the diffuse X-ray background yield an \\emph{upper} bound on the mass of a putative sterile neutrino in the range $ 3-8\\,\\mathrm{keV} $ \\cite{Xray}, the latest Lyman-$\\alpha$ analysis \\cite{lyman}  yields \\emph{lower} bounds in the range $ 10-13\\,\\mathrm{keV} $   in tension with the X-ray constraints.  More recent constraints from Lyman-$\\alpha$ yield a lower limit for the mass of a WDM candidate $ m_{WDM} \\gtrsim 1.2\\,\\mathrm{keV} \\,(2\\sigma) $ for an early decoupled \\emph{thermal} relic and $m_{WDM} \\gtrsim 5.6\\,\\mathrm{keV} \\,(2\\sigma)$ for sterile neutrinos \\cite{viel}. Strong upper limits on the mass and mixing angles of sterile neutrinos have been recently discussed \\cite{beacom}, however, there are uncertainties as to whether WDM candidates can explain large cores in dSphs \\cite{strigari}. It has been recently argued \\cite{palazzo} that if sterile neutrinos are produced non-resonantly \\cite{dw} the combined X-ray and Lyman-$\\alpha$ data suggest that  these \\emph{cannot} be the \\emph{only} WDM component, with an upper limit for  their fractional relic abundance  $ \\lesssim 0.7 $. Recent \\cite{boyarski2} constraints on a radiatively decaying DM particle from the EPIC spectra of (M31) by XMM-Newton confirms this result and  places a stronger lower mass limit $ m < 4\\,\\mathrm{keV} $.  All these results suggest  that DM could be a mixture of several components with sterile neutrinos as viable candidates.  \\medskip  {\\bf Motivation and goals:} Although the $\\Lambda\\mathrm{CDM}$ paradigm describes large scale structure formation remarkably well, the \\emph{possible} small scale discrepancies mentioned above  motivate us to  study new constraints that different dark matter components must fulfill to be suitable candidates. Cosmological bounds on  dark matter components primarily focused on standard model neutrinos \\cite{bond,TG}, heavy relics that decoupled in local thermodynamic  equilibrium (LTE) when non-relativistic \\cite{LW,kt,dominik} or \\emph{thermal ultrarelativistic relics} \\cite{madsen,madsenbec,madsenQ,salu,hogan}. More recently, cosmological precision data were used to constrain the (HDM) abundance of low mass particles \\cite{pastor, steen,raffelt,mena} assuming these to be thermal relics.  \\medskip  The main results of this article are:  \\medskip  (\\textbf{a:}) We consider particles that  decouple {\\bf in or  out of LTE} during the radiation dominated era with  an \\emph{arbitrary} (but homogeneous and isotropic) distribution  function.  Particles which decouple being ultrarelativistic  eventually become non-relativistic because of redshift of physical momentum.   We establish a direct connection between the microphysics of decoupling \\emph{in or out of LTE} and the constraints that the particles must fulfill to be suitable DM candidates \\emph{in terms of  the distribution functions at decoupling}.  \\medskip  (\\textbf{b:}) We introduce a  primordial coarse grained phase space density $$ \\mathcal{D}  \\equiv \\frac{n(t)}{\\big\\langle \\vec{P}^2_f \\big\\rangle^\\frac32} \\; , $$ where $ n(t) $ is the number of particles per unit physical volume and $\\Big\\langle\\vec{P}^2_f\\Big\\rangle$ is the average of the physical momentum with the distribution function of the decoupled particle. $ \\mathcal D $ is a Liouville invariant after decoupling and only depends on the distribution functions at decoupling. In the non-relativistic regime $ \\mathcal D $ is simply related  to the phase densities considered in refs. \\cite{dalcanton1,TG,hogan,madsenQ} and   can only {\\bf decrease} by collisionless phase mixing or self-gravity dynamics \\cite{theo}.  In the non-relativistic regime we obtain \\be \\label{Dint} \\mathcal{D} = \\frac1{3^\\frac32 \\;  m^4} \\; \\frac{\\rho_{DM}}{\\sigma^3_{DM}} \\ee where $ \\sigma_{DM} $ is the primordial  one-dimensional velocity dispersion and $ \\rho_{DM} $ the dark matter density. Combining the result for the primordial phase space density $ \\mathcal D $ determined by the mass and the distribution function of the decoupled particles, with the recent compilation of photometric and kinematic data on dSphs satellites in the Milky-Way \\cite{gilmore} yields {\\bf lower} bounds on the DM particle mass $ m $ whereas  {\\bf upper} bounds on the DM mass are obtained using the value of the observed dark matter density today. Therefore the combined analysis of observational data from (dSphs), N-body simulations and the present DM density allows us to establish both \\emph{upper and lower} bounds on the mass of the DM candidates.  We thus provide a link between the microphysics of decoupling, the observational aspects of dark matter halos and the DM mass value.  \\medskip  (\\textbf{c:}) Recent $N$-body simulations \\cite{numQ} indicate that the phase-space density decreases a factor $ \\sim 10^2 $ during gravitational clustering. This result combined with eq.(\\ref{Dint}) and the observed values on dSphs satellites \\cite{gilmore} yield $$ m_{cored}\\sim \\frac2{g^\\frac14} \\; \\mathrm{keV} \\quad , \\quad m_{cusp} \\sim\\frac8{g^\\frac14} \\; \\mathrm{keV} \\; . $$ for the   masses of \\emph{thermal relics} DM candidates, where `cored' and 'cusp' refer to the type of profile used in the dShps description and $  1\\leq g \\leq 4 $ is the number of internal degrees of freedom of the DM particle. Wimps with masses $ \\sim 100 \\, \\textrm{GeV} $ decoupling in LTE at temperatures $ T_d \\sim 10\\,\\mathrm{MeV} $ lead to primordial phase space densities many orders of magnitude larger than those observed in (dSphs).  The results of $N$-body simulations, which yield relaxation by $2-3$ orders of magnitude\\cite{numQ} suggest a potential tension for WIMPs as DM candidates. However, the $N$-body simulations in ref.\\cite{numQ} begin with initial conditions with values of the phase space density much lower than the primordial one. Hence it becomes an important question whether the enormous relaxation required from the primordial values to those of observed in dSphs can be inferred from numerical studies with suitable (much larger) initial values of the phase space density.    \\medskip  (\\textbf{d:}) We study the possibility that the DM particle is a light Boson that undergoes  Bose-Einstein Condensation (BEC) prior to decoupling while still ultrarelativistic. (This possibility  was addressed in \\cite{madsenbec}). We analyze in detail the constraints on such BEC DM candidate from velocity dispersion and phase space arguments, and contrast the BEC DM properties to those of the hot or warm thermal relics.  \\medskip  (\\textbf{e:}) Non-equilibrium scenarios that describe various possible WDM candidates are studied in detail. These scenarios describe particle production \\cite{boydata} and incomplete thermalization \\cite{dvd}, resonant \\cite{dw} and non-resonant \\cite{este} production of sterile  neutrinos and a model recently proposed \\cite{strigari} to describe cores in dSphs.   Our analysis of the DM candidates is based on their masses, statistics and  properties at decoupling (being it in LTE or not). We combine observations on dSphs \\cite{gilmore} and $N$-body simulations \\cite{numQ}, with theoretical analysis using the non-increasing property of the phase space density \\cite{TG,dalcanton1,hogan,theo}.  The results from the combined  analysis of the primordial phase space densities, the observational data on dSphs \\cite{gilmore} and the $N$-body simulations in ref.\\cite{numQ} are the following: \\begin{itemize} \\item{\\textbf{(i):}  conventional thermal relics, and sterile neutrinos produced resonantly or non-resonantly  with mass in the range  $ m \\sim \\mathrm{few}\\,\\mathrm{keV} $ that decouple when ultrarelativistic lead to a primordial phase space density of the same order of magnitude as in cored dShps and disfavor cusped satellites for which the data \\cite{gilmore} yields a much larger phase space density. }   \\item{\\textbf{(ii):} CDM from wimps that decouple when non-relativistic with $ m \\gtrsim 100~\\mathrm{GeV} $ and kinetic decoupling at $ T_d \\sim 10~\\mathrm{MeV} $ \\cite{dominik} yield phase space densities at least   eighteen to fifteen orders of magnitude  [see eqs.(\\ref{rss}), (\\ref{cusp}) and (\\ref{denw})] larger than the typical  average in dSphs \\cite{gilmore}. Results from $N$-body simulations, albeit with initial conditions with much smaller values of the phase space density,  yield a dynamical relaxation by a factor $ 10^2-10^3 $ \\cite{numQ}. If these results are confirmed by simulations with larger initial values there may be a potential tension between the primordial phase space density for \\emph{thermal relics} in the form of WIMPs with $ m\\sim 100\\,\\mathrm{GeV},  \\; T_d \\sim 10 $ MeV  and those observed in dShps.}   \\item{\\textbf{(iii):} Light bosonic particles decoupled while ultrarelativistic and which form a BEC lead to phase space densities consistent with cores and also consistent with cusps if $ T_c/T_d \\gtrsim 10 $. However if these thermal relics  satisfy the observational bounds,  they must decouple when $ g_d \\; g^{-\\frac34} \\; (T_d/T_c)^\\frac98 > 130 $, namely {\\it above} the electroweak scale.  } \\end{itemize}  \\medskip  Section II analyzes the {  generic} dynamics of decoupled particles for {  any} distribution function, {  with or without} LTE at decoupling, and for {\\bf different} species of particles. In section III we consider light thermal relics which decoupled in LTE as DM components: fermions and  bosons, including the possibility of a Bose-Einstein condensate. Section IV deals with coarse grained phase space densities which are Liouville invariant and the new bounds obtained with them by using the observational dSphs data and recent results from $N$-body simulations, bounds from velocity dispersion, and the generalized Gunn-Tremaine bound. In Section V we study the case of particles that decoupled {  out of equilibrium} and the consequences  on the dark matter constraints. Section VI summarizes our conclusions. ",
        "conclusions": "We have obtained new constraints on light DM candidates that decoupled while ultrarelativistic in or out of LTE in terms of their distribution functions. The only assumption is that these distribution functions are homogeneous  and isotropic. A Liouville invariant coarse grained primordial phase space density is introduced that allows to combine phase space density arguments with a recent compilation of photometric and kinematic data on dSphs galaxies to yield \\emph{new constraints} on the mass, velocity dispersion and phase space density of DM candidates. The new constraint on the mass range is \\be \\frac{62.36~\\mathrm{eV}}{\\mathcal{D}^{\\frac{1}{4}}} \\; \\Bigg[10^{-4}  \\; \\frac{\\rho}{\\sigma^3} \\; \\frac{\\big(\\mathrm{km}/\\mathrm{s}\\big)^3}{M_\\odot/\\mathrm{kpc}^3} \\Bigg]^{\\frac14} \\leq   m \\leq 2.695 \\; \\frac{2  \\; g_d \\; \\zeta(3)}{g \\; \\int^\\infty_0 y^2  \\; f_{d }(y) \\; dy} \\; \\mathrm{eV}   \\; , \\ee where the primordial phase space density is given by \\be \\mathcal{D} = \\frac{g}{2 \\; \\pi^2} \\frac{\\Bigg[\\int_0^\\infty y^2  \\; f_d(y)  \\; dy \\Bigg]^{\\frac52}}{\\Bigg[\\int_0^\\infty y^4  \\; f_d(y) \\; dy \\Bigg]^{\\frac32}} \\; , \\ee $ f_d(p_c/T_d) $ is the distribution function at decoupling, $ g $ the number of internal degrees of freedom of the particle, and $ \\rho/\\sigma^3 $ is the phase space density obtained from observations. The \\emph{upper bound} arises from requesting that the DM candidate has a density $ \\leq \\rho_{DM} $ today, and the \\emph{lower} bound arises from requesting that the phase space density in halos $ \\rho/\\sigma^3 $ be \\emph{smaller} than or equal to the primordial phase space density of the collisionless non-relativistic (today) DM component $$ \\rho_{DM}/\\sigma^3_{DM}=3^\\frac32  \\; m^4  \\; \\mathcal{D} . $$ We have studied the consequences of Bose-Einstein condensation of light ultrarelativistic particles when chemical freeze out occurs well before kinetic decoupling at $ T_d < T_c $ with $ T_c $ the critical temperature below which a non-vanishing condensate fraction exists. We find that the presence of the condensate hastens the onset of the non-relativistic regime and that Bose-Einstein condensed particles can effectively act as a CDM component {\\it even} when they decoupled being ultrarelativistic. The reason for this unusual behavior is that the   particles in the condensate all have vanishing velocity dispersion.  For \\emph{thermal relics} we find \\be \\mathcal{D} = g \\times \\left\\{  \\begin{array}{l} 1.963\\times 10^{-3}~~\\mathrm{Fermions},~\\mu_d=0 \\\\ 3.657\\times 10^{-3}~~\\mathrm{Bosons~no~BEC} \\\\ 3.657\\times 10^{-3}\\,\\Big(\\frac{T_c}{T_d} \\Big)^{ \\frac{15}{2}}~~\\mathrm{Bosons~with~BEC},~T_c>T_d  \\\\ 8.442\\times 10^{-2}\\, g_d\\,Y_\\infty~~\\mathrm{non-relativistic~Maxwell-Boltzmann.} \\end{array} \\right.\\label{LTDD2} \\ee The combination of data in ref. \\cite{gilmore} from dSphs when applied to \\emph{light thermal relics}  yields the mass range \\bea && \\frac{444~ \\mathrm{eV} }{g^{\\frac14}}  \\leq m \\leq \\frac{g_d}{g}~ 4.253 ~ \\mathrm{eV}  ~~\\mathrm{fermions~with}~ \\mu_d=0  \\; ,\\nonumber \\\\ && \\frac{ 380~\\mathrm{eV}}{g^{\\frac14}}  \\leq m \\leq \\frac{g_d}{g}~ 2.695 ~\\mathrm{eV}   ~~\\mathrm{bosons~with} ~ \\mu_d=0 \\; \\mathrm{and~no~BEC} \\; \\, \\nonumber \\\\   && \\frac{380~\\mathrm{eV}}{g^{\\frac14}}\\Bigg[\\frac{T_d}{T_c}\\Bigg]^\\frac{15}{8} \\leq m \\leq \\frac{g_d}{g}~ 2.695 ~\\Bigg[\\frac{T_d}{T_c}\\Bigg]^3  \\; \\mathrm{eV}  ~~\\mathrm{BEC} \\; . \\label{massrangeF2} \\eea with the implication that if these particles are suitable DM candidates, they must decouple at high temperature when the effective number of ultrarelativistic degrees of freedom is $ g_d > 100 $. Namely, in absence of a BEC, thermal decoupling must occur above the electroweak scale. In the BEC case, for $ T_d\\ll T_c $, the fulfillment of the bound requires very large $ g_d $. Namely, in the presence of a BEC thermal decoupling occurs at a scale much larger than the electroweak scale for $ T_d\\ll T_c $.  Assuming that the DM particle is the only component with the density $ \\rho_{DM} $ today, we obtained an independent bound from velocity dispersion which for the favored cored profiles \\cite{gilmore} yield the lower mass bound \\be \\frac{m}{\\mathrm{keV}} \\geq  \\frac{0.855}{g^{\\frac13}_d}\\, \\Bigg[ \\frac{\\int_0^\\infty y^4  \\; f_d(y) \\; dy}{\\int_0^\\infty y^2 \\;  f_d(y) \\; dy}\\Bigg]^{\\frac12} \\label{masdisp2} \\; . \\ee For light thermal relics this bound implies that $ m  \\gtrsim 0.6-1.5\\,\\mathrm{keV} $ with a suppression factor $ T_d/T_c $ in the BEC case.  For light thermal relics that decoupled while ultrarelativistic we find the primordial phase space density \\be \\frac{\\rho_{DM}}{\\sigma^3_{DM}} \\sim 10^6 ~\\frac{ \\mathrm{eV}/\\mathrm{cm}^3} {\\Big( \\mathrm{km}/\\mathrm{s} \\Big)^3}~\\Bigg( \\frac{m}{\\mathrm{keV}}\\Bigg)^3  \\; g_d \\; \\Bigg\\{\\begin{array}{l} 0.177~~~\\mathrm{Fermions} \\\\ 0.247~~~\\mathrm{Bosons ~ without ~ BEC} \\\\ 0.247\\,(T_c/T_d)^\\frac92 ~~~\\mathrm{Bosons ~ with  ~ BEC} \\; . \\end{array} \\ee An enhancement factor $ (T_c/T_d)^\\frac92 $ appears in the r.h.s. in the presence of a BEC.  For wimps with kinetic decoupling temperature $ 10 $MeV \\cite{dominik}, we find \\be\\label{wimpo} \\frac{\\rho_{wimp}}{\\sigma^3_{wimp}} \\sim 10^{24} \\; \\frac{ \\mathrm{eV}/\\mathrm{cm}^3}{\\Big( \\mathrm{km}/\\mathrm{s} \\Big)^3}~\\Bigg( \\frac{m}{100\\,\\mathrm{GeV}}\\Bigg)^3   \\; g_d  \\; . \\ee The observational data compiled in ref. \\cite{gilmore} assuming a favored cored profile suggests \\be \\left(\\frac{\\rho_s}{\\sigma^3_s}\\right)_{cored} \\sim 5\\times 10^6 ~\\frac{ \\mathrm{eV}/\\mathrm{cm}^3} {\\Big( \\mathrm{km}/\\mathrm{s}\\Big)^3}\\,. \\label{core2} \\ee If the distribution of dark matter is cusped, ref. \\cite{gilmore} gives the value for the density $ \\rho_s \\sim 2 \\; \\mathrm{TeV}/\\mathrm{cm}^3 $ yielding \\be \\left(\\frac{\\rho_s}{\\sigma^3_s}\\right)_{cusped} \\sim 2\\times 10^9 ~\\frac{ \\mathrm{eV}/\\mathrm{cm}^3} {\\Big( \\mathrm{km}/\\mathrm{s} \\Big)^3} \\,.\\label{cusp2} \\ee Therefore, for $ g_d \\gtrsim 10 $ the primordial phase space density for \\emph{thermal relics} with $ m \\sim \\mathrm{keV} $ favors a {\\bf cored distribution}.  Notice that a bosonic thermal relic that features a BEC can behave as CDM with small velocity dispersion and a primordial phase space density consistent with cusped distributions if $ T_d \\ll T_c $. However,  these BEC DM candidates must decouple at a temperature scale {\\it higher} than the electroweak.  Recent results from $N$-body simulations suggests that the phase space density relaxes by a factor $ \\sim 10^2 $ during gravitational clustering for $ 0 \\leq z \\leq 10 $ \\cite{numQ}. Combining these numerical results with the observational results on dSphs \\cite{gilmore} and the present DM density, we conclude that the mass of \\emph{thermal relics} that decoupled when ultrarelativistic is \\be \\label{masarango2} m_{cored}\\sim \\frac2{g^\\frac14} \\; \\mathrm{keV} \\quad , \\quad m_{cusp} \\sim\\frac8{g^\\frac14} \\; \\mathrm{keV} \\; . \\ee The decoupling temperature for the DM candidate that would favor cusped profiles must be near a grand unified scale for a large symmetry group with $g_d \\gtrsim 2000$ which effectively results in a colder relic today with a far smaller velocity dispersion.  \\medskip  The \\emph{enormous} discrepancy between the primordial phase space density for WIMPs of $ m\\sim 100\\,\\mathrm{GeV}; \\; T_d \\sim 10\\, \\mathrm{MeV} $, eq.(\\ref{wimpo}) and the phase space densities in dSphs, either cored (eq.\\ref{core2} ) or cusped (eq.\\ref{cusp2}) cannot   be explained by the two orders of magnitude of gravitational relaxation of phase space densities found with recent $N$-body simulations \\cite{numQ}, although these initialize the simulation with much smaller values of the primordial phase space density.  \\medskip  We have studied a scenario for decoupling out of equilibrium motivated by previous studies of particle production and thermalization via an UV cascade. The distribution function obtained from previous studies \\cite{dvd}, remarkably describes the non-equilibrium distribution functions for sterile neutrinos produced either resonantly \\cite{este} or non-resonantly \\cite{dw} as well as a recently proposed model for halo structure \\cite{strigari}. Our bounds in terms of arbitrary distribution functions lead to the following bounds on the mass, phase space density and velocity dispersion of these light relics that decoupled out of LTE:  \\begin{itemize}  \\item{For sterile neutrinos produced non-resonantly via the Dodelson-Widrow mechanism \\cite{dw}   we find \\be \\frac{1.04 \\; \\mathrm{keV}}{g^\\frac14} \\leq m \\leq \\frac{g_d}{g} \\;  46.5 \\, \\mathrm{eV} \\quad , \\quad \\frac{\\rho_{DM}}{\\sigma^3_{DM}} = 0.57 \\; g \\times 10^5 \\; \\Big[\\frac{m}{\\mathrm{keV}}\\Big]^3 \\; \\frac{M_\\odot/\\mathrm{kpc}^3}{\\big(\\mathrm{km}/\\mathrm{s}\\big)^3} \\quad , \\quad \\sigma_{DM} =   \\frac{0.187 }{g^{\\frac13}_d} \\; \\Big(\\frac{\\mathrm{keV}}{m}\\Big) \\; \\Big(\\frac{\\mathrm{km}}{s}\\Big) \\; . \\label{dw2} \\ee The upper and lower bound on the mass can only be compatible if the sterile neutrino decouples with $g_d \\gtrsim 20-30$. For $m \\sim \\mathrm{keV}$ the primordial phase space density is compatible with cored but not with cusped profiles in the dShps data \\cite{gilmore}. Combining these bounds with the results from $N$-body simulations on the relaxation of the phase space density \\cite{numQ} and with the observational constraint eq.(\\ref{rss}) \\cite{gilmore}, we obtain the value \\be m\\sim \\frac{4}{g^{\\frac13}} \\;   \\mathrm{keV} \\label{boundste2} \\ee for the mass of sterile neutrinos produced non-resonantly by the Dodelson-Widrow mechanism.  } \\item{For  sterile neutrinos produced by a net-lepton number driven resonant conversion \\cite{este} we find \\be \\frac{289\\,\\mathrm{eV}}{g^\\frac14} \\leq m \\leq \\frac{g_d}{g}\\; 81.4\\; \\mathrm{eV} \\quad , \\quad \\frac{\\rho_{DM}}{\\sigma^3_{DM}} = 9.6 \\; g \\times 10^6\\,\\Big[\\frac{m}{\\mathrm{keV}}\\Big]^4 \\; \\frac{M_\\odot/\\mathrm{kpc}^3}{\\big(\\mathrm{km}/\\mathrm{s}\\big)^3} \\quad , \\quad \\sigma_{DM} = \\frac{0.028 }{g^{\\frac13}_d} \\; \\Big(\\frac{\\mathrm{keV}}{m}\\Big) \\; \\Big(\\frac{\\mathrm{km}}{s}\\Big) \\; . \\label{sterQ2} \\ee The small velocity dispersion is a consequence of the distribution function being skewed towards small momentum. Again for $ m \\sim \\mathrm{keV} $, the primordial phase space density is compatible with cored but not cusped profiles in the dShps data \\cite{gilmore}. For sterile neutrinos produced by resonant conversion, a similar analysis as for the  previous case yields \\be m\\sim \\frac{0.8}{g^\\frac14} \\; \\mathrm{keV} \\; . \\label{boundeste} \\ee } \\item{For the model  proposed in ref. \\cite{strigari} we find \\be \\frac{475 \\; \\mathrm{eV}}{(\\beta \\; g)^\\frac14} \\leq m \\leq \\frac{g_d}{\\beta g} \\; 1.94 \\; \\mathrm{eV} \\quad , \\quad \\frac{\\rho_{DM}}{\\sigma^3_{DM}} = 1.33 \\; \\beta \\; g \\times 10^6 \\; \\Big[\\frac{m}{\\mathrm{keV}}\\Big]^4 \\; \\frac{M_\\odot/\\mathrm{kpc}^3}{\\big(\\mathrm{km}/\\mathrm{s}\\big)^3} \\quad , \\quad \\sigma_{DM} =   \\frac{0.187}{g^{\\frac13}_d} \\; \\Big(\\frac{\\mathrm{keV}}{m}\\Big) \\; \\Big(\\frac{\\mathrm{km}}{s} \\Big)\\; . \\label{bounstri}\\ee } \\end{itemize} It is noteworthy that the $N$-body results of ref. \\cite{numQ} which yield phase space relaxation by a factor $ \\sim 10^2 $ bring the values of the primordial phase space density of the above cases within the range consistent with the phase space densities for cored profiles in dSphs \\cite{gilmore} for $ m \\sim \\mathrm{keV} $. On the contrary, in the   case of WIMPs with $m\\sim 100\\,\\mathrm{GeV},T_d\\sim 10\\,\\mathrm{MeV}$, relaxation by {\\bf many} orders of magnitude is necessary for their phase space densities to be compatible with the observed values both for cores and for cusps.  Therefore the  bounds  eqs.(\\ref{boundste2})-(\\ref{boundeste}) confirm that $ \\sim \\mathrm{keV} $ relics that decouple out of equilibrium while ultrarelativistic via the mechanisms described above yield values for phase space densities that are in agreement with cores in the DM distribution.  \\medskip  The results obtained in this article for the new mass bounds, primordial phase space densities and velocity dispersion in term of arbitrary, but homogeneous and isotropic distribution functions establish a link between the microphysics of decoupling and observable quantities. They also warrant   deeper scrutiny  of the non-equilibrium aspects of sterile neutrinos\\cite{boyho} for a firmer assessment of their potential as DM candidates."
    },
    "0710/0710.4003_arXiv.txt": {
        "abstract": "The stellar evolution code YREC is outlined with emphasis on its applications to helio- and asteroseismology. The procedure for calculating calibrated solar and stellar models is described. Other features of the code such as a non-local treatment of convective core overshoot, and the implementation of a \\cws{parametrized} description of turbulence in stellar models, are considered in some detail.  The code has been extensively used for other astrophysical applications, some of which are briefly mentioned at the end of the paper. ",
        "introduction": "The aim of this paper is to provide an overview of the Yale Rotating Stellar Evolution Code (YREC), as it has been applied in the last few years to research in helio- and asteroseismology. Although YREC contains extensions to model the effects of rotation in an oblate coordinate system, we describe here the ``non-rotating'' version.  In addition to a general description, we shall emphasize three features of the code which have been implemented because of their special relevance to seismology.  The first feature is the procedure utilized for the automatic calculation of calibrated solar and stellar models whose pulsational properties are to be investigated.  The second feature is the treatment of convective core overshoot. Finally, the third feature is the implementation in stellar models of the effects of turbulence on the structure of the surface layers of stars with a convective envelope.  The \\cws{parametrization} of turbulence to one dimension is based on three-dimensional radiative hydrodynamical (3D HRD) simulations of the highly superadiabatic layer (SAL) in the atmosphere.  The interaction of turbulent convection and radiation in these thin transition regions is poorly known.  Oscillation frequencies are sensitively affected by the structure of transition regions between radiative and convective layers.  Seismology thus offers a unique opportunity to explore a long standing problem in stellar physics.  Like most stellar evolution codes, YREC is a continuously evolving research tool to which many have contributed. As a result, different versions of YREC are in use at several institutions, which have been applied to a variety of research purposes. Some of the most significant applications of YREC are listed in the text and at the end of this paper (see Sect.~\\ref{other}). The rotating version of YREC, originally developed by \\citet{1988PhDT........14P}, includes a 1.5D treatment of rotation, extending the work of \\citet{1970stro.coll...20K} and \\citet{1981ApJ...243..625E}, and using the formalism of \\citet{1980PhDT.........5L}. A 2D version of YREC has also recently been implemented, specifically to address some fundamental aspects of solar magnetic activity \\citep{2006ApJS..164..215L}.  Sect.~\\ref{num} outlines the numerical scheme adopted to solve the classical differential equations of stellar structure and evolution. The treatment of the boundary conditions, of special importance for seismology, are described in Sect.~\\ref{bc}.  The constitutive physics, i.e. the equation of state and radiative and conductive opacities, are reviewed in Sect.~\\ref{eos}, and the nuclear processes are described in Sect.~\\ref{nuc}.  Stellar physics topics such as superadiabatic convection, element diffusion, convective core overshoot, and turbulence in the outer layers, all of which also have important seismological signatures, are covered in Sect.~\\ref{sbc}, Sect.~\\ref{diff}, Sect.~\\ref{ov} and Sect.~\\ref{turb}, respectively.  The operation of the code is described in Sect.~\\ref{runyrec}, with emphasis on helio- and asteroseismic applications. Seismic diagnostics applications are described in Sect.~\\ref{dia}. The role played by YREC in the research on solar neutrinos and helio-seismology is summarized in Sect.~\\ref{solar}. Studies of advanced evolutionary phases and applications to stellar population studies are listed in Sect.~\\ref{other}. ",
        "conclusions": ""
    },
    "0710/0710.4145_arXiv.txt": {
        "abstract": "Precision infrared photometry from Spitzer has enabled the first direct studies of light from extrasolar planets, via observations at secondary eclipse in transiting systems. Current Spitzer results include the first longitudinal temperature map of an extrasolar planet, and the first spectra of their atmospheres. Spitzer has also measured a temperature and precise radius for the first transiting Neptune-sized exoplanet, and is beginning to make precise transit timing measurements to infer the existence of unseen low mass planets. The lack of stellar limb darkening in the infrared facilitates precise radius and transit timing measurements of transiting planets.  Warm Spitzer will be capable of a precise radius measurement for Earth-sized planets transiting nearby M-dwarfs, thereby constraining their bulk composition.  It will continue to measure thermal emission at secondary eclipse for transiting hot Jupiters, and be able to distinguish between planets having broad band emission {\\it vs.} absorption spectra. It will also be able to measure the orbital phase variation of thermal emission for close-in planets, even non-transiting planets, and these measurements will be of special interest for planets in eccentric orbits. Warm Spitzer will be a significant complement to Kepler, particularly as regards transit timing in the Kepler field.  In addition to studying close-in planets, Warm Spitzer will have significant application in sensitive imaging searches for young planets at relatively large angular separations from their parent stars. ",
        "introduction": "The Spitzer Space Telescope (\\citet{Werner2004}) was the first facility to detect photons from known extrasolar planets (\\citet{Charbonneau2005, Deming2005}), inaugurating the current era wherein planets orbiting other stars are being studied directly. Cryogenic Spitzer has been a powerful facility for exoplanet characterization, using all three of its instruments. Spitzer studies have produced the first temperature map of an extrasolar planet (\\citet{Knutson2007a}), and the first spectra of their atmospheres (\\citet{Grillmair2007, Richardson2007}).  Spitzer will continue to study exoplanets when its store of cryogen is exhausted.  `Warm Spitzer' (commencing $\\sim$ spring 2009) will remain at T $\\sim$~35K (passively cooled by radiation), allowing imaging photometry at 3.6 and 4.5 $\\mu$m, at full sensitivity. The long observing times that are projected for the warm mission will facilitate several pioneering exoplanet studies not contemplated for the cryogenic mission. ",
        "conclusions": ""
    },
    "0710/0710.0370_arXiv.txt": {
        "abstract": "This paper describes the MegaPipe image processing pipeline at the Canadian Astronomical Data Centre. The pipeline combines multiple images from the MegaCam mosaic camera on CFHT and combines them into a single output image.  MegaPipe takes as input detrended MegaCam images and does a careful astrometric and photometric calibration on them.  The calibrated images are then resampled and combined into image stacks.  The astrometric calibration of the output images is accurate to within 0.15 arcseconds relative to external reference frames and 0.04 arcseconds internally.  The photometric calibration is good to within 0.03 magnitudes.  The stacked images and catalogues derived from these images are available through the CADC website. ",
        "introduction": "The biggest barrier to using archival MegaCam \\markcite{megacam}({Boulade} {et~al.} 2003) images is the effort required to process them. While the individual images are occasionally useful by themselves, more often the original scientific program called for multiple exposures on the same field in order to build up depth and get rid of image defects. Therefore, the images must be combined.  A typical program calls for 5 or more exposures on a single field. Each MegaCam image is about 0.7Gb (in 16-bit integer format), making image retrieval over the web tedious. Because of the distortion of the MegaPrime focal plane, the images must be resampled. This involves substantial computational demands. During this processing, which is often done in a 32-bit format, copies of the images must be made, increasing the disk usage. In summary, the demands in terms of CPU and storage are non-trivial.  Presumably Moore's law (1965), \\markcite{moore} will make these concerns negligible, if not laughable, in ten years time.  However, at the moment, they present a technological barrier to easy use of MegaCam data.  The Elixir pipeline \\markcite{elixir}({Magnier} \\& {Cuillandre} 2004) at CFHT processes each MegaCam image. It does a good job of detrending (bias-subtracting, flat-fielding, de-darking) the images. However, the astrometric solution Elixir provides is only good to 0.5-1.0 arcseconds. To combine the images, they must be aligned to better than a pixel. One arcsecond accuracy is insufficient. Therefore, it is necessary for a user to devise some way of aligning the images to higher accuracy. This is not an easy task, and rendered more difficult by the distortion of the MegaPrime focal plane. The problem is not intractable and there do exist a number of software solutions to the problem, but it remains an obstacle to easy use of MegaCam data.  In short, while the barriers to using archival MegaCam data are not insurmountable, they make using these data considerably less attractive. MegaPipe aims to increase usage of MegaCam data by removing these barriers.  This paper describes the MegaPipe image processing pipeline.  MegaPipe combines MegaCam images into stacks and extracts catalogues.  The procedure can be broken down into the following steps: \\begin{itemize} \\item Image grouping (Section \\ref{grouping}) \\item Astrometric calibration (Section \\ref{astrom}) \\item Photometric calibration (Section \\ref{photom}) \\item Image stacking (Section \\ref{comb}) \\item Catalogue generation (Section \\ref{cat}) \\end{itemize} Sections \\ref{qualastro} and \\ref{qualphoto} discuss checks on the astrometric and photometric properties of the output images. Section \\ref{sec:prod} describes the production and distribution of the images. ",
        "conclusions": ""
    },
    "0710/0710.5591.txt": {
        "abstract": "We study X-ray spectra of Cyg X-3 from \\sax, taking into account absorption and emission in the strong stellar wind of its companion. We find the intrinsic X-ray spectra are well modelled by disc blackbody emission, its upscattering by hot electrons with a hybrid distribution, and by Compton reflection. These spectra are strongly modified by absorption and reprocessing in the stellar wind, which we model using the photoionization code {\\tt cloudy}. The form of the observed spectra implies the wind is composed of two phases. A hot tenuous plasma containing most of the wind mass is required to account for the observed features of very strongly ionized Fe. Small dense cool clumps filling $\\la$0.01 of the volume are required to absorb the soft X-ray excess, which is emitted by the hot phase but not present in the data. The total mass-loss rate is found to be (0.6--$1.6)\\times 10^{-5}\\msun$ yr$^{-1}$. We also discuss the feasibility of the continuum model dominated by Compton reflection, which we find to best describe our data. The intrinsic luminosities of our models suggest that the compact object is a black hole. ",
        "introduction": "Cyg X-3 is a high-mass X-ray binary (XRB) system with a short orbital period, $P=4.8$ h. It is located at a distance of $d\\sim 9$ kpc in the Galactic plane (Dickey 1983, assuming the 8 kpc distance to the Galactic Centre; Predehl et al.\\ 2000). Due to strong interstellar absorption, its optical counterpart is not observable, though infrared observations indicate it is a Wolf-Rayet (WR) star (van Kerkwijk et al.\\ 1996). In spite of its discovery already in 1966 (Giacconi et al.\\ 1967), it remains poorly understood. In particular, it remains uncertain whether its compact object is a black hole or a neutron star. Currently, there are two other known X-ray binaries containing WR stars, IC 10 X-1 (Prestwich et al.\\ 2007) and NGC 300 X-1 (Carpano et al.\\ 2007a, b). In both cases, there is dynamical evidence that the compact object is a black hole.  Cyg X-3 is the brightest radio source among X-ray binaries (Mccollough et al.\\ 1999). It shows very strong radio outbursts and resolved jets (Marti, Paredes \\& Peracaula 2000; Mioduszewski et al.\\ 2001). In the X-rays, it exhibits a wide range of variability patterns. In particular, transitions between the hard and soft spectral state occur on the timescale of months to years.  The understanding of the radiative processes underlying the X-ray spectra of Cyg X-3 remains rather rudimentary. Zdziarski \\& Gierlinski (2004) have shown an overall similarity of the spectral states of Cyg X-3 to canonical states of black hole binaries. However, details, in particular the energy of the cutoff in the hard state, appear to differ significantly. Many different models have been fitted to spectra of Cyg X-3 by various authors, but each of them appeared to fit the data well, which precluded determination of the correct one. The models assumed different system components, geometry, radiative processes and the absorbing medium. They included different combinations of absorbers, Gaussian lines, blackbody, power-law with or without a cutoff, and bremsstrahlung, and different model combinations were proposed for the hard and soft states, see, e.g., White \\& Holt (1982), Nakamura et al.\\ (1993), Rajeev et al.\\ (1994). Later, the power-law models were replaced by more physical models of Comptonization, see Vilhu et al.\\ (2003), Hjalmarsdotter et al.\\ (2004) and Hjalmarsdotter et al.\\ (2008, hereafter H08). The interpretation of the intrinsic unabsorbed spectra, remains, however, not clear. In particular, H08 show that {\\it INTEGRAL\\/} data of Cyg X-3 in the hard state could be well modelled by three models with very different shapes of the intrinsic continuum,  yet the model most promising from a physical point of view yielded the worst statistical fit.  A key issue in the search for understanding the spectra of Cyg X-3 appears to be the nature of the complex intrinsic absorption, most likely caused by the wind from the companion star. In this paper, we study the effects of the WR stellar wind on the X-ray spectra. In Section \\ref{data}, we present the data used for our study. Then, we discuss our theoretical models in Section \\ref{s:model}. In Section \\ref{results}, we describe the complex interactions between the X-rays and the wind, and present results of our fits to the data. Sections \\ref{discussion} and \\ref{conclusions} contain a discussion of our results and our conclusions, respectively. ",
        "conclusions": "\\label{conclusions}  There are a number of essential features of our model which are new in the history of the studies of Cyg X-3. We have used a realistic model of the WR star, with a heavy small inner core and dense wind, and with the abundances typical to helium stars. We have used an accurate model of Comptonization together with accurate radiative transfer calculations to model the intrinsic X-ray emission in the vicinity of the compact object reprocessed by the surrounding wind. We have used high quality \\sax\\/ data, which are characterized by both relatively good spectral resolution and broad-band coverage.  We have used the X-rays emission as a probe of the wind properties. We have determined the wind mass-loss rate to be (0.6--$1.6)\\times 10^{-5}\\msun$ yr$^{-1}$, which is consistent with mass-loss rates of WR stars. We have disproved a previous claim that if Cyg X-3 contains a WR star its wind would be completely opaque to X-rays. We have also determined the structure of the wind illuminated by the X-rays, and found it contains both a hot phase (containing most of its mass) and cold dense clumps, occupying less than 1 per cent of the wind volume. The hot phase was needed to account for the strong $\\sim$9 keV edge seen in the data, and the cold phase was needed to absorb a significant soft-excess emission produced by the hot phase. We have also found that the density enhancement in the clumps averaged over both phases is consistent with the clumping factor found in WR stellar winds using other methods. The lines emitted by the wind found in the \\sax\\/ data are consistent with the high-resolution measurement by {\\it Chandra}.  Our fits imply that the observed X-ray continua contain very strong Compton reflection components. This would point out to the presence of a geometrically thick torus in the accretion flow, which would obscure most of the intrinsic X-ray source but allow us to see its reflection. The reflecting material is found to be weakly ionized (from the presence of an edge at $\\sim$7 keV), but no associated fluorescent Fe K$\\alpha$ (at 6.4 keV) is found. This presents us with a puzzle remarkably similar to that encounted in a number of Narrow-Line Seyfert 1 galaxies. On the other hand, the finding of the dominance of reflection may be an artifact of not including enough details in our spectral modelling. In particular, an additional partial covering by a weakly ionized medium could be responsible for the $\\sim$7-keV edge in the data. Based on the bolometric luminosity of the inferred intrinsic continua, the compact object in Cyg X-3 most likely a black hole."
    },
    "0710/0710.4623_arXiv.txt": {
        "abstract": "{ We present the analysis and results of 12.5 hours of high-energy gamma-ray observations of the EGRET-detected pulsar PSR B1951+32 using the Solar Tower Atmospheric Cherenkov Effect Experiment (STACEE). STACEE is an atmospheric Cherenkov detector, in Albuquerque, New Mexico, that detects cosmic gamma rays using the shower-front-sampling technique. STACEE's sensitivity to astrophysical sources at energies around 100 GeV allows it to investigate emission from gamma-ray pulsars with expected pulsed emission cutoffs below 100 GeV. We discuss the observations and analysis of STACEE's PSR 1951+32 data, accumulated during the 2005 and 2006 observing seasons. }  \\begin{document} ",
        "introduction": "Young energetic pulsars (rapidly rotating, highly magnetized neutron stars, that produce non-thermal photons) remain as yet the only Galactic objects unambiguously detected within the EGRET energy range. To date, more than 1500 pulsars have been observed at radio energies \\cite{manchester05}. About 70 of these are seen in X-rays \\cite{kaspi06} but only 6 (Crab, Geminga, Vela, PSR B1951+32, PSR1706-44, and PSR B1055-52) were detected by EGRET \\cite{nolan96}. Of these, only one, PSR B1951+32, has been observed up to $\\sim$20~GeV without any apparent cutoff in its differential energy spectrum \\cite{ramanamurthy95}. Accordingly, PSR 1951+32 is often considered the best pulsar candidate for pulsed-emission detection in the very-high-energy (VHE) regime; i.e.\\ at energies above $\\sim$50~GeV.  As is the case for most pulsars, PSR B1951+32 was first detected at radio energies. It was discovered with a 39.5~ms period in the radio synchrotron nebula CTB 80 \\cite{kulkarni98}. From the radio observations it was deduced that the pulsar has a characteristic age of 1.1$\\times$10$^5$~yr with a surface magnetic field of 4.9$\\times$10$^{11}$~G, and a rotational energy loss rate of 3.7$\\times$10$^{36}$~ergs~s$^{-1}$. It has also been observed in X-rays. While the pulsar emits a single pulse at radio frequencies and in X-rays, the EGRET observations display a double-pulsed profile with neither of the gamma-ray peaks coinciding with the radio peak. Interestingly, there is no evidence for any interpeak emission in the gamma-ray data \\cite{ramanamurthy95}.  At TeV energies, only upper limits on the pulsed emission, and on the emission from the surrounding synchrotron nebula, exist \\cite{srinivasan97}. A recent report by the MAGIC collaboration \\cite{albert07} presents observations above 75~GeV with no evidence for pulsed emission. %  Models that attempt to explain the non-thermal high-energy pulsed emission from pulsars generally fall into one of two broad categories; the Polar Cap \\cite{muslimov03} or Outer Gap \\cite{hirotani01} models. While both models can explain the observed gamma-ray emission at EGRET energies, they differ in their predictions for detectable emission above 20~GeV. In both scenarios the pulsed emission is attributed to particle acceleration in the pulsar's magnetosphere, with spectral cutoffs predicted in the 20-100~GeV energy range.  The Polar Cap model localizes the emission site to a region close to the magnetic poles of the neutron star, where the magnetic field is strong. The Outer Gap model, on the other hand, contends that gamma-ray production occurs far from the neutron star surface in a region of relatively weak magnetic field, in so called ``outer gaps'' near the null surface of the outer magnetosphere.  The emission sites dictate the energy of the expected spectral cutoff, insofar as the maximum energy of the curvature-radiated photons escaping the magnetosphere is limited by pair-production in the pulsar's magnetic field. Since the magnetic field is stronger near the polar cap, the Polar Cap model anticipates a lower-energy cutoff than the Outer Gap model.  Any detection of TeV emission would clearly favor the Outer Gap model and would thereby significantly contribute to our understanding of pulsar emission processes. Thus far, no pulsed TeV emission has been detected from any pulsar, although ever-lower upper limits are constraining the emission models. As a gamma-ray observatory operating at the lower end of the VHE regime (around 100 GeV) the Solar Tower Atmospheric Cherenkov Effect Experiment (STACEE) is suited to pulsar observations. STACEE observations of PSR B1951+32, undertaken in 2005 and 2006, are reported here. ",
        "conclusions": "No evidence for pulsed gamma-ray emission above 117 GeV was found in the selected gamma-ray data set used in this work (Figure 1). To calculate flux upper limits, we used the method of Helene \\cite{helene83} to estimate 99.0\\% upper limits for excess events within the pulsed phase profile seen by EGRET at lower energies \\cite{ramanamurthy95}. That is, emission is assumed to occur in the phase range of the main pulse, phase 0.12--0.22, and the interpulse, phase 0.48--0.74. A differential flux upper limit of 4.52~$\\times 10^{-6}$~MeV~cm$^{-2}$~s$^{-1}$ was determined at the energy threshold of 117 GeV, by extrapolation of the EGRET spectrum to STACEE energies.    \\bigskip {\\bf Acknowledgements:} Many thanks go to the staff of the National Solar Tower Test Facility, who have made this work possible. This work was funded in part by the US National Science Foundation, the Natural Sciences and Engineering Research Council of Canada, Fonds Quebecois de la Recherche sur la Nature et les Technologies, the Research Corporation, and the University of California at Los Angeles."
    },
    "0710/0710.0402.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary % {The cosmic X-ray background is the summed radiation from the growth of the massive black holes that reside in the centres % of galaxies.  By studying the population of sources that produce the X-ray background through X-ray surveys at various depths, % we can determine when, where and how this growth occurs.  } % aims heading (mandatory) {} {X-ray sources at intermediate fluxes (a few $\\times 10^{-14}\\, {\\rm erg\\, cm^{-2}\\, s^{-1}}$) with sky density of $\\sim 100\\, {\\rm deg}^{-2}$, are responsible for a significant fraction of the cosmic X-ray background at various energies below 10 keV. The aim of this paper is to provide an unbiased and quantitative description of the X-ray source population at these fluxes and in various X-ray energy bands.} % methods heading (mandatory) {We present the XMM-Newton Medium sensitivity Survey (XMS), including a total of 318 X-ray sources found among the serendipitous content of 25 XMM-Newton target fields. The XMS comprises four largely overlapping source samples selected at soft (0.5-2 keV), intermediate (0.5-4.5 keV), hard (2-10 keV) and ultra-hard (4.5-7.5 keV) bands, the first three of them being flux-limited. } % results heading (mandatory) {We report on the optical identification of the XMS samples, complete to 85-95\\%. At the flux levels sampled by the XMS we find that the X-ray sky is largely dominated by Active Galactic Nuclei. The fraction of stars in soft X-ray selected samples is below 10\\%, and only a few per cent for hard selected samples. We find that the fraction of optically obscured objects in the AGN population stays constant at around 15-20\\% for soft and intermediate band selected X-ray sources, over 2 decades of flux. The fraction of obscured objects amongst the AGN population is larger ($\\sim 35-45\\%$) in the hard or ultra-hard selected samples, and constant across a similarly wide flux range. The distribution in X-ray-to-optical flux ratio is a strong function of the selection band, with a larger fraction of sources with high values in hard selected samples.  Sources with X-ray-to-optical flux ratios in excess of 10 are dominated by obscured AGN, but with a significant contribution from unobscured AGN. } % conclusions heading (optional), leave it empty if necessary {} ",
        "introduction": "Supermassive black holes (SMBHs, i.e., with masses $\\sim 10^6-10^9\\, {\\rm M}_{\\odot}$) have been detected in the centers of virtually all nearby galaxies \\citep{Merrit01,Tremaine02}. In many of these galaxies -including our own-, the SMBH is largely dormant, i.e., the luminosity is many orders of magnitude below the Eddington limit. Only $\\sim 10\\%$ of today's galaxies (at most) host active galactic nuclei (AGN), and a very large fraction of them are in fact inconspicuous at most wavelengths because of obscuration \\citep{Fabian99}.  %SMBH %masses correlate well with host galaxy properties (stellar or gas %central velocity dispersion, bulge luminosity), and are estimated %to contain around 0.4-0.6\\% of the bulge masses %\\citep{Magorrian98}.  It is generally believed that the seeds of these SMBHs were the remnants of the first generation of massive stars in the history of the Universe.  These early black holes may have had masses of tens of ${\\rm M}_{\\odot}$ at most. The growth of these relic black holes to their current sizes is very likely dominated by accretion, with additional contributions by other phenomena like black hole mergers and tidal capture of stars \\citep{Marconi04}. According to current synthesis models, the integrated X-ray emission produced by the growth of SMBHs by accretion over the history of the Universe is recorded in the X-ray background (XRB). Thus the XRB can be used to constrain the epochs and environments in which SMBHs developed.  There are currently a number of existing or on-going surveys in various X-ray energy bands (see \\citet{Brandt05} for a recent compilation). In the pre-Chandra and pre-{\\it XMM-Newton} era the Einstein Extended Medium Sensitivity Survey \\citep{Maccacaro82,Gioia90,Stocke91} pioneered the procedure of determining typical X-ray to optical flux ratios for different classes of X-ray sources to facilitate the identification processes and has set the standards for serendipitous X-ray surveys. $ROSAT$ produced a number of surveys in the soft 0.5-2 keV X-ray band at various depths, e.g., the $ROSAT$ Bright Survey \\citep{Schwope00}, the intermediate flux RIXOS survey \\citep{Mason00} and the $ROSAT$ deep surveys \\citep{McHardy98,Georgantopoulos96,Hasinger98,Lehmann01} among others. These surveys show that AGN dominate the high Galactic latitude soft X-ray sky at virtually all relevant fluxes. The majority of these AGN are of spectroscopic type 1, which means that we are witnessing the growth of SMBH through unobscured lines of sight. In a moderate fraction of the sources identified, however, there is evidence for obscuration as their optical spectra lack broad emission lines (type 2 AGN).  \\citet{Ueda03} discuss the results from a large area X-ray survey in the 2-10 keV band with ASCA and those from HEAO-1 and $Chandra$, where a larger fraction of the sources identified correspond to type 2 AGN.  With {\\it Chandra} and {\\it XMM-Newton} coming into operation X-ray surveys, particularly at energies above a few keV, have been significantly boosted. Thanks to the high sensitivity and large field of view of the EPIC cameras \\citep{Turner01,Struder01} on board {\\it XMM-Newton} \\citep{Jansen01}, X-ray surveys requiring large solid angles have been dominated by this instrument.  The Bright Source Survey-BSS \\citep{Dellaceca04} contains 400 sources brighter than $\\sim 7\\times 10^{-14}\\, {\\rm erg}\\, {\\rm cm}^{-2}\\, {\\rm s}^{-1}$ either in 0.5-4.5 keV or 4.5-7.5 keV. The BSS samples\\footnote{{\\tt http://www.brera.mi.astro.it/\\~{}xmm/}}, which have been identified to $\\sim 90\\%$ \\citep{Caccianiga07}, show an X-ray sky dominated by AGN, where the fraction of obscured objects varies with the selection band (sample selection at harder energies reveals a higher fraction of obscured objects as expected).  Deep surveys have also been conducted by {\\it XMM-Newton}, for example in the Lockman Hole down to $\\sim 10^{-15} \\, {\\rm erg}\\, {\\rm cm}^{-2}\\, {\\rm s}^{-1}$ \\citep{Hasinger01,Mateos05b}. However, thanks to its much better angular resolution, the {\\it Chandra} deep surveys are photon counting limited and far from confusion and are consequently much more competitive at fainter fluxes \\citep{Alexander03,Tozzi06}. Optical identification of these deep surveys is largely incomplete, a fact that is driven by the intrinsic faintness and red colour of most of the counterparts to the faintest X-ray sources. In the intermediate flux regime, however, the identified fractions are large and nearing completion. It is interesting to note that deep surveys start to find a population of galaxies not necessarily hosting active nuclei as an important ingredient. In addition, the AGN population is found to contain an important fraction of obscured objects.  The wide range of intermediate X-ray fluxes, between say $10^{-15}\\, \\, {\\rm erg}\\, {\\rm cm}^{-2}\\, {\\rm s}^{-1}$ and $10^{-13}\\, {\\rm erg}\\, {\\rm cm}^{-2}\\, {\\rm s}^{-1}$ have also been the subject of a number of on-going surveys. Besides bridging the gap between wide and deep surveys, intermediate fluxes sample the region around the break in the X-ray source counts \\citep{Carrera07}, and therefore their sources are responsible for a large fraction of the X-ray background. Among these, we highlight the {\\it XMM-Newton} survey in the well-studied (at many bands) COSMOS field, which covers $2\\, \\deg^2$ to fluxes $\\sim 10^{-15}\\, {\\rm erg\\,  cm^{-2}\\,  s^{-1}}$ \\citep{Hasinger07}. The optical identification is still on-going, reaching 40\\% \\citep{Brusa07}. At fluxes around $10^{-14}\\, {\\rm erg}\\, {\\rm cm}^{-2}\\, {\\rm s}^{-1}$, the HELLAS2XMM survey \\citep{Baldi02,Fiore03}, now extended to cover $1.4\\, \\deg^2$, contains over 220 X-ray sources, optically identified to 70\\% completeness \\citep{Cocchia07}.  Other surveys in this flux range include the {\\it XMM-Newton} survey in the Marano field \\citep{Krumpe07}, which is 65\\% identified over a modest solid angle of $0.28\\, \\deg^2$. Also the XMM-2dF survey (Tedds et al., in preparation), which contains almost 1000 X-ray sources optically identified in the Southern Hemisphere, is an important contributor in this regime. {\\it Chandra} has also triggered surveys at intermediate fluxes, most notably the {\\it Chandra} Multiwavelength Survey \\citep{Kim04a,Kim04b,Green04}, covering $1.7\\, \\deg^2$ and identified to $\\sim 40\\%$ completeness \\citep{Silverman05}.  In the realm of this variety of X-ray surveys that yield a qualitative picture of the X-ray sky, the {\\it XMM-Newton} Medium sensitivity Survey (XMS) discussed in this paper, finds its role in three important ways: a) it deals with very large samples, selected at various X-ray bands where {\\it XMM-Newton} is sensitive, from 0.5 to 10 keV; b) the samples that we consider have been identified almost in full, from 85\\% to 95\\% completeness and c) three out of the four samples that we explore are strictly flux limited in three energy bands (0.5-2 keV, 0.5-4.5 keV and 2-10 keV). Armed with these unique features, the XMS is a very powerful tool to derive a {\\it quantitative} characterization of the population of X-ray sources selected in various bands, and also to study and characterize minority populations, all of it at specific intermediate X-ray fluxes where a substantial fraction of the X-ray background below 10 keV is generated. The power of the XMS is enhanced by the fact that to some extent it is a representative sub-sample of the {\\it XMM-Newton} X-ray source catalogue 2XMM\\footnote{Pre-release under {\\tt http://xmm.vilspa.esa.es/xsa}}, containing 150,000 entries.  Specific goals that have driven the construction of the XMS whose results are presented in this paper include: a) quantify the fraction of stars versus extragalactic sources at intermediate X-ray fluxes and at different X-ray energy bands; b) quantify the fraction of AGN that are classified as obscured by optical spectroscopy at intermediate X-ray fluxes and for samples selected in different energy bands; c) find the redshift distribution for the various classes of extragalactic sources and compare soft and hard X-ray selected samples; d) study the distribution of the X-ray-to-optical flux ratio for the various classes of X-ray sources, also as a function of X-ray selection band. The X-ray spectral properties of the sources of the XMS were already discussed in \\citet{Mateos05a}.  Further goals that we will achieve with the XMS in forthcoming papers include: e) determine the fraction of ``red QSOs'' at intermediate X-ray fluxes and as a function of X-ray selection band; f) relate X-ray spectral properties (like photoelectric absorption) to optical colours of the counterpart; g) quantify the fraction of radio-loud AGN in the samples selected at various X-ray energies; h) construct Spectral Energy Distributions for the various classes of sources in the XMS.  Results on these further aspects will be presented in a forthcoming paper (Bussons-Gordo et al., in preparation).  The paper is organized as follows: in Section~\\ref{sec:XMS} we define the XMS along with the 4 samples that constitute it, including the X-ray source list; in Section~\\ref{sec:imaging} we discuss the multi-band optical imaging conducted on the {\\it XMM-Newton} target fields and the process for selecting candidate counterparts; this is continued in Section~\\ref{sec:identification} where we discuss the identification of the XMS sources in terms of optical spectroscopy, and list photometric and spectroscopic information on each XMS source. Section~\\ref{sec:XMSpopulations} presents the first scientific results from the XMS, specifically a description of the overall source populations, the fraction of stars in the various samples, the fraction of optically obscured AGN, and the X-ray to optical flux ratio of the different source populations. Section~\\ref{sec:conclusions} summarizes our main results.  To clarify the terminology used in this paper, an AGN not displaying broad emission lines in its optical spectrum is termed as type 2 or obscured, and type 1 or unobscured otherwise.  The property of being absorbed or unabsorbed refers only to the detection or not of photoelectric X-ray absorption. Throughout this paper, we used a single power law X-ray spectrum to convert from X-ray source count rate to flux in physical units, with a photon spectral index $\\Gamma=1.8$ for the XMS-S and XMS-X samples and $\\Gamma=1.7$ for the XMS-H and XMS-U samples. These are the average values obtained by \\citet{Carrera07}, which -as opposed to what we do here- used the specific value of $\\Gamma$ for each individual source and energy range. When computing luminosities, we also use the above spectra for K-correction and the concordance cosmology parameter values: $H_0=70\\, {\\rm km}\\, {\\rm s}^{-1}\\, {\\rm Mpc}^{-1}$, $\\Omega_m=0.3$ and $\\Omega_\\Lambda=0.7$. All quoted uncertainties in parameter estimates are shown at 90\\% confidence level for one interesting parameter. ",
        "conclusions": "\\label{sec:conclusions}  In this paper we have presented the {\\it XMM-Newton} Medium sensitivity Survey XMS, and extracted a number of robust quantitative conclusions about the population of high Galactic latitude X-ray sources at intermediate flux levels. We have argued that given the completeness of our identifications and the relatively large size of the XMS samples, these conclusions can be safely exported to a much larger X-ray source catalogue like 2XMM.  Our conclusions can be summarized as follows:  \\begin{enumerate}  \\item The high galactic latitude X-ray sky at intermediate flux levels is dominated by AGN, which includes type-1 and type-2 AGN as well as the so-called XBONG which are likely to host a low luminosity or obscured nucleus (or both).  The stellar content is less than 10\\% in soft X-ray selected samples, and drops to below 5\\% at around soft X-ray fluxes $\\sim 10^{-14}\\, {\\rm erg}\\, {\\rm cm}\\, {\\rm s}^{-1}$. The stellar content in hard X-ray selected samples does not exceed a few per cent at most. Selection in 0.5-4.5 keV produces intermediate results.  \\item Given the limited sensitivity of {\\it XMM-Newton} above a few keV -which is due to the roll over of effective area- current surveys conducted in the so-called ultra-hard band (4.5-7.5 keV) do not bring any new source population or any significant difference with respect to 2-10 keV selected surveys. Much longer exposure times would be needed to unveil any new heavily obscured population with {\\it XMM-Newton}.  \\item Obscured AGN represent $\\sim 20\\%$ of the soft X-ray selected population of AGN, all the way from  $\\sim 10^{-13}\\, {\\rm erg}\\, {\\rm cm}^{-2}\\, {\\rm s}^{-1}$ down to  $\\sim 10^{-15}\\, {\\rm erg}\\, {\\rm cm}^{-2}\\, {\\rm s}^{-1}$, with no compelling evidence for an increase of this fraction towards fainter fluxes within this range.  \\item Likewise, obscured AGN represent $\\sim 35\\%$ ($45\\%$ if all unidentified sources are obscured AGN) of the hard X-ray selected population of AGN, with no hint of an increase down to a hard X-ray flux  $\\sim 10^{-14}\\, {\\rm erg}\\, {\\rm cm}^{-2}\\, {\\rm s}^{-1}$.  \\item The fraction of X-ray sources with X-ray to optical flux ratio $>10$ (or $X/O>1$ using the notation of this paper) is a mere 3\\% in soft X-ray selected samples, but grows to 20\\% in hard X-ray selected samples.  \\item Those sources with $X/O>1$ are mostly obscured AGN, but a fraction of around  20\\% of them in the hard band are unobscured type-1 AGN. This means that $X/O>1$ alone cannot be used as a proxy for obscured X-ray sources.  \\end{enumerate}"
    },
    "0710/0710.3659_arXiv.txt": {
        "abstract": "}{}  \\textheight 240mm \\textwidth 170mm \\hoffset= 0mm \\voffset=0cm \\topmargin -2cm \\oddsidemargin 0mm \\evensidemargin 0mm   \\begin{document} \\fontsize{12}{12} \\selectfont \\title{Luminosity Dependence of the Quasar Clustering from SDSS DR5} \\author{Ganna Ivashchenko} \\date{}  43024 objects, which were primarily identified as quasars in SDSS DR5 and have spectroscopic redshifts were used to study the luminosity dependence of the quasar clustering with the help of two different techniques. The obtained results reveal that brighter quasars are more clustered, but this dependence is weak, which is in agreement with the results by Porciani \\& Norberg \\cite{porciani2006} and theoretical predictions by Lidz et al. \\cite{lidz}.  \\textbf{Key words:} quasars, large-scale structure, clustering, luminosity.  PACS numbers: 98.65.Aj, 98.54.-h, 98.62.Ve. ",
        "introduction": "\\indent \\indent Determining the distribution of extragalactic objects is one of the most important problems of modern cosmology, because they are the only tracers of the dark matter, except anisotropy of the CMB, that helps us to realize the matter distribution on the redshifts of about 1000. Unfortunately even in the local Universe it is difficult to see the most faint galaxies, and when we go farther to larger redshifts, the most part of objects we can observe there with the modern ground-based telescopes used for large surveys are quasars as the most luminous objects. The matter distribution that we could reconstruct with the help of quasars is just like a light sketch, but it is all we have for today.  The quasars are nonuniform objects: they have various luminosities and are on the different stages of their evolution. Thus the reasonable question arises: how their clustering could depend on their physical properties, for example on their luminosity? According to different numerical simulations of galaxy mergers that incorporate black hole growth, this dependence has to exist because more luminous quasars are considered to be born in denser environment. In the main part of such models, in which the host halo mass correlates with the instantaneous luminosity of the quasars (see e.g. \\cite{kaufman}), there should be a strong luminosity dependence of the quasar clustering. In the other type of such models, in which the host halo mass correlates with the peak luminosity of quasars, the luminosity dependence of the quasar clustering should be weaker because all of the quasars we see now, are considered to be similar objects but on different stages of their evolution (see \\cite{lidz} and references therein). It is worth to note, that the redshift distribution of the quasars is not the same for different luminosities due to possible evolution effects.  The largest quasar surveys are 2dF and SDSS. In contrast to 2dF, which is finished for today and has 2QZ catalogue as a result \\cite{2dF-XII}, SDSS \\cite{SDSS5} is in progress and the area covered by it is increasing. Only a part of objects primarily classified as quasars were justified by spectra analysis and were included into SDSS Quasar Catalogue IV \\cite{dr4}. However even photometric classification of quasars with color diagrams \\cite{phot} is sufficient for using these objects for statistical purposes \\cite{myers2, myers1}.  Some attempts to find luminosity dependence of the quasar clustering have been made with different samples. E.g. Adelberg \\& Steidel \\cite{adelberg}, who worked with their own survey, pointed to luminosity independent quasar clustering. 2dF-team, that studied 2QZ survey and found the little redshift evolution in the amplitude of the power spectrum \\cite{2dF-IV}, \\cite{2dF-XI} and significant increase in clustering amplitude at high redshifts \\cite{2dF-XIV, 2dF-II}, detected only marginal evidence for quasars with brighter apparent magnitudes to have a stronger clustering amplitude \\cite{2dF-IX}. Porciani and Norberg \\cite{porciani2006} also found weak luminosity dependence of the clustering in 2QZ survey. Furthermore they noted that samples with different redshifts show different trends in luminosity dependence: in the redshift bin $1.7<z<2.1$ the brightest quasars seem to be more clustered (have larger bias parameter); in the redshift bin $1.3<z<1.7$ the bias parameter seems to follow a U-shape; and the low redshift bin ($0.8<z<1.3$) does not show any particular trend. Myers et al. \\cite{myers2}, working with $\\sim$300,000 photometrically classified quasars from the 4th Data Release of the SDSS, detected no significant luminosity dependence and pointed out that a 3$\\sigma$ detection of such effect requires a sample several times larger.  The present work deals with the study of the luminosity dependence of the quasar clustering for quasars from the Fifth Data Release of SDSS. The sample and its peculiarities are described in Section 2. The results of estimation of the correlation length as a function of luminosity are presented in Section 3. As we do not have such large sample, which we need, according to Myers et al. \\cite{myers2}, for precise measurements of the correlation function, we used another technique, which is described in Section 4. Finally, Section 5 is dedicated to discussion of obtained results. ",
        "conclusions": "\\indent \\indent As we can see from the Table 2 to obtain any reliable results with the first method one need much larger samples, which are unavailable for today. Moreover the correlation length is the spatial scale of clustering on the one hand and a measure of the clustering amplitude on the other hand. Thus it is not easy to interpret the first technique results unambiguously. That is why the second method seems to be better for this purpose because it is more direct and does not require such large samples.  From the results of the second method we can say that brighter quasars reside in closer pairs than faint ones. Note, that such splitting of f(r) curves on larges scales could also be the consequence of the different redshift distributions of the quasars with different luminosities. But we know, that the bright quasars represent higher redshifts and the quasar density decreases with the redshift, thus the inverse effect would be present. Anyway the difference between the mean redshift for subsamples of quasars with different luminosities (see the last column of Table 1) within each redshift interval is negligibly small and cannot affect the results.  Anyway this problem requires further investigations and improvement of the methods. It would be interesting to compare for example, the same $f(r)$ function for the fifth nearest neighbour or cross-correlation with galaxies. The last possibility could increase the sample, but this could be applied only for the lower redshift intervals, than those used in the present work.  Summing up the obtained results one can speal about luminosity dependence of the quasar clustering, but this dependence is not strong, which is in agreement with the results by Porciani \\& Norberg \\cite{porciani2006} and theoretical predictions by Lidz et al. \\cite{lidz}. But even if this dependence is weak, we do not have to neglect it. Note that when we estimate e.g. the correlation length of the quasars on the low redshift we obtain the mean correlation length averaged over all the quasars with different luminosities. But on the high redshifts the obtained results correspond only to the correlation length for bright quasars and do not reflect the whole picture. That is why possible effects of luminosity dependence of quasar clustering should be taken into account when studying the redshift evolution of it. But even so one should keep in mind that the quasars (as another AGNs) do not reflect the whole distribution of extragalactic objects because quasars are considered to reside in the strongest peaks of the matter density. This fact could explain larger values of the correlation function for quasars than for galaxies (see \\cite{blake}, \\cite{gonzalez} for comparision)."
    },
    "0710/0710.1140_arXiv.txt": {
        "abstract": "We use narrowband imaging (\\fwhm{} $=70$\\,\\AA{}) to select a sample of emission line galaxies between $0.20 \\lesssim z \\lesssim 1.22$ in two fields covering 0.5 sq.~deg. We use spectroscopic follow-up to select a sub-sample of \\ha{} emitting galaxies at $z\\sim0.24$ and determine the \\ha{} luminosity function and star formation density at $z\\sim0.24$ for both of our fields. Corrections are made for imaging and spectroscopic incompleteness, extinction and interloper contamination on the basis of the spectroscopic data. When compared to each other, we find the field samples differ by $\\Delta \\alpha = 0.2$ in faint end slope and $\\Delta \\log [ L^* (\\mathrm{\\ergs}) ] = 0.2$ in luminosity. In the context of other recent surveys, our sample has comparable faint end slope, but a fainter $L^*$ turn-over. We conclude that systematic uncertainties and differences in selection criteria remain the dominant sources of uncertainty between \\ha{} luminosity functions at this redshift.  We also investigate average star formation rates as a function of local environment and find typical values consistent with the field densities that we probe, in agreement with previous results. However, we find tentative evidence for an increase in star formation rate with respect to the local density of star forming galaxies, consistent with the scenario that galaxy-galaxy interactions are triggers for bursts of star formation. ",
        "introduction": "It is now widely accepted that the amount of star formation in Universe as a whole has increased since the formation of the first galaxies, peaking around redshifts $z\\sim2-3$ and subsequently declining by a factor of ten \\citep[e.g.][and references therein]{Hopkins04}. Cosmic star formation history provides strong constraints on models of galaxy formation and evolution \\citep{Pei99,Somerville01}, because it directly traces the accumulation of stellar mass and metal fraction \\citep{Pei95,Madau96} to their present-day values \\citep{Cole01,Panter03}. Its rapid decline over the past 8 Gyr is consistent with ``downsizing'' scenarios \\citep{Cowie96} in which the more massive galaxies have produced their stellar mass at earlier times than the less massive galaxies \\citep{Heavens04,Juneau05,Thomas05,Fardal06}. The star formation history of the universe has also been used to constrain allowable stellar initial mass functions \\citep{Baldry03,Hopkins06} and cosmic supernova rates \\citep{GalYam04,Daigne06}.  Star forming galaxies exhibit a strong UV continuum courtesy of newly formed OB stars in sites of star formation. This newborn population can be inferred from the UV directly \\citep[e.g.][]{Treyer98,Lilly96} or through a host of indirect calibrators spread across the electromagnetic spectrum \\citep{RosaGonzalez02,Condon92,Schaerer00}. At low redshifts the most direct calibrator -- and of the optical calibrators the least affected by internal extinction -- is the \\ha{} recombination line, which emits when stimulated by ionising UV radiation \\citep[e.g.][]{Kennicutt98}.  Narrowband surveys at optical wavelengths have long been recognised as a powerful way of yielding large samples of emission line galaxies, including those selected by \\ha{} at redshifts $z\\lesssim0.4$ \\citep{Ly07,Pascual07,Jones01}. They are advantageous in that they select galaxies in exactly the same quantity that they seek to measure, and are optimised for the detection of the faint emission line signatures indicative of star formation. Narrowband surveys also have the advantage of a simplified selection function, with filters that probe only a very narrow redshift slice, thereby yielding a volume limited sample at a common distance. Many recent emission line surveys have targeted \\lya{} at high redshift \\citep{Ajiki03,Hu04,Rhoads04,Gawiser06}, as well as \\ha{}, \\hb{}, \\oiii{} and \\oii{} at lower redshifts \\citep{Fujita03,Hippelein03,Ly07}.  Here we describe a survey for \\ha{} emission line galaxies at $z\\sim0.24$, found as a by-product of the Wide Field Lyman Alpha Search \\citep[WFILAS;][]{Westra05,Westra06}. The resulting sample has been utilised to determine the \\ha{} luminosity function at $z\\sim0.24$ and its associated co-moving star formation density. In Section~\\ref{sec:candsel} we describe the selection of candidates using narrow- and broadband imaging. In Section~\\ref{sec:spectroscopy} we detail follow-up spectroscopy used to identify the nature of the emission and test completeness of the sample. In Section~\\ref{sec:lumfieSFD} we derive the \\ha{} luminosity function for galaxies at $z\\sim0.24$ and explore its variation with the local environment in Section~\\ref{sec:environment}. A summary and concluding remarks are made in Section~\\ref{sec:conclusionsLowz}.  Throughout this paper we assume a flat Universe with $(\\Omega_{\\rm m}, \\Omega_{\\Lambda}) = (0.3,0.7)$ and a Hubble constant $H_0 = 70$\\,\\kmsMpc. All quoted magnitudes are in the {\\it AB} system \\citep{Oke83}\\footnote{$m_{AB} = -2.5 \\log f_\\nu - 48.590$, where $m_{AB}$ is the {\\it AB} magnitude and $f_\\nu$ is the flux density in ergs\\,s$^{-1}$\\,cm$^{-2}$\\,Hz$^{-1}$}. ",
        "conclusions": "\\label{sec:conclusionsLowz}  In this paper we report the results of a survey for \\ha{} emitting galaxies at $z\\sim0.24$. We used two fields from the Wide Field Imager Lyman Alpha Search (WFILAS). It consists of imaging in three narrowband filters (\\fwhm{} $=70$\\,\\AA{}), an encompassing intermediate band filter (\\fwhm{} $=220$\\,\\AA{}), supplemented with broadband $B$ and $R$. The narrowband filters cover a redshift range of $0.23 \\lesssim z \\lesssim 0.26$ for \\ha{} galaxies. These galaxies were selected by having an excess flux in one of the narrowband over to the other two, while also being detected in the intermediate and broadband $R$ filters. This yielded a total of 707 candidate emission line galaxies (after the removal of stellar contaminants) for both fields.  Of the 372 and 335 candidates, we observed 301 and 255 through spectroscopic follow-up for the CDFS and S11 fields, respectively.  We have identified emission in 189 and 117 candidates and confirmed that around half of these galaxies are \\ha{} at $z\\sim0.24$. A significant number of galaxies were also found at $z\\sim0.21$ by means of their \\sii{} emission. Other galaxies found were \\oii{} and \\hb{}/\\oiii{} emitters at $z\\sim1.2$ and $z\\sim0.6-0.7$, respectively.  Through use of the spectroscopy, we refined our colour selection to account for galaxies with a single emission line, leading to a measure of the fraction of \\ha{} galaxies as a function of narrowband flux in both of these regions of the sky. We also used the spectroscopy to determine a generic extinction correction using the Balmer decrement.  We have determined the \\ha{} luminosity function at $z\\sim0.24$ separately for both of our fields after correcting for imaging and spectroscopic incompleteness, extinction and contamination from interlopers. We find small differences in their slope and turn-over luminosity while their normalisations were the same. When compared to recent \\ha{} surveys, there is remarkable agreement between the luminosity function of our CDFS field with that one the Fabry-Perot imaging survey of \\citet{Hippelein03}. Differences between our fields were of the order expected by cosmic variance but less than the scatter between the \\ha{} luminosity functions of recent surveys. We surmise that while cosmic variance is a major contributor to this scatter, it is differences in methodology between surveys (mainly differences in selection criteria) that dominate discrepancies between \\ha{} luminosity functions and its related observables at $z\\sim0.24$. A survey that covers $10-20\\times$ the volume of one of our fields is required to get the uncertainty due to cosmic variance to the levels of \\citet{Gallego95}.  We estimated the star formation density for both our fields to be $\\log \\dot{\\rho} = -1.93 ^{+0.08} _{-0.10}$ and $-2.24 ^{+0.11} _{- 0.14}$ ($\\dot{\\rho}$ in \\Msunyr{}) for the CDFS and S11 fields, respectively, down to our survey limit of $\\log F_\\mathrm{line} = -16.0$ ($F_\\mathrm{line}$ in \\lineunits{}) or $\\log L_\\mathrm{line} = 40.6$ ($L_\\mathrm{line}$ in \\ergs{}). These values are comparable to other surveys at this redshift when calculated to the same flux limit. Correcting for AGN would decrease these values by 0.02 to 0.04 depending on exactly how much of the \\ha{} flux is contributed by the active nucleus rather than by normal star formation.  Furthermore, we determined the star formation density in the hypothetical case where \\sii{} emitters at $z\\sim0.21$ were classified as \\ha{} to illustrate the problems associated with solely relying colour selections. The star formation density $\\log \\dot{\\rho}$ of the S11 field does not change by much (+0.02). On the other hand, the star formation density in the CDFS increases by 0.12, due to the large number of foreground \\sii{} galaxies at $z\\sim0.21$.  We explored the amount of star formation with respect to the local environment and found that the star formation rates were typical for the field galaxy densities probed, in agreement with the results of previous work. However, we also found tentative evidence of an increase in star formation rate per galaxy with increasing density of the star forming galaxies. This supports scenarios where merger events are triggers for enhanced star formation, provided it can be demonstrated to be occurring on the smallest scales. We explored this trend by examining the spatial distribution of our fields individually and found that it was largely attributable to one field. A formal study of the clustering statistics of this field is required to confirm this and will be the subject of a future study."
    },
    "0710/0710.4881_arXiv.txt": {
        "abstract": "Inclusive neutrino-nucleus cross sections are calculated using a consistent relativistic mean-field theoretical framework. The weak lepton-hadron interaction is expressed in the standard current-current form, the nuclear ground state is described with the relativistic Hartree-Bogoliubov model, and the relevant transitions to excited nuclear states are calculated in the relativistic quasiparticle random phase approximation. Illustrative test calculations are performed for charged-current neutrino reactions  on $^{12}$C, $^{16}$O, $^{56}$Fe, and $^{208}$Pb, and results compared with previous studies and available data. Using the experimental neutrino fluxes, the averaged cross sections are evaluated for nuclei of interest for neutrino detectors. We analyze the total neutrino-nucleus cross sections, and the evolution of the contribution of the different  multipole excitations as a function of neutrino energy. The cross sections for reactions of supernova neutrinos on $^{16}$O and $^{208}$Pb target nuclei are analyzed as functions of the temperature and chemical potential. ",
        "introduction": "Introduction}  Neutrino-nucleus reactions at low energies play an important role in many phenomena in nuclear and particle physics, as well as astrophysics. These reactions present extremely subtle physical processes, not only because they involve the weak interaction, but also because they are very sensitive to the structure of nuclear ground states and excitations, i.e. to the solution of the nuclear many-body problem that includes the strong and electromagnetic interactions. The use of microscopic nuclear structure models in a consistent theoretical framework is therefore essential for a quantitative description of neutrino-nucleus reactions~\\cite{Hay.99}. Detailed predictions of neutrino-nucleus cross sections are crucial for the interpretation of neutrino experiments, detection of neutrinos produced in supernova explosions and understanding the underlying nature of these explosions~\\cite{sns1}. Neutrino-nucleus reactions which occur in a type II supernova could also contribute to the nucleosynthesis~\\cite{Heg.05,Lan.06}, but more data on cross sections are necessary for a more complete understanding of this process, as well as the supernova dynamics.  Data on neutrino-nucleus cross sections have been obtained by the LSND~\\cite{Alb.95,Ath.97} and KARMEN~\\cite{Bod.92,Bod.94,Mas.98} collaborations, and at LAMPF~\\cite{Kra.92,Koe.92}, but only for $^{12}$C and $^{56}$Fe target nuclei. New experimental programs are being planned which will provide essential data on the neutrino-nucleus reactions, and also help to improve the reliability of present cross-section calculations. These include the spallation neutron source (SNS) at ORNL, where neutrinos produced by pion decay at rest will enable measurements of cross sections for a wide range of target nuclei~\\cite{Avi.03,Efr.05}, and the promising ``beta-beams'' method for the production of pure electron neutrino-beams by using the $\\beta$-decay of boosted radioactive ions~\\cite{Zuc.02,Vol.04}. Cross sections for neutrino energies in the range of tens of MeV could be measured, and these reactions are particularly interesting for supernova studies~\\cite{McL.04}.  Weak interaction rates at low energies have been analyzed employing a variety of microscopic approaches, principally in the frameworks of the shell model~\\cite{Hax.87,Eng.96,Hay.00}, the random phase approximation (RPA)~\\cite{Aue.97,Sin.98,Vol.00,Vol.02}, continuum RPA (CRPA)~\\cite{Kol.92,Kol.95,Jac.99,Jac.02,Bot.05}, hybrid models of CRPA and shell model~\\cite{Kol.99,Kol.03}, and the Fermi gas model~\\cite{Wal.75,Gai.86,Kur.90}. The shell model provides a very accurate description of ground state wave functions. The description of high-lying excitations, however, necessitates the use of large model spaces and this often leads to computational difficulties, making the approach applicable essentially only to light and medium-mass nuclei. For systematic studies of weak interaction rates throughout the nuclide chart, microscopic calculations must therefore be performed using models based on the RPA.  Hybrid models combine the shell-model and CRPA in such a way that occupation probabilities of single-particle states in a specific nucleus are determined by shell model calculations, and then inserted into the CRPA~\\cite{Kol.99}. In general the CRPA employs different interactions for the calculation of the nuclear ground state (for instance, the Woods-Saxon potential), and in the residual CRPA interaction (G matrix from the Bonn potential, or the Landau-Migdal interaction), and thus additional parameters are required in order to adjust the calculated rates to data. A more consistent approach, based on the quasiparticle RPA with Skyrme effective interactions, has been employed in calculations of weak interaction rates~\\cite{Vol.00}. Although in this framework the residual RPA interaction is derived from the same energy functional which determines the nuclear ground state, some terms are usually omitted, and pairing correlations require the adjustment of additional factors. A fully consistent theoretical framework for the description of charge-exchange excitations in open shell nuclei, based on Skyrme effective interactions, has only recently been developed  \\cite{Fra.05}, but not yet employed in the analysis of neutrino-nucleus reactions. Neutrino-nucleus cross sections have also been calculated using the ab-initio no-core shell model based on realistic NN and three-body interactions~\\cite{Hay.03}, and the shell model with an improved Hamiltonian which properly takes into account the spin-isospin interactions~\\cite{Suz.06}. The importance of improved calculations of neutrino-nucleus cross sections for neutrino-oscillation studies has been demonstrated in the recent reanalysis of the LSND experiment~\\cite{Sam.06}, using the particle-number projected quasiparticle RPA~\\cite{Krm.05}, which has shown an enhancement of the neutrino-oscillation probability when compared to previous studies.  Although the general expressions for the transition matrix elements relevant for the calculation of neutrino-nucleus cross sections have been known since many years~\\cite{Con.72,Don.80}, it is only more recently that systematic calculations have been performed in open-shell nuclei by making use of modern effective interactions in the description of both nuclear ground states and excitations~\\cite{Vol.00}. Reliable prediction of weak interaction rates in nuclei necessitates a fully consistent description of the structure of ground states and multipole excitations. Among the relevant charge-exchange excitations, the isobaric analog state (IAS) and Gamow-Teller resonance (GTR) have been the subject of extensive experimental and theoretical studies. Much more limited are the data and theoretical predictions for properties of excitations of higher multipolarities at finite momentum transfer.  In this work we analyze charged-current neutrino-nucleus reactions by employing a fully consistent microscopic approach based on relativistic energy density functionals, and also including pairing correlations in the description of open-shell target nuclei. An essential advantage over most current approaches is the use of a single universal effective interaction in calculations of both ground-state properties and multipole excitations of nuclei in various mass regions of the chart of nuclides.  Of particular interest for the present study are rates for neutrino-nucleus reactions in the low-energy range below 100 MeV, which play an important role in many astrophysical processes, including stellar nucleosynthesis. A quantitative description of nucleosynthesis of heavy elements during the r-process necessitates accurate predictions of neutrino-nucleus cross sections not only in stable nuclei, but also in nuclei away from the valley of $\\beta$-stability. Since nuclei are used as detectors for solar and supernovae neutrinos, as well as in neutrino oscillation experiments, it is important to describe the neutrino detector response in a consistent and fully microscopic theory. Finally, a quantitative estimate of neutrino-nucleus reaction rates will provide information relevant for feasibility studies and simulations of a low-energy beta beam facility, which could be used to produce neutrino beams of interest for particle physics, nuclear physics and astrophysics~\\cite{Vol.04}.  In Sec.~\\ref{sectionII} we outline the basic formalism used in the evaluation of neutrino-nucleus cross sections. In Sec.~\\ref{sectionIII} the relativistic Hartree-Bogoliubov model (RHB), and the proton-neutron relativistic quasiparticle random phase approximation (PN-RQRPA) are briefly reviewed. Section~\\ref{sectionIV} includes several test cases, and the calculated cross sections are compared with results of previous theoretical studies and the available data. Cross sections for supernova neutrinos are analyzed in Sec.~\\ref{sectionV}, and finally Sec.~\\ref{sectionVI} summarizes the results of the present investigation and ends with an outlook for future studies. ",
        "conclusions": ""
    },
    "0710/0710.4553_arXiv.txt": {
        "abstract": "We present detailed images of diffuse UV intergalactic light (IGL), situated in a 60\\,kpc halo that surrounds the radio galaxy MRC\\,1138-262 at z=2. We discuss the nature of the IGL and rule out faint cluster galaxies, nebular continuum emission, synchrotron, inverse Compton, synchrotron self-Compton emission and scattering of galactic stellar light as possible sources of the IGL. Dust scattered quasar light is an unlikely possibility that cannot be ruled out entirely. We conclude that the source of the IGL is most likely to be a young stellar population distributed in a halo encompassing the radio and satellite galaxies, undergoing  star formation at a rate greater than 57$\\pm8$\\Msunpyr. Within 70\\,kpc of the radio core, approximately 44\\% of the star formation that is traced by UV light occurs in this diffuse mode. The average  UV colour of the IGL is bluer than the average galaxy colour, and there is a trend for the IGL to become bluer with increasing radius from the radio galaxy. Both the galaxies and the IGL show a UV colour--surface brightness relation which can be obtained by variations in either stellar population age or extinction. These observations show a different, but potentially important mode of star formation, that is diffuse in nature. Star formation, as traced by UV light,  occurs in two modes in the high redshift universe: one in the usual Lyman break galaxy clump-like mode on kpc scales, and the other in a diffuse mode over a large region surrounding massive growing galaxies.  Such a mode of star formation can easily be missed by high angular resolution observations that are well suited for detecting high surface brightness compact galaxies. Extrapolating from these results, it is possible that a significant amount of star formation occurs in large extended regions within the halos of the most massive galaxies forming at high redshift. ",
        "introduction": "Distant powerful radio galaxies are important probes of the formation and evolution of massive galaxies, because they are among the most luminous and massive galaxies known in the early Universe, and are likely progenitors of dominant cluster galaxies \\cite[e.g][]{Pentericci1997,Villar-Martin2006,Seymour2007}. They are generally embedded in giant (cD-sized) ionized gas halos \\cite[e.g][]{Reuland2003} and surrounded by galaxy overdensities \\citep{Pentericci2000,Venemans2007}. With radio lifetimes (few $\\times10^7$yr) being much smaller than cosmological timescales, the statistics are consistent with every dominant cluster galaxy having gone through a luminous radio phase during its evolution.   The deepest image of a distant radio galaxy with the {\\it Hubble Space Telescope (HST)} is a recent {\\it Advanced Camera for Surveys (ACS)} observation of MRC\\,1138-262 at z$ = 2.156$ \\citep{Miley2006}. This system has been dubbed the Spiderweb galaxy as its complex morphology resembles a spider's web.  The {\\it ACS} image shown in Fig.~1 shows tens of UV bright galaxies surrounding a central galaxy (marked with a white cross). As the radio core is cospatial with the central galaxy we refer to this central galaxy as the radio galaxy, and all the surrounding galaxies are referred to as the satellite galaxies. All of these galaxies may merge together to form a single massive galaxy at z=0, therefore the whole complex system is referred to as the Spiderweb system This system is one of the most intensively studied distant radio galaxies \\citep{Pentericci1997,Pentericci1998,Pentericci2000} and exhibits the properties expected for a progenitor of a dominant cluster galaxy. The infrared luminosity provides an upper limit to the stellar mass of 10$^{12}$\\Msun\\ \\citep{Seymour2007}, so it is one of the most massive galaxies known at z$>$2. The host galaxy is surrounded by a giant Ly$\\alpha$ halo and the high rotation measure of the radio source means that it is embedded in a dense medium with an ordered magnetic field \\citep{Pentericci1997}. The radio galaxy is associated with a $>$3\\,Mpc-sized structure of galaxies with an estimated mass $>4\\times10^{14}$\\Msun\\ \\citep{Venemans2007}, indicating that it is the predecessor of a local rich cluster. Mechanical feedback from the active galactic nucleus (AGN) is sufficient to expel significant fractions of the interstellar medium of the massive gas-rich galaxy  \\citep{Nesvadba2006}. Therefore the AGN may quench the star formation and allow the radio galaxy to evolve onto the red sequence as suggested by popular models of massive galaxy formation \\citep{Croton2006}. The study of the formation of the most massive galaxies through these and other mechanisms has so far mainly been limited to models and simulations \\cite[e.g][]{Dubinski1998,Gao2004,DeLucia2007}. Our aim is to use these observations of a high-redshift radio galaxy to study brightest cluster galaxy formation.  In this work we present detailed images of widespread UV intergalactic light (IGL) situated in a 60\\,kpc halo that lies between the radio galaxy and the UV bright satellite galaxies. We examine the possible origins of this light and present in-situ star formation as the most probable. In section \\ref{method} we describe the observations and data reduction, in section \\ref{results} we describe the distribution and colour of the IGL, and discuss possible origins in section \\ref{nature}. Section \\ref{discussion} discusses the implications of widespread star formation in a halo around a massive forming brightest cluster galaxy. Throughout this work we use $H_0=71$, $\\Omega_M=0.27$, and $\\Omega_\\Lambda=0.73$ \\citep{Spergel2003}. All magnitudes are AB magnitudes. At a redshift of 2.156, the linear scale is 8.4\\,kpc/\". ",
        "conclusions": "The radio galaxy MRC\\,1138-262 is surrounded by a halo of IGL that extends across 60\\,kpc. After examining nebular continuum emission, synchrotron, inverse Compton, synchrotron self-Compton emission, scattering, and stripping of stars as possible sources of the IGL, we conclude the most likely origin of the IGL is in-situ star formation.  The minimum star formation rate (i.e. uncorrected for dust) of the IGL is 57$\\pm8$\\Msunpyr, which is comparable to the total star formation rate derived from the UV luminosity integrated over all the galaxies within 70 kpc of the radio galaxy (and including the radio galaxy itself). Applying a minimum dust correction of $E(B-V)\\sim0.1$ imply by the red colours of the the IGL increases the star formation rate of the IGL to 142\\Msunpyr, and that of the whole Spiderweb system to more than 325\\Msunpyr. We estimate that the total observed star formation rate can produce approximately $\\sim$7\\% of the \\lya\\ emission in the halo surrounding the galaxy. The radial colour gradient of the IGL indicates that there is a smooth range in extinction or stellar population age from the outer to the inner parts of the halo, where the inner parts are older or dustier than the outer region.  While the presence of large quantities of ionized gas found around several remarkable species of high redshift galaxies has been attributed to a number of proposed mechanisms that include AGN feedback, infall, cooling flows, starburst superwinds and mergers, these observations of MRC\\,1138-262 show that any successful model should be able to accommodate the mode of extended star formation present in this paper. Our data suggest that the formation of the most massive galaxies is connected with that of their gaseous envelopes and star forming halos, part of which may precede the intracluster light or cD envelopes, or perhaps may contribute to a satellite population. A significant amount of star formation might be occurring in the form of extended low surface brightness features, beyond the typical UV detection limits, as well as in largely obscured extended halos as detected at infrared and sub-mm wavelengths."
    },
    "0710/0710.0589.txt": {
        "abstract": "Large-scale asymmetries in the stellar mass distribution in galaxies are believed to trace non-equilibrium situations in the luminous and/or dark matter component. These may arise in the aftermath of events like mergers, accretion, and tidal interactions. These events are key in the evolution of galaxies. In this paper we quantify the large-scale lopsidedness of light distributions in 25155 galaxies at $z < 0.06$ from the Sloan Digital Sky Survey Data Release 4 using the $m = 1$ azimuthal Fourier mode.  We show that the lopsided distribution of light is primarily due to a corresponding lopsidedness in the stellar mass distribution. Observational effects, such as seeing, Poisson noise, and inclination, introduce only small errors in lopsidedness for the majority of this sample. We find that lopsidedness correlates strongly with other basic galaxy structural parameters: galaxies with low concentration, stellar mass, and stellar surface mass density tend to be lopsided, while galaxies with high concentration, mass, and density are not. We find that the strongest and most fundamental relationship between lopsidedness and the other structural parameters is with the surface mass density. We also find, in agreement with previous studies, that lopsidedness tends to increase with radius. Both these results may be understood as a consequence of several factors. The outer regions of galaxies and low-density galaxies are more susceptible to tidal perturbations, and they also have longer dynamical times (so lopsidedness will last longer). They are also more likely to be affected by any underlying asymmetries in the dark matter halo. ",
        "introduction": "%\\citet{rz95}: 18 galaxies. Galaxy disks exhibit a wide variety of shapes visible in the near-IR.  They are due to distortions in surface density distributions, not mass-to-light ratios.  Stellar velocities in nonaxisymmetric galaxies differ by 3-6\\% from those in symmetric ones.  %\\citet{zr97}: 60 galaxies, magnitude-limited.  30\\% of galaxies have $A_1 > 0.2$.  Lopsided mass distributions remain long enough to indicate past interactions, not just ongoing ones.  Often the companion is not obvious, either merged or fled.  %\\citet{rr98}: 54 early-type disk galaxies.  Low SFRs to reduce possibility of asymmetric stellar distributions and increase possibilty of light->mass asymmetry tracing.  Traces mass dependence because VRI bands show similar lopsidedness and it's OK to go shorter than I and K bands.  20\\% of galaxies have lopsidedness above 0.19.   %\\citet{rrk00}: Correlations between lopsidedness and recent SF, and current SF were found, lopsided = star-forming.  Significant fractions of stellar content can be created in a short time from minor mergers and on a similar timescale (1 Gyr).  It has long been recognized that galaxies show large-scale asymmetries in their structure \\citep{bl+80}. Lopsided galaxies have such asymmetries where one side of their disk is more massive and/or more extended than the opposite side.  This \u0093lopsidedness\u0094 can be traced in the spatial structure of the stars \\citep{rz95} and/or the HI gas \\citep{rs94} and/or in the large-scale kinematics of this material \\citep{ss+99}.  There are a variety of mechanisms or events that have been proposed to produce the observed lopsidedness. All of them involve a time-dependent non-equilibrium dynamical state, in most cases triggered through an external process. Such external processes are a natural consequence of the standard Lambda Cold Dark Matter cosmological framework. This implies that galaxies assemble hierarchically (a process that is on-going). Examples that can lead to lopsidedness include a minor merger (\\citealt{wm+96}; \\citealt{zr97}), the tidal interaction resulting from a close encounter between roughly equal-mass galaxies \\citep{kl+02}, and the asymmetric accretion of intergalactic gas into the disk (\\citealt{bc+05}; \\citealt{kk+05}). Other mechanisms involve the dark matter halo: stars and gas orbiting in a lopsided dark matter halo (\\citealt{w94}; \\citealt{j97}; \\citealt{j99}) or a stellar/gas disk that is offset with respect to the center of the dark matter halo (\\citealt{ls98}; \\citealt{ns+01}). These also involve past tidal interactions and/or mergers that have perturbed the dark matter halo, but such perturbations may be quite long-lived. Finally, dynamical processes internal to the disk that lead to mildly lopsided distributions have also been investigated (\\citealt{st+90}; \\citealt{st96}; \\citealt{mt97}).  A variety of programs to study lopsidedness have been undertaken over the past decade. Most of these investigations have studied the lopsided distribution of the stellar component through analysis of optical and near-infrared images. \\citet{zr97} studied a magnitude-limited sample of 60 field spiral galaxies. They measured lopsidedness as the radially averaged, azimuthal $m=1$ Fourier amplitude $A_1$ of the light (see Section \\ref{sec:error} below) and computed lopsidedness between 1.5 and 2.5 scale lengths in the galactic disks. The value of $A_1$ indicates the typical large-scale variation in mass density from side to opposite side at the same distance from the galactic center. The mass density typically varies from between $1 \\pm A_1$ times the average density at the same radius. They found that $\\sim 30$\\% of field spiral galaxies exhibited significant lopsidedness ($A_1 > 0.2$).  \\citet{rr98} followed up this work by studying lopsidedness in 54 early-type galaxies and found that $\\sim 20$\\% had $A_1 > 0.19$. \\citet{cbj00} studied a sample of 113 $z < 0.01$ galaxies (elliptical, spiral, and irregular) but used a 180-degree rotational asymmetry measure $A_{180}$.  They found that asymmetry is strongly dependent on morphological type, with lower asymmetry in elliptical and lenticular galaxies and higher asymmetry in late-type disk and irregular galaxies. More recently, \\citet{bc+05} have measured the Fourier $A_1$ parameter for 149 galaxies in the Ohio State University Bright Galaxy Survey. They confirmed that a large fraction of galaxies have significant lopsidedness in their stellar disks, with late-type galaxies being more lopsided.  Lopsidedness in the {\\it light} distribution can be produced by either a corresponding asymmetry in the underlying {\\it mass} distribution in the stellar population or by large-scale variations in the mass-to-light ratio (e.g., from star formation and dust obscuration). \\citet{rz95} investigated this issue with a sample of 18 face-on spiral galaxies imaged in the K$^\\prime$ (2.2$\\mu$ m) band where the effect of young stars or dust is minimized. They found that about a third of the sample showed significant lopsidedness (similar to results from optical investigations). Similarly, \\citet{rr98} found that lopsidedness in early-type disk galaxies is nearly identical when observed in the $V$, $R$, and $I$ bands. They concluded that an asymmetric mass distribution then accounts for the majority of the asymmetry in the light distribution in these galaxies.  Lopsidedness has also been studied in the distribution of HI gas. Since the HI can frequently be traced to significantly larger radii than the stars, these investigations are highly complementary to the optical image analysis. Due to the time-consuming nature of HI interferometric mapping, only modest size samples have been analyzed in this way \\citep{ss+99}. On the other hand, \\citet{rs94} have examined the global HI line profiles for roughly 1700 galaxies, and shown that at least 50\\% are significantly asymmetric (confirming that the large-scale HI distribution is frequently lopsided). HI maps also show that \u0096 apart from a lopsided distribution of the gas \u0096 the HI rotation curves are often asymmetric \\citep{ss+99}. The connection between the phenomena of structural and kinematic lopsidedness in galaxies is not yet clear \\citep{ss+99}.  Despite these diverse investigations and the abundance of proposed models, the origin of lopsidedness remains unsettled. For models involving tidal interactions or minor mergers, there is an expected link between lopsidedness and the local environment. The evidence in this regard has been mixed (e.g., \\citealt{wp04}; \\citealt{bc+05}; \\citealt{aj+06,aj+07}; \\citet{dc+07}).  %\\citet{dc+07} have undertaken %the most comprehensive investigation so far of this issue. Using the %rotational asymmetry measure $A_{180}$ to study a sample of over 3000 %galaxies, they find that close pairs of galaxies are more asymmetric %than other galaxies and that the asymmetry increases as the pair %separation decreases. They conclude that these global asymmetries %trace recent tidal interactions or mergers.  The investigations summarized above have all involved relatively small samples of galaxies, making it difficult to assess the overall distribution of asymmetry or lopsidedness as a function of the basic parameters that characterize the structure of galaxies. This is the first of three papers in which we use the wealth of data available from the Sloan Digital Sky Survey (SDSS) to extend these studies of small samples (of-order one hundred galaxies) to large samples (tens of thousands). In this paper, we describe our sample selection and methodology. We also relate lopsidedness to the basic structural properties of the galaxies. In Paper II we will investigate the connection between lopsidedness and both star formation and black hole growth in galaxies. Finally, in Paper III we will examine the connection between lopsidedness and the local galaxy environment.  In \\S\\ref{sec:data}, we begin by presenting an initial low-redshift sample from the SDSS and describe the observations and properties for its galaxies.  Next, we explain our lopsidedness calculation. In \\S\\ref{sec:syserr}, we address the major data quality issues that limit the reliability of the measurements for portions of the sample. On this basis, we apply cuts on the observational parameters to weed out the problematic cases for our subsequent analysis. Next, \\S\\ref{sec:lopprops} describes the lopsidedness of galactic light distributions in different optical/near-IR bands, its correspondence with lopsided mass distributions, and its radial dependence.  We then examine the relationship between lopsidedness and the basic structural properties of galaxies in \\S\\ref{sec:lopsfh}.  Finally, we summarize our findings in \\S\\ref{sec:summary}. ",
        "conclusions": "}  We have measured large-scale galactic asymmetry for a large sample of low-redshift ($z <$ 0.06) galaxies drawn from the Sloan Digital Sky Survey.  Our use of lopsidedness, a radially averaged $m=1$ azimuthal Fourier mode, has proven useful for a large fraction of the sample. Images of a minority of galaxies in the sample have poor observational properties that cause significant systematic errors in the lopsidedness calculation, and these galaxies were removed from the sample via cuts on angular size, signal-to-noise, and ellipticity/inclination.  Those cuts removed a higher fraction of the high-mass, high-mass-density, and high-concentration galaxies than those with low values of these structural properties.  Nonetheless, the resulting sample is well represented by the galaxies of the same range of structural properties as the original sample.  We find that there are no systematic differences between the $(g-i)$ colors of the brighter and fainter sides of lopsided galaxies. This implies that there is no systematic difference in the mass/light ratio \\citep{k+07}, and hence that the lopsided light distributions are primarily caused by lopsided distributions in the stellar mass. We have verified this through analysis of the relationship between color and mass/light ratio for both model galaxy spectral energy distributions and SDSS galaxy data. However, for our sample the lopsidedness in the $g$-band tends to be slightly greater than in the $r$- and $i$-bands. Thus, some of the lopsidedness in the light does arise from the effects of star-formation and/or dust extinction (which will more strongly affect the $g-$band light).  Lopsidedness is a structural property that depends strongly on other structural properties.  Galaxies with progressively lower concentration, stellar mass, or stellar mass density tend to have progressively higher lopsidedness. We show that the strongest and most fundamental correlation is between lopsidedness and stellar mass density. We also find that lopsidedness increases systematically with increasing radius, particularly for late-type galaxies.  Lopsidedness can be induced through tidal stress associated with interactions with a companion galaxy or through accretion or minor mergers (e.g. \\citealt{zr97}; \\citealt{bc+05}). Galaxies with low density will be most affected by tidal stress, and the effects of a tidal perturbation will last longer in such systems due to the longer dynamical times. The same arguments pertain to the outer parts of galaxies. Thus, the two above results make good physical sense. Alternatively, if the dark matter halo is lopsided, its effects on the structure of the stellar disk will be more pronounced in the outer region and in galaxies with low mass and low density (where dark matter is more dynamically important). The relatively large values of lopsidedness we measure to be commonplace ($A_1 > 0.1$) appear to be too large to be generated by internally generated dynamical processes (e.g., \\citealt{mt97}).  Our overall goal in this investigation has been to use lopsidedness as a way of quantifying the signature of moderate or weak global dynamical perturbations.  The next step will be to determine the connections between such perturbations and both the on-going/recent star formation and the growth of supermassive black holes in galaxies. These connections can help constrain the processes and conditions that guide the formation and evolution of the galaxies.  In future papers we will address these questions using the present sample of galaxies.  %are also thought to induce bursts of star formation.  Using two well %studied SFH indicators, we have shown that lopsidedness is typically %associated with bursty and fast star formation, and quiescent star  %formation is typical in symmetric galaxies.  We will extend this %discussion in future work.  %Interactions are also thought to induce increased star formation rates %in galaxies.  We have found that lopsided galaxies produce stars at %much faster rates per unit stellar mass than symmetric galaxies. %Lopsidedness is then another way to qualitatively describe the stellar %population age of galaxies.  Taking lopsidedness as a signature of a %recent interaction, the link between star formation and lopsidedness %may be that they are symptoms of the same interactions.  We will %extend this discussion in future work.   JB acknowledges the receipt of a FCT post-doctoral grant BPD/14398/2003.  We would like to thank Vivienne Wild for reading a draft of the manuscript.    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/.  The SDSS is managed by the Astrophysical Research Consortium for the Participating Institutions. The Participating Institutions are the American Museum of Natural History, Astrophysical Institute Potsdam, University of Basel, University of Cambridge, Case Western Reserve University, University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington."
    },
    "0710/0710.4547_arXiv.txt": {
        "abstract": "We present an extensive data set of $\\sim 150$ localized features from \\Cassit{} images of Saturn's Ring~A, a third of which are demonstrated to be persistent by their appearance in multiple images, and half of which are resolved well enough to reveal a characteristic ``propeller'' shape.  We interpret these features as the signatures of small moonlets embedded within the ring, with diameters between 40 and 500~meters.  The lack of significant brightening at high phase angle indicates that they are likely composed primarily of macroscopic particles, rather than dust.  With the exception of two features found exterior to the Encke Gap, these objects are concentrated entirely within three narrow ($\\sim 1000$ km) bands in the mid-A Ring that happen to be free from local disturbances from strong density waves.  However, other nearby regions are similarly free of major disturbances but contain no propellers.  It is unclear whether these bands are due to specific events in which a parent body or bodies broke up into the current moonlets, or whether a larger initial moonlet population has been sculpted into bands by other ring processes. ",
        "introduction": "Saturn's main rings (particularly Ring~A) were determined by the Voyager Radio Science experiment \\citep{Zebker85} to be primarily composed of a distribution of icy particles of diameter $D \\gtrsim 1$~cm.  A steep cutoff in the size-distribution was discerned at $D \\sim 20$~m, but particles larger than that value could not be probed due to limitations imposed by the radio experiment's carrier wavelength.  At the large end of the particle-size distribution are the two known moonlets embedded in gaps in the outer A~Ring, Pan and Daphnis, of diameters $\\sim 28$~km and $\\sim 8$~km, respectively \\citep{PorcoSci07}.  Nothing was known about the distribution of intermediate-size ring particles, between 20~m and 8~km in diameter, until the first evidence of ``missing-link'' particles was found in very-high-resolution images taken by the \\Cassit{} spacecraft during its insertion into Saturn orbit \\citep{Propellers06}.  These intermediate-size moonlets are not directly seen, but rather the propeller-shaped disturbances they create in the ring continuum.  This morphology had previously been predicted by numerical simulations \\citep{SS00,SSD02,Seiss05}.  The sizes of the perturbing moonlets, subject to some ambiguities in interpretation, were given as $D \\sim 100$ meters.  Their surface densities are quite low relative to smaller particles, corroborating the steep cutoff reported by \\citet{Zebker85}.  We here present and analyze a data set of 158 localized features in the A~Ring, many of which are well-enough resolved to reveal the characteristic propeller shape.  Four of these objects are those reported by \\citet{Propellers06}, and another eight were first noted by \\citet{Sremcevic07}.  Recently, \\citet{Espo07} have found evidence for similarly-sized moonlets in the narrow and highly disturbed F~Ring; however, this is not directly applicable to our analysis because there is no particular reason to expect the particle-size distribution of the F~Ring to be simply related to that of the A~Ring.  Section~\\ref{Propellers} gives further background on the nature and interpretation of propellers.  Section~\\ref{Observations} summarizes the imaging sequences we used in compiling our data set, and Section~\\ref{Analysis} describes the process by which features were identified and characterized with one of two models (``resolved'' or ``unresolved'').  Our results are summarized in Section~\\ref{Results}.  Section~\\ref{Interpretation} contains further discussion of the interpretation of propeller features.   A full and unabridged presentation of our data set is given as an Appendix.  \\begin{figure}[!t] \\begin{center} \\includegraphics[width=10cm,keepaspectratio=true]{f1.pdf} \\caption{Particle trajectories under Hill's equations \\citep[see, e.g.,][ch.~3.13]{MD99}, showing the propeller-shaped chaotic zones as well as associated moonlet wakes.  Radial ($r$)) and azimuthal ($\\ell$) coordinates are each shown in units of Hill radii; the perturbing mass is located at the origin, and the $L_1$ and $L_2$ Lagrange points are at $\\ell = 0$, $r = \\pm 1$  The direction towards Saturn is down, and the orbital direction is to the right.  Scattered trajectories deflected by $>5R_H$ are not shown.  This simple model, which neglects inter-particle collisions as well as self-gravity, is presented for conceptual purposes only. \\label{Illustration}} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0710/0710.1648_arXiv.txt": {
        "abstract": "We have carried out a search for radio emission from six X-ray dim isolated neutron stars (XDINSs) observed with the Robert C. Byrd Green Bank Radio Telescope (GBT) at 820~MHz. No bursty or pulsed radio emission was found down to a $4\\sigma$ significance level. The corresponding flux limit is 0.01--0.04~mJy depending on the integration time for the particular source and pulse duty cycle of 2\\%. These are the most sensitive limits yet on radio emission from these objects. ",
        "introduction": "The \\emph{ROSAT} mission discovered a group of seven nearby low-luminosity isolated neutron stars, termed the X-ray dim isolated neutron stars (XDINSs). They share very similar properties and are characterized by soft blackbody-like spectra in the range $\\sim 40$--100 eV, very faint optical counterparts ($V>25$), long spin periods of 3--12~s (see Table~\\ref{table}). For a recent review see \\citet{haberl2007}.  So far, no confident detections of pulsed radio emission were found from XDINSs. The aim of this project is to search for pulsed and bursty radio emission from XDINSs to link them in their evolutionary scenarios with other classes of neutron stars, such as magnetars and rotating radio transients (RRATs). All three populations of neutron stars have similar properties, such as period, period derivative, age, and magnetic field so that connections between them are plausible. ",
        "conclusions": "The sporadicity of the RRATs' radio emission led to immediate suggestions that they are related to other classes of traditionally ``radio-quiet'' neutron stars such as XDINSs and magnetars. \\citet{popov2006} have shown that the implied birthrate of RRATs is more  consistent with that of XDINSs than that of magnetars. As shown in the P-\\.P diagram on the Figure~\\ref{ppdot}, RRATs and XDINSs also have similar periods and period derivatives, implied ages and magnetic fields. However, the RRATs spin-down properties are also consistent with those of the normal pulsar population and X-ray observations of one RRAT~\\citep{mmclaugh2007} reveal properties similar to those of both normal radio pulsars and XDINSs.  \\begin{figure} \\includegraphics[scale=0.44]{p-pdot_bw.ps} \\caption{P-\\.P diagram. Lines on the top mark the period values for the objects with unknown yet period derivatives: seven RRATs (solid), one magnetar SGR~1627$-$41 and two candidates AX~J1845$-$03 and CXO~J164710.2$-$455216 (dashed), and one XDINS RX~J1605.3+3249 (longer dashed). } \\label{ppdot} \\end{figure}  The RRATs are powerful sources of isolated radio bursts, and we have not detected such bursts, or any periodic emission, from the six XDINSs we have observed. Because the distances to the XDINSs are believed to be much smaller than those to the RRATs, we should have had high sensitivity to RRAT-like radio emission. However, our non-detection of such emission does not necessarily mean that there is no relationship between these two source classes.  XDINSs may simply be ``radio-quiet'', but it is also likely that perhaps the narrow radio beams from these XDINSs are simply misaligned with our line-of-sight. It is possible that searches at lower frequencies, where radio emission beams are believed to be wider, may be more sensitive to radio emission from XDINSs. Indeed, Malofeev and co-authors~\\citep{malofeev2005, malofeev2007} reported detection of radio emission from RX~J1308.6+2127 and  RX~J2143.7+0654 at the low frequency of 111~MHz. On the other hand, XDINSs could have very steep spectral indices. If the detection of Malofeev and co-authors is real, our non-detection of radio emission from these two XDINSs at 820~MHz sets a lower limit on the spectral index of 3.6. Finally, it is possible that our non-detection of radio emission from these XDINSs is due to the large amount of contamination from RFI. Our search highlights the importance of improved excision algorithms for impulsive, broadband terrestrial interference.    \\begin{theacknowledgments} SZ thanks STFC (ex-PPARC) for support through an AF. The Robert C. Byrd Green Bank Telescope (GBT) is operated by the National Radio Astronomy Observatory which is a facility of the U.S. National Science Foundation operated under cooperative agreement by Associated Universities, Inc. \\end{theacknowledgments}"
    },
    "0710/0710.1683.txt": {
        "abstract": "The lensing cross section of triaxial halos depends on the relative orientation between a halo's principal axes and its line of sight.  Consequently, a lensing subsample of randomly oriented halos is not, in general, randomly oriented. Using an isothermal mass model for the lensing galaxies and their host halos, we show that the lensing subsample of halos that produces doubles is preferentially aligned along the lines of sight, whereas halos that produce quads tend to be projected along their middle axes. These preferred orientations result in different projected ellipticity distributions for quad, doubles, and random galaxies.   We show that $\\approx 300$ lens systems must be discovered to detect this effect at the $95\\%$ confidence level.  We also investigate the importance of halo shape for predicting the quad-to-double ratio and find that the latter depends quite sensitively on the distribution of the short-to-long axis ratio, but is otherwise nearly independent of halo shape. Finally, we estimate the impact of the preferred orientation of lensing galaxies on their projected substructure mass fraction, and find that the observed alignment between the substructure distribution and the mass distribution of halos result in a negligible bias. ",
        "introduction": "Statistics of lensing galaxies have been used as cosmological and galaxy formation probes since early in the modern history of gravitational lensing \\citep[][]{turneretal84}.  Lensing rates can be used to constrain dark energy \\citep[][]{fukugitaetal92, chae03,mitchelletal05, chae07,ogurietal07}, to probe the structure of lensing galaxies \\citep[][]{keeton01d,kochanekwhite01, chae05}, and to probe galaxy evolution \\citep[][]{chaemao03, ofeketal03,rusinkochanek05}.   While the use of lensing statistics as a cosmological probe has had mixed success, particularly early on, it remains a unique probe with entirely different systematics from more traditional approaches. Consequently, lensing statistics are likely to remain a fundamental cross-check of our understanding of cosmology and galaxy evolution.  One of the difficulties that confronts the study of lensing statistics is that, in general, the halo population that produces gravitational lenses can in fact be a highly biased subsample of the general halo population.  For instance, it has long been known that while early type galaxies compose only $\\approx 30\\%$ of all luminous galaxies, the majority of lensing galaxies are in fact early type since these tend to be more massive and reside in more massive halos than their late counterparts.  By the same token, lensing early type galaxies tend to have higher luminosity and velocity dispersions than non-lensing early type galaxies \\citep[][]{moelleretal06, boltonetal06}.  Overall, then, when interpreting lensing statistics, one ought to always remember that by selecting lensing galaxies one is automatically introducing an important selection effect that can significantly bias the distribution of any galaxy observable that has an impact on the lensing probabilities. Here, we consider one such source of bias, the triaxiality of galaxy halos.\\footnote{Throughout this work, we will be using the term galaxy and halo more or less interchangeably.  The reason for this is that we are primarily focused on the impact of halo triaxiality on the lensing cross section, and the latter depends only on the {\\it total} matter density.  Consequently, differentiating between halo and galaxy would only obfuscate presentation and introduce unnecessary difficulties.  For instance, while modeling the total matter distribution as isothermal is a reasonable approximation, neither the baryons nor the dark matter by itself is isothermally distributed.  Thus, it is much simpler to adopt an isothermal model, and refer to the baryons plus dark matter as a single entity, than to try to differentiate between the two.  Likewise, when discussing triaxiality, what is important in this work is the triaxiality of the total matter distribution.}   That halo triaxiality can have important consequences for lensing statistics has been known for several years. For instance, \\citet[][]{ogurikeeton04} have shown that triaxiality can significantly enhance the optical depth of large image separation lenses.  Similar conclusions have been reached concerning the formation of giant arcs by lensing clusters \\citep[see e.g.][and references therein]{ogurietal03, rozoetal06c, hennawietal07}. Curiously, however, little effort has gone into investigating how observational properties of lensing galaxies can be different from those of the galaxy population as a whole due to the triaxial structure of galactic halos. This work addresses this omission.  The first observable we consider is the projected axis ratio of lensing galaxies. Roughly speaking, given that non-zero ellipticities are needed in order to produce quad systems, one would generically expect lenses that lead to this image configuration to be more elliptical than the overall galaxy population. Likewise, lensing galaxies that produce doubles should, on average, be slightly more circular than a random galaxy. There can, however, be complications for these simple predictions due to halo triaxiality.   For instance, given a prolate halo, projections along the long axis of the lens will result in highly concentrated, very circular profiles.  Will the increase in Einstein radius of such projections compensate for the lower ellipticity of the system, implying most quads will be projected along their long axis, or will it be the other way around? Clearly, the relation between ellipticity and lensing cross sections is not straightforward once triaxiality of the lensing galaxies is taken into account, but it seems clear that there should be some observable difference between the ellipticity distribution of lensing galaxies and that of all early types. Interestingly, no such difference has been observed \\citep[][]{keetonetal97,rusintegmark01}, which seems to fly in the face of our expectations \\citep[though see also the discussion in][]{keetonetal98}. Is this actually a problem, or will a quantitative analysis show that the consistency of the two distributions is to be expected?  Here, we explicitly resolve this question, and demonstrate that current lens samples are much too small to detect the expected differences.  Having considered the ellipticity distribution of random and lensing galaxies, it is then a natural step to investigate the impact of halo triaxiality on predictions of the quad-to-double ratio.  Specifically, it is well known that the quad-to-double ratio is sensitive to the ellipticity distribution of lensing galaxies \\citep[][]{keetonetal97}, so if lensing can bias the distribution of ellipticities in lensing galaxies, then it should also affect the predicted quad-to-double ratios.  This is an important point because it has been argued that current predictions for the quad-to-double ratio are at odds with observations. More specifically, the predicted quad-to-double ratio for the CLASS \\citep[Cosmic Lens All-Sky Survey,][]{myersetal03,browneetal03} sample of gravitational lenses is too low relative to observations \\citep[][]{rusintegmark01,hutereretal05}. Curiously, however, recent work on the quad-to-double ratio observed in the SQLS \\citep[Sloan Digital Sky Survey Quasar Lens Search,][]{ogurietal06,inadaetal07}. suggests that the exact opposite is true for the latter sample, namely, theoretical expectations are too high relative to observations \\citep[][]{oguri07}.   In either case, it is of interest to determine how exactly does triaxiality affects theoretical predictions, especially since the aforementioned difficulties with the CLASS sample has led various authors to offer possibilities as to how one might boost the expected quad-to-double ratios.  Specifically, one can boost the quad-to-double ration in the class sample either from the effect of massive satellite galaxies near the lensing galaxies \\citep[][]{cohnkochanek04}, or through the large-scale environment of the lensing galaxy \\citep[][]{keetonzabludoff04}.  Clearly, we should determine whether halo triaxiality can be added to this list.  This brings us then to the final problem we consider here, namely whether the substructure population of lensing galaxies is different from that of non-lensing galaxies. Specifically, we have argued that lensing galaxies will not be isotropically distributed in space.  Since the substructure distribution of a dark matter halo is typically aligned with its parent halo's long axis \\citep[][]{zentneretal05,libeskindetal05,agustssonbrainerd06,azzaroetal06}, it follows that the projected distribution of substructures for lensing galaxies may in fact be different for lensing halos than for non-lensing halos.  Such an effect could be quite important given the claimed tension between the Cold Dark Matter (CDM) predictions for the substructure mass fraction of halos \\citep[see][]{maoetal04} and their observed values \\citep[][]{dalalkochanek02a,kochanekdalal04}. Likewise, such a bias would impact the predictions for the level of astrometric and flux perturbations produced by dark matter substructures in gravitational lenses \\citep[][]{rozoetal06,chenetal07}.  Here, we wish to estimate the level at which the projected substructure mass fraction of lensing halos could be affected due to lensing biasing.  The paper is organized as follows: in section \\ref{sec:biases} we derive the basic equations needed to compute how observable quantities will be biased in lensing galaxy samples due to halo triaxiality.  Section \\ref{sec:model} presents the model used in this work to quantitatively estimate the level of these biases, and discusses how lensing halos are oriented relative to the line of sight as a function of the halos' axes ratios.  Section \\ref{sec:axis} investigates the projected axis ratio distributions of lensing versus non-lensing galaxies, and demonstrates that present day lensing samples are too small to detect the triaxiality induced biases we have predicted.  Section \\ref{sec:ratio} discusses the problem of the quad to double ratio, and section \\ref{sec:subs} demonstrates that halo triaxiality biases the projected substructure mass fraction in lensing halos by a negligible amount.  Section \\ref{sec:caveats} discusses a few of the effects we have ignored in our work and how these may alter our results, and finally section \\ref{sec:summary} summarizes our work and presents our conclusions.  %---------------------------------------------- ",
        "conclusions": "\\label{sec:summary}  The triaxial distribution of mass in galactic halos implies that the probability that a galaxy becomes a lens is dependent on the relative orientation of the galaxy's major axis to the line of sight.   Consequently, a subsample of randomly oriented galaxies that act as strong lenses will {\\it not} be randomly oriented in space.  The relative orientation and the strength of the alignment depends on the shape of the matter distribution, and on the type of lens under consideration: prolate doubles have a high probability of being project along their long axis, whereas the distribution of oblate doubles is nearly isotropic. Prolate quads are most often projected along their middle axis, though the degree to which alignment occurs is not as strong as for prolate doubles.  Interestingly, highly prolate quads are also more likely to be projected along their short axis than along the long axis, though this very quickly changes as halos become more triaxial and less prolate.  Oblate quads strongly avoid projections along the short axis of the lens, but projections along the other two axis are almost equally likely.  An important consequence of the differences in the distribution of halo orientations for quad lenses, double lenses, and the galaxy population as a whole is that the ellipticity distribution of these various samples must be different, even if the distribution of halo shapes is the same.   Specifically, we predict that quad lenses are typically more elliptical than random galaxies, and that the ellipticity distribution of doubles is very slightly more circular than that of random galaxies.  While current data do not show any indication of these trends, we have shown that $\\approx 300\\ (1,400)$ lenses are necessary to obtain a $2\\sigma\\ (5\\sigma)$ detection of the effect.  The fact that halo triaxiality affects the ellipticity distribution of lensing galaxies also means that halo triaxiality needs to be properly taken into account in lensing statistics.   Consequently, we estimate how the biased lensing cross sections of galaxies depend on halo shape, and find that they are nearly independent of the halo shape parameter $T$.  Instead, the mean biased cross section of a lens depends almost exclusive on the distribution on the short-to-long axis ratio $q_2$ (often denoted by $s$).  Finally, given that the distribution of substructures in numerical simulations is observed to be preferentially aligned with the long axis of the host halos, we estimate how the preferred orientation of lensing galaxies affects their predicted substructure mass fraction.  We find that biases due to non-isotropic distribution of halos relative to the line of sight have an insignificant impact on the mean substructure mass fraction of lensing galaxies.  {\\bf Acknowledgements: } ER would like to thank Christopher Kochanek for numerous discussions and valuable comments on the manuscript which have greatly improved both the form and content of this work.  The authors would also like to thank to Emilio Falco for kindly providing the isophotal axis ratio data that was needed for producing Figure \\ref{fig:cumq}, and to Charles Keeton for a careful reading of the manuscript. ER was funded by the Center for Cosmology and Astro-Particle Physics (CCAPP) at The Ohio State University.  ARZ has been funded by the University of Pittsburgh, the National Science Foundation (NSF) Astronomy and Astrophysics Postdoctoral Fellowship program through grant AST 0602122, and by the Kavli Institute for Cosmological Physics at The University of Chicago.  This work made use of the National Aeronautics and Space Administration Astrophysics Data System.  % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- %--------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % --------------------------------------------------------------------------------"
    },
    "0710/0710.5407_arXiv.txt": {
        "abstract": "{Based on the analogy with non-minimal $SU(2)$ symmetric Wu-Yang mono\\-pole with regular metric, the solution describing a non-minimal $U(1)$ symmetric Dirac monopole is obtained. In order to take into account the curvature coupling of gravitational and electromagnetic fields, we reconstruct the effective metrics of two types: the so-called associated and optical metrics. The optical metrics display explicitly that the effect of birefringence induced by curvature takes place in the vicinity of the non-minimal Dirac monopole; these optical metrics are studied analytically and numerically. }  \\email 1 {Alexander.Balakin@ksu.ru}% \\email 2 {Alexei.Zayats@ksu.ru}% ",
        "introduction": "The Dirac monopole as a specific static spherically symmetric solution to the minimal Einstein-Maxwell equations has became a subject of discussion in tens papers, reviews and books (see, e.g., \\cite{Dirac,000,00,0,1,3,2}). The motion of massive and massless particles, which possess electric charge or are uncharged, is studied in detail (see, e.g., \\cite{4,5} and references therein). In the paper \\cite{BaZa07} we introduced and discussed the $SU(2)$ symmetric Wu-Yang monopole of the new type, namely, the non-minimal monopole with regular metric. Since the non-minimal Wu-Yang monopole is effectively Abe\\-lian, it is naturally to consider the corresponding analog of that solution in the framework of non-minimal electro\\-dynamics. Mention that non-minimal models with mag\\-netic charge have been discussed earlier (see, e.g., \\cite{Horn,MHS}), but the exact analytical regular solution obtained here as direct reduction to the $U(1)$ symmetry is the new one.  The second novelty of the presented paper is the investigation of photon dynamics in the vicinity of the Dirac monopole accounting the curvature coupling of the gravitational and electromagnetic fields. In the presence of non-minimal interaction (induced by curva\\-ture) the master equations for electromagnetic and gravitational fields in vacuum can be rewritten as the master equations in some effective anisotropic (quasi)me\\-dium \\cite{B1,BL05}. This means that two effective (optical) metrics can be introduced \\cite{HehlObukhov,Perlick,BZ05}, so that the photon propagation in vacuum interacting with curvature is equivalent to the photon motion in the effective space-time with the first or second optical metric, depending on the photon polarization. Even if the real space-time has a regular metric, the optical metrics can be singular, admitting the interpretation in terms of the so-called ``trapped surfaces'' and ``inaccessible zones'' \\cite{AG1,AG2}. We discuss this problem in Sect. 3. Numerical modeling of the photon orbits, presented in Sect. 4, supplement our conclusions. ",
        "conclusions": "Qualitative and numerical analysis of the photon orbits in the vicinity of non-minimal Dirac monopole with regular metric has demonstrated the following interesting features.  1. Propagation of the electromagnetic waves in the vicinity of non-minimal Dirac monopole is characterized by birefringence, induced by curvature, i.e., the phase velocities of waves depend on their polarization. Two different optical metrics should be introduced to describe two principal states of polarization.  2. The metric of the non-minimal Dirac monopole, obtained and discussed in this paper, is regular, thus, all the singularities of the optical metrics have a dynamic origin and are supported by the non-minimal (curvature induced) interaction of the gravitational and electromagnetic fields.  3. The points of self-intersection, the points of the closest approach, the reverse points, etc. in the photon trajectories can be recognized and catalogued for different combinations of the values of the impact parameter and the parameter of non-minimal coupling. We discussed here only the principal pictures.   \\Acknow This work was partially supported by the DFG through project No. 436RUS113/487/0-5. The authors are grateful to Claus L\\\"ammerzahl for valuable remarks.  \\small"
    },
    "0710/0710.5294_arXiv.txt": {
        "abstract": "{Thanks to remarkable progress, radial velocity surveys are now able to detect terrestrial planets at habitable distance from low-mass stars. Recently, two planets with minimum masses below 10~$M_{\\oplus}$have been reported in a triple system around the M-type star Gliese 581. These planets are found at orbital distances comparable to the location of the boundaries of the habitable zone of their star. } {In this study, we assess the habitability of planets Gl~581c and Gl 581d (assuming that their actual masses are close to their minimum masses) by estimating the locations of the habitable-zone boundaries of the star and discussing the uncertainties affecting their determination. An additional purpose of this paper is to provide simplified formulae for estimating the edges of the habitable zone. These may be used to evaluate the astrobiological potential of terrestrial exoplanets that will hopefully be discovered in the near future.} {Using results from radiative-convective atmospheric models and constraints from the evolution of Venus and Mars, we derive theoretical and empirical habitable distances for stars of F, G, K, and M spectral types. } {Planets Gl~581c and Gl~581d are near to, but outside, what can be considered as the conservative habitable zone. Planet `c' receives 30\\% more energy from its star than Venus from the Sun, with an increased radiative forcing caused by the spectral energy distribution of Gl~581. This planet is thus unlikely to host liquid water, although its habitability cannot be positively ruled out by theoretical models due to uncertainties affecting cloud properties and cloud cover. Highly reflective clouds covering at least 75\\% of the day side of the planet could indeed prevent the water reservoir from being entirely vaporized. Irradiation conditions of planet `d' are comparable to those of early Mars, which is known to have hosted surface liquid water.  Thanks to the greenhouse effect of CO$_2$-ice clouds, also invoked to explain the early Martian climate, planet `d' might be a better candidate for the first exoplanet known to be potentially habitable. A mixture of several greenhouse gases could also maintain habitable conditions on this planet, although the geochemical processes that could stabilize such a \\textit{super-greenhouse} atmosphere are still unknown. } {} ",
        "introduction": "The M-type star Gl~581 hosts at least 3 planets, which were detected using radial velocity measurements by Bonfils et al. \\citeyearpar{2005A&A...443L..15B} (planet 'b') and Udry et al. \\citeyearpar{2007A&A...469L..43U} (planets `c' and `d'). The properties of this star and its planets are given in Table~\\ref{tab:gl581}. Before this discovery, only two exoplanets were known to have a minimum mass below 10~$M_{\\oplus}$, which is usually considered as a boundary between terrestrial and giant planets, the latter having a significant fraction of their mass in an H$_2$-He envelope. The first one was GJ~876d, a very hot planet ($P\\leq 2$~days) with a minimum mass of 5.9~M$_{\\oplus}$ \\citep{2005ApJ...634..625R}. The other one is OGLE-05-390L b, found to be a $\\sim$5.5~M$_{\\oplus}$ cold planet at 2.1~AU from its low-mass parent star thanks to a microlensing event \\citep{2006Natur.439..437B,2006ApJ...651..535E}. Neither of these two planets is considered as habitable, even with very loose habitability criteria. In the case of Gl~581, and as already mentioned by Udry et al. (2007), the locations of planet `c' and `d' must be fairly close to the inner and outer edges, respectively, of the habitable zone (HZ). In this paper, we investigate the atmospheric properties that would be required to make the habitability of these planets possible.  Because of its equilibrium temperature of $\\sim$300~K when calculated with an albedo of 0.5, it has been claimed that the second planet of this system, Gl~581c, is potentially habitable (Udry et al. 2007), with climatic conditions possibly similar to those prevailing on Earth. After a brief discussion about the relationship between the equilibrium temperature and habitability, we summarize in this paper what are usually considered as the boundaries of the circumstellar HZ and the uncertainties on their precise location.  In Sect.~\\ref{sec:HZstars} we provide parameterizations to determine such limits as a function of the stellar luminosity and effective temperature.  These can be used to evaluate the potential habitability of the terrestrial exoplanets that should soon be discovered. We then discuss the specific case of the system around Gl~581.   \\begin{table}[!t] \\caption{Properties of the star Gl~581 and its 3 detected planets, from Udry et al. (2007).} \\begin{center} \\begin{tabular}{lllll} \\hline Star & $T_{\\rm eff}$ (K) &$M$/M$_{\\odot}$ & $R$/R$_{\\odot}$ & $L$/L$_{\\odot}$ \\\\ \\hline Gl~581 & 3200 &0.31 & 0.38 & 0.0135 \\\\ \\hline &&& \\\\ &&& \\\\ \\hline Planets  & $a$~(AU) & $M_{\\rm min}$/M$_{\\oplus}$ & $R_{\\rm min}$/R$_{\\oplus}$ & stellar flux\\\\ &   &$^{*}$ & $^{**}$  & $S/S_{0}$$^{***}$\\\\ \\hline b & 0.041 & 15.6 & 2.2-2.6 & 8.1 \\\\ \\rowcolor{lightgray} c & 0.073 & 5.06 & 1.6-2.0 & 2.55 \\\\ \\rowcolor{lightgray} d & 0.253 & 8.3 & 1.8-2.2 & 0.21 \\\\ \\hline \\end{tabular} \\end{center} The potential habitability of planets `c' and `d', highlighted in grey, is discussed in this paper. \\\\ $^{*}$ $M_{\\rm min}=M \\sin i$, where $i$ is the orbital inclination.\\\\ $^{**}$ Radius for a rocky and ocean planet, respectively \\citep{2007Sotin,2007ApJ...665.1413V}.\\\\ $^{***}$ $S_0$ is the solar flux at 1 AU: 1360 W m$^{-2}$.\\\\ \\label{tab:gl581} \\end{table} ",
        "conclusions": "According to our present knowledge, based on available models of planetary atmospheres, and assuming that the actual masses of the planets are the minimum masses inferred from radial velocity measurements, Gl~581c is very unlikely to be habitable, while Gl~581d could potentially host surface liquid water, just as early Mars did.  Because of the uncertainties in the precise location of the HZ boundaries, planets at the edge of what is thought to be the HZ are crucial targets for future observatories able to characterize their atmosphere. At the moment, our theory of habitability is only confirmed by the divergent fates of Venus and the Earth. We will have to confront our models with actual observations to better understand what makes a planet habitable. The current diversity of exoplanets (planets around pulsars, hot Jupiters, hot Neptunes, super-Earths, etc) has already taught us that Nature has a lot more imagination when building a variety of worlds than we expected from our former models inspired by the Solar System.  It is obvious that the idealized model of a habitable planet atmosphere, where the two important constituents are CO$_2$ and H$_2$O, CO$_2$ being controlled by the carbonate-silicate cycle, is likely to represent only a fraction of the diversity of terrestrial planets that exist at habitable distances from their parent star. As an example, planets fully covered by an ocean may be common, either because they are richer in water than Earth or because the distribution between surface and mantle water is different, or perhaps simply because, for a given composition, the mass-to-surface ratio and thus the water-to-surface ratio increases with the planetary mass, as noted by Lissauer \\citeyearpar{1999Lissauer}. Without emerged continents, it is not at all clear that the carbonate-silicate cycle could operate. The planets around Gl~581 can fall into this category since they are significantly more massive than the Earth (especially the $>$8~$M_{\\oplus}$ planet Gl~581d) and also because they may have started their formation in the outer and more water-rich region of the protoplanetary disk.  Darwin/TPF-I and TPF-C could eventually reveal what the actual properties of the atmosphere of Gl~581c and Gl~581d are. From their thermal light curves we could infer if a thick atmosphere is making the climate more or less uniform on both the day and night hemispheres of these planets, despite a (nearly?) synchronized rotation \\citep{2004ASPC..321..170S}. Visible and mid-IR water vapor bands could be searched in the atmosphere of Gl~581d to confirm its habitability. Mid-IR spectra of this planet could also reveal other greenhouse gases at work. Spectral observations of Gl~581c could potentially distinguish between a Venus-like atmosphere dominated by CO$_2$ or an H$_2$O-rich atmosphere. The detection of O$_2$ on this planet would generate a fascinating debate about its possible origin: as either a leftover of H$_2$O photolysis and H escape or a biological release. There is certainly no doubt that Gl~581c and Gl~581d are prime targets for exoplanet characterization missions."
    },
    "0710/0710.0364_arXiv.txt": {
        "abstract": "We explore the cosmological consequences of Modified Gravity (MOG), and find that it provides, using a minimal number of parameters, good fits to data, including CMB temperature anisotropy, galaxy power spectrum, and supernova luminosity-distance observations without exotic dark matter. MOG predicts a bouncing cosmology with a vacuum energy term that yields accelerating expansion and an age of $\\sim$13 billion years. ",
        "introduction": "The preferred model of cosmology today, the $\\Lambda$CDM model, provides an excellent fit to cosmological observations, but at a substantial cost: according to this model, {\\em about 95\\% of the universe is either invisible or undetectable, or possibly both} \\cite{Komatsu2008}. This fact provides a strong incentive to seek alternative explanations that can account for cosmological observations without resorting to dark matter or Einstein's cosmological constant.  For gravitational theories designed to challenge the $\\Lambda$CDM model, the bar is set increasingly higher by recent discoveries. Not only do such theories have to explain successfully the velocity dispersions, rotational curves, and gravitational lensing of galaxies and galaxy clusters, the theories must also be in accord with cosmological observations, notably the acoustic power spectrum of the cosmic microwave background (CMB), the matter power spectrum of galaxies, and the recent observation of the luminosity-distance relationship of high-$z$ supernovae, which is seen as evidence for ``dark energy''.  Modified Gravity (MOG) \\citep{Moffat2006a} has been used successfully to account for galaxy cluster masses \\citep{Brownstein2006b}, the rotation curves of galaxies \\citep{Brownstein2006a}, velocity dispersions of satellite galaxies \\citep{Moffat2007}, and globular clusters \\citep{Moffat2007a}. It was also used to offer an explanation for the Bullet Cluster \\citep{Brownstein2007} without resorting to cold dark matter.  Remarkably, MOG also meets the challenge posed by cosmological observations. In this paper, it is demonstrated that MOG produces an acoustic power spectrum, a matter power spectrum, and a luminosity-distance relationship that are in good agreement with observations, and require no dark matter nor Einstein's cosmological constant.  In the arguments presented here, we rely on simplified analytical calculations. We are not advocating these as substitutes for an accurate numerical analysis. However, a thorough numerical analysis requires significant time and resources; before these are committed, it is useful to be able to demonstrate if a theory is viable, and if the additional effort is warranted.  In the next section, we review the key features of MOG. This is followed by sections presenting detailed calculations for the luminosity-distance relationship of high-$z$ supernovae, the acoustic power spectrum of the CMB, and the galaxy power spectrum. A concluding section summarizes our results and maps out future steps. ",
        "conclusions": "In this paper, we demonstrated how Modified Gravity can account for key cosmological observations using a minimum number of free parameters. Although MOG permits the running of its coupling and scaling constants with time and space, we made very little use of this fact. Throughout these calculations, we used consistently the value of $\\alpha\\simeq 19$ for the MOG coupling constant, consistent with a flat universe with $\\Omega_b\\simeq 0.05$ visible matter content, no dark matter, nor Einstein's cosmological constant. In nearly all cosmological calculations, we set the MOG scaling constant $\\mu$ to the inverse of the radius of the visible universe, which is a natural choice. The only exception is the mass power spectrum calculation, where the scaling constant enters in conjunction with the wave number $k$, and describes gravitational interactions between nearby concentrations of matter, not on the cosmological scale.  The theory requires {\\em no other parameters} to obtain the remarkable fits to data that have been demonstrated here.  At all times, $\\lim\\limits_{r\\rightarrow 0}G=G_N$, i.e., the effective gravitational constant at short distances remains Newton's constant of gravitation. For this reason, the predictions of MOG are {\\em not contradicting our knowledge of the processes of the initial nucleosynthesis}, taking place at redshifts of $z\\simeq 10^{10}$, since the interactions take place over distance scales that are much shorter than the horizon scale.  Our calculations relied on analytical approximations. This is dictated by necessity, not preference. We recognize that numerical methods, including high-accuracy solutions of coupled systems of differential equations, as in {\\tt CMBFAST} \\citep{Seljak1996}, or $N$-body simulations, can provide superior results, and may indeed help either to confirm or to falsify the results presented here. Nevertheless, our present work demonstrates that at the very least, MOG provides a worthy alternative to $\\Lambda$CDM cosmology."
    },
    "0710/0710.2157_arXiv.txt": {
        "abstract": "We investigate the incidence of major mergers creating ${\\rm M}_{\\rm star}>10^{11} {\\rm M}_{\\sun}$ galaxies in the dense environments of present-day groups and clusters more massive than ${\\rm M}_{\\rm halo}=2.5\\times10^{13}{\\rm M}_{\\sun}$. We identify 38 pairs of massive galaxies with mutual tidal interaction signatures selected from $>5000$ galaxies with ${\\rm M}_{\\rm star}\\geq5\\times10^{10} {\\rm M}_{\\sun}$ that reside in a halo mass-limited sample of 845 groups. We fit the images of each galaxy pair as the line-of-sight projection of symmetric models and identify mergers by the presence of residual asymmetric structure associated with both progenitors, such as nonconcentric isophotes, broad and diffuse tidal tails, and dynamical friction wakes. At the resolution and sensitivity of the SDSS, such mergers are found in 16\\% of the high-mass, galaxy-galaxy pairs with $\\leq1.5$ $r$-band magnitude differences and $\\leq30$ kpc projected separations. Relying on automated searches of major pairs from the SDSS spectroscopic galaxy sample will result in missing 70\\% of these mergers owing to spectroscopic incompleteness in high-density regions. We find that 90\\% of these mergers are between two nearly equal-mass progenitors with red-sequence colors and centrally-concentrated morphologies, in agreement with numerical simulations that predict that an important mechanism for the formation of massive elliptical galaxies is the dissipationless (gas-poor or so-called dry) major merging of spheroid-dominated galaxies. We identify seven additional ${\\rm M}_{\\rm star}>10^{11} {\\rm M}_{\\sun}$ mergers with disturbed morphologies and semi-resolved double nuclei. Mergers at the centers of massive groups are more common than between two satellites, but both types are morphologically indistinguishable and we tentatively conclude that the latter are likely located at the dynamical centers of large subhalos that have recently been accreted by their host halo, rather than the centers of distinct halos seen in projection. We find that the frequency of central and satellite merging diminishes with group mass in a manner that is consistent with dynamical friction. Based on reasonable assumptions, the centers of these massive halos are gaining stellar mass at a rate of 1--9\\% per Gyr on average. Compared to the merger rate for the overall population of luminous red galaxies, we find that the rate is 2--9 times greater when restricted to these dense environments. Our results imply that the massive end of the galaxy population continues to evolve hierarchically at a measurable level, and that the centers of massive groups are the preferred environment for the merger-driven assembly of massive ellipticals. ",
        "introduction": "Understanding the formation of the most-massive galaxies (${\\rm M}_{\\rm star}>10^{11}{\\rm M}_{\\sun}$) remains an important challenge in astrophysics. The tip of the stellar mass function is dominated by elliptical galaxies with intrinsically spheroidal mass distributions that are supported by anisotropic stellar motions \\citep{kormendy96,burstein97}. Numerical simulations have long demonstrated that ``major'' mergers between smaller galaxies of comparable mass could produce the observed shapes and dynamics of ellipticals \\citep{toomre77,barnes96,naab03,cox06}. Moreover, massive ellipticals are found in greater abundance in high-density structures like large groups and clusters of galaxies \\citep[e.g.,][]{dressler80a,postman84,hashimoto99,smith05}, which naturally grow through the hierarchical merging of dark-matter halos over cosmic time as expected in the $\\Lambda$CDM cosmological model \\citep{blumenthal84,davis85,cole00}. There is, therefore, a clear expectation for galaxy-galaxy and halo-halo merging to be physically linked \\citep{maller06,hopkins06b,delucia07}. Indeed, modern galaxy formation models predict that massive ellipticals form by major dissipationless (so-called ``dry'') merging of likewise spheroidal and gas-poor progenitors \\citep{boylan06,naab06a}, that a large fraction of today's massive ellipticals had their last major merger since redshift $z=0.5$ \\citep[e.g.,][]{delucia06}, and that the most-massive systems form at the centers of large dark-matter halos \\citep{dubinski98,aragon98}. Yet, direct evidence for the major-merger assembly of massive galaxies at present times has been lacking, and finding such systems is needed to place constraints on their rates, progenitor properties, and environmental dependencies. To this end we look for close pairs of massive interacting galaxies within a complete and well-defined sample of over 5000 galaxies with $z\\leq0.12$ and ${\\rm M}_{\\rm star}\\geq5\\times10^{10}{\\rm M}_{\\sun}$, selected from galaxy groups in the Sloan Digital Sky Survey (SDSS) with dark-matter halo masses above ${\\rm M}_{\\rm halo}=2.5\\times10^{13}{\\rm M}_{\\sun}$.   Ellipticals galaxies make up the bulk of the massive end of the red-sequence population with optical colors indicative of their non-star-forming and old stellar nature. Despite a quiet star-formation history over the last 6--8 billion years \\citep{bell05a}, the total stellar mass density on the red sequence has roughly doubled over this interval \\citep{bell04b,blanton06,borch06,faber07,brown07} and now accounts for more than half of the present-day budget \\citep{hogg02,bell03b}, providing strong observational evidence for the ongoing hierarchical growth of the massive galaxy population. These results were derived from red galaxy number densities over a wide range of stellar masses above and below $10^{11}{\\rm M}_{\\sun}$. Owing to the scarcity of the highest-mass galaxies, cosmic variance, and systematic uncertainties in stellar mass estimates, any increase in the number density of ${\\rm M}_{\\rm star}>10^{11}{\\rm M}_{\\sun}$ galaxies is poorly constrained, resulting in controversy over whether this population has continued to grow slowly \\citep[e.g.,][]{brown07} or has been effectively static \\citep[e.g.,][]{scarlata07}, since $z\\sim1$.  Besides number density evolution, mergers of sufficiently massive galaxies could provide a more clear indication for some continued stellar mass growth in the high-mass galaxy population. The existence of a handful of massive red mergers over the redshift interval $0.1<z<0.9$ \\citep{vandokkum99,tran05,bell06a,lotz06,rines07} proves that the growth is non-zero at high stellar masses and implies that this mechanism does contribute to the assembly of galaxies at the top of the food chain. Yet, the importance of this process and the related rate of mass growth are highly uncertain given the tiny samples over this large cosmic time interval. Indirect measures such as the presence of faint tidal debris or shells around many local massive ellipticals \\citep{vandokkum05,mihos05}, the isophotal properties of giant ellipticals \\citep{kang07}, the lack of evolution of the stellar mass-size relation of red spheroids since $z=1$ \\citep{mcintosh05b}, and the lack of morphological evolution on the red sequence since $z=0.7$ \\citep{bell04a} provide a variety of limits to the importance of dissipationless mergers. Perhaps the most powerful method for obtaining estimates for the stellar mass growth rate via major merging is based on small-scale clustering statistics that provide an accurate measurement of close pair frequencies in real space \\citep{masjedi06,bell06b,masjedi07}. However, this method likely yields an overestimate of the merger frequency because it assumes that all close pairs will merge. All estimates of merger-driven growth rates are limited by uncertainties in the time interval over which a pair will merge, and over what duration an object could be identified as interacting. \\citet{masjedi07} find a very small growth rate (1--2\\% per Gyr) at $z\\sim0.25$ for major mergers involving at least one progenitor drawn from the SDSS Luminous Red Galaxy sample \\citep[LRG,][]{eisenstein01}; LRGs have typical masses of several times $10^{11}{\\rm M}_{\\sun}$. To date there remains no direct evidence of ongoing merger-driven assembly of massive galaxies at $z<0.1$, and the LRG result implies that this formation process is no longer important. These facts motivate a thorough search for the existence/nonexistence of ongoing examples in the present-day universe.  While the aforementioned statistical method for finding close physical pairs is powerful, it does not isolate actual merging systems and thus provides no information on the progenitor properties of massive merger remnants. Recent numerical simulations and models make a range of predictions regarding the progenitor morphologies at the time of the last major merger \\citep{khochfar03,naab06a,kang07}, yet robust observational constraints are missing for ${\\rm M}_{\\rm star}>10^{11} {\\rm M}_{\\sun}$ systems. Many studies have identified major-merger candidates by either close pairs \\citep{carlberg94,patton00,carlberg00,patton02,bundy04,lin04} or disturbed morphologies \\citep{lefevre00,conselice03,lotz06}, but these samples mostly contain major mergers between lower-luminosity galaxies that tend to be gas-rich spiral disks. Numerical simulations show that such dissipative merging of disk galaxies will not produce massive pressure-supported ellipticals \\citep[e.g.,][]{naab06d}. As mentioned above, only circumstantial evidence and a small number of red galaxy pairs with $z<0.9$ support the existence of mergers likely to produce massive ellipticals. Our understanding of the progenitors is therefore very limited. Here we present a thorough census of 38 massive merger pairs from SDSS, providing an order-of-magnitude increase in the number of such detections at $z<0.5$ and allowing an improved understanding of their progenitor properties.  While many estimates of major merger rates are found in the literature, to date no measure of the environmental dependence of merger-driven mass growth has been attempted. In the standard cosmological model, there is a trade off between the expansion of the universe and the gravitational collapse of dark and luminous matter. Therefore, the rate at which stellar mass is assembled at the centers of the largest dark matter halos over recent cosmic history is a fundamental aspect of the ongoing formation of large-scale structure, and the rate that high-mass galaxies form by mergers as a function of halo mass constrains galaxy formation theories. Some theories predict that the mergers producing massive ellipticals occur preferentially in groups rather than in high-density cluster or low-density field environments because the smaller velocity dispersions allow more galaxy interactions \\citep{cavaliere92}; also dynamical friction is more efficient in lower-mass halos \\citep[e.g.][]{cooray05d}. Others predict that the brightest cluster galaxies (BCGs) grow by hierarchical merging (``galactic cannibalism'') at the centers of the dark-matter potential wells of large clusters \\citep{ostriker75,merritt85,dubinski98,cooray05d}. A handful of low-redshift BCGs show multiple nuclei suggesting cannibalism in the form of multiple minor mergers \\citep{lauer88}, but there are no observations of major mergers at the centers of clusters. In this paper we make use of the statistically large SDSS group catalog \\citep{yang05a,weinmann06a} to show that major mergers occur in present-day dense environments, and to explore the halo-mass dependence and central/satellite identity of merger-driven massive galaxy assembly.  Throughout this paper we calculate comoving distances in the $\\Lambda$CDM concordance cosmology with $\\Omega_{\\rm m} = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and assume a Hubble constant of $H_0 = 70 $\\,km\\,s$^{-1}$\\,Mpc$^{-1}$. SDSS magnitudes are in the AB system. ",
        "conclusions": "\\label{sec:Disc}  We find the first direct observational evidence for an important population of galaxy-galaxy mergers with total stellar masses above $10^{11} {\\rm M}_{\\sun}$ in the local universe. These objects provide an unprecedented census of the progenitor properties for the merger-driven assembly of high-mass galaxies, which we compare to recent predictions from numerical models of galaxy formation and evolution. Moreover, the existence of these mergers prove that a measurable amount of stellar mass growth continues in the massive galaxy population at present times, and we compare estimates based on this sample with other estimates in the literature. Finally, we have identified mergers restricted to reside in large SDSS groups and clusters with $z\\leq0.12$, thus allowing the first constraints on the halo-mass dependencies of recent massive merger activity. While it is well-established that massive galaxies are more common in such high-density environments, we are missing much more than 50\\% of the population with ${\\rm M}_{\\rm star}<4\\times10^{11} {\\rm M}_{\\sun}$ in the local volume, as Table \\ref{mstar_counts} shows. Therefore, we must keep this caveat in mind when interpreting the conditions for which our results hold. In an upcoming study, we are examining the role of major mergers as a function of stellar mass over the full range of environments hosting galaxies more massive than ${\\rm M}_{\\rm star}=5\\times10^{10}{\\rm M}_{\\sun}$.  \\subsection{Massive Merger Progenitors: Observations Meet Theories} \\label{sec:Disc1}  Establishing the luminosity dependence of elliptical (E) galaxy properties \\citep{davies83,bender88,bender92} set the stage for theories regarding the types of merger progenitors that would produce the characteristics of low and high-mass early-type galaxies (ETGs)\\footnote{The distinction between elliptical and early-type galaxies is often blurred in the literature. We consider Es to be a morphological subset of ETGs, which are concentrated and spheroid-dominated systems including Es, lenticulars (S0s), and Sa spirals. When referencing other authors we remain faithful to their choice of nomenclature.} galaxies \\citep{bender92,kormendy96,faber97}. We concentrate on modern numerical simulations and semi-analytic models that attempt to reproduce the kinematic, photometric, and structural properties observed in massive Es through major merging \\citep{naab99,naab03,khochfar03,khochfar05,naab06a,boylan06,kang07}. For this discussion we make the straight-forward assumption that the major mergers that we have identified will produce remnants that are {\\it not unlike} the ${\\rm M}_{\\rm star}>10^{11} {\\rm M}_{\\sun}$ galaxy population already in place. We can only guess at remnant properties (see Fig. \\ref{fig:cm.remn}), but in general, massive galaxies on the red-sequence are typically early-type.  As we show in Figure \\ref{fig:prog.mass}, the progenitor masses are comparable for the most part, and quantitatively consistent with the LRG-LRG merger mass spectrum from \\citet{masjedi07} under the assumption that companions merge on dynamical friction time scales. $N$-body simulations \\citep[e.g.][]{naab99} have long shown that ${\\rm M}_1/{\\rm M}_2\\approx1$ are necessary to produce the lack of significant rotation observed in massive Es. Yet, a near unity mass ratio alone is not sufficient to produce the predominance of boxy and anisotropic Es found at high luminosity \\citep{naab03,naab06a}. To match the decreasing fraction of rotational support and increasing fraction of boxiness in more luminous Es, the role of gas dissipation must be significantly reduced at high masses \\citep{bender92,khochfar05,naab06d,kang07}, and recent ETG-ETG merger simulations have demonstrated this numerically \\citep{naab06a}. Figures \\ref{fig:prog.type} and \\ref{fig:cm.prog} show that 90\\% of the progenitors in this study have concentrated light profiles and red-sequence colors, both common attributes of ETGs, with little or no cold gas content. In addition, the tidal signatures of the bulk of these massive mergers (see Figs. \\ref{fig:dpairs1} \\& \\ref{fig:dpairs2}) match those of observed \\citep{bell06a} and simulated \\citep{naab06a} major dissipationless (or gas-poor) merging of ETGs. Thus, our sample represents a more than order-of-magnitude increase in the number of such known systems with $z<0.2$, and demonstrates that dissipationless merging is indeed an important channel for the formation of massive galaxies.  Finally, we compare the observed high fraction of ETG-ETG mergers ($f_{\\rm ETG-ETG}=0.9$) with several semi-analytic predictions. Recall that we have looked for signs of interaction in $>200$ major pairs from a total sample of $>5000$ massive galaxies (i.e., sampM), yet only 10\\% of the 38 mergers we identify could possibly form a ${\\rm M}_{\\rm star}>10^{11}{\\rm M}_{\\sun}$ remnant by other than an ETG-ETG merger. The progenitor morphologies of this study best match the predictions of \\citet{khochfar03}, who find $f_{\\rm E-E}=0.75$ for the last major merger of $4L^{\\ast}$ remnants, independent of environment. We find much larger ETG-ETG fractions than \\citet{naab06a} who predict only 20-35\\% (also independent of environment) over the estimated mass range of our merger remnants ($11.1<\\log_{10}{({\\rm M}_{\\rm star}/{\\rm M}_{\\sun})}<11.7$), and \\citet{kang07} who predict $f_{\\rm ETG-ETG}<0.1$ for $\\log_{10}{({\\rm M}_{\\rm star}/{\\rm M}_{\\sun})}>11$. We note that these predicted progenitor morphologies for present-day Es are based on the final major mergers that could occur over a large redshift range out to $z\\sim1$, which could be different in nature to those that occur in the short time interval that we observe. Moreover, we focus on high-density environments known to have very few massive late-type (blue) galaxies \\citep{butcher78a}, which might explain the low number of ``mixed'' (early-late or elliptical-spiral) mergers that we find. Hence, for these models to be consistent with our data, either (1) the $f_{\\rm ETG-ETG}$ of present-day major mergers depends on halo mass (i.e., environment), or (2) the relative importance of major mixed mergers has decreased significantly since $z=1$.  \\subsection{Estimating Stellar Mass Accretion Rates}  The existence of massive dissipationless mergers at low redshift is direct observational evidence that the growth of ${\\rm M}_{\\rm star}>10^{11}{\\rm M}_{\\sun}$ galaxies continues at present times in agreement with many cosmologically-motivated simulations \\citep{khochfar05,delucia06,kaviraj07,kang07}. Moreover, even under conservative assumptions that limit the amount of companion mass that is added to massive CEN galaxies, all of our sample will still result in remnants with ${\\rm M}_{\\rm star}>10^{11}{\\rm M}_{\\sun}$. Previously, the observational evidence for recent merger-based assembly of $z\\sim0$ massive Es was limited to luminous/massive galaxy clustering statistics \\citep{masjedi06,bell06b,masjedi07} or post-merger signatures that cannot distinguish between minor and major merging; e.g., tidal shells \\citep{malin83}, fine structure \\citep{schweizer92}, faint tidal features \\citep{vandokkum05,mihos05}, or kinematic/photometric properties \\citep[e.g.,][]{kang07}. With the merger sample presented here we can quantify directly the amount of growth, occurring in dense environments, at the high-mass end of the stellar mass function.  Going from the observed merger counts to an inferred merger rate is limited mostly by the uncertainty in the merger timescale ($t_{\\rm merg}$) that one assumes. Numerical models show that the time interval for two galaxies to interact and finally merge into a single remnant depends critically on the orbital parameters, progenitor mass ratios and densities, and the degree to which the merger is dissipationless. For major mergers of massive galaxies a number of different $t_{\\rm merg}$ have been put forth in the literature based on simple orbital timescale arguments. For example, \\citet{masjedi06} derived a reasonable lower limit of $t_{\\rm merg}=0.2$ Gyr for a close ($d_{1,2}=10$ kpc) pair of LRG galaxies with a velocity dispersion of $\\sigma=200$ \\kms. Naturally, bound pairs with $d_{1,2}>10$ kpc separation will take longer to merge. \\citet{bell06b} made a similar calculation for somewhat less-massive galaxies typically separated by $d_{1,2}=15$ kpc and estimated $t_{\\rm merg}=0.4$ Gyr and argued for at least a factor of two uncertainty in this time. The mergers in this study have an average projected separation of 15.5 kpc (see Fig. \\ref{fig:pair.mergVnon}), so in what follows, we adopt $t_{\\rm merg}=0.4^{+0.4}_{-0.2}$ Gyr with conservative error bars that encompass the range of uncertainties discussed in the literature.  Here, we compute the rate of stellar mass accretion by major merging onto massive galaxies in large groups. First, we find that the total mass accreted onto the centers of the $N_{\\rm CEN}=845$ halos that we study is $\\sum{ f{\\rm M}_{{\\rm s},i}}=3.9(3.5)\\times10^{12} {\\rm M}_{\\sun}$, if we include (exclude) the four SAT-SAT mergers at their host's dynamical center (see \\S \\ref{sec:cenmerging}). ${\\rm M}_{{\\rm s},i}$ is the stellar mass of the secondary (SAT) galaxy in the $i^{\\rm th}$ CEN-SAT merger, and $f$ is the fraction of ${\\rm M}_{{\\rm s},i}$ that winds up as part of the CEN galaxy. The rate of stellar mass buildup per massive CEN galaxy is therefore \\begin{equation} \\dot{{\\rm M}}_{\\rm CEN} = \\frac{\\sum{ f{\\rm M}_{{\\rm s},i} }}{N_{\\rm CEN}} \\times \\frac{1}{t_{\\rm merg}} , \\label{eq:4} \\end{equation} or between $1.0^{+1.0}_{-0.5}\\times10^{10}{\\rm M}_{\\sun}{\\rm Gyr}^{-1}$ and $1.2^{+1.1}_{-0.6}\\times10^{10}{\\rm M}_{\\sun}{\\rm Gyr}^{-1}$, depending on which sample of CEN-SAT mergers that we consider. The lopsided error bars result from the range of accretion rates for $t_{\\rm merg}=0.4^{+0.4}_{-0.2}$ Gyr, as described above. If we divide all of these accretion rates by $2.69\\times10^{11} {\\rm M}_{\\sun}$, the average stellar mass of the 845 CEN galaxies in this study, we find that each CEN is growing by 1--9\\% per Gyr. Finally, these values can be decreased by assuming $f<1$ in (\\ref{eq:4}), but as we discuss in \\S \\ref{sec:cmplane}, $f=0.5$ represents a likely lower limit.  Likewise, the total stellar mass accreted onto all galaxies in sampM is $\\sum{ f{\\rm M}_{{\\rm s},i}}+\\sum{ {\\rm M}_{{\\rm s},j}}=5.1\\times10^{12} {\\rm M}_{\\sun}$, where ${\\rm M}_{{\\rm s},j}$ is the mass of the secondary (SAT) galaxy in the $j^{\\rm th}$ SAT-SAT merger. Therefore, the growth per ${\\rm M}_{\\rm star}\\geq5\\times10^{10} {\\rm M}_{\\sun}$ galaxy in high-mass groups is \\begin{equation} \\dot{{\\rm M}}_{(\\geq5\\times10^{10} {\\rm M}_{\\sun})} = \\frac{\\sum{ f{\\rm M}_{{\\rm s},j} } + \\sum{ {\\rm M}_{{\\rm s},j}} }{ (N_{\\rm CEN}+N_{\\rm SAT}-N_{\\rm s,sampM})} \\times \\frac{1}{t_{\\rm merg}} , \\end{equation} where $N_{\\rm s,sampM}=12$ is the number of secondary SAT galaxies in sampM that are involved in major mergers and must be subtracted to avoid double counting. We find $\\dot{{\\rm M}}_{(\\geq5\\times10^{10} {\\rm M}_{\\sun})}=2.4^{+2.4}_{-1.2}\\times10^{9}{\\rm M}_{\\sun}{\\rm Gyr}^{-1}$; if we assume $f=0.5$ for CEN-SAT mergers only we find $\\dot{{\\rm M}}_{(\\geq5\\times10^{10} {\\rm M}_{\\sun})}=1.6^{+1.5}_{-0.7}\\times10^{9}{\\rm M}_{\\sun}{\\rm Gyr}^{-1}$. Given that the average stellar mass of sampM galaxies is $1.04\\times10^{11} {\\rm M}_{\\sun}$, we find that every massive galaxy is growing by 1--5\\% per Gyr. Even though SAT-SAT mergers may occur as frequently as CEN-SAT mergers in these massive groups, the centers are where much of the mass growth takes place. It is clear from Figure \\ref{fig:prog.mass} that mostly only ${\\rm M}_{\\rm star}>10^{11} {\\rm M}_{\\sun}$ galaxies build up in mass by major mergers in groups with ${\\rm M}_{\\rm halo}>2.5\\times10^{13}{\\rm M}_{\\sun}$. In contrast, we find few mergers among the $5\\times 10^{10}<{\\rm M}_{\\rm star}<10^{11} {\\rm M}_{\\sun}$ galaxies in these high-mass groups, which make up the bulk (60\\%) of sampM. This suggests that if major merging is playing an important role in the strong mass growth observed on the red sequence below M* \\citep{bell04b,blanton06,borch06,faber07,brown07}, it is occurring in lower-mass groups than we study here.   Rather than mass growth rates we can use the same line of reasoning to estimate massive galaxy-galaxy merging rates of $(21+4)/845/t_{\\rm merg}=0.074^{+0.074}_{-0.037}{\\rm Gyr}^{-1}$ for CEN-SAT and $(38+7)/(845+4531-12)/t_{\\rm merg}=0.021^{+0.021}_{-0.011}{\\rm Gyr}^{-1}$ for all galaxies in sampM. For these estimates we included the seven additional major mergers (4 CEN, 3 SAT) we identified by their highly-disturbed appearance. \\citet{masjedi06} found a strict upper limit to the LRG-LRG rate of only $0.006{\\rm Gyr}^{-1}$. We estimate that LRGs have a stellar mass range of $11.4<\\log_{10}{({\\rm M}_{\\rm star}/{\\rm M}_{\\sun})}<12.0$, based on typical red-sequence colors and luminosities between $4L^{\\ast}$ and $25L^{\\ast}$. Within these mass limits, we find a merger rate of $5/462/t_{\\rm merg}=0.027^{+0.027}_{-0.014}{\\rm Gyr}^{-1}$ on the red sequence, or 2--9 times the LRG-LRG rate. In Table \\ref{mstar_counts}, we show that the high-mass groups that we study contain $>70\\%$ of the very-massive, red galaxy population in the $z\\leq0.12$ volume of DR2, with the vast majority being CENs. Yet, the same group selection contains only 30\\% of the population of $11.4<\\log_{10}{({\\rm M}_{\\rm star}/{\\rm M}_{\\sun})}<11.6$ systems. These numbers show that a significant portion of the local counterparts to LRGs are found in groups with ${\\rm M}_{\\rm halo}<2.5\\times10^{13}{\\rm M}_{\\sun}$. Therefore, we conclude that LRG-LRG merging occurs more frequently in the more massive groups."
    },
    "0710/0710.1542_arXiv.txt": {
        "abstract": "{The initial-final mass relationship of white dwarfs, which is poorly constrained, is of paramount importance for different aspects in modern astrophysics. From an observational perspective, most of the studies up to now have been done using white dwarfs in open clusters.} {In order to improve the initial-final mass relationship we explore the possibility of deriving a semi-empirical relation studying white dwarfs in common proper motion pairs. If these systems are comprised of a white dwarf and a FGK star, the total age and the metallicity of the progenitor of the white dwarf can be inferred from the detailed analysis of the companion.} {We have performed an exhaustive search of common proper motion pairs containing a DA white dwarf and a FGK star using the available literature and crossing the SIMBAD database with the Villanova White Dwarf Catalog.  We have acquired long-slit spectra of the white dwarf members of the selected common proper motion pairs, as well as high resolution spectra of their companions.  From these observations, a full analysis of the two members of each common proper motion pair leads to the initial and final masses of the white dwarfs.} {These observations have allowed us to provide updated information for the white dwarfs, since some of them were misclassified.  In the case of the DA white dwarfs, their atmospheric parameters, masses, and cooling times, have been derived using appropriate white dwarf models and cooling sequences.  From a detailed analysis of the FGK stars spectra we have inferred the metallicity. Then, using either isochrones or X-ray luminosities we have obtained the main-sequence lifetime of the progenitors, and subsequently their initial masses.} {This work is the first one in using common proper motion pairs to improve the initial-final mass relationship, and has also allowed to cover the poorly explored low-mass domain. As in the case of studies based on white dwarfs in open clusters, the distribution of the semi-empirical data presents a large scatter, which is higher than the expected uncertainties in the derived values.  This suggests that the initial-final mass relationship may not be a single-valued function.} ",
        "introduction": "White dwarfs are the final remnants of low- and intermediate-mass stars. About 95\\% of main-sequence stars will end their evolutionary pathways as white dwarfs and, hence, the study of the white dwarf population provides details about the late stages of the life of the vast majority of stars.  Since white dwarfs are long-lived objects, they also constitute useful objects to study the structure and evolution of our Galaxy (Liebert et al.~2005a; Isern et al. 2001). For instance, the initial-final mass relationship (IFMR), which connects the properties of a white dwarf with those of its main-sequence progenitor, is of paramount importance for different aspects in modern astrophysics. It is required as an input for determining the ages of globular clusters and their distances, for studying the chemical evolution of galaxies, and also to understand the properties of the Galactic population of white dwarfs.  Despite its relevance, this relationship is still poorly constrained, both from the theoretical and the observational points of view.  The first attempt to empirically determine the initial-final mass relationship was undertaken by \\cite{wei77}, who also provides a recent review on this subject (Weidemann 2000). It is still not clear how this function depends on the mass and metallicity of the progenitor, its angular momentum, or the presence of a strong magnetic field. The total age of a white dwarf can be expressed as the sum of its cooling time and the main-sequence lifetime of its progenitor. The latter depends on the metallicity of the progenitor of the white dwarf, but it cannot be determined from observations of single white dwarfs.  This is because white dwarfs have such strong surface gravities that gravitational settling operates very efficiently in their atmospheres, and any information about their progenitors (e.g. metallicity)  is lost in the very early evolutionary stages of the cooling track. Moreover, the evolution during the AGB phase of the progenitors is essential in determining the size and composition of the atmospheres of the resulting white dwarfs, since the burning processes that take place in H and He shells determine their respective thicknesses and their detailed chemical compositions, which are crucial ingredients for determining the evolutionary cooling times.  A promising approach to circumvent the problem, and also to directly test the initial-final mass relationship, is to study white dwarfs for which external constraints are available.  This is the case of white dwarfs in open and globular clusters (Ferrario et al.~2005, Dobbie et al.~2006) or in non-interacting binaries, for instance, common proper motion pairs (Wegner 1973, Oswalt et al.~1988).  Focusing on the latter, it is sound to assume that the members of a common proper motion pair were born simultaneously and with the same chemical composition. Since the components are well separated (100 to 1000 AU), mass exchange between them is unlikely and it can be considered that they have evolved as isolated stars. Thus, important information of the white dwarf, such as its total age or the metallicity of the progenitor, can be inferred from the study of the companion.  In particular, if the companion is an F, G or K type star the metallicity can be derived with high accuracy from detailed spectral analysis. On the other hand, the age can be obtained using different methods. In particular, we will use stellar isochrones when the star is moderately evolved, or the X-ray luminosity if the star is very close to the ZAMS.  The purpose of this work is to present our spectroscopic analysis of both members of some common proper motion pairs containing a white dwarf, and the semi-empirical intial-final mass relationship that we have derived from this study. The paper is organized as follows.  In \\S 2 we present the observations done so far and describe the data reduction.  Section 3 is devoted to discuss the classification and the analysis of the observed white dwarfs, whereas in \\S 4 we present the analysis of the companions. This is followed by \\S 5 where we present our main results and finally in \\S 6 we elaborate our conclusions. ",
        "conclusions": "We have studied a sample of common proper motion pairs comprised of a white dwarf and a FGK star. We have performed high signal-to-noise low resolution spectroscopy of the white dwarf members, which led us to carry out a full analysis of their spectra and to make a re-classification when necessary.  From the fit of their spectra to white dwarf models we have derived their atmospheric parameters. Then, using different cooling sequences --- namely those of \\cite{sal00} and \\cite{fon01} --- their masses and cooling times were obtained.  Simultaneously, we have performed independent high resolution spectroscopic observations of their companions.  Using the available photometry we have obtained their effective temperatures.  Then, from a detailed analysis of their spectra and using either isochrones or X-ray luminosities, we have derived their metallicities and ages (i.e., the metallicities of the progenitors of the white dwarfs and their total ages).  These observations allowed us to obtain the initial and final masses of six white dwarfs in common proper motion pairs, four of them corresponding to initial masses below $2\\,\\rm M_{\\sun}$, a range which has not been previously covered by the open cluster data.  Our semi-empirical relation shows significant scatter, compatible with the results obtained by \\cite{fer05} and \\cite{dob06}, which are mainly based on open cluster data.  However, the dispersion of the results is higher than the error bars, which leaves some open questions that should be studied in detail (e.g., rotation or magnetic fields).  We have shown that common proper motion pairs containing white dwarfs can be useful to improve the initial-final mass relationship, since they cover a wide range of ages, masses and metallicities, and they are also representative of the disk white dwarf population.  We have seen that the accuracy in the total ages depends almost exclusively on the evolutionary state of the low-mass companions. Such relative accuracy becomes poor when the star is close to the ZAMS. However, this limitation may not be critical to many common proper motion pairs. Planned deep surveys like GAIA, LSST or the Alhambra Survey will discover thousands of new white dwarfs, some of them belonging to wide binaries. In the meantime, our most immediate priority is to further extend the sample of wide binaries valid for this study.  We are working in the search for more wide binaries of our interest in the NLTT catalog (Gould \\& Chanam\\'e 2004) and also in the LSPM-north catalog (L\\'epine \\& Bongiorno 2007). Detailed study of the current and future common proper motion pairs of this type should help to explain the scatter in the semi-empirical initial-final mass relationship and to discern whether this is a single-valued function. If consistency between observations and theoretical calculations is found, this would have a strong impact on stellar astrophysics, since this relationship is used in many different areas, such as chemical evolution of galaxies, the determination of supernova rates or star formation and feedback processes in galaxies."
    },
    "0710/0710.4151_arXiv.txt": {
        "abstract": "The discovery of dark energy (DE) as the physical cause for the accelerated expansion of the Universe is the most remarkable experimental finding of modern cosmology. However, it leads to insurmountable theoretical difficulties from the point of view of fundamental physics. Inflation, on the other hand, constitutes another crucial ingredient, which seems necessary to solve other cosmological conundrums and provides the primeval quantum seeds for structure formation. One may wonder if there is any deep relationship between these two paradigms. In this work, we suggest that the existence of the DE in the present Universe could be linked to the quantum field theoretical mechanism that may have triggered primordial inflation in the early Universe. This mechanism, based on quantum conformal symmetry, induces a logarithmic, asymptotically-free, running of the gravitational coupling. If this evolution persists in the present Universe, and if matter is conserved, the general covariance of Einstein's equations demands the existence of dynamical DE in the form of a running cosmological term, $\\CC$, whose variation follows a power law of the redshift. ",
        "introduction": "Modern Cosmology incorporates the notion of dark energy (DE) as an experimental fact that accounts for the physical explanation of the observed accelerated expansion\\,\\cite{SNe,WMAP3Y}. Although the nature of the DE is not known, one persistent possibility is the 90-years-old cosmological constant (CC) term, $\\CC$, in Einstein's equations. In recent times, one is tempted to supersede this hypothesis with another, radically different, one: viz. a slowly evolving scalar field $\\phi$ (``quintessence'') whose potential, $V(\\phi)\\gtrsim 0$, could explain the present value of the DE and whose equation of state (EOS) parameter $\\omega_{\\phi}= p_{\\phi}/\\rho_{\\phi}\\simeq -1+\\dot\\phi^2/V(\\phi)$ is only slightly larger than $-1$ (hence insuring a negative pressure mimicking the $\\CC$ case)\\,\\cite{weinberg}. The advantage to think this way is that the DE can then be a dynamical quantity taking different values throughout the history of the Universe. However, this possibility can not explain why the DE is entirely due to such an \\textit{ad hoc} scalar field and why the contributions to the vacuum energy from the other fields (e.g. the electroweak Standard Model ones) must not be considered. In short, it does not seem to be such a wonderful idea to invent the field $\\phi$ and simply replace $\\rL=\\CC/8\\pi\\,G$ (the energy density associated to $\\CC$, where $G$ is Newton's constant) with $\\rho_{\\phi}\\simeq V(\\phi)$. One has to explain, too, why the various contributions (including the additional one $V(\\phi)$!) must conspire to generate the tiny value of the DE density at present -- the ``old CC problem''\\,\\cite{weinberg}. While we cannot solve this problem at this stage, the dynamical nature of the DE makes allowance for this possibility. Furthermore, since there is no obvious gain in the quintessence idea, we stick to the CC approach, although we extend it to include the possibility of a dynamical (``running'') $\\CC$ term\\,\\cite{JHEPCC1,SSRev}. The obvious question now is: where this dynamics could come from?  One possibility is that it could originate from the fundamental mechanism of inflation\\,\\cite{inflation}, which presumably took place in the very early Universe and could have left some loose end or remnant -- kind of ``fossil'' -- in our late Universe, which we don't know where to fit in now. However, what mechanism of inflation could possibly do that? There is in principle a class of distinct possibilities, in particular see \\cite{Sahni98,Mongan01}, but our very source of inspiration here is the quantum theory of the conformal factor, which was extensively developed in\\,\\cite{Conformal1}. For a recent discussion, see e.g. \\cite{Conformal2,Conformal3} and references therein. More specifically, we start from the idea of ``tempered anomaly-induced inflation'', which was first proposed in \\cite{shocom,DESY} (see also \\cite{PST}). It leads essentially to a modified form of the original Starobinsky model\\,\\cite{Starobinsky}. In the present paper, we push forward the possibility that the mechanism that successively caused, stabilized, slowed down (``tempered'') and extinguished the fast period of inflation in our remote past could have left an indelible imprint in the current Universe, namely a very mild (logarithmically) running Newton's coupling $G$. We show that, if matter is covariantly conserved, this necessarily implies an effective renormalization group (RG) running of the ``cosmological constant'' energy density, $\\rL=\\rL(a)$, which takes the form of a cubic law of $a^{-1}=1+z$ during the matter dominated epoch ($a$ being the scale factor and $z$ the cosmological reshift). ",
        "conclusions": "In this work we have suggested that the presence of dynamical dark energy (DE) in the current Universe is actually a consistency demand of Einstein equations under the two assumptions of: i) matter conservation, and ii) the existence of a period of primordial inflation in the early Universe, especially when realized as ``tempered anomaly-induced inflation''. Based essentially on the previous works\\,\\cite{shocom,DESY,PST} and on the general setting of the quantum theory of the conformal factor\\,\\cite{Conformal1,Conformal2}, we have found that if the inflationary mechanism is caused by quantum effects on the effective action of conformal quantum field theory in curved space-time, then the gravitational coupling $G$ becomes a running quantity of the scale factor, $G(a)=G_0/(1-\\f\\ln a)$, $\\f$ being the coefficient of the $\\beta$-function for the conformal Newton's coupling. The effect of this coupling on the inflationary dynamics is to efficiently ``temper'' the regime of stable inflation presumably into the FLRW regime. The rigorous high energy calculation of $\\f$ in QFT in curved space-time shows that both fermions and bosons produce non-negative contributions ($\\f\\geq 0$). As a consequence, $G$ becomes an asymptotically-free coupling of the scale factor. Intriguingly enough, we have suggested the possibility that this running might persist in the present Universe and, if so, it could provide a \\textit{raison d'\\^etre} for the existence of the (dynamical) DE, which would appear in the form of running cosmological vacuum energy $\\rL$. In fact, the logarithmic evolution of $G$ induces a power-law running of $\\rL$, which is essentially driven by the soft-decoupling terms $\\sim H^2\\,M_i^2$ (hence by the heaviest particle masses). The result is a Universe effectively filled with a mildly-dynamical DE, which can be perfectly consistent with the present observations.  To summarize, from the point of view of the ``RG-cosmology'' under consideration, the current Universe appears as FLRW-like while still carrying some slight imprints of important physical processes that determined the early stages of the cosmic evolution. Most conspicuously, the smooth dynamics of $G$ and $\\rL$ can be thought of as ``living fossils'' left out of the quantum field theoretical mechanism that triggered primordial inflation. Remarkably, this framework fits with previous attempts to describe the renormalization group running of the cosmological term\\,\\cite{JHEPCC1,RGTypeIa,IRGA03,SSS1,croat,SS12,LXCDM12,Bilic07} and could provide an attractive link between all stages of the cosmic evolution.  It is reassuring to find that there is a large class of RG models behaving effectively the same way. Differences between them could probably be resolved at the level of finer tests, such as those based on cosmological perturbations and structure formation. For example, in references\\,\\cite{FSS1,GOPS} it is shown that the study of cosmological perturbations within models of running cosmological constant puts a limit on the amount of running, which is more or less stringent depending on the peculiarities of the model. Similarly, a particular study of perturbations would be required in the present framework (which includes the variation of both $\\Lambda$ and $G$) to assess the implications on the parameter $\\f$. This study is beyond the scope of the present work.  \\vspace{1cm}  \\textit{Acknowledgements}. I am very grateful to Ilya Shapiro for discussions on different aspects of this work and for the fruitful collaboration maintained on the cosmological constant ptoblem over the years. I thank also Ana Pelinson for interesting discussions in the early stages of this work. The author has been supported in part by MECYT and FEDER under project 2004-04582-C02-01, and also by DURSI Generalitat de Catalunya under 2005SGR00564 and the Brazilian agency FAPEMIG. I am thankful for the warm hospitality at the Dept. of Physics of the Univ. Federal de Juiz de Fora, where part of this work was carried out."
    },
    "0710/0710.1318_arXiv.txt": {
        "abstract": "We present preliminary results on the calculation of synthetic spectra obtained with the stellar model atmospheres developed by Cardona, Crivellari, and Simonneau. These new models have been used as input within the {\\sc Synthe} series of codes developed by Kurucz. As a first step we have tested if {\\sc Synthe} is able to handle these models which go down to $\\log{\\tau_{\\rm Ross}}= -13$. We have successfully calculated a synthetic solar spectrum in the wavelength region 2000--4500~\\AA\\ at high resolution ($R=522\\,000$). Within this initial test we have found that layers at optical depths with $\\log{\\tau_{\\rm Ross}} < -7$ significantly affect the mid-UV properties of a synthetic spectrum computed from a solar model. We anticipate that these new extended models will be a valuable tool for the analysis of UV stellar light arising from the outermost layers of the atmospheres. \\end {abstract} ",
        "introduction": "\\label{sec:1} A set of spectral energy distributions (SEDs) is a very useful tool to analyze stellar spectra and the integrated spectral properties of stellar systems (via some evolutionary population synthesis code, see e.g. Buzzoni 1995). Observational atlases of stellar SEDs generally lack of homogeneous and complete coverage of the main stellar parameters ($T_{\\rm eff}, \\log{g}$ and [M/H]), therefore many of the recent population analyses rely on results of stellar atmosphere modelling, a practise that has been eased with the development of faster computers and sophisticated computational codes. In order to calculate a theoretical stellar SED it is necessary to have a model atmosphere which describes the physical quantities at different depths and, ideally, a complete set of opacities that account for the absorption of the radiation passing throughout the atmosphere.  During the last decades several groups  have developed computational codes capable to calculate model atmospheres and spectra at high resolution, some of whom have allowed the public use of their codes, among others, the codes {\\sc Atlas9}, {\\sc Atlas12} and {\\sc Synthe} built by Kurucz\\footnote{http://kurucz.harvard.edu/} (1993a,b) and {\\sc Tlusty} and {\\sc Synspec} constructed by Hubeny \\& Lanz\\footnote{http://nova.astro.umd.edu/} (1992). In particular, {\\sc Atlas9} and {\\sc Synthe} codes have been used by our group to calculate the UVBLUE grid of theoretical SEDs and to investigate its potential for stellar and populations studies in the ultraviolet (UV) wavelength interval (see Rodriguez-Merino et al. 2005 for more details).  The UV wavelength range has historically been challenging in many branches of modern astrophysics since the observed stellar spectra are not well reproduced by predictions of theoretical models. One possible reason is the missing opacity problem (see Holweger 1970; Gustafsson et al. 1975), but another probable reason is actually that most of the model atmospheres do not provide the atmosphere structure near the stellar surface, where most of the UV radiation emerges. Therefore, it is crucial to calculate new models which describe the outermost layers of the stellar atmospheres.  In this work we briefly describe the structure of a new model atmosphere for the Sun  which incorporates layers down to $\\log{\\tau_{\\rm Ross}}=-13$. That is, optical depths more than five orders of magnitude thinner compared to classical models currently in use. This model has been couplet to {\\sc Synthe} codes for testing their compatibility and exploring the effects of such layers on the UV flux. ",
        "conclusions": "\\label{sec:4} The main result of this work is that {\\sc Synthe} series of codes is capable of treating the CCS model atmospheres, which reach very low values of $\\tau_{\\rm Ross}$. The analysis of the effects of extending the atmosphere indicates that at mid-UV wavelengths the effects are significant while negligible in the blue. The following steps are to extend the analysis to models with atmospheric parameters different of the Sun, to complement the opacity (both continuous and of lines) as well as to introduce more chemical species. We are in the process of including convection for intermediate and cool star models. A detailed comparison at high resolution with an observed solar atlas (Kurucz et al. 1984) is also underway."
    },
    "0710/0710.0128_arXiv.txt": {
        "abstract": "% Preliminary VLBA polarisation results on 6 ``blazars'' from 6.5-cm to 7-mm are presented here. Observing at several different wavelengths, separated by short and long intervals, enabled reliable information about the magnetic (B) field structure to be obtained and for the effect of Faraday Rotation to be determined and corrected. For all sources the magnitude of the core Rotation Measure (RM) derived from the shorter wavelength data was greater than that derived from the longer wavelength data, consistent with a higher electron density and/or B-field strength closer to the central engine. A transverse RM gradient was detected in the jet of 0954+658, providing evidence for the presence of a helical B-field surrounding the jet. The RM in the core region of 2200+420 (BL Lac) displays sign changes in different wavelength intervals (on different spatial scales); we suggest an explanation for this in terms of modest bends in a helical B-field surrounding the jet. ",
        "introduction": "The Faraday effect causes a rotation of the plane of linear polarisation, described by: $\\Delta\\chi = RM \\lambda^2$, with the rotation measure (RM) determined by the integral of the electron density and the dot product of the magnetic (B) field and the path length along the line of sight (LoS). A positive/negative RM indicates that the LoS B-field is pointing towards/away from the observer. $\\Delta\\chi$ is the change in the electric vector position angle (EVPA). Previous results indicated the presence of different RM signs in the core regions of 6 blazars in different wavelength intervals (O'Sullivan \\& Gabuzda 2006). This has two main possible origins: (1) the LoS B-field changes with distance from the centre of activity, or (2) since the previous observations were not simultaneous, it could be due to an intrinsic change in the overall jet B-field structure between observing epochs. Our new 8 wavelength observations are designed to test these possibilities. ",
        "conclusions": "Our results for 2200+420 confirm the presence of an RM sign reversal in the core region. Since the dominant jet B-field is transverse to the jet and remains transverse while the jet bends, we will suppose a helical B-field surrounds the jet. The observed RM sign reversal can be explained by a slight bend of the jet, due, for example, to a collision with material in the parent galaxy or some instabilities inherent in the jet itself. (A longitudinal jet B-field with a change in the angle to the line of sight could also cause a RM sign reversal, but this does not correspond to the observed B-field.)  A side-on view of a helical B-field (Fig.1 Top right) will have a RM that will be equally strong on both sides of the jet, hence, a zero net RM will be observed for an unresolved jet. This would occur when the source is viewed at $1/\\Gamma$ in the observer's frame. For a tail-on view of a helical B-field (ie. $\\theta > 1/\\Gamma$) (Fig.1 Middle), the dominant RM will be from the bottom half of the jet and a negative RM will be observed because the dominant LoS B-field will be pointing away from us. (Assuming the jet is not fully resolved in the transverse direction.) Conversely, for a head-on view of a helical B-field (ie. $\\theta < 1/\\Gamma$) (Fig.1 Bottom), a positive RM will be observed.  Therefore, regions with different RM signs in the jets of AGN can be explained within a helical B-field model as places where the jet is observed at angles greater than or less than $~1/\\Gamma$, due to bends in the jet. Since VLBI resolution is usually not sufficient to completely resolve the true optically thick core, the VLBI ``core'' consists of emission from the true core and some of the optically thin inner-jet. So if bends occur on scales smaller than the observed VLBI core, ``core'' RMs with different signs could be derived from observations at different wavelengths (ie. probing different scales of the inner-jet). In our future work, we will attempt to reconstruct the 3D path of the jet through space using the combined information from the observed distributions of the total intensity, linear polarisation, spectral index and rotation measure."
    },
    "0710/0710.5682_arXiv.txt": {
        "abstract": "{  {\\it Context: } One of the most debated issues about sub-mJy radio sources, which are responsible for the steepening of the 1.4 GHz source counts, is the origin of their radio emission. Particularly interesting, from this point of view, is the possibility of combining radio spectral index information with other observational properties to assess whether the sources are triggered by star formation or nuclear activity.  {\\it Aims:} The aim of this work is to study the  optical and near infrared  properties of a complete sample of 131 radio sources with $S>0.4$ mJy, observed at both 1.4 and 5 GHz  as part of the ATESP radio survey. The availability of multi--wavelength radio and optical information is exploited to infer the physical properties of the faint radio population.  {\\it Methods:} We use deep multi--colour (UBVRIJK) images, mostly taken in the framework of the ESO \\emph{Deep Public Survey}, to optically identify and derive photometric redshifts for the ATESP radio sources. Deep optical coverage  and extensive colour information are available for 3/4 of the region covered by the radio sample. Typical depths of the images are $U\\sim 25$, $B\\sim 26$, $V\\sim 25.4$, $R\\sim 25.5$, $I\\sim 24.3$, $19.5\\leq K_{s}\\leq 20.2$, $J\\leq 22.2$. We also add shallower optical imaging and spectroscopy obtained previously in order to perform a global analysis of the radio sample.  {\\it Results:} Optical/near infrared counterparts are found for $\\sim 78\\%$ (66/85) of the radio sources in the region covered by the deep multi--colour imaging, and for 56 of these reliable estimates of the redshift and type are derived. We find that many of the sources with flat radio spectra are characterised by high radio--to--optical ratios ($R>1000$), typical of classical powerful radio galaxies and quasars. Flat--spectrum sources with low $R$ values are preferentially identified with early type galaxies, where the radio emission is most probably triggered by low--luminosity active galactic nuclei. Considering both early type galaxies and quasars as sources with an active nucleus, such sources largely dominate our sample (78\\%). Flat--spectrum sources associated with early type galaxies are quite compact ($d<10-30$ kpc), suggesting core-dominated radio emission. ",
        "introduction": "\\label{sec:introduction}  The faint (sub-mJy) radio population consists of a mixture of different classes of objects. Since the early seventies it  has been known that the strongest sources are almost exclusively associated with either active galactic nuclei (AGNs) or giant ellipticals, the latter of which are also known as radio galaxies (99\\% above 60 mJy, \\citealt{Windhorst90}). More recent work on mJy and sub-mJy sources has revealed that faint sources are also found to be associated with normal elliptical, spiral and star-forming galaxies, with the early type galaxies being the dominant component \\citep{Gruppioni1999, Georgakakis1999, Magliocchetti2000, Prandoni2001b, Afonso2006}, while at $\\mu$Jy levels star-forming galaxies prevail (see e.g. \\citealt{Richards1999}).  In spite of the progress made in our understanding of the faint radio population, many questions remain open. For example, the relative fractions of the different types of objects are still quite uncertain, and our knowledge of their dependence on limiting flux density is still incomplete. The reason is, of course, that very little is known about the faint ends of the various luminosity functions, and even less is known about the cosmological evolution of different kinds of objects. This uncertainty is due to the incompleteness of optical identification and spectroscopy, since faint radio sources usually have very faint optical counterparts. Clearly {\\it very}\\, deep ($I\\apprge 25$) optical imaging and spectroscopy, for reasonably large deep radio samples, are critical if one wants to investigate these radio source populations.  Since the radio emission comes from different types of objects an important question is what are the physical processes that trigger this emission. It is natural to assume that in the case of star-forming galaxies the emission traces the history of galaxy formation and subsequent evolution by merging and interaction, while the emission in  AGNs will reflect black hole accretion history. To make matters more complicated, both processes may be present at the same time.   Although research in this field proceeds slowly due to very time--consuming spectroscopy much progress has been made in recent years thanks to strong improvement in the photometric redshift technique. Several multi--colour/multi--object spectroscopy surveys overlapping deep radio fields have recently been undertaken, including the Phoenix Deep Survey \\citep{Hopkins1998,Georgakakis1999,Afonso2006} and the Australia Telescope ESO Slice Project (ATESP) survey \\citep{Prandoni2000a,Prandoni2000b,Prandoni2001b}. In other cases, deep multi--colour/multi--wavelength surveys have been complemented by deep radio observations (see e.g. the VLA--VIRMOS, \\citealt{Bondi2003}; and the COSMOS, \\citealt{Schinnerer2006}).   Multi--frequency radio observations are also important in measuring the radio spectral index, which may help to constrain the origin of the radio emission in the faint radio sources. This approach is especially meaningful when high resolution radio images are available and radio source structures can be inferred. However, multi--frequency radio information is available for very few, and small, sub-mJy radio samples.  The largest sample with multi--frequency radio coverage available so far is a complete sample of 131 radio sources with $S>0.4$ mJy, extracted from a square degree region observed at both 1.4 and 5 GHz as part of the ATESP radio survey \\citep{Prandoni2000a,Prandoni2000b,Prandoni2006}.  The $1.4-5$~GHz radio spectral index analysis of the ATESP radio sources was presented in the first paper of this series (\\citealt{Prandoni2006}, hereafter Paper I). We found a flattening of the radio spectra with decreasing radio flux density. At mJy levels most sources have steep spectra ($\\alpha \\sim -0.7$, assuming $S\\sim \\nu^{\\alpha}$), typical of synchrotron radiation, while at sub-mJy flux densities a composite population is present, with up to $\\sim 60\\%$ of the sources showing flat ($\\alpha > -0.5$) spectra and a significant fraction ($\\sim 30\\%$) of inverted-spectrum ($\\alpha>0$) sources. This flattening at sub-mJy fluxes confirms previous results based on smaller samples (\\citealt{Donnelly1987,Gruppioni1997,Ciliegi2003}). Flat spectra in radio sources usually indicate the presence of a self-absorbed nuclear core, but they can also be produced on larger scales by thermal emission from stars.  It is possible to combine the spectral index information with other observational properties and infer the nature of the faint radio population. This is especially important with respect to the class of flat/inverted--spectrum sources as it permits us to study the physical processes that trigger the radio emission in those sources. This kind of analysis needs information about the redshifts and types of the galaxies hosting the radio sources. A detailed radio/optical study of the sample above is possible, thanks to the extensive optical/infrared coverage mostly obtained in the ESO \\emph{Deep Public Survey} (DPS, \\citealt{Mignano2007,Olsen2006}).   We give a brief discussion of all the data collected so far in Sect.~\\ref{sec:datacoverage}, followed by a more detailed analysis of the DPS optical data in Sect.~\\ref{sec:dpsanalysis}, where we derive the UBVRI colour catalogue and photometric redshifts for the DPS galaxies in the region covered by the ATESP survey, assessing the reliability of the photometric redshifts themselves. In Sects.~\\ref{sec:optid} and \\ref{sec:radiozphot}, respectively, we use the DPS UBVRIJK optical data to identify the ATESP radio sources and to derive photometric redshifts.  A radio/optical analysis of the optically identified radio sources is presented in Sect.~\\ref{sec:comp}, while in Sect.~\\ref{sec:nature} we discuss the nature of the mJy and sub--mJy population on the basis of all the radio and optical data available to the ATESP sample. The main results are briefly summarised in Sect.~\\ref{sec:summary}.\\\\  Throughout this paper we use the $\\Lambda$CDM model, with $H_0=70$, $\\Omega_m=0.3$ and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "\\label{sec:summary}  In this paper we have discussed the nature of the faint, sub-mJy, radio population, using a sample of 131 radio sources that were observed at 1.4 and 5 GHz with the ATCA (the ATESP--DEEP1 sample). A smaller sample of 85 radio sources is covered by deep multi--colour images. These were optically identified down to very faint magnitudes, which was possible thanks to the availability of very deep multi--colour optical material (in U, B, V, R, I, and sometimes J and K bands). The high percentage of identifications ($\\sim 78\\%$) makes this a sample that is well suited for follow up studies concerning the composition of the sub-mJy population and, in general, the cosmological evolution of the various classes of objects associated with faint radio sources.  We summarise our main results here. \\begin{itemize} \\item For  85\\% of the identification sample we succeeded in deriving reliable photometric redshifts, based on the available accurate colours (UBVRIJK). \\item Based on spectral types determined either directly from spectroscopy or from the photometry (or both), we find that at the sub-mJy level the large majority of sources are associated with objects that have early type (64\\%) and  AGN (14\\%) spectra; these are of course what we would normally call radio galaxies and quasars. \\item Although earlier work (based on shallower optical imaging and spectroscopy) revealed the presence of a conspicuous component of late type and star-burst objects, such objects appear to be important only at brighter magnitudes ($I<19$), and are rare at fainter magnitudes ($19<I<23.5$). \\item From an overall comparison of the radio spectral index with other radio and optical properties of the entire ATESP--DEEP1 sample, we find that most sources with flat radio spectra have high radio-to-optical ratios, as expected for classical radio galaxies and quasars. Flat-spectrum sources with low radio-to-optical ratios are preferentially associated with ETS, in which the radio emission is most plausibly triggered by nuclear activity as well, while star-forming galaxies are associated to steep-spectrum radio sources. \\item ETS with flat or inverted spectra are mostly compact, with linear size $<10-30$ kpc, suggesting core-dominated radio emission. Their low radio luminosities (in the range $10^{22}-10^{24}$ W/Hz at 1.4 GHz) and the absence of emission lines in their spectra (when available) suggest that they are FRI sources, although these would normally have steeper spectra and be more extended. They may therefore represent specific phases in the life of a radio source, or may be similar to the low power compact radio sources discussed by \\cite{Giroletti2005}. \\end{itemize}"
    },
    "0710/0710.2541_arXiv.txt": {
        "abstract": "Preliminary results are presented from a simple, single-antenna experiment designed to measure the all-sky radio spectrum between 100 and 200~MHz. The system used an internal comparison-switching scheme to reduce non-smooth instrumental contaminants in the measured spectrum to 75~mK. From the observations, we place an initial upper limit of $450$~mK on the relative brightness temperature of the redshifted 21~cm contribution to the spectrum due to neutral hydrogen in the intergalactic medium (IGM) during the epoch of reionization, assuming a rapid transition to a fully ionized IGM at a redshift of 8. With refinement, this technique should be able to distinguish between slow and fast reionization scenarios.  To constrain the duration of reionization to $\\Delta z>2$, the systematic residuals in the measured spectrum must be reduced to 3~mK. ",
        "introduction": "The transition period at the end of the cosmic ``Dark Ages'' is known as the epoch of reionization (EOR). During this epoch, radiation from the very first luminous sources---early stars, galaxies, and quasars---succeeded in ionizing the neutral hydrogen gas that had filled the intergalactic medium (IGM) since the recombination event following the Big Bang. Reionization marks a significant shift in the evolution of the Universe. For the first time, gravitationally-collapsed objects exerted substantial feedback on their environments through electromagnetic radiation, initiating processes that have dominated the evolution of the visible baryonic Universe ever since. The epoch of reionization, therefore, can be considered a dividing line when the relatively simple evolution of the early Universe gave way to more complicated and more interconnected processes. Although the Dark Ages are known to end when the first luminous sources ionized the neutral hydrogen in the IGM, precisely when this transition occurred remains uncertain.  The best existing constraints on the timing of the reionization epoch come from two sources: the cosmic microwave background (CMB) anisotropy and absorption features in the spectra of high-redshift quasars. The amplitude of the observed temperature anisotropy in the CMB is affected by Thomson scattering due to electrons along the line of sight between the surface of last scattering and the detector, and thus, it is sensitive to the ionization history of the IGM through the electron column density. In addition, if there is sufficient optical depth to CMB photons due to free electrons in the IGM after reionization, some of the angular anisotropy in the unpolarized intensity can be converted to polarized anisotropy. This produces a peak in the polarization power spectrum at the angular scale size equivalent to the horizon at reionization with an amplitude proportional to the optical depth \\citep{1997ApJ...488....1Z}. Measurements by the WMAP satellite of these effects indicate that the redshift of reionization is $z_r\\approx11\\pm4$ \\citep{2007ApJS..170..377S}, assuming an instantaneous transition.  Lyman-$\\alpha$ absorption by neutral hydrogen is visible in the spectra of many high-redshift quasars and, thus, offers the second currently feasible probe of the ionization history of the IGM. Continuum emission from quasars is redshifted as it travels through the expanding Universe to the observer.  Neutral hydrogen along the line of sight creates absorption features in the continuum at wavelengths corresponding to the local rest-frame wavelength of the Lyman-$\\alpha$ line.  Whereas CMB measurements place an integrated constraint on reionization, quasar absorption line studies are capable of probing the ionization history in detail along the sight-lines. There is a significant limitation to this approach, however. The Lyman-$\\alpha$ absorption saturates at very low fractions of neutral hydrogen (of order $x_{HI} \\approx 10^{-4}$). Nevertheless, results from these studies have been quite successful and show that, while the IGM is highly ionized below $z\\lesssim6$ (with typical $x_{HI}\\lesssim10^{-5}$), a significant amount of neutral hydrogen is present above, although precisely how much remains unclear \\citep{2001ApJ...560L...5D, 2001AJ....122.2850B, 2002AJ....123.1247F, 2003AJ....125.1649F, 2004Natur.427..815W, 2006AJ....132..117F}.  The existing CMB and quasar absorption measurements are somewhat contradictory.  Prior to these studies, the reionization epoch was assumed generally to be quite brief, with the transition from an IGM filled with fully neutral hydrogen to an IGM filled with highly ionized hydrogen occurring very rapidly.  These results, however, open the possibility that the ionization history of the IGM may be more complicated than previously believed \\citep{2003ApJ...595....1H, 2003ApJ...591...12C, 2003MNRAS.344..607S, 2004ApJ...604..484M}.  Direct observations of the 21~cm (1420~MHz) hyperfine transition line of neutral hydrogen in the IGM during the reionization epoch would resolve the existing uncertainties and reveal the evolving properties of the IGM.  The redshifted 21~cm signal should appear as a faint, diffuse background in radio frequencies below $\\nu<200$~MHz for redshifts above $z>6$ (according to $\\nu=1420/[1+z]$~MHz). For diffuse gas in the high-redshift ($z\\approx10$) IGM, the expected unpolarized differential brightness temperature of the redshifted 21~cm line relative to the pervasive CMB is readily calculable from basic principles and is given by \\citep[their \\S~2]{2004ApJ...608..622Z} \\begin{equation} \\begin{array}{rl} \\label{eqn_intro_temp} \\delta T_{21}(\\vec{\\theta}, z) \\approx~& 23~(1+\\delta)~x_{HI} \\left ( 1 - \\frac{T_\\gamma}{T_S} \\right ) \\\\ & \\times \\left ( \\frac{\\Omega_b~h^2}{0.02} \\right ) \\left [ \\left ( \\frac{0.15}{\\Omega_m~h^2} \\right ) \\left ( \\frac{1+z}{10} \\right ) \\right ]^{1/2} \\mbox{mK}, \\end{array} \\end{equation} where $\\delta(\\vec{\\theta},z)$ is the local matter over-density, $x_{HI}(\\vec{\\theta},z)$ is the neutral fraction of hydrogen in the IGM, $T_\\gamma(z) = 2.73~(1+z)$~K is the temperature of CMB at the redshift of interest, $T_S(\\vec{\\theta},z)$ is the spin temperature that describes the relative population of the ground and excited states of the hyperfine transition, and $\\Omega_b$ is the baryon density relative to the critical density, $\\Omega_m$ is the total matter density, and $h$ specifies the Hubble constant according to $H_0=100~ h$~km~s$^{-1}$~Mpc$^{-1}$. From Equation~\\ref{eqn_intro_temp}, we see that perturbations in the local density, spin temperature, and neutral fraction of hydrogen in the IGM would all be revealed as fluctuations in the brightness temperature of the observed redshifted 21~cm line.  The differential brightness temperature of the redshifted 21~cm line is very sensitive to the \\hi spin temperature. When the spin temperature is greater than the CMB temperature, the line is visible in emission. For $T_S \\gg T_\\gamma$, the magnitude of the emission saturates to a maximum (redshift-dependent) brightness temperature that is about 25 to 35~mK for a mean-density, fully neutral IGM between redshifts 6 and 15, assuming a $\\Lambda$CDM cosmology with $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, $\\Omega_b=0.04$, and $h=0.7$. At the other extreme, when the spin temperature is very small and $T_S \\ll T_\\gamma$, the line is visible in absorption against the CMB with a potentially very large (and negative) relative brightness temperature.  A number of factors are involved in predicting the typical differential brightness temperature of the redshifted 21~cm line as a function of redshift.  In particular, the spin temperature must be treated in detail, including collisional coupling between the spin and kinetic temperatures of the gas, absorption of CMB photons, and heating by ultra-violet radiation from the first luminous sources. We direct the reader to \\citet{2006PhR...433..181F} for a good introduction to the topic. The results of several efforts to predict the evolution of the differential brightness temperature of the redshifted 21~cm line have yielded predictions that are generally consistent in overall behavior, but vary highly in specific details \\citep{1997ApJ...475..429M, 1999A&A...345..380S, 2004ApJ...608..611G, 2006MNRAS.371..867F}. These models tend to agree that, for a finite period at sufficiently high redshifts ($z\\gtrsim20$), the \\hi hyperfine line should be seen in absorption against the CMB, with relative brightness temperatures of up to $|\\delta T_b|\\lesssim100$~mK.  This is because the IGM initially cools more rapidly than the CMB following recombination \\citep{1994ApJ...427...25S, 1997ApJ...475..429M}. During this period, fluctuations in the differential brightness temperature of the redshifted 21~cm background should track the underlying baryonic matter density perturbations \\citep{1972A&A....20..189S, 1979MNRAS.188..791H, 1990MNRAS.247..510S, 2002ApJ...572L.123I, 2003MNRAS.341...81I, 2004PhRvL..92u1301L, 2005ApJ...626....1B}. Eventually, however, the models indicate that the radiation from the first generations of luminous sources will elevate the spin temperature of neutral hydrogen in the IGM above the CMB temperature and the redshifted 21~cm line should be detected in emission with relative brightness temperatures up to the expected maximum values (of order $25$~mK). Finally, during the reionization epoch, the neutral hydrogen becomes ionized, leaving little or no gas to produce the \\hi emission, and the apparent differential brightness temperature of the redshifted 21~cm line falls to zero as reionization progresses. As the gas is ionized, a unique pattern should be imprinted in the redshifted 21~cm signal that reflects the processes responsible for the ionizing photons and that evolves with redshift as reionization progresses \\citep{1997ApJ...475..429M, 2000ApJ...528..597T, 2003ApJ...596....1C, 2004ApJ...608..622Z, 2004ApJ...613...16F}. The details of the specific timing, duration, and magnitude of these features remains highly variable between theoretical models due largely to uncertainties about the properties of the first luminous sources.  Measuring the brightness temperature of the redshifted 21~cm background could yield information about both the global and the local properties of the IGM. Determining the average brightness temperature over a large solid angle as a function of redshift would eliminate any dependence on local density and temperature perturbations and constrain the evolution of the product $\\overline{x_{HI}(1-T_\\gamma/T_S)}$, where we use the bar to denote a spatial average. During the reionization epoch, it is, in general, believed to be a good approximation to assume that $T_S\\gg T_\\gamma$ and, therefore, that the brightness temperature is proportional directly to $\\bar{x}_{HI}$. Global constraints on the brightness temperature of the redshifted 21~cm line during the EOR, therefore, would directly constrain the neutral fraction of hydrogen in the IGM. Such constraints would provide a basic foundation for understanding the astrophysics of reionization by setting bounds on the duration of the epoch, as well as identifying unique features in the ionization history (for example if reionization occurred in two phases or all at once).  They would also yield improvements in estimates of the optical depth to CMB photons and, thus, would help to break existing degeneracies in CMB measurements between the optical depth and properties of the primordial matter density power spectrum \\citep{2006PhRvD..74l3507T}. \\begin{figure} \\centering \\includegraphics[width=20pc]{f1_color.eps} \\caption[Photograph of EDGES deployed at Mileura Station]{ \\label{f_edges_photos} EDGES deployed at Mileura Station in Western Australia. The left panel shows the full antenna and ground screen in the foreground and the analog-to-digital conversion and data acquisition module in the background.  The right panel is a close-up view of the amplifier and switching module connected directly to the antenna (through the balun).} \\end{figure} For these reasons, several efforts are underway to make precise measurements of the radio spectrum below $\\nu<200$~MHz ($z>6$).  In this paper, we report on the initial results of the Experiment to Detect the Global EOR Signature (EDGES). In \\S~\\ref{s_edges_method}, we describe the specific approach used for EDGES to address the issue of separating the redshifted 21 signal from the foreground emission. We then give an overview of the EDGES system in \\S~\\ref{s_edges_system}, followed by the results of the first observing campaign with the system in \\S~\\ref{s_edges_results}, along with a discussion of the implications for future single-antenna measurements. ",
        "conclusions": "In principle, useful measurements of the redshifted 21~cm background can be carried out with a small radio telescope. These measurements would be fundamental to understanding the evolution of the IGM and the EOR. In particular, the global evolution of the mean spin temperature and mean ionization fraction of neutral hydrogen in the high redshift IGM could be constrained by very compact instruments employing individual radio antennas.  We have reported preliminary results to probe the reionization epoch based on this approach from the first observing campaign with the EDGES system. These observations were limited by systematic effects that were an order of magnitude larger than the anticipated signal and, thus, ruled out only an already unlikely range of parameter space for the differential amplitude of the redshifted 21~cm brightness temperature and for the duration of reionization. Nevertheless, the results of this experiment indicate the viability of the simple global spectrum approach.  Building on the experiences of these initial efforts, modifications to the EDGES system are underway to reduce the residual systematic contribution in the measured spectrum and to expand the frequency coverage of the system down to 50~MHz or lower in order to place constraints on the anticipated transition of the hyperfine line from absorption to emission as the IGM warms before the EOR. Constraining the redshift and intensity of this feature would be very valuable for understanding the heating history of the IGM and, since the transition has the potential to produce a step-like feature in the redshifted 21~cm spectrum with a magnitude over 100~mK (up to a factor of 4 larger than the amplitude of the step during the reionization epoch), it may be easier to identify than the transition from reionization---although the sky noise temperature due to the Galactic synchrotron foreground increases significantly at the lower frequencies, as well. Through these and other global spectrum efforts, the first contribution to cosmic reionization science from measurements of the redshifted 21~cm background will hopefully be achieved in the near future."
    },
    "0710/0710.1175_arXiv.txt": {
        "abstract": "{The physical mechanism responsible for the short outbursts in a recently recognized class of High Mass X--ray Binaries, the Supergiant Fast X--ray Transients (SFXTs), is still unknown. Two main hypotheses have been proposed to date: the sudden accretion by the compact object of small ejections originating in  a clumpy wind from the supergiant donor, or outbursts produced at (or near) the periastron passage in wide and eccentric orbits, in order to explain the low ($\\sim$10$^{32}$~erg~s$^{-1}$)  quiescent emission. Neither  proposed mechanisms seem to explain the whole phenomenology of these sources. } {Here we propose a new explanation for the outburst mechanism, based on the X--ray observations of the unique SFXT known to display periodic outbursts, IGR~J11215-5952. } {We performed three Target of Opportunity observations with \\sw, \\xmm\\ and \\inte\\ at the time of the fifth  outburst, expected on 2007 February 9. \\sw\\ observations of the February 2007 outburst have been reported elsewhere. Another ToO with \\sw\\ was performed in July 2007, in order to monitor the supposed ``apastron'' passage. } {\\xmm\\ observed the source on 2007 February 9, for 23~ks, at the peak of the outburst, while \\inte\\ started the observation two days later, failing to detect the source, which had already undergone the decaying phase of the fast outburst. The Swift campaign performed in July 2007 reveals a second outburst occurring on 2007 July 24, as bright as that observed about 165~days before. } {The new X--ray observations allow us to propose an alternative hypothesis for the outburst mechanism in SFXTs, linked to the possible presence of a second wind component, in the form of an equatorial disk from the supergiant donor. We discuss the applicability of the model to the short outburst durations of all other Supergiant Fast X--ray Transients, where a clear periodicity in the outbursts has not been found yet. The new outburst from \\src\\ observed in July suggests that the true orbital period is  $\\sim$165~days, instead of 329~days, as previously thought. } ",
        "introduction": "The Galactic plane monitoring performed by the \\inte\\ satellite in the last 5 years has allowed the discovery of a number of new High Mass X--ray Binaries (HMXBs). Several of these new sources are intrinsically highly absorbed and were difficult to discover with previous missions (e.g. IGR~J16318--4848, \\citealt{Walter2003}). Others are transient HMXBs (associated with OB supergiant) displaying short outbursts (few hours, typically less than a day; \\citealt{Sguera2005}), and form the recently recognized new class of  Supergiant Fast X--ray Transients (SFXTs).  \\object{IGR~J11215--5952} is a hard X--ray transient  discovered by \\inte\\  during a fast outburst in  April 2005 \\citep{Lubinski2005}. The short duration of this outburst led \\citet{Negueruela2005a} to propose that \\src\\ could be a new member of the class of Supergiant Fast X-ray Transients. The optical counterpart is indeed a B-type supergiant, \\object{HD~306414} located at a distance of 6.2~kpc (\\citealt{Negueruela2005b}, \\citealt{Masetti2006}, \\citealt{Steeghs2006}).  From the analysis of archival \\inte\\ observations and the discovery of two previously unnoticed outbursts, a recurrence period in the X--ray activity of $\\sim$330~days has been found \\citep*[hereafter Paper I]{SidoliPM2006}, likely linked to the orbital period of the binary system. This periodicity was later confirmed by the fourth outburst from \\src\\ observed with $RXTE/PCA$  on 2006 March 16--17, 329~days after the previous one \\citep{Smith2006a}. The \\inte\\ spectrum was well fitted by a hard power-law with a high energy cut-off around 15~keV (Paper~I). Assuming a  distance of 6.2 kpc, the peak fluxes of the outbursts correspond to a luminosity of $\\sim 3 \\times$10$^{36}$~erg~s$^{-1}$ (5--100~keV; Paper~I). The $RXTE/PCA$ observations showed a possible pulse period of  $\\sim$195\\,s \\citep{Smith2006b}, later confirmed during the February 2007 outburst, yielding P=$186.78\\pm0.3$\\,s \\citep{Swank2007}.  All these findings confirmed \\src\\ as a member of the class of SFXTs, and  the first one displaying periodic outbursts. Based on the known periodicity, an outburst was expected for 2007 February 9. This  allowed us to obtain several Target of Opportunity (ToO) observations  with $Swift/XRT$,  $XMM-Newton$ and $INTEGRAL$. The $Swift/XRT$ results of the February 2007 outburst have been reported in \\citealt{Romano2007} (hereafter Paper~II). Here we report the results of the \\xmm\\ and \\inte\\ observations of the February 2007 outburst, and of a $Swift/XRT$ campaing performed in July 2007, in order to monitor the supposed apastron passage. ",
        "conclusions": "\\label{conclusions}     We have  proposed here a new explanation for the short outbursts in SFXTs, based on our results of a monitoring campain of \\src. An equatorial wind from the  supergiant companion is suggested, based on the narrow and steep shape of the X--ray lightcurve observed during the latest outburst. The short outburst is suggestive of a  deviation from coplanarity ($\\theta$$>$0) of the equatorial plane of the companion with the orbital plane (as, e.g., in PSR~B1259-63, \\citealt{Wex1998} or in PSR~J0045--7319, \\citealt{Kaspi1994}), and of a some degree of eccentricity (e$>$0).  Both orbital eccentricity and no-coplanarity can be explained by a substantial supernova ``kick'' (e.g. \\citealt{Eggleton2001}) at birth. This could suggest that SFXTs are likely young systems, probably younger than persistent HMXBs.  We are aware that the derived quantities for the proposed supergiant equatorial disk are highly uncertain and still speculative. This uncertainty clearly emerges from the fact that the observed X--ray lightcurve can be reproduced with both orbital periods (329~days and 164.5~days), different eccentricities  and different wind properties. A determination of the orbital eccentricity of the system, for example, would be  essential in better constraining the expected properties of the wind which match with the level of X--ray emission during the outbursts. These considerations highlight the need for an as complete as possible deep monitoring of the X--ray emission along the orbit, together with optical observations in order to constrain the supergiant wind properties.    \\begin{figure*} \\vbox{ \\includegraphics[height=6.5cm,angle=0]{8137fig10a.ps} \\includegraphics[height=6.5cm,angle=0]{8137fig10b.ps} \\includegraphics[height=6.5cm,angle=0]{8137fig10c.ps} \\includegraphics[height=6.5cm,angle=0]{8137fig10d.ps} \\includegraphics[height=6.5cm,angle=0]{8137fig10e.ps} \\includegraphics[height=6.5cm,angle=0]{8137fig10f.ps}} \\caption[]{Results of the proposed model compared with the $Swift/XRT$ lightcurve ({\\em upper panels}) the expected luminosity variations ({\\em middle panels}) and wind density  along the whole  neutron star orbit ({\\em lower panels}) in two cases: {\\em on the left} geometry case ``a'' (with orbital period of 164.5~days and an assumed eccentricity of 0.4) is shown (see Fig.~\\ref{fig:geom}), while {\\em on the right} the results for case ``b'' are reported (orbital period of 329~days, circular orbit and two similar outbursts per orbit). Both models assume a blue supergiant with a mass of 39~M$_\\odot$ and radius of 42~R$_\\odot$, a  ``polar wind'' component \\textbf{(``PW'')} with a terminal velocity of 1800~km~s$^{-1}$. The X--ray lightcurve observed with $Swift/XRT$ is better reproduced assuming a ``polar wind'' mass loss rate of 5$\\times$10$^{-6}$~M$_\\odot$~yrs$^{-1}$ for case ``a'' and 9$\\times$10$^{-7}$ ~M$_\\odot$~yrs$^{-1}$ for case ``b''. The second wind component, in form of an equatorial disk (``ED''), has a variable velocity ranging from 750~km~s$^{-1}$ to 1400~km~s$^{-1}$ (for case ``a''), and from 850~km~s$^{-1}$ to 1600~km~s$^{-1}$ for case ``b''. The wind density profiles assumed here are reported in the two lower panels. For a magnetic field of 10$^{12}$~G the centrifugal barrier is open along all the orbit in both cases. Dashed line in the first upper right figure shows the model prediction based on a ten times smaller wind density with respect to that assumed to reproduce the peak luminosity (maintaining fixed all the other parameters). } \\label{fig:model} \\end{figure*}   \\onltab{2}{        % \\begin{table*} \\begin{center} \\caption{Observation log.} \\label{igr112:tab:alldata2} \\begin{tabular}{lllll} \\hline \\hline \\noalign{\\smallskip} Sequence         & Start time (MJD)     & Start time  (UT)             & End time   (UT)                 & Net Exposure$^{\\mathrm{a}}$    \\\\ &                      & (yyyy-mm-dd hh:mm:ss)           & (yyyy-mm-dd hh:mm:ss)        &(s)       \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 00030881024     & 54256.7449\t  &\t  2007-06-05 17:52:35\t  &\t  2007-06-05 21:15:58\t  &\t  1519    \\\\ 00030881025     & 54263.4918\t  &\t  2007-06-12 11:48:11\t  &\t  2007-06-12 16:57:58\t  &\t  1522    \\\\ 00030881026     & 54270.3808\t  &\t  2007-06-19 09:08:17\t  &\t  2007-06-19 11:05:57\t  &\t  1843    \\\\ 00030881027     & 54277.1485\t  &\t  2007-06-26 03:33:52\t  &\t  2007-06-26 05:27:57\t  &\t  2109    \\\\ 00030881028     & 54284.8571\t  &\t  2007-07-03 20:34:14\t  &\t  2007-07-03 23:59:57\t  &\t  2347    \\\\ 00030969001     & 54294.5627\t  &\t  2007-07-13 13:30:15\t  &\t  2007-07-14 07:10:56\t  &\t  2126    \\\\ 00030881030     & 54298.0425\t  &\t  2007-07-17 01:01:12\t  &\t  2007-07-17 02:58:57\t  &\t  2069    \\\\ 00030881031     & 54305.0012   \t  &       2007-07-24 00:01:42     &       2007-07-24 13:16:55     &       1551    \\\\ 00030881032     & 54312.2343      &       2007-07-31 05:37:23     &       2007-07-31 10:26:56     &       1695    \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{center} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] The exposure time is spread over several snapshots (single continuous pointings at the target) during each observation. \\end{list} \\end{table*} } %"
    },
    "0710/0710.2189.txt": {
        "abstract": "{Minimal walking technicolor models can provide a nontrivial solution for cosmological dark matter, if the lightest technibaryon is doubly charged. Technibaryon asymmetry generated in the early Universe is related to baryon asymmetry and it is possible to create excess of techniparticles with charge ($-2$). These excessive techniparticles are all captured by $^4He$, creating \\emph{techni-O-helium} $tOHe$ ``atoms'', as soon as $^4He$ is formed in Big Bang Nucleosynthesis. The interaction of techni-O-helium with nuclei opens new paths to the creation of heavy nuclei in Big Bang Nucleosynthesis. Due to the large mass of technibaryons, the $tOHe$ ``atomic'' gas decouples from the baryonic matter and plays the role of dark matter in large scale structure formation, while structures in small scales are suppressed. %Due to Nuclear interactions with matter slow down cosmic techni-O-helium in Earth below the threshold of underground dark matter detectors, thus escaping severe CDMS constraints. On the other hand, these nuclear interactions %of techni-O-helium are not sufficiently strong to exclude this form of Strongly Interactive Massive Particles by constraints from the XQC experiment. Experimental tests of this hypothesis are possible in search for $tOHe$ in balloon-borne experiments (or on the ground) and for its charged techniparticle constituents in cosmic rays and accelerators. The $tOHe$ ``atoms'' can cause cold nuclear transformations in matter and might form anomalous isotopes, offering possible ways to exclude (or prove?) their existence.} ",
        "introduction": "The question of the existence of new quarks and leptons is among the most important in the modern high energy physics. This question has an interesting cosmological aspect. If these quarks and/or charged leptons are stable, they should be present around us and the reason for their evanescent nature should be found.  Recently, at least three elementary particle frames for heavy stable charged quarks and leptons were considered: (a) A heavy quark and heavy neutral lepton (neutrino with mass above half the $Z$-boson mass) of a fourth generation \\cite{N,Legonkov}, which can avoid experimental constraints \\cite{Q,Okun}, and form composite dark matter species \\cite{I,lom,KPS06,Khlopov:2006dk}; (b) A Glashow's ``Sinister'' heavy tera-quark $U$ and tera-electron $E$, which can form a tower of tera-hadronic and tera-atomic bound states with ``tera-helium atoms'' $(UUUEE)$ considered as dominant dark matter \\cite{Glashow,Fargion:2005xz}; (c) AC-leptons, based on the approach of almost-commutative geometry \\cite{5,book}, that can form evanescent AC-atoms, playing the role of dark matter \\cite{5,FKS,Khlopov:2006uv}.  In all these recent models, the predicted stable charged particles escape experimental discovery, because they are hidden in elusive atoms, composing the dark matter of the modern Universe. It offers a new solution for the physical nature of the cosmological dark matter. Here we show that such a solution is possible in the framework of walking technicolor models \\cite{Sannino:2004qp,Hong:2004td,Dietrich:2005jn,Dietrich:2005wk,Gudnason:2006ug,Gudnason:2006yj} and can be realized without an {\\it ad hoc} assumption on charged particle excess, made in the approaches (a)-(c).    This approach differs from the idea of dark matter composed of primordial bound systems of superheavy charged particles and antiparticles, proposed earlier to explain the origin of Ultra High Energy Cosmic Rays (UHECR) \\cite{UHECR}. To survive to the present time and to be simultaneously the source of UHECR, superheavy particles should satisfy a set of constraints, which in particular exclude the possibility that they possess gauge charges of the standard model.  The particles considered here, participate in the Standard Model interactions and we show how the problems, related to various dark matter scenarios with composite atom-like systems, can find an elegant solution on the base of the minimal walking technicolor model.  The approaches (b) and (c) try to escape the problems of free charged dark matter particles \\cite{Dimopoulos:1989hk} by hiding opposite-charged particles in atom-like bound systems, which interact weakly with baryonic matter. However, in the case of charge symmetry, when primordial abundances of particles and antiparticles are equal, annihilation in the early Universe suppresses their concentration. If this primordial abundance still permits these particles and antiparticles to be the dominant dark matter, the explosive nature of such dark matter is ruled out by constraints on the products of annihilation in the modern Universe \\cite{Q,FKS}. Even in the case of charge asymmetry with primordial particle excess, when there is no annihilation in the modern Universe, binding of positive and negative charge particles is never complete and positively charged heavy species should retain. Recombining with ordinary electrons, these heavy positive species give rise to cosmological abundance of anomalous isotopes, exceeding experimental upper limits. To satisfy these upper limits, the anomalous isotope abundance on Earth should be reduced, and the mechanisms for such a reduction are accompanied by effects of energy release which are strongly constrained, in particular, by the data from large volume detectors.  These problems of composite dark matter models \\cite{Glashow,5} revealed in \\cite{Q,Fargion:2005xz,FKS,I}, can be avoided, if the excess of only $-2$ charge $A^{--}$ %($(\\bar U \\bar U \\bar u)$, $(\\bar U \\bar U \\bar U)$) % or neutral %$(\\bar U u)$ particles is generated in the early Universe. Here we show that in walking technicolor models, technilepton and technibaryon excess is related to baryon excess and the excess of $-2$ charged particles can appear naturally for a reasonable choice of model parameters. It distinguishes this case from other composite dark matter models, since in all the previous realizations, starting from \\cite{Glashow}, such an excess was put by hand to saturate the observed cold dark matter (CDM) density by composite dark matter.  After it is formed in Big Bang Nucleosynthesis, $^4He$ screens the $A^{--}$ charged particles in composite $(^4He^{++}A^{--})$ ``atoms''. These neutral primordial nuclear interacting objects saturate the modern dark matter density and play the role of a nontrivial form of strongly interacting dark matter \\cite{Starkman,McGuire:2001qj}. The active influence of this type of dark matter on nuclear transformations seems to be incompatible with the expected dark matter properties. However, it turns out that the considered scenario is not easily ruled out \\cite{FKS,I} and challenges the experimental search for techni-O-helium and its charged techniparticle constituents.  The structure of the present paper is as follows. Starting with a review of possible dark matter candidates offered by the minimal walking technicolor model, we reveal the possibility for the lightest techniparticle(s) to have electric charge $\\pm 2$ (Section II). % In the framework of the considered approach, the %excess of techniparticles relative to their antiparticles can be %generated in early Universe and related to baryon asymmetry. In Section III we show how the minimal technicolor model can provide substantial excess of techniparticles with electric charge $-2$. %can be generated in the form of antitechnibaryons %$(\\bar{U}\\bar{U})^{-2}$, of technileptons $(\\zeta^{-2}$ or of %their mixture. In Section IV we show how all these $-2$ charge particles can be captured by $^4He$, after its formation in the Standard Big Bang Nucleosynthesis (SBBN), % and neutral %$^4He^{+2}(\\bar{U}\\bar{U})^{-2}$, $^4He^{+2}\\zeta^{-2}$ making neutral techni-O-helium ``atoms\" that can account for the modern dark matter density. Techni-O-helium catalyzes a path for heavy element formation in SBBN, but we stipulate in Section IV a set of arguments, by which the considered scenario can avoid immediate contradiction with observations. Gas of heavy techni-O-helium ``atoms\" decouples from the plasma and radiation only at a temperature about few hundreds eV, so that small scale density fluctuations are suppressed and gravitational instability in this gas develops more close to warm dark matter, rather than to cold dark matter scenario (subsection A of Section V). We further discuss in Section V the possibility to detect charged techniparticle components of cosmic rays (subsection B), effects of techni-O-helium catalyzed processes in Earth (subsection C), and possibilities of direct searches for techni-O-helium (subsection D). The problems, signatures, and possible experimental tests of the techni-O-helium Universe are considered in Section VI. Details of our calculations are presented in the Appendices 1 and 2. % In spite of technilepton $(\\zeta^{-2}$ excess, in the %difference with the case of technibaryons, the frozen out %abundance of antiparticles $(\\bar \\zeta^{+2}$ is not negligible. %Since $(\\bar \\zeta^{+2}$ represents a form of anomalous helium, it %might cause the problem of anomalous helium overproduction. %Cosmological evolution of technileptons and mechanisms of %suppression for primordial $(\\bar \\zeta^{+2}$ abundance are %considered in Appendix 2. ",
        "conclusions": "Discussion} In this paper we explored the cosmological implications of a walking technicolor model with doubly charged technibaryons and technileptons. The considered model escapes most of the drastic problems of the Sinister Universe \\cite{Glashow}, related to the primordial $^4He$ cage for $-1$ charge particles and a consequent overproduction of anomalous hydrogen \\cite{Fargion:2005xz}. These charged $^4He$ cages pose a serious problem for composite dark matter models with single charged particles, since their Coulomb barrier prevents successful recombination of positively and negatively charged particles. The doubly charged $A^{--}$ techniparticles considered in this paper, bind with $^4He$ in the techni-O-helium neutral states. %catalyzers of $AC$ binding and AC-leptons may thus escape this trap.  To avoid overproduction of anomalous isotopes, an excess of $-2$ charged techniparticles over their antiparticles should be generated in the Universe. In all the previous realizations of composite dark matter scenarios, this excess was put by hand to saturate the observed dark matter density. In our paradigm, this abundance of techibaryons and/or technileptons is connected naturally to the baryon relic density.  A challenging problem that we leave for future work is the nuclear transformations, catalyzed by techni-O-helium. The question about their consistency with observations remains open, since special nuclear physics analysis is needed to reveal what are the actual techni-O-helium effects in SBBN and in terrestrial matter. Another aspect of the considered approach is more clear. For reasonable values of the techiparticle mass, the amount of primordial $^4He$, bound in this atom like state is significant and should be taken into account in comparison to observations.  The destruction of techni-O-helium by cosmic rays in the Galaxy releases free charged techniparticles, which can be accelerated and contribute to the flux of cosmic rays. In this context, the search for techniparticles at accelerators and in cosmic rays acquires the meaning of a crucial test for the existence of the basic components of the composite dark matter. At accelerators, techniparticles would look like stable doubly charged heavy leptons, while in cosmic rays, they represent a heavy $-2$ charge component with anomalously low ratio of electric charge to mass.  To conclude, walking technicolor cosmology can naturally resolve most of problems of composite dark matter. Therefore, the model considered in this paper with stable $-2$ charged particles might provide a realistic physical basis for a composite dark matter scenario."
    },
    "0710/0710.5560_arXiv.txt": {
        "abstract": "We describe the time- and position-dependent point spread function (PSF) variation of the Wide Field Channel (WFC) of the Advanced Camera for Surveys (ACS) with the principal component analysis (PCA) technique. The time-dependent change is caused by the temporal variation of the $HST$ focus whereas the position-dependent PSF variation in ACS/WFC at a given focus is mainly the result of changes in aberrations and charge diffusion across the detector, which appear as position-dependent changes in elongation of the astigmatic core and blurring of the PSF, respectively. Using $>400$ archival images of star cluster fields, we construct a ACS PSF library covering diverse environments of the $HST$ observations (e.g., focus values). We find that interpolation of a small number ($\\sim20$) of principal components or ``eigen-PSFs'' per exposure can robustly reproduce the observed variation of the ellipticity and size of the PSF. Our primary interest in this investigation is the application of this PSF library to precision weak-lensing analyses, where accurate knowledge of the instrument's PSF is crucial. However, the high-fidelity of the model judged from the nice agreement with observed PSFs suggests that the model is potentially also useful in other applications such as crowded field stellar photometry, galaxy profile fitting, AGN studies, etc., which similarly demand a fair knowledge of the PSFs at objects' locations. Our PSF models, applicable to any WFC image rectified with the Lanczos3 kernel, are publicly available. ",
        "introduction": "} Even in the absence of atmospheric turbulence, the finite aperture of Hubble Space Telescope ($HST$) causes light from a point source to spread at the focal plane with the diffraction pattern mainly reflecting the telescope's aperture and optical path difference function. Although the point-spread-function (PSF) of $HST$ is already far smaller than what one can achieve with any of the current ground-based facilities, astronomers' endless efforts to push to the limits of their scientific observations with $HST$ ever increase the demand for the better knowledge of the instrument's PSF. Especially, since the installation of the Advanced Camera for Surveys (ACS) on $HST$, there have been concentrated efforts to carefully monitor and understand the instrument's PSFs, and to utilize the unparalleled resolution and sensitivity of ACS in gravitational weak-lensing (e.g., Jee et al. 2005a; Heymans et al. 2005; Schrabback et al. 2007; Rhodes et al. 2007).  Modeling the PSFs of ACS has proven to be non-trivial because of its complicated time- and position-dependent variation. The time-dependent change occurs due to the variation in the $HST$ focus, which relates to the constant shrinking of the secondary mirror truss structure and the thermal breathing of $HST$. The former is the main cause of the long-term focus change, and the secondary mirror position has been occasionally adjusted to compensate for this shrinkage (Hershey 1997). The latter is responsible for the short-term variation of the $HST$ focus and is affected by the instrument's earth heating, sun angle, prior pointing history, roll angle, etc. Even at a fixed focus value of $HST$, the PSFs of ACS also significantly change across the detector from the variation of the CCD thickness and the focal plane errors, which appear as position-dependent changes in charge diffusion and elongation of the astigmatic cores, respectively.  The strategies to model these PSF variations can be categorized into two types: an empirical approach based on real stellar field observations and a theoretical prediction based on the understanding of the instrument's optics. The first method treats the optical system of the instrument nearly as a blackbox and mainly draws information from observed stellar images. Although the PSF variation pattern can be most straightforwardly described by the variation of the pixel intensity as a function of position (e.g., Anderson \\& King 2006), frequently orthogonal expansion of the observed PSFs (e.g., Lauer 2002; Bernstein \\& Jarvis 2002; Refregier 2003) have been utilized to make the description compact and tractable. On the other hand, the second approach mainly relies on the careful analysis of the optical configurations of the instrument and receives feedbacks from observations to fine-tune the existing optics model. The TinyTim software (Krist 2001) is the unique package of this type applicable to most instruments of $HST$.  In this paper, we extend our previous efforts of the first kind (Jee et al. 2005a; 2005b; 2006; 2007) to describe the time- and position-dependent PSF variations of ACS/WFC now with the principal component analysis (PCA). In our previous work, we used ``shapelets'' (Bernstein \\& Jarvis 2002; Refregier 2003) to perform orthogonal expansion of the PSFs. Shapelets are the polar eigenfunctions of two-dimensional quantum harmonic oscillators, which form a highly localized orthogonal set. Although the decomposition of the stars with shapelets is relatively efficient and has proven to meet the desired accuracy for cluster weak-lensing analyses, the scheme is less than ideal in some cases. One important shortcoming is that it is too localized to capture the extended features of PSFs (Jee et al. 2007; also see \\textsection\\ref{section_basis_function}). In principle, the orthonormal nature of shapelets should allow us to represent virtually all the features of the target image when the number of basis functions are sufficiently large. However, this is not a viable solution not only because the convergence is slow, but also because the orthonormality breaks down in pixelated images for high orders as the function becomes highly oscillatory within a pixel.  The PCA technique provides us with a powerful scheme to obtain the optimal set of basis functions from the data themselves. Unlike ``shapelets'', the basis functions derived from the PCA are by nature non-parametric, discrete, and highly customized for the given dataset. Therefore, it is possible to summarize the multi-variate statistics, with a significantly small number of basis functions (i.e., much smaller than the dimension of the problem). For example, PCA has been applied to the classification of object spectra in large area surveys (Connolly et al. 1995; Bromley et al. 1998; Madgwick et al. 2003). It has been shown that only a small number ($10\\sim 20$) of the basis functions or $eigenspectra$ are needed to reconstruct the sample. The application of PCA to the PSF decomposition is used by the Sloan Digital Sky Survey (SDSS) to model the PSF variations (e.g., Lupton et al. 2001; Lauer 2002). Jarvis \\& Jain (2004) used the PCA technique to describe the variation in the PSF pattern in the CTIO 75 square-degree survey for cosmic shear analyses. They fit the ``rounding'' kernel component with PCA, not the PSF shape directly. This scheme is motivated by their shear measurement technique (i.e., reconvolution to remove systematic PSF anisotropy). However, in the current study we choose to fit the PSF shapes directly because this is more general in the sense that the rounding kernel components are not uniquely determined for a given PSF. In addition, our PSF library generating the PSF shapes directly has more uses in other studies.  We aim to construct a high-quality PSF library for the broadband ACS filters (F435W, F475W, F555W, F606W, F625W, F775W, F814W, and F850LP) from $>400$ archival stellar images, which sample a wide range of the $HST$ environments (e.g., the focus values). Our PSF models describe ACS PSFs in rectified images, specifically, drizzled using the Lanczos3 kernel with an output pixel scale of 0.05\\arcsec (see \\textsection\\ref{section_drizzling_kernel} for the justification of this choice). The results from this work are made publicly available on-line via the ACS team web site\\footnote{The full PSF library of ACS will become available at http://acs.pha.jhu.edu/$\\sim$mkjee/acs\\_psf/.}.  We will present our works as follows. The justification and the basic mathematical formalism of PCA are briefed in \\textsection\\ref{section_PCA}. In \\textsection\\ref{section_application_acs}, we demonstrate how the technique can be applied to ACS data with some test results. Focus dependency of the ACS PSFs, comparison with TinyTim, and strategies to find matching templates are discussed in \\textsection\\ref{section_discussion} before we conclude in \\textsection\\ref{section_conclusion}. ",
        "conclusions": "} We showed that the time- and position-dependent ACS/WFC PSF can be robustly described through PCA. The PCA technique allows us to perform orthogonal expansion of the observed PSFs with as few as 20 eigen-PSFs derived from the data themselves. This method is superior to our previous shapelet-based decomposition of the PSFs, capturing more details of the diffraction pattern of the instrument PSF. By interpolating the position-dependent variation of the eigen-PSFs with 5th order polynomials, we are able to recover the observed pattern of the PSF ellipticity and width variation. Although the TinyTim software provides a good approximation of the observed PSFs, we demonstrate that there are some important mismatches between the TinyTim prediction and the real PSFs, which cannot be attributed to CTE degradation of WFC over time. The CTE charge trailing effect should be negligible for these bright high S/N stars, and we do not observe any long-term variation of the pattern (i.e., increasing elongation in parallel read-out direction with time) due to the CTE degradation. Because typical science observations require integration of one or more orbits in broadband filters, the background levels are high ($\\sim200$ $e^{-}$ for integration of one orbit). These high background photons are supposed to fill the charge traps and thus mitigate the CTE effects. Therefore, we argue that the CTE-induced elongation is not likely to limit the application of our PSF models extracted from short-exposure observations to long-exposure science images.  We have compiled WFC PSFs from $>400$ stellar field observations, which span a wide range of $HST$ focus values. Although the current paper mainly deals with the ACS/WFC PSF issue in the context of weak-lensing analysis, we believe that our PSF model can be used in a wide range of the astronomical data analyses where the knowledge of the position-dependent WFC PSF is needed (e.g., crowded field stellar photometry, robust profile fitting of small objects, weak-lensing analyses, etc.).  ACS was developed under NASA contract NAS5-32865, and this research was supported by NASA grant NAG5-7697."
    },
    "0710/0710.0345_arXiv.txt": {
        "abstract": "We present results of a population synthesis study aimed at examining the role of spin-kick alignment in producing a correlation between the spin period of the first-born neutron star and the orbital eccentricity of observed double neutron star binaries in the Galactic disk. We find spin-kick alignment to be compatible with the observed correlation, but not to alleviate the requirements for low kick velocities suggested in previous population synthesis studies. Our results furthermore suggest low- and high-eccentricity systems may form through two distinct formation channels distinguished by the presence or absence of a stable mass transfer phase before the formation of the second neutron star.  The presence of highly eccentric systems in the observed sample of double neutron stars may furthermore support the notion that neutron stars accrete matter when moving through the envelope of a giant companion. ",
        "introduction": "Recent observations of single and binary pulsars have sparked new questions and challenged accepted ideas on the formation of neutron stars (NSs) and the nature of supernova (SN) kicks. In particular, measurements of pulsar radio emission polarization and pulsar wind nebulae symmetry axis directions have provided increasingly compelling evidence for the alignment of pulsar proper motions and pulsar rotation axes \\citep[e.g.][]{2005MNRAS.364.1397J, 2006ApJ...639.1007W, 2006ApJ...644..445N, 2007ApJ...660.1357N, 2007ApJ...664..443R, 2007arXiv0708.4251J}. Assuming the proper motion and pulsar spin axis directions are representative of the natal kick and NS progenitor rotation axis, the alignment suggests that natal kicks are preferentially aligned with the progenitor's rotation axis. Moreover, the increasing sample of double neutron star (DNS) binaries has revealed a possible correlation between the spin period $P_{\\rm spin}$ of the first born NS and the binary orbital eccentricity $e$ (see Fig.~\\ref{f1}) \\citep{2005ASPC..328...43M, 2005ApJ...618L.119F}. Such a relation arises naturally for symmetric SN explosions, but is highly constraining for asymmetric explosions. Our aim in this paper is to examine the role of spin-kick alignment in establishing the observed $P_{\\rm spin}$--$e$ correlation.  \\begin{figure} \\includegraphics[height=.24\\textheight]{observedpspinns1e.ps} \\caption{Orbital eccentricities and spin periods of observed DNSs in the Galactic disk. The dashed line represents the best-fitting log-linear curve $e=-1.7 + 1.2 \\log P_{\\rm spin}$.} \\label{f1} \\end{figure} ",
        "conclusions": "We investigated the role of spin-kick alignment in establishing a correlation between the spin period of the first-born NS and the orbital eccentricity of DNSs using the StarTrack binary population synthesis code and a simple prescription for the spin-up of a NS due to mass accretion. For DNSs forming through the standard evolutionary channel, spin-kick alignment is compatible with the observed $P_{\\rm spin}$--$e$ relation, but does not alleviate the requirement for low kick velocities proposed in previous population synthesis studies based on isotropic kick distributions. This is puzzling considering the stringent lower limits on the kick velocities imparted to the second-born NS in PSR\\,B1534+12 and PSR\\,B1913+16.  Moreover, if low kicks are imparted to the second-born NS in DNSs, the standard formation channel cannot explain the formation of the high-eccentricity systems PSR\\,B1913+16 and PSR\\,J1811-1736. A possible resolution would be a dichotomous formation channel where low-eccentricty DNSs are formed through the standard formation channel with low kicks imparted to the second-born NS, and high-eccentricity systems are formed through a formation channel where no mass transfer occurs after the common envelope phase of the second NS's progenitor and  \"normal\" kicks typical of isolated radio pulsars are imparted to the second-born NS. The low-eccentricity systems on the $P_{\\rm spin}$--$e$ relation can be formed with either polar or isotropic kicks. However, to obtain a $P_{\\rm spin}$--$e$ relation at high eccentricities, some spin-kick alignment is required to counteract the increased spread in the post-SN orbital eccentricities introduced by the larger kicks. We will explore this in the continuation of this investigation.     \\begin{theacknowledgments} This work is partially supported by a Packard Foundation Fellowship, a NASA BEFS grant (NNG06GH87G), and a NSF CAREER grant (AST-0449558) to VK. \\end{theacknowledgments}"
    },
    "0710/0710.3085_arXiv.txt": {
        "abstract": "We present the results of recent observations of phase-dependent variations in brightness designed to characterize the atmospheres of hot Jupiters.  In particular, we focus on recent observations of the transiting planet HD~189733b at 8~\\micron~using the \\emph{Spitzer Space Telescope}, which allow us to determine the efficiency of the day-night circulation on this planet and estimate the longitudinal positions of hot and cold regions in the atmosphere.  We discuss the implications of these observations in the context of two other successful detections of more sparsely-sampled phase variations for the non-transiting systems $\\upsilon$~And~b and HD~179949b, which imply a potential diversity in the properties of the atmospheres of hot Jupiters.  Lastly, we highlight several upcoming \\emph{Spitzer} observations that will extend this sample to additional wavelengths and more transiting systems in the near future. ",
        "introduction": "There are currently more than 20 known transiting planetary systems, of which the majority are gas-giant planets orbiting extremely close ($<$0.05~A.U.) to their parent stars\\footnote{See http://vo.obspm.fr/exoplanetes/encyclo/encycl.html for the latest count}.  These planets, known as ``hot Jupiters'', receive $>$10,000 times more radiation from their stars than Jupiter does from the sun, heating them to temperatures as high as 2000~K \\citep{har07}.  The picture is further complicated by the fact that most of these planets are expected to be tidally locked, with permanent day and night sides.  As a result, the equilibrium compositions and circulation patterns in these atmospheres are expected to differ significantly from those of the gas-giant planets in the solar system.  It is not surprising, then, that simple models for the atmospheres of these planets \\citep[see, for example, ][]{seag05,bar05,fort06b,bur07b} can sometimes differ significantly in both their assumptions and the conclusions that they reach; this is particularly true of the circulation models discussed in \\S\\ref{models}  Fortunately, it is possible to test the predictions of these models using observations of transiting systems.  From the wavelength-dependent depth of the transit, when the planet passes in front of the star, we can search for features in the transmission spectrum of the planet near the day-night terminator \\citep{char02,vid03,vid04,knut07a,bar07,tin07,ehren07}. Similarly, by measuring the depth of the secondary eclipse, when the planet moves behind the star, we can characterize the properties of the emission from the day side of the planet \\citep{char05,dem05,dem06,dem07,demory07,grill07,rich07,har07,knut07b,knut07c}.  Lastly, by measuring the changes in brightness over a significant fraction of a planet's orbit we can directly constrain the longitudinal temperature distribution across the planet's atmosphere \\citep{har06,cow07,knut07b}.  This last type of observation is particularly informative for hot Jupiters, which may have extreme changes in temperature between the day and night sides of the planet.  It is also a difficult observation to make in practice as it requires a very high precision over time scales of several days; this is why all of the successful detections to date have been made from space.  Although the shape of the ingress and egress during secondary eclipse also contain information about the brightness distribution on the day side of the planet \\citep{will06,rau07}, these variations are smaller and take place on shorter time scales, and there have not been any clear detections of this effect.  In this paper we focus on observations of phase variations, discussing the predictions of the current generation of atmospheric circulation models, the results of the three successful measurements of phase variations to date, and upcoming plans for more such observations in the near future. ",
        "conclusions": "The above discussion highlights the significant advances that have been made in the year since \\citet{har06} reported the first detection of the phase variation of an extrasolar planet.  However, the current sample is limited, and there is clearly much to be gained from observations of additional transiting systems, for which the planetary radius, orbital inclination, and dayside flux are all well-constrained.  There are a number of exciting programs scheduled for the upcoming \\emph{Spitzer} observing cycle which will address this issue.  First, we have a program which will observe HD 189733b continuously at 24~\\micron~over the same half orbit as \\citet{knut07b}, which should probe a different region of the atmosphere than the previous 8~\\micron~observations.  As a part of this same program we will also observe HD~209458b over half an orbit at both 8 and 24~\\micron, which will allow us to directly compare the wavelength-dependent properties of the atmospheric circulation for these two planets.  To date all of the observations we have described have examined planets with effectively circular orbits; however, there is another promising set of upcoming observations by G. Laughlin and G. Bakos that will observe two highly eccentric planets (HD 80606b and HAT-P-2b) continuously over $\\sim30$ hours centered on periastron passage.  Although only one of these planets, HAT-P-2b, is transiting, both observations should detect the rapid increase in brightness as the planetary atmosphere is flash-heated during this passage, providing a direct measurement of the radiative time constant.  Of the two methods described above, it is the more time-intensive, continuous observations of transiting systems that have provided the most detailed information about circulation in the atmospheres of hot Jupiters to date.  However, the detections of phase variations for $\\upsilon$~And~b and HD~179949b highlight the advantages of more sparsely-sampled observations in facilitating a much wider-ranging survey of a large number of planetary systems.  These two approaches are inherently complementary, and we expect that both will make valuable contributions to our understanding of these unusual planets in the near future."
    },
    "0710/0710.4170_arXiv.txt": {
        "abstract": "{We present the analysis and results of recent high-energy gamma-ray observations of the high energy-peaked BL Lac (HBL) object 1ES 1218+304 with the Solar Tower  Atmospheric Cherenkov Effect Experiment (STACEE). 1ES 1218+304 is an X-ray bright HBL at a redshift z=0.182. It has been predicted to be a gamma-ray emitter above 100 GeV, detectable by ground-based Cherenkov telescopes. Recently, this source has been detected by MAGIC and VERITAS, confirming these predictions. STACEE's sensitivity to astrophysical sources at energies above 100 GeV allows it to explore high energy sources such as X-ray bright active galaxies and gamma-ray bursts. We present results from STACEE observations of 1ES 1218+304 in the 2006 and 2007 observing seasons. }     \\begin{document} ",
        "introduction": "Active galaxies of the ``blazar'' class include BL Lac objects and flat-spectrum radio quasars (FSRQs), and are characterized by non-thermal continuum emission that extends from radio to high energy gamma rays. The spectral energy distributions (SEDs) of these sources typically have two broad peaks, one at low energies (radio to X-ray) and the other at higher energies (keV to TeV). In the framework of relativistic jet models, these objects are highly beamed sources, emitting plasma in relativistic motion (e.g. \\cite{Urry&Padovani1995}). In blazar SEDs, the low energy peak is explained as synchrotron emission from high energy electrons in the jet, while the high energy peak is probably due to inverse Compton emission. Several competing ``leptonic'' and ``hadronic'' jet model explanations exist for the high energy emission (e.g. see \\cite{Boettcher2002} \\& \\cite{Muckeetal2003} for reviews), and further broadband observation of blazars are needed to distinguish between these models.  Since the discovery of blazars as high energy gamma-ray sources by the Energetic Gamma-Ray Experiment Telescope (EGRET) on board the {\\sl Compton Gamma Ray Observatory} (CGRO) \\cite{Hartman1999} and the first detection of a TeV ($10^{12}$ eV) blazar by the Whipple Observatory (Mrk 421 \\cite{Punch1992}), the search has been on for more TeV blazars. The number count of TeV blazars is growing, with the advent of new generation atmospheric Cherenkov telescopes (ACTs) \\cite{Hessblazar19xx}. To date, almost all confirmed blazars detected at TeV energies are high-frequency-peaked BL Lac objects (HBLs), as opposed to quasars that constitute the majority of the EGRET detections. HBLs are a sub-class of BL Lac objects, characterized by lower luminosity than FSRQs and a synchrotron peak in the X-ray band \\cite{Urry&Padovani1995}. Blazars are categorized into different sub-classes based on the peak frequencies and the relative power in the low and high energy peaks of their SEDs \\cite{Fossati1998}. Given the high synchrotron peak frequencies of HBLs, indicating the presence of high energy electrons, these sources have been predicted to be good candidates for TeV emission, based on synchrotron self-Compton (SSC) emission models \\cite{Costamante&Ghisellini2002} as well as hadronic models \\cite{Mannheim1993}. Several of the ``extreme'' synchrotron BL Lacs \\cite{Costamante2001} have been detected at TeV energies, confirming these predictions.  1ES 1218+304 is an X-ray bright (flux at 1 keV $> 2$ $\\mu$Jy) HBL, categorized as an ``extreme'' BL Lac, and predicted to be a TeV source. The source was recently detected by both MAGIC \\cite{Albert2006} and VERITAS \\cite{Fortin2007}, at energies $>100$ GeV. At a redshift of $z=0.182$, 1ES 1218+304 is one of the most distant blazars detected to date. The source was never detected by EGRET, indicating that ACTs are sensitive to a different population of gamma-ray blazars than EGRET.  The Solar Tower Atmospheric Cherenkov Effect Experiment (STACEE) is a ground-based experiment that is sensitive to gamma rays above 100 GeV. STACEE observations of AGN are motivated by the need to understand particle acceleration and emission mechanisms in blazars, as well as their interaction with the extragalactic background light (EBL). Despite what is already known, a great deal remains to be discovered regarding the physics of blazars.  STACEE's extragalactic observing program has included both HBLs, as well as LBLs \\cite{Mukherjee2006, Mukherjee2005}. Recent observations of 1ES 1218+304 with STACEE were motivated by the detection of TeV emission from the source by MAGIC, providing further evidence that X-ray bright HBLs tend to be strong VHE sources. STACEE observations of 1ES 1218+304 were carried out in the 2006 and 2007 observing seasons. In this paper we present a summary of STACEE observations of the HBL 1ES 1218+304. ",
        "conclusions": "STACEE observed the high frequency-peaked BL Lac object 1ES 1218+304 in 2006 and 2007. After all cuts and padding 28.3 hr of data yielded an ON-source excess with a significance of $2.3\\sigma$ consistent with no detected flux. Simulated effective areas (Figure 3) were used to derive flux upper limits. For the combined 2006 and 2007 data sets the differential flux upper limit at the 99\\% confidence level was derived to be $< 5.2 \\times 10^{-6}$ m$^{-2}$ s$^{-1}$ TeV$^{-1}$ at 150 GeV, the energy threshold of STACEE. The upper limit was calculated assuming that the differential flux of photons follows a power law with an index of $-3.0$, as measured by MAGIC \\cite{Albert2006}. The STACEE upper-limit is shown overlaid on the MAGIC spectrum in Figure 4. These numbers are consistent with the spectrum measured by MAGIC \\cite{Albert2006}, and with the recent VERITAS results indicating a weak gamma-ray source at $\\sim5$\\% of the Crab flux \\cite{Fortin2007}.  \\begin{figure} \\begin{center} \\includegraphics*[width=0.48\\textwidth,height=2.5in]{rmukherjee_0403_fig4.ps} \\end{center} \\caption{Gamma-ray spectrum of 1ES 1218+304, as measured by MAGIC (figure from \\cite{Albert2006}). The STACEE 99\\% flux upper limit is overlaid on the plot at 150 GeV, the energy threshold of STACEE, as obtained from detector simulations. The upper limit was calculated assuming the power-law spectral index of $-3.0$ measured by MAGIC. } \\label{fig4} \\end{figure}  \\bigskip \\vskip 2in  {\\bf Acknowledgements: } Many thanks go to the staff of the National Solar Tower Test Facility, who have made this work possible. This work was funded in part by the US National Science Foundation, the Natural Sciences and Engineering Research Council of Canada, Fonds Quebecois de la Recherche sur la Nature et les Technologies, the Research Corporation, and the University of California at Los Angeles. R. M. acknowledges support from NSF grant 0601112."
    },
    "0710/0710.1080_arXiv.txt": {
        "abstract": "HCN and CO line diagnostics provide new insight into the OH megamaser (OHM) phenomenon, suggesting a dense gas trigger for OHMs. We identify three physical properties that differentiate OHM hosts from other starburst galaxies: (1) OHMs have the highest mean molecular gas densities among starburst galaxies; nearly all OHM hosts have $\\bar{n}({\\rm H}_2)=10^3$--$10^4$~cm$^{-3}$ (OH line-emitting clouds likely have $n({\\rm H}_2)>10^4$~cm$^{-3}$). (2) OHM hosts are a distinct population in the nonlinear part of the IR-CO relation. (3) OHM hosts have exceptionally high dense molecular gas fractions, $L_{\\rm HCN}/L_{\\rm CO}>0.07$, and comprise roughly half of this unusual population. OH absorbers and kilomasers generally follow the linear IR-CO relation and are uniformly distributed in dense gas fraction and $L_{\\rm HCN}$, demonstrating that OHMs are independent of OH abundance. The fraction of non-OHMs with high mean densities and high dense gas fractions constrains beaming to be a minor effect: OHM emission solid angle must exceed $2\\pi$ steradians. Contrary to conventional wisdom, IR luminosity does not dictate OHM formation; both star formation and OHM activity are consequences of tidal density enhancements accompanying galaxy interactions.  The OHM fraction in starbursts is likely due to the fraction of mergers experiencing a temporal spike in tidally driven density enhancement. OHMs are thus signposts marking the most intense, compact, and unusual modes of star formation in the local universe. Future high redshift OHM surveys can now be interpreted in a star formation and galaxy evolution context, indicating both the merging rate of galaxies and the burst contribution to star formation. ",
        "introduction": "OH megamasers (OHMs) are rare luminous 18~cm masers associated with major galaxy merger-induced starbursts. The hosts of OHMs are (ultra)luminous IR galaxies ([U]LIRGs), and the OHM fraction in (U)LIRGs peaks at about 1/3 % in the highest luminosity mergers \\citep{darling02a}. It is not known whether all major mergers experience an OHM stage or what detailed physical conditions produce OHMs, but it is clear that OHMs are a radically different phenomenon from the aggregate OH maser emission associated with ``normal'' (Galactic) modes of star formation in galaxies. \\citet{lo05} posed a key question: why do $80\\%$ of LIRGs show no OHM activity?  To reframe the question:  given two merging systems with similar global IR and radio continuum properties in the same morphological stage of merging, why does one show OHM emission while the other does not? What is the difference between the two systems?  Perhaps there is no difference and the fraction of OHMs among mergers simply reflects beaming or OH abundance. Or perhaps OHM activity depends on small-scale conditions that are decoupled from global properties of mergers.  The provenance of OHM emission vis-\\`{a}-vis the host galaxy has been extensively investigated in the radio through X-ray bands by comparing samples of OHM galaxies to similarly selected non-masing control samples. For example, \\citet{darling02a} and \\citet{baan06} studied radio and IR properties vis-\\`{a}-vis the AGN versus starburst contributions to OHM activity, \\citet{baan1992} and \\citet{darling02a} investigated the OHM fraction in (U)LIRGs versus star formation rate and IR color, \\citet{baan1998} and \\citet{darling06} used optical spectral classification to distinguish populations and to quantify AGN fraction in OHM hosts, and \\citet{vignali05} conducted an X-ray study of the contribution of AGNs to OHM hosts. While some of these studies pointed to minor differences in statistical samples of OHM hosts versus nonmasing systems, they could not identify on a case-by-case basis which systems would harbor OHMs and which would not based on any observable quantity except the OH line itself.  Theoretical modeling of OHM formation has seen a recent renaissance: \\citet{parra2005} model the $\\sim50$~pc molecular torus in III~Zw~35 and show how OHM emission is a stochastic amplification of unsaturated emission by multiple overlapping clouds, and \\citet{lockett07} show how the general excitation of OHMs is fundamentally different from Galactic OH maser emission and predict that a single excitation temperature governs all 18~cm OH lines.  While the physics of OHMs is crystallizing, and models predict that beaming is not likely to be the dominant factor in the OHM fraction among (U)LIRGs, it remains unclear on a case-by-case basis what conditions found in starbursts drive or prohibit OHM formation.   Here we describe a dense gas trigger for OHM formation, at last identifying physical observable properties that differentiate OHMs from nonmasing mergers.  We identify OHMs, OH absorbers, OH kilomasers, and OH non-detections in the Gao \\& Solomon (2004a; hereafter GS04a) sample (\\S \\ref{sec:sample}) and employ CO($1-0$) and HCN($1-0$) molecular gas tracers to show that while OH absorbers appear nearly uniformly distributed in $L_{\\rm IR}$ and $L_{\\rm HCN}$, OHMs represent the {\\it majority} of the nonlinear population in the IR-CO relation (\\S \\ref{sec:results}).  In combination with a Kennicutt-Schmidt-based star formation model of CO line emission by \\citet{krumholz07}, we identify a high mean molecular density driving OHM emission, and from the HCN/CO ratio we find that OHM galaxies are exclusively high dense gas fraction starbursts (\\S \\ref{sec:results}). Now that we can at last observe quantities that are highly predictive of OHM activity, we can employ OHMs at high redshifts as probes of major galaxy mergers and extreme star formation (\\S \\ref{sec:discussion}). ",
        "conclusions": "We have identified three closely related physical properties that differentiate OHMs from other starburst galaxies: OHM hosts have the highest mean molecular gas densities, they are a distinct population in the nonlinear part of the IR-CO relation, and they reside in galaxies with exceptionally high dense molecular gas fractions. We conclude that molecular gas must be concentrated and massive in order to reach the mean density required to form an OHM in a galactic nucleus. IR luminosity is not a condition for OHM formation; both star formation and OHM activity are consequences of the tidal density enhancements accompanying galaxy interactions. The fraction of OHMs in dense starbursts constrains OHM beaming to be a minor effect: OHM solid angle emission must be greater than $2\\pi$ steradians. These conclusions are in good agreement with the stochastic cloud-cloud overlap amplification model by \\citet{parra2005}. The rather uniform distribution of OH absorbers in IR, HCN, and CO luminosity suggests that OH abundance is not a significant factor in OHM formation. The main caveat to these conclusions is that the sample of OHMs with HCN observations remains small, and should be expanded, particularly to higher redshifts to include ``typical'' OHMs.  OHMs are signposts of the most intense, compact, and unusual modes of star formation in the local universe, and surveys for OHMs will now provide detailed information about the detected host galaxies and their mode of star formation. The missing datum required for a complete interpretation of OHM surveys, however,  is the OHM lifetime."
    },
    "0710/0710.4941_arXiv.txt": {
        "abstract": "The bulk composition of an exoplanet is commonly inferred from its average density. For small planets, however, the average density is not unique within the range of compositions. Variations of a number of important planetary parameters---which are difficult or impossible to constrain from measurements alone---produce planets with the same average densities but widely varying bulk compositions. We find that adding a gas envelope equivalent to 0.1\\%-10\\% of the mass of a solid planet causes the radius to increase 5-60\\% above its gas-free value. A planet with a given mass and radius might have substantial water ice content (a so-called ocean planet) or alternatively a large rocky-iron core and some H and/or He.  For example, a wide variety of compositions can explain the observed radius of GJ 436b, although all models require some H/He. We conclude that the identification of water worlds based on the mass-radius relationship alone is impossible unless a significant gas layer can be ruled out by other means. ",
        "introduction": "Out of over 250 exoplanets known to date, over 20 are known to transit their stars. Transiting planets are important because we can derive the precise mass and radius, and can begin to determine other planetary properties, such as the bulk composition.  Much attention has been given to ``ocean planets'' or ``water worlds'', planets composed mostly of solid water \\citep{kuch2003, lege2004}. If a water world is found close to a star, it would be strong evidence for migration because insufficient volatiles exist near the star for {\\it in situ} formation. The proposed identification of water worlds is through transits. From a measured mass and radius a low-density water planet could potentially be identified.  We examine the possibility that water worlds cannot be uniquely identified based on the mass and radius of a transiting planet.  An alternative interpretation could be a rocky planet with a thick hydrogen-rich atmosphere.  Most authors have assumed that solid planets in the 5 to 10 $M_{\\oplus}$ range have an insignificant amount of hydrogen \\citep{vale2006, vale2007a, fort2007, seag2007, sels2007, soti2007}.  Exoplanets have, however, contradicted our basic assumptions before. Notable examples include: the existence of hot Jupiters; the predominance of giant planets in eccentric orbits; and the gas-rock hybrid planet HD~149026b with its $\\sim 60 M_{\\oplus}$ core and $\\sim 30 M_{\\oplus}$ H/He envelope \\citep{sato2005}.  We adopt the idea that a wide range of atmospheric formation and loss mechanisms exist and can lead to a range of atmosphere masses on different exoplanets. We explore the mass-radius relationship for the lowest-mass exoplanets yet detected ($\\sim$~5-20 $M_{\\oplus}$) in order to identify potential ambiguities that result from the presence of a massive atmosphere. We explore atmospheres ranging from $\\sim 10^{-3} M_{\\oplus}$ (10 $\\times$ Venus' atmospheric mass) to $\\sim 1 M_{\\oplus}$ (the estimated mass of Uranus' and Neptune's H/He \\citep{guil2005, podo1995, hubb1991}), with a focus on the smaller mass range. We also explore potential compositions for the transiting Neptune-size planet GJ~436b \\citep{butl2004, gill2007a}. ",
        "conclusions": "Our fiducial planet consists of a 30\\% Fe core and a 70\\% MgSiO$_3$ mantle, roughly analogous to Earth. We used a H/He mixture with helium mass fraction $Y=0.28$ (the He mass fraction of the solar nebula).  We chose $T_{eq}=300$~K, based on the observation that a planet around an M-dwarf at an orbital distance of 0.1 AU has a similar equilibrium temperature to Earth (assuming similar albedos). We set $T_{eff}=30$~K, similar to Earth and Uranus\\footnote{Earth has 44 $\\times 10^{12}$~ W \\citep{poll1993} and Uranus has $340 \\times 10^{12}$~W of energy flow \\citep{pear1990}.}. For the atmospheric parameters, we fixed $\\mu_0 = \\cos 60^{\\circ}$ and $\\gamma=0.1$ to represent radiation absorbed deep in the atmosphere\\footnote{In comparison $\\gamma=10$ would correspond to absorption high in the atmosphere.}. We later investigate variations on $Y$, $T_{eff}$, $T_{eq}$ and $\\gamma$.  Figure~1 shows a plot of the mass-radius relationship for fiducial planets of masses 5, 10, 15, and 20 $M_{\\oplus}$.  For each planet mass, we added atmospheres ranging in mass from 0.001-1 $M_{\\oplus}$.  A robust finding for all models is that a small amount of gas creates a large radius increase.  While this result is expected, the radius increase is far more dramatic than anticipated.  For example, an H/He atmosphere of $\\sim 0.001$ by mass---only ten times greater than Venus' atmospheric mass fraction---is required for a noticeable radius increase. As seen in Figure~1, adding a hydrogen-helium atmosphere with just 0.1\\% of the mass of a 10 $M_{\\oplus}$ rocky planet results in a 5\\% increase in the planetary radius---within a measurement precision that has been obtained for currently known transiting planets.  \\begin{figure} \\plotone{f1} \\caption{The increase in radius due to adding H/He to a solid planet. A H/He layer of 0.002-1 $M_{\\oplus}$ is added to a solid planet of 5, 10, 15, or 20 $M_{\\oplus}$, with fiducial model parameters (30\\% Fe and 70\\% MgSiO$_3$). The black points are for atmospheres at 0.01 $M_{\\oplus}$ and every 0.1 $M_{\\oplus}$ afterwards. The mass-radius relationship of solid planets with no gas is plotted for comparison. The water (blue), rock (red), and iron (green) curves are taken from \\citet{seag2007} and represent homogeneous solid planets. Intermediate compositions for differentiated planets are, from top down: dashed-blue, 75\\% H$_2$O, 22\\% MgSiO$_3$, and 3\\% Fe; dashed-dotted-blue, 48\\% H$_2$O, 48.5\\% MgSiO$_3$, and 6.5\\% Fe; dotted-blue, 25\\% H$_2$O, 52.5\\% MgSiO$_3$, and 22.5\\% Fe; dashed-red, 67.5\\% MgSiO$_3$ and 32.5\\% Fe; and dotted-red, 30\\% MgSiO$_3$ and 70\\% Fe. In general, the addition of a gas layer of up to $\\sim 5$\\% of the solid planet mass will inflate the radius of a rocky-iron planet through the range of radii corresponding to water planets with different water mass fractions. \\label{fig:gas1}} \\end{figure}  As a second example, adding a gas layer of H/He equal to 1\\% of the mass of our fiducial planets increases the radius by $\\sim$ 20\\% of the original planet radius, or by about 0.35 $R_{\\oplus}$ for the planet masses we considered.  Our major finding is that that exoplanets with a significant H/He layer cannot be distinguished from water worlds, based on $M_p$ and $R_p$ alone.  For our fiducial solid exoplanets, adding up to 5\\% H/He by mass (for 10 $M_{\\oplus}$ planets) is sufficient to push the planet's radius through the entire range of radii corresponding to solid planets with no gas, including planets with up to 100 percent water composition. While we have not completed an exhaustive study of possible compositions, we find the non-uniqueness of water planets to be valid for any conditions we investigated.  This generic finding holds for a wide range of assumptions of assumed temperatures.  Taking our fiducial model, we vary $T_{eq}$, $T_{eff}$, and $\\gamma$ individually.  For a 10 $M_{\\oplus}$ solid planet with an additional 0.1 $M_{\\oplus}$ H/He atmosphere, increasing $T_{eq}$ from 300 to 500 K increases the radius by about ~1\\% (Figure~2). For the same planet varying $T_{eff}$ from 10~K to 50~K results in an 8\\% increase in radius. While large, this value is comparable to expected radius uncertanties for these planets \\citep{gill2007a, gill2007b, demi2007}. Varying the altitude where radiation is absorbed (specified by $\\gamma$) has a much smaller effect on the planet radius. Varying $\\gamma$ from 0.1 to 10 causes the radius to decrease by 0.2\\%.  \\begin{figure} \\plotone{f2} \\caption{The effects on the radius of varying the equilibrium temperature (top) and effective temperature (bottom) for a 10 $M_{\\oplus}$ planet with otherwise fiducial parameters. $T_{eq}$ values of 300, 400, and 500~K are plotted to simulate the effect of uncertainty in the orbital parameters and albedo on the expected radius. $T_{eff}$ values of 10, 30, and 50 K are plotted on the same scale as the $T_{eq}$ plot, to show uncertainties in the planet's interior temperature. Uncertainties in the internal energy of a planet lead to large variations in radii for a given mass, showing how temperature is a large uncertainty in the interpretation of a planet's internal composition. \\label{fig:temp}} \\end{figure}   A corollary of our main result is that when a planet has a significant H/He atmosphere there is a wide degeneracy in allowable internal composition. This is not just compositional, but also relates to the trade off of temperature and mass of H/He gas.  It could be argued that specifying a planet's composition implies a particular internal thermal profile derived from a consistent cooling history. As addressed in \\S2, the many unknowns and free input parameters for rocky planet interiors---such as the possible differences between atmospheric and interior compositions, equation of states, and the effect of tides on the planet's cooling history---prevent a self-consistent solution for the present time.  How could a 5--20 $M_{\\oplus}$ exoplanet get a substantial H/He layer? Two different scenarios may produce them: direct capture of gas from the protoplanetary disk (possibly modified by the escape of some fraction of the original gas) or outgassing during accretion.  A planet may capture and retain up to 1 to 2 $M_{\\oplus}$ of H/He if the planetary core did not grow quickly enough to capture more before the gas in the disk evaporated (as is the paradigm for Uranus and Neptune). Alternatively, for short-period exoplanets, a 1 to 2~$M_{\\oplus}$ H/He envelope may result after substantial loss of an initially massive gas envelope from irradiated evaporation \\citep[e.g.][]{bara2006}.  \\citet{alib2006} consider atmospheric evaporation during migration, and conclude that the $10 M_{\\oplus}$ innermost planet in HD~69830, at 0.08 AU, kept $\\sim$2 $M_{\\oplus}$ of H/He over the 4 Gyr lifetime of the star.  Little attention has been given to the mass and composition of exoplanet atmospheres from outgassing. Venus' atmosphere is $10^{-4} M_{\\oplus}$; if Venus had a surface gravity high enough to prevent H escape, its atmosphere would be over $10^{-3} M_{\\oplus}$.  Even more massive H-rich atmospheres are possible. If a massive iron-silicate planet formed with enough water, the iron may react with the water during differentiation, liberating hydrogen gas \\citep{ring1979, waen1994}.  L. Elkins-Tanton et al. (in prep.) estimate that the maximum H component is about six percent by mass for a terrestrial-composition planet. For a $10 M_{\\oplus}$ planet this would result in a $0.6 M_{\\oplus}$ H envelope.  For short-period, low-mass planets, theoretical arguments of atmospheric escape may be the best way to identify a water world based on the mass and radius measurements alone \\citep{lege2004, sels2007}. Indeed, our assumption of H/He atmospheres for exoplanets relies on the condition that atmospheric mass loss has not evaporated all of the H/He. In the absence of hydrodynamic escape, the exospheric temperature (and not the atmospheric $T_{eff}$) drives the thermal Jeans escape of light gases. Earth and Jupiter both have exobase temperatures of 1000~K \\citep{depa2001}, significantly above their $T_{eff}$ of 255~K and 124.4~K respectively \\citep{cox2000}. Uranus and Neptune have exobase temperatures around 750~K \\citep{depa2001}. We note that because GJ 436 must have at least 1~$M_{\\oplus}$ of H/He, its exospheric temperature is not too high.  On the other extreme, planets of $5 M_{\\oplus}$ would require very low exospheric temperatures ($\\sim 300$~K) to retain a massive atmosphere over the course of billions of years. Nevertheless, a young $5 M_{\\oplus}$ Earth-mass planet with a captured atmosphere could still have a H/He atmosphere and an old $5 M_{\\oplus}$ planet could retain a substantial He fraction, making its compositional identification ambiguous.  \\begin{figure} \\plotone{f3} \\caption{Density vs. radius for three different potential compositions of GJ436b. From top to bottom: solid (black) curve, 19.0 $M_{\\oplus}$ core (30\\% Fe, 70\\% MgSiO$_3$) with 3.2 $M_{\\oplus}$ H/He (Y=0.28); dotted (red) curve, 20.0 $M_{\\oplus}$ core (100\\% (Mg,Fe)SiO$_3$) with 2.2 $M_{\\oplus}$ H (Y=0); dashed (blue) curve, 20.5 $M_{\\oplus}$ core (90\\% H$_2$O, 10\\% MgSiO$_3$) with 1.7 $M_{\\oplus}$ H/He (Y=0.28). All three planets have the same total radius (4.3 $R_{\\oplus}$) and total mass (22.2 $M_{\\oplus}$). \\label{fig:GJinterior}} \\end{figure}   We now turn to a qualitative study of GJ436b, to show that the interpretation that GJ436b is a water world akin to Uranus and Neptune \\citep{gill2007b} is not the only possibility.  We consider the GJ~436b values $M_p = 22.2 M_{\\oplus}$ and $R_p = 4.3 R_{\\oplus}$ from \\citet{demi2007}. The internal structure in Figure~3 shows how three planets with very different internal compositions can have the same total mass and radius.  We first explore a planet similar to our fiducial model: a 22.2 $M_{\\oplus}$ solid planet with Earth-like iron/rock mass ratio (30/70), $T_{eff}=30$~K, and $T_{eq}=600$~K in rough agreement with the orbital parameters (assuming an albedo of 0.1). By adding $\\sim$3.2$M_{\\oplus}$ of hydrogen-helium to the 19.0 $M_{\\oplus}$ solid planet, we are able to reproduce GJ~436b's radius. We note that the mass of gas is 15\\% of the solid mass, likely too much to have originated from outgassing, and so capture must be at least partially invoked to explain such a massive atmosphere.  The second composition for GJ~436b we considered is for water worlds, one with a 50\\% water mantle (by mass) and 50\\% silicate core, and another with 90\\% water mantle and a 10\\% silicate core. These planets also need some H/He to match the known radius, 12\\% and 8\\% by mass, respectively.  The third model approximates planets with atmospheres created from outgassing, considering an extreme scenario where all of the available water has oxidized iron, leaving a 100\\% (Mg,Fe)SiO$_3$ solid planet core. To match the observed radius, a 22.2 $M_{\\oplus}$ planet requires $\\sim 2.2 M_{\\oplus}$ of H alone, a case that assumes no initial trapping and subsequent outgassing of He. The pure-hydrogen atmosphere is 10\\% of the mass of the solid planet. This is above the theoretical maximum of outgassing based on observed abundances of metallic iron in chondritic meteorites from our solar system (see Elkins-Tanton et al. in prep).  Although not an exhaustive study, the range of interior compositions illustrates the variety of possibilities, though all models require some H/He.  While our study is preliminary, we make the robust point that H-rich thick atmospheres will confuse the interpretation of planets based on a measured mass and radius.  This point is independent of the uncertainties retained by our model including $T_{eq}$, $T_{eff}$, the mass fraction of H/He, and the mixing ratio of H and He. We find that the identification of water worlds based on the mass-radius relationship alone is impossible unless a significant gas layer can be ruled out by other means. Spectroscopy is the most likely means and may become routine with transit transmission and emission spectroscopy, and eventually with spectroscopy by direct imaging."
    },
    "0710/0710.4036_arXiv.txt": {
        "abstract": "We present near-infrared (H- and K-band) integral-field observations of the circumnuclear star formation rings in five nearby spiral galaxies. The data, obtained at the {\\it Very Large Telescope} with the SINFONI spectrograph, are used to construct maps of various emission lines that reveal the individual star forming regions (\"hot spots\") delineating the rings. We derive the morphological parameters of the rings, and construct velocity fields of the stars  and the emission line gas. We propose a qualitative, but robust, diagnostic for relative hot spot ages based on the intensity ratios of the emission lines \\brg , \\hel , and \\fetwo . Application of this diagnostic to the data presented here provides tentative support for a scenario in which star formation in the rings is triggered predominantly at two well-defined regions close to, and downstream from, the intersection of dust lanes along the bar with the inner Lindblad resonance. ",
        "introduction": "In many spiral galaxies of early- and intermediate Hubble type (Sa-Sc),  active star formation is organized in a ring-like structure. These  star formation rings offer a unique opportunity to study massive star formation in  external galaxies: they produce up to 2/3 of the bolometric luminosity of their host galaxies \\citep[e.g. NGC\\,7469; ][]{gen95}, and often contain a large fraction of the entire star formation activity of the galaxy.  The general picture of why molecular gas assembles in a ring is well understood as a natural consequence of a non-axisymmetric gravitational  potential, nearly always due to the presence of a stellar bar or oval distortion \\citep[e.g.][]{com85,kna95,hel96,but96}. Because of its dissipative nature, molecular gas accumulates around the radii at which the stellar orbits experience dynamical resonances with the rotating bar potential. Depending on the pattern speed of the bar and the rotation curve of the galaxy, there can be one or multiple such resonances. The high gas densities found in the rings, combined with a variety of excitation mechanisms such as ultra-violet radiation from young stars and mechanical shocks due to energetic outflows from massive stars and/or an active galactic nucleus (AGN) help reveal the physical state of the interstellar matter (ISM), be it molecular, atomic, or ionized gas.  Besides being fascinating laboratories in their own right, star formation rings are important also for the secular evolution of disk galaxies \\citep{kor04}. This is particularly true for the innermost of the dynamical resonances which is called either the \"inner Lindblad resonance\" (ILR) or - in cases where a compact massive  object leads to an additional dynamical resonance - the ``nuclear Lindblad resonance'' \\citep*[NLR; ][]{fuk98}. Either of these resonances can produce gas rings with radii of  a few hundred pc or less, depending on the enclosed mass, and the rotation speed of the galaxy disk. On these spatial scales, a number of processes can cause the gas to lose  angular momentum and to subsequently flow towards the nucleus. Examples for such processes include torques due to the stellar potential \\citep{gar05}, dynamical friction between giant molecular clouds that form within the ring due to self-gravity of the gas \\citep*{fuk00}, the formation of spiral density waves \\citep{eng00}, or mechanical energy released by star formation in the ring via stellar winds and/or supernova explosions.  Understanding the gas dynamics and star formation processes of nuclear\\footnote{Throughout this paper, the term ``nuclear ring'' is used  to imply that it is the innermost (star forming) gas ring that can be resolved with the resolution of our data, typically a few hundred pc in diameter.} rings in disk galaxies  is therefore crucial for developing models of gas accumulation at the very nucleus of a galaxy and the evolution of any compact massive object (CMO) which can exist in the form of a nuclear star cluster \\citep[NC,][]{boe02} and/or a  supermassive black hole (SMBH). While some theoretical models exist on the gas behavior around a CMO \\citep[e.g.][]{fuk00}, observational data to constrain these models are rare, mostly because of the limited  spatial resolution of mm-observations. Only recently has it become possible to study  the molecular gas flows within a few tens of pc from the nucleus in a small number of nearby  galaxies \\citep[e.g.][]{sch03,sch06,sch07}.  In order to increase the number of well-studied nuclear rings,  we have begun a project to study the near-infrared (NIR) properties of five such objects,  using the SINFONI integral-field spectrograph on the Very Large Telescope (VLT). This paper discusses the morphologies, star  formation rates, and kinematic properties of the rings. In a subsequent paper, we will make use of the spectroscopic information contained in the SINFONI data  to investigate in more detail the star formation process, stellar populations, and gas excitation  mechanism(s) in individual galaxies, both in the rings and the galaxy nuclei.  This paper is structured as follows. In \\S\\ref{sec:data}, we describe the sample selection, observational details, and data reduction techniques  common to all galaxies. The continuum and emission line morphology of the rings as well as the velocity fields of stars and gas for the individual galaxies are presented in \\S\\ref{sec:diagnostics}. We discuss the results in the context of competing models for the propagation of star formation in the rings in \\S\\ref{sec:discuss}, and summarize our analysis in \\S\\ref{sec:summary}. ",
        "conclusions": "} Based on high-resolution NIR integral-field observations of five nuclear star formation rings, we have presented their emission line morphologies and velocity structure both in gas and stars. We have introduced a new method to derive relative hot spot ages along the rings using the relative strengths of the \\hel , \\brg , and \\fetwo\\ lines. We employ this method to investigate the plausibilty of two competing scenarios for the way star formation is induced in nuclear rings, namely i) the ``popcorn'' model in which hot spots appear stochastically around the ring, and ii) the ``pearls on a string'' scenario in which star formation is triggered predominantly at two overdensity regions on either side of the ring. Only the latter predicts a well-ordered age sequence of hot spots along either half of a ring.  The data presented in this study provide tentative support for the ``pearls on a  string'' scenario, in that three out of five sample galaxies show some evidence for an age gradient of hot spots along the ring, while the remaining two galaxies have incomplete information and thus are consistent with either model.  Of course, it might well be that star formation proceeds differently in some rings than in others. However, given that nuclear rings appear to form via a common mechanism (i.e. the gas response to a bar-shaped potential), it is not unreasonable to expect that they also follow a common path for inducing star formation. The small number of objects described here is clearly insufficient to provide reliable statistics, and similar studies of larger galaxy samples are needed to decide whether a particular mechanism governs the star formation in the majority of nuclear rings. Be that as it may, the proposed method of using NIR line ratios to estimate relative ages of young star clusters has been demonstrated to be a powerful tool for studies along this line."
    },
    "0710/0710.5801_arXiv.txt": {
        "abstract": "We discuss the use of Sloan Digital Sky Survey (SDSS) {\\it ugriz} point-spread function (PSF) photometry for setting the zero points of {\\it UBVRI} CCD images. From a comparison with the Landolt (1992) standards and our own photometry we find that there is a fairly abrupt change in {\\it B, V, R, \\& I} zero points around $g, r, i \\sim 14.5$, and in the $U$ zero point at $u \\sim 16$.  These changes correspond to where there is significant interpolation due to saturation in the SDSS PSF fluxes. There also seems to be another, much smaller systematic effect for stars with $g, r \\gtrsim 19.5$. The latter effect is consistent with a small Malmquist bias. Because of the difficulties with PSF fluxes of brighter stars, we recommend that comparisons of {\\it ugriz} and {\\it UBVRI} photometry should only be made for unsaturated stars with $g, r$ and $i$ in the range 14.5 -- 19.5, and $u$ in the range 16 -- 19.5. We give a prescription for setting the $UBVRI$ zero points for CCD images, and general equations for transforming from {\\it ugriz} to {\\it UBVRI}. ",
        "introduction": "When CCD images of a field are taken it is necessary to determine the photometric zero points from stars of known magnitudes. It is, however, not unusual for there to be no stars with $UBVRI$ photometry available. Fortunately, the Sloan Digital Sky Survey (SDSS) now provides homogenous $ugriz$ photometry for stars in a large fraction of the northern sky out of the plane of the Milky Way.  Technical details of the SDSS are given in \\citep{york00} and \\citep{stoughton02}.  The $ugriz$ system \\citep{fukugita96} is significantly different from the widely used $UBVRI$ Johnson-Cousins system \\citep{cousins76}, so it is necessary to transform between the two systems. A number of papers \\citep{fukugita96,smith02,karaali03,karaali05,bilir05,jordi06,rodgers06,ivezic07,davenport07,bilir07} have considered the transformations between $ugriz$ and $UBVRI$ (see Section 6 for a discussion of these transformations).  During the course of using SDSS $ugriz$ photometry to establish the zero points for comparison stars for photometry of active galactic nuclei (AGN), we noticed that the zero points were different for the fainter stars in a field than for the brighter stars.  The difference was in the sense that stars with $g \\lesssim 14$ were systematically brighter than predicted from the SDSS magnitudes. The difference did not seem to depend on the color of the stars and a check of the CCD used showed no evidence for non-linearity.  A subsequent comparison of magnitudes of Landolt standards \\citep{landolt92} revealed a similar zero-point difference for stars brighter or fainter than $r \\sim 14$.  In this note we report results of our investigation of the limitations of using SDSS photometry for bright stars, and give a prescription for setting zero points in CCD images taken through {\\it UBVRI} filters. ",
        "conclusions": "The transformations we give in equations (1) -- (5) are consistent with the range of previously published transformations. Because our linear transformation equations are derived for a practical purpose of calibrating $UBVRI$ photometry, and are available for each of the Johnson-Cousins filters individually, our transformations are different in nature from those previously published.  We briefly summarize here previously published transformations and discuss how they differ from the ones given above. \\cite{fukugita96} give synthetic transformations from $UBVRI$ to $u'g'r'i'z'$. \\cite{smith02} gave transformations between $UBVRI$ and $u'g'r'i'z'$ magnitudes observed with the Photometric Telescope (PT) at Apache Point Observatory for some filters and for colors. \\cite{rodgers06} give improved color transformations between $u'g'r'i'z'$ and $UBVRI$ for main-sequence stars. They also consider higher-order color terms. It is important to note the difference between $u'g'r'i'z'$ and $ugriz$. This is discussed in \\cite{smith07}. Additional technical details concerning the difference between the two systems as well as transformations between them are discussed in \\cite{tucker06}. \\cite{jordi06} give color transformations between $ugriz$ as observed with the SDSS 2.5-m telescope (rather than the PT) and $UBVRI$. Additional transformations are given by \\cite{jester05}, \\cite{karaali03}, \\cite{karaali05},\\cite{bilir05}, \\cite{davenport07}, and \\cite{bilir07}. Some of the transformations including \\cite{ivezic07} consider polynomials in the color terms, but we found no need for higher-order terms for the restricted range of colors we consider. Note that the above cited transformations consider only colors, transform from $UBVRI$ to $ugriz$, are derived for the $u'g'r'i'z'$ system, or give transformations only for select Johnson-Cousins filters.  In this note, our aim has been to give a practical means of photometrically calibrating $UBVRI$ CCD images. Researchers who are interested in astrophysical applications of SDSS photometry (such as the determination of spectroscopic parallaxes or fitting theoretical isochrones to HR diagrams) are referred to the above mentioned papers because the $ugriz$ to $UBVRI$ transformations depend on the luminosity class and metallicity of the stars.  We have minimized these effects for zero-point setting by using a fairly tight color selection."
    },
    "0710/0710.5030_arXiv.txt": {
        "abstract": "Optical high-resolution  spectra of the R Coronae Borealis star V CrA at light maximum and  during minimum light are discussed. Abundance analysis confirms previous results showing that V CrA has the composition of the small subclass of R Coronae Borealis (RCB) stars know as `minority' RCBs, i.e.,  the Si/Fe and S/Fe ratios are 100 times their solar values. A notable novel result for RCBs is the detection of the 1-0 Swan system $^{12}$C$^{13}$C bandhead indicating that $^{13}$C is abundant: spectrum synthesis shows that $^{12}$C/$^{13}$C is about  3 to 4. Absorption line profiles are variable at maximum light with some lines showing evidence of splitting by about 10 km s$^{-1}$. A spectrum obtained as the star was recovering from a deep minimum shows the presence of  cool C$_2$ molecules with a rotational temperature of about  1200K, a temperature suggestive of gas in which carbon is condensing into soot. The presence of rapidly outflowing gas is shown by blue-shifted absorption components of the Na\\,{\\sc i}  D and K\\,{\\sc i} 7698 \\AA\\ resonance lines. ",
        "introduction": "R Coronae Borealis stars (here, RCBs) are a rare class of peculiar variable stars. The two defining characteristics of RCBs are (i) a propensity to fade at unpredictable times by up to about 8 magnitudes as a result of obscuration by clouds of soot, and (ii) a supergiant-like atmosphere that is very H-deficient and He-rich. The subject of this paper, V Coronae Australis (V CrA) is even a peculiar member of this class of peculiar  stars. It is a `minority' R CrB. The distinction between majority and minority members was made first by Lambert \\& Rao (1994) on the basis of an abundance analysis of  warm RCBs. The minority RCBs are quite severely deficient in iron relative to the majority RCBs and to the Sun but some elements, particularly Si and S, have near-solar abundances in the minority (and majority) RCBs. This combination results in some very unusual abundance ratios, for example, the Si/Fe and S/Fe ratios of minority RCBs are approximately 100 times the solar ratios. V CrA also seems to be an especially  lively producer of dust (Feast et al. 1997). The realization that V CrA is an unusual RCB led us to occasional spectroscopic  monitoring of its optical spectrum.  The current  paper discusses high-resolution spectroscopic observations obtained on seven occasions between 1989 and 2003  when the star was either at or near maximum light or in decline by 3 to 5 magnitudes. A suitable spectrum at maximum light is subjected here to an abundance analysis. The previous analysis (Asplund et al. 2000) was based on a spectrum obtained during a shallow light minimum. This spectrum may have been  contaminated  by  phenomena expected at minimum light (i.e., some lines may have been filled in partially by emission). In addition, our new spectra cover a broader bandpass at higher resolution and signal-to-noise ratio than the earlier spectrum. Other spectra  show  for the first time for V CrA line splitting indicative of the presence of an atmospheric shock. Spectra taken at minimum light are discussed in the context of our detailed studies of 1995-96 and 2003 minima of R CrB (Rao et al 1999; Rao, Lambert \\& Shetrone 2006).     \\begin{figure} \\epsfxsize=8truecm \\epsffile{lcvcra1.ps} \\caption{The  visual (red dots) light curve of V CrA showing the several light  minima during 1988--1998 period. Dates on which four spectroscopic observations were obtained are indicated by a dashed line (also see Table 1). Upper limits to the brightness are shown by green unfilled circles. The observations are from the AAVSO database.} \\end{figure} ",
        "conclusions": "Differences  between the spectra of RCBs at minimum light encourage us to continue our  monitoring of V CrA and other RCBs. Studies of the initial stages of a decline should reveal clues to the trigger that sets off a decline. In this era when photometric observations by amateur astronomers are reported on the internet almost instantaneously, spectroscopic and other follow-up observations are limited by access to suitable telescopes. The advent of queue scheduling is smoothing the path to obtaining the follow-up observations. Development of a consortium of observers would also ease the situation. Observations of the RCBs in the deepest of minima are likely to provide novel data on their extended atmospheres, especially on the enigmatic broad lines which, the Na D lines  apart, have been studied in detail with high-quality spectra only in the case of R CrB. For RCBs, aside from the three brightest stars R CrB, RY Sgr, and V854 Cen, a large telescope will be needed to acquire quality spectra at the faint magnitudes expected to reveal the broad lines."
    },
    "0710/0710.2426_arXiv.txt": {
        "abstract": "Sgr A* is a source of strongly variable emission in several energy bands. It is generally agreed that this emission comes from the material surrounding the black hole which is either falling in or flowing out. The activity must be driven by accretion but the character of accretion flow in this object is an open question. We suggest that the inflow is dominated by the relatively low angular momentum material originating in one of the nearby group of stars. Such material flows in directly towards the black hole up to the distance of order of ten Schwarzschild radii or less, where it hits the angular momentum barrier which leads naturally to a flow variability. We study both the analytical and the numerical solutions for the flow dynamics, and we analyze the radiation spectra in both cases using the Monte Carlo code to simulate the synchrotron, bremsstrahlung and the Compton scattering. Our model roughly reproduces the broad band spectrum of Sgr A* and its variability if we allow for a small fraction of energy to be converted to non-thermal population of electrons. It is also consistent (for a range of viewing angles) with the strong constraints on the amount of circumnuclear material imposed by the measurements of the Faraday rotation. ",
        "introduction": "%  The character of the accretion flow onto a black hole depends on the initial angular momentum of the material. This angular momentum is specified by the outer boundary conditions which depend on the relative motion of the donor with respect to the black hole. This angular momentum corresponds to a certain circularization radius, i.e. the radius where this angular momentum is equal to the local Keplerian value. In binary systems the material comes from the secondary star and in general is possess high angular momentum due to the orbital motion. In Low mass X-ray binaries the flow proceeds through an inner Lagrange point and the circularization radius is a significant fraction of a Roche radius around a black hole, of order of $10^4 R_g$ ($R_g = GM/c$). In high mass X-ray binaries the accretion flow comes from the intercepted focused wind, so the circularization radius is smaller but still large, of order of $10^3 R_g$. In such case the inflowing material form an accretion disk around a black hole and the inflow proceeds due to the angular momentum transfer. Apart from the outermost region and the region close to the ISCO (innermost stable circular orbit), the distribution of the angular momentum is relatively smooth and not much different from the Keplerian law. The exact departures from Keplerian motion depends on the disk temperature (or more exactly, on the pressure distribution).   In active galactic nuclei (AGN) the source of material is less specified. The material comes either from the stars (in the form of stellar winds) or from the gaseous phase of the galactic material. Bright AGN (quasars, Seyfert 1 galaxies) show the presence of accretion disks similar to the disks in binary systems so we can conclude that the angular momentum reaching the galactic center is high. In sources showing water maser activity we observe the outer parts of the disk directly, and in most sources the motion of the disk material is Keplerian. However, in weakly active galaxies like Sgr A* or giant elliptical galaxies we see no direct evidence of a disk. In Sgr A* the presence of the cold disk is actually excluded by the lack of eclipses of the stars which move very close to the central black hole and are systematically monitored since several years.  Since in weakly active galaxies there are no direct observational arguments for any value of the angular momentum of the donated material and the location of material sources, three types of models are being considered: \\begin{itemize} \\item high angular momentum flow, with circularization radius of order of hundreds-thousands of $R_g$ \\item low angular momentum flow, with circularization radius of order of a few  $R_g$ \\item spherical and quasi-spherical accretion, without angular momentum barrier. \\end{itemize}  The high angular momentum flow solutions for weakly active galaxies generally belong to ADAF (advection dominated accretion flow) family \\citep{1977ApJ...214..840I,1994ApJ...428L..13N}, with possibly additional effects like outflows \\citep{1999MNRAS.303L...1B} and convection. In this case the flow is not exactly Keplerian since the pressure gradients are important, but the local ratio of the angular momentum to the Keplerian angular momentum in most part of the flow is not wildly different from unity, and the angular momentum transfer (through viscosity) or angular momentum loss (through magnetic wind) at all radii is essential. Stationary solutions usually exist, and asymptotically the density of the flow approaches zero at infinity.  In spherical and quasi-spherical flow there is no angular momentum barrier so the loss of angular momentum is not the necessary condition for the accretion to occur. Examples of such solutions are: purely spherical Bondi flow or flows where the angular momentum density is below the minimum angular momentum at the circular orbit around a black hole which is given by \\begin{equation} l_{min} = 3 \\sqrt{3} GM/c \\end{equation} in case of Schwarzschild black hole; more general formula for a Kerr black hole can be found in \\citep{1972ApJ...178..347B}. In Bondi solution \\citep{1952MNRAS.112..195B,2003ApJ...591..891B} the outer boundary condition are specified by the density and the temperature of the uniform medium surrounding black hole at large distances. The flow velocity is zero at infinity, the inflow becomes transonic at the Bondi radius, and the supersonic flow reaches the black hole horizon. The Bondi radius depends significantly on the gas properties (e.g. politropic index; Bondi radius is of order of thousands of $R_g$ for relativistic flow with $\\gamma = 4/3$ but is approaches zero if $\\gamma \\rightarrow 5/3$, typical for perfect fluid non-relativistic solution), but the accretion rate is much less sensitive to those assumptions.Purely Bondi flow has generally very low radiative efficiency so it cannot reproduce the observed luminosity in most weakly active galaxies \\citep{2006A&A...450...93M}. If the accreting material at the outer boundary condition has certain angular momentum $l < l_{min}$, the dynamics of the flow is slightly modified in comparison with Bondi flow and the flow is not spherically symmetric any more but the stationary solution for the flow always exists.    The intermediate case of low angular momentum the situation is the most complex as initially the flow behaves as the Bondi flow but close to the black hole the flow starts suddenly to feel the angular momentum barrier \\citep{1981ApJ...246..314A}. In this case analytical stationary solutions frequently do not exist. In numerical solutions the flow is variable and does not reach a stationary solution in the computing time. If the angular momentum of the donated material is also a subject of changes (e.g. the result of the stellar motion), a truly stationary solution indeed can never be reached for physical reasons.  In the case of Sgr A* the available spatial resolution is the highest and we can have the best insight into the sources of material \\citep{2007IAUS..238..173G}. Therefore, in the present paper we concentrate specifically on this source and we argue that the low angular momentum flow is an interesting and promising option for the flow description. ",
        "conclusions": "%  Low angular momentum accretion flow is a promising scenario for the accretion onto Sgr A* due to its natural variability pattern. The flow is slightly more energetically efficient than the purely spherical Bondi flow and can reproduce both the required level of the luminosity and is consistent with the data on Faraday rotation measure. The overall broad band spectra are also roughly reproduced if a fraction of energy is allowed to be converted the non-thermal population of electrons. The current results are therefore encouraging, and the further work is in progress.  \\ack%  The present work was supported by the Polish Grant 1P03D~008~29 and the Polish Astroparticle Network 621/E-78/SN-0068/2007."
    },
    "0710/0710.5354_arXiv.txt": {
        "abstract": "We simulate the collisional formation of a  ring galaxy and we integrate its evolution up to 1.5 Gyr after the interaction. About $100-200$ Myr after the collision, the simulated galaxy is very similar to observed ring galaxies (e.g. Cartwheel). After this stage, the ring keeps expanding and fades. Approximately $0.5-1$ Gyr after the interaction, the disc becomes very large ($\\sim{}100$ kpc) and flat. Such extended discs have been observed only in giant low surface brightness galaxies (GLSBs). We compare various properties of our simulated galaxies (surface brightness profile, morphology, HI spectrum and rotation curve) with the observations of four well-known GLSBs (UGC6614, Malin 1, Malin 2 and NGC7589). The simulations match quite well the observations, suggesting that ring galaxies could be the progenitors of GLSBs. This result is crucial for the cold dark matter (CDM) model, as it was very difficult, so far, to explain the formation of GLSBs within the CDM scenario. ",
        "introduction": "Ring galaxies are one of the most intriguing categories of peculiar galaxies. About 280 galaxies have been classified as ring-like in the  Catalogue of Southern Peculiar Galaxies and Associations (CPGA; Arp \\& Madore 1987). They have commonly been divided in two different classes, P-type and O-type ring galaxies (Few \\& Madore 1986). The former consists of galaxies where the nucleus is often off-centre and the ring is quite knotty and irregular, while the latter includes  objects where the nucleus is central and the ring is regular. Even if for some objects such classification is ambiguous, the number of P-type and  of O-type ring galaxies are roughly comparable.  The differences among the two classes are probably connected with the formation mechanism of such galaxies: for P-type ring galaxies a collisional origin has been proposed (Lynds \\& Toomre 1976; Theys \\& Spiegel 1976; Appleton \\& Struck-Marcell 1987a, 1987b; Hernquist \\& Weil 1993; Mihos \\& Hernquist 1994;  Appleton \\& Struck-Marcell 1996; Struck 1997; Horellou \\& Combes 2001), as most of them have at least one nearby companion (Few \\& Madore 1986); whereas O-type ring galaxies can be resonant (R)S galaxies (de Vacouleurs 1959).  Simulations of galaxy collisions leading to the formation of P-type ring galaxies show that the ring phase is quite short-lived (Hernquist \\& Weil 1993; Mihos \\& Hernquist 1994; Horellou \\& Combes 2001; Mapelli et al. 2007, hereafter M07): the simulated disc galaxy develops a ring similar to the observed ones  $\\approx{}100$ Myr after the collision with the intruder; but the ring remains dense and clearly visible only for the first  $\\approx{}300$ Myr (M07).  Thus, ring galaxies are expected to rapidly evolve into something else. Up to now, only  analytic models for ring waves have been adopted to study the late stages of the ring galaxy life (Struck-Marcell \\& Lotan 1990;  Appleton \\& Struck-Marcell 1996). Neither observations nor simulations have been carried on to investigate  the fate of ring galaxies after the ring phase, leaving a lot of uncertainties. This paper aims at describing the subsequent stages of the evolution of a ring galaxy (up to $\\sim{}1.5$ Gyr), by  means of SPH/$N$-body simulations.    Our simulations suggest a possible link between old ring galaxies and the so-called giant low surface brightness galaxies (GLSBs). The GLSBs are low surface brightness galaxies (LSBs) characterized by unusually large extension of the stellar and gaseous discs (up to $\\sim{}100$ kpc; Bothun et al. 1987; Impey \\& Bothun 1989; Bothun et al. 1990; Sprayberry et al. 1995; Pickering et al. 1997; Moore \\& Parker 2007) and (often) by the presence of a  normal stellar bulge (Sprayberry et al. 1995; Pickering et al. 1997).  The existence of GLSBs has always been puzzling. Galaxy formation simulations within the cold dark matter (CDM) model have serious difficulties in producing realistic disc galaxies. In such simulations, too much angular momentum is lost during the assembly of objects, producing discs that tend to be too compact and dense. While a 'Milky Way-like' galaxy can be reproduced by high-resolution CDM simulations (provided that its merging history is fairly quiet and that heating by supernovae is properly accounted for; Governato et al. 2007), the extended discs of GLSBs require that much more angular momentum is preserved during the hierarchical build up. Such extended discs are thus beyond the reach of current galaxy formation simulations, and it is unclear whether improving the realism of such simulations will solve the problem.  Hoffman, Silk \\& Wyse (1992) proposed that  GLSBs form from rare density peaks in voids. However, most of currently known GLSBs do not appear to be connected with voids and often show interacting companions (Pickering et al. 1997). A more promising scenario is the formation of GLSBs from  massive disc galaxies due to a bar instability: a large scale bar can redistribute the disc matter and significantly increase the disc scale-length (Noguchi 2001; Mayer \\& Wadsley 2004).  In this case the redistribution of angular momentum by the bar instability could counteract the natural tendency of hierarchical assembly to remove angular momentum from the disc material (Kaufmann et al. 2007). Bar instabilities, however, normally do not increase the disc scale length by more than a factor of 2-2.5 (Debattista et al. 2006; Kaufmann et al. 2007).  In this paper we show that also the propagation of the ring in an old collisional ring galaxy can lead to the redistribution of  mass and angular momentum in both the stellar and gas component out to a distance of $\\sim{}100-150$ kpc from the centre of the galaxy, producing features (e.g. the surface brightness profile, the star formation, the HI emission spectra and the rotation curve) which are typical of GLSBs. ",
        "conclusions": "In this paper we presented a numerical model of ring galaxy evolution. About $100-200$ Myr after the collision with the intruder, the target disc galaxy evolves into a ring galaxy, similar to Cartwheel. Afterward, the ring progressively expands and fades. After $\\approx{}0.5-1.0$ Gyr the ring is no longer distinguishable from the disc, its surface density is more than 1 order of magnitude lower than in the Cartwheel phase, and the disc extends up to $\\sim{}100$ kpc.  We showed that simulated ring galaxies in the late stages of their dynamical evolution ($\\gtrsim{}500$ Myr) are very similar to the observed GLSBs. The $R$-band surface brightness profile, the SFR, the HI spectra and also the rotation curves of four GLSBs are well reproduced by the simulations. This result is unlikely due to a simple coincidence.  If all GLSBs were originated by the evolution of P-type ring galaxies, their current number density would be comparable to the observed number density of P-type ring galaxies, i.e. $\\sim{}5.4\\times{}10^{-6}\\,{}h^3\\,{}{\\rm Mpc}^{-3}$ (where $h$ is the Hubble constant; Few \\& Madore 1986). Since we know $\\sim{}17$ galaxies which can be classified as GLSBs (Sprayberry et al. 1995; Pickering et al. 1997) and which are within $\\sim{}340$ Mpc from our Galaxy, the lower limit of the current number density of GLSBs is  $\\sim{}2.7\\times{}10^{-7}\\,{}h^3\\,{}{\\rm Mpc}^{-3}$, i.e. about one order of magnitude lower than the density of ring galaxies. This can imply either that a large fraction of GLSBs have not been detected yet, or that only the $\\sim{}$5 per cent of ring galaxies ends up into a GLSB. The former scenario is quite realistic, as magnitude-limited surveys are strongly biased against GLSBs (Bothun et al. 1997). The latter hypothesis is also likely, as ring galaxies can end their life in other ways. For example, if the intruder is not sufficiently massive, the density wave is not strong enough to produce a GLSB, or, if the relative velocity is not sufficiently high, the companion can come back and merge or disrupt the propagating wave. Furthermore, recycled dwarf galaxies might also form from the debris of old collisional ring galaxies (e.g. the case of NGC~5291; Bournaud et al. 2007).   Furthermore, our model can explain the formation of GLSBs within the context of the CDM scenario,  in which very extended discs are hardly explained as a result of normal hierarchical assembly. Hence it will be quite important to check the consistency of our model with future observations. For example, the available HI data (Pickering et al. 1997) show that most of GLSBs have strong non-circular motions. From the published data it is not possible to understand whether these motions are consistent with the expansion and/or the falling back of gas in the ring. Thus, it is very important to make new radio observations of GLSBs or, at least, re-examine the archival data.  Another important point to address is how many GLSBs have a nearby companion which could be the  intruder. Among the 4 GLSBs considered here, NGC~7589 has a well-known interacting companion (Pickering et al. 1997). Furthermore, UGC~6614 is  surrounded by other galaxies with approximately the same redshift (e.g. KUG 1136+173, AGC~211143 and CGCG~097$-$034), but no studies have been done to establish whether they are in the same group as UGC6614. Finally, the surroundings of Malin~1 are populated by many companion candidates, but no redshift measurements are available at present (Sprayberry et al. 1995).  A further issue raised by this paper is whether there are objects in the intermediate stage between ring galaxies and GLSBs. The ring galaxy Arp~10 has been thought to be a relatively old ring galaxy, because its rings are not as prominent as those of other ring galaxies (e.g. Cartwheel) and because its SFR, from H$\\alpha{}$ observations, appears quite low (Charmandaris, Appleton \\& Marston 1993). However, Bizyaev, Moiseev \\& Vorobyov (2007) have shown that the SFR of Arp~10, from far-infrared observations , is $\\sim{}10-21\\,{}M_\\odot{}$ yr$^{-1}$ (analogous to the one of Cartwheel; Mayya et al. 2005), and that the collision which produced the ring occurred only $\\sim{}85$ Myr ago. Then, both current observational data and theoretical models are not sufficient to support the idea that Arp~10 is an old ring galaxy\\footnote{Even our simulations cannot solve the problem, as they indicate that the surface brightness profile of Arp~10 can be matched by a $\\sim{}80$ Myr old ring galaxy (in agreement with Bizyaev et al. 2007), as well as by a $\\sim{}300$ Myr old ring galaxy. In particular, the surface brightness profile of Arp~10 is matched by a $80$ Myr old ring galaxy with the same initial conditions as run A but with $R_d=8.8$ kpc. Nice agreement with the observations is obtained also for a $300$ Myr old ring galaxy with the same initial conditions as runs B in table 1 of M07, but with $R_d=13.2$ kpc.}.  On the other hand, one of the four GLSBs considered in this paper, UGC6614, shows (both in H$\\alpha{}$, in HI and in optical) structures which look like the remnants of the inner ring. Thus, UGC6614 could be in an intermediate, connecting stage between GLSBs and ring galaxies. It would be crucial to study more deeply UGC6614, as well as to search for other galaxies which could represent the intermediate phase between GLSBs and ring galaxies."
    },
    "0710/0710.5162_arXiv.txt": {
        "abstract": "We present recent results from a Keck study of the composition of the Galactic bulge, as well as results from the bulge Bulge Radial Velocity Assay (BRAVA). Culminating a 10 year investigation, Fulbright, McWilliam, \\& Rich (2006, 2007) solved the problem of deriving the iron abundance in the Galactic bulge, and find enhanced alpha element abundances, consistent with the earlier work of McWilliam \\& Rich (1994). We also report on a radial velocity survey of {\\sl 2MASS}-selected M giant stars in the Galactic bulge, observed with the  CTIO 4m Hydra multi-object spectrograph.  This program is to test dynamical models of the bulge and to search for and map any dynamically cold substructure in the Galactic bulge.   We show initial results on fields at $-10^{\\circ} < l <+10^{\\circ}$ and $b=-4^{\\circ}$.  We construct a longitude-velocity plot for the bulge stars and the model data, and find that contrary to previous studies, the bulge does not rotate as a solid body; from $-5^{\\circ}<l<+5^{\\circ}$ the rotation curve has a slope of $\\approx 100\\  km\\ s^{-1}$  and flattens considerably at greater $l$ and reaches a maximum rotation of $45\\ {km\\ s^{-1}}$ (heliocentric) or $\\sim 70\\ {km\\ s^{-1}}$ (Galactocentric).  This rotation is slower than that predicted by the dynamical model of Zhao (1996). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.1431_arXiv.txt": {
        "abstract": "We present the first high spatial resolution near-infrared direct and polarimetric observations of \\pars{}, obtained with the VLT/NACO instrument. We complemented these measurements with archival infrared observations, such as HST/WFPC2 imaging, HST/NICMOS polarimetry, Spitzer IRAC and MIPS photometry, Spitzer IRS spectroscopy as well as ISO photometry. Our main conclusions are the following: (1) we argue that \\pars{} is probably an FU\\,Orionis-type object; (2) \\pars{} is not associated with any rich cluster of young stars; (3) our measurements reveal a circumstellar envelope, a polar cavity and an edge-on disc; the disc seems to be geometrically flat and extends from approximately 48 to 360 AU from the star; (4) the SED can be reproduced with a simple model of a circumstellar disc and an envelope; (5) within the framework of an evolutionary sequence of FUors proposed by \\citet{green} and \\citet{quanz}, \\pars{} can be classified as an intermediate-aged object. ",
        "introduction": "The most important process of low-mass star formation is the accretion of circumstellar material onto the young star. According to current models, the mass accumulation is time-dependent with alternating episodes of high and low accretion rates \\citep{vorobyov,boley}. The high phase may correspond to FU\\,Orionis objects (FUors), a unique class of low-mass pre-main sequence stars that have undergone a major outburst in optical light of $4\\,$mag or more \\citep{herbig77}. Currently about 20 objects have been classified as FU\\,Ori-type \\citep[for a list see][]{fuors}. Records of the outbursts are not available for all of these objects: some were classified as FUors since they share many spectral properties with the FUor prototypes. According to the most widely accepted model \\citep{kh91}, the observed spectral energy distribution (SED) is reproduced by a combination of an accretion disc and an envelope. The energy of the outburst originates from the dramatic increase of the accretion rate. On the other hand, \\citet{herbig2003} favour a model in which the outburst occurs in an unstable young star rotating near breakup velocity.  \\pars{} is a system consisting of a central object, HBC\\,687 ($\\alpha_{2000}$ = 19$^{\\rm h}$ 29$^{\\rm m}$ 0\\fs87, $\\delta_{2000}$ = 9\\degr{} 38\\arcmin{} 42\\farcs{}69), and an extended nebula first listed in the catalogue of \\citet{par}. Though no optical outburst was ever observed, \\pars{} was classified as a FUor on the basis of its optical spectroscopic and far-infrared properties \\citep{sn92}. According to \\citet{henning} \\pars{} is situated in a molecular cloud named ``Cloud\\,A'', very close to the Galactic plane. Indeed, the distance of 400\\,pc and the radial velocity of $v_{lsr} = +27 \\pm 15$\\,km s$^{-1}$ \\citep{sn92} agree very well with the respective values for Cloud\\,A \\citep[$v_{lsr} = +27 \\pm 4$\\,km s$^{-1}$, $d = 500 \\pm 100$\\,pc,][]{dame85}.  Recently \\citet{quanz} questioned the FUor nature of \\pars{} referring to mid-infrared spectral properties, which better resemble those of post-AGB stars (this issue will be discussed in Sect.~\\ref{sec:discussion}).  The elongated nebulosity around the central source extends ${\\approx}\\,1\\,$arcmin to the north, while it is less developed to the south. It is visible in the optical and the near-infrared, though the source becomes unresolved at $10.8$ and $18.2\\,\\mu$m \\citep{polom}. Polarimetric observations by \\citet{draper} revealed a centrosymmetric polarization pattern and an elongated band of low polarization perpendicular to the axis of the bright nebula. \\citet{bm90} interpreted the polarization map of \\pars{} in terms of multiple scattering in flattened, optically thick structures and derived an inclination angle of 80-85\\degr{} and a size of 30 $\\times$ 8 arcsec for this disc-like structure. The existence of a disc is also supported by the discovery of a short bipolar outflow oriented along the polar axis of the nebula \\citep{sn92}. Emission at submm wavelengths was detected by \\citet{henning} and \\citet{polom}. The surroundings of \\pars{} have been searched for companions several times, but close-by stars seen in J, H and K-band images proved to be field stars \\citep{li}, and no close-by sources have been found at longer wavelengths so far \\citep{polom}.  Recently several studies were published on the circumstellar environment of FUors using high-resolution infrared and interferometric techniques (e.g.~\\citealt{green, malbet1998, malbet2005, millan-gabet, quanz}). Such observations could prove or disprove the basic assumptions of FUor models. In this paper we present the highest resolution imaging and polarimetric data yet on \\pars{} from the VLT/NACO instrument, allowing the inspection of the circumstellar material and disc at spatial scales of ${\\approx}\\,0.07$ arcsec.  The edge-on geometry of \\pars{} represent an ideal configuration to separate the components of the circumstellar environment (disc, envelope) and understand their role in the FUor phenomenon. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{\\pars{}: an FU\\,Orionis-type star}  Although no optical outburst was ever observed, \\pars{} shares many properties characteristic of FUors. Its spectral type in the literature ranges from A5 to F8 supergiant \\citep[e.g.~][]{sn92}, similar to the typical FUor spectral types (F--G supergiant). In addition, it shows strong infrared excess and drives a bipolar outflow, similarly to many FUors. However, in a recent paper by \\citet{quanz} the pre-main sequence nature and the FUor status of \\pars{} were questioned, mainly because its PAH emission bands are untypical for young stars. They suggest that \\pars{} is either an intermediate mass FUor object, or an evolved star sharing typical properties with post-AGB stars.  An important parameter which may discriminate between a pre-main sequence object and a post-AGB star is the luminosity. The typical luminosity of post-AGB stars is in the order of $10^3-10^4\\,$L$_{\\odot}$ \\citep{kwok}, while FUors are typically fainter than a few $100\\,$L$_{\\odot}$. The luminosity of \\pars{}, $10\\,$L$_{\\odot}$ (see Sect.~\\ref{sec:sed}), is at least a factor of 100 lower than that of post-AGB stars. Even adopting the largest distance estimate mentioned in the literature (1800\\,pc, \\citealt{sn92}), the luminosity of \\pars{} is still too low. The observed proper motion of the Herbig-Haro knots, however, prefers the lower distance (400\\,pc), otherwise the outflow velocities (${\\approx}\\,2400\\,$km s$^{-1}$ at 1800\\,pc) would be unusually high for a young stellar object \\citep[typically up to $600\\,$km s$^{-1}$, see e.g.~][]{mundt, hessman}. Moreover, the existence of Herbig-Haro outflows are usually tracers of low-mass star formation \\citep{hh}.  In the following discussion, we consider \\pars{} to be a FUor and we discuss its properties in the context of young eruptive stars. Nevertheless we note that most of our results on the morphology and circumstellar structure are valid regardless of the nature of the central object.   \\subsection{\\pars{}: an isolated young star?} \\label{sec:isolated}  \\begin{figure} \\centering \\includegraphics[angle=0,width=0.95\\linewidth]{kospal_fig10.ps} \\caption{2MASS colour-colour diagram of 188 sources in an area of $5.3\\,{\\times}\\,5.3$ arcmin centred on \\pars. The main-sequence and giant branch are marked by solid lines \\citep{koornneef}, the reddening path with dashed lines \\citep{cardelli} and the T\\,Tauri locus with dotted line \\citep{meyer}.} \\label{fig:cc2mass} \\end{figure}  \\begin{figure} \\centering \\includegraphics[angle=0,width=0.95\\linewidth]{kospal_fig11.ps} \\caption{IRAC colour-colour diagram of 100 sources located in the same area as in Fig.~\\ref{fig:cc2mass}. The gray square in the upper right corner marks the approximate domain of Class II sources \\citep{allen}.} \\label{fig:ccirac} \\end{figure}  \\begin{table} \\caption{VLT/NACO photometry for the two closest stars. Uncertainties are about 0.1 mag.} \\label{tab:twostars} \\centering \\begin{tabular}{c c c c c} \\hline $\\alpha_{2000}$                       & $\\delta_{2000}$                  & H mag & K$_S$ mag & L$^{\\prime}$ mag \\\\ \\hline 19$^{\\rm h}$ 29$^{\\rm m}$ 0\\fs96      & 9\\degr{} 38\\arcmin{} 42\\farcs11  & 20.2  & 19.1      & $>$14.6          \\\\ 19$^{\\rm h}$ 29$^{\\rm m}$ 0\\fs70      & 9\\degr{} 38\\arcmin{} 44\\farcs78  & 20.6  & 19.5      & $>$14.4          \\\\ \\hline \\end{tabular} \\end{table}  \\pars{} is situated close to the Galactic plane ($l=45.8^{\\circ}$, $b=-3.8^{\\circ}$), in a molecular cloud called Cloud\\,A. This cloud was identified by \\citet{dame85} in their CO survey of molecular clouds in the northern Milky Way. The cloud occupies an area of 8 square degrees in the sky, and the only known young star associated with it is the T\\,Tauri star AS\\,353 \\citep{dame85}.  In order to check whether FUors are usually associated with star forming regions, we searched the literature and found that most FUors are located in areas of active star formation \\citep[e.g.][]{henning}. To find out whether there is star formation in the vicinity of \\pars{}, we searched for pre-main sequence stars. For this purpose we constructed a 2MASS J$-$H vs.~H$-$K$_S$ and an IRAC [3.6]$-$[4.5] vs.~[5.8]$-$[8.0] colour-colour diagram for sources found in our $5\\farcm3\\,{\\times}\\,5\\farcm3$ IRAC field of view (Fig.~\\ref{fig:cc2mass}, \\ref{fig:ccirac}). Our selection criteria in the case of 2MASS was S/N ${>}\\,10$ and uncertainties ${<}\\,0.1\\,$mag in all J, H and K$_S$ bands, while in the case of IRAC S/N ${>}\\,3$ and detectability at all four bands were required.  The 2MASS diagram revealed that most of the nearby objects are reddened main sequence or giant stars. On the IRAC diagram, however, there are three objects (apart from \\pars{} itself), which display infrared excess at $8\\,\\mu$m (marked by A, B and C in Fig.~\\ref{fig:ccirac}). According to the classification of \\citet{allen}, Class II sources exhibit colours of [3.6]$-$[4.5] ${>}\\,0.0$ and [5.8]$-$[8.0] ${>}\\,0.4$, thus one of these stars (A) might be a Class II source, while B and C are more likely Class III/main sequence sources. The nature of source A and its possible relationship to \\pars{} is yet to be investigated. Nevertheless, \\pars{} seems to be rather isolated compared to most FUors and certainly not associated with any rich cluster of young stellar objects.  We also searched for possible close companions of \\pars{} in the WFPC2 and NACO direct images. In order to establish a detection limit for source detection, we measured the sky brightness on the NACO images (before sky-subtraction), and estimated a limiting magnitude for each filter. The resulting values are $22.8$, $21.6$, and $15.2\\,$mag in H, K$_S$ and L$^{\\prime}$, respectively. In case of the HST/WFPC2 image, the larger ($80\\,{\\times}\\,80$ arcsec) field of view made it possible to estimate a limiting magnitude using star counts; the resulting value is $23.5\\,$mag. Due to the bright reflection nebula, the detection limit is somewhat lower close to the star. The two closest objects we found are the following: one star to the southeast, at a distance of $1.4$ arcsec (560\\,AU at 400\\,pc), and another one to the northwest, at a distance of $3.3$ arcsec (1320\\,AU at 400\\,pc). These sources are marked with arrows in Fig.~\\ref{fig:images}. Neither of the stars are visible at $3.8$ or $0.814\\,\\mu$m, although we can give an upper limit for their L$^{\\prime}$ brightness. Their positions and photometry are given in Table~\\ref{tab:twostars}. As these sources are very red, they can equally be heavily reddened background stars, or stars with infrared excess (indicating that they might be associated with \\pars{}). Supposing that they are reddened main sequence stars, one can estimate an extinction of A$_V\\,{\\approx}\\,10\\,{-}\\,15\\,$mag. Further multifilter observations may help to clarify the nature of these objects and their possible relationship to \\pars{}.   \\subsection{The circumstellar environment of \\pars{}}  \\begin{figure} \\centering \\includegraphics[angle=0,width=0.45\\linewidth]{kospal_fig12.ps} \\caption{Sketch of the morphology of circumstellar material around \\pars{}, overlaid on the HST/WFPC2 image. The central star is surrounded by an edge-on disc. Perpendicular to the disc, the star drives a bipolar outflow that excavates an outflow cavity in the dense circumstellar material. Light from the central star illuminates the walls of the cavity.} \\label{fig:morph} \\end{figure}  \\begin{figure} \\centering \\includegraphics[angle=90,width=0.95\\linewidth]{kospal_fig13.ps} \\caption{Brightness profiles of \\pars{} at $0.8\\,\\mu$m and in the H and K$_S$ bands. Dotted lines mark Hubble's relation.} \\label{fig:metszet} \\end{figure}  \\begin{figure} \\centering \\includegraphics[angle=0,width=0.95\\linewidth]{kospal_fig14.ps} \\caption{South-north cut at $0.6$ arcsec east from the star. {\\it Solid line:} total intensity; {\\it dashed line:} polarized intensity.} \\label{fig:cutok} \\end{figure}  \\begin{figure*} \\centering \\includegraphics[angle=90,width=0.95\\linewidth]{kospal_fig15.ps} \\caption{VLT/NACO polarization map of \\pars{} overlaid on H-band total intensity contours. The circle in the upper right corner displays the FWHM of the polarization measurement. Polarization vectors are displayed at full resolution, only showing the central low polarization band, where the polarization vectors are aligned. Such arrangement is expected when multiple scattering occurs in an edge-on disc.} \\label{fig:kiskep} \\end{figure*}  The appearance and polarization properties of the nebula around \\pars{} can be understood in the following way: the star drives an approximately north-south oriented bipolar outflow, which had excavated a conical cavity in the dense circumstellar material (Fig.~\\ref{fig:morph}). The star illuminates this cavity and the light is scattered towards us mainly from the walls of the cavity. The outflow direction is perpendicular to an almost edge-on dense circumstellar disc. This picture is supported by the following facts: (a) the centrosymmetric polarization pattern is characteristic of reflection nebulae with single scattering; (b) the morphology and limb brightening suggest a hollow cavity (as opposed to an ``outflow nebula'', where the lobes are composed of dense material ejected by the central source); and (c) the low-polarization lane across the star strongly suggests the presence of an edge-on circumstellar disc, where multiple scattering occurs. In general the \\pars{} system shows similarities to the NGC\\,2261 nebula associated with R\\,Mon. This object also consists of a northern cometary nebula and a southern jetlike feature \\citep{warren}.   \\subsubsection{Envelope/Cavity} \\label{sec:env}  We characterised the opening of the upper lobe by marking the ridge along the northeastern and northwestern arcs which we interpret as the walls of the cavity. As viewed from the star northwards, the cavity starts as a cone with an opening angle of ${\\approx}\\,60^{\\circ}$, giving the nebula in Fig.~\\ref{fig:images} a characteristic equilateral triangle-shape. Farther away from the star the cavity deviates from the conical shape, becomes narrower. The whole cavity occupies an area of $8\\,000\\,{\\times}\\,24\\,000\\,$AU (at a distance of $400\\,$pc). The sharp outer boundary of the nebula implies a significant density contrast between the cavity and the surrounding envelope.  In Fig.~\\ref{fig:metszet} we plotted radial brightness profiles at $0.8\\,\\mu$m and in the H and K$_S$ bands, starting from the star northwards. The slope of the intensity profiles in the inner ${\\sim}\\,1.6''$ ($640\\,$AU) follows closely Hubble's relation \\citep{hubble}, i.e.~the brightness is proportional to r$^{-2}$. This trend can be clearly followed above the noise out to about $5\\,$arcsec in our H and K$_S$ images, giving an estimate of $2000\\,$AU for the outer size of the envelope. The fact that the brightness profiles follow Hubble's relation implies that the nebula is produced by isotropic single scattering, in accordance with the centrosymmetric polarization pattern and the high degree of polarization. Around ${\\sim}\\,1.6''$ the profiles becomes steeper, probably due to decreased density in the inside of the cavity. Similar steepening in the outer part of the nebula was already mentioned by \\citet{li}. The profiles are similar at all observed wavelengths and do not show significant dependence on the position angle.   \\subsubsection{Disc} \\label{sec:disc}  As mentioned in Sect.~\\ref{sec:pol}, both the NACO and NICMOS polarimetric images show a lane of low polarization oriented nearly east-west across the star (Fig.~\\ref{fig:pol}). The drop in the degree of polarization can also be clearly seen in Fig.~\\ref{fig:cutok}, where a north-south cut at $0\\farcs6$ east to the star is plotted. Following \\citet{bm90}, we interpret the low polarization by multiple scattering in an edge-on disc (possible other explanations for the origin of the low polarization areas are discussed in e.g.~\\citealt{lr}). According to models of such circumstellar structures \\citep[e.g.][]{whitney, fischer} the polarization vectors are oriented parallel to the disc plane. As can be seen in Fig.~\\ref{fig:kiskep}, despite the low degree of polarization, the predicted alignment of the vectors can be clearly seen in the case of \\pars{} too. The NICMOS polarization map shows a similar effect. The dark lane in Fig.~\\ref{fig:pol} can be followed inwards to as close as $48\\,$AU from the star. This is an upper limit for the inner radius of the circumstellar disc.  An interesting feature of the mentioned models is two depolarized areas on either side of the central star, which mark the outer end points of an edge-on disc. The depolarization is due to a transition from the linearly aligned to the centro-symmetic polarization pattern. These depolarized areas can be seen in Fig.~\\ref{fig:vector} bottom, on both sides at about $0.9''$ from the star. At the distance of \\pars{} this corresponds to $360\\,$AU, and can be adopted as the outer radius of the dense part of the disc, where multiple scattering at near-infrared wavelengths is dominant. It is interesting that the NACO map does not show clear depolarized areas at the same position, but seems to resolve the transition in vector orientation, while the larger beam of NICMOS averaged the differently oriented vectors, resulting in depolarized spots.  The polarized intensity and degree of polarization images in Fig.~\\ref{fig:pol} suggest a slight asymmetry in the dense disc: the eastern (left) side is straight, while the western (right) side shows a kink and also has a different position angle than that on the other side. The thickness of the disc can be measured on the area where the polarization vectors are aligned (Fig.~\\ref{fig:kiskep}). This approximately corresponds to the area where the degree of polarization is below ${\\approx}\\,10\\%$. The resulting thickness is approximately $0.1''$ ($40\\,$AU) to the east and is somewhat larger, $0.2''$ ($80\\,$AU), to the west. One should note that, since these values are close to the spatial resolution of the polarimetric images, these numbers should be considered as upper limits for the thickness of the circumstellar disc around \\pars{}. They are upper limits also because if the inclination is not exactly 90$^{\\circ}$, the thickness of the disc can be even less. The thickness does not show significant increase with radial distance, suggesting the picture of a flat, rather than a flared disc, at least considering the dense, multiple-scattering part. The smooth brightness and polarization distribution between the disc and the surrounding envelope (Fig.~\\ref{fig:cutok}), however, implies that there is a continuous density transition between the two components. From the ratio of the horizontal to vertical sizes of the disc a lower limit of $84^{\\circ}$ for the inclination of the system (the angle between the normal of the disc and the line of sight) can be derived.  \\citet{fischer} computed a grid of polarization maps of young stellar objects with the aim of helping the interpretation of polarimetric imaging observations. They consider five different models, four with massive, self-gravitating discs and one with a massless Keplerian disc. Since the mass of circumstellar material of \\pars{} derived from submillimetre observations is relatively low (${<}\\,0.3\\,$M$_{\\odot}$, \\citealt{henning, polom, sw, hillen}), the most appropriate model for our case is a Keplerian disc. Indeed the polarization pattern as computed by \\citet{fischer} for an inclination of 87$^{\\circ}$ (their Fig.~1) looks remarkably similar to our Fig.~\\ref{fig:vector} (the differences might be explained by the narrower cavity and flatter disc of \\pars{}). Thus, the geometry and structure assumed by \\citet{fischer} in their Keplerian model could be a good starting point for further radiative transfer modelling of \\pars{}.   \\subsection{Modelling the circumstellar environment} \\label{sec:modelling}  \\begin{table} \\centering \\caption{Model parameters. Parameters in italics are fixed, while the others were fitted.} \\label{tab:par} \\begin{tabular}{lcc} \\hline Parameter & Variable & Value \\\\ \\hline Inner disc radius                     & $R_1$         & 3.5\\,$R_\\odot$ \\\\ {\\it Outer disc radius}               & $R_2$         & {\\it 360\\,AU}  \\\\ Temperature at 1 AU                   & $T_{\\rm d,0}$ & 285\\,K  \\\\ {\\it Power-law index for temperature} & $q_{\\rm d}$   & {\\it 0.75}  \\\\ Power-law index for surface density   & $p_{\\rm d}$   & 1.6 \\\\ {\\it Disc mass}                       & $M_{\\rm d}$   & {\\it 0.02\\,M$_\\odot$} \\\\ {\\it Inclination}                     & $i$           & {\\it 86$^\\circ$} \\\\ Inner envelope radius                 & $R_3$         & 5.4\\,AU \\\\ {\\it Outer envelope radius}           & $R_4$         & {\\it 2000\\,AU} \\\\ Temperature at 5 AU                   & $T_{\\rm e,0}$ & 368\\,K \\\\ {\\it Power-law index for temperature} & $q_{\\rm e}$   & {\\it 0.4} \\\\ Power-law index for surface density   & $p_{\\rm e}$   & 0.4 \\\\ Envelope mass                         & $M_{\\rm e}$   & 0.02\\,M$_\\odot$ \\\\ {\\it Interstellar Extinction}         & $A_{V}$       & {\\it 2\\,mag}\\\\ \\hline \\end{tabular} \\end{table}  Our observations provide some direct measurements of the geometry of the circumstellar structure (disc size and thickness, inclination, envelope size). In the following we discuss the consistency of this picture with the observed SED. Our approach is to construct a simple disc+envelope model, in which we fix those parameters whose values are known from our NACO observations or from other sources (outer disc radius from this work; power-law index for disc temperature from \\citealt{shakura}; disc mass was set in order to ensure that the whole disc is optically thick; inclination from this work; outer envelope radius from this work, the power-law index for envelope temperature is a typical value for optically thin envelopes containing larger than interstellar grains, e.g.~\\citealt{hartmannkonyv}, Eqn.~4.13; interstellar extinction from \\citealt{hillen}). Then we check whether the SED can be fitted by tuning the remaining parameters. We adopted an analytical disk model \\citep{adams}, which has been successfully used to model FUors \\citep{quanz2, v1647ori_midi}. Our model consists of two components, an optically thick and geometrically thin accretion disc \\citep{shakura} and an optically thin envelope (no cavity is assumed). No central star is included in the simulation, partly because in outbursting FUors the star's contribution is negligible compared to that of the inner disc \\citep{hk96}, and partly because of the edge-on geometry where the star is obscured by the disc. This assumption is supported by the fact that the shape of the SED at optical wavelengths is broader than a stellar photosphere. The model also does not take into account internal extinction and light scattering, thus it cannot reproduce any of the near-IR imaging and polarimetric observations.  The temperature and surface density distribution in the disc are described by power-laws: \\begin{equation} T(r)=T_{{\\rm d}, 0}\\left(\\frac{r}{1\\,{\\rm AU}}\\right)^{-q_{\\rm d}}, \\end{equation} \\begin{equation} \\Sigma(r)=\\Sigma_{{\\rm d},0}\\left(\\frac{r}{1\\,{\\rm AU}}\\right)^{-p_{\\rm d}}. \\end{equation} Similar power-laws were assumed for the envelope. The observed flux at a specific frequency is given by \\begin{eqnarray} F_\\nu &=& \\frac{\\cos{i}}{D^2}\\int_{\\rm R1}^{\\rm R2}2\\pi r(1-e^{\\frac{-\\Sigma_{{\\rm d}}\\kappa_\\nu}{\\cos{i}}})B_\\nu(T_{\\rm d}){\\rm d}r + \\\\ & & \\frac{1}{D^2}\\int_{\\rm R3}^{\\rm R4}2\\pi r(1-e^{-\\Sigma_{{\\rm e}}\\kappa_\\nu})B_\\nu(T_{{\\rm e}}){\\rm d}r. \\end{eqnarray} The first term describes the emission of the accretion disc, the second term describes the radiation of the optically thin envelope. For the dust opacity we used a constant value of $\\kappa_{\\nu}\\,{=}$ 1 cm$^2$g$^{-1}$ at $\\lambda\\,{>}\\,1300\\,\\mu$m, $\\kappa_{\\nu}\\,{=}\\,\\kappa_{1300\\mu\\rm m}\\left(\\frac{\\lambda}{1300\\mu\\rm m}\\right)^{-1}$ between $1300$ and $100\\,\\mu$m and again a constant value of $\\kappa_{\\nu}\\,{=}\\,\\kappa(100\\,\\mu\\rm m)$ at $\\lambda\\,{<}\\,100\\,\\mu$m.  We fitted the SED via $\\chi^2$ minimisation using a genetic optimization algorithm PIKAIA \\citep{charbonneau}. This algorithm performs the maximization of a user defined function, for which purpose we used the inverse $\\chi^2$. Since there are many photometric measurements in the mid-infrared domain, but just a few in the far-infrared, the mid-infrared region has a higher weight during the fit, compared to the far-infrared domain. Therefore, in order to ensure an equally good fit at all wavelengths, we divided the SED into four regions and weighted the $\\chi^2$ of each domain with the inverse of number of photmetric points the region contained. Then the final $\\chi^2$ was the sum of the $\\chi^2$ of all regions. The regions we used were: $0.3-3$, $3-30$, $30-300$ and $300-3000\\,\\mu$m. The parameters of the best-fit model, which gives a weighted $\\chi^2$ of 0.67, are listed in Table~\\ref{tab:par}.  The fitted model SED as well as the disc and envelope components are overplotted in Fig.~\\ref{fig:sed}.  The model SED is consistent with the observed fluxes. This shows that the picture of a thin accretion disc and an envelope is consistent with both the measured SED and the geometry and disc/envelope parameters inferred from our polarimetric observations. Detailed modelling of the silicate, PAH, and ice spectral features, as well as the correct treatment of internal extinction and scattering would require radiative transfer modelling, which will be the topic of a subsequent paper.  \\subsection{The evolutionary status of \\pars{}} \\label{sec:general}  The geometry of FUor models discussed in the literature (e.g.~\\citealt{hk96, tbb}) usually consist of a central star surrounded by an accretion disc and an infalling envelope with a wind-driven polar hole. These assumptions are supported by the fact that they fit well the SED \\citep{green, quanz}, the interferometric visibilities \\citep{mg,v1647ori_midi} and the temporal evolution of the SED \\citep{fuors}. In this paper we present the first direct imaging of these circumstellar structures in a FUor. Our polarimetric measurements of \\pars{} show the existence of a circumstellar disc which extends from at least 48 to 360 AU. The most striking feature of the disc is its flatness over the whole observed range. The short-wavelength part of the SED could be well reproduced using a radial temperature profile of $r^{-0.75}$ (Sect.~\\ref{sec:modelling}). This profile is expected from both a geometrically thin accretion disc and a flat reprocessing disc.  An envelope was also seen in the polarization maps of \\pars{} and it was also a necessary component for the SED modelling. Envelopes are involved in many FUor models and in this paper we present a direct detection of this model component. Our images reveal that the envelope can be followed inwards as close to the star as the disc. FUor models often assume a polar cavity in the envelope, created by a strong outflow or disc wind. The direct images of \\pars{} clearly show the presence of such a cavity and we also detected a bipolar outflow in the \\pars{} system.  In the recent years, as new interferometric and infrared spectroscopic observations were published for FUors, the group turned out to be more inhomogeneous in physical properties than earlier assumed, when mainly optical photometry and spectroscopy had been available. \\citet{quanz} proposed that some differences might be understood as an evolutionary sequence. They suggest that FUors constitute the link between embedded Class I objects and the more evolved Class II objects. Members of the group exhibiting silicate absorption at $10\\,\\mu$m are younger and more embedded (Category 1, e.g.~V346\\,Nor); while objects with pure silicate emission are more evolved (Category 2, e.g.~FU\\,Ori and Bran\\,76). There are objects showing a superposition of silicate absorption and emission, which are probably in an intermediary evolutionary stage (e.g.~RNO\\,1B). \\citet{green} also sorted FUors, based on the ratio of the far-infrared excess and the luminosity of the central accretion disc, $f_d$ (Equ.~7 in their paper). A large relative excess ($f_d\\,{>}\\,5\\%$) indicates an envelope of large covering fraction (V1057\\,Cyg and V1515\\,Cyg), while low relative excess means a tenuous or completely missing envelope (Bran\\,76 and FU\\,Ori). This is also an evolutionary sequence, as young, more embedded objects have large envelopes, while around more evolved stars, the envelope has already dispersed. $f_d$ can also be used to calculate the opening angle of the envelope, thus a prediction of this scheme is that the opening angle is becoming wider during the evolution, probably due to strong outflows during the repeated FUor outbursts.  The two classification schemes are not inconsistent and one can merge them into the following evolutionary sequence: (1) the {\\it youngest objects} exhibit silicate absorption and large far-infrared excess (V346\\,Nor, probably also OO\\,Ser and L1551\\,IRS\\,5 belong here); (2) {\\it intermediate-aged objects}, where the silicate feature is already in emission but there is still a significant far-infrared excess (V1057\\,Cyg, V1515\\,Cyg, probably also RNO\\,1B and V1647\\,Ori); (3) the most {\\it evolved objects} show pure silicate emission and low far-infrared excess (FU\\,Ori, Bran 76). We note, however, that this classification has some weak points. As \\citet{quanz} already mentioned, an edge-on geometry in a more evolved system may appear as a younger one. Moreover, during an outburst and the subsequent fading phase, certain spectral features as well as the global shape of the SED may change.  \\pars{} can be placed in this evolutionary scheme, though one should keep in mind that because of the nearly edge-on geometry, the classification of this object is somewhat uncertain. \\pars{} displays silicate emission (Fig.~\\ref{fig:pah}). Integrating the flux of the two components in our simple model (Sect.~\\ref{sec:modelling}), and correcting the apparent disc luminosity for inclination effect ($i=86^{\\circ}$, Table~\\ref{tab:par}) using Equ.~6 of \\citet{green}, we obtained a large relative far infrared excess of $f_d\\,{=}\\,75\\%$. These two properties place \\pars{} into the intermediate-aged category (though because of its inclination, it may actually seem younger than it is). Following Equ.~7 of \\citet{green}, from the $f_d$ value, we also computed the opening angle of the envelope. The resulting opening angle of $60^{\\circ}$ agrees well with the angle measured in the direct NACO images (Sect.~\\ref{sec:env}).  \\pars{} was placed into the evolutionary scheme using two parameters: the silicate feature and the relative far infrared excess. In the following we discuss whether its other physical characteristics match with those of other FUors. \\begin{itemize} \\item[{\\it (i)}] Our observations revealed that the circumstellar disc of \\pars{} is very flat. Due to lack of similar direct measurements for other FUors, we can only speculate that perhaps all FUors with envelopes have such flat discs. On the other hand, the most evolved FUor, FU\\,Ori, seems to have no envelope but its disc is probably flared \\citep{kh91, green, quanz2}. This might suggest that disc flaring develops at later stages, when illumination from the central source may heat the disc surface more directly. \\item[{\\it (ii)}] In a nearly edge-on system like \\pars{}, one expects to see the $10\\,\\mu$m silicate feature in absorption. The fact that \\pars{} has silicate emission indicates that the line of sight towards the central region is not completely obscured. Using the optical depth of the $15.2\\,\\mu$m CO$_2$ ice feature, we calculated an $A_V\\,{=}\\,8\\,$mag ($A_V\\,{=}\\, 38.7 A_{15.2 \\mu\\rm{}m}$, \\citealt{savage}). This value is surprisingly low compared to V1057\\,Cyg ($A_V\\,{\\sim}\\,50-100\\,$mag, \\citealt{kh91}). This indicates a much more tenuous envelope, which is also supported by the low envelope mass of $0.02\\,\\rm{}M_{\\odot}$ in our modelling. \\item[{\\it (iii)}] Following \\citet{quanz}, we analysed the profile of the $15.2\\,\\mu$m CO$_2$ ice feature of \\pars{}. The inset in Fig.~\\ref{fig:pah} shows that the feature has a characteristic double-peaked sub-structure, very similar to HH\\,46\\,IRS, an embedded young source \\citep{boogert}. HH\\,46\\,IRS is a reference case for processed ice. The presence of processed ice in \\pars{} indicates heating processes and the segregation of CO$_2$ and H$_2$O ice, already at this evolutionary stage. Other FUors exhibiting this kind of profile are L1551\\,IRS\\,5, RNO\\,1B and RNO\\,1C \\citep{quanz}. \\end{itemize}  The evolutionary state of a young stellar object can also be estimated following the method proposed by \\citet{chen}. According to their Equ.~(1) we calculated a bolometric temperature of T$_{\\rm bol}=410\\,$K for the measured SED. We compared this value with the distribution of corresponding values among young stellar objects in the Taurus and $\\rho\\,$Ophiuchus star forming regions (Chen et al. 1995). From this check we can conclude that \\pars{} seems to be a class I object, and its age is ${\\sim}\\,10^5\\,$yr. However, \\citet{green} argued that the apparent SED of the disc component depends on the inclination. Thus we computed T$_{\\rm bol}$ also for a face-on disc configuration and obtained T$_{\\rm bol}=1160\\,$K, corresponding to a Class II object. In fact, \\pars{} is probably close to the Class I / Class II border, in accordance with the proposal of \\citet{quanz}."
    },
    "0710/0710.3328_arXiv.txt": {
        "abstract": "We analyze the electromotive force (EMF) terms and basic assumptions of the linear and nonlinear dynamo theories in our three-dimensional (3D) numerical model of the Parker instability with cosmic rays and shear in a galactic disk. We also apply the well known prescriptions of the EMF obtained by the nonlinear dynamo theory (Blackman \\& Field 2002 and Kleeorin et al. 2003) to check if the EMF reconstructed from their prescriptions corresponds to the EMF obtained directly from our numerical models. We show that our modeled EMF is fully nonlinear and it is not possible to apply any of the considered nonlinear dynamo approximations due to the fact that the conditions for the scale separation are not fulfilled. ",
        "introduction": "\\label{sec:intro}  It seems that the issue of the magnetic field amplification in galaxies may be well explained  by the two main physical mechanisms: the Parker instability (PI), which takes into account the cosmic rays (CR) and the shear \\cite[e.g.][and references therein]{hanasz03,hanasz04}, and the magneto-rotational instability \\cite[MRI, e.g.][]{dziourkevitch04,kitchatinov04}. The possible scenario of the magnetic field evolution could be presented as follows: when the protogalaxy starts to rotate differentially, the MRI mechanism occurs, and this results in a very efficient magnetic field amplification even to the level of $\\mu$G with the global e-folding time of 100~Myr or even less \\citep{dziourkevitch04}. Simultaneously, the quadrupole symmetry of the large-scale magnetic field is being created. MRI also causes the turbulent motions in the galactic disks. The process of supernovae (SN) explosions, that arises in young galactic objects \\cite[e.g.][and references therein]{widrow02} suppresses the MRI mechanism. In the same time, the cosmic rays produced in SN remnants may induce the Parker instability process \\cite[e.g.][]{parker92,hanasz04}. Hence, we may conclude that during the early stage of galaxy evolution the MRI process is being replaced by the PI mechanism.  The local simulations of the large-scale magnetic field took into account the turbulent dynamo theory and the  magnetic back-reaction onto the turbulent motions \\citep{piddington70,piddington72a,piddington75a,kulsrud95}. The authors drew the conclusion that it was difficult to obtain the amplification of the total magnetic field \\cite[e.g.][and references therein]{widrow02}. It might be explained by the equipartition of the random magnetic field component with the random turbulent motions. Such process suppresses the dynamo action at later times. Moreover, the magnetic field amplification might be easily stopped even by the presence of the weak large-scale magnetic field \\citep{cattaneo91,vainshtein92,cattaneo94,cattaneo96,ziegler96}. \\cite{blackman99} explained that papers analyzing the analytically strong suppression of the dynamo coefficient $\\alpha$ should distinguish between different state orders of the turbulent quantities. However, their analysis did not completely solve the problem of the quenching of the dynamo coefficients. In their next paper \\citep{blackman00}, they proved that the results from the \\cite{cattaneo96} model were based on the assumption about the periodicity of the boundary conditions. Nevertheless, the quenching effects could also appear even when the open boundary conditions were applied \\cite[e.g.][]{brandenburg01b}.  The classical dynamo theory does not conserve the total magnetic helicity \\cite[see e.g.][BF02]{blackman02}. In media characterized by the high magnetic Reynolds number ($R_m\\gg1$) the total helicity should remain constant in closed regions \\cite[e.g.][]{berger84,brandenburg02b,brandenburg05b,subramanian02}. The permanent helicity is also an additional factor, which suppresses the dynamo activity \\citep[e.g.][]{brandenburg02b}. The following papers \\citep{blackman00,kleeorin00,blackman03,kleeorin99,kleeorin02,kleeorin00, kleeorin03,rogachevskii00,rogachevskii01} presented the two methods that allowed the modeling of the dynamo action evading the problem of the constant helicity. The first method is based on the ejection of the magnetic helicity through boundaries. The second one uses the creation of the negative and positive helicity at the large and small scales respectively \\cite[see][]{brandenburg02b,kleeorin02,kleeorin03}. \\cite{blackman02} in their next paper analyzed the nonlinear prescription of both dynamo coefficient $\\alpha$ and $\\beta$. Both factors were obtained without any linearization and took into account all terms in the equation of the evolution of the fluctuating part of the magnetic field (see Eq.~\\ref{eqn:fluct_field_evol}). The results were similar to the dynamo coefficients obtained by \\citep{pouquet76}, but the units were different (without the time integration). They also solved the small- and large-scale helicity dynamo equations numerically simultaneously with the equation for the EMF time evolution. That allowed them to obtain the growth of the large-scale magnetic field in the kinematic phase.  The new form of the dynamo coefficients for anisotropic turbulent motions with the presence of the large-scale magnetic field was presented by \\cite{rogachevskii01}. They calculated dynamo coefficients according to the \\cite{raedler80} EMF prescription \\cite[see also][]{kowal05}, which neglected all quadratic  terms in the mean field in the EMF. \\cite{kleeorin03} used those forms of the dynamo coefficients to solve numerically the dynamo equation in the local thin-disc approximation. The authors took into account the quenching of both coefficients, $\\alpha$ and $\\beta$, and helicity flux through the boundary. They found that it was  possible to obtain the growth of the large-scale magnetic field when $\\alpha$-quenching was only analyzed \\citep{kleeorin02}. If the model included also $\\beta$-quenching, no growing solution of the dynamo \\citep{kleeorin03} could be obtained.  Both in Blackman-Field and in Kleeorin-Rogachevskii approaches there is an $\\alpha_{\\rm m}$ term that quantifies the small scale helicity current. This term depends on the current helicity flux, and there are different theories for this flux. In Kleeorin et al. a heuristically motivated expression for the flux was used, in \\cite{brandenburg05b}  the \\cite{vishniac01} flux was used, and in \\cite{shukurov06} a simple advective flux was used. In all these cases the current helicity flux allows the field to saturate at high levels.  The latest research on the turbulent enforcement in the solar convective zone calculated in the local cube with the shear has shown that the open boundaries help to obtain the amplification of the large-scale magnetic field even without the helicity of turbulent motions \\cite[end references therein]{brandenburg05a,brandenburg05b}. We have to stress that their result, an increase of the total magnetic energy, was obtained without any assumption considering the additional EMF of dynamo. The authors applied isotropic and homogeneous turbulence with and without the helical forcing in their model. On the other hand, Brandenburg and his collaborators \\citep{brandenburg05a} interpreted their results in terms of the mean field dynamo theory. The time evolution of $\\alpha$ was obtained from the calculated electromotive force \\citep{brandenburg04,kowal05,brandenburg02a,brandenburg01a}. \\cite{brandenburg04} explained that thanks to the flux of the current helicity flowing out of the cube the process of $\\alpha$-quenching tends to be not as disastrous as forseen. The only conditions are the intermediate level of $R_m$ and open boundaries. When the high value of $R_m$ is applied to the model, the comparatively lower value of the large-scale magnetic field strength are obtained \\citep{brandenburg04}. However, in astrophysical objects $R_m$ is always high. That is why this result seems to be peculiar. On the other hand, the authors made it clear that the total magnetic energy grows mainly in the kinematic phase of the dynamo and their results do not depend strongly on $R_m$. Futhermore, they calculated the value of $\\alpha$ coefficient based on the modeled EMF. The obtained factor was similar to the same coefficient calculated according to the \\cite{kleeorin00,kleeorin02,kleeorin03} prescriptions. We believe that the fact that such similarity occurs results from the isotropic and homogeneous turbulence in both models \\citep{brandenburg05a,brandenburg05b,brandenburg04}. It may also happen due to the fact that the authors applied standard dynamo approximations, which neglect the quadratic terms in the mean field in EMF.  The previously mentioned results indicate that the realistic physical simulations are of the great importance when the MRI \\citep{dziourkevitch04} and the PI \\citep{hanasz04,hanasz03,kowal05} processes are considered. Both models, which meet enumerated requirements, showed that it was possible to amplify galactic magnetic field efficiently. \\cite{hanasz03} and \\cite{hanasz04,hanasz05} presented that the following two processes: the Parker instability driven by cosmic rays from supernovae and the shear from the differential rotation enable the magnetic field amplification (with the e-folding time scale of 250~Myr or even 140~Myr). The model also applied realistic gravity according to the \\cite{ferriere98} prescriptions.  The idea of obtaining the dynamo coefficients ($\\alpha$ and $\\beta$) from the electromotive force calculated in the local numerical simulations proved to be essential for many other authors too \\cite[e.g.][etc]{ziegler96,brandenburg02a}. We included the calculations of the dynamo coefficients, which we obtained from the calculated EMF in our previous paper \\citep{kowal05}. The application of the statistical methods provided us with the acceptable values of the dynamo $\\alpha$-tensor. On the other hand, the values of $\\beta$ coefficients were negative. Such values are inconsistent with the R\\\"adler prescription \\citep{raedler80}. This may be caused either by the applied statistical method, which does not take into account physical differentiation, or by the linear EMF approximation \\citep{kowal05}. For this reason we decided to analyze that matter in our present study. We search for the conditions, which should be fullfiled in order to make linearization of the electromotive force in the mean field dynamo theory possible. We would like to examine the following problems: the scale separation, the ratios of the terms in the equation for the small-scale magnetic field evolution \\cite[e.g.][]{raedler80}, the magnitude of the turbulent kinetic energy in comparison to the large-scale magnetic one. Next, we plan to apply the estimations of the dynamo coefficients derived by \\cite{blackman00,rogachevskii01} and \\cite{brandenburg04} to our models. We would like to check if their approximations fit into the calculated electromotive force in our models \\cite{hanasz04}.   In this part of this work we do not include the explicit analysis involving the conservation of the magnetic helicity. The investigation of the magnetic helicity conservation in our models is already advanced, but it is complex enough to be described in separate paper, which is under preparation. We cannot apply the approximations of \\cite{ruediger93} and \\cite{kitchatinov94} derived from the EMF quenching by the usage of the Second Order Correlation Approximations. They assumed that the Strouhal number $S$ is essentially smaller than 1 ($S\\ll1$, where $S = \\tau_c \\times \\rm v/ \\rm l_c$). This assumption is not fullfiled in our numerical experiments, where $S$ is about 1. Finally, we discuss our results.  \\begin{table*} \\begin{center} \\caption{Parameters of models examined in this paper. (*) The conversion rate, presented as 10\\% in \\cite{hanasz04}, was in fact equal to 100\\%, due to a trivial calculation mistake. Therefore, the overall injection rate of cosmic ray energy was equivalent to a realistic one, corresponding to SN rate=~20~kpc$^{-2}$Myr$^{-1}$, with the energy conversion factor = 10\\%. \\label{table-params}} \\begin{tabular}{|l|c|c|c|c|} \\hline Model   & A & B & C & D\\\\ \\hline Domain sizes [kpc] & 0.5 $\\times$ 1 $\\times$ 1.2 & \\multicolumn{3}{c}{0.5 $\\times$ 1 $\\times$ 4} \\vline \\\\ Resolution         & 50  $\\times$ 100 $\\times$ 120& \\multicolumn{3}{c}{50 $\\times$ 100 $\\times$ 400} \\vline \\\\ Vertical gravity at $R$ [kpc] = & 8.5 & \\multicolumn{3}{c}{5} \\vline\\\\ Gas column density [cm$^{-2}$] & ... & \\multicolumn{3}{c}{27$\\times$10$^{20}$} \\vline\\\\ Angular velocity $\\Omega$ [Myr$^{-1}$]  & 0.05 & \\multicolumn{3}{c}{0.05}  \\vline \\\\ SN rate [kpc$^{-2}$Myr$^{-1}$]           & $2$    & \\multicolumn{3}{c}{130} \\vline \\\\ Initial $\\alpha=e_{\\rm mag}/e_{\\rm gas}$ & $10^{-8}$ & \\multicolumn{3}{c}{$10^{-4}$}\\vline \\\\ Diffusion coefficients $K_\\parallel$, $K_\\perp$ [cm$^2$s$^{-1}$]  & $3 \\times 10^{27}$, $3 \\times 10^{26}$ & \\multicolumn{3}{c}{$3 \\times 10^{27}$, $3 \\times 10^{26}$}\\vline \\\\ Conversion rate of SN kinetic to CR energy  &  10\\%$^{(*)}$  & \\multicolumn{3}{c}{10\\%}\\vline \\\\ \\hline Resistivity $\\eta$ [cm$^2$s$^{-1}$]    & $3\\times 10^{24}$ & $0\\times 10^{24}$ & $3\\times 10^{24}$ & $30\\times 10^{24}$  \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": " \\begin{enumerate} \\item {Neither the velocity nor the magnetic field scale separation occurs in our model.}  \\item { The electromotive forces in the  cosmic-ray driven dynamo model are nonlinear, but none of the two examined nonlinear approaches is capable of reproducing electromotive forces in the numerical experiments correctly.}  \\item {Various nonlinear prescriptions of the dynamo coefficients have been proposed by other outhors, however, they are not capable of reconstructing the electromotive force resulting from experiments of cosmic-ray driven dynamo. Moreover the reconstructions of the magnetic field  produce too fast or too slow growth of the magnetic energy in comparison with the results of cosmic-ray driven dynamo numerical experiments.}  \\end{enumerate}  Extension of the present work, including considerations of magnetic helicity conservation will be presented in the forthcoming paper."
    },
    "0710/0710.1094_arXiv.txt": {
        "abstract": "We report a measurement of the supernova (SN) rates (Ia and core-collapse) in galaxy clusters based on the 136 SNe of the sample described in \\citet{C99} and \\citet{M05}.  Early-type cluster galaxies show a type Ia SN rate (0.066 SNuM) similar to that obtained by \\citet{sharon07} and more than 3 times larger than that in field early-type galaxies (0.019 SNuM). This difference has a 98\\% statistical confidence level. We examine many possible observational biases which could affect the rate determination, and conclude that none of them is likely to significantly alter the results. We investigate how the rate is related to several properties of the parent galaxies, and find that cluster membership, morphology and radio power all affect the SN rate, while galaxy mass has no measurable effect. The increased rate may be due to galaxy interactions in clusters, inducing either the formation of young stars or a different evolution of the progenitor binary systems.  We present the first measurement of the core-collapse SN rate in cluster late-type galaxies, which turns out to be comparable to the rate in field galaxies. This suggests that no large systematic difference in the initial mass function exists between the two environments. ",
        "introduction": "\\label{sec:intro}  Type Ia Supernovae (SNe Ia) are believed to be the result of the thermonuclear explosion of a C/O white dwarf (WD) in a binary system due to mass exchange with the secondary star. This conclusion follows from a few fundamental arguments: the explosion requires a degenerate system, such as a white dwarf; the presence of SNe Ia in old stellar systems implies that at least some of their progenitors must come from old, low-mass stars; the lack of hydrogen in the SN spectra requires that the progenitor has lost its outer envelope; and, the released energy per unit mass is of the order of the energy output of the thermonuclear conversion of carbon or oxygen into iron. Considerable uncertainties about the explosion model remain within this broad framework, such as the structure and the composition of the exploding WD (He, C/O, or O/Ne), the mass at explosion (at, below, or above the Chandrasekhar mass) and the flame propagation (detonation, deflagration, or a combination of the two). The key observations constraining the explosion models are the light curve and the evolution of the spectra.  Large uncertainties also remain regarding the nature of the progenitor binary system, its evolution through one or more common envelope phases, and its configuration (single or double-degenerate) at the moment of the explosion (see \\citealt{yungelson05}, for a review). Solving the problem of the progenitor system is of great importance for modern cosmology as SNe dominate metal production, (e.g., \\citealt{matteucci86}), are expected to be important producer of high-redshift dust \\citep{maiolino01, maiolino04a,maiolino04b,bianchi07}, and are essential to understand the feedback process during galaxy formation (e.g., \\citealt{scannapieco06}). The nature of the progenitor systems can be probed by studying the SN rate in different stellar populations, and constraining the delay time distribution (DTD) between star formation and SN explosion.   \\smallskip  In 1983, Greggio \\& Renzini computed the expected DTD for a single-degenerate system. The computation was later refined by many authors and extended to double-degenerate systems \\citep{tornambe86,tornambe89,tutukov94,yungelson00,matteucci01,belczynski05,greggio05}. The DTD can be convolved with the star formation history (SFH) of each galaxy to obtain its SN rate. The observation of the SN rates per unit mass in galaxies of different types \\citep{M05,sullivan06} and in radio-loud early-type galaxies \\citep{dellavalle05} has proved to be an effective way to constrain the DTD. The SN rates per unit mass show that SNe Ia must come from both young and old progenitors \\citep{M05,sullivan06}. The dependence of the SN rate on the radio power of the parent galaxy is well reproduced by a ``two channel'' model \\citep{mannucci06}, in which about half of the SNe Ia, the so-called ``prompt'' population, explode soon after the formation of the progenitors, on time scales shorter than $10^8$ yr, while the other half (the ``tardy'' population) explode on a much longer time scale, of the order of $10^{10}$ yr. Several attempts to compare the evolution of SN rate with redshift with that of the SFR have also been presented (see, among many others, \\citealt{galyam04,dahlen04,cappellaro05,neill06,barris06,botticella07} and \\citealt{poznanski07}), but the large uncertainties on both quantities prevent strong conclusions (see, for example, \\citealt{forster06}).  \\smallskip  In principle, an accurate measurement of the DTD could identify the progenitor binary system. In practice, both the large number of free parameters involved in the theoretical computations of the DTD, and the complex SFHs of most of the galaxies make this identification much more uncertain. To solve the problem of the complexity of the SFH, it is interesting to measure the SN Ia rate in galaxy clusters. Most of the stellar mass of these systems is contained in elliptical galaxies, whose stellar populations are dominated by old stars. Despite the problem that even a small amount of new stars could give a significant contribution to the SN rate (see the discussion in sect.~\\ref{sec:discussion}), the reduction in the uncertainty in the SFH is of great help to derive the DTD.  \\smallskip  The cluster SN rate is also of great importance to study the metallicity evolution of the universe. The gravitational potential well of galaxy clusters is deep enough to retain in the intracluster medium (ICM) all the metals which are produced in galactic or intergalactic SNe. As a result, the metallicity of the ICM is a good measure of the integrated past history of cluster star formation and metal production. As discussed by \\citet{renzini93}, the measured amount of iron is an order of magnitude too high to be produced by SNe Ia exploding at the current rate. Explanations of this effect include the presence of higher SN rates in the past \\citep{matteucci06}, the importance of the intracluster stellar population \\citep{zaritsky04}, or evolving properties of star formation processes \\citep{maoz04,lowenstein06}. The observed abundance ratios in the ICM can be used to constrain the ratio between the total numbers of Ia and CC SNe, as recently done by \\citet{deplaa07}. Constraints on the SN Ia models can also be derived from the radial distribution of metallicity \\citep{dupke02}. \\citet{calura07} used the observed cosmic evolution of iron abundances in \\citet{balestra07} to constrain the history of SN explosion, iron formation and gas stripping in galaxy clusters. They found good agreement with the observations, especially when the ``two channel'' model of SNe Ia by \\citet{mannucci06} is used.    \\smallskip  There are strong motivations for measuring also the cluster rates of the other physical class of SNe, the core-collapse (CC) group. Type II and type Ib/c SNe are attributed to this group because there is a general consensus that these explosions are due to the collapse of the core of a massive (about 8--40~\\msun) star. Thus, CC SNe are expected to be good tracers of star formation in moderately dusty environments (see \\citealt{mannucci07}). Their rate per unit mass is also very sensitive to the initial mass function (IMF), because SN explosions are due to massive stars while most of the mass is locked in low-mass stars. As a consequence, studying the CC SN rate as a function of environment is a sensitive test for any systematic difference in IMF.  \\begin{table} \\caption{ Measured type Ia SN rates in early-type cluster galaxies \\label{tab:history} } \\begin{tabular}{llcc} \\hline \\hline Reference         & ~~$z$ ~~($z$ range)  & N$_{SN}$& Rate\\\\ &                     &         & (SNuB) \\\\ \\hline This work         & 0.02 (0.005--0.04)  &  20 & $0.28^{+0.11}_{-0.08}$\\\\ \\citet{crane77}   & 0.023 (0.020--0.026)&   8 & $\\sim$0.10  \\\\ \\citet{barbon78}  & 0.023 (0.020--0.026)&   5 & $\\sim$0.16  \\\\ \\citet{germany04} & 0.05 (0.02--0.08)   &  23 & unpubl. \\\\ \\citet{sharon07}  & 0.15 (0.06--0.19)   &   6 & 0.27$^{+0.16}_{-0.11}$\\\\ \\citet{galyam02}  & 0.25 (0.18--0.37)   &   1 & 0.39$^{+1.65}_{-0.37}$\\\\ \\citet{galyam02}  & 0.90 (0.83--1.27)   &   1 & 0.80$^{+0.92}_{-0.41}$\\\\ \\hline \\end{tabular} \\end{table}   \\subsection{The observed cluster supernova rate}  Prompted by all these motivations, several groups have measured the SN Ia rate in galaxy clusters, but the results are still quite sparse. The first published values are due to \\citet{crane77} and \\citet{barbon78} (see Table~\\ref{tab:history} for a summary, including the results of our work, discussed below), before a clear distinction between type Ia and Ib/c had been introduced. They used a sample of 5--8 SNe in the Coma Cluster and constrained the SN rate to be of the order of 0.15 SNuB (SN per century per $10^{10}$ \\lsun\\ in the B band). The SN rate as a function of galaxy environment was also addressed by \\citet{caldwell81} to derive information on SN progenitors.  Modern searches for cluster SNe begin with \\citet{norgaard89} who discovered a SN Ia in a cluster at $z=0.31$. Starting from the late '90s, the Mount Stromlo 1.3 m telescope was used to monitor a few tens of Abell Clusters \\citep{reiss98}. Three years of monitoring resulted in the detection of 23 candidate SNe Ia in cluster galaxies \\citep{germany04}, but a rate based on this sample was never published.  The first rates for cluster galaxies based on modern searches were published by \\citet{galyam02}. These authors used archive images from the Hubble Space Telescope (HST) of 9 galaxy clusters, and discovered 6 SNe, 2 of which are associated with the clusters, at $z=0.18$ and  $z=0.83$. The derived rates were affected by large statistical uncertainties due to the small number of detected SNe, but were consistent with a moderate increase of the rate with redshift compared to the rate in local elliptical galaxies. A sample of 140 low-redshift Abell clusters were monitored by the Wise Observatory Optical Transient Search (WOOTS, \\citealt{galyam07}) using the Wise 1m telescope. The seven detected cluster SNe were used to constrain the fraction of intergalactic stars and SNe \\citep{galyam03} and to measure the cluster SN rate \\citep{sharon07}. This latter work obtains a value of the SN rate per unit mass of $0.098^{+0.058}_{-0.039}$ SNuM (SN per century per $10^{10}$ \\msun\\ of stellar mass), which is larger than, but still consistent with, the value of $0.038^{+0.014}_{-0.012}$ SNuM, derived by \\citet{M05} for local ellipticals. Finally, a SN search in clusters is ongoing at the Bok Telescope on Kitt Peak \\citep{sand07}.  All of the previous published SN Ia rates are based on a small number of SNe and, as a consequence, have large statistical errors. Also, a cluster rate for CC SNe has never been published because many of the cited samples only contain Ia SNe. In this work, we use the SNe in the \\citet{C99} sample to study the SN rate as a function of galaxy environment.  Throughout this paper we use the ``737'' values of the cosmological parameters: $(h_{100},\\Omega_m,\\Omega_\\Lambda)=(0.7,0.3,0.7)$.  \\begin{figure} \\includegraphics[width=9cm]{galdist.ps} \\caption{ \\label{fig:galdist} Surface density of galaxies (upper panel) and fraction of early-type galaxies (lower panel) as a function of the projected distance from the closest cluster. Above 3 Mpc, the average for field galaxies is shown. The vertical dotted lines show the two projected distances, 0.5 and 1.5 Mpc, used to define cluster galaxies (see text). } \\end{figure} ",
        "conclusions": "\\label{sec:discussion}  The interpretation of the possible difference in SN Ia rate between cluster and field early-type galaxies is not straightforward. As the observed rate is the convolution of the SFH with the DTD, the differences could be due to either of these functions.  \\begin{enumerate}  \\item The first possibility is that the rate difference is due to differences in the stellar populations. \\citet{M05}, \\citet{sullivan06}, and \\citet{aubourg07} have shown that the type Ia SN rate has a strong dependence on the parent stellar population, with younger stars producing more SNe. The difference in SN rate could be related to this effect, i.e., to a higher level of recent star formation in cluster ellipticals. Only a very small amount of younger stars is needed, because the amplitude of the DTD at short times can be hundreds of times larger than at long times. As an example, the \\citet{greggio83} single-degenerate model has 300 times more amplitude at $10^8$ yr than at $10^{10}$ yr, and this means that a recently formed stellar population contributing 0.3\\% of the mass can provide as many SNe as the remaining 99.7\\% of old stars. For the ``two channel'' model by \\citet{mannucci06}, the amount of young stars needed can be even lower, at the 0.1\\% level, as this DTD amplitude ratio between $10^7$ and $10^{10}$ years is as large as 1000.  \\smallskip  The presence of traces of star formation in early-type galaxies is not inconsistent with other observations. Many ellipticals show signs of recent interactions or star formation activity: faint emission lines \\citep{sarzi06}, tidal tails \\citep{vandokkum05}, dust lanes \\citep{vandokkum95,colbert01}, HI gas \\citep{morganti06}, molecular gas \\citep{welch03}, and very blue UV colors \\citep{kaviraj06,schawinski07,haines07,kaviraj07}. Even if the interpretation of most of these effects is matter of debate (for example, \\citealt{serego07} have found only small amounts of HI gas in cluster ellipticals), the observations suggest a widespread, low-level presence of star formation.  The dependence of this presence with environment is not settled yet. \\citet{ferreras06} have found evidence for recent star formation, at the percent level, in ellipticals in compact groups, but not in field ellipticals. In contrast, \\citet{verdugo07} and \\citet{haines07} have found higher levels of present star formation in field rather then cluster early-type galaxies.   \\smallskip  Some studies (see, for example, \\citealt{sanchez06}, \\citealt{bernardi06} and \\citealt{collobert06}), have found younger ages in field early-type galaxies with respect to cluster galaxies (but \\citealt{serego06} have found no difference). Taken at face value, this would seem to contradict the star formation interpretation of the SN rate, but this is not necessary the case. Field ellipticals could be younger that cluster ellipticals, but nevertheless they could show a lower level of {\\em present} star formation. The difference in the age of the {\\em dominant} stellar population of early-type galaxies, of the order of 1 Gyr for ages of about 12 Gyr, might not be directly related to the amount of star formation in the last few $10^8$ years. Such a contribution cannot be detected in the integrated colors of the galaxies. The expected differences are at the 0.05 mag level for the (B--K) color, assuming the younger stars are not associated with dust, and even smaller (0.02 mag for $A_V$=1), allowing for dust extinction.  It is usually assumed that early-type galaxies can form new stars only after merging with a small, gas rich galaxy, because usually they do not host much interstellar gas. The average amount of stars formed is proportional to the merger (or encounter) rate, to the typical amount of gas in the accreted galaxy, and to the efficiency of star formation in the accreted gas. It is possible that one or more of these quantities are larger for cluster galaxies than for field galaxies because of the different galaxy volume density and galaxy-galaxy encounter velocity.  \\smallskip  If this is the correct interpretation, the ``prompt'' population of SNe Ia would be associated with the explosion of CC SNe from the same young stellar populations. If a SN Ia is to explode within $10^8$ yr of the formation of its progenitor, the primary star of the progenitor binary system must have a mass above 5.5 \\msun\\ to allow for the formation of a white dwarf in such a short time. \\citet{mannucci06} have shown that reproducing the observed SN rates by using the ``bimodal'' DTD in that paper implies that about 7\\% of all stars between 5.5 and 8 \\msun\\ explode as ``prompt'' SNe Ia, while the ``tardy'' population corresponds to a lower explosion efficiency, about 2\\%, and on a much longer timescale (see also \\citealt{maoz07} for various estimates of these efficiencies). For a Salpeter IMF and assuming that 100\\% of the stars between 8 and 40 \\msun\\ end up at CC SNe, we expect 1.3 CC SNe for each ``prompt'' type Ia. Assuming that the difference between cluster and field early-type galaxies is due to the ``prompt'' SNe Ia, the rate of this population is of the order of 0.066-0.019=0.047 SNuM (see Table~\\ref{tab:massrate}). Converting this rate to an observed number, about 2 CC SNe are expected in the cluster early-type galaxies of our sample, consistent with our null detection at about 1.3$\\sigma$ level. We conclude that the non detection of CC in the early-type galaxies belonging to our sample and the corresponding upper limits to the CC rate are consistent with the hypothesis of a ``prompt'' Ia component. We also note that some CC SNe have been discovered in the recent past in prototypical early-type galaxies. \\citep{pastorello07}.   \\item A second possible interpretation is that the higher rate in cluster early-type galaxies is related to differences in the DTD. If the stars in ellipticals are 9-12 Gyrs old (see, for example, \\citealt{mannucci01}), the SN rate is dominated by the tail of the DTD at long times. Differences in the environments could produce small differences in the shape of this function, for example because of the higher numbers of encounters.  A interesting possibility is also that the changes in the DTD are related to differences in metallicity between cluster and field early-type galaxies, as discussed by \\cite{sanchez06,bernardi06,collobert06} and \\cite{prieto07}. The differences between cluster and field galaxy metallicity presented by these papers are neither large nor always in the same direction. Nevertheless systematic, although not large, differences in metallicity could be present and produce significant changes in the DTD, for example, by affecting the efficiency of mass loss during the complex life of a binary system.  \\end{enumerate}   Table~\\ref{tab:history} lists the different measurements of the SN rate in early-type cluster galaxies. The evolution of this rate can be compared with the history of star formation of the parent galaxies to derive the DTD. Currently published cluster SN rates at z$>$0.2 are too uncertain to permit any strong conclusions. However, current and future searches for SNe are expected to change this situation and allow for the derivation of meaningful constraints (see, for example, \\citealt{sharon06}).  \\smallskip  To summarise, we have used a sample of 136 SNe in the local universe to measure the SN rate as a function of environment. For the first time, we measure the CC SN rate in clusters. We find it is very similar to the CC SN rate in field galaxies, suggesting that the IMF is not a strong function of the environment. For Ia SNe, the rates in clusters and in the field are similar for all galaxy types except for the early-type systems, where we detect a significant excess in clusters. This excess is not related to other properties of those galaxies, such as mass, morphology, or radio loudness. Environments itself appears to be important. We interpret this effect as possibly due to galaxy-galaxy interaction in clusters, either producing a small amount of young stars (of the order of the percent in mass over one Hubble time), or affecting the evolution of the properties of the binary systems.  \\bigskip  {\\bf Acknowledgments}  We thank Sperello di Serego and the MEGA group (Arcetri Extragalactic Meeting) for useful discussions about the properties of elliptical galaxies. 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. DM, MD, and AG thank the Kavli Institute for Theoretical Physics for its hospitality. This research was supported in part by the National Science Foundation under Grant No. PHY05-51164."
    },
    "0710/0710.1788_arXiv.txt": {
        "abstract": "{The {\\it ROSAT} X-ray source \\cal\\ has recently been identified as a likely compact object whose properties suggest it could be a very nearby radio millisecond pulsar at $d = 80 - 260$\\,pc.} {We investigated this hypothesis by searching for radio pulsations using the Westerbork Synthesis Radio Telescope.} {We observed \\cal\\ at 385 and 1380\\,MHz, recording at high time and frequency resolution in order to maintain sensitivity to millisecond pulsations.  These data were searched both for dispersed single pulses and using Fourier techniques sensitive to constant and orbitally modulated periodicities.} {No radio pulsations were detected in these observations, resulting in pulsed radio luminosity limits of $L_{400}^{\\rm max} \\approx 0.3 (d/250 {\\rm pc})^2$\\,mJy kpc$^2$ and $L_{1400}^{\\rm max} \\approx 0.03 (d/250 {\\rm pc})^2$\\,mJy kpc$^2$ at 400 and 1400\\,MHz respectively.} {The lack of detectable radio pulsations from \\cal\\ brings into question its identification as a nearby radio pulsar, though, because the pulsar could be beamed away from us, this hypothesis cannot be strictly ruled out.} ",
        "introduction": "Recently, \\citet*{rfs07}, hereafter RFS07, have identified the X-ray source \\cal\\ (from the {\\it ROSAT} All-Sky Survey Bright Source Catalog) as having an X-ray to optical flux ratio $F_{\\rm X}(0.1-2.4 {\\rm keV})/F_{\\rm V} > 8700$ ($3\\sigma$).  Such a high ratio is strong evidence that \\cal, dubbed ``Calvera'' by RFS07, is a compact object.  RFS07 consider several specific source classes to explain \\cal's X-ray spectrum and luminosity: an X-ray dim isolated neutron star (INS), an anomalous X-ray pulsar (AXP), a compact central object (CCO), or a nearby radio pulsar.  Based on careful comparison of \\cal's properties with the canonical features displayed by these different classes of neutron star, RFS07 conclude that the most likely explanation is that \\cal\\ is a very nearby radio pulsar ($d = 80-260$\\,pc) similar to the radio millisecond pulsars (MSPs) residing in the globular cluster 47~Tuc.  We have investigated this interpretation by conducting sensitive searches for radio pulsations using the Westerbork Synthesis Radio Telescope (WSRT) in the Netherlands.  No pulsations were found by these searches, and we use these non-detections to place strong limits on the pulsed radio luminosity of \\cal.  These limits bring into question the interpretation of this object as a nearby radio pulsar. ",
        "conclusions": "No plausible astronomical radio pulsations or bright, dispersed single pulses were detected in any of our observations.  Using the radiometer equation modified for pulsar signals \\citep{dtws85}, we can place limits on \\cal's flux density in these observations (Table~\\ref{obs.tab}). We find that \\cal\\ has a maximum flux density at 400\\,MHz $S_{400}^{\\rm max} \\approx 4$\\,mJy and a maximum flux density at 1400\\,MHz $S_{1400}^{\\rm max} \\approx 0.3$\\,mJy, where the fractional uncertainty on these limits is roughly 50\\%. RFS07 argue that \\cal\\ may be a radio pulsar with a nearby distance $d = 80-260$\\,pc.  Assuming the pulsar is isolated, we were sensitive to spin periods encompassing the observed range for radio pulsars ($P_{\\rm spin} \\sim 1$\\,ms$-10$\\,s), with a factor of roughly $2-5$ degredation in sensitivity due to red-noise at the longest periods. Using an assumed distance $d = 250$\\,pc, we can convert our flux density limits to pseudo luminosity ($L \\equiv Sd^2$) limits.  We find $L_{400}^{\\rm max} \\approx 0.3 (d/250 {\\rm pc})^2$\\,mJy kpc$^2$ and $L_{1400}^{\\rm max} \\approx 0.03 (d/250 {\\rm pc})^2$\\,mJy kpc$^2$.  A simple check of the ATNF pulsar catalog\\footnote{Available at http://www.atnf.csiro.au/research/pulsar/psrcat} \\citep{mhth05} shows that $\\lesssim 1$\\% of the known pulsars have a luminosity below these limits. This suggests that if \\cal\\ is a radio pulsar, then it is either especially weak, not beamed towards the Earth, or significantly further away than 250\\,pc.  Assuming \\cal\\ is a radio MSP, what is the probability that its radio beam will pass the Earth (i.e. what is the beaming fraction of such pulsars)? A period-dependent beaming fraction proportional to $P_{\\rm spin}^{-0.5}$ has been shown for normal, un-recycled pulsars \\citep[see e.g.][]{ran93}. When extrapolated to millisecond spin periods, this relation implies a $\\sim 100$\\% beaming fraction for MSPs.  However, \\citet{kxl+98} find observational evidence that this relationship does not apply directly to MSPs and that the beaming fraction of MSPs is more like $50-90$\\%. These estimates are consistent with considerations of the MSP population in 47~Tuc. \\citet{hge+05} used a deep {\\it Chandra} observation of the cluster to perform a population analysis which concluded that there are likely $\\sim 25$ radio MSPs ($< 60$ at 95\\% confidence) residing in the cluster, independent of beaming.  Comparing this with optical and radio studies of the MSP population by \\citet{egh+03} and \\citet{mdca04} respectively, who both conclude that the total MSP population in 47~Tuc is $\\sim 30$, \\citet{hge+05} conclude that the beaming fraction is $\\gtrsim 37$\\%.  While these various studies suggest that the beaming fraction of MSPs is significantly larger than for normal pulsars, and could be quite high, we cannot strongly rule out beaming as the cause of our non-detection of \\cal.  It is also possible, though we feel unlikely, that we have not seen \\cal\\ because it is in a compact binary orbit and/or is highly accelerated by a binary companion.  Our luminosity limits are for a coherent search and do not include these effects. However, given that we have employed acceleration searches, which are generally sensitive in cases where the total integration time $T_{\\rm int}$ is less than roughly 1/10 of the orbital period $P_{\\rm orb}$, \\cal\\ would have to be in a $\\lesssim 2$-hr orbit {\\it and} fairly weak (or eclipsed) not to have been detected in our searches of 15-min data sections. The shortest orbital period of any known radio MSP is 1.6\\,hr \\citep[][PSR~J0024$-$7204R in the GC 47~Tuc]{clf+00}.  In conclusion, given the caveats we have discussed, primarily pertaining to extreme orbital or spin parameters, we believe that these searches were sensitive enough to detect any nearby radio pulsar coincident with \\cal, assuming it is beamed towards the Earth.  We have checked for catalogued radio sources in the FIRST, NVSS, and WENSS surveys and find no counterpart to \\cal.  Future, deeper observations detecting a radio, optical, or X-ray pulsar wind nebula could address whether this source is a nearby radio pulsar beamed away from Earth.  Alternate scenarios for \\cal's nature, for instance that it might be the first unhosted CCO (RFS07), should continue to be investigated."
    },
    "0710/0710.0351_arXiv.txt": {
        "abstract": "{Modern observations and models of various astrophysical objects suggest that many of their physical parameters fluctuate substantially at different spatial scales. The rich variety of the emission processes, including Transition Radiation but not limited to it, arising in such turbulent media constitutes the scope of Stochastic Theory of Radiation. We review general approaches applied in the stochastic theory of radiation and specific methods used to calculate the transition radiation produced by fast particles in the magnetized randomly inhomogeneous plasma. The importance of the theory of transition radiation for astrophysics is illustrated by one example of its detailed application to a solar radio burst, including specially designed algorithms of the spectral forward fitting.}  \\def\\gsim{\\ \\raise 3pt \\hbox{$>$} \\kern -8.5pt \\raise -2pt \\hbox{$\\sim$}\\ } \\def\\lsim{\\ \\raise 3pt \\hbox{$<$} \\kern -8.5pt \\raise -2pt \\hbox{$\\sim$}\\ } ",
        "introduction": "The phenomenon of transition radiation was discovered theoretically by two Nobel Prize winning (2003 and 1958 respectively) physicists \\cite{Gin_Fr}. Ginzburg and Frank (1946) considered a simplest case when a charged particle passed through a boundary between two dielectrically different media and so generated waves due to a variation of the dielectric constant at the boundary. Remarkably, no acceleration of the particle is necessary to produce the emission due to transition through the boundary.  It is easy to understand that a similar effect of electromagnetic emission will take place if a medium is uniformly filled by turbulence that produces fluctuations of the dielectric constant throughout the whole volume rather than at an isolated boundary. Many astrophysical sources, especially those under strong energy release, are believed to be filled by turbulent, randomly inhomogeneous plasma and fast, nonthermal particles. In this situation, an efficient contribution of the transition radiation to the overall electromagnetic emission should be produced. Therefore, distinguishing this contribution from competing mechanisms is important. Below we describe the fundamentals of the transition radiation produced in a magnetized turbulent plasma, and demonstrate its high potential for astrophysical applications. ",
        "conclusions": "Nita et al. (2005) proved that the dm continuum component of this solar radio burst is produced by RTR and derived the level of the microturbulence in the plasma to be $\\left<\\Delta n^2\\right>/n^2 = 10^{-5}$. This finding is potentially very important for other cosmic objects. Indeed, the obtained microturbulence level is not particularly strong and much stronger turbulence is expected in many cases, especially, when there is a strong release of the energy at the source. Sometimes, such energy release gives rise to a relativistic expansion of the source, so the emission spectrum is Doppler-boosted and RTR produced at the local plasma frequency can be observed at the Earth even from relatively tenuous sources with low plasma frequency.  In this study we present more evidence in favor of RTR generation at the dm continuum solar bursts and use this emission component to derive additional plasma parameters. In particular, we determine the mean plasma frequency and its dispersion at the source in the course of time. Interestingly, these two parameters do not change much during the time of the dm burst. We note that these parameters are obtained from the total power spectra recorded without spatial resolution. Fig. 2 demonstrates that the radio sources at various frequencies do not coincide exactly. Therefore, in cases where a sequence of spatially resolved spectra are available we would be able to study the structure of the flaring plasma density in much greater detail as well as the distribution of the microturbulence over the source.  Generally speaking, the RTR contribution is also informative about the fast electrons producing it. In the example presented in Fig. 3 we show the emission decay constants, which can be associated with the fast electron life times. In our case we used exponential fragments of the light curves at the late decay phase of the emission, since no exponential phase was found in the early decay phase. We found the life time to be within 10-40 sec, which corresponds to the electrons of 300 keV or larger in the case of dense flare plasma available in this event. On the other hand, we can expect that most of the RTR emission (around the peak of the burst) is produced by the electrons with E=100-200 keV (Nita et al. 2005). This apparent contradiction can be easily resolved if we recall that the lower energy electrons have the life time of only a few seconds in the given dense plasma, so they die even before the light curves reach the exponential decay stage, and we observe the RTR contribution from preferentially higher energy electrons late in the event."
    },
    "0710/0710.5574_arXiv.txt": {
        "abstract": "We study the effect of quasar feedback on distributions of baryons and properties of intracluster medium in galaxy groups using high-resolution numerical simulations. We use the entropy-conserving Gadget code that includes gas cooling and star formation, modified to include a physically-based model of quasar feedback. For a sample of ten galaxy group-sized dark matter halos with masses in the range of $1$ to $5\\times 10^{13} M_{\\odot}/h$, star formation is suppressed by more than 50\\% in the inner regions due to the additional pressure support by quasar feedback, while gas is driven from the inner region towards the outer region of the halos. As a result, the average gas density is 50\\% lower in the inner region and 10\\% higher in the outer region in the simulation, compared to a similar simulation with no quasar feedback. Gas pressure is lowered by about 40\\% in the inner region and higher in the outer region, while temperature and entropy are enhanced in the inner region by about 20-40\\%. The total group gas fraction in the two simulations generally differs by less than 10\\%.  We also find a small change of the total thermal Sunyaev-Zeldovich distortion, leading to 10\\% changes in the microwave angular power spectrum at angular scales below two arcminutes. % ",
        "introduction": "Galaxy clusters and groups are the largest gravitationally bound objects in the universe, and they dominate the total baryon content of the universe. Their spatial distribution and mass function contain information about the formation and evolution of large-scale structure, which in turn constrain a variety of fundamental cosmological properties including normalization of the matter power spectrum, the cosmic baryon density, and dark matter properties. However, in order to use them as a cosmological probe, it is necessary to understand their astrophysical properties, and in particular their baryon physics. This issue is of particular current interest due to upcoming arcminute-resolution microwave sky surveys like ACT \\citep{kosowsky06,fowler07} and SPT \\citep{ruhl05}, which will image galaxy clusters via the Sunyaev-Zeldovich distortions to the cosmic microwave blackbody spectrum from the hot electrons in the cluster gas \\citep{sz80}.   The majority of baryons in clusters and groups are in the form of hot intracluster gas rather than than individual galaxies. Properties of the Intracluster Medium (ICM) have been studied through a combination of X-ray and radio observations \\citep{nulsen05, heinz02, fabian2000}. Although the dark matter distribution in galaxy clusters follow a self-similar relation \\citep{pointecouteu05,vikhilin06}, the hot gas does not \\citep{sanderson03,popesso05}. Additional non-gravitational sources of heating are required to explain the observations. One interesting and plausible possibility is the energy radiated from quasars or Active Galactic Nuclei (AGN) and deposited into the ICM \\citep{kaiser91,valageassilk99,nath02, evan05, evan06}, which we study in this work.  The best arena in which to study the impact of various feedback mechanisms is galaxy groups. Massive clusters with deeper gravitational potential wells are likely to have their global thermodynamic and morphological properties less affected by feedback. In comparison, galaxy groups have shallower potential wells while still having enough gas to display the effect of feedback on the ICM. Galaxy groups have recently been observed in X-rays at redshifts as large as $z=0.6$ \\citep{willis05}. In the optical band, \\cite{sdss07} have compiled group catalogs from the SDSS Data Release 5 catalog. Evidence for heating by a central AGN or radio source in galaxy groups and clusters has been the subject of several recent papers \\citep{croston05, jetha06,sanderson05}. These observations show excess entropy in cluster cores, which suggests that  some heating process must act to offset cooling.  In recent years, cosmological simulations including dark matter and gas have been able to follow the evolution of individual galaxy groups and clusters. A number of studies have investigated the cluster baryon fraction and its evolution in numerical simulations. Adiabatic simulations that do not include radiative cooling find cluster baryon fractions around $0.85$ of the universal baryon fraction \\cite{evrard90,metzler_evrard94,navarro_etal95,lubin_etal96,eke_etal98,frenk_etal99,mohr_etal99,bialek_etal01}. Preheating the gas reduces the fraction further \\citep{bialek_etal01,borgani_etal02,muanwong_etal02,kay_etal03}. When cooling, star formation and other feedback processes are included, the baryon fraction is higher than that obtained from adiabatic simulations \\citep{muanwong_etal02,kay_etal03,valdarnini03,ettori04,nagai07}. This leads to an ``overcooling'' problem and indicates an additional feedback mechanism.       In the current study, we analyze the effect of quasar feedback on the baryon distribution and thermodynamics of hot gas in galaxy groups. We also study its implication for the Sunyaev-Zeldovich angular power spectrum, which receives a dominant contribution from high-redshift halos. Ref.~\\cite{ks02} showed that the thermal SZ angular power spectrum provides a strong constraint on the normalization of the matter power spectrum, $\\sigma_8$. Upcoming SZ surveys like ACT or SPT will have sufficient sensitivity to determine $\\sigma_8$ with an accuracy limited by uncertainty in the theoretical model. Also, the kinematic SZ effect is a measure of bulk motions in the universe and may be a competitive probe for studying cosmology \\citep{sehgal04, bk06, penn06, dedeo05, maturi07, roncarelli07}. But one of the major sources of uncertainty in modeling the kSZ effect is the gas fraction and its evolution. So understanding both the thermal and kinematic SZ signals requires detailed understanding of feedback mechanisms in galaxy clusters and groups.  The mechanisms and effects of feedback are also a long-standing question in astrophysics, with particular bearing on the process of galaxy formation.  To this end, we have analyzed a sample of ten galaxy groups at $z=1$ from numerical cosmological simulations of gas and dark matter which have been extended to include a self-consistent model for the evolution of massive black holes and their baryon feedback.  At redshift $z>1$, the quasar mode of black hole accretion is expected to be the dominant feedback mechanism, compared to the radio-loud accretion mode which becomes important at lower redshifts \\citep{sijacki07}. The size of our simulations prevents studying feedback in galaxy clusters, but rather restricts us to less massive galaxy groups. But as already mentioned, galaxy groups with shallow potential wells provide the best place to study non-gravitational heating and its implications for the properties of hot gas. High-redshift galaxy groups are also a major contributor to the thermal SZ power spectrum, which peaks around $z\\approx 1$, when galaxy groups are more numerous than massive clusters \\cite{ks02}.  Following this introduction, Section II describes our simulation and its implementation of quasar feedback. In Section III we study the effect of numerical resolution on our results; in Section IV we describe our results and compare them with a simulation that do not include quasar feedback. Finally, in Section V we summarize our results and discuss directions for future work, including motivations and prospects for studying more massive galaxy clusters and more realistic feedback modeling for quasars and AGN. ",
        "conclusions": "  1. Compared to the no-feedback case, star formation is suppressed by 30-40\\% in the inner regions of the halos because of the additional pressure support provided by quasar feedback.  2. Quasar feedback redistributes hot gas, driving it from the inner region towards the outer part of the halos. As a result, gas density is 20\\% less in the inner part and 10\\% to 15\\% greater in the outer region when compared to the simulation without feedback. However, the gas fraction in the two simulation differs by only 5\\% to 10\\%, and gas fractions tends to increase mildly with increasing halo mass.   3. The ratio of gas mass to stellar mass increases by a factor of 3.5 in the simulation including quasar feedback and a factor of 2.5 in the simulation without quasar feedback in the region $0.2 R_{200m}<R<0.5 R_{200m}$. This contradicts the common assumption that this ratio is constant at all radii.  4. Both temperature and entropy increase by 30\\% to 50\\% in the halo core region because of the additional thermal energy radiated by quasars.  5. Pressure decreases by 30\\% in the inner region and increases by 15\\% to 20\\% at radii larger than 0.4 $R_{200m}$ due to the increased gas density in this region. This leads to a change of about 6\\% in the mean Sunyaev-Zeldovich $Y$-distortion. The resulting SZ angular power spectrum will be larger by around 10\\% for $l>5000$. We find little dependence of the SZ enhancement with halo mass.  The effects of quasar feedback on the intracluster medium will be most evident in the group-sized haloes considered here, with their relatively shallow gravitational potential wells. Observationally, the most interesting haloes are larger in mass by a factor of ten, galaxy clusters: these are the haloes which are most readily detected via their SZ, X-ray, or optical signals. The gas fractions do not show any particular trend with increasing halo mass, and star fractions increase very weakly with mass over the halo mass range studied here. So it is reasonable to expect that the results of this work will hold for cluster-sized halos as well. Nevertheless, given the  substantial impact of quasar feedback on various properties of the intracluster medium which the current study suggests, it is imperative to study cluster-sized halos as well. This requires larger-volume simulations, as the number density of clusters decreases with cluster mass. To this end, we are currently running a simulation of box size $50\\,{\\rm Mpc}/h$; results will be reported elsewhere.  This is the first attempt to study the impact of quasar feedback on the baryon fraction and thermodynamics of the intracluster medium in a cosmological hydrodynamic simulation. Both the gas and star fractions in our simulation are consistent with current observational limits \\citep{allen_etal04, ettori99}. Note that we have studied only quasar feedback at redshifts greater than unity. However, active galactic nuclei also inject energy into the ICM via a ``radio mode'' which is believed to be the dominant feedback mechanism at lower redshift \\citep{sijacki05, sijacki07}. Thus our results should be treated as a conservative estimate of the total impact of AGN feedback for galaxy groups at low redshifts.  Gas pressure in cosmological halos, particularly those with masses ranging from galaxy groups to galaxy clusters, determines the important thermal Sunyaev-Zeldovich signal which will soon be measured with high precision. The gas fraction is important for connecting kinematic SZ signals of cluster gas momentum with theoretical predictions about cluster velocity or total momentum. This paper takes the first step towards quantifying the impact of quasars on these quantities, which turns out to be significant but not dominating. Much work remains to be done, both through larger simulations which contain many galaxy-cluster-sized haloes, and in enhancing the realism of the quasar feedback models. We hope the results here plus the exciting observational prospects in the near future will open the door to further advances in this area."
    },
    "0710/0710.2481_arXiv.txt": {
        "abstract": "The stellar populations in the bulges of S0s, together with the galaxies' dynamics, masses and globular clusters, contain very interesting clues about their formation. I present here recent evidence suggesting that S0s are the descendants of fading spirals whose star formation ceased. ",
        "introduction": "Combining published data with high-quality VLT/FORS spectroscopy of sample of Fornax S0s (Bedregal et al.\\ 2006a)  we have carried out a combined study of the Tully-Fisher relation and the stellar populations of these galaxies.  Despite the relatively small sample and the considerable technical challenges involved in determining the true rotation velocity $V_{\\rm rot}$ from absorption line spectra of galaxies with significant non-rotational support (see Mathieu et al.\\ 2002), some very interesting results arise. S0s lie systematically below the spiral galaxy Tully-Fisher relation in both the optical and near-infrared (Figure~1). If S0s are the descendants of spiral galaxies, this offset can be naturally interpreted as arising from the luminosity evolution of spiral galaxies that have faded since ceasing star formation. Moreover, the amount of fading implied by the offset of individual S0s from the spiral relation seems to correlate with the luminosity-weighted age of their stellar population, particularly at their centres (Figure~2). This correlation suggests a scenario in which the star formation clock stopped when gas was stripped out from a spiral galaxy and it began to fade into an S0. The stronger correlation at small radii indicates a final last-gasp burst of star formation in this region. See Bedregal, Arag\\'on-Salamanca  \\& Merrifield (2006b) for details.  \\begin{figure} \\begin{center} \\includegraphics[height=3.3in,width=4.0in,angle=0]{Aragon-Salamanca_fig1.ps} \\end{center} \\caption{$B$-band Tully-Fisher relation (TFR) for S0 galaxies using different samples from the literature (open symbols) and our VLT Fornax data (filled circles). The solid and dashed lines show two independent determinations of the TFR relation for local spirals. On average (dotted line), S0s are $\\sim3$ times fainter than spirals at similar rotation velocities (Bedregal, Arag\\'on-Salamanca \\& Merrifield 2006b). } \\label{fig:fig1} \\end{figure}   \\begin{figure} \\begin{center} \\includegraphics[height=1.65in,angle=0]{Aragon-Salamanca_fig2.eps} \\end{center} \\caption{ For our VLT Fornax data we plot the shift in magnitudes from the $B$-band spiral TFR versus the stellar population age at the galaxy centre (left panel), at $1\\,R_e$ (middle panel) and at $2\\,R_e$ (right panel). The lines show models for fading spirals. Note that the correlation is strongest for the central stellar populations of the galaxies, suggesting that the last episode of star formation took place there (Bedregal, Arag\\'on-Salamanca \\& Merrifield 2006b). } \\label{fig:fig2} \\end{figure} ",
        "conclusions": "The stellar populations, dynamics and globular clusters of S0s  provide evidence consistent with these galaxies being the descendants of fading spirals whose star formation ceased.  However, caution is needed since significant problems could still exist with this picture (see, e.g., Christlein \\& Zabludoff 2004; Boselli \\& Gavazzi 2006). Moreover, the  number of galaxies studied here is still small, and  it would be highly desirable to extend this kind of studies to much larger samples covering a broad range of galaxy masses and environments.   \\begin{figure} \\begin{center} \\includegraphics[height=2.9in,angle=0]{Aragon-Salamanca_fig3.eps} \\end{center} \\caption{ Log$_{10}$ of the luminosity-weighted ages is Gyr vs.\\ the globular cluster specific frequency ($S_N$) of S0s. The line shows the evolution expected for a fading galaxy according to the stellar population models of Bruzual \\& Charlot (2003). The correlation between the fading of the galaxies (or increase in $S_N$) and the spectroscopically-determined age of their stellar populations is clearly consistent with the predictions of a simple fading model. Note that the $S_N$ value for NGC3115B is very unreliable and almost certainly severely overestimated due to contamination from the GC systems of neighbouring galaxies. See Barr et al.\\ (2007) for details. } \\label{fig:fig3} \\end{figure}"
    },
    "0710/0710.0484_arXiv.txt": {
        "abstract": "We study the propagation of relativistic jets originating from AGNs within the Interstellar/Intergalactic Medium of their host galaxies, and use it to build a model for the suppression of stellar formation within the expanding cocoon. ",
        "introduction": "\\label{sec:jetcocoon} Relativistic jets from Supermassive Black Holes hosted within the bulges of spirals and nuclei of early-type galaxies inject significant amounts of energy into the medium within which they propagate, \\begin{figure} \\centering \\includegraphics[scale=0.6,angle=90]{antonucciodelogu_fig1.eps} \\caption{Expansion of the cocoon within the ISM. The jet has an input mechanical power $\\textrm{P}_{jet} = 10^{46} \\textrm{ergs}\\cdot\\textrm{cm}^{-3}\\cdot\\textrm{sec}^{-1}$, and the ISM density is: $\\, n_{e}^{ism} = 1 e^{-} cm^{-3}$.}\\label{fig:cocoon} \\end{figure} creating an extended, underdense and hot cocoon. We have performed a series of simulations of these jets and cocoons using an AMR code, FLASH 2.5: a typical output is shown in fig.~\\ref{fig:cocoon}.\\\\ We find that a slight modification of the exact models of \\cite{1991MNRAS.250..581F} and \\cite{1997MNRAS.286..215K}, approximates well the simulations. ",
        "conclusions": ""
    },
    "0710/0710.2815_arXiv.txt": {
        "abstract": "We present evolutionary tracks of binary systems with high mass companion stars and stellar-through-intermediate mass BHs. Using Eggleton's stellar evolution code, we compute the luminosity produced by accretion from the donor during its entire evolution. We compute also the evolution of the optical spectrum of the binary system taking the disc contribution and irradiation effects into account. The calculations presented here can be used to constrain the properties of the donor stars in Ultraluminous X-ray Sources by comparing their position on the HR or color-magnitude diagrams with the evolutionary tracks of massive BH binaries. This approach may actually provide interesting clues also on the properties of the binary system itself, including the BH mass.  We found that, on the basis of their position on the color-magnitude diagram, some of the candidate counterparts considered can be ruled out and more stringent constraints can be applied to the donor masses. ",
        "introduction": "Point like off nuclear X-ray sources with luminosities well in excess of the Eddington limit for a stellar mass black hole, have been discovered in a large number of nearby galaxies (e.g.~\\citealt{2002ApJS..143...25C}; \\citealt{sgtw04}; \\citealt{2005ApJS..157...59L}). These Ultraluminous X-ray Sources (ULXs) are too dim to be low luminosity AGNs and too bright to be normal X-ray binaries (XRBs) emitting below the Eddington limit. In this paper we define a ULX as a source with bolometric luminosity in excess of the Eddington limit for a stellar mass black hole of $20\\msun$ and less luminous than a few $10^{41}\\rm erg/s$. This definition implies that a Galactic XRB radiating isotropically at or below the Eddington limit can not be a ULX.  In spiral or starburst galaxies ULXs turn out to be associated with star forming regions, emission nebulae and stellar clusters (\\citealt{zfrm02}, \\citealt{pm02}). These facts along with, in some case, the detection of stellar optical counterparts (\\citealt{2001MNRAS.325L...7R}; \\citealt{2002MNRAS.335L..67G}; \\citealt{2002ApJ...580L..31L}; \\citealt{2004ApJ...602..249L};\\citealt{kwz04};\\citealt{zamp1}; \\citealt{2005ApJ...629..233K}; \\citealt{mucetal05, mucetal07}; \\citealt{scpmw04}) strongly indicate an association between ULXs and young massive stars, although the nature of the accreting compact object remains unclear.  A certain number of ULXs, however, appear as isolated X-ray sources with no obvious counterpart at any wavelength (\\citealt{sm04}) and without a clear association with star forming regions or emission nebulae. Some low-luminosity ULXs may possibly be present also in elliptical galaxies (\\citealt{jcbg03}), but their nature and the actual evidence for their existence is not well established (see e.g. \\citealt{iba04}, \\citealt{aglc04}, \\citealt{2005ApJ...622L..89G}).  The nature of the ULXs in spiral galaxies is less controversial. Different models have been proposed to explain the large luminosities reached by these sources. One of the favored models consists of an intermediate mass black hole (IMBH) with a mass in the range $10^{2}-10^{3}M_{\\rm \\odot}$, accreting from an high mass donor star. The presence of an IMBH can account for most of the observational properties of ULXs in a rather straightforward way. For instance, the observed cool disc spectra of some ULX can be explained with the fact that the innermost stable circular orbit of an IMBH is larger than that of a stellar mass black hole (e.g. \\citealt{mfm04}). The detection of a $\\sim$50-160 mHz quasi periodic oscillations in the power density spectrum of M82 X-1 and NGC5408 X-1 (\\citealt{2003ApJ...586L..61S}, \\citealt{2004ApJ...614L.113F}, \\citealt{2006MNRAS.365.1123M}; \\citealt{stroh07}), the very high luminosity of some ULXs ($\\sim 10^{41} \\ergs$) along with their cool discs, and the energy content and morphology of the nebulae around some of them (\\citealt{pm02}) all suggest an IMBH interpretation. The main problem with this interpretation resides in the formation mechanism of such an extreme object. In fact, if IMBHs with masses in excess of $\\sim 100 M_\\odot$ exist, they will require a new formation root with respect to the stellar black holes in our Galaxy and to the supermassive black holes in Active Galactic Nuclei. Until now two scenarios have been proposed to form a black hole in the intermediate mass range: the runaway collision of massive stars in dense open clusters (\\citealt{2004Natur.428..724P}, \\citealt{2004ApJ...604..632G}) and the primordial collapse of a very high mass star with zero metallicity (\\citealt{2000ApJ...540...39A}, \\citealt{2001ApJ...551L..27M}). Both the mechanisms however suffer of a certain degree of uncertainty related to the incomplete knowledge of the behavior of very massive stars. Therefore we have no final evidence that an IMBH can really form. Furthermore, the interpretation of the soft components observed in some ULXs in terms of cool accretion discs is not univocal (e.g. \\citealt{2005ApJ...631L.109C}, \\citealt{2005ApJ...635..198D}, \\citealt{2005MNRAS.357.1363R}, \\citealt{fk06}, \\citealt{gonc06}, \\citealt{2006MNRAS.368..397S}).  Other interpretations in terms of stellar (or quasi-stellar) mass black holes have been proposed. A mechanical (\\citealt{kdwfe01}) or a relativistic beaming (\\citealt{kfm02}) can reproduce the observed luminosities of ULXs up to a few $10^{40}\\ergs$ with a beaming factor around $\\sim 10.$ At most, as suggested by \\citealt{2005MNRAS.357..275K}, accretion from helium rich matter from a geometrically thick disc can generate luminosities up to $\\sim 5\\times10^{40}\\ergs$. However, luminosities in excess of $5\\times 10^{40}\\ergs$ ($\\sim 5$\\% of the ULX population) and the isotropy of the ionized nebulae around some ULXs can not be easily explained in terms of beaming models. On the other hand, photon bubbles disc instabilities (\\citealt{b02}, \\citealt{b06}) and emission from a slim disc (\\citealt{wmm01}, \\citealt{ezkmw03}) can produce genuine isotropic super-Eddington luminosities around 10 times the Eddington limit.  As shown by \\cite{2005MNRAS.356..401R}, a normal binary with a stellar mass black hole and a donor star with initial mass $\\simgreat 10\\mdot$, can in principle explain a large sample of ULXs if a super Eddington luminosity around $10$ is allowed (see also \\citealt{2003MNRAS.341..385P}). However, the typical temperature of a slim disc in this regime ($\\sim$ 1--2 keV) is inconsistent with the observation of some cool disc sources which are, at the same time, the most luminous ULXs. Furthermore, luminosities in excess of a few $10^{40}\\ergs$ are not attainable with slim discs around stellar mass black holes. On the other hand, the photon bubble instability model makes no predictions on the observable accretion disc temperature and therefore cannot be compared directly with observations.  Thanks to the high precision astrometry of the {\\it Chandra} X-ray observatory, we know that some ULXs have optical counterparts which match with the X-ray source position (\\citealt{2002ApJ...580L..31L}, \\citealt{2004ApJ...602..249L}, \\citealt{zamp1}, \\citealt{kwz04}, \\citealt{sm04}, \\citealt{2005ApJ...629..233K}, \\citealt{mucetal05, mucetal07}, \\citealt{Liu07}). Most of them are suspected to be main sequence (MS) high mass stars or supergiant stars of uncertain spectral type. The ambiguity arises because, until now, optical spectra of these stars either are not available, or they are too noisy and with peculiar spectral features (\\citealt{2004ApJ...602..249L}), or there is more than an optical counterpart in the X-ray error box (\\citealt{sm04}, \\citealt{mucetal05}).  In this paper, we present the evolutionary tracks of binary systems with high mass companion stars and stellar-through-intermediate mass BHs and compare them with the properties of the donor stars in ULX binary systems. In \\S 2 we present the model adopted to evolve the binary system. In \\S 3, we summarize the properties of the four ULXs considered and their optical counterparts. Our results are presented in \\S 4 and compared with the properties of ULX counterparts in \\S 5. Conclusions follow in \\S 6. ",
        "conclusions": "In this paper we have shown that the effects of the binary evolution on the optical properties of a donor star in a ULX binary system can significantly alter its colors and evolution. We included also the contribution of the optical emission of the accretion disc and the X-ray irradiation of the donor and disc surfaces. The evolutionary track of a donor in an accreting binary on the CM diagram is very different with respect to that of a single isolated star. These important differences can not be overlooked when trying to identify the donor mass of a ULX. We calculated tracks for stellar and intermediate mass black holes, and demonstrate the brighter nature of the IMBH as an intrinsic phenomenon produced by the low mass ratio in the binary.  The photometric data of four ULX counterparts have been compared with the evolutionary tracks of massive donors on the CM diagram. We find that the counterparts 2, 3 and 4 of NGC 4559 X-7 are consistent only with donors of $10$--$15\\msun$ undergoing a case B or C mass transfer episode. The only possibility for objects 5 and 8 is a very massive ($\\sim50\\msun$) MS donor or a H-shell burning donor with mass between 10-15 and 30$\\msun$. Finally, object 1 is compatible only with a very massive ($M\\simgreat 50\\msun$) donor in a H-shell burning phase. Our results are quite different with respect to what obtained in previous analysis. \\cite{scpmw04} compared theoretical evolutionary tracks of single stars between 9 and 25$\\msun$ with the position in the CM diagram of the counterparts of NGC4559 X-7. They find that the colors and luminosity of six out of seven stars are consistent with the tracks of main sequence, blue or red supergiant with masses in the range 10-15 $M_\\odot$ and ages $\\sim 20$ Myr. Their favored counterpart, object 1, was identified with a main sequence or blue supergiant of $\\sim 20 M_\\odot$ and an age of $\\sim 10$ Myr, although \\cite{scpmw04} recognize that X-ray irradiation may affect its colors and hence its classification. Irradiation effects were taken into account by \\cite{cpw07}, giving a donor mass of 5-20$M_{\\odot}$ for object $1$. In their work, however, \\cite{cpw07} considered irradiation on single stars and did not take into account the effects of the markedly different evolution of a donor star in a binary system.  As far as NGC1313 X-2 is concerned, we definitely rule out the candidate counterpart C2 and identify C1 as a H-shell burning donor of 10$\\msun<$M$<$15$\\msun$ around a stellar mass black hole or a $\\simless 15 \\msun$ donor undergoing RLOF during MS (or the H-shell burning phase) in an IMBH binary system. On statistical grounds alone, it is then more likely that NGC1313 X-2 hosts a $\\sim 100 \\msun$ BH rather than a stellar mass BH as the H-shell burning phase is much shorter than the MS phase. Our result for object C2 leaves only C1 as the likely counterpart of NGC 1313 X-2, in agreement with the evidence coming from the refined X-ray astrometry of the field recently reported by \\cite{Liu07}.  On the basis of single star isochrone fitting of the parent stellar population, \\cite{pak06} and \\cite{ramsey06} estimate a maximum MS mass of 8--9$\\msun$ for object C1. \\cite{Liu07} find consistency with either a $\\sim 8\\msun$ star of very low metallicity or an O spectral type, solar metellicity star of $\\sim 30\\msun$. Our estimated range of donor masses for object C1 appears more in agreement with the value (10-18$\\msun$) reported by \\cite{mucetal07}, probably because they also included the contribution of the disc and irradiation effects. In Holmberg II X-1 there are several possibilities, and we can only rule out stars with M$<$10$\\msun$ (\\citealt{cpw07} give a lower bound of 5$\\msun$), while for M81 X-9 we are left with the unique possibility that the donor is a $\\sim 10\\msun$ star in the H-shell burning phase.  We conclude therefore that in all previous work where the binary evolution and/or irradiation effects ware not taken into account, the donor mass has been systematically overestimated. We find that mass transfer occurring at late stages is very unlikely, not only for the short timescale of the contact phase ($10^{3}$--$10^{5}$yrs) but also for the incompatibility of the optical colors with the majority of the observed counterparts. If mass transfer sets in during MS, two contact phases occur (segments A--B and C--D). The mass transfer timescales are $t_{A-B}\\sim 10^{6}-10^{7}\\rm\\,yrs$ and $t_{C-D}\\sim 10^{3}-10^{5}\\rm\\,yrs$ (depending on the donor mass). Therefore assuming a flat distribution of periods between 1 and $\\sim$10 days, as during MS the contact phase is reached for $P_{orb}\\simless 2$days, we expect a factor $\\sim \\frac{10}{2}\\times \\frac{t_{A-B}}{t_{C-D}}\\simeq 500$--$5000$ more main sequence systems than H-shell burning donors."
    },
    "0710/0710.4328_arXiv.txt": {
        "abstract": "The initial conditions and relevant physics for the formation of the earliest galaxies are well specified in the concordance cosmology. Using ab initio cosmological Eulerian adaptive mesh refinement radiation hydrodynamical calculations, we discuss how very massive stars start the process of cosmological reionization. The models include non-equilibrium primordial gas chemistry and cooling processes and accurate radiation transport in the Case B approximation using adaptively ray traced photon packages, retaining the time derivative in the transport equation. Supernova feedback is modeled by thermal explosions triggered at parsec scales. All calculations resolve the local Jeans length by at least 16 grid cells at all times and as such cover a spatial dynamic range of $\\sim$10$^6$. These first sources of reionization are highly intermittent and anisotropic and first photoionize the small scales voids surrounding the halos they form in, rather than the dense filaments they are embedded in. As the merging objects form larger, dwarf sized galaxies, the escape fraction of UV radiation decreases and the \\ion{H}{2} regions only break out on some sides of the galaxies making them even more anisotropic.  In three cases, SN blast waves induce star formation in overdense regions that were formed earlier from ionization front instabilities.  These stars form tens of parsecs away from the center of their parent DM halo. Approximately 5 ionizing photons are needed per sustained ionization when star formation in 10$^6$ \\Ms~halos are dominant in the calculation.  As the halos become larger than $\\sim$$10^7 \\Ms$, the ionizing photon escape fraction decreases, which in turn increases the number of photons per ionization to 15--50, in calculations with stellar feedback only.  Radiative feedback decreases clumping factors by 25 per cent when compared to simulations without star formation and increases the average temperature of ionized gas to values between 3,000 and 10,000 K. ",
        "introduction": "It is clear that quasars are not responsible to keep the universe ionized at redshift 6. The very brightest galaxies at those redshifts alone also provide few photons.  The dominant sources of reionization so far are observationally unknown despite remarkable advances in finding sources at high redshift \\citep[e.g.][]{Shapiro86, Bouwens04, Fan06, Thompson07, Eyles06} and hints for a large number of unresolved sources at very high redshifts \\citep{Spergel07, Kashlinsky07} which is still a topic of debate \\citep{Cooray07, Thompson07}. At the same time, ab initio numerical simulations of structure formation in the concordance model of structure formation have found that the first luminous objects in the universe are formed inside of cold dark matter (CDM) dominated halos of total masses $2 \\times 10^5 - 10^6 \\Ms$ \\citep{Haiman96, Tegmark97, Abel98}.  Fully cosmological ab initio calculations of \\citet{Abel00, Abel02} and more recently \\citet{Yoshida06} clearly show that these objects will form isolated very massive stars. Such stars will be copious emitters of ultraviolet (UV) radiation and are as such prime suspect to get the process of cosmological reionization started.  In fact, one dimensional calculations of \\citet{Whalen04} and \\citet{Kitayama04} have already argued that the earliest \\ion{H}{2} regions will evaporate the gas from the host halos and that in fact most of the UV radiation of such stars would escape into the intergalactic medium. Recently, \\citet{Yoshida07a} and \\citet{Abel07} demonstrated with full three-dimensional radiation hydrodynamical simulations that indeed the first \\ion{H}{2} regions break out of their host halos quickly and fully disrupt the gaseous component of the cosmological parent halo.  All of this gas finds itself radially moving away from the star at $\\sim30\\kms$ at a distance of $\\sim100$ pc at the end of the stars life. At this time, the photo-ionized regions have now high electron fractions and little destructive Lyman-Werner band radiation fields creating ideal conditions for molecular hydrogen formation which may in fact stimulate further star formation above levels that would have occurred without the pre-ionization. Such conclusion have been obtained in calculations with approximations to multi dimensional radiative transfer or one dimensional numerical models \\citep{Ricotti02a, Nagakura05, OShea05, Yoshida06, Ahn07, Johnson07}. These early stars may also explode in supernovae and rapidly enrich the surrounding material with heavy elements, deposit kinetic energy and entropy to the gas out of which subsequent structure is to form. This illustrates some of the complex interplay of star formation, primordial gas chemistry, radiative and supernova feedback and readily explains why any reliable results will only be obtained using full ab initio three dimensional hydrodynamical simulations. In this paper, we present the most detailed such calculations yet carried out to date and discuss issues important to the understanding of the process of cosmological reionization.  It is timely to develop direct numerical models of early structure formation and cosmological reionization as considerable efforts are underway to \\begin{enumerate} \\item Observationally find the earliest galaxies with the James Webb Space Telescope \\citep[JWST;][]{Gardner06} and the Atacama Large Millimeter Array \\citep[ALMA;][]{Wilson05}, \\item Further constrain the amount and spatial non-uniformity of the polarization of the cosmic microwave background radiation \\citep{Page07}, \\item Measure the surface of reionization with LOFAR \\citep{Rottgering06}, MWA \\citep{Bowman07}, GMRT \\citep{Swarup91} and the Square Kilometer Array \\citep[SKA;][]{Schilizzi04}, and \\item Find high redshift gamma ray bursts with SWIFT \\citep{Gehrels04} and their infrared follow up observations. \\end{enumerate}  We begin by describing the cosmological simulations that include primordial star formation and accurate radiative transfer.  In \\S\\ref{sec:SF}, we report the details of the star formation environments and host halos in our calculations.  Then in \\S\\ref{sec:reion}, we describe the resulting start of cosmological reionization, and investigate the environments in which these primordial stars form and the evolution of the clumping factor.  We compare our results to previous calculations and further describe the nature of the primordial star formation and feedback in \\S\\ref{sec:discussion}.  Finally we summarize our results in the last section.    \\begin{deluxetable*}{lccccccc} \\tablecolumns{8} \\tabletypesize{} \\tablewidth{\\textwidth} \\tablecaption{Simulation Parameters\\label{tab:sims}}  \\tablehead{ \\colhead{Name} & \\colhead{$l$} & \\colhead{Cooling model} & \\colhead{SF} & \\colhead{SNe} & \\colhead{N$_{\\rm{part}}$} & \\colhead{N$_{\\rm{grid}}$} & \\colhead{N$_{\\rm{cell}}$} \\\\ \\colhead{} & \\colhead{[Mpc]} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{} } \\startdata  SimA-Adb & 1.0 & Adiabatic & No & No & 2.22 $\\times$ 10$^7$ & 30230 & 9.31 $\\times$ 10$^7$ (453$^3$) \\\\ SimA-HHe & 1.0 & H, He & No & No & 2.22 $\\times$ 10$^7$ & 40601 & 1.20 $\\times$ 10$^8$ (494$^3$) \\\\ SimA-RT & 1.0 & H, He, \\hh & Yes & No & 2.22 $\\times$ 10$^7$ & 44664 & 1.19 $\\times$ 10$^8$ (493$^3$) \\\\ SimB-Adb & 1.5 & Adiabatic & No & No & 1.26 $\\times$ 10$^7$ & 23227 & 6.47 $\\times$ 10$^7$ (402$^3$) \\\\ SimB-HHe & 1.5 & H, He & No & No & 1.26 $\\times$ 10$^7$ & 21409 & 6.51 $\\times$ 10$^7$ (402$^3$) \\\\ SimB-RT & 1.5 & H, He, \\hh & Yes & No & 1.26 $\\times$ 10$^7$ & 24013 & 6.54 $\\times$ 10$^7$ (403$^3$) \\\\ SimB-SN & 1.5 & H, He, \\hh & Yes & Yes & 1.26 $\\times$ 10$^7$ & 24996 & 6.39 $\\times$ 10$^7$ (400$^3$)  \\enddata \\tablecomments{Col. (1): Simulation name. Col. (2): Box size. Col. (3): Cooling model.  Col. (4): Star formation. Col. (5): Supernova feedback. Col. (6): Number of dark matter particles. Col. (7): Number of AMR grids. Col. (8): Number of unique grid cells.} \\end{deluxetable*} ",
        "conclusions": "\\label{sec:discussion}  We have studied the details of massive metal-free star formation and its role in the start of cosmological reionization.  We have treated star formation and radiation in a self-consistent manner, allowing for an accurate investigation of the evolution of cosmic structure under the influence of early Pop III stars.  Stellar radiation from these stars provides thermal, dynamical, and ionizing feedback to the host halos and IGM.  Although Pop III stars are not thought to provide the majority of ionizing photons needed for cosmological reionization, they play a key role in the early universe because early galaxies that form in these relic \\ion{H}{2} regions are significantly affected by Pop III feedback.  Hence it is important to consider primordial stellar feedback while studying early galaxy formation.  In this section, we compare our results to previous numerical simulations and semi-analytic models of reionization and then discuss any potential caveats of our methods and possible future directions of this line of research.   \\begin{figure}[t] \\begin{center} \\epsscale{1.15} \\plotone{f14_color} \\caption{\\label{fig:filtering} The Jeans mass $M_J$ and filtering mass $M_F$ that can form bound objects.  The squares denote the total mass of star forming halos in all three simulations.} \\end{center} \\end{figure}  \\subsection{Comparison to Previous Models}  \\subsubsection{Filtering Mass} \\label{sec:filtering}  One source of negative feedback is the suppression of gas accretion into potential wells when the IGM is preheated.  The lower limit of the mass of a star forming halo is the Jeans filtering mass \\begin{equation} \\label{eqn:filtering} M_F^{2/3}(a) = \\frac{3}{a} \\int_0^a \\> da^\\prime M_J^{2/3}(a^\\prime) \\left[ 1 - \\left(\\frac{a^\\prime}{a}\\right) \\right], \\end{equation} where $a$ and $M_J$ are the scale factor and time dependent Jeans mass in the \\ion{H}{2} region \\citep{Gnedin98, Gnedin00b}.  Additionally, the virial shocks are weakened if the accreting gas is preheated and will reduce the collisional ionization in halos with $\\tvir \\gsim 10^4$ K.  To illustrate the effect of Jeans smoothing, we take the large \\ion{H}{2} region of SimB-SN because it has the largest ionized filling fraction, which is constantly being heated after $z = 21$. Temperatures in this region fluctuates between 1,000~K and 30,000~K, depending on the proximity of the currently living stars.  In Figure \\ref{fig:filtering}, we show the resulting filtering mass of regions with an ionization fraction greater than $10^{-3}$ along with the total mass of star forming halos.  \\citet{Gnedin00b} found the minimum mass of a star forming halo is better described by $M_F$ instead of $M_J$.  Our simulations are in excellent agreement for halos that are experiencing star formation after reincorporation of their previously expelled gas.  The filtering mass is the appropriate choice for a minimum mass in this case as the halo forms from preheated gas.  However for halos that have already assembled before they become embedded in a relic \\ion{H}{2} region, the appropriate minimum mass $M_{\\rm{min}}$ is one that is regulated by the LW background \\citep{Machacek01, Wise05} and photo-evaporation \\citep[e.g.][]{Efstathiou92, Barkana99, Haiman01, Mesigner06}.  This is evident in the multitude of star forming halos below $M_F$.  With the exception of star formation induced by SN blast waves or I-fronts, this verifies the justification of using $M_{\\rm{min}}$ and $M_F$ for Pop III and galaxy formation, respectively, as a criterion for star forming halos in semi-analytic models.  \\begin{figure*}[t] \\begin{center} \\epsscale{1.15} \\plotone{f15} \\caption{\\label{fig:aniso} Density (left) and temperature (right) slices of an anisotropic \\ion{H}{2} region in the most massive halo of SimB-RT.  The star has lived for 2.5 Myr out of its 2.7 Myr lifetime.  The field of view is 900 proper parsecs.} \\end{center} \\end{figure*}   \\subsubsection{Star Formation Efficiency} \\label{sec:SFeff}  Semi-analytic models rely on a star formation efficiency $f_\\star$, which is the fraction of collapsed gas that forms stars, to calculate quantities such as emissivities, chemical enrichment, and IGM temperatures.  Low-mass halos that form a central star have $f_\\star \\sim 10^{-3}$ whose value originates from a single 100~\\Ms~star forming in a dark matter halo of mass 10$^6$~\\Ms~\\citep{Abel02, Bromm02, Yoshida06}.  Pop II star forming halos are usually calibrated with star formation efficiencies from local dwarf and high-redshift starburst galaxies and are usually on the order of a few percent \\citep[e.g.][]{Taylor99, Gnedin00a}.  This leads to the question: how efficient is star formation in these high-redshift halos while explicitly considering feedback?  This is especially important when halos start to form multiple massive stars and when metallicities are not sufficient to induce Pop II star formation.  The critical metallicity for a transition to Pop II is still unclear.  Recently, \\citet{Jappsen07a} showed that metal line cooling is dynamically unimportant in diffuse gas until metallicities of $10^{-2} \\> Z_\\odot$.  On the other hand, dust that is produced in SNe can generate efficient cooling down in dense gas with $10^{-6} \\> Z_\\odot$ \\citep{Schneider06}.  If the progenitors of the more massive halos did not result in a pair-instability SN, massive star formation can continue until it becomes sufficiently enriched.  Hence our simulations can probe the efficiency of this scenario of massive metal-free star formation.  It has also been suggested that the cosmological conditions that lead to the collapse of a metal-poor molecular cloud ($Z/Z_\\odot \\approx 10^{-3.5}$) may be more important than some critical metallicity in determining the initial mass function of a given stellar system \\citep{Jappsen07b}.  We calculate $f_\\star$ with the ratio of the sum of the stellar masses to the total gas mass of unique star-forming halos.  For example at the final redshift of 15.9 in SimA-RT, the most massive halo and its progenitors had hosted 11 stars and the gas mass of this halo is $1.8 \\times 10^6 \\Ms$, which results in $f_\\star = 6.1 \\times 10^{-4}$ for this particular halo.  Expanding this quantity to all star forming halos, $f_\\star/10^{-4} = 5.6, 6.7, 7.4$ for SimA-RT, SimB-RT, and SimB-SN, respectively.  We note that our choice of $M_\\star = 170 \\Ms$ in SimB-SN increases $f_\\star$ by 70\\%.  Our efficiencies are smaller than the isolated Pop III case because halos cannot form any stars once the first star expels the gas, and 40 -- 75 million years must pass until star can form again when the gas is reincorporated into the halo.     By regarding the feedback created by Pop III stars and associated complexities during the assembly of these halos, the $f_\\star$ values of $\\sim$$6 \\times 10^{-4}$ that are explicitly determined from our radiation hydrodynamical simulations provide a more accurate estimate on the early star formation efficiencies.  \\subsubsection{Intermittent \\& Anisotropic Sources}  Our treatment of star formation and feedback produces intermittent star formation, especially in low-mass halos.  If one does not account for this, star formation rates might be overestimated in this phase of star formation.  Kinetic energy feedback is the main cause of this behavior.  As discussed in sections \\ref{sec:haloMass} and \\ref{sec:kinetic}, shock waves created by D-type I-fronts and SN explosions expel most of the gas in halos with masses $\\lsim 10^7$ \\Ms.  A period of quiescence follows these instances of star formation.  Then stars are able to form after enough material has accreted back into the halo.  Only when the halo becomes massive enough to retain most of the outflows and cool efficiently through \\lya~and \\hh~radiative processes, star formation becomes more regular with successive stars forming.  The central gas structures in the host halo are usually anisotropic as it is acquiring material through accretion along filaments and mergers.  At scales smaller than 10 pc, the most optically thick regions produce shadows where the gas radially behind the dense clump is not photo-ionized or photo-heated by the source.  This produces cometary and so-called elephant trunk structures that are also seen in local star forming regions and have been discussed in detail since \\citet{Pottasch58}.  At a larger distance, the surrounding cosmic structure is composed of intersecting or adjacent filaments and satellite halos that breaks spherical symmetry.  The filaments and nearby halos are optically thick and remain cool and thus the density structures are largely unchanged.  The entropy of dense regions are not increased by stellar radiation and will feel little negative feedback from an entropy floor that only exists in the ionized IGM \\citep[cf.][]{Oh03}.  Ray-tracing allows for accurate tracking of I-fronts in this inhomogeneous medium.  Radiation propagates through the least optically thick path and generates champagne flows that have been studied extensively in the context of present day star formation \\citep[e.g.][]{Franco90, Churchwell02, Shu02, Arthur06}.  In the context of massive primordial stars, these champagne flows spread into the voids and are impeded by the inflowing filaments.  The resulting \\ion{H}{2} regions have ``butterfly'' morphologies \\citep{Abel99, Abel07, Alvarez06a, Mellema06, Yoshida07a}.  We also point out that sources embedded in relic \\ion{H}{2} largely maintain or increase the ionization fraction.  Here the already low optical depth of the recently ionized medium (within a recombination time) allows the radiation to travel to greater distances than a halo embedded in a completely neutral IGM.  The \\ion{H}{2} regions become increasingly anisotropic in higher mass halos.  We show an example of the morphology of a \\ion{H}{2} region near the end of the star's lifetime in a dark matter halo with mass $1.4 \\times 10^7 \\Ms$ in Figure \\ref{fig:aniso}.  \\subsection{Potential Caveats and Future Directions}  Although we have simulated the first generations of stars with radiation hydrodynamic simulations, our methods have neglected some potentially important processes and made an assumption about the Pop III stellar masses.  One clear shortcoming of our simulations is the small volume and limited statistics of the objects studied here.  However, it was our intention to focus on the effects of Pop III star formation on cosmological reionization and on the formation of an early dwarf galaxy instead of global statistics.  The star formation only simulations (SimA-RT and SimB-RT) converge to the similar averaged quantities, e.g. ionized fraction, temperatures, star formation rates, at the final redshift.  The evolution of these quantities differ because of the limited number of stars that form in the simulations, which then causes the evolution to depend on individual star formation times.  This variance should be expected in the small volumes that we simulate and should not diminish the significance of our results.  We have verified even in a 2.5-$\\sigma$ peak that Pop III stars cannot fully reionize the universe, which verified previous conclusions that low-luminosity galaxies provide the majority of ionizing photons. Furthermore, it is beneficial to study Pop III stellar feedback because it regulates the nature of star formation in these galaxies that form from pre-heated material.  Further radiation hydrodynamics simulations of primordial star and galaxy formation with larger volumes while still resolving the first star forming halos of mass $\\sim$$3 \\times 10^5 \\Ms$ will improve the statistics of early star formation, especially in more typical overdensities, i.e. 1-$\\sigma$ peaks, some of which could survive to become dwarf spheroidal galaxies at $z = 0$.  In this work, we treated the LW radiation field as optically thin, but in reality, \\hh~produces a non-zero optical depth above column densities of $10^{14} \\> \\mathrm{cm}^{-2}$ \\citep{Draine96}. Conversely, Doppler shifts of the LW lines arising from large velocity anisotropies and gradient may render \\hh~self-shielding unimportant up to column densities of $10^{20} - 10^{21} \\> \\mathrm{cm}^{-2}$ \\citep{Glover01}.  If self-shielding is important, it will lead to increased star formation in low-mass halos even when a nearby source is shining.  Moreover, \\hh~production can also be catalyzed ahead of I-fronts \\citep{Ricotti01, Ahn07}.  In these halos, LW radiation will be absorbed before it can dissociate the central \\hh~core.  On the same topic, we neglect any type of soft UV or LW background that is created by sources that are cosmologically nearby ($\\Delta z / z \\sim 0.1$).  A soft UV background either creates positive or negative feedback, depending on its strength \\citep{Mesigner06}, and a LW background increases the minimum halo mass of a star-forming halo \\citep{Machacek01, Yoshida03, OShea07, Wise07b}.  However in our calculations, the lack of self-shielding, which suppresses star formation in low-mass halos, and the neglect of a LW background, which allows star formation in these halos, may partially cancel each other. Hence one may expect no significant deviations in the SFRs and reionization history if one treats these processes explicitly.  To address the incident radiation and the resulting UV background from more rare density fluctuations outside of our simulation volume, it will be useful to bridge the gap between the start of reionization on Mpc scales to larger scale (10 -- 100 Mpc) simulations of reionization, such as the work of \\citet{Sokasian03}, \\citet{Iliev06}, \\citet{Zahn07}, and \\citet{Kohler07}.  Radiation characteristics from a volume that has similar overdensities as our Mpc-scale simulations can be sampled from such larger volumes to create a radiation background that inflicts the structures in our Mpc scale simulations. Inversely, perhaps the small-scale evolution of the clumping factor, filtering mass, and average temperature and ionization states can be used to create an accurate subgrid model in large volume reionization simulations.  Another potential caveat is the continued use of primordial gas chemistry in metal enriched regions in the SN runs.  Our simulations with SNe give excellent initial conditions to self-consistently treating the transition to low-mass star formation.  In future work, we plan to introduce metal-line and dust cooling models \\citep[e.g. from][] {Glover07, Smith07} to study this transition.  The one main assumption about Pop III stars in our calculations is the fixed, user-defined stellar mass.  The initial mass function (IMF) of these stars is largely unknown, therefore we did not want to introduce an uncertainty by choosing a fiducial IMF.  It is possible to calculate a rough estimate of the stellar mass by comparing the accretion rates and Kelvin-Helmholtz time of the contracting molecular cloud \\citep{Abel02, OShea05}.  Protostellar models of primordial stars have also shown that the zero-age main sequence (ZAMS) is reached at 100 \\Ms~for typical accretion histories after the star halts its adiabatic contraction \\citep{Omukai03, Yoshida06}. Furthermore, we have neglected HD cooling, which may become important in halos embedded in relic \\ion{H}{2} regions and result in lower mass ($\\sim$$30 \\Ms$) metal-free stars \\citep{OShea05, Greif06, Yoshida07b}. Based on accretion histories of star forming halos, one can estimate the ZAMS stellar mass for each halo and create a more self-consistent and ab initio treatment of Pop III star formation and feedback.    We conducted three radiation hydrodynamical, adaptive mesh refinement simulations that supplement our previous cosmological simulations that focused on the hydrodynamics and cooling during early galaxy formation.  These new simulations concentrated on the formation and feedback of massive, metal-free stars.  We used adaptive ray tracing to accurately track the resulting \\ion{H}{2} regions and followed the evolution of the photo-ionized and photo-heated IGM.  We also explored on the details of early star formation in these simulations.  Theories of early galaxy formation and reionization and large scale reionization simulations can benefit from the useful quantities and characteristics of the high redshift universe, such as SFR and IGM temperatures and ionization states, calculated in our simulations. The key results from this work are listed below.  \\medskip  1. SFRs increase from $5 \\times 10^{-4}$ at redshift 30 to $6 \\times 10^{-3}$ \\sfr~at redshift 20 in our simulations.  Afterwards the SFR begins to have a bursting nature in halos more massive than $10^7 \\Ms$ and fluctuates around $10^{-2}$ \\sfr.  These rates are larger than the ones calculated in \\citet{Hernquist03} because our simulation volume samples a highly biased region that contains a 2.5-$\\sigma$ density fluctuation.  The associated emissivity from these stars increase from 1 to $\\sim$100 ionizing photons per baryon per Hubble time between redshifts 15 and 30.  2. In order to provide a comparison to semi-analytic models, we calculate the star formation efficiency to be $\\sim$$6 \\times 10^{-4}$ averaged over all redshifts and the simulation volume.  For Pop III star formation, this is a factor of two lower than stars that are not affected by feedback \\citep{Abel02, Bromm02, Yoshida06, OShea07}.  3. Shock waves created by D-type I-fronts expel most of the gas in the host halos below $\\sim$$5 \\times 10^6 \\Ms$.  Above this mass, significant outflows that are still bound to the halo are generated. This feedback creates a dynamical picture of early structure formation, where star formation is suppressed in halos because of this baryon depletion, which is more effective than UV heating or the radiative dissociation of \\hh.  4. We see three instances of induced star formation in halos with masses $\\sim 3 \\times 10^6 \\Ms$.  Here a star forms as a SN blast wave overtakes an overdensity created by an ionization front instability. \\hh~formation is catalyzed by additional free electrons in the relic \\ion{H}{2} region and in the SN blast wave \\citep{Ferrara98}.  5. As star formation occurs regularly in the simulation after redshift 25, four (six) ionizing photons are needed per sustained hydrogen ionization.  As the most massive halo becomes larger than $\\sim$$10^7 \\Ms$ in the simulations without SNe, \\ion{H}{2} regions become trapped and ionizing radiation only escapes into the IGM in small solid angles.  Hence the number of photons per effective ionization increases to 15 (50).  In SimB-SN, stellar radiation from induced star formation have an escape fraction of nearly unity, which occur four times in the calculation.  This allows the IGM to remain ionized at a volume fraction 3 times higher than without SNe.  Similarly, the ionizing photon to ionization ratio also stays elevated at 10:1 instead of decreasing in the calculations with star formation only.  6. Our simulations that include star formation and \\hh~formation capture the entire evolution of the clumping factor that is used in semi-analytic models to calculate the effective enhancement of recombinations in the IGM.  We showed that clumping factors in the ionized medium fluctuate around the 75\\% of the values found in adiabatic simulations.  They evolve from unity at high redshifts and steadily increase to $\\sim$4 and 3.5 with and without SNe at $z = 17$, respectively.  Photo-evaporation from stellar feedback causes the decrease of the clumping factor.  7. We calculated the Jeans filtering mass with the volume-averaged temperature only in fully and partially ionized regions, which yields a better estimate than the temperature averaged over both ionized and neutral regions.  The filtering mass depends on the thermal history of the IGM, which mainly cools through Compton cooling.  It increases by two orders of magnitude to $\\sim$$3 \\times 10^7 \\Ms$ at $z \\sim 15$. It describes the minimum mass a halo requires to collapse after hosting a Pop III star.  For halos forming their first star, the minimum halo mass is regulated by the LW background \\citep{Machacek01} and photo-evaporation \\citep[e.g.][]{Haiman01}.  \\medskip  Pop III stellar feedback plays a key role in early star formation and the beginning of cosmological reionization.  The shallow potential wells of their host halos only amplify their radiative feedback.  Our understanding of the formation of the oldest galaxies and the characteristics of isolated dwarf galaxies may benefit from including the earliest stars and their feedback in galaxy formation models. Although these massive stars only partially reionized the universe, their feedback on the IGM and galaxies is crucial to include since it affects the characteristics of low-mass galaxies that are thought to be primarily responsible for cosmological reionization.  Harnessing observational clues about reionization, observations of local dwarf spheroidal galaxies, and numerical simulations that accurately handle star formation and feedback may provide great insight on the formation of the first galaxies, their properties, and how they completed cosmological reionization."
    },
    "0710/0710.4602_arXiv.txt": {
        "abstract": "We use semi-analytic techniques to evaluate the burst sensitivity of designs for the {\\it EXIST} hard X-ray survey mission.  Applying these techniques to the mission design proposed for the Beyond Einstein program, we find that with its very large field-of-view and faint gamma-ray burst detection threshold, {\\it EXIST} will detect and localize approximately two bursts per day, a large fraction of which may be at high redshift. We estimate that {\\it EXIST}'s maximum sensitivity will be $\\sim 4$~times greater than that of {\\it Swift}'s Burst Alert Telescope.  Bursts will be localized to better than 40~arcsec at threshold, with a burst position as good as a few arcsec for strong bursts. {\\it EXIST}'s combination of three different detector systems will provide spectra from 3~keV to more than 10~MeV.  Thus, EXIST will enable a major leap in the understanding of bursts, their evolution, environment, and utility as cosmological probes. ",
        "introduction": "In its quest to find black holes throughout the universe, the {\\it Energetic X-ray Imaging Survey Telescope (EXIST)} will detect, localize and study a large number of gamma-ray bursts, events thought to result from the birth of stellar-mass black holes.  We present the methods used to calculate {\\it EXIST}'s capabilities as a gamma-ray burst detector; we use the {\\it EXIST} design evaluated by the National Research Council's `Committee on NASA's Beyond Einstein Program:  An Architecture for Implementation' (2007; see also Grindlay 2007).  The combination of large detector area, broad energy coverage, and wide field-of-view (FOV) will result in the detection of a substantial number of bursts with a flux distribution extending to fainter fluxes than that of previous missions. Thus {\\it EXIST} should detect high redshift bursts, perhaps even bursts resulting from the death of Pop~III stars. {\\it EXIST}'s imaging detectors will localize the bursts, while the combination of detectors, both imaging and non-imaging, will result in well-determined spectra from 3~keV to well over 10~MeV.  In this paper we first describe the {\\it EXIST} mission design (\\S 2), emphasizing aspects relevant to burst detection. Then we present the sensitivity methodology (\\S 3), which we apply to the individual coded mask sub-telescopes (\\S 4). {\\it EXIST} will consist of arrays of these detectors with overlapping FOVs, and the overall mission sensitivity results from adding the sensitivity of the individual sub-telescopes (\\S 5).  Imaging using counts accumulated over different timescales increases the sensitivity (\\S 6). Finally, we combine these different calculations to evaluate {\\it EXIST}'s overall capabilities to study bursts (\\S 7). ",
        "conclusions": "From the preceding analysis, we can draw several conclusions on {\\it EXIST}'s impact on the study of gamma-ray bursts.  First we estimate the {\\it EXIST} burst detection rate. The BATSE observations provide the cumulative burst rate as a function of the peak flux value $\\psi_B$ averaged over $\\Delta t=1$~s in the $\\Delta E=$50--300~keV band (Band 2002): \\begin{equation} N_B \\sim 550 \\left[ {{\\psi_B}\\over \\hbox{0.3 ph cm$^{-2}$ s$^{-1}$}} \\right]^{-0.8} \\hbox{ bursts yr$^{-1}$ sky$^{-1}$ } \\quad . \\end{equation} The HET threshold sensitivity for a single sub-telescope on-axis is $\\psi_B\\sim 0.12$~ph~cm$^{-2}$~s$^{-1}$ for $E_p > 100$~keV. Using the BATSE rate in eq.~10 and integrating over the solid angle distribution in Figure~6 gives a burst detection rate for the HET of $\\sim 400$ bursts per year. Note that this rate is over the BATSE-specific values of $\\Delta E$ and $\\Delta t$, and {\\it EXIST} will use at least two different values of $\\Delta E$ (see \\S 4.1) and a variety of $\\Delta t$ values (see \\S 6). Consequently this rate should be increased by approximately 50\\% to account for the soft, faint, long duration bursts to which BATSE was less sensitive than {\\it EXIST}'s HET will be; we therefore expect the HET array to detect $\\sim 600$~bursts per year.  The value of $\\psi_B$ for an LET varies more with the burst spectral parameters than for an HET, and therefore estimates of the LET burst detection rate based on the BATSE rate are much more uncertain.  For a single LET $\\psi_B\\sim 0.3$~ph~cm$^{-2}$~s$^{-1}$ on axis at $E_p=100$~keV, which gives a burst detection rate of $\\sim$180 bursts per year using eq.~10 and the LET distribution in Figure~6. This rate should be increased by a factor of 2 to account for the different energy band $\\Delta E$ and accumulation times $\\Delta t$. We use a larger adjustment factor for the LETs than for the HETs because the LETs' energy band will overlap less with BATSE's than the HETs'.  We therefore expect the LET array to detect $\\sim350$~bursts per year.  Next we simulate the spectra that the {\\it EXIST} suite of detectors will observe.  Figure~9 shows a count spectrum (counts s$^{-1}$ keV$^{-1}$) for a moderately strong burst as it might be observed by the LETs (lefthand set of curves), HETs (middle set) and the CsI active shields for the HETs (righthand set; based on Garson et al. 2006a). The solid curves show the signal count rate, while the dashed curves provide the estimated background.  Thus {\\it EXIST} will facilitate spectral-temporal studies.  Particularly important to physical burst emission models is determining $E_p$, which is typically of order 250~keV (Kaneko et al. 2006).  In addition, correlations of $E_p$ with other burst properties, such as the `isotropic' energy (the Amati relation---Amati 2006) or total energy (the Ghirlanda relation---Ghirlanda, Ghisellini \\& Lazzati 2004), have been proposed. `Pseudo-redshifts' calculated from the observables related to the burst-frame parameters in these relations can be used in burst studies when spectroscopic redshifts are not available, and can guide ground observers in allocating telescope time to observing potential high redshift bursts.  The recently proposed Firmani relation (Firmani et al. 2006) correlates $E_p$, the peak luminosity, and a measure of the burst duration, all of which are related to observables in the gamma ray band. Thus pseudo-redshifts will be estimated using the Firmani relation based on {\\it EXIST} data alone, independent of observations by other facilities.  With well determined broadband spectra down to 3~keV, {\\it EXIST} will be capable of determining whether the Band function (Band et al. 1993) suffices to describe burst spectra.  For example, Preece et al. (1996) found evidence in the BATSE data for the presence of additional emission below 10~keV.  By scaling from the EXIST survey's source localization (Grindlay 2007), we find that bursts should be localized at threshold by the HETs and LETs to better than 40~arcsec and 8~arcsec, respectively; this localization should scale as $\\sim(\\sigma-1)^{-1}$.  Because the HETs are more sensitive to the LETs, the HET localization is relevant to the faintest bursts EXIST will detect.  {\\it EXIST}'s burst capabilities calculated above will constitute a major leap beyond current detectors, and should increase the number of high redshift bursts detected.  On average, high redshift bursts should be fainter, softer and longer than low redshift bursts (although the broad burst luminosity function and great variety in burst lightcurves and spectra obscure this trend).  Figure~10 compares the detector sensitivities of the HET (solid curve) and LET (dashed curve) arrays to the BAT on {\\it Swift} (dot-dashed curve) and BATSE's Large Area Detector (LAD---dot-dot-dashed curve).  As discussed above, the sensitivity is the threshold peak flux $F_T$ integrated over the 1--1000~keV band as a function of the spectrum's $E_p$; $\\alpha=-1$ and $\\beta=-2$ are assumed. In addition, the figure shows families of identical bursts at different redshifts (the curves with the points marked by `+'). Each family is defined by the value of $E_p$ in the burst frame; here again $\\alpha=-1$ and $\\beta=-2$ are assumed. In each family the burst would be observed to have $F_T$=7.5 ph cm$^{-2}$ s$^{-1}$ if it were at $z=1$. The points marked by `+' are spaced every $\\Delta z=1/2$; thus the uppermost points are at $z=1$ and the lowermost points are at $z=10$. The pulses in burst lightcurves become narrower (shorter) at higher energy, an effect that is generally proportional to $E^{0.4}$ (Fenimore et al. 1995). Since the observed lightcurve originated in a higher energy band, pulses should become narrower with redshift, reducing the peak flux when integrated over a fixed accumulation time; the plotted families include this effect.  Finally, in \\S 6 we showed that forming images on long timescales increases the sensitivity to long duration bursts, as might result from cosmological time dilation."
    },
    "0710/0710.1840_arXiv.txt": {
        "abstract": "{Detailed studies of Be stars in environments with different metallicities like the Magellanic Clouds or the Galactic bulge are necessary to understand the formation and evolution mechanisms of the circumstellar disks. However, a detailed study of Be stars in the direction of the bulge of our own galaxy has not been performed until now.} {The aim of this work is to report  the first systematic  search for Be star candidates in the direction of the Galactic Bulge. We present the full catalogue, give a brief description of the stellar variability seen, and show some light curve examples.} {We searched for stars matching specific criteria of magnitude, color and variability in the $I$ band. Our search was conducted on the 48 OGLE II fields of the Galactic Bulge. } {This search has resulted in 29053 Be star candidates, 198 of them showing periodic light variations. Nearly 1500 stars  in this final sample are almost certainly Be stars, providing an ideal sample for spectroscopic multiobject follow-up studies. } {} ",
        "introduction": "Be stars are non-supergiant fast rotator B stars whose spectra have, or had at some time, one or more Balmer lines in emission (Collins 1987).  This  emission is originated from a flattened circumstellar  disk and  can come and go episodically on time scales of days to decades.  The responsible mechanisms for the production and dynamics of the circumstellar gas are still not constrained. Possible mechanisms include non-radial pulsations, wind-compressed disk model,  magnetic activity and binarity (Porter $\\&$ Rivinius 2003, and  references therein).\\\\ Be stars are variable in brightness on three time scales that are often superimposed. Many of them (especially early type Be stars) show short-term photometric variability on time-scales of 0.2 to 2 days and amplitudes up 0.1 magnitudes, caused by non-radial pulsation or rotation (Percy et al. 2002, 2004). Some have mid-term variations on times scales form weeks to months, probably due to density waves within the disk (Sterken et al. 1996). Their amplitudes go to up to 0.2 magnitudes. They show also long-term variations  from years to decades, with amplitudes up 0.8 magnitudes (Mennickent, Vogt, \\& Sterken 1994; Pavlovski et al. 1997; Hubert  $\\&$ Floquet 1998; Percy $\\&$ Bakos 2001). Stagg (1987) found that this type of variability occurres in at least half of the Be stars. A few of Be stars are close binaries  and other present ejection process due to magnetic activity resulting in outbursts (Hubert et al. 1997).\\\\ Many Galactic Be stars have been surveyed for photometric variability in order to detect and confirm short-term or long-term variations and to find correlations between them and get clues on the physical processes in Be stars. For instance, Hubert $\\&$ Floquet (1998) investigated the short-period variability of Be stars using analysis based on the Fourier and CLEAN algorithms on the Hipparcos photometry. Percy et al. (2002, 2004) analyzed a large sample of stars with Hipparcos photometry using a form of autocorrelation function. Typical problems in these studies are the gaps in the time distribution of the measurements and the limitations of the used algorithms.\\\\ In recent years, many Be-star like variables have been discovered in the Magellanic Clouds, showing a big variety of light curves, some of them reminding those of the Galactic Be stars and others never observed in that type of stars. Keller et al. (2002) concluded that most of these blue variables should be Be stars. Searches for Be stars in the Magellanic Clouds were performed by Mennickent et al. (2002), Keller et al. (1999, 2002) and Sabogal et al. (2005) on the basis of selection criteria applied to different photometric databases (OGLE-II and MACHO),  and took into account amplitude of variability and ranges of color-magnitudes in the selection process. De Wit et al. (2006) investigated a subsample of the blue variables found by Mennickent et al. (2002) in the Small Magellanic Cloud and found that the photometric variability of these Be stars is due to variations in the amount of Brehmstrahlung due to the evolution of the circumstellar gas from a disk-shaped envelope towards a ring-like structure. \\\\ The study of Be stars  is relevant to make contributions to several important branches of stellar physics. In particular, detailed studies of Be stars in environments with different metallicities like the Magellanic Clouds or the Galactic bulge is crucial to understand the formation and evolution mechanisms of the circumstellar disks. However, a detailed study of Be stars in the direction of the bulge of our own galaxy has not been performed until now.\\\\ A very large number of stars was observed in the region of the Galactic Bulge during the course of the second phase of the Optical Gravitational Lensing Experiment (OGLE II) (Udalski, Kubiak \\& Szymanski 1997). We have performed a search for Be star candidates into this database. Here we present the  results of this search. ",
        "conclusions": "In this paper we have provided a catalogue of 29053 Be star candidates, 198 of them periodic, in the direction of the Galactic bulge that were selected basically on photometric criteria. Most of these Be star candidates are probably members of the Galactic disk and trace the gradient of metallicity towards the Galactic centre. They are ideal targets for future observing programs based on multiobject spectroscopy, narrow band photometry  or H$\\alpha$ imaging surveys. These programs could eventually establish their true nature and broke the residual degeneracy with variable red giants in the red part of the (V-I) color distribution."
    },
    "0710/0710.4524_arXiv.txt": {
        "abstract": "It is well know that the coronagraphic observations of halo CMEs are subject to projection effects.  Viewing in the plane of the sky does not allow us to determine the crucial parameters defining geoeffectivness of CMEs, such as the velocity, width or source location. We assume that halo CMEs at the beginning phase of propagation have  constant velocities, are symmetric and propagate with constant angular widths. Using these approximations and determining projected velocities and difference between times when CME appears on the opposite sides of the occulting  disk we are able to get the necessary parameters. We present consideration for the whole halo CMEs from SOHO/LASCO catalog until the end of 2000. We show that the halo CMEs are in average much more faster and wider than the all CMEs from the SOHO/LASCO catalog. ",
        "introduction": "Space Weather is significantly controlled by coronal mass ejections (CMEs) which can affect the Earth in a different way. CMEs originating close to the central meridian, directed toward the Earth, excite the biggest scientific concern. In coronagraphic observations they appear as enhancement surrounding the entire occulting disk and they were called `halo CME'. Since the first identification by Howard et al. (1982) plenty of them were detected and now they are routinely recorded by the high sensitive SOHO/LASCO coronagraphs. In spite of large advantage over previous instruments, the SOHO/LASCO observations are still affected by a projection effect (Gopalswamy et al. 2000b). Viewing in the plane of the sky does not allow us to determine the crucial parameters defining geoeffectivness of CMEs, such as the velocity, width or source location. Prediction of the arrival of CME in the vicinity of Earth is critically important in space weather investigations. Basing on interplanetary shocks detected by Wind and the corresponding CMEs detected by SOHO, Gopalswamy et al. (2000a) developed and next (Gopalswamy, 2001) improved an empirical model to predict the arrival of CMEs at 1AU. The critical element affecting this model is the initial CME speed. The better prediction could be achieved if real initial velocities are used instead projected velocities determined from LASCO observations. Similarly, attempts made to estimate the projection effect based on the location of the solar source employ ad hoc assumptions on parameter such as the width of CMEs (Sheeley et al. 1999, Leblanc and Dulk, 2000). In the present paper we try to determine these crucial parameters defining geoeffectivness of CMEs, such as the velocity, width or source location. We assume that halo CMEs at the beginning phase of propagation have a  constant velocities, are symmetric and propagate with  constant angular widths. Using these approximations and determining projected velocities and difference between times when CME appears on the opposite sides of the occulting disk we are able to get necessary parameters. We present results for the whole halo CMEs from SOHO/LASCO catalog until the and of 2000. ",
        "conclusions": "In this paper we present possibility to estimate the crucial parameters determining geoeffectiveness of the halo CMEs.  The clue of this method is based on the difference between times when the halo CME appears at the opposite sides (first and finally appearance) of the occulting disk. We considered the whole events form SOHO/LASCO catalog until the end of 2000. We were able to determine the real velocity, width and source location for 73th CMEs from our sample. Unfortunately, 58 events were symmetric or too faint to do necessary considerations. Results are listed in the three successive tables. This list could be use for further statistical examination or to prediction of the arrival of CME in the vicinity of Earth. Presented results suggest that the halo CMEs represent a specific class of CMEs which are very wide and fast. Using our results we have to remember that the simple model has several shortcomings: (i) CMEs may be accelerating, moving with constant speed or decelerating at the beginning phase of propagation. This means the constant velocity we assumed may not hold. (ii) CMEs may expand in addition to radial motion. Then the measured sky-plane speed is a sum of the expansion speed and the projected radial speed. This also would imply that the CMEs may not be a rigid cone as we assumed (Gopalswamy et al. 2001) (iii) The cone symmetry also may not hold. CME originating from loop structure could be elongated. All these limits can be overcome by stereoscopic observations only. Unfortunately, at the present time they are not available yet. It is necessary to develop the model to get the better fit to observations. The first step to improve our model could be achieved by consideration of acceleration and expansion of CMEs. \\\\ \\\\ {\\small \\bf Acknowledgments}\\\\ {\\scriptsize This paper was done  during work of Grzegorz Michalek at Center for Solar and Space Weather, Catholic University of America in Washington.\\\\ In this paper we used data from SOHO/LASCO CME catalog. This CME catalog is generated and maintained by the Center for Solar Physics and Space Weather, The Catholic University of America in cooperation with the Naval Research Laboratory and NASA. SOHO is a project of international cooperation between ESA and NASA.\\\\ Work done by Grzegorz Michalek was partly supported by {\\it Komitet Bada\\'{n} Naukowych} through the grant PB 258/P03/99/17.} \\vspace{10mm}"
    },
    "0710/0710.4238_arXiv.txt": {
        "abstract": "{Since most high- and intermediate-mass protostars are at great distance and form in clusters, high linear resolution observations are needed to investigate their physical properties.} {To study the gas in the innermost region around the protostars in the proto-cluster IRAS\\,05358+3543, we observed the source in several transitions of methanol and other molecular species with the Plateau de Bure Interferometer and the Submillimeter Array, reaching a linear resolution of 1100~AU.} {We determine the kinetic temperature of the gas around the protostars through an LVG and LTE analysis of their molecular emission; the column densities of CH$_3$OH, CH$_3$CN and SO$_2$ are also derived. Constrains on the density of the gas are estimated for two of the protostellar cores.} {We find that the dust condensations are in various evolutionary stages. The powerhouse of the cluster, mm1a, harbours a hot core with $T\\sim 220~(75<T<330)$~K. A double-peaked profile is detected in several transitions toward mm1a, and we found a velocity gradient along a linear structure which could be perpendicular to one of the outflows from the vicinity of mm1a. Since the size of the double-peaked emission is less than 1100~AU, we suggest that mm1a might host a massive circumstellar disk. The other sources are in earlier stages of star formation. The least active source, mm3, could be a starless massive core, since it is cold ($T<20$~K), with a large reservoir of accreting material ($M\\sim 19~M_\\odot$), but no molecular emission peaks on it.} {} ",
        "introduction": "The last decade has seen  significant progress in the understanding of how high-mass stars form. Large samples of massive young stellar objects (YSOs) were studied with single-dish telescopes, to investigate their physical properties through the analysis of their (sub)mm continuum and molecular emission \\citep[e.g.,][]{1996A&A...308..573M,1998A&A...336..339M,2000A&A...355..617M,1997MNRAS.291..261W,1998MNRAS.301..640W,1999MNRAS.309..905W,2000A&A...357..637H,2001ApJ...552L.167Z,2005ApJ...625..864Z,2002ApJ...566..931S,2002ApJ...566..945B,2002A&A...383..892B,2004A&A...426...97F,2004A&A...417..115W,2005A&A...434..257W}. However, an intrinsic feature of high-mass stars is that they form in clusters, and that most of them are at large (several kpc) distances. Therefore, single-dish studies, as valuable as they are, lack the necessary spatial resolution to resolve single protostars and study the inner regions where high-mass star formation takes place. Interferometric observations started shedding light into the complex nature of high-mass star forming regions with the adequate spatial resolution \\citep[e.g.,][]{1997A&A...325..725C,1999A&A...345..949C,2005A&A...434.1039C,1999ApJ...514L..43W,2002A&A...387..931B,2005ApJ...628..800B,2006ApJ...649..888H,2000ApJ...535..833S}. However, the number of massive YSOs studied at high resolution is still too small to establish the general properties of the dense cores where massive stars form on a statistical base.  In this paper, we present an interferometric analysis of the high-mass star forming region \\object{IRAS\\,05358+3543} at (sub)mm wavelengths in several molecular transitions.  IRAS\\,05358+3543 (also known in literature as \\object{S233IR}) is part of a sample of 69 high-mass protostellar objects studied in great detail in recent years \\citep{2002ApJ...566..931S,2002ApJ...566..945B,2002A&A...383..892B,2002A&A...390..289B,2004A&A...417..115W,2005A&A...434..257W,2005A&A...442..949F}. At a distance of 1.8~kpc \\citep{1990ApJ...352..139S}, IRAS 05358+3543 has a bolometric luminosity of 6300~L$_{\\sun}$; strong high-mass star formation activity is evidenced by maser emission \\citep[see][]{1991ApJ...380L..75M,1995A&AS...112..299T,2000A&A...362.1093M} and outflow activity \\citep{1990ApJ...352..139S,2002A&A...387..931B}. Previous interferometric observations by \\citet{2002A&A...387..931B} resolved three dust condensations (mm1, mm2 and mm3) within an area of $9'' \\times 4''$ ($\\sim17\\,100\\times7\\,200$~AU), and revealed at least three outflows in CO and SiO, the most prominent of which is more than a parsec in length, and massive $(M >10~M_ {\\sun})$.  Two of the three identified outflows originate from the vicinity of mm1, which is probably the main powerhouse in the region.  To zoom in on the innermost region around the protostars, and study the physical properties of the individual potentially star-forming cores, we carried out a comprehensive program to observe the region at high spatial resolution with the Plateau de Bure Interferometer\\footnote{IRAM is supported by INSU/CNRS (France), MPG (Germany) and IGN (Spain).} at 97~GHz and 241~GHz, and the Submillimeter Array\\footnote{The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics and is funded by the Smithsonian Institution and the Academia Sinica.} at 338~GHz.  The new observations reach a resolution down to ~$0.6''$, corresponding to $\\sim 1100$~AU at the distance of the source. \\citet{05358-beuther} studied the continuum emission of this dataset. They identified four compact protostellar sources in the region; mm1 is resolved into two continuum peaks, mm1a and mm1b, with a projected linear separation of $\\sim 1700$~AU.  A mid-infrared source \\citep{2006MNRAS.369.1196L}, and a compact 3.6~cm continuum source \\citep{05358-beuther} coincide with mm1a, which is also associated with the class II methanol masers detected by \\citet{2000A&A...362.1093M}. The previously identified source mm2 resolves into several sub-sources; however, only one of them (mm2a, according to the nomenclature used by \\citealt{05358-beuther}) is a protostellar source, while the others are probably caused by the outflows in the region. The third source mm3 remains a single compact core even at the highest spatial resolution. In Table~\\ref{cores}, we report the positions of the four sources identified in the continuum emission by \\citet{05358-beuther}.  In this paper, we discuss the spectral line observations complementing the continuum data discussed by \\citet{05358-beuther}. In section \\S\\ref{observations}, the different observations are presented.  In section \\S\\ref{obs-res}, we discuss our results, and analyse the extended emission of low excitation molecular transitions (\\S\\ref{extended}), as well as the molecular spectra at the positions of the dust condensations(\\S\\ref{proto}). Finally, in section \\S\\ref{analysis} we derive the physical parameters of the gas around the protostars from the analysis of their spectra. In the following sections, we use the term protostar for young massive stellar objects which are still accreting material from the surroundings, independently whether they already started burning hydrogen or not. \\begin{table} \\centering \\caption{Positions of the four dust condensations in IRAS\\,05358+3543 \\citep[from][]{05358-beuther}.\\label{cores}} \\begin{tabular}{lcc} \\multicolumn{1}{c}{Source} &\\multicolumn{1}{c}{R.A. [J2000]}&\\multicolumn{1}{c}{Dec. [J2000]}\\\\ \\hline \\hline mm1a&05:39:13.08&35:45:51.3\\\\ mm1b&05:39:13.13&35:45:50.8\\\\ mm2a&05:39:12.76&35:45:51.3\\\\ mm3&05:39:12.50&35:45:54.9\\\\   \\hline \\hline \\end{tabular} \\end{table} ",
        "conclusions": "Our new interferometric data resolve at least four cores in the high-mass protocluster IRAS\\,05358+3543. By analysing  the molecular spectrum of each condensations, we characterised the properties of the gas surrounding it. Our main results are summarised in the following:  \\begin{itemize} \\item the main powerhouse of the region, mm1a, harbours a hot core with $T\\sim 220$~K, and the central heating source has a chemical timescale of $10^{4.6}$~yr. Our data suggest that mm1a might host a massive circumstellar disk; \\item  although the properties of the mm continuum emission of mm1b are very similar to the one of mm1a, the two sources differ significantly in the cm and mid-infrared spectrum. This core could be in an earlier stage of star formation than mm1a, since no molecular emission is detected toward it, with the only exception of the $5 \\to 4$~C$^{34}$S transition; \\item given the low abundance of methanol, mm2 could be a low--intermediate mass protostar; \\item strong emission is detected in several molecular species to the north-west of mm2, at a position where no continuum emission is detected. We suggest that this is caused by the interaction of the outflows with the ambient molecular cloud; \\item the least active source, mm3, could be a starless massive core, since it is cold ($T<20$~K), with a large reservoir of accreting material ($M\\sim 19~M_\\odot$), but no molecular emission peaks on it. \\end{itemize}"
    },
    "0710/0710.5714_arXiv.txt": {
        "abstract": "Fusion reactions in the crust of an accreting neutron star are an important source of heat, and the depth at which these reactions occur is important for determining the temperature profile of the star.  Fusion reactions depend strongly on the nuclear charge $Z$.  Nuclei with $Z\\le 6$ can fuse at low densities in a liquid ocean.  However, nuclei with $Z=8$ or 10 may not burn until higher densities where the crust is solid and electron capture has made the nuclei neutron rich.  We calculate the $S$ factor for fusion reactions of neutron rich nuclei including $^{24}$O + $^{24}$O and $^{28}$Ne + $^{28}$Ne.  We use a simple barrier penetration model.  The $S$ factor could be further enhanced by dynamical effects involving the neutron rich skin.  This possible enhancement in $S$ should be studied in the laboratory with neutron rich radioactive beams.  We model the structure of the crust with molecular dynamics simulations.  We find that the crust of accreting neutron stars may contain micro-crystals or regions of phase separation.  Nevertheless, the screening factors that we determine for the enhancement of the rate of thermonuclear reactions are insensitive to these features.  Finally, we calculate the rate of thermonuclear $^{24}$O + $^{24}$O fusion and find that $^{24}$O should burn at densities near $10^{11}$ g/cm$^3$.  The energy released from this and similar reactions may be important for the temperature profile of the star. ",
        "introduction": "Nuclei accreting onto a neutron star undergo a variety of reactions.  First at low densities, conventional thermonuclear fusion takes place, see for example \\cite{rpash}.  Next as nuclei are buried to higher densities, the rising electron Fermi energy induces a series of electron captures \\cite{gupta}.  Finally at very high densities, nuclei can fuse via pycnonuclear reactions.   These reactions are induced by the quantum zero point motion \\cite{pycno}.   The energy released, and the densities at which reactions occur, are important for determining the temperature profile of neutron star crusts.  Superbursts are very energetic X-ray bursts from accreting neutron stars that are thought to involve the unstable thermonuclear burning of carbon \\cite{superbursts, superbursts2}.  However, some simulations do not reproduce the conditions needed for carbon ignition because they have too low temperatures \\cite{superignition}.  An additional heat source, from fusion or other reactions, could raise the temperature and allow carbon ignition at densities that reproduce observed burst frequencies.  Recently the cooling of two neutron stars has been observed after extended outbursts \\cite{Wijnands, cackett}.  These outbursts heated the crusts out of equilibrium and then the cooling time was measured as the crusts returned to equilibrium.  The surface temperature of the neutron star in KS 1731-260 decreased with an exponential time scale of 325 $\\pm$ 100 days while MXB 1659-29 has a time scale of 505 $\\pm$ 59 days \\cite{cackett}.  These cooling times depend on the thermal conductivity of the crust and the initial temperature profile.  Comparing these observations of relatively rapid cooling to calculations by Rutledge et al. \\cite{rutledge} and Shternin et al. \\cite{shternin} suggests that the crust has a high thermal conductivity.  However, if the initial temperature profile of the crust is peaked near the surface, then this peak could quickly diffuse to the surface and lead to rapid cooling.  Therefore, cooling time scales are also sensitive to the initial temperature profile, and this depends on heating from nuclear reactions at moderate densities in the crust.  Gupta et al. have calculated heating from electron capture reactions in the outer crust \\cite{gupta}.  While they find more heating than previous works, they still find no more than 0.4 MeV per nucleon total heating from all of the electron captures on any mass number $A$ system.  Haensel and Zdunik have calculated pycnonuclear fusion reactions at great densities in the inner crust \\cite{haensel}.  However, if reactions occur deep in the inner crust, most of the heat may flow in to the core instead of out towards the surface.  As a result, there may be a smaller impact on the temperature profile of the outer crust.  A low crust thermal conductivity, for example from an amorphous solid, could help explain superburst ignition.  This could better insulate the outer crust and allow higher carbon ignition temperatures.  However, a low thermal conductivity appears to be directly contradicted by the observed short crust cooling times.  Furthermore, our molecular dynamics simulations in ref. \\cite{horowitz} and further results we present in Section \\ref{MD} find a regular crystal structure, even when the system has a complex composition with many impurities.  We do not find an amorphous phase.  These results will be discussed further in a later publication.  We conclude that the thermal conductivity of the crust is high.  If the thermal conductivity is high, one may need additional heat sources, at moderate densities, in order to explain superburst ignition.  Although Gupta et al. find additional heating from electron captures to excited nuclear states, simple nuclear structure properties may provide a natural limit to the total heating from electron captures \\cite{brownprivate}.  Haensel and Zdunik \\cite{haensel,haensel2007} consider heating from pycnonuclear reactions using a simple one component plasma model.  They find that fusion reactions may not take place until relatively high densities above $10^{12}$ g/cm$^3$.  However, their use of a one component plasma could be a significant limitation.  Fusion reactions depend strongly on the nuclear charge $Z$.  Therefore, the reaction rate may be highest for the rare impurities that have the lowest $Z$, instead of for nuclei of average charge.  In this paper, we go beyond Haensel and Zdunik and consider a full mixture of complex composition instead of assuming one average charge and mass.  We focus on thermonuclear and pycnonuclear reactions at densities around $10^{11}$ g/cm$^3$.  This is near the base of the outer crust.  Heat released at this density could be important for superburst ignition and for crust cooling times.  Nuclei at this density are expected to be neutron rich.  Furthermore, the other nearby ions strongly screen the Coulomb barrier and greatly enhance the rate of thermonuclear reactions.  We begin by describing the initial composition.  This includes neutron rich light nuclei such as $^{24}$O and $^{28}$Ne.  We calculate cross sections and $S$ factors for $^{24}$O + $^{24}$O and $^{28}$Ne + $^{28}$Ne fusion using a simple barrier penetration model.  Note that the dynamics of the neutron rich skins of these nuclei can enhance the cross section over that predicted by our simple barrier penetration model.  This is a very interesting and open nuclear structure question, see for example \\cite{subbarrier}.  Next, we use classical molecular dynamics simulations to determine the structure of the crust and screening factors for the enhancement of thermonuclear reactions.  There are many previous calculations of screening factors for the one component plasma \\cite{ocpscreening} and for binary ion mixtures, see for example \\cite{bimscreening}.  However, we are not aware of any previous calculations for a crystal of a complex multicomponent composition.   Finally, we calculate reaction rates and conclude that $^{24}$O is expected to fuse at densities near $10^{11}$ g/cm$^3$ while $^{28}$Ne should react at densities near $10^{12}$ g/cm$^3$.  Heat from these reactions may be important for determining the temperature profile of accreting neutron stars. ",
        "conclusions": "\\label{summary} Fusion reactions in the crust of an accreting neutron star are an important source of heat, and the depth at which these reactions occur is important for determining the temperature profile of the star.  Fusion reactions depend strongly on the nuclear charge $Z$.  Nuclei with $Z\\le 6$ can fuse at low densities in a liquid ocean.  However, nuclei with $Z=8$ or 10 may not burn until higher densities where the crust is solid and electron capture has made the nuclei neutron rich.  In Section \\ref{crosssections} we calculated the $S$ factor for fusion reactions of neutron rich nuclei including $^{24}$O + $^{24}$O and $^{28}$Ne + $^{28}$Ne.  We used a simple barrier penetration model.  We find that $S$ for $^{24}$O+$^{24}$O is over eight orders of magnitude larger than that for $^{16}$O+$^{16}$O.  The $S$ factor could be further enhanced by dynamical effects involving the neutron rich skin of $^{24}$O.  For example, the skins of the two nuclei could deform to form a neck that would reduce the Coulomb barrier.  This possible enhancement in $S$ should be studied in the laboratory with neutron rich radioactive beams.  In Section \\ref{MD} we modeled the structure of the crust with molecular dynamics simulations.  We find that the crust of accreting neutron stars may contain micro-crystals or regions of phase separation.  Nevertheless, the screening factors that we determined for the enhancement of the rate of thermonuclear reactions are insensitive to these features.  Finally, we calculated in Section \\ref{reaction rates} the rate of thermonuclear $^{24}$O + $^{24}$O fusion and find that $^{24}$O should burn at densities near $10^{11}$ g/cm$^3$.  This is a lower density than some previous estimates.  The 0.52 MeV per nucleon energy released may be important for the temperature profile of the star.  In future work, we will use our molecular dynamics results to study other properties of the crust such as its thermal conductivity.  In addition, we will use these MD results to calculate pycnonuclear reaction rates for the fusion of $^{28}$Ne and other heavier nuclei."
    },
    "0710/0710.2628_arXiv.txt": {
        "abstract": "We observed four southern AXPs in 1999 near 1400 MHz with the Parkes 64-m radio telescope to search for periodic radio emission. No Fourier candidates were discovered in the initial analysis, but the recent radio activity observed for the AXP XTE J1810$-$197 has prompted us to revisit these data to search for single radio pulses and bursts. The data were searched for both persistent and bursting radio emission at a wide range of dispersion measures, but no detections of either kind were made. These results further weaken the proposed link between rotating radio transient sources and magnetars. However, continued radio searches of these and other AXPs at different epochs are warranted given the transient nature of the radio emission seen from XTE J1810$-$197, which until very recently was the only known radio-emitting AXP. ",
        "introduction": "The detection of pulsed radio emission from the anomalous X-ray pulsar (AXP) XTE J1810$-$197 in 2006 \\cite{crh+06}, and more recently from a second AXP, 1E 1547.0$-$5408 \\cite{crh+07}, has renewed interest in searching for radio emission from these objects. In both of these cases, the radio activity is believed to be connected to the X-ray variability of the sources and is transient in nature (or at least highly variable). Given this transient behavior and that both persistent periodic emission and single pulses were detected from both AXPs, renewed searches of archival radio search data of AXPs at different epochs may reveal previously undetected radio signals from these sources. ",
        "conclusions": "We found no convincing radio signals in either the folding or single pulse searches. The derived upper limits on the radio emission from our AXP targets are presented in Table \\ref{tbl-1}. These are the most stringent radio upper limits to date for these sources. The estimated 1400 MHz luminosity limits on the periodic radio emission ($\\approxlt 1$ mJy kpc$^{2}$) are 2-3 times lower than those established for XTE J1810$-$197 prior to outburst. However, it is still conceivable that weak radio pulses are being emitted, but that they are below our detection threshold.  The luminosity limits presented here for the periodic emission are lower by about two orders of magnitude than the 1400 MHz luminosity of the periodic pulsed radio emission from XTE J1810$-$197 soon after the radio emission was first detected ($\\sim 80$ mJy kpc$^{2}$) \\citep{crh+06}. Thus we would expect to be able to easily detect comparably strong radio emission if it were beamed toward us.  Our luminosity limits on single radio pulses from our sources range from 22 to 69 Jy kpc$^{2}$ in the most conservative case, which is below the 1400 MHz luminosity of $\\approxgt 100$ Jy kpc$^{2}$ derived from the pulse strengths reported for XTE J1810$-$197 in its radio discovery paper \\citep{crh+06}. Since single pulses were detected from almost every rotation of XTE J1810$-$197, it is likely that we would have detected a large number of comparable pulses during our observations if such pulses were beamed toward us.  Our non-detection of single pulses further weakens the hypothesis that rotating radio transients (RRATs) \\citep{mlm+06} and magnetars are linked. This has been weakened by two other recent results. First, the X-ray detection of the RRAT J1819$-$1458 shows that its emission is more typical of middle-aged pulsars than it is of magnetars \\citep{rbg+06}. Second, the nearby, rotation-powered pulsar PSR B0656+54 would probably have been identified as an RRAT if it were farther away \\citep{wsr+06}.  We conclude from our results that any periodic or bursting radio emission from the four target AXPs is either very weak (below our detection thresholds), not beamed toward us, or non-existent or sporadic at the epoch of observation. This last possibility is suggested by the connection between the X-ray and radio activity observed for the two known radio-emitting magnetars to date. Continued radio searches of AXPs are therefore warranted given the apparent transient nature of the radio emission. Further details of this work and a more complete discussion of the results are presented in a recent journal article \\cite{chk07}.   \\begin{table} \\begin{tabular}{lcccc} \\hline & \\tablehead{1}{r}{b}{1E 1048.1$-$5937} & \\tablehead{1}{r}{b}{AX J1845$-$0258\\tablenote{AXP candidate only}} & \\tablehead{1}{r}{b}{1E 1841$-$045} & \\tablehead{1}{r}{b}{1RXS J170849.0$-$400910}   \\\\ \\hline Spin period (s)     & 6.45 & 6.97 & 11.77 & 11.00 \\\\ Ephemeris reference & \\cite{kgc+01} & \\cite{tkk+98}\\tablenote{No period derivative available} & \\cite{gvd99} & \\cite{gk02} \\\\ Galactic longitude, latitude (deg) & 288.26, $-$0.52 & 29.52, 0.07 & 27.39, $-$0.01 & 346.47, 0.03 \\\\ $T_{\\rm sky}$ (K)\\tablenote{1374 MHz sky temperature estimated from \\cite{hss+82} assuming a spectral index of $-2.6$} & 9.1 & 12.3 & 13.2 & 16.3 \\\\ Observation MJD & 51378 & 51391 & 51382 & 51379 \\\\ Observation date & 1999 Jul 19 & 1999 Aug 1 & 1999 Jul 23 & 1999 Jul 20 \\\\ $S_{1400}$ (mJy)\\tablenote{1400 MHz flux density limit on pulsed emission estimated using the modified radiometer equation and an assumed duty cycle of 2.7\\%} & $\\approxlt 0.02$ & $\\approxlt 0.02$ & $\\approxlt 0.02$ & $\\approxlt 0.02$ \\\\ $S_{1400}$ single (mJy)\\tablenote{Range of single-pulse 1400 MHz flux limits  for pulse time-scales 0.25-240 ms} & $\\approxlt 875$-50 & $\\approxlt 975$-60 & $\\approxlt 1000$-60 & $\\approxlt 1085$-65 \\\\ Distance (kpc)\\tablenote{Taken from \\cite{bri+06}. Question marks indicate significant uncertainty in the value} & $\\sim 5$? & $\\sim 8$? & $\\sim 7$ & $\\sim 8$? \\\\ $L_{1400}$ (mJy kpc$^{2}$)\\tablenote{1400 MHz luminosity limit on pulsed emission, assuming a 1 sr beaming fraction} & $\\approxlt 0.5$ & $\\approxlt 1.3$ & $\\approxlt 1.0$ & $\\approxlt 1.3$ \\\\ $L_{1400}$ single (Jy kpc$^{2}$)\\tablenote{Range of 1400 MHz luminosity limits on single pulses for pulse time-scales 0.25-240 ms} & $\\approxlt 22$-1.3 & $\\approxlt 62$-3.7 & $\\approxlt 49$-2.9 & $\\approxlt 69$-4.1 \\\\ \\hline \\end{tabular} \\caption{Radio Search Parameters and Results} \\label{tbl-1} \\end{table}"
    },
    "0710/0710.0288_arXiv.txt": {
        "abstract": "The channeling effect of low energy ions along the crystallographic axes and planes of NaI(Tl) crystals is discussed in the framework of corollary investigations on  WIMP Dark Matter candidates. In fact, the modeling of this existing effect implies a more complex evaluation of the luminosity yield for low energy recoiling Na and I ions. In the present paper related phenomenological arguments are developed and possible implications are discussed at some extent. ",
        "introduction": "It is known that ions (and, thus, also recoiling nuclei) move in a crystal in a different way than in amorphous materials. In particular, ions moving (quasi-) parallel to crystallographic axes or planes feel the so-called ``channeling effect'' and show an anomalous deep penetration into the lattice of the crystal \\cite{chan,nel63,oth2}; see Fig. \\ref{fg:schema_chan}.  For example, already on 1957, a penetration of $^{134}$Cs$^+$ ions into a Ge crystal was observed to a depth of about 1000 \\AA \\, \\cite{Bre56}, larger than that expected in the case the ions would cross amorphous Ge ($\\simeq 50$ \\AA). Afterwards, high intensities of H$^+$ ions at 75 keV transmitted through thick (3000-4000 \\AA) single-crystal gold films in the $<110>$ directions were detected \\cite{nel63}. Other examples for keV range ions have been shown in ref. \\cite{Ras05} where 3 keV P$^+$ ions moving into layers of 500 \\AA \\, of various crystals were studied.  The channeling effect is also exploited in high energy Physics e.g. to extract high energy ions from a beam by means of bent crystals or to study diffractive Physics by analysing scattered ions along the beam direction (see e.g. ref. \\cite{Baur00}).  Recently \\cite{droby2} it has been pointed out the possible role which this effect can play in the evaluation of the detected energy of recoiling nuclei in crystals, such as the NaI(Tl)\\footnote{For completeness, it is worth to note that luminescent response for channeling in NaI(Tl) was already studied in ref. \\cite{lunt} for MeV-range ions.}.  \\begin{figure} [!t] \\centering \\includegraphics[width=9.0cm] {fig1.eps} \\caption{Simplified schema of the channeling effect in the NaI(Tl) lattice. The axial channeling occurs when the angle of the motion direction of an ion with the respect to the crystallographic axis is less than a characteristic angle, $\\Psi_c$, depicted there (see for details Sec. 2). Two examples for channeled and unchanneled ions are also shown (dashed lines).} \\label{fg:schema_chan} \\end{figure}  In fact, the channeling effect can occur in crystalline materials due to correlated collisions of ions with target atoms. In particular, the ions through the open channels have ranges much larger than the maximum range they would have if their motion would be either in other directions or in amorphous materials. Moreover, when a low-energy ion goes into a channel, its energy losses are mainly due to the electronic contributions. This implies that a channeled ion transfers its energy mainly to electrons rather than to the nuclei in the lattice and, thus, its quenching factor (namely the ratio between the detected energy in keV electron equivalent [keVee] and the kinetic energy of the recoiling nucleus in keV) approaches the unity.  It is worth to note that this fact can have a role in corollary analyses in the Dark Matter particle direct detection experiments, when WIMP (or WIMP-like) candidates are considered. In fact, since the routine calibrations of the detectors are usually performed by using $\\gamma$ sources (in order to avoid induced radioactivity in the materials), the quenching factor is a key quantity to derive the energy of the recoiling nucleus after an elastic scattering. Generally, for scintillation and ionization detectors this factor has been inferred so far by inducing tagged recoil nuclei through neutron elastic scatterings \\cite{Mis_neut}; however, as it will be discussed in Sec. 3, the usual analysis carried out on similar measurements does not allow to account for the channeled events. A list of similar values for various nuclei in different detectors can be found e.g. in ref. \\cite{RNC}. In particular, commonly in the interpretation of the dark matter direct detection results in terms of WIMP (or WIMP-like) candidates the quenching factors are assumed to be constant values without considering e.g. their energy dependence, the properties of each specific used detector and the experimental uncertainties. An exception was in the DAMA/NaI corollary model dependent analyses for WIMP (or WIMP-like) candidates \\cite{RNC,ijmd,epj06,ijma2} where at least some of the existing uncertainties on the $q_{Na}$ and $q_I $ values, measured with neutrons, were included.  In this paper the possible impact of the channeling effect in NaI(Tl) crystals is discussed in a phenomenological framework and comparisons on some of the corollary analyses carried out in terms of WIMP (or WIMP-like) candidates \\cite{RNC,ijmd,epj06,ijma2}, on the basis of the 6.3 $\\sigma$ C.L. DAMA/NaI model independent evidence for particle Dark Matter in the galactic halo\\footnote{We remind that various possibilities for some of the many possible astrophysical, nuclear and particle Physics scenarios have have been analysed by DAMA itself both for some WIMP/WIMP-like candidates and for light bosons \\cite{RNC,ijmd,epj06,ijma2,ijma}, while other corollary analyses are also available in literature, such as e.g. refs. \\cite{Bo03,Bo04,Botdm,khlopov,Wei01,foot,Saib,droby1,droby2}. Many other scenarios can be considered as well.}, are given. ",
        "conclusions": "In this paper the channeling effect of recoiling nuclei induced by WIMP and WIMP-like elastic scatterings in NaI(Tl) crystals has been discussed. Its possible effect in a reasonably cautious modeling has been presented as applied to some given simplified scenarios in corollary quests for the candidate particle for the DAMA/NaI model independent evidence. This further shows the role of the existing uncertainties and of the correct description and modeling of all the involved processes as well as their possible impact in the investigation of the candidate particle. Some of them have already been addressed at some extent, such as the halo modeling \\cite{halo,RNC,ijmd}, the possible presence of non-thermalized components in the halo (e.g. caustics \\cite{siki} or SagDEG \\cite{epj06} contributions), the accounting for the electromagnetic contribution to the WIMP (or WIMP-like) expected energy distribution \\cite{ijma2}, candidates other than WIMPs (e.g. \\cite{ijma} and in literature), etc..  Obviously, many other arguments can be addressed as well both on DM candidate particles and on astrophysical, nuclear and particle physics aspects; for more see \\cite{RNC,ijmd,ijma,epj06,ijma2} and in literature. In particular, we remind that different astrophysical, nuclear and particle Physics scenarios as well as the experimental and theoretical associated uncertainties leave very large space also e.g. for significantly lower cross sections and larger masses."
    },
    "0710/0710.0846_arXiv.txt": {
        "abstract": "s{ The Alpha Magnetic Spectrometer (AMS), to be installed on the International Space Station (ISS) in 2008, is a cosmic ray detector with several subsystems, one of which is a proximity focusing Ring Imaging \\CK\\ (RICH) detector. This detector will be equipped with a dual radiator (aerogel+NaF), a lateral conical mirror and a detection plane made of 680 photomultipliers and light guides, enabling precise measurements of particle electric charge and velocity. Combining velocity measurements with data on particle rigidity from the AMS Tracker it is possible to obtain a measurement for particle mass, allowing the separation of isotopes. \\\\ A Monte Carlo simulation of the RICH detector, based on realistic properties measured at ion beam tests, was performed to evaluate isotope separation capabilities. Results for three elements --- H (Z=1), He (Z=2) and Be (Z=4) --- are presented. }  \\vspace{-0.9cm} ",
        "introduction": " ",
        "conclusions": "AMS will provide a major improvement on existing data for isotopic abundances in cosmic rays. Simulation results indicate that the separation of light isotopes using the combination of RICH data and tracker rigidity measurements is feasible. The dual radiator configuration of NaF and aerogel makes isotope separation of light elements possible for energies in the range from 0.5 to 10 GeV/nucleon, approximately. Best mass resolutions are $\\sim$~2\\% at 3 GeV/nucleon for aerogel, and $\\sim$~3\\% at 1 GeV/nucleon for NaF. Techniques presented here may also be applied in the separation of antimatter isotopes which is of great importance in dark matter studies.  \\vspace{-0.2cm}"
    },
    "0710/0710.2558_arXiv.txt": {
        "abstract": "We highlight the potential importance of gaseous TiO and VO opacity on the highly irradiated close-in giant planets.  The atmospheres of these planets naturally fall into two classes that are somewhat analogous to the M- and L-type dwarfs.  Those that are warm enough to have appreciable opacity due to TiO and VO gases we term the ``pM Class'' planets, and those that are cooler, such that Ti and V are predominantly in solid condensates, we term ``pL Class'' planets.  The optical spectra of pL Class planets are dominated by neutral atomic Na and K absorption.  We calculate model atmospheres for these planets, including pressure-temperature profiles, spectra, and characteristic radiative time constants.  Planets that have temperature inversions (hot stratospheres) of $\\sim$2000 K and appear ``anomalously'' bright in the mid infrared at secondary eclipse, as was recently found for planets \\hh\\ and \\hd, we term the pM Class.  Molecular bands of TiO, VO, H$_2$O, and CO will be seen in emission, rather than absorption.  This class of planets absorbs incident flux and emits thermal flux from high in their atmospheres.  Consequently, they will have large day/night temperature contrasts and negligible phase shifts between orbital phase and thermal emission light curves, because radiative timescales are much shorter than possible dynamical timescales.  The pL Class planets absorb incident flux deeper in the atmosphere where atmospheric dynamics will more readily redistribute absorbed energy.  This leads to cooler day sides, warmer night sides, and larger phase shifts in thermal emission light curves.  We briefly examine the transit radii for both classes of planets.  The boundary between these classes is particularly dependent on the incident flux from the parent star, and less so on the temperature of the planet's internal adiabat (which depends on mass and age), and surface gravity.  Around a Sun-like primary, for solar composition, this boundary likely occurs at $\\sim$0.04-0.05 AU, but uncertainties remain.  We apply these results to pM Class transiting planets that are observable with the \\emph{Spitzer Space Telescope}, including \\hd, WASP-1b, TrES-3b, TrES-4b, \\hh, and others.  The eccentric transiting planets HD 147506b and HD 17156b alternate between the classes during their orbits.  Thermal emission in the optical from pM Class planets is significant red-ward of 400 nm, making these planets attractive targets for optical detection via Kepler, COROT, and from the ground.  The difference in the observed day/night contrast between $\\upsilon$ Andromeda b (pM Class) and \\he\\ (pL Class) is naturally explained in this scenario. ",
        "introduction": "The blanket term ``hot Jupiter\" or even the additional term ``very hot Jupiter'' belies the diversity of these highly irradiated planets.  Each planet likely has its own unique atmosphere, interior structure, and accretion history.  The relative amounts of refractory and volatile compounds in a planet will reflect the parent star abundances, nebula temperature, total disk mass, location of the planet's formation within the disk, duration of its formation, and its subsequent migration (if any).  This accretion history will give rise to differences in core masses, total heavy elements abundances, and atmospheric abundance ratios.  Given this incredible complexity, it is worthwhile to first look for physical processes that may be common to groups of planets.  In addition to a mass and radius, one can further characterize a planet by studying its atmosphere.  The visible atmosphere is a window into the composition of a planet and contains clues to its formation history \\cp[e.g.,][]{Marley07b}.  Of premier importance in this class of highly irradiated planets is how stellar insolation affects the atmosphere, as this irradiation directly affects the atmospheric structure, temperatures, and chemistry, the planet's cooling and contraction history, and even its stability against evaporation.  Since irradiation is perhaps the most important factor in determining the atmospheric properties of these planets, we examine the insolation levels of the 23 known transiting planets.  We restrict ourselves to those planets more massive than Saturn, and hence for now exclude treatment of the ``hot Neptune'' GJ 436b, which is by far the coolest known transiting planet.  \\mbox{Figure~\\ref{flux}} illustrates the stellar flux incident upon the planets as a function of both planet mass (\\mbox{Figure~\\ref{flux}}\\emph{a}) and planet surface gravity (\\mbox{Figure~\\ref{flux}}\\emph{b}).  In these plots diamonds indicate transiting planets and triangles indicate other interesting hot Jupiters, for which \\emph{Spitzer Space Telescope} data exist, but which do not transit.  The first known transiting planet, \\hd, is seen to be fairly representative of these planets in terms of incident flux.  Planets OGLE-TR-56b and OGLE-TR-132b are somewhat separate from the rest of the group because they receive the highest stellar irradiation.  Both orbit their parent stars in less than 2 days and are prototypes of what has been called the class of ``very hot Jupiters'' \\cp{Konacki03,Bouchy04} with orbital periods less than 3 days.  However, orbital period is a poor discriminator between ``very hot'' and merely ``hot,'' as \\he\\ clearly shows.  Labeled a ``very hot Jupiter'' upon its discovery, due to its short 2.2 day period \\citep{Bouchy05}, \\he\\ actually receives a comparatively modest amount of irradiation due to its relatively cool parent star.  Therefore, perhaps a classification based on incident flux, equilibrium temperature, or other attributes would be more appropriate.  In this paper we argue that based on the examination of few physical processes that two classes of hot Jupiter atmospheres emerge with dramatically different spectra and day/night contrasts.  Equilibrium chemistry, the depth to which incident flux will penetrate into a planet's atmosphere, and the radiative time constant as a function of pressure and temperature in the atmosphere all naturally define two classes these irradiated planets.  Our work naturally builds on the previous work of \\ct{Hubeny03} who first investigated the effects of TiO and VO opacity on close-in giant planet atmospheres as a function of stellar irradiation.  These authors computed optical and near infrared spectra of models with and without TiO/VO opacity.  In general they found that models with TiO/VO opacity feature temperature inversions and molecular bands are seen in emission, rather than absorption.  Two key questions from the initial \\ct{Hubeny03} investigation were addressed but could not be definitely answered were: 1) if a relatively cold planetary interior would lead to Ti/V condensing out deep in the atmosphere regardless of incident flux, thereby removing gaseous TiO and VO, and, 2) if this condensation did not occur, at what irradiation level would TiO/VO indeed be lost at the lower atmospheric temperatures found at smaller incident fluxes.  Later \\ct{Fortney06} investigated model atmospheres of planet \\hh\\ including TiO/VO opacity at various metallicities.  Particular attention was paid to the temperature of the deep atmosphere \\emph{P-T} profiles (as derived from an evolution model) in relation to the Ti/V condensation boundary.  Similar to \\ct{Hubeny03}, they found a temperature inversion due to absorption by TiO/VO and computed near and mid-infrared spectra that featured emission bands.  Using the \\emph{Spitzer} InfraRed Array Camera (IRAC) \\ct{Harrington07} observed \\hh\\ in secondary eclipse with \\emph{Spitzer} at 8 $\\mu$m and derived a planet-to-star flux ratio consistent with a \\ct{Fortney06} model with a temperature inversion due to TiO/VO opacity.  At that point, looking at the work of \\ct{Fortney06} and especially \\ct{Hubeny03}, \\ct{Harrington07} could have postulated that all objects more irradiated than \\hh\\ may possess inversions due to TiO/VO opacity, but given the single-band detection of \\hh, caution was in order.  More recently, based on the four-band detection of flux from \\hd\\ by \\ct{Knutson08}, \\ct{Burrows07c} find that a temperature inversion, potentialy due to TiO/VO opacity, is necessary to explain this planet's mid-infrared photometric data.  Based on their new \\hd\\ model and the previous modeling investigations these authors posit that planets warmer than \\hd\\ may features inversions, while less irradiated objects such as \\he\\ do not, and discuss that photochemical products and gaseous TiO/VO are potential absorbers which may lead to this dichotomy.  We find, as has been previously shown, that those planets that are warmer than required for condensation of titanium (Ti) and vanadium (V)-bearing compounds will possess a temperature inversion at low pressure due to absorption of incident flux by TiO and VO, and will appear ``anomalously'' bright in secondary eclipse at mid-infrared wavelengths.  Thermal emission in the optical will be significant \\cp{Hubeny03,Lopez07}.  Furthermore, here we propose that these planets will have large day/night effective temperature contrasts.  We will term these very hot Jupiters the ``pM Class,'' meaning gaseous TiO and VO are the prominent absorbers of optical flux.  The predictions of equilibrium chemistry for these atmospheres are similar to dM stars, where absorption by TiO, VO, H$_2$O, and CO is prominent \\cp{Lodders02b}.  Planets with temperatures below the condensation curve of Ti and V bearing compounds will have a gradually smaller mixing ratio of TiO and VO, leaving Na and K as the major optical opacity sources \\cp*{BMS}, along with H$_2$O, and CO.  We will term these planets the pL class, similar to the dL class of ultracool dwarfs.  These planets will have relatively smaller secondary eclipse depths in the mid infrared and significantly smaller day/night effective temperature contrasts.  As discussed below, published \\emph{Spitzer} data are consistent with this picture.  The boundary between these classes, at irradiation levels (and atmospheric temperatures) where Ti and V may be partially condensed is not yet well defined.  In this paper we begin by discussing the observations to date.  We then give an overview of our modeling methods and the predicted chemistry of Ti and V.  We calculate pressure-temperature (\\emph{P-T}) profiles and spectra for models planets.  For these model atmospheres we then analyze in detail the deposition of incident stellar flux and the emission of thermal flux, and go on to calculate characteristic radiative time constants for these atmospheres.  We briefly examine transmission spectra before we apply our models to known highly irradiated giant planets.  Before our discussion and conclusions we address issues of planetary classification. ",
        "conclusions": "Though 1D radiative-convective equilibrium model atmospheres we have addressed the class of atmospheres, the pM Class, for which TiO and VO are extremely strong visible absorbers \\cp{Hubeny03}.  This absorbed incident flux drives these planets to have hot ($\\sim$2000+ K) stratospheres.  Therefore, these planets will appear very bright in the mid-infrared, with brightness temperatures larger than their equilibrium temperatures.  This is the case for \\hh\\ \\cp{Harrington07}, and \\hd\\ \\ct{Knutson08}.  In addition, these planets will have large day/night temperature contrasts because radiative time constants at photospheric pressures are much shorter than reasonable advective timescales.  The hottest point of the planet should be the substellar point, which absorbs the most flux, leading to perhaps negligible phase shift between the times of maximum measured thermal emission and when the day side is fully visible.  This appears to be the situation for planet $\\upsilon$ And b, observed by \\ct{Harrington06}.  Given that its irradiation level is intermediate between \\hd\\ and \\hh\\ we find that this planet is pM Class.  Due to the fast radiative times, the day side of these planets may have \\emph{P-T} profiles that do not deviate much from radiative equilibrium models.  Although atmospheric dynamics will surely be vigorous, winds will be unable to advect gas before it cools to space.  For these planets, high irradiation, the presence of gaseous TiO and VO, a hot stratosphere, the location of the hottest atmosphere at the substellar point, and a large day/night temperature contrast all go hand-in-hand.  On the other hand, we have shown that in the pL Class, dominated by absorption by H$_2$O, Na, and K, photospheric pressures and temperatures prevail such that advective timescales and radiative timescales are similar \\cp[see also][]{Seager05}.  Since atmospheric dynamics will be important for the redistribution of energy, the consequences for the structure and thermal emission of these atmospheres will be quite complex.  The efficiency of energy redistribution will vary with planetary irradiation level, surface gravity, and rotation rate.  pL Class planets will have smaller day/night temperature contrasts and measurable phase shifts in thermal emission light curves that will be wavelength-dependent.  In addition, these planets may show variability in secondary eclipse depth \\cp[e.g.~][]{Rauscher07}.  However, without a better understanding of the dynamics it is difficult to make detailed predictions at this time.  Secondary eclipse depths should range somewhere between values expected for a ``full redistribution'' model and inefficient redistribution.  The published secondary eclipse data for pL Class planets \\T\\ and \\he\\ are all consistent with this prediction \\cp{Fortney05,Fortney07b}.  In addition, the 8 $\\mu$m light curves for 51 Peg b and \\hd\\ \\cp{Cowan07} and \\he\\ \\cp{Knutson07b} are consistent with this prediction as well.  Examination of \\mbox{Figure~\\ref{flux}} shows that HD 179949b is a pM Class planet, and indeed \\ct{Cowan07} found the largest phase variation it their small sample for this planet, but the unknown orbital inclination makes definitive conclusions difficult.  Additional observational results will soon help to test the models presented here.  We find that transiting planets WASP-1b, TrES-4b, TrES-3b, OGLE-Tr-10b, and TrES-2b will be in the pM Class, along with non-transiters $\\upsilon$ And b and HD 179949.  The low-irradiation boundary of this class is not yet clear, and planet \\hd\\ shows that temperature inversions persist to irradiation levels where TiO/VO are expected to begin being lost to condensation \\cp{Burrows07c}.  At still lower irradiation levels, the limited data for \\he\\ lead us to conclude it is pL Class \\cp{Fortney07b}.  Secondary eclipse data for XO-2b, HAT-P-1b, and WASP-2b will be important is determining how temperature inverstions (and the TiO/VO abundances) wane with irradiation level.  Just as in dM and dL stars, the condensation of Ti and V is expected to be gradual process, so we fully expect transition objects between the distinct pM and pL class members.  We will soon have additional information that will help shed light on the atmosphere of \\hd\\.  The 8 $\\mu$m light curve for \\hd\\ obtained by \\ct{Cowan07} shows little phase variation.  However, if our theory connecting TiO/VO opacity and temperature inversions to large day/night contrasts is correct, we expect to see a large variation. Soon H.~Knutson and collaborators will obtain half-orbit light curves for \\hd\\ at 8 and 24 $\\mu$m.  The quality should be comparable to that obtained by \\citet{Knutson07b} for \\he, and will put our theory to the test.  HD 147506, with an eccentricity of 0.517 \\cp{Bakos07b}, and HD 17156, with an eccentricity of 0.67 \\cp{Fischer07,Barbieri07}, will be extremely interesting cases as the flux they receive varies by factors of 9 and 26, respectively, between periapse and apoapse.  They should each spend part of their orbits as pL class and part as pM class.  This makes predictions difficult, but large day/night temperatures differences at apoapse are likely.  There has recently been considerable discussion on the relative merits of multi-dimensional dynamical models and 1D radiative-convective model atmospheres for these highly irradiated planets.  Both kinds of studies provide interesting predictions.  While it could be claimed that 1D radiative-convective models are unrealistic because they ``lack dynamics,'' they do include very detailed chemistry, vast opacity databases, and advanced non-gray radiative transfer, which all dynamics models for these planets lack.  Given the large diversity in predictions among the various dynamical models, which use a host of simplifications, the next step will be combining dynamics and radiative transfer, an idea which has been mentioned or advanced by a number of authors \\cp{Seager05,Barman05,Fortney06b,Burrows06,Dobbs07}.  Eventually we will be able to give up 1D $f$-type parameters to treat the incident flux in a more realistic fashion for these exotic atmospheres.  We recently started working toward this goal in \\ct[][see also \\citealp{Dobbs07}]{Fortney06b}; work continues, and we believe it will have a promising future.  We think that the predictions we have made here with a 1D model will provide a framework for understanding the observations to come.  Additional observations for both pL and pM class planets, along with additional theoretical and modeling efforts, should further clarify our understanding."
    },
    "0710/0710.3148_arXiv.txt": {
        "abstract": "In this paper, warm inflationary models on a brane scenario are studied. Here we consider slow-roll inflation and high-dissipation regime in a high-energy scenario. General conditions required for these models to be realizable are derived. We describe scalar and tensor perturbations for these scenarios. Specifically we study power-law potentials considering a dissipation parameter to be a constant on the one hand and $\\phi$ dependent on the other hand. We use recent astronomical observations to restrict the parameters appearing in our model. ",
        "introduction": "It is well known that many long-standing problems of the Big Bang model, namely the horizon problem, flatness, homogeneity and the numerical density of monopoles, may find a natural explanation in the frame of the inflationary universe model \\cite{Guth1981,Inflation,Inflationary}. Perhaps the most relevant feature of the inflationary universe model is that it provides a causal interpretation for the origin of the observed anisotropy in the cosmic microwave background (CMB) radiation, and also the distribution of large-scale structures \\cite{WMAP1,WMAP3}. But the inflationary universe model has problems too. One of the problems in it is how to attach the observed universe to the end of the inflationary epoch; there are three possible solutions to this problem: reheating \\cite{KolbandTurner}, preheating \\cite{TBKLS} and warm inflation \\cite{Berera1995,deOliveira1998,Berera1999,Bellini1998}. In this work we focus on the latter.  In standard inflationary universe models, the acceleration of the universe is driven by a scalar field (named inflaton) with a specific scalar potential. These kinds of model are divided into two regimes, the slow-roll and reheating ephocs.  In the slow-roll period the universe inflates and all interactions between the inflaton scalar field and any other field are typically neglected. Subsequently, a reheating period is invoked to end the period of inflation. After reheating, the universe is filled with radiation \\cite{afterReheating,Inflation}, and then the universe gets connected with the Big Bang model.  Warm inflation is an alternative mechanism to have successful inflation and avoid the reheating period \\cite{Berera1995}. In this kind of model, dissipative effects are important during inflation, so that radiation production occurs concurrently with the inflationary expansion. The inflaton interacts with a thermal bath via a friction term, where phenomenologically the decay of the scalar field is described by means of an interaction Lagrangian. For instance, the authors of Ref.\\cite{new} take the interaction terms of the form $\\frac{1}{2}\\lambda^2\\phi^2\\chi^2$ and $g\\chi\\bar{\\psi}\\psi$, where the inflationary period presents a two-stage decay chain $\\phi\\rightarrow\\chi\\rightarrow\\psi$. In this case, they reported that the damping term $\\Gamma$ becomes $\\frac{\\lambda^3g^2\\phi}{256\\pi^2}$. Note that if the scalar field changes a little bit, then the friction coefficient $\\Gamma$ remains almost constant. From the point of view of statistical mechanics, the interaction between quantum fields and a thermal bath could be illustrated by a general fluctuation-dissipation relation \\cite{new2}. Warm inflation was criticized on the basis that the inflaton cannot decay during the slow-roll phase \\cite{new3}. However, in recent years, it has been shown that the inflaton can indeed decay during the slow-roll phase (see \\cite{new4} and references therein) whereby it now rests on solid theoretical grounds. On the other hand, the inclusion of the damping term into the model needs a very special scheme: for instance, in Ref.\\cite{Profe} it was found that when $\\Gamma$ is constant, in the slow-roll approximation a wrong value for the number of e-fold was obtained: meanwhile, in the power-law approach this problem is absent.  Warm inflation ends when the universe heats up to become radiation dominated. At this epoch the universe stops inflating and `smoothly' enters into a radiation dominated Big Bang phase \\cite{Berera1995}. The matter components of the universe are created by the decay of either the remaining inflationary field or the dominant radiation field \\cite{TaylorandBerera2000}.  In standard inflationary universe models the quantum fluctuations associated to the inflaton scalar field generate the density perturbations seeding the structure formation at a late time in the evolution of the universe. Instead, in warm inflation models, the density fluctuations arise from thermal rather than quantum fluctuations \\cite{Berera2004,Berera1995}. These fluctuations have their origin in the hot radiation and influence the inflaton through a friction term in the equation of motion of the inflaton scalar field \\cite{Berera1996}.  Several aspects of warm inflationary universe models have been studied in the past few years. The goal of the present work is to investigate warm inflationary models on brane scenarios, where the total energy density $\\rho=\\rho_{\\phi}+\\rho_{\\gamma}$ is found on the brane \\cite{Profe2}. The universe is filled with a self-interacting scalar field of energy density $\\rho_{\\phi}$ and a radiation field with energy density $\\rho_{\\gamma}$.  The motivation for introducing brane scenarios is the increasing interest in higher dimensional cosmological models, motivated by superstring theory, where the matter fields (related to open string modes) are confined to a lower dimensional brane, while gravity (closed string modes) can propagate in the bulk \\cite{Strings}. Shiromizu et al. \\cite{Shiromizu} have found the four-dimensional Einstein's equations projected onto the brane. These projections introduce some differences in the fundamental field equations, such as the Friedmann equation and, therefore, in the equations that describes the linear perturbations theory \\cite{Brane}. Our aim is quantify the modifications of the warm inflation model in the brane scenario for arbitrary inflaton potentials on the brane. In order to do this we study the linear theory of cosmological perturbations for a warm inflationary model on a brane. The perturbations are expressed in term of different parameters appearing in our model: these parameters will be constrained from the WMAP three-year data \\cite{WMAP3}.  In section II we describe the dynamics of our model and we establish some approximations that we use in this work. In section III we investigate the linear theory of perturbations. Here we calculate the scalar perturbations in the longitudinal gauge and also the tensor perturbations. In section IV we take a power-law potential and we investigate the high-energy and high-dissipation regime considering a power-law dissipation coefficient $\\Gamma$. Finally, in section V, we give some conclusions. We have used units in which $c=\\hbar=1$. ",
        "conclusions": "In this paper we have considered a warm inflationary scenario on a brane. We have restricted ourselves to a high-dissipation and high-energy regime. In the slow-roll approximation we have found a general relationship between radiation and scalar field densities; see Eq.(\\ref{eq21}).  In relation to the perturbations we have considered that the there does no exist any interaction between the brane and the bulk, i.e., we have neglected back-reaction due to metric perturbations in the fifth dimension. We note that a full investigation is required to discover when back-reaction will have a significant effect in the perturbations. With the above-mentioned restriction we have obtained explicitly the contributions of the adiabatic and entropy modes. We have shown that the dissipation parameter plays a crucial role in producing the entropy mode (see Eq.(\\ref{eq62b})). A general relation for the density perturbations is given in Eq.(\\ref{eq41}). The tensor pertubations are generated via stimulated emission into the existing thermal background (see Eq.(\\ref{eq52})) and the tensor-scalar ratio is modified by a temperature-dependent factor.  We have studied a power-law potential for different dependence of the dissipation coefficient $\\Gamma$. From the normalization of the WMAP three-year data, the potential becomes of the order of $V(\\phi_0)\\sim10^{-20}M_4^4$ when it leaves the horizon at the scale of $k_0=0.002\\text{Mpc}^{-1}$. As in the situation studied in Ref.\\cite{Chaotic}, the value of the potential depends on $M_5$. Here we have considered $M_5\\sim10^{-5}M_4$. To fulfill the approximations in our model we have restricted the range of the parameters in which warm inflation on a brane can occur (see FIGS. \\ref{fig:HT}, \\ref{fig:VM5}). In order to show some explicit results we have chosen the following set of values for the parameters, $T_r\\simeq T=3.5\\times10^{-6}M_4$, $r=1000$ and $s=100$. First, we have considered a chaotic potential and a constant dissipation coefficient; for this case we have found that the spectrum is driven toward an scale-invariant spectrum and the running of the spectral index is extremely small. The situation was similar when we considered a variable dissipation coefficient. In the case of a power-law potential with $n=4$ and a constant dissipation coefficient we have found that the value to $n_s$ is smaller than the previous case but out of the range given by WMAP three-year data. On the other hand, the running of the spectral index is closer to the value given by WMAP three-year data but one order of magnitud smaller.  Finally the best situation is found when we consider a constant dissipation coefficient and a $n=6$ potential. In this case both parameters, $n_s$ and $\\alpha_s$, are in the ranges specified by WMAP.  It is necessary to note that with our approximations we have found that for the case in which $n=6$ and $m=0$ is favored in the light of the recent results reported by WMAP three-year data. On the other hand, we have considered some approximations; as a result, we were not able to find any real solution in the allowed range of parameters. For example, in the slow-roll approximation, in the case $n=4$ and, when we try to consider a variable coefficient of dissipation $\\Gamma$, we did not find any real solution. However, we think that we could find a solution if in place of using this approximation we assume a power-law for the scale factor. We intend to return to this point (and others) in the near future."
    },
    "0710/0710.1651_arXiv.txt": {
        "abstract": "We report on an investigation of the environments of the SLACS sample of gravitational lenses. The local and global environments of the lenses are characterized using SDSS photometry and, when available, spectroscopy. We find that the lens systems that are best modelled with steeper than isothermal density profiles are more likely to have close companions than lenses with shallower than isothermal profiles. This suggests that the profile steepening may be caused by interactions with a companion galaxy as indicated by N-body simulations of group galaxies. The global environments of the SLACS lenses are typical of non-lensing SDSS galaxies with comparable properties to the lenses, and the richnesses of the lens groups are not as strongly correlated with the lens density profiles as the local environments. Furthermore, we investigate the possibility of line-of-sight contamination affecting the lens models but do not find a significant over-density of sources compared to lines of sight without lenses. ",
        "introduction": "The Sloan Lens ACS Survey \\citep[SLACS;][]{bolton} is a large sample of strong gravitational lenses derived from the Sloan Digital Sky Survey (SDSS). The lens sample has proved to be particularly useful due to the high quality of the data; all of the SLACS lenses have known lens and source redshifts, stellar velocity dispersions for the lensing galaxies have been measured from the SDSS spectra, all of the systems have {\\em Hubble Space Telescope} ({\\em HST}) ACS imaging in two bands for accurate non-parametric lens modelling, and all of the sources are extended and therefore provide additional constraints on the mass profile of the lensing galaxy. Furthermore, the density of background sources in the {\\em HST} imaging allows a weak lensing analysis of the ensemble sample of lenses \\citep{gavazzi}; the major drawback of the lens sample is that it is unlikely to find any variable sources that would provide time delays.  \\citet{koopmansSLACS} have used the SLACS lenses to show that early-type galaxies have isothermal total inner density profiles with very little intrinsic scatter (approximately 6 per cent) assuming a uniformity in the environments of the lenses that has not been rigorously tested.  While the SLACS lenses seem to lie on the Fundamental Plane \\citep{bolton,treu,bolton07} and do not differ noticeably from other SDSS galaxies with similar luminosities and stellar velocity dispersions, the local environments of the lenses might affect the mass profiles of the lensing galaxies \\citep[e.g.,][]{rusin,dobke,augerb}. If some of the lenses are being perturbed by neighbouring galaxies the intrinsic scatter of the density slope for {\\em isolated} early-type galaxies might be even smaller than 6 per cent. Furthermore, it has been suggested that line-of-sight (LOS) contamination significantly affects the SLACS lenses \\citep{guimaraes} and the density slope might be expected to be shallower than originally reported.  We report on a spectroscopic and photometric evaluation of the environments and lines of sight of the 15 SLACS lenses investigated by \\citet{koopmansSLACS}. A weighting scheme is used to determine the effective number of potential perturbing companions to each lens galaxy. We also characterize the `richness' of the global environment of each lens field and quantify the number of galaxies along the LOS to the lens. Throughout this paper the term `global environment' is used to describe the group, cluster, or field in which the lens resides while the `local environment' describes the environment within $\\approx 100$ \\hinv kpc of the lensing galaxy. A $\\Lambda$CDM cosmology with $\\Omega_M = 0.27$ and $\\Omega_\\Lambda = 0.73$ is used to determine all physical distances, which are measured in \\hinv units. ",
        "conclusions": "We find that the global environments of the SLACS lenses are typical of other massive early-type galaxies found in the SDSS. Two of the steeper than isothermal systems lie in very over-dense regions but the remaining lenses all have richness values that fall near the peak of the richness distribution (Figure \\ref{figure_comp_rich}). Furthermore, the suggestion that the SLACS lenses are affected by LOS contamination does not seem to be merited by the data. While there are other galaxies along the lines of sight to the lens systems, the LOS densities do not significantly deviate from the densities along comparable lines of sight. We therefore expect that the parameter estimates should not be affected as proposed by \\citet{guimaraes}.  The SDSS photometric data indicate that the SLACS lenses are only slightly more likely to be associated with companion galaxies than comparable lenses selected from the SDSS, though the uncertainty of the photometric identifications makes the difference negligible. We therefore conclude that SLACS lenses do lie in typical environments both globally and locally. However, we also find that lens systems with steeper than isothermal density slopes are preferentially associated with companion galaxies compared to lenses with shallower density slopes. N-body simulations suggest that interactions with neighbouring galaxies can induce a steepening in the density slope \\citep{dobke} and there are other lens systems with companion galaxies that are found to be best modelled with steeper than isothermal profiles \\citep[e.g.,][]{rusin,augerb}. The interaction-induced steepening is a transient effect and the density profile of the galaxy will return to isothermal approximately 0.5-2 Gyr after the encounter with the neighbour \\citep{dobke}. This may account for the large range of $N_w$ for lenses with nearly isothermal profiles; isothermal lenses with a companion may be in the relaxed state before or after an encounter with the neighbour galaxy. This stripping mechanism may also account for local observations of dark matter deficient galaxies \\citep[e.g.,][]{romanowsky,proctor}.  We note that one steeper-profile system, SDSSJ1250+0523, is not photometrically associated with any neighbouring galaxies; this perhaps illustrates the limits of using photometry to find perturbing companions or demonstrates that other factors also influence the slope of the density profile. Additionally, we have not found an environmental bias to account for the shallower lenses. However, if a lens galaxy is embedded in a cluster, the joint profile of the cluster and galaxy would tend to be modelled with a shallower than isothermal profile if a single-component power law is used (ignoring interactions between the two halos). This effect is dependent on the location of the lens with respect to the centre of the cluster, which our photometric analysis is unable to address. More complete field spectroscopy would better characterize the global environments of these lens systems and allow correlations between the mass slopes and cluster centre offsets to be investigated. Furthermore, a spectroscopic investigation of the local environments of the complete sample of SLACS early-type lenses would confirm the correlation indicated by our photometric analysis and provide strong evidence for truncation caused by galaxy interactions."
    },
    "0710/0710.0711_arXiv.txt": {
        "abstract": "{ We present results of $H\\alpha$ imaging for 42 galaxies in the nearby low-density cloud Canes Venatici I populated mainly by late-type objects. Estimates of the $H\\alpha$ flux and integrated star formation rate ($SFR$) are now available for all 78 known members of this scattered system, spanning a large range in luminosity, surface brightness, $HI$ content and $SFR$. Distributions of the CVnI galaxies versus their $SFR$, blue absolute magnitude and total hydrogen mass are given in comparison with those for a population of the nearby virialized group around M81. We found no essential correlation between star formation activity in a galaxy and its density environment. A bulk of CVnI galaxies had enough time to generate their baryon mass with the observed $SFR$. Most of them possess also a supply of gas sufficient to maintain their observed $SFR$s during the next Hubble time.  } ",
        "introduction": "The distribution over the sky of 500 galaxies of the Local volume  with distances  within 10 Mpc shows considerable inhomogeneities due to the presence of groups and voids. Apart from several virialized groups, like the group around the galaxy M81, an amorphous association of nearby galaxies in the Canes Venatici constellation was noted by many authors (Karachentsev 1966, de Vaucouleurs 1975, Vennik 1984, Tully 1988). The boundaries of it are rather uncertain. Roughly, in a circle of radius $\\sim20\\degr$ around the galaxy NGC~4736 there are about 80 known galaxies with $D<10$ Mpc, which corresponds to a density contrast on the sky $\\Delta N/N\\sim4$. Inside of this complex the distribution of galaxies is also inhomogeneous, showing some clumps differing in their location on the sky and distances in depth. About 70\\% of the population of the cloud accounts for irregular dwarf galaxies, whose masses are obviously insufficient to keep such a system in the state of virial equilibrium. As data on distances of galaxies accumulated, a possibility appeared of studying the kinematics of the cloud in details. As was shown by Karachentsev et al. (2003), the CVnI cloud is in a state close to the free Hubble expansion, having a characteristic crossing time of about 15 Gyr.  Being a scattered system with rare interactions between galaxies, the nearby CVnI cloud is a unique laboratory for studying star formation processes in galaxies running independently, without a noticeable external influence. Kennicutt et al. (1989), Hoopes et al. (1999), van Zee (2000), Gil de Paz et al. (2003), James et al. (2004) and Hunter \\& Elmegreen (2004) conducted observations in the $H\\alpha$ line of three dozen galaxies of this complex, which made it possible to determine the star formation rate ($SFR$) in them. However, more than half of other members of the cloud proved to be out of vision of these authors. Our task consisted in the completion of  $H\\alpha$-survey of the population of the CVnI cloud. The results of our observations and their primary analysis are presented in this paper. ",
        "conclusions": "A systematic survey of $H{\\alpha}$-emission in the nearest scattered cloud CVnI shows that in most of its galaxies the process of active star formation is on despite the low density contrast of this cloud and rather rare interaction between its members. By making full use of our $H{\\alpha}$-survey, we can estimate the mean density of $SFR$ in the cloud. A conical volume in the distance range from 2 to 10 Mpc, resting in a sky region of $\\sim1500$ square degrees, makes 153 Mpc$^3$. In this volume we have a summary value $\\Sigma(SFR)=18.6M_{\\sun}$yr$^{-1}$, which yields an average density of star formation rate $\\dot{\\rho}_{SFR}=0.12M_{\\sun}$yr$^{-1}$Mpc$^{-3}$. The obtained value turns out to be a little less than $\\dot{\\rho}_{SFR}=0.165M_{\\sun}$yr$^{-1}$Mpc$^{-3}$ for a ``sell of homogeneity'' embracing the group M81 (Karachentsev \\& Kaisin, 2007). According to Nakamura et al. (2004), Martin et al. (2005) and Hanish et al. (2006), the average star formation rate per 1 Mpc$^3$ at the present epoch (z=0) is (0.02--0.03)$M_{\\sun}$yr$^{-1}$Mpc$^{-3}$. Consequently, the CVnI cloud has excess of $SFR$ density  4--6 times higher in comparison with the mean global quantity, which roughly corresponds to the cloud density contrast on the sky, $\\Delta N/N\\sim4$. Herefrom we conclude that the CVnI cloud is characterised by a usual norm of star formation in its galaxies."
    },
    "0710/0710.3120.txt": {
        "abstract": "We present a new measurement of the volumetric rate of Type Ia supernova  up to a redshift of 1.7, using the Hubble Space Telescope (HST) GOODS data combined with an additional HST dataset covering the North GOODS field collected in 2004. We employ a novel technique that does not require spectroscopic data for identifying Type Ia supernovae (although spectroscopic measurements of redshifts are used for over half the sample); instead we employ a Bayesian approach using only photometric data to calculate the probability that an object is a Type Ia supernova. This Bayesian technique can easily be modified to incorporate improved priors on supernova properties, and it is well-suited for future high-statistics supernovae searches in which spectroscopic follow up of all candidates will be impractical.  Here, the method is validated on both ground- and space-based supernova data having some spectroscopic follow up. We combine our volumetric rate measurements with low redshift supernova data, and fit to a number of possible models for the evolution of the Type Ia supernova rate as a function of redshift. The data do not distinguish between a flat rate at redshift $>$ 0.5 and a previously proposed model, in which the Type Ia rate peaks at redshift $\\sim$ 1 due to a significant delay from star-formation to the supernova explosion.  Except for the highest redshifts, where the signal to noise ratio is generally too low to apply this technique, this approach yields smaller or comparable uncertainties than previous work. ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%5 The empirical evidence for the existence of dark energy came from observations of Type Ia supernovae~\\citep{bib:riess, bib:P99, bib:P03}, which are believed to arise from the thermonuclear explosion of a progenitor white dwarf after it approaches the Chandrasekhar mass limit~\\citep{bib:chan}. However, the physics of Type Ia supernova production is not well understood. The two most plausible scenarios for the white dwarf to accrete the necessary mass are the single degenerate case, where the white dwarf is located in a binary system; and the double degenerate case, where two white dwarfs merge.  The Type Ia supernova rate is correlated with the star formation history (SFH), and thus a measurement of the rate as a function of redshift helps constrain the possible type Ia progenitor models.  In addition to its importance for understanding Type Ia supernovae as astronomical objects, a good grasp of the Type Ia supernova rate to high redshifts is important for the next generation of proposed space-based supernova cosmology experiments, such as SNAP~\\citep{bib:snap}. It is therefore of great practical interest to determine the rate of Type Ia supernovae at redshifts $>$ 1.  The subject of Type Ia supernova rates has been addressed by many authors in the past. Existing rate measurements have been mostly limited to redshift ranges $<$ 1: the results of~\\cite{bib:cappellaro},~\\cite{bib:hardin},~\\cite{bib:madgwick}, and~\\cite{bib:blanc} measure the rates at redshifts $\\leq$ $\\sim$0.1;~\\cite{bib:neill},~\\cite{bib:tonry}, and~\\cite{bib:pain}, at intermediate redshifts of 0.47, 0.50, and 0.55, respectively; and~\\cite{bib:barris}, up to a redshift of 0.75.  The only published measurement of the rates at redshifts $>$ 1 is that of~\\cite{bib:dahlen}, who analyzed the GOODS dataset.  There are several important differences that distinguish our work from that of~\\cite{bib:dahlen}. First, we augment the GOODS sample with the HST data collected during the Spring-Summer 2004 high redshift supernova searches. Second, our methods of calculating the control time (the time during which a supernova search is potentially capable of finding supernova candidates) and the efficiency to identify a supernova are based on a detailed Monte Carlo simulation technique using a library of supernova templates. Third, we adopt a novel approach to typing supernovae, using photometric data and a Bayesian probability method described in~\\cite{bib:ourpaper}. The Bayesian technique is able to perform classification using only photometric data, and therefore does not require spectroscopic follow up.  Optionally, photometric or spectroscopic redshifts can be used to improve the classification accuracy. Our initial requirements on potential supernova candidates are more stringent in terms of the number of points on the light curve and the signal to noise of those points than those of~\\cite{bib:dahlen}; thus some of the candidates they identified will fail our cuts. However, we are able to reliably separate Type Ia supernova from other supernovae types based on their Bayesian probability, with an efficiency that is readily quantifiable, thus allowing us to use larger data samples. Our approach therefore avoids the problems that arise in estimating the efficiency for the decision to schedule spectroscopic follow up based on a potentially low signal-to-noise initial detection.   The Bayesian classification technique uses photometric data, and does not require any spectroscopic followup.  This is an advantage for future large-area surveys  (such as the Dark Energy Survey, Pan-STARRS, and LSST) that will discover thousands of supernova candidates, but are unlikely  to be able to obtain spectroscopic data for all of them, to distinguish Type Ia supernovae from core collapse supernovae and other variable objects. The technique described here can be considered a prototype of the kind of analysis that could be performed on these future large data sets to identify Type Ia supernovae for cosmological studies. There is a clear trade-off involved in using photometric measurements alone: if the quality of the photometric data is poor, then the efficiency of this technique to identify Type Ia supernovae is reduced; on the other hand, this technique enables larger samples of Type Ia from imaging surveys to be identified for cosmological studies, without the need for time-consuming spectroscopic follow up.   Note that although the method is able to perform the supernova typing with photometric data alone (\\emph{i.e.}, it does not require spectroscopic data, either redshifts or types), it is certainly able to use the extra information that is available, and in fact 70\\% of the supernova candidates discussed in the present work have redshifts which were obtained spectroscopically. It is also worth noting that while in this paper we only analyze the Type Ia supernova rates, the  Bayesian classification technique can be used to classify other types as well, making it possible to measure the rates of non-Type Ia supernovae in a similar fashion. These analyses will be presented in future publications.    The paper is organized as follows.  In section~\\ref{sec:data} we describe the data samples used in the analysis.  In section~\\ref{sec:candselection} we describe the supernova candidate selection and typing process. In section 4 we calculate the control time, survey area, and search efficiency, and determine the volumetric Type Ia supernova rate from our data sample. A comparison of the rates with those reported in the literature is given in section~\\ref{sec:comp}, and fits of the rates to different models relating the Type Ia supernova rates to the SFH are given in section~\\ref{sec:sf}. A summary is given in section~\\ref{sec:concl}.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{sec:concl} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% We have analyzed the rates of Type Ia supernovae up to a redshift of 1.7 using two samples collected with the HST: the GOODS data, and the 2004 ACS sample collected in the Spring-Summer 2004 covering the GOODS North field.  Using only the data from two broadband filters, F775W and F850LP, we applied a novel technique for identifying Type Ia supernovae based on a Bayesian probability approach. This method allows us to automatically type supernova candidates in large samples, properly taking into account all known sources of systematic error. We also make use of the best currently available full spectral templates for five different supernova types for the candidate typing, as well as for calculating the efficiency of our supernova search, and the control time. These templates will undoubtedly be improved over the next several years as more supernova data becomes available.  Current and upcoming supernova surveys will not only provide a better understanding of individual supernova types, but may also uncover new types of supernovae, which can then be added to the Bayesian classification framework. Likewise, a better understanding of the many parameters that affect supernova observations will improve  the classification scheme, which will result in better constraints on the measured rates. The calculations of the supernova finding efficiency, the control time, and the survey area are all done taking into account the specific observing configurations pertinent for the surveys, such as exposure times, cadences, and the orientations of the GOODS tiles.   We carried out a comparison of the predicted and observed numbers of supernovae in redshift bins of $\\Delta \\bar{z}$ = 0.1, for two different models of the Type Ia supernova rates:  a redshift-independent rate and a power-law redshift-dependent rate. We find that the available data fit both models equally well.   For comparison with previous work, particularly that of~\\cite{bib:dahlen}, who also analyzed a large subset of the data used here, we calculated the volumetric Type Ia supernova rates in four redshift bins, 0.2 $\\leq$ $\\bar z$ $<$ 0.6, 0.6 $\\leq$ $\\bar z$ $<$ 1.0, 1.0 $\\leq$ $\\bar z$ $<$ 1.4, and 1.4 $\\leq$ $\\bar z$ $<$ 1.7. We find that our results are generally consistent with those of~\\cite{bib:dahlen}. Due to the larger of number supernova candidates which this Bayesian classification technique makes available, we obtain smaller or equal uncertainties in all the bins up to $z$ = 1.7. In the highest redshift bin we obtain a larger uncertainty because the signal to noise ratio is generally too low to apply this technique.    We fitted the resulting rates to two leading models used in recent literature: the two-component model and a Gaussian time delay model.  The former model implies an increase in the Type Ia supernova rates at highest redshifts; while the latter, a decrease. We find that the statistics of the present sample does not definitively discriminate between the two scenarios -- only one supernova in this work and two supernovae in~\\cite{bib:dahlen} contribute to the important highest-redshift bin. Significantly larger surveillance time would be required to arrive at a conclusive statement on the trends for the Type Ia rates at high redshifts.   In the future, several ambitious new surveys are planned that will collect photometric data for thousands of supernovae in order to improve the constraints on dark energy.  Individual spectroscopic follow up for every supernova candidate is likely to be impractical in these surveys. The Bayesian classification method described here has the ability to classify supernovae using photometric measurements alone, and is a promising technique for these future surveys.   %%%%%%%%"
    },
    "0710/0710.4129.txt": {
        "abstract": "{} {To compute the chemical evolution of spiral bulges hosting Seyfert nuclei, based on updated chemical and spectro-photometrical evolution models for the bulge of our Galaxy, to make predictions about other quantities measured in Seyferts, and to model the photometric features of local bulges. The chemical evolution model contains updated and detailed calculations of the Galactic potential and of the feedback from the central supermassive black hole, and the spectro-photometric model covers a wide range of stellar ages and metallicities.} {We computed the evolution of bulges in the mass range $2\\times 10^{9}-10^{11}M_{\\odot}$ by scaling the efficiency of star formation and the bulge scalelength as in the inverse-wind scenario for elliptical galaxies, and considering an Eddington limited accretion onto the central supermassive black hole.} {We successfully reproduced the observed relation between the mass of the black hole and that of the host bulge. The observed nuclear bolometric luminosity emitted by the supermassive black hole is reproduced only at high redshift or for the most massive bulges; in the other cases, at $z \\simeq 0$ a rejuvenation mechanism is necessary. The energy provided by the black hole is in most cases not significant in the triggering of the galactic wind. The observed high star formation rates and metal overabundances are easily achieved, as well as the constancy of chemical abundances with the redshift and the bulge present-day colours. Those results are not affected if we vary the index of the stellar IMF from $x=0.95$ to $x=1.35$; a steeper IMF is instead required in order to reproduce the colour-magnitude relation and the present $K$-band luminosity of the bulge.} {We show that the chemical evolution of the host bulge, with a short formation timescale of $\\sim 0.1$ Gyr, a rather high efficiency of star formation ranging from $11$ to $50$ Gyr$^{-1}$ according to the bulge mass and an IMF flatter with respect to the solar neighbourhood, combined with the accretion onto the black hole is sufficient to explain the main observed features of Seyfert galaxies.} ",
        "introduction": "The outstanding question of the co-evolution of Active Galactic Nuclei (AGNs) and their host galaxies has received considerable attention in the past decades, since various pieces of evidence pointed to a link between the formation of supermassive black holes (BHs) and the formation and evolution of their host spheroids: for example, the usual presence of massive dark objects at the centre of nearby spheroids (Ford et al. 1997; Ho 1999; Wandel 1999); the correlation between the BH mass and the stellar velocity dispersion of the host (for quiescent galaxies, Ferrarese \\& Merritt 2000; Gebhardt et al. 2000a; Tremaine et al. 2002; for active galaxies, Gebhardt et al. 2000b; Ferrarese et al. 2001; Shields et al., 2003; Onken et al. 2004; Nelson et al. 2004) or its mass (Kormendy \\& Richstone 1995; Magorrian et al. 1998; Marconi \\& Hunt 2003; Dunlop et al. 2003); the similarity between light evolution of quasar (QSO) population and the star formation history of galaxies (Cavaliere \\& Vittorini 1998; Haiman et al. 2004); the establishment of a good match among the optical QSO luminosity function, the luminosity function of star-forming galaxies and the mass function of dark matter halos (DMHs) at $z\\sim 3$ (Haenhelt et al., 1998).  The most widely accepted explanation for the luminosity emitted by an AGN, is radiatively efficient gas accretion onto a central supermassive BH. %Some theoretical model link quasar activity with major merger events %in a hierarchical galaxy formation process (Kauffmann \\& Haenhelt %2000; Kauffmann et al., 2003). The outflows from AGNs can profoundly affect the evolution of the host galaxy, e.g. by quenching or inducing the star formation (e.g., see Ciotti \\& Ostriker 2007, and references therein).  The mutual feedback between galaxies and QSOs was used as a key to solve the shortcomings of the semianalytic models in galaxy evolution, e.g. the failure to account for the surface density of high-redshift massive galaxies (Blain et al., 2002; Cimatti et al., 2002) and for the $\\alpha$-enhancement as a function of mass (Thomas et al., 2002), since it could provide a way to invert the hierarchical scenario for the assembly of galaxies and star formation (see e.g. Monaco et al., 2000; Granato et al., 2004; Scannapieco et al., 2005).  The study of the chemical abundances of the QSOs was first undertaken by Hamann \\& Ferland (1993), who combined chemical evolution and spectral synthesis models to interpret the N~{\\scshape v}/C~{\\scshape iv} and N~{\\scshape v}/He~{\\scshape ii} broad emission line ratios, and found out that the high metallicities and the abundance ratios of the broad-line region are consistent with the outcomes of the models for giant elliptical galaxies (Arimoto \\& Yoshii, 1987; Matteucci \\& Tornamb\\`e, 1987; Angeletti \\& Giannone, 1990), where the timescales of star formation and enrichment are very short and the initial mass function (IMF) is top-heavy.  In the same year, Padovani \\& Matteucci (1993) and Matteucci \\& Padovani (1993) employed the chemical evolution model of Matteucci (1992) to model the evolution of radio-loud QSOs, which are hosted by massive ellipticals, following in detail the evolution of several chemical species in the gas.  They supposed that the mass loss from dying stars after the galactic wind provides the fuel for the central BH and modeled the bolometric luminosity as $L_{bol}=\\eta\\dot{M}c^2$, with a typical value for the efficiency of $\\eta = 0.1$ and were successful in obtaining the estimated QSO luminosities and the observed ratio of AGN to host galaxy luminosity.  Then, they studied the evolution of the chemical composition of the gas lost by stars in elliptical galaxies and spiral bulges for various elements (C, N, O, Ne, Mg, Si and Fe), and found out that due to the high star-formation rate (SFR) of spheroids at early times the standard QSO emission lines were naturally explained. The relatively weak observed time dependence of the QSO abundances for $t \\gtrsim 1$ Gyr was also predicted.  The model of Matteucci \\& Padovani (1993) still followed the classic wind scenario, where the efficiency of star formation decreases with increasing galactic mass and which was found to be inconsistent with the correlation between spheroid mass and $\\alpha$-enhancement (Matteucci 1994).  Moreover, Padovani \\& Matteucci (1993) pointed out that if all mass lost by stars in the host galaxy after the wind were accreted by the central BH, the final BH mass would be up to two orders of magnitude larger than observed.  Other works (Fria\\c ca \\& Terlevich 1998; Romano et al. 2002; Granato et al., 2004), which had a more refined treatment of gas dynamics, limited their analysis of chemical abundances to the metallicity $Z$ and the [Mg/Fe] ratio and their correlation with the galactic mass.  All these studies were mainly devoted to studying the co-evolution of radio-loud QSOs and their host spheroids, which are elliptical galaxies.  Now we want to extend the approach of Padovani \\& Matteucci (1993) to AGNs hosted by spiral bulges, with a more recent chemical evolution model for the bulge with the introduction of the treatment of feedback from the central BH and a more sophisticated dealing of the accretion rate.  Since Seyfert nuclei are preferentially hosted by disk-dominated galaxies (Adams, 1977; Yee, 1983; MacKenty, 1990; Ho et al. 1997) our study can be applied to this class of objects.  The paper is organized as follows: in \\S 2 we illustrate the chemical and photometrical evolution model, in \\S 3 we show our calculations of the potential energy and of the feedback from supernovae (SNe) and from the AGN, in \\S 4 we discuss our results concerning the black hole masses and luminosities, the chemical abundances and the photometry, and in \\S 5 we draw some conclusions. ",
        "conclusions": ""
    },
    "0710/0710.5628_arXiv.txt": {
        "abstract": "Dust has been detected in the recurrent nova RS\\,Ophiuchi on several occassions. I model the historical mid-infrared photometry and a recent Spitzer Space Telescope spectrum taken only half a year after the 2006 eruption. The dust envelope is little affected by the eruptions. I show evidence that the eruptions and possibly the red giant wind of RS\\,Oph may sculpt the interstellar medium, and show similar evidence for the recurrent dwarf nova T\\,Pyxidis. ",
        "introduction": "\\subsection{Asymptotic Giant Branch stars and red supergiants}  Stars with an initial mass between $\\sim1$ and $\\sim40$ M$_\\odot$ become hydrogen/helium-shell-burning Asymptotic Giant Branch (AGB) stars or core-helium-burning red supergiants (RSG). These phases are characterised by cool, molecular atmospheres giving rise to M spectral types (S or C for some chemically peculiar AGB stars), strong radial pulsations of this atmosphere on timescales of a year or more, and a high luminosity ($L>2,000$ L$_\\odot$). The pulsation may lift the atmosphere high enough such that dust can condense \\citep{BowenWillson1991}. Radiation pressure on the grains then drives a dust wind which, via collisions with molecular hydrogen, drags the gas along with it. At a typical wind speed $v_\\infty\\sim5$ to 30 km s$^{-1}$ these stars lose mass at rates from $\\dot{M}\\sim10^{-7}$ to over $10^{-4}$ M$_\\odot$ yr$^{-1}$ \\citep{vanLoonEtal1999}. In AGB stars this leads to the removal of the mantle and the premature death of the star, leaving the truncated core behind as a cooling white dwarf. The more massive RSGs either explode or evolve back to higher surface temperatures, and it is not yet clear how much mass they will have shed during the preceding RSG stage.  Mass-loss rates of cool, luminous stars are most easily estimated from the reprocessed radiation emitted by the dust grains at infrared (IR) wavelengths, which is particularly conspicuous in the $\\lambda\\sim10$ to 40 $\\mu$m region. The main problem with this method is that the dust constitutes only a minor fraction of the total mass in the wind, typically $\\sim1$:$200$ \\citep{Knapp1985} but poorly known for all but the dustiest stars. Carbon monoxide has strong rotational transitions at mm wavelengths which can be detected in relatively nearby stars, but despite being the most abundant molecule after H$_2$ this too is a trace species of which the mass fraction is uncertain --- and modelling the CO line requires knowledge of the temperature profile throughout the envelope which depends, amongst other things, on the dust content.  \\subsection{First ascent Red Giant Branch stars}  Low-mass stars ($M_{\\rm initial}<2$ M$_\\odot$) evolve along a first ascent Red Giant Branch (RGB) as hydrogen-shell burning stars with an inert helium core. They reach a maximum luminosity of only $L_{\\rm tip}\\sim2,000$ L$_\\odot$, and it becomes problematic for them to drive a wind. The situation is worsened by the fact that many RGB stars do not pulsate strongly and slowly enough to sufficiently increase the scaleheight, and the dust condensation is therefore unlikely to reach completeness. With a low, uncertain dust:gas mass ratio and possibly higher dilution of the dust envelope as it kinematically decouples from the bulk gas, measured mass-loss rates will tend to be too low. The threshold below which this happens is not well known, partly because of uncertainties in the opacity of the grains as they nucleate and grow: \\citet{GailSedlmayr1987} estimate $\\dot{M}\\gg 10^{-6}$ M$_\\odot$ yr$^{-1}$ but \\citet{NetzerElitzur1993} place it at $\\dot{M}>10^{-7}$ M$_\\odot$ yr$^{-1}$. \\citet{JudgeStencel1991} show that RGB mass loss must be driven by another mechanism, but that this appears to be similarly efficient as a dust-driven wind. The warmer stars further down the RGB have more prominent chromospheres, which may be accompanied by the generation of Alfv\\'en or acoustic waves that provide a possible alternative mechanism for driving a wind.  Mass-loss rates estimated from the IR emission for the brightest, dustiest RGB stars are typically $\\dot{M}\\sim10^{-7}$ to $10^{-6}$ M$_\\odot$ \\citep{vanLoonEtal2006,OrigliaEtal2007}. Mass-loss rates from $\\dot{M}\\sim10^{-9}$ to $10^{-6}$ M$_\\odot$ yr$^{-1}$ have been determined from the blue-displaced cores of strong optical absorption lines and emission in some of these lines, but these are very uncertain too because of the sensitivity to the precise excitation and ionization conditions in the wind. For the most luminous RGB stars the IR and optical estimates tend to agree within an order of magnitude \\citep{McDonaldvanLoon2007}.  \\begin{figure}[!t] \\plotone{Jacco_vanLoon_fig1.eps} \\caption{RS\\,Oph compared to isochrones from \\citet{GirardiEtal2000}.} \\end{figure} ",
        "conclusions": "Dust in RS\\,Oph has been detected in between several eruptions, and at only half a year since the 2006 eruption. The silicate dust species and mass-loss rate of $\\dot{M}\\sim2$ to $3\\times10^{-8}$ M$_\\odot$ yr$^{-1}$ are not at odds with the expectations for well-developed dust-driven winds of single stars. But because RS\\,Oph has a relatively warm photosphere, weak pulsation and not very high luminosity, this might in fact suggest that the mass-loss rate {\\it is} enhanced by the effects of the companion white dwarf. The eruptions do {\\it not} seem to have a dramatic effect on the dust envelope. An accretion rate onto the white dwarf surface of $\\dot{M}>10^{-7}$ M$_\\odot$ yr$^{-1}$ would probably require an enhanced flow through the Lagrangian point L$_1$.  Considering the pattern in the time intervals between historic eruptions, including the 1907 eruption \\citep{Schaefer2004}, one might heuristically expect the next eruption to be due around 2015 (Fig.\\ 10).  \\begin{figure}[!b] \\plotone{Jacco_vanLoon_fig10.eps} \\caption{Recent outbursts of RS\\,Oph and a heuristic prediction for the next.} \\end{figure}"
    },
    "0710/0710.2839_arXiv.txt": {
        "abstract": "{% Neutron stars contain persistent, ordered magnetic fields that are the strongest known in the Universe. However, their magnetic fluxes are similar to those in magnetic A and B stars and white dwarfs, suggesting that flux conservation during gravitational collapse may play an important role in establishing the field, although it might also be modified substantially by early convection, differential rotation, and magnetic instabilities. The equilibrium field configuration, established within hours (at most) of the formation of the star, is likely to be roughly axisymmetric, involving both poloidal and toroidal components. The stable stratification of the neutron star matter (due to its radial composition gradient) probably plays a crucial role in holding this magnetic structure inside the star. The field can evolve on long time scales by processes that overcome the stable stratification, such as weak interactions changing the relative abundances and ambipolar diffusion of charged particles with respect to neutrons. These processes become more effective for stronger magnetic fields, thus naturally explaining the magnetic energy dissipation expected in magnetars, at the same time as the longer-lived, weaker fields in classical and millisecond pulsars.} ",
        "introduction": "The purpose of this talk is to present and discuss some of the physical processes that are likely to be relevant in determining the structure and evolution of magnetic fields in neutron stars, almost regardless of their (so far largely unknown) internal composition and state of matter. For this reason, I do not discuss the fascinating, exotic issues of quark matter, superfluidity, superconductivity, and the like, but emphasize the much more pedestrian concepts of stable stratification and non-ideal magnetohydrodynamics (MHD) processes such as ambipolar diffusion. A review with a very different focus has recently been given by Geppert (2006). ",
        "conclusions": "The extremely strong magnetic fields found in neutron stars are nevertheless weak enough to be balanced by small perturbations in the density, pressure, and chemical composition of the stably stratified, degenerate, multi-species fluid found in their interior. The field likely originates from the fairly strong magnetic flux inherited from the progenitor star, possibly modified by a combination of convection, differential rotation, and magnetic instabilities acting during the short, protoneutron star stage following collapse. The equilibrium field set up during this stage likely involves linked toroidal and poloidal field components. It can decay through dissipative processes such as weak interactions and ambipolar diffusion, which change the chemical composition of the matter, allowing it to move. These processes become particularly effective at high field strengths, possibly accounting for the energy release inside magnetars, which also keeps the stellar interior hot enough for the magnetic field to rearrange substantially. This forces changes as well in the crust of the star, which can be broken by strong fields, whereas weaker ones can evolve by a combination of Hall drift and resistive dissipation. The internal dissipation processes are ineffective for pulsar-strength fields, which can easily survive for a pulsar lifetime. None of these processes depends essentially on the exotic properties of a neutron star, such as Cooper pairing or quark matter, which can however modify the time scales for these processes to occur."
    },
    "0710/0710.5244_arXiv.txt": {
        "abstract": "In this paper we present a deep and homogeneous i-band selected multi-waveband catalogue in the COSMOS field covering an area of about $0.7\\sq\\degr$. Our catalogue with a formal 50\\% completeness limit for point sources of $i\\sim 26.7$ comprises about 290~000 galaxies with information in 8 passbands. We combine publicly available u, B, V, r, i, z, and K data with proprietary imaging in H band. We discuss in detail the observations, the data reduction, and the photometric properties of the H-band data. We estimate photometric redshifts for all the galaxies in the catalogue. A comparison with 162 spectroscopic redshifts in the redshift range $ 0 \\lsim z \\lsim 3$ shows that the achieved accuracy of the photometric redshifts is \\mbox{$\\Delta z / (z_{spec}+1) \\lsim 0.035$} with only $\\sim 2$\\% outliers. We derive absolute UV magnitudes and investigate the evolution of the luminosity function evaluated in the restframe UV (1500~\\AA).  { There is a good agreement between the LFs derived here and the LFs derived in the FORS Deep Field. We see a similar brightening of M$^\\ast$ and a decrease of $\\phi^\\ast$ with redshift.} The catalogue including the photometric redshift information is made publicly available. ",
        "introduction": "\\label{sec:intro}  In the last decade our knowledge about the evolution of global galaxy properties over a large redshift range has improved considerably.  The 2dF Galaxy Redshift Survey (2dFGRS; \\citealt{colless:1}), the Sloan Digital Sky Survey (SDSS; \\citealt{stoughton:1}), and the 2MASS survey \\citep{2MASS} have provided very large local galaxy samples with spectroscopic and/or photometric information in various passbands.  Thanks to these data sets we are now able to assess very accurate local ($z\\sim 0.1$) reference points for many galaxy evolution measurements like the luminosity function, the star formation activity, the spatial clustering of galaxies, the stellar population, the morphology, etc.  In the redshift range between $0.2 \\lsim z \\lsim 1 $ pioneering work has been done in the context of the Canada France Redshift Survey \\citep{lilly:3}, the Autofib survey \\citep{ellis:1} and in the Canadian Network for Observational Cosmology survey \\citep{yee:1}. They provide accurate distances and absolute luminosities by spectroscopic followup of optically selected galaxies, thus being able to probe basic properties of galaxy evolution. Moreover the K20-survey \\citep{cimatti:2} as well as the MUNICS survey \\citep{drory:2,feulner:1} extend the analysis into the near infrared regime (for $0.2 \\lsim z \\lsim 1.5 $).  An important step towards probing the galaxy properties also in the high redshift regime around $z \\sim 3$ and $z\\sim 4$ was the work of \\citet{steidel_lbg:1} and \\citet{steidel_lbg:2}. They used colour selection to discriminate between low and high redshift galaxies \\citep[see ][for a review]{giavalisco:3}. The so-called Lyman-break galaxies (LBGs, mainly starburst galaxies at high redshift) are selected by means of important features in the UV spectrum of star-forming galaxies.   The next milestones in pushing the limiting magnitude for detectable galaxies to fainter and fainter limits were the space based Hubble Deep Field North (HDFN; \\citealt{HDF96}) and Hubble Deep Field South \\citep[HDFS; ][]{HDFS00,HDFS00a} \\citep[see ][for a review]{ferguson:1}. Although of a limited field of view of about $5\\sq\\arcmin$ only, the depth of the HDFs allowed the detection of galaxies up to a redshift of 5 and even beyond.  In the past years the space based HDFs were supplemented by many more multi-band photometric surveys like the NTT SUSI deep Field (NDF; \\citealt{arnouts_ntt}), the Chandra Deep Field South (CDFS; \\citealt{arnouts_cdfs}), the William Herschel Deep Field (WHDF; \\citealt{mccracken:1,metcalf:1}), the Subaru Deep Field/Survey (SDF; \\citealt{maihara:1,ouchi:2}), the COMBO-17 survey \\citep{combo17:1}, FIRES \\citep{labbe:1}, the FORS Deep Field (FDF; \\citealt{fdf_data}), the Great Observatories Origins Deep Survey (GOODS; \\citealt{giavalisco:1}), the Ultra Deep Field (UDF and UDF-Parallel ACS fields; \\citealt{giavalisco:2,bunker:1, bouwens:2004}), the VIRMOS deep survey \\citep{lefevre:3}, GEMS \\citep{rix:1}, the Keck Deep Fields \\citep{sawicki:2}, and the Multiwavelength Survey by Yale-Chile (MUSYC; \\citealt{gawiser:1, quadri:1}).  With the advent of all these deep multi-band photometric surveys the photometric redshift technique (essentially a generalisation of the drop-out technique) can be used to identify high-redshift galaxies. Photometric redshifts are often determined by means of template matching algorithm that applies Bayesian statistics and uses semi-empirical template spectra matched to broad-band photometry (see also \\citealt{baum:1, koo:1, brunner:1, Soto:1, benitez:1, bender:1, borgne:1, firth:1}).  Redshifts of galaxies that are several magnitudes fainter than typical spectroscopic limits can be determined reliably with an accuracy of \\mbox{$\\Delta z / (z_{spec}+1)$} of 0.02 to 0.1.  In this context the COSMOS survey (\\citet{scoville:1}; see also http://www.astro.caltech.edu/cosmos/ for an overview) combines deep to very-deep multi-waveband information in order to extend the analysis of deep pencil beam surveys to a much bigger volume, thus being able to drastically increase the statistics and detect also very rare objects. For this, the survey covers an area of about $2\\sq\\degr$ with imaging by space-based telescopes (Hubble, Spitzer, GALEX, XMM, Chandra) as well as large ground based telescopes (Subaru, VLA, ESO-VLT, UKIRT, NOAO, CFHT, and others).  In this paper we combine publicly available u, B, V, r, i, z, and K COSMOS data with proprietary imaging in the H band to derive a homogeneous multi-waveband catalogue suitable for deriving accurate photometric redshifts. In Section~\\ref{sec:imaging_data} we give an overview of the near-infrared (NIR) data acquisition and we describe our 2-pass data reduction pipeline used to derive optimally (in terms of signal-to-noise for faint sources) stacked images in Section~\\ref{sec:imaging_reduction}. We also present NIR galaxy number counts and compare them with the literature.\\\\ In Section~\\ref{sec:isel_catalogue} we present the deep multi-waveband i-band selected catalogue and discuss its properties, whereas the data reduction of the spectroscopic redshifts is described in Section~\\ref{sec:spec}.  In Section~\\ref{sec:photoz} we present the photometric redshift catalogue, discuss the accuracy of the latter and show the redshift distribution of the galaxies.  In Section~\\ref{sec:uvlf} we derive the redshift evolution of the restframe UV luminosity function and luminosity density at 1500~\\AA\\ from our i-selected catalogue before we summarise our findings in Section~\\ref{sec:summary_conclusion}.  We use AB magnitudes and adopt a $\\Lambda$ cosmology throughout the paper with \\mbox{$\\Omega_M=0.3$}, \\mbox{$\\Omega_\\Lambda=0.7$}, and \\mbox{$H_0=70 \\, \\mathrm{km} \\, \\mathrm{s}^{-1} \\, \\mathrm{Mpc}^{-1}$}. ",
        "conclusions": "\\label{sec:summary_conclusion}  In this paper we present the data acquisition and reduction of NIR Js, H, and K' bands in the COSMOS field. We describe a 2-pass reduction pipeline to reduce NIR data. The 2-pass pipeline is optimised to avoid flat-field errors introduced if the latter are constructed from science exposures. Moreover we present and implement a method to stack images of different quality resulting in an optimal S/N ratio for faint sky dominated {point sources}. The Js and K' band cover an area of about \\mbox{$200\\sq\\arcmin$} (1 patch) whereas the H band covers about $0.85\\sq\\degr$ (15 patches) in total. The 50\\% completeness limits are 22.67, $\\sim 21.9$, and 21.76 in the Js, H, and K' band, respectively. The number counts of all NIR bands nicely agree with the number counts taken from literature.  Furthermore we present a deep and homogeneous i-band selected multi-waveband catalogue in the COSMOS field by combining publicly available u, B, V, r, i, z, and K bands with the H band. The clean catalogue with a formal 50\\% completeness limit for point sources of $i\\sim 26.7$ comprises about 290~000 galaxies with information in 8 passbands and covers an area of about $0.7\\sq\\degr$ (12 patches).  We exclude all objects with corrupted magnitudes in only one of the filters from the catalogue in order to have a catalogue as homogeneous as possible.  Photometric redshifts for all objects are derived and a comparison with 162 spectroscopic redshifts in the redshift range $ 0 \\lsim z \\lsim 3$ shows that the achieved accuracy of the photometric redshifts is \\mbox{$\\Delta z / (z_{spec}+1) \\lsim 0.035$} with only $\\sim 2$\\% outliers. Please note that in order to break the degeneracy between high redshift and low redshift solutions we included also the GALEX FUV and NUV filters in the photometric redshift estimation which considerably reduced the number of outliers.  The multi-waveband catalogue including the photometric redshift information is made publicly available. The data can be downloaded from {\\verb http://www.mpe.mpg.de/~gabasch/COSMOS/ }  We derive absolute UV magnitudes and a comparison in a magnitude-redshift diagram with the FDF shows good agreement. Moreover we investigate the evolution of the luminosity function evaluated in the restframe UV (1500~\\AA). We find a substantial brightening of M$^\\ast$ and a decrease of $\\phi^\\ast$ with redshift: from \\mbox{$\\langle z \\rangle\\sim 0.5$} to \\mbox{$\\langle z \\rangle\\sim 4.5$} the characteristic magnitude increases by about 3 magnitudes, whereas the characteristic density decreases by about 80 -- 90\\%.  We compare the redshift evolution of the UV luminosity density in the COSMOS field and the FDF up to a redshift of $z\\sim 5$. Below a redshift of $z\\sim 2.5$ the mean UV luminosity density in COSMOS is systematically higher by about $1\\sigma$ if compared to the FDF. At $2.5 \\lsim z \\lsim 5$ both UV luminosity densities agree very well.  It is worth noting the remarkably good agreement between the UV LF as well as the UV LD despite the fact, that the FDF is about 60 times smaller than the COSMOS field analysed here."
    },
    "0710/0710.2464_arXiv.txt": {
        "abstract": "{} {We studied the temporal and spectral evolution of the synchrotron emission from the high energy peaked BL~Lac object \\1e.} {Two recent observations have been performed by the \\xmm\\,\\,and \\sw\\,satellites; we carried out X-ray spectral analysis for both of them, and photometry in optical-ultraviolet filters for the \\sw\\,\\,one. Combining the results thus obtained with archival data we built the long-term X-ray light curve, spanning a time interval of 26~years, and the Spectral Energy Distribution (SED) of this source.} {The light curve shows a large flux increasing, about a factor of six, in a time interval of a few years. After reaching its maximum in coincidence with the \\xmm\\,\\,pointing in December~2000 the flux decreased in later years, as revealed by \\sw\\,. The very good statistics available in the 0.5-10 keV \\xmm\\,\\,X-ray spectrum points out a highly significant deviation from a single power law. A log-parabolic model with a best fit curvature parameter of 0.25 and a peak energy at $\\sim$~1~keV describes well the spectral shape of the synchrotron emission. The simultaneous fit of \\sw\\,\\,UVOT and XRT data provides a milder curvature ($b\\sim0.1$) and a peak at higher energies ($\\sim15$~keV), suggesting a different state of source activity. In both cases UVOT data support the scenario of a single synchrotron emission component extending from the optical/UV to the X-ray band.} {New X-ray observations are important to monitor the temporal and spectral evolution of the source; new generation $\\gamma$-ray telescopes like AGILE and GLAST could for the first time detect its inverse Compton emission.} ",
        "introduction": "BL~Lac objects are thought to be radio-loud Active Galactic Nuclei (AGNs) observed in a direction very close to the axis of a relativistic jet outflowing from the inner nuclear region (\\cite{Urry95}). This interpretation could explain most of the characteristics of these sources like compact and flat-spectrum radio emission, superluminal motion revealed by VLBI imaging, high and variable radio and optical polarization, non-thermal continuum emission extending from radio to $\\gamma$-ray frequencies, an almost featureless optical spectrum and the fast variability at all frequencies.  BL~Lac objects are generally characterized by a double bump structure in the broad band Spectral Energy Distribution (SED). The low frequency bump is attributed to synchrotron radiation emitted by relativistic electrons in the jet; inverse Compton scattering by the same electron population on the synchrotron radiation is thought to be at the origin of the high frequency bump. The peak of the first bump may vary in a rather wide range of frequencies: from the IR/optical band for the low energy peaked BL~Lacs (LBLs) to the UV/X-ray band for the high energy peaked BL~Lacs (HBLs) (\\cite{Giommi94}, \\cite{Padovani}).  \\1e, also named BZB J1210+3929 in the recent Multifrequency Catalogue of Blazars (\\cite{BZcat}) is one of the X-ray selected BL~Lacertae objects of the \\ein\\,\\,Medium Sensitivity Survey (\\cite{EMSS}). It was discovered as a serendipitous source located about five arc-minutes north of one of the most intensively studied AGNs, the bright Seyfert galaxy NGC~4151. For this reason \\1e has been observed on many occasions by all the imaging X-ray instruments that have operated since the \\ein\\,\\,observatory. Despite its relatively high redshift (z=0.615) HST was able to detect the bright (M$_{\\rm R}$\\,=\\,$-$24.4) host galaxy which is of elliptical type (\\cite{Scarpa00}).  In this paper we report the long-term X-ray light curve which spans over 26 years. We also report and compare the spectral analysis of the most recent observations carried out by the \\xmm\\,\\,(\\cite{jansen}) and \\sw\\,\\,(\\cite{gehrels}) satellites; we discuss the possibility to model the synchrotron emission of this source with a single log-parabolic model. The peak of the synchrotron component lies in the X-ray band, and for this reason the source can safely be classified as an HBL. We finally report the Spectral Energy Distribution (SED) of \\1e compiled from non-simultaneous multi-frequency archival data. ",
        "conclusions": "\\label{Discussion}  Adding to multi-frequency archival data the results of our analysis we built the radio to X-ray SED of \\1e (Fig.\\ref{SED}) which follows the characteristic trend of extreme HBL sources (\\cite{Padovani}) with the synchrotron emission covering the entire frequency interval and peaking well into the X-ray band. The inverse Compton component is not detected and therefore is expected at higher energies.  The \\sw\\,\\,satellite is sensitive both in the X-ray and in the optical-ultraviolet band. This provides the opportunity to test if the Spectral Energy Distribution of a source like \\1e can be explained by a single synchrotron component, or if multiple emission components are present. For this purpose we first fitted a log-parabolic model to simultaneous \\sw\\,\\,XRT and UVOT fluxes (solid line in Fig.~\\ref{combo}) and obtained a rather low value for the curvature parameter ($b\\sim0.1$) in contrast with the one obtained by fitting XRT data only, as reported in Table~\\ref{tab1}; the energy peak lies at about 15~keV or more. Comparing these results with those from \\xmm, \\sw\\,\\,evidently caught \\1e in a different state of activity, characterized by a milder curvature and a peak of the synchrotron component shifted to higher energies. Anyway firm conclusions cannot be drawn due to a relatively poor statistics and to the fact that the found $E_p$ value likely lies at energies higher than 10~keV, outside the range of \\sw\\,\\,XRT.  Aware of the limits of using non simultaneous observations for a variable source, yet we veri\\-fied if \\sw\\,\\,UVOT and \\xmm\\,\\,MOS1 data are compatible with a log-parabolic model. We took points representing XSPEC model of observation~I and fit them on a wider frequency interval (dot-dashed line in Fig.\\,\\ref{combo}): UVOT points are sistematically above the extrapolation of the log-parabolic model. We took care to compare this result with the one obtained by fitting directly X-ray rebinned data and obtained an almost coincident result. Then we fitted both data sets with a log-parabola (dotted line in Fig.\\,\\ref{combo}) and estimated the curvature parameter: we obtained $b=0.18\\pm0.01$, a value at about 2$\\sigma$ of the one derived fitting only \\xmm\\,\\,data (see Table \\ref{tab1}). At this point we tested the log-parabolic model with $b$ fixed at 0.18 and found $\\chi^2_r/d.o.f. = 1.02/376$, a value practically coincident with the one obtained leaving $b$ free to vary, with only a small difference in the $a$ parameter. This result encouraged us in concluding that a single synchrotron component is in a good agreement with the emission observed.  \\1e is certainly worth monitoring in the next years. New observations in the X-ray band would add essential information to the light curve behaviour: if any kind of regularity in flux increasing and fading should emerge, it would be possible to put physical constraints to establish the nature of the mechanism responsible for these variations. Moreover, observations with instruments which extend to higher energies than XRT, like those on board \\textit{Suzaku}, would extend the known Spectral Energy Distribution and would allow us to obtain a better parametrization of the spectral curvature. Finally, \\1e may be an interesting target for AGILE and the forthcoming GLAST $\\gamma$-ray mission which could detect for the first time the inverse Compton component of this HBL source."
    },
    "0710/0710.0182_arXiv.txt": {
        "abstract": "Using high-resolution N-body simulations, we examine whether a major dry merger mitigates the difference in the radial density distributions between red and blue globular clusters (GCs). To this end, we study the relation between the density slope of the GCs in merger progenitors and that in a merger remnant, when the density distribution is described by $n_{\\rm GC}\\propto r^{-\\alpha}$. We also study how our results depend on the merger orbit and the size of the core radius of the initial GC density distribution. We find that a major dry merger makes the GC profile flatter, and the steeper initial GC profile leads to more significant flattening, especially if the initial slope is steeper than $\\alpha\\sim3.5$. Our result suggests that if there is a major dry merger of elliptical galaxies whose red GCs have a steeper radial profile than the blue GCs, as currently observed, and their slopes are steeper than $\\alpha\\sim3.5$, the difference in the slopes between two populations becomes smaller after dry mergers. Therefore, the observed slopes of red and blue GCs can be a diagnostic of the importance of dry merger. The current observational data show that the red and blue GCs have more comparable and shallower slopes in some luminous galaxies, which may indicate that they have experienced dry mergers. ",
        "introduction": "Globular clusters (GCs) in elliptical galaxies have been intensively studied in consideration of explaining the formation of both their host elliptical galaxies and GCs themselves \\citep[see][for a review]{araa06}. GCs are attractive as tracers of the star formation history of their host galaxies \\citep[e.g.][]{yi04,strader06}, because some properties of GC systems are correlated with the properties of their host galaxies \\citep[e.g.,][]{brodie91,djorgovski92}. It is thought that the formation of GCs is triggered by starburst accompanying gas-rich galaxy merging \\citep[e.g.][]{schweizer87,ashman92} or starburst that might happen with multiple dissipational collapses \\citep{forbes97}. Forming young star clusters are found and they are expected to become star clusters like current old GCs in local galaxy mergers \\citep{schweizer06}.  An important aspect of GC systems in elliptical galaxies is a color bimodality \\citep[e.g.,][]{zepf93,geisler96,gebhardt99,larsen01,peng06}. It is also found that red GCs are more centrally concentrated than blue GCs \\citep[e.g.][]{forte05,bassino06b, tamura06}, which could put additional constraints on their formation scenario \\citep{bekki02}. The radial profile of each GC subpopulation is well described by a power-law distribution, especially at the outer radii, and the red GCs have a steeper slope than the blue GCs. Although the origin of this color bimodality is still uncertain, it is probably closely related to the formation history of their host elliptical galaxies \\citep[e.g.,][]{yoon06,strader07,kundu07}. Classically, three scenarios have been proposed to explain the color bimodality; major gas-rich mergers, in situ formation of multiple dissipational collapses, and dissipationless accretion. An explanation suggested by \\citet{ashman92} is that red GCs are metal-rich clusters which might be formed by gas-rich disk-disk mergers. Therefore, the red GCs might be younger than blue GCs. Meanwhile, \\citet{forbes97} explain that blue GCs might be formed in the first stage of dissipational collapse and red GCs might be formed after the truncation of the blue GC formation. Another explanation given by \\citet{cote98} includes accretion of blue GCs from small galaxies into the already formed red GCs. More recently, \\citet{beasley02} demonstrate that the bimodality can be explained in elliptical galaxy formation based on a hierarchical clustering scenario.  On the other hand, recent research suggests that in the late stage of evolution, early-type galaxies might have experienced dry merging where merger progenitors do not have much gas, nor accompany star formation. The number density evolution of red galaxies has been discussed in observations of COMBO-17 \\citep{bell04} and DEEP2 surveys \\citep{faber05}, and such studies suggest that the density change can be understood by the dominance of dry merging after z $\\simlt$ 1 \\citep[but see also][]{yamada05,cimatti06,bundy07,scarlata07}. The evolution of galaxy clustering also implies late effects on the evolution of massive red galaxies from dry merging \\citep{white07}. Moreover, the observations show that dry merging does occur \\citep{vandokkum05,tran05,rines07}, while the observed features of galaxies are well explained in cosmological simulations \\citep[e.g.][]{kawata06}. Recent theoretical studies of dry merger simulations of ellipticals show that merger remnants maintain their properties on the fundamental plane and other scaling relations \\citep[e.g.][]{nipoti03,boylan05,robertson06,ciotti07}. The dry merging of binary ellipticals also can explain the formation of boxy-type ellipticals \\citep{naab06}.  \\citet{bekki06a} demonstrate that the observed correlation between a spatial distribution of GCs and the total luminosity of ellipticals can be explained by sequential dissipationless major mergers, because the radial density profile of GCs progressively flatten after each major dry merger. This study raises the important question of whether or not the slopes of the density profiles of red and blue GCs persist after major dry merging. For example, we now consider that the density profiles of red and blue GCs in progenitor elliptical galaxies are described by $n_{\\rm GC, red}\\propto r^{-\\alpha_{\\rm p,red}}$ and $n_{\\rm GC, blue}\\propto r^{-\\alpha_{\\rm p,blue}}$, and these profiles in a major dry merger remnant become $n_{\\rm GC, red}\\propto r^{-\\alpha_{\\rm r,red}}$ and $n_{\\rm GC, blue}\\propto r^{-\\alpha_{\\rm r,blue}}$. The current observations suggest that $\\alpha_{\\rm r,red}>\\alpha_{\\rm r,blue}$. However, if a dry merger flattens the red GC density profile more than the blue GC density profile, $\\alpha_{\\rm p,red}-\\alpha_{\\rm r,red}>\\alpha_{\\rm p,blue}-\\alpha_{\\rm r,blue}$, in the remnant galaxy the difference between $\\alpha_{\\rm r,red}$ and $\\alpha_{\\rm r, blue}$ becomes smaller. In this case, the observed difference in the slopes of the density profiles of red and blue GCs in nearby ellipticals can be a valuable diagnostic for the importance of the dry merger in the evolution of ellipticals. To clarify this issue, the question becomes how the flattening of GC profiles during dry merging depends on the initial distributions of the GCs.   We use numerical simulations of major dry mergers to study the dependence of GC distributions in merger remnants on the initial distributions in merger progenitors. Then, we can compare $\\alpha_{\\rm p,red}-\\alpha_{\\rm r,red}$ and $\\alpha_{\\rm p,blue}-\\alpha_{\\rm r,blue}$, for different sets of $\\alpha_{\\rm p,red}$ and $\\alpha_{\\rm p,blue}$. Since major dry mergers between two equal-mass merger progenitors must leave the most significant effects on merger remnants, compared with minor mergers, we study only equal-mass mergers in this paper. In \\S2, we explain details of our dry merger simulations and initial GC distributions. The changes in GC spatial distributions due to major mergers are shown in \\S3. We discuss the implication of our results in \\S4 that is followed by conclusion. ",
        "conclusions": "Our main result is that steeper initial GC profiles experience stronger flattening as shown in Figure \\ref{fig:alpha}. $\\alpha_{\\rm p} \\approx 3.5$ is in boundary between strong and weak flattening. The results imply that if the initial slopes of both red and blue GCs are steeper than $\\alpha_{\\rm p} \\approx 3.5$, the difference in the slopes between two populations of GCs will become much smaller, independent of merger orbits. In particular, even only one dry merger can make the slope flatter dramatically. Moreover, it also makes both slopes be around $\\alpha_{\\rm r} \\approx 3.5 - 4.5$. For example, if $\\alpha_{\\rm p,red}=5$ and $\\alpha_{\\rm p, blue}=4$, Figure \\ref{fig:alpha} suggests that $\\alpha_{\\rm r,red} \\approx 4.3$ and $\\alpha_{\\rm p, blue} \\approx 3.6$. The difference in the slope between red and blue GCs becomes significantly small. Therefore, the difference in the slopes of red and blue GCs can be a constraint of the number of major dry mergers.  On the other hand, if the initial GC distribution has a slope shallower than $\\alpha_{\\rm p} \\approx 3.5$, the slope changes very little, and almost no change in the case of $\\alpha_{\\rm p} \\approx 2$. Therefore, if the initial distributions of red and blue GCs in merger progenitors are steeper than $\\alpha_{\\rm p} = 2$, and the galaxies experienced a number of major dry mergers, it leads the distributions of both red and blue GCs to become close to $\\alpha_{\\rm r} \\sim 2$ for both red and blue GCs, and the difference in the slope of spatial distributions becomes difficult to be measured.  It is also worth noting that if the initial core size is larger for blue GCs than for red GCs, but their initial slopes are the same, a major dry merger can make the slope for the blue GCs shallower than for the red GCs, as shown in Figure \\ref{fig:r_c_dependence1}. Therefore, the final slope is not a simple function of the initial slope.  Our results suggest that dry mergers make the slopes of the density profiles for red and blue GCs shallower and similar. Several studies \\citep[e.g.][]{kissler97,vandenbergh98,forbes05,lauer07, emsellem07} claim that the various observed properties for ellipticals show a transition around ${\\rm M_{V} \\sim -21}$, i.e. ${\\rm \\sim 10^{11}\\ M_{\\sun}}$. Some of these bimodalities have been interpreted in the picture of dry merger hypothesis for the growth of massive ellipticals \\citep[e.g.][]{capetti06}. For example, \\citet{emsellem07} show that slowly rotating ellipticals might be mainly affected by dry mergers. The slowly rotating ellipticals are more luminous than ${\\rm M_{B} \\approx -20.5}$, while less luminous ellipticals are fast rotators. It may indicate that less luminous galaxies have not experienced any major dry  mergers. Therefore, it is interesting to examine the observed distributions of red and blue GCs as a function of the stellar mass of ellipticals. If more luminous galaxies have experienced dry mergers, our results predict that the slopes for red and blue GCs in bright galaxies are shallower and similar. Unfortunately, the current observational samples are not enough to test it statistically. Below we provide some discussion based on the current limited measurements.  NGC 1399 is one of ellipticals whose spatial distributions of red and blue GCs are well-observed \\citep{dirsch03,bassino06b}. When assuming ${\\rm M/L_{V} = 5}$ and $B-V=0.9$ for an elliptical galaxy, as used in \\citet{forbes05}, the stellar mass of NGC 1399 is about ${\\rm 5 \\times\\ 10^{11}\\ M_{\\odot}}$ (${\\rm M_{B} = -21.8}$) \\citep{araa06}. The derived power-law indices of projected radial density profiles are $1.9\\pm0.06$ and $1.6\\pm0.10$ for red and blue GCs, respectively \\citep{bassino06b}. If the distributions are spherically symmetric, the three-dimensional radial distributions have power-law slopes of $\\alpha \\approx 2.9$ and $\\approx 2.6$ for red and blue GCs, respectively. Note that the real slope might be slightly steeper, because the distribution of GCs is not infinite. The slope difference between red and blue GC distributions is $\\Delta \\alpha \\approx 0.3$ while both GC populations show $\\alpha \\approx 3$. In addition to NGC 1399, NGC 1407 is a brightest group galaxy with ${\\rm M_V=-21.86}$, and has a bimodal color distribution of GCs. \\citet{forbes06} report that the projected slopes of the red and blue GC density profiles are respectively $1.50\\pm0.05$ and $1.65\\pm0.29$, i.e., expected three dimensional slopes of $\\alpha_{\\rm red}=2.50$ $\\alpha_{\\rm blue}=2.65$, which are statistically the same as each other.  On the other hand, NGC 1374 and NGC 1379 are less luminous than NGC 1399, having ${\\rm M_{V} = -20.4}$ and $-20.6$, respectively. The red and blue GCs of the two galaxies are studied in \\citet{bassino06a}. The derived projected density distributions of red and blue GCs in NGC 1374 have slopes of 3.2 and 2.3. Hence, the expected slopes in their three-dimensional density profiles are $\\alpha_{\\rm red}\\sim 4.2$ and $\\alpha_{\\rm blue}\\sim3.3$. The $\\Delta \\alpha$ in NGC 1379 is $\\sim 0.6$, and the red and blue GCs have projected power-law slopes of $\\sim 2.9$ and $\\sim 2.3$, respectively. The low-luminosity galaxies have a larger difference in the density slopes of the red and blue GCs, and their slopes are steeper than the more luminous galaxies discussed above. Combining with our results, these four sample suggests that luminous galaxies, like NGC 1399 and NGC 1407, may have experienced some dry mergers, while less luminous galaxies, such as NGC 1374 and NGC 1379, may have not been formed through dry mergers, which is consistent with a scenario that the significance of dry mergers depends on the stellar mass.  However, we also find some giant galaxy that does not follow the above trend. Recently, \\citet{tamura06} measured the density profiles of red and blue GCs for  giant ellipticals: M87 and NGC 4552. We have fitted the projected density profile in Figure 5 of their paper by a power-law profile, and the derived power-law index is 2.4 and 1.4 for red and blue GCs in M87 and 1.8 and 1.2 for red and blue GCs in NGC 4552, respectively. Although these giant ellipticals have relatively shallower slopes, the difference in slopes of red and blue GCs is significant. According to our result, these giant ellipticals are unlikely to have experienced significant dry mergers. Some giants might form without dry mergers, or some other factor, such as a number of minor mergers, might affect the slopes of GCs.  Again, so far there are not many studies which derive the density distributions of both red and blue GCs. This problem limits the application of our results to the current observational data. However, we expect that future observations will improve the statistics for nearby ellipticals, and our results will add an important framework to interpret the galaxies' evolutionary history. It will be also valuable to compare spatial distributions of blue and red GCs with other expected properties from dry mergers such as surface brightness profile."
    },
    "0710/0710.5857_arXiv.txt": {
        "abstract": "The classical picture of GUT baryogenesis has been strongly modified by theoretical progress concerning two nonperturbative features of the standard model: the phase diagram of the electroweak theory, and baryon and lepton number changing sphaleron processes in the high-temperature symmetric phase of the standard model. We briefly review three viable models, electroweak baryogenesis, the Affleck-Dine mechanism and leptogenesis and discuss the prospects to falsify them. All models are closely tied to the nature of dark matter, especially in supersymmetric theories. In the near future results from LHC and gamma-ray astronomy will shed new light on the origin of the matter-antimatter asymmetry of the universe. ",
        "introduction": "The cosmological matter-antimatter asymmetry can be dynamically generated if the particle interactions and the cosmological evolution satisfy Sakharov's conditions \\cite{sak67}, \\begin{itemize} \\item baryon number violation, \\item $C$ and $C\\!P$ violation, \\item deviation from thermal equilibrium. \\end{itemize} Although the baryon asymmetry is just a single number, it provides an important connection between particle physics and cosmology. In his seminal paper, 40 years ago, Sakharov not only stated the necessary conditions for baryogenesis, he also proposed a specific model. The origin of the baryon asymmetry were $C\\!P$ violating decays of superheavy `maximons' with mass $\\mathcal{O}(M_\\mathrm{P})$ at an initial temperature $T_i \\sim M_\\mathrm{P}$. The $C\\!P$ violation in maximon decays was related to the $C\\!P$ violation observed in $K^0$-decays, and the violation of baryon number led to a proton lifetime $\\tau_p > 10^{50}\\ \\mathrm{years}$, much larger than current estimates in grand unified theories.  \\begin{figure}[t] \\includegraphics[height=6cm]{fig3v}\\hspace{1cm} \\includegraphics[height=6cm,width=8cm]{pot} \\caption{{\\it Left:} Critical temperature $T_c$ of the electroweak transition as function of $R_{HW}=m_H/m_W$; from \\cite{cfh98}. {\\it Right:} Effective potential of the Higgs field $\\varphi$ at temperature $T>T_c$. \\label{fig:ew}} \\end{figure}  At present there exist a number of viable scenarios for baryogenesis. They can be classified according to the different ways in which Sakharov's conditions are realized. In grand unified theories baryon number ($B$) and lepton number ($L$) are broken by the interactions of gauge bosons and leptoquarks. This is the basis of classical GUT baryogenesis (cf.~\\cite{kt90}). In a similar way, lepton number violating decays of heavy Majorana neutrinos lead to leptogenesis \\cite{fy86}. In the simplest version of leptogenesis the initial abundance of the heavy neutrinos is generated by thermal processes. Alternatively, heavy neutrinos may be produced in inflaton decays or in the reheating process after inflation. Because in the standard model baryon number, $C$ and $C\\!P$ are not conserved, in principle the cosmological baryon asymmetry can also be generated at the electroweak phase transition \\cite{krs85}. A further mechanism of baryogenesis can work in supersymmetric theories where the scalar potential has approximately flat directions. Coherent oscillations of scalar fields can then generate large asymmetries \\cite{ad85}.  The theory of baryogenesis crucially depends on nonperturbative properties of the standard model, first of all the nature of the electroweak transition. A first-order phase transition yields a departure from thermal equilibrium. Fig.~1 shows the phase diagram of the electroweak theory, i.e. the critical temperature in units of the Higgs mass, $T_c/m_H$, as function of the Higgs mass in units of the W-boson mass, $R_{HW}=m_H/m_W$ \\cite{cfh98,lr98}. For small Higgs masses the phase transition is first-order; above a critical Higgs mass, $m_H > m_H^c \\simeq 72$~GeV, it turns into a smooth crossover \\cite{bp94,klx96}. This upper bound for a first-order transition has to be compared with the lower bound from LEP, $m_H > 114$~GeV. Hence, there is no departure from thermal equilibrium at the electroweak transition in the standard model.  The second crucial nonperturbative aspect of baryogenesis is the connection between baryon number and lepton number in the high-temperature, symmetric phase of the standard model. Due to the chiral nature of the weak interactions $B$ and $L$ are not conserved \\cite{tho76}. At zero temperature this has no observable effect due to the smallness of the weak coupling. However, as the temperature reaches the critical temperature $T_c$ of the electroweak phase transition, $B$ and $L$ violating processes come into thermal equilibrium \\cite{krs85}. The rate of these processes is related to the free energy of sphaleron-type field configurations which carry topological charge. In the standard model they lead to an effective interaction of all left-handed fermions \\cite{tho76} (cf.~Fig.~2), \\begin{equation} O_{B+L} = \\prod_i \\left(q_{Li} q_{Li} q_{Li} l_{Li}\\right)\\; , \\end{equation} which violates baryon and lepton number by three units, \\begin{equation} \\Delta B = \\Delta L = 3\\;. \\label{sphal1} \\end{equation} The sphaleron transition rate in the symmetric high-temperature phase has been evaluated by combining an analytical resummation with numerical lattice techniques \\cite{bmr00}. The result is, in accord with previous estimates, that $B$ and $L$ violating processes are in thermal equilibrium for temperatures in the range \\begin{equation} T_{EW} \\sim 100\\ \\mbox{GeV} < T < T_{SPH} \\sim 10^{12}\\ \\mbox{GeV}\\;. \\end{equation}  \\begin{figure}[t] \\includegraphics[height=6cm]{lg1} \\caption{One of the 12-fermion processes which are in thermal equilibrium in the high-temperature phase of the standard model. \\label{fig:sphal} } \\end{figure}  Sphaleron processes have a profound effect on the generation of the cosmological baryon asymmetry. An analysis of the chemical potentials of all particle species in the high-temperature phase yields the following relation between the baryon asymmetry and the corresponding $L$ and $B-L$ asymmetries, \\begin{equation}\\label{basic} \\langle B\\rangle_T = c_S \\langle B-L\\rangle_T = {c_S\\over c_S-1} \\langle L\\rangle_T\\;. \\end{equation} Here $c_S$ is a number ${\\cal O}(1)$. In the standard model with three generations and one Higgs doublet one has $c_s= 28/79$.  We conclude that lepton number violation is necessary in order to generate a cosmological baryon asymmetry\\footnote{In the case of Dirac neutrinos, which have extremely small Yukawa couplings, one can construct leptogenesis models where an asymmetry of lepton doublets is accompanied by an asymmetry of right-handed neutrinos such that the total lepton number is conserved and $\\langle B-L \\rangle_T = 0$ \\cite{dlx00}.}. However, it can only be weak, because otherwise any baryon asymmetry would be washed out. The interplay of these conflicting conditions leads to important contraints on neutrino properties and on possible extensions of the standard model in general.  \\begin{figure}[t] \\includegraphics[height=6cm]{ewbg} \\caption{Sketch of nonlocal electroweak baryogenesis. From \\cite{ber02}. } \\end{figure} ",
        "conclusions": "40 years after Sakharov's work on the cosmological matter-antimatter asymmetry we have several viable models of baryogenesis, the most predictive ones being electroweak baryogenesis and leptogenesis. In fact, based on our theoretical understanding of the electroweak phase diagram, electroweak baryogenesis in the standard model has already been excluded by the LEP bound on the Higgs mass. Supersymmetric electroweak baryogenesis will soon be tested at the LHC.  Detailed studies of the nonequilibrium leptogenesis process have led to the preferred neutrino mass window $10^{-3}\\ {\\rm eV} < m_i < 0.1\\ {\\rm eV}$ in the simplest scenario with hierarchical heavy neutrinos. The consistency with the experimental evidence for neutrino masses has dramatically increased the popularity of the leptogenesis mechanism. It is exciting that new experiments and cosmological observations will probe the absolute neutrino mass scale in the coming years. However, more work is needed on the full quantum mechanical treatment of leptogenesis, in particular the flavour dependence.  All baryogenesis mechanisms are closely related to the nature of dark matter. A discovery of the standard supergravity scenario at LHC could be consistent with electroweak baryogenesis but would rule out the simplest version of thermal leptogenesis. On the other hand, evidence for gravitino dark matter can be consistent with leptogenesis. Finally, the discovery of macroscopic dark matter like Q-balls would point towards nonperturbative dynamics of scalar fields in the early universe and therefore favour Affleck-Dine baryogenesis."
    },
    "0710/0710.2193_arXiv.txt": {
        "abstract": "Rapid mass transfer in a binary system can drive the accreting star out of thermal equilibrium, causing it to expand. This can lead to a contact system, strong mass loss from the system and possibly merging of the two stars. In low metallicity stars the timescale for heat transport is shorter due to the lower opacity. The accreting star can therefore restore thermal equilibrium more quickly and possibly avoid contact.  We investigate the effect of accretion onto main sequence stars with radiative envelopes with different metallicities.  We find that a low metallicity ($Z<10^{-3}$) $4\\Msun$ star can endure a 10 to 30 times higher accretion rate before it reaches a certain radius than a star at solar metallicity.  This could imply that up to two times fewer systems come into contact during rapid mass transfer when we compare low metallicity. This factor is uncertain due to the unknown distribution of binary parameters and the dependence of the mass transfer timescale on metallicity. In a forthcoming paper we will present analytic fits to models of accreting stars at various metallicities intended for the use in population synthesis models. ",
        "introduction": "The majority of stars are found in binaries, many of which can interact, for example by exchanging mass, resulting in an evolution very distinct from isolated stars. Although the fraction of stars in binaries might be different for earlier generations of stars, formed in metal-poor environments, they are certainly worth a systematic study.  In this work we discuss the effect of accretion onto main sequence stars as function of metallicity. In a second contribution to these proceedings we discuss how the ranges for different cases of mass transfer depend on metallicity \\citep{caseABC07}.  Mass transfer takes place when one of the stars exceeds a certain critical radius, the Roche lobe radius. The first phase of mass transfer is usually so fast that the accreting star is driven out of thermal equilibrium and expands \\citep{Benson70,Yungelson73}, potentially so much that the two stars come into contact. Contact binaries are not well understood, but they probably involve strong mass loss from the system and merging of the two stars.  Instead of using full binary evolution models, we choose to reduce the large parameter space of binaries by studying the behavior of models of isolated stars under controlled conditions.  We follow the approach of \\citet{Kippenhahn+Meyer77} and \\citet{Neo+ea77}, who studied the evolution of the radius of accreting main sequence stars. We extend their work, to study the effect of metallicity, using up-to-date input physics.  Here we present the results of an exploratory study and its possible implications for binaries at low metallicity. In a forth coming paper we will present analytic fits to a finer grid of models, intended for the use in population synthesis models. Expansion due to thermal timescale mass transfer is commonly neglected or taken into account using a simple approximate criterion. Our fits will provide an improvement, which is easy to implement.   \\begin{figure} \\includegraphics[ angle=-90, width=\\columnwidth]{deminkea_poster1_fig1.ps} \\caption{Radius versus mass of an initially 4 \\Msun star at solar metallicity accreting with different accretion rates. \\label{R_vs_m} } \\end{figure} ",
        "conclusions": "We find that a $4\\Msun$ low metallicity star ($Z<10^{-3}$) can endure a 10 to 30 times higher accretion rate before it expands to a certain radius compared to solar metallicity stars.  This suggests that at low metallicity fewer binaries come into contact during rapid mass transfer. A rough estimate based on very simplified assumptions indicate that the effect could be up to a factor 2, i.e. only half as many may binaries evolve into contact at low metallicity(Z < 10{-3}) compared to solar metallicity.  This factor is uncertain and probably only an upper limit as it depends on the initial distribution of binary parameters, on how the typical mass transfer rate depends on metallicity \\citep[it is probably higher at low metallicity, e.g.][]{Langer_ea00} and the binary parameters, on the specific entropy of the accreted material and on the efficiency of mass transfer \\citep[see for example][]{DeMink+ea07}.  In a forthcoming paper we will present analytic fits to our models intended for the use in population synthesis models.    \\begin{theacknowledgments} We would like to thank Peter Eggleton for providing his stellar evolution code, John Eldridge for the opacity tables and Rob Izzard for interesting discussions and suggestions. \\end{theacknowledgments}"
    },
    "0710/0710.3706_arXiv.txt": {
        "abstract": " ",
        "introduction": "The Magellanic Clouds (MCs) are interacting SBm galaxies similar to many that exist in Universe. They are the largest neighbouring satellites of the Milky Way, reflecting a typical environment of a large galaxy surrounded by satellites. They contain stars which are as old as the Universe as well as newly forming and this extended range of star formation is a highly valuable source to understand the process of formation and evolution of galaxies in general.  The MCs are overall more metal poor than the Galaxy and therefore may hold information about the Universe at its early stages. They are located at a fairly well known distance, which makes it easier to measure details of their stellar component and structure. They are also fortunately located in a region of sky only lightly affected by Galactic reddening, which translates into the capability of detecting their faint stellar populations.  The MCs belong to a complex system, the Magellanic System, which has in total four distinct components: the Large Magellanic Cloud (LMC), the Small Magellanic Cloud (SMC), the Bridge connecting the two Clouds and the Stream attached to the SMC. The latter two are predominantly formed of gas and are of tidal origin. ",
        "conclusions": "The Magellanic System has yet many challenging aspects that new surveys, with the increased quality of the coming data and new theoretical models and their ability to explain detail observations, aim to resolve in the next decade.  Prior to new facilities like GAIA, JWST and ALMA we need to exploit data from VISTA and similarly powerful telescopes at other wavelengths. Surveys like VMC will provide unique and high quality data for science and training of young astronomers."
    },
    "0710/0710.3530_arXiv.txt": {
        "abstract": "We present the first measurements of the angular correlation function of galaxies selected in the far (\\fuvcenter) and near (\\nuvcenter) Ultraviolet from the \\galex survey fields overlapping SDSS DR5 in low galactic extinction regions. The area used covers $120$ sqdeg (\\galex - MIS) down to magnitude AB $=22$, yielding a total of 100,000 galaxies. The mean correlation length is $\\sim3.7 \\pm 0.6$~Mpc and no significant trend is seen for this value as a function of the limiting apparent magnitude or between the \\galex bands. This estimate is close to that found from samples of blue galaxies in the local universe selected in the visible, and similar to that derived at $z\\simeq3$ for LBGs with similar rest frame selection criteria. This result supports models that predict anti-biasing of star forming galaxies at low redshift, and brings an additional clue to the downsizing of star formation at $z<1$. ",
        "introduction": "In the current paradigm of structure formation, the bulk of the most massive systems form in a cold dark matter-dominated universe by the merging of less massive units formed earlier. In parallel to this hierarchical evolution, recent observations point to the so-called ``downsizing'', namely the fact that in galaxies having high baryonic masses the bulk of stars formed at high redshift ($z \\gtrsim 1$), while in galaxies having low baryonic masses the bulk of stars formed at lower redshift \\citep[][and also \\citet{DeLucia_2006} and \\citet{Neistein_2006} for results from simulations]{Cowie_1996, Heavens_2004, Bundy_2006, Jimenez_2005}. The star formation efficiency shows a strong decline at $0<z<1$, as measured by the evolution of the star formation rate density \\citep{Hopkins_2006, Lilly_1996, Schiminovich_2005, Sullivan_2000, Wilson_2002}. These epochs also see the bulk of the build-up of the bimodality in galaxy properties of the local universe, which is apparent in their color distribution \\cite{Baldry_2004}, morphologies \\citep{Kauffmann_2004}, spectral class \\citep{Madgwick_2002} and spatial distribution \\citep{Budavari_2003}. Understanding the full picture is complex as this evolution is the result of the interplay of several physical processes \\citep{Faber_2005} and combine the effects of initial galaxy formation conditions (``nature'') with galaxy evolution events (``nurture'') \\citep{Kauffmann_2004}. In this context, tracers that measure over cosmic time galaxy populations selected with homogeneous physical criteria are of primary interest. They help compare observations to simulation predictions over a large range of redshifts with reduced uncertainties, and allow a study of the redshift evolution of galaxy properties derived from different surveys.  The ultraviolet (UV) range of the spectrum meets these conditions: UV luminosities provide a good measure of recent star formation within galaxies \\citep{Kennicutt_1998}, modulo attenuation by dust, and has been widely used at high redshifts to study the properties of the Lyman Break Galaxies (LBGs) \\citep{Giavalisco_2001, Shapley_2003,Steidel_1995}. As large amounts of data are now becoming available at lower redshifts as part of the GALEX surveys \\citep{Martin_2005}, the restframe UV spectral domain is presently well sampled over the full $0<z<6$ redshift range. Furthermore, comparison of results from high and low $z$ UV-selected samples is eased by the fact that the UV luminosity density fractions\\footnote{We define the UV luminosity density fraction of a given sample as the ratio of the UV luminosity density encompassed by the sample over the total UV luminosity density at the same redshift.} probed at high and low $z$ are similar \\citep[][hereafter Paper II]{Heinis_2007}, due to the strong luminosity evolution of the UV luminosity function \\citep{Arnouts_2005}. Noticeably, during the epochs probed by GALEX the properties of active star forming galaxies show a very fast evolution.  The wealth of UV-selected data now available at low redshifts enables statistical studies in the context of the downsizing of star formation, and in particular searches for links between the star formation properties and galaxy environment in terms of galaxy or dark matter density. Here we focus on the evolution with redshift of the link of star formation with Dark Matter and particularly the evolution of the class of Dark Matter halos hosting actively star forming galaxies since $z\\sim1$. This can be achieved by the study of the clustering of galaxies: at high redshift, LBGs studies show that UV selected galaxies inhabit high galaxy density regions \\citep{Steidel_1998}, and are strongly biased with respect to the underlying Dark Matter, with more actively star forming galaxies being more biased \\citep[][see \\citet{Giavalisco_2002} for a review on the properties of LBGs]{Adelberger_2005, Giavalisco_2001, Foucaud_2003}. We propose to extend such studies to low redshifts using similar selection criteria. This paper is the first in a series and presents the methods and first results of angular clustering measurements from GALEX data. The following section presents the datasets and the derivation of the redshift distributions. Section \\ref{sec_acf_methods} presents two methods to derive the angular correlation function from a set of fields, and a discussion about the behavior of these methods with respect to photometry inhomogeneity. In section \\ref{sec_acf} we present our results on the angular correlation functions and correlation lengths. To provide the crucial link to Dark Matter halos, we use the analytical \\citet{Mo_2002} formalism that we present in sec. \\ref{sec_mo_white}. We end with a short discussion in sec. \\ref{sec_discussion}.     Throughout the paper a ${\\Lambda}CDM$ cosmology is assumed with matter density $\\Omega_m = 0.3$, vacuum energy density $\\Omega_\\Lambda=0.7$, and a Hubble parameter $h=0.7$ where $H_0 = 70 $km s$^{-1}$ Mpc$^{-1}$. All correlation length values taken from the literature have been converted accordingly using equation (4) in \\citet{Magliocchetti_2000}. ",
        "conclusions": ""
    },
    "0710/0710.0207.txt": {
        "abstract": "% Aims -- Methods --Results -- Conclusions The first steps of planet formation are marked by the growth and crystallization of sub--micrometer--sized dust grains accompanied by dust settling toward the disk midplane. In this paper we explore whether the first steps of planet formation are affected by the presence of medium--separation stellar companions. We selected two large samples of disks around single and binary T~Tauri stars in Taurus that are thought to have only a modest age spread of a few Myr. The companions of our binary sample are at projected separations between $\\sim$10 and 450\\,AU with masses down to about 0.1\\,M$_{\\sun}$. We used the strength and shape of the 10\\,\\micron{} silicate emission feature as a proxy for grain growth and for crystallization respectively. The degree of dust settling was evaluated from the ratio of fluxes at two different mid--infrared wavelengths. We find no statistically significant difference between the distribution of 10\\,\\micron{} silicate emission features from single and binary systems. In addition,  the distribution of disk flaring is indistinguishable between the single and binary system samples. These results show that  the first steps of planet formation are not affected by the presence of a companion at tens of AU. ",
        "introduction": "Two--third of the G stars in the  solar neighborhood are members of multiple--star systems (e.g. \\citealt{dm91}). These binaries and multiple systems  are often found to harbor giant planets (e.g. \\citealt{bondes07}). Similarly, young low--mass pre--main sequence stars are very frequently members of multiple systems, mostly binaries \\citep{mathieu00,duchene07}. This suggests that planet formation around single stars such as our Sun may be atypical and urges us to understand the effects of stellar companions on planet formation. We  tackle this question from an observational point of view. Because numerical simulations of grain agglomeration suggest short timescales for the formation of planetesimals (only a few 10$^4$\\,yr, e.g. \\citealt{beckwith00}), it is crucial to know how  stellar companion(s) affect the dust processing in the first few million years.  Grain growth and the settling of dust grains towards the disk midplane are thought to represent the first steps in the planet--formation process (e.g. \\citealt{ls93} for a review). The study of disks around intermediate--mass stars also indicates a link between grain growth and crystallinity. High crystallinity was found only when grains larger than the dominant sub--micron interstellar grains were present \\citep{van05}. In the context of these findings and dust evolution models (e.g. \\citealt{dd04}), older disks are expected to have more processed dust (larger grains and crystals) than younger disks. In addition, their disk structure should be flatter because of the gradual settling of large dust grains towards the disk midplane. However, recent observations show that the degree of dust processing can be very different even for coeval disks around stars of similar spectral type in the same star--forming region (e.g. \\citealt{prz03,apai05}). This demonstrates that dust evolution is not uniquely controlled by stellar age and luminosity but at least one additional parameter is present.  There are two studies suggesting that stellar multiplicity could play a major role in the initial dust processing. \\citet{meeus03} found that among three coeval T~Tauri disks in the Chamaeleon~I star--forming region the closest binary system (projected separation of $\\sim$120\\,AU) sports the strongest contribution from large ($\\sim$\\,2\\,\\micron) grains and has the highest crystalline mass fraction. Similarly,  \\citet{sterzik04} pointed out that the disk of a young brown dwarf with a companion at $\\ga$\\,30\\,AU shows  more processed dust than the disk around a single brown dwarf. Although the small samples inhibit any firm conclusions, these results suggest that companions might trigger rapid dust evolution. Intuitively this may happen in different ways. A companion could speed up dust evolution by dynamically stirring the circumstellar dust grains and leading to an enhanced collision and grain growth rate (e.g. \\citealt{dubrulle95}).  In addition, the dynamical stirring may lead to an increased  mixing that could also expose larger amounts of dust to temperatures high enough ($\\ge$\\,800\\,K) to be crystallized.  In this paper we compare two carefully constructed samples of  disks around single and binary\\footnote{a few of the systems in our study are triple or quadruple systems. For simplicity, we refer to the whole sample as binaries)} stars with a narrow age spread to test the hypothesis that binary systems have disks with more processed dust and flatter structures. In Sect.~\\ref{S:sample} we describe our samples. The data reduction of the {\\it Spitzer} spectra and of the 24\\,\\micron{} MIPS photometry is presented in Sect.~\\ref{S:redu}. We summarize our results in Sect.~\\ref{S:results} and discuss in Sect.~\\ref{S:discussion} their implications on planet formation in single and binary systems.  %{\\bf COMMENT: note that stellar multiplicity does not preclude planet formation} %{\\bf COMMENT: note no significant difference in NIR excess between binaries and singles PPVI, nor in the frequencies of IRAs detections. But differences in mm fluxes (smaller masses)--similar temperatures surface densities, similar planet formation}  %{\\bf COMMENT: from PPV Duchene et al. among filed-stars the frequency of high-order multiple systems is 4 times lower than that of binaries.Shall we shift to binaries rather than multiple systems?}  %{\\bf COMMENT: there is a statistically significant difference in the mass distribution of planets around binariee and single stars: massive planets in short orbits (0.05AU) are found  mainly around tight binaries.} ",
        "conclusions": "\\label{S:discussion} Our study shows a large diversity in the 10\\,\\micron{} silicate emission features and SED slopes of T~Tauri disks. We found that neither the dust processing nor the disk flaring correlates with the multiplicity of the sources. These results are particularly interesting for two aspects that will be discussed in the following.  \\subsection{Medium--separation binaries and planet formation} A stellar companion induces tidal forces in a disk that become particularly strong at resonance points. Resonant interactions result in the excitation of density waves that can truncate a disk and act to modify the binary eccentricity (see e.g. \\citealt{lubow00} for a review). Theoretical calculations of binary--disk interactions predict that circumstellar disks will be truncated at 0.2--0.5 times the binary semimajor axis $a$, with the exact values depending on eccentricity, mass ratio, and disk viscosity \\citep{art94}. These theoretical expectations are supported by millimeter observations of  binaries tracing the optically thin dust emission and thus the total disk mass in the system. There is evidence for a diminished millimeter flux (hence disk mass) among the 1--100\\,AU binaries in comparison to wider binaries or single stars (e.g. \\citealt{mathieu00} for a review). This result is qualitatively consistent with the circumstellar disks of medium--separation binaries being tidally truncated at 0.2--0.5$a$. Two--thirds of our sources have stellar companions between 0.1\\arcsec--1\\arcsec{}, with a mean projected separation of 0.4\\arcsec{} or 56\\,AU at the distance of Taurus. Therefore, the typical truncation radius for disks in our sample is  $>$11--28\\,AU, well outside the location of Jupiter and Saturn in our Solar System. Even in other systems these outer radii are found to be devoid of giant planets  (e.g. \\citealt{kasper07}). This fact suggests that the formation of terrestrial and giant planets may proceed undisturbed in disks around medium--separation binaries even if these disks are constrained in size. %Although disks in binaries are constrained in size, the formation of terrestrial and giant planets could then proceed undisturbed if the circumstellar material is similar in surface density and temperature to that in disks around single stars. These parameters are difficult to constrain from observations, but there are indications of a similar pathway for planet formation in disks around single and medium--separation binaries.   Early investigations of young TTSs found no significant difference in the frequency of near-- and mid--infrared excess emission between single and binary star systems (e.g. \\citealt{simonprato95,jensen96}). With the 60\\micron{} IRAS flux probing dust $\\la$10\\,AU from the central star, these  measurements demonstrate that binary systems as often have disks as single stars do. Recently \\citet{monin07} analyzed the separation distributions of binaries with and without disks and found no  statistical difference. Since most of their binaries have projected separations $>$20\\,AU, their result shows that medium--and wide-- separation binaries do not have a significant effect on the circumstellar disk lifetime. Our work indicates that these disks also evolve in a similar way. The extent of dust processing in the disk surface layer and the degree of dust settling in binary disk systems do not statistically differ from those in disks around single stars. This suggests that the first few Myr of disk evolution in the terrestrial (and maybe out to the giant) planet--forming region are not affected by medium--separation stellar companions. Whether the disk evolution proceeds undisturbed for tens of millions of years until planets are fully formed cannot yet be assessed observationally. \\citet{bouwman06} estimate a mean disk dispersion timescale of $\\sim$5\\,Myr for close ($\\le$4AU) binaries in contrast to a timescale of $\\approx$9\\,Myr for single star systems. They argue that the time available to form planets in close binary systems is considerably shorter than that in disks around single stars, which may inhibit planet formation.  The only two medium--separation binaries in their sample hint for a disk dispersal timescale comparable to that of single stars suggesting a similar disk evolution for single and medium--separation binary systems over the first $\\sim$\\,10\\,Myr.   Exoplanet surveys offer us a glimpse into the frequency and properties of giant planets in multiple star systems. Recently \\citet{eu07} reported 42 planets orbiting binary  and multiple stars (see, their Table 1). \\citet{bondes07} analyze a subsample of radial velocity planet host stars with uniform planet detectability and  demonstrate that the overall frequency of giant planets in binaries is not statistically different from that of planets in single stars. However, they find indications for a lower frequency of radial velocity planets in the subgroup of close-- and medium--separation binaries ($<\\,50-100$\\,AU). In a complementary study, \\citet{desbar07} find that the mass distribution of planets in binaries  with separations $<\\,300-500$\\,AU is statistically different from that around wider binaries and single stars: Massive planets in short--period orbits are found predominantly around close-- and medium--separation binaries. Taken together, the results from the frequency and properties of exoplanets suggest that a stellar companion with separation less than a few hundred AU affects giant planet formation and/or the subsequent  migration. Numerical simulations seem to support this notion. \\citet{kley00} shows that a fairly eccentric ($e_{\\rm bin}=0.5$) stellar companion at 50--100\\,AU enhances the growth rate of a Jupiter mass planet embedded in a circumstellar disk and makes its inward migration more rapid. Recently, \\citet{kn07} confirm these trends by following the evolution of a 30\\,M$_\\earth$ protoplanet in a disk truncated by a stellar companion at 18.5\\,AU and $e_{\\rm bin}=0.36$, like the $\\gamma$~Cep binary system. Our study shows that the early evolution of protoplanetary disks surrounding binary stars is similar to that in single stars indicating that that the differences in the exoplanet properties arise in the later stages of their formation and/or migration.  Whether terrestrial planet formation is also affected by medium--separation binaries cannot be yet addressed observationally. Our study shows that the initial dust processing is not impacted by the presence of a stellar companion. Based on the fact that the build--up of planetesimals as large as the $\\sim$500--km Vesta has occurred in the first 3.8$\\pm$1.3~Myr of the Solar nebula \\citep{kleine02}, it is reasonable to speculate on the basis of our study that the formation of planetesimals in binary and single systems proceed along, if not on identical avenues.  Another indication supporting this suggestion comes from the finding of a similar incidence of debris disks in Gyr--old single and binary stars \\citep{trilling06}. If the debris dust is produced by colliding asteroids, then the similar  rate of debris dust in binaries implies that planetesimal formation is not inhibited by the presence of stellar companions. Recent simulations of the later stages of terrestrial planet formation show that rocky planets can form in a wide variety of binary systems \\citep{quinta07}. The binary periastron is the most important parameter in limiting the number of forming planets and their range of orbits. \\citet{quinta07} show that binaries with periastron $\\ga$10\\,AU, comprising most of the medium--separation binaries investigated in this paper, can form terrestrial planets over the entire range of orbits allowed for single stars. As a result more than 50\\% of the binary systems in the Milky Way \\citep{dm91} are wide enough to allow the formation of  Earth--like planets.  \\subsection{The diversity in silicate features and SEDs} Although  small sample statistics suggested a correlation between stellar multiplicity and initial dust processing \\citep{meeus03,sterzik04,sic07}, our study demonstrates that  medium--separation stellar companions do not appreciably affect the growth and crystallization of dust in circumstellar disks. Given the criteria applied to select our samples, we can also exclude that age, spectral type, and stellar environment can account for the large variety of observed silicate emission features and SED slopes in our study. There may be several other factors contributing to this diversity that will be fully explored in an upcoming contribution. In the following we briefly mention two of them:  Turbulence in circumstellar disks not only drives the accretion of gas onto the central star but also replenishes the disk atmosphere with more grains that can be larger in size. If the grains inferred from the 10\\,\\micron{} silicate emission feature reflect the level of disk turbulence, the strength of the features should depend on the stellar accretion rates. \\citet{sic07} note that stars with strong features tend to have large accretion rates in their sample of several Myr old intermediate-- and low--mass stars. This trend may be the result of turbulence determining the grain population in the disk atmosphere.  Alternatively, the trend could be due to the more massive stars (that have typically larger accretion rates) in their sample heating larger disk area and thus producing stronger silicate emission features (see, e.g.  \\citealt{kessler07}).  The tentative correlation seen in the sample of \\citet{sic07}  needs to be confirmed using  a larger and more homogeneous sample of stars with well--determined accretion rates.  %Stellar high--energy photons (UV and X--rays) may also affect  the distribution of grain sizes and the composition of solids in circumstellar disks. Low--mass TTs are strong emitters of relative hard KeV X-rays (Wolk et al. 2005) that could evaporate the grain mantle or the complete grain. Voit (1992) calculated that silicate ?? grains with a radius $<$10\\AA evpaorates completed when subject to an X-ray energu of zz, which could be reached in the disk atmospheres for grains closer than yy from the central star. This could produce a grain size distribution biased toward larger grains. I DON'T UDERSTAND THE DRAINE A EQUANTION IT GIVES UNREASONABLE RESULTS!!! It says that 50um gains can survive (not evaporate as close as 2 stellar radii.) %Low-mass TTSs also show for about 25\\% of their time intense X--ray flares that can carry energies as high as a few MeV (Wolk et al. 2005). Such flares could melt and crystallize grains and may be responsible for the  flash melted chondrules in meteorites  (e.g. Shu et al. 2001). To asses the effect of stellar high--energy photons on small grains it would be interesting to correlate stellar UV/X-ray variability with changes in the strength and shape of the 10\\,\\micron{} silicate emission features.   Different initial conditions for the collapsing cores  may also leave their imprints on the formation and evolution of circumstellar disks. This possibility has been explored by \\citet{dullemond06} to explain crystallization of dust grains in the early stages of disk evolution. In their model the level of  crystallinity depends crucially on the rotation rate of the collapsing cloud core because this determines the radius at which the infalling matter reaches the disk: rapidly rotating clouds would evolve into disks with low crystallinity, while slowly rotating clouds into disks with high crystallinity.   In this paper we explored the effect of a stellar companion on the initial growth and settling of dust grains in circumstellar disks. We constructed two large samples of disks around single and binary TTSs with a narrow age spread and a spectral type distribution for the single stars identical to that of the primary stars in the binary sample. We used the strength of the 10\\,\\micron{} silicate emission feature derived from IRS/{\\it Spitzer} spectra as a proxy for grain growth and the SED slope of circumstellar disks  as a proxy for dust settling. Our results can be summarized as follows: \\\\ \\smallskip -- there is no statistically significant difference between the distribution of 10\\,\\micron{} silicate emission features from single and binary systems. \\\\ -- the distribution of disk flaring is indistinguishable between the single and binary system samples. \\\\ \\smallskip These results show that stellar companions at projected separations of $\\ga$\\,10\\,AU do not appreciably affect the degree of crystallinity nor the degree of grain growth. Based on the combination of these and other results we argue that the formation of planetesimals and possibly terrestrial planets is not inhibited in a circumstellar disk perturbed by a medium--separation stellar companion.    %%%%%%%%%% Single star spectra \\begin{figure*} \\includegraphics[angle=90,scale=0.7]{f8.ps} \\caption{Infrared spectra for the sample of single stars.\\label{s1}} \\end{figure*} \\begin{figure*} \\includegraphics[angle=90,scale=0.7]{f9.ps} \\caption{Infrared spectra for the sample of single stars.\\label{s2}} \\end{figure*} \\begin{figure*} \\includegraphics[angle=90,scale=0.7]{f10.ps} \\caption{Infrared spectra for the sample of single stars.\\label{s3}} \\end{figure*} \\begin{figure*} \\includegraphics[angle=90,scale=0.7]{f11.ps} \\caption{Infrared spectra for the sample of single stars.\\label{s4}} \\end{figure*}  %%%%%%%%%% Binary star spectra \\begin{figure*} \\includegraphics[angle=90,scale=0.7]{f12.ps} \\caption{Infrared spectra for the sample of binary stars.\\label{b1}} \\end{figure*} \\begin{figure*} \\includegraphics[angle=90,scale=0.7]{f13.ps} \\caption{Infrared spectra for the sample of binary stars.\\label{b2}} \\end{figure*} \\begin{figure*} \\includegraphics[angle=90,scale=0.7]{f14.ps} \\caption{Infrared spectra for the sample of binary stars.\\label{b3}} \\end{figure*} \\begin{figure*} \\includegraphics[angle=90,scale=0.7]{f15.ps} \\caption{Infrared spectra for the sample of binary stars.\\label{b4}} \\end{figure*}"
    },
    "0710/0710.1473_arXiv.txt": {
        "abstract": "In this contribution we study integrated properties of dynamically segregated star clusters. The observed core radii of segregated clusters can be 50\\% smaller than the ``true'' core radius. In addition, the measured radius in the red filters is smaller than those measured in blue filters. However, these difference are small ($\\lesssim10\\%$), making it observationally challenging to detect mass segregation in extra-galactic clusters based on such a comparison. Our results follow naturally from the fact that in nearly all filters most of the light comes from the most massive stars. Therefore, the observed surface brightness profile is dominated by stars of similar mass, which are centrally concentrated and have a similar spatial distribution. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.0063_arXiv.txt": {
        "abstract": "An investigation of the dynamo instability close to the threshold produced by an ABC forced flow is presented. We focus on the on-off intermittency behavior of the dynamo and the counter-effect of the Lorentz force in the non-linear stage of the dynamo. The Lorentz force drastically alters the statistics of the turbulent fluctuations of the flow and reduces their amplitude. As a result much longer burst (on-phases) are observed than what is expected based on the amplitude of the fluctuations in the kinematic regime of the dynamo. For large Reynolds numbers, the duration time of the ``On'' phase follows a power law distribution, while for smaller Reynolds numbers the Lorentz force completely kills the noise and the system transits from a chaotic state into a ``laminar'' time periodic flow. The behavior of the On-Off intermittency as the Reynolds number is increased is also examined. The connections with dynamo experiments and theoretical modeling are discussed.  \\pacs{47.65.-d,47.20.Ky,47.27.Sd,52.65.Kj} ",
        "introduction": "Dynamo action, the self amplification of magnetic field due to the stretching of magnetic field lines by a flow, is considered to be the main mechanism for the generation of magnetic fields in the universe \\cite{mhdbooks}. To that respect many experimental groups have successfully attempted to reproduce dynamos in liquid sodium laboratory experiments \\cite{gailitis2000,gailitis2001,gailitis2004,muller2000,stieglitz2001,monchaux2007,berhanu2007}. The induction experiments \\cite{odier1998,peffley2000a,peffley2000b,frick2002,bourgoin2002,nornberg2006a,nornberg2006b,stepanov2006,volk2006,bourgoin2006} studying the response of an applied magnetic field inside a turbulent metal liquid represent also a challenging science. With or without dynamo instability the flow of a conducting fluid forms complex system, with a large degree of freedoms and a wide branch of non linear behaviors.  In this work we focus on one special behavior: the On-Off intermittency or blowout bifurcation \\cite{pomeau1980,platt1993}. On-off intermittency is present in chaotic dynamical systems for which there is an unstable invariant manifold in the phase space such that the unstable solutions have a growth rate that varies strongly in time taking both positive and negative values. If the averaged growth rate is sufficiently smaller than the fluctuations of the instantaneous growth rate, then the solution can exhibit on-off intermittency where bursts of the amplitude of the distance from the invariant manifold are observed (when the growth rate is positive) followed by a decrease of the amplitude (when the growth rate is negative). (See \\cite{sweet2001a,sweet2001b} for a more precise definition).  On-Off intermittency has been observed in different physical experiments including electronic devices, electrohydrodynamic convection in nematics, gas discharge plasmas, and spin-wave instabilities \\cite{phys-onoff}. In the MHD context, near the dynamo instability onset, the On-Off intermittency has been investigated by modeling of the Bullard dynamo \\cite{leprovost2006}. Using direct numerical simulation \\cite{sweet2001a,sweet2001b} were able to observe On-Off intermittency solving the full MHD equations for the ABC dynamo, (here we present an extended work of this particular case). On-Off intermittency has also been found recently for a Taylor-Green flow \\cite{ponty2007b}. Finally, recent liquid metal experimental results (VKS) \\cite{pinton2007} show some intermittent behavior, with features reminiscent of on-off self-generation that motivated our study.   For the MHD system we are investigating the evolution of the magnetic energy $E_b=\\frac{1}{2}\\int {\\bf b}^2 dx^3$ is given by $\\partial_t  { E_b} = \\int {\\bf b( \\cdot b \\nabla) u - \\eta (\\nabla b)^2} dx^3$. If the velocity field has a chaotic behavior in time the right hand side of the equation above can take positive or negative values and can be modeled as multiplicative noise. A simple and proved very useful way to model the behavior of the magnetic field during the on-off intermittency is using a stochastic differential equation (SDE-model) \\cite{pomeau1980,platt1993,fujisaka,Yu1990,ott1004,Platt1994,Heagy1994,Venka1995,Venka1996,aumaitre2006,aumaitre2005}:  \\begin{equation} \\partial_t E_b = (a+\\xi) E_b - NL(E_b) \\label{SDE} \\end{equation} where $ E_b$ is the magnetic energy, $a$ is the long time averaged growth rate, $\\xi$  models the noise term typically assumed to be white (see however \\cite{aumaitre2006,aumaitre2005}) and of amplitude $D$ such that $\\langle \\xi(t)\\xi(t') \\rangle = 2D\\delta(t-t')$. $NL$  is a non-linear term that guaranties the saturation of the magnetic energy to finite values typically taken to be $NL(X)=X^3$ for investigations of supercritical bifurcations or $NL(X)=X^5-X^3$ for investigations of subcritical bifurcations. Alternative, an upper no-flux boundary is imposed at $E_b=1$. In all these cases (independent of the non-linear saturation mechanism) the above SDE leads to the stationery distribution function that for $0<a<D$ has a singular behavior at $E_b=0$: $P(E_b)\\sim E_b^{a/D-1}$ indicating that the systems spends a lot of time in the neighborhood of $E_b=0$. This is singularity is the signature of On-Off intermittency. Among other predictions of the SDE model here we note that the distribution of the duration time of the ``off\" phases follows a power law behavior $PDF(\\Delta T_{off})\\sim \\Delta T_{off}^{-1.5}$, all moments of the magnetic energy follow a linear scaling with $a$, $\\langle E_b^m \\rangle \\sim a$, and for $a=0$ the set of the burst has a fractal dimension $d=1/2$ \\cite{Platt1994,Heagy1994,Venka1995,Venka1996}.  In this dynamical system eq.(\\ref{SDE}) however the noise amplitude or the noise proprieties do not depend on the amplitude of the magnetic energy. However, in the MHD system, when the non-linear regime is reached, the Lorentz force has a clear effect on the the flow such as the decrease of the small scale fluctuation, and the decrease of the local Lyapunov exponent \\cite{cattaneo1996,zienicke1998}. Some cases, the flow is altered so strongly that the MHD dynamo system jumps into an other attractor, that cannot not sustain any more the dynamo instability \\cite{brummell}. Although the exact mechanism of the saturation of the MHD dynamo is still an open question that might not have a universal answer, it is clear that both the large scales and the turbulent fluctuations are altered in the non-linear regime and need to be taken into account in a model.  \\begin{figure} \\includegraphics[width=8cm]{fig1.eps} \\caption{A typical example of a burst. The top panel shows the evolution of the kinetic energy (top line) and magnetic energy (bottom line). The bottom panel shows the evolution of the magnetic energy in a log-linear plot. During the on phase of the dynamo the amplitude of the noise of the kinetic energy fluctuations is significantly reduced. The runs were for the parameters $Gr=39.06$ and $G_M=50.40$. \\label{fig1}} \\end{figure}  Figure \\ref{fig1} demonstrates this point, by showing the evolution of the kinetic and magnetic energy as the dynamo goes through On- and Off- phases. During the On phases although the magnetic field energy is an order of magnitude smaller than the kinetic energy both the mean value and the amplitude of the observed fluctuations of the kinetic energy are significantly reduced. As a result the On-phases last a lot longer than what the SDE-model would predict. With our numerical simulations, we aim to describe which of the On-Off intermittency proprieties are affected through the Lorentz force feed-back.  This paper is structured as follows. In the next section \\ref{Nmethod} we discuss the numerical method used. In section \\ref{Onset} we present the table of our numerical runs and discuss the dynamo onset. Results for small Reynolds numbers investigating the transition from a laminar dynamo to on-off intermittency are presented in \\ref{Route}, and the results on fully developed on-off intermittency behavior are given in section \\ref{Onoff}. Conclusions, and implications on modeling and on  the laboratory experiments are given in the last section \\ref{Cons}. ",
        "conclusions": "\\label{Cons}  In this work we have examined how the on-off intermittency behavior of a near criticality dynamo is changed as the kinematic Reynolds is varied, and what is the effect of the Lorentz force in the non-linear stage of the dynamo. The predictions of \\cite{Platt1994,Heagy1994,Venka1995,Venka1996}, linear scaling of the averaged magnetic energy with the deviation of the control parameter from its critical value, fractal dimensions of the bursts, distribution of the ``off\" time intervals, and singular behavior of the pdf of the magnetic energy that were tested numerically in \\cite{sweet2001a,sweet2001b} were verified for a larger range of Kinematic Grashof numbers when On-Off intermittency was present. Note however that all these predictions are based on the statistics of the flow in the kinematic stage of the dynamo. However it was found that the Lorentz force can drastically alter the On-Off behavior of the dynamo in the non-linear stage by quenching the noise. For small Grashof numbers the Lorentz force can trap the original chaotic system in the linear regime in to a time periodic state resulting to no On-Off intermittency. At larger Grashof numbers $Gr>20$ On-Off intermittency was observed but with long durations of the ``on\" phases that have a power law distribution. These long ``on\" phases result in a pdf that peaks at finite values of $E_b$. This peak can be attributed to the presence of a subcritical instability or to the quenching of the hydrodynamic ``noise\" at the nonlinear stage or possibly a combination of the two. In principle the SDE model (eq.\\ref{SDE}) can be modified to include these two effects: a non-linear term that allows for a subcritical bifurcation and a $E_b$ dependent amplitude of the noise. There many possibilities to model the quenching of the noise, however the nonlinear behavior might not have a universal behavior and we do not attempt to suggest a specific model.  The relative range of the On-Off intermittency was found to decrease as the Reynolds number was increased possibly reaching an asymptotic regime. However the limited number of Reynolds numbers examined did not allow us to have a definite prediction for this asymptotic regime. This question is of particular interest to the dynamo experiments \\cite{gailitis2000,gailitis2001,gailitis2004,muller2000,stieglitz2001,monchaux2007,berhanu2007} that until very recently \\cite{pinton2007} have not detected On-Off intermittency . There are many reasons that could explain the absence of detectable On-Off intermittency in the experimental setups, like the strong constrains imposed on the flow \\cite{gailitis2004,muller2000} that do not allow the development of large scale fluctuations or the Earths magnetic field that imposes a lower threshold for the amplitude of the magnetic energy. Numerical investigations at higher resolution and a larger variety of flows or forcing would be useful at this point to obtain a better understanding."
    },
    "0710/0710.0580_arXiv.txt": {
        "abstract": "{Using lattice techniques we investigate the generation of long range cosmological magnetic fields during a cold electroweak transition. We will show how magnetic fields arise, during bubble collisions, in the form of magnetic strings. We conjecture that these magnetic strings originate from the alignment of magnetic dipoles associated with EW sphaleron-like configurations. We also discuss the early thermalisation of photons and the turbulent behaviour of the scalar fields after tachyonic preheating.}  \\FullConference{The XXV International Symposium on Lattice Field Theory\\\\ July 30 - August 4 2007\\\\ Regensburg, Germany}    \\begin{document} ",
        "introduction": "There have been many theoretical attempts to explain the origin of large scale cosmological magnetic fields (LSMF)~\\cite{giova}.  The main difficulty resides in understanding their correlation scale which ranges from the size of galaxies to clusters and super-clusters with an amplitude of the order of micro-gauss, pointing to a primordial origin. Following the work initiated in~\\cite{lat05}, we address this issue in the context of a cold electroweak transition taking place after a period of hybrid inflation. The EW transition has been in the heart of many proposals to address magnetogenesis, linking it in many cases with the generation of the baryon asymmetry. The results presented here resemble the mechanism proposed by Vachaspati connecting the appearance of magnetic fields to that of sphalerons and Z-strings~\\cite{vachas}.   The model we have considered is a hybrid inflation model with the bosonic field content of the Standard Model coupled, via the Higgs field, to a singlet inflaton: \\begin{eqnarray} {\\cal L} = - \\frac{1}{4}G^a_{\\mu\\nu}G^{\\mu\\nu}_a - \\frac{1}{4}F^{Y}_{\\mu\\nu}F_{Y}^{\\mu\\nu}+ {\\rm Tr}\\Big[(D_\\mu\\Phi)^\\dag D^\\mu\\Phi\\Big ]+ \\frac{1}{2}(\\partial_\\mu\\chi)^2 - V(\\Phi,\\chi)\\\\ {\\rm V} (\\Phi,\\chi) = {\\rm V}_0 + \\frac{1}{2}(g^2\\chi^2-m^2)\\,|\\Phi|^2 + \\frac{\\lambda}{4} |\\Phi|^4 + \\frac{1}{2} \\mu^2 \\chi^2 \\, \\end{eqnarray} The couplings are fixed to the standard model values for several ratios of the Higgs to W masses. For concreteness, we have fixed the inflaton to Higgs coupling by the relation: $g^2= 2\\lambda $. To solve the time evolution of the system, starting at the end of inflation, we have performed a numerical evolution based on a suitable classical approximation (more details can be found in Refs.~\\cite{marga2,lat05,articulo}). In this work we discuss the results for $\\mh = 4.65\\ \\mw$.  Results for more realistic values will be presented in~\\cite{articulo}. ",
        "conclusions": "We have analysed numerically the proposal that long range magnetic fields could be generated during a cold electroweak transition after a period of low scale hybrid inflation.  The generation mechanism is mainly based on two facts:  \\begin{itemize}  \\item At the SSB stage bubble-like structures, associated to local maxima in the Higgs-field norm, appear. Points outside the bubble front, remaining close to the false vacuum, form string like structures.  \\item Bubble collisions give rise to sphaleron-like configurations attached to the location of zeroes of the Higgs field. For  $\\theta_W \\ne 0$ these sphalerons behave as magnetic dipoles.  \\end{itemize}  These two ingredients together lead to an allignment of the sphalerons' dipoles forming magnetic string networks as we have observed in our simulations. Some results concerning the coherence and intensity of the generated magnetic fields have been presented. A further analysis of the time evolution of the magnetic field, as well as the study of the dependence of the $\\mh$ to $\\mw$ ratio will be presented in~\\cite{articulo}.  We have also discussed some features of the late time behaviour of the system, among them, thermalisation of photon radiation and turbulence in the scalar fields."
    },
    "0710/0710.2250_arXiv.txt": {
        "abstract": "Networks of reactions on dust grain surfaces play a crucial role in the chemistry of interstellar clouds, leading to the formation of molecular hydrogen in diffuse clouds as well as various organic molecules in dense molecular clouds. Due to the sub-micron size of the grains and the low flux, the population of reactive species per grain may be very small and strongly fluctuating. Under these conditions rate equations fail and the simulation of surface-reaction networks requires stochastic methods such as the master equation. However, the master equation becomes infeasible for complex networks because the number of equations proliferates exponentially. Here we introduce a method based on moment equations for the simulation of reaction networks on small grains. The number of equations is reduced to just one equation per reactive specie and one equation per reaction. Nevertheless, the method provides accurate results, which are in excellent agreement with the master equation. The method is demonstrated for the methanol network which has been recently shown to be of crucial importance. ",
        "introduction": "Chemical networks in interstellar clouds consist of gas-phase and grain-surface reactions \\citep{Hartquist1995,Tielens2005}. Reactions that take place on dust grains include the formation of molecular hydrogen \\citep{Gould1963,Hollenbach1971b} as well as reaction networks producing ice mantles and various organic molecules. Unlike gas phase reactions in cold clouds that mainly produce unsaturated molecules, surface processes are dominated by hydrogen-addition reactions that result in saturated, hydrogen-rich molecules, such as H$_2$CO, CH$_3$OH, NH$_3$ and CH$_4$. In particular, recent experiments show that methanol cannot be efficiently produced by gas phase reactions \\citep{Geppert2006}. On the other hand, there are indications that it can be efficiently produced on ice-coated grains \\citep{Watanabe2005}. Therefore, the ability to perform simulations of the production of methanol and other complex molecules on grains is of great importance \\citep{Garrod2006}. Unlike gas-phase reactions, simulated using rate equation models \\citep{Pickles1977,Hasegawa1992}, grain-surface reactions require stochastic methods such as the master equation \\citep{Biham2001,Green2001}, or Monte Carlo (MC) simulations \\citep{Charnley2001}. This is due to the fact that under interstellar conditions, of extremely low gas density and sub-micron grain sizes, surface reaction rates are dominated by fluctuations which cannot be accounted for by rate equations \\citep{Tielens1982,Charnley1997,Caselli1998,Shalabiea1998}. A significant advantage of the master equation over MC simulations is that it consists of differential equations, which can be easily coupled to the rate equations of gas-phase chemistry. Furthermore, unlike MC simulations that require the accumulation of statistical information over long times, the master equation provides the probability distribution from which the reaction rates can be obtained directly. However, the number of equations increases exponentially with the number of reactive species, making the simulation of complex networks infeasible \\citep{Stantcheva2002,Stantcheva2003}. The recently proposed multi-plane method dramatically reduces the number of equations, by breaking the network into a set of fully connected sub-networks \\citep{Lipshtat2004}, enabling the simulation of more complex networks. However, the construction of the multi-plane equations for large networks turns out to be difficult.  In this Letter we introduce a method based on moment equations which exhibits crucial advantages over the multi-plane method. The number of equations is further reduced to the smallest possible set of stochastic equations, including one equation for the population size of each reactive specie (represented by a first moment) and one equation for each reaction rate (represented by a second moment). Thus, for typical sparse networks the complexity of the stochastic simulation becomes comparable to that of the rate equations. Unlike the master equation (and the multi-plane method) there is no need to adjust the cutoffs - the same set of equations applies under all physical conditions. Unlike the multi-plane equations, the moment equations are linear and for steady state conditions can be easily solved using algebraic methods. Moreover, for any given network the moment equations can be easily constructed using a diagrammatic approach, which can be automated \\cite{Barzel2007}. ",
        "conclusions": "In summary, we have introduced a method, based on moment equations, for the simulation of chemical networks taking place on dust-grain surfaces in interstellar clouds. The method provides highly efficient simulations of complex reaction networks under the extreme conditions of low gas density and sub-micron grain sizes, in which the reaction rates are dominated by fluctuations and stochastic simulations are required. The number of equations is reduced to one equation for each reactive specie and one equation for each reaction, which is the lowest possible number for such networks. This method enables us to efficiently simulate networks of any required complexity without compromising the accuracy. It thus becomes possible to incorporate the complete network of surface reactions into gas-grain models of interstellar chemistry. To fully utilize the potential of this method, further laboratory experiments are needed, that will provide the activation energy barriers for diffusion, desorption and reaction processes not only for hydrogen but for all the molecules involved in these networks.  We thank A. Lipshtat for helpful discussions. This work was supported by the Israel Science Foundation and the Adler Foundation for Space Research."
    },
    "0710/0710.1409_arXiv.txt": {
        "abstract": "{% The origin of low-luminosity Type IIP supernovae is unclear: they have been proposed to originate either from massive ($\\sim 25~M_{\\sun}$) or low-mass ($\\sim 9~M_{\\sun}$) stars. }{% We wish to determine parameters of the low-luminosity Type IIP supernova 2003Z, to estimate a mass-loss rate of the presupernova, and to recover a progenitor mass. }{% We compute the hydrodynamic models of the supernova to describe the light curves and the observed expansion velocities. The wind density of the presupernova is estimated using a thin shell model for the interaction with circumstellar matter. }{% We estimate an ejecta mass of $14.0\\pm1.2~M_{\\sun}$, an explosion energy of $(2.45\\pm0.18)\\times10^{50}$ erg, a presupernova radius of $229\\pm39~R_{\\sun}$, and a radioactive $^{56}$Ni amount of $0.0063\\pm0.0006~M_{\\sun}$. The upper limit of the wind density parameter in the presupernova vicinity is $10^{13}$ g\\,cm$^{-1}$, and the mass lost at the red/yellow supergiant stage is $\\leq 0.6~M_{\\sun}$ assuming the constant mass-loss rate. The estimated progenitor mass is in the range of $14.4-17.4~M_{\\sun}$. The presupernova of SN~2003Z was probably a yellow supergiant at the time of the explosion. }{% The progenitor mass of SN~2003Z is lower than those of SN~1987A and SN~1999em, normal Type IIP supernovae, but higher than the lower limit of stars undergoing a core collapse. We propose an observational test based on the circumstellar interaction to discriminate between the massive ($\\sim 25~M_{\\sun}$) and moderate-mass ($\\sim 16~M_{\\sun}$) scenarios. } ",
        "introduction": "\\label{sec:intro} Type II plateau supernovae (SNe~IIP) with the plateau of $\\sim 100$ days in the light curve are believed to be an outcome of a core collapse of the $9-25~M_{\\sun}$ stars (e.g., Heger et al. \\cite{HFWLH_03}). This paradigm assuming the Salpeter mass spectrum suggests that about 66\\% of all SNe~IIP should be produced by progenitors, i.e., stars on the main sequence, in the $9-15~M_{\\sun}$ range. At present this general picture remains unconfirmed. It could be verified via the determination of ejecta masses from the hydrodynamic modeling for a sufficiently large sample of SNe~IIP. Unfortunately, only few SNe~IIP have the well-observed light curves and spectra needed to reliably reconstruct the basic SN parameters. It is not therefore surprising that up to now ejecta mass has been determined using the detailed hydrodynamic simulations for only SN~1987A and SN~1999em. It is noteworthy that, from the point of view of explosion mechanism, SN~1987A is a normal SN~IIP; it has the ejecta mass, the explosion energy, and the amount of ejected $^{56}$Ni comparable to those of SN~1999em.  Recently an interesting subclass of SNe~IIP, so called low-luminosity SNe~IIP, was selected observationally (Pastorello et al. \\cite{PZT_04}). This family is characterized by luminosities, expansion velocities, and radioactive $^{56}$Ni masses which are significantly lower than those of normal SNe~IIP. Two views on the origin of low-luminosity SNe~IIP have been proposed. Turatto et al. (\\cite{TMY_98}) have suggested the origin from massive stars, $\\geq~25~M_{\\sun}$, which presumably form black holes and eject a low amount of $^{56}$Ni; alternatively, these SNe might originate from low-mass stars,  $\\sim 9~M_{\\sun}$, which are expected to eject a low amount of $^{56}$Ni (Chugai \\& Utrobin \\cite{CU_00}; Kitaura et al. \\cite{KJH_06}).  The investigation of the origin of these SNe~IIP began with some confusion. The point is that SN~1997D, the first low-luminosity SN~IIP, was detected long after the explosion and, therefore, was erroneously claimed to possess a short ($\\sim 50$ days) plateau (Turatto et al. \\cite{TMY_98}). This, in turn, provoked a conclusion that SN~1997D originated from a low-mass ($\\sim 9~M_{\\sun}$) main-sequence star (Chugai \\& Utrobin \\cite{CU_00}). The subsequent discovery of several low-luminosity SNe~IIP with a long plateau of $\\sim 100$ days (Pastorello et al. \\cite{PZT_04}) falsified this conclusion. Until now there were no attempts to model hydrodynamically low-luminosity SNe~IIP on the basis of new observational data. The estimate of the ejecta masses of low-luminosity SN~1999br and SN~2003Z was made only using a semi-analytical model (Zampieri et al. \\cite{ZPT_03}; Zampieri \\cite{Zam_05}). The semi-analytical model, being a sensible tool for the first order estimates, cannot, however, substitute the hydrodynamic simulations.  In this paper we study the well-observed low-luminosity Type IIP SN~2003Z (Pastorello \\cite{Pas_03}; Knop et al. \\cite{KHBD_07}). Our approach is based on the hydrodynamic modeling of the light curves and expansion kinematics. This paper is organized as follows. The observational data are presented in Sect.~\\ref{sec:obsdat}. Section~\\ref{sec:model} describes the modeling procedure. Results, specifically, the SN~2003Z parameters and progenitor mass are presented in Sect.~\\ref{sec:results}. The implications of the results are discussed in Sect.~\\ref{sec:discon}.  A distance to SN~2003Z of 21.68 Mpc is adopted using the Hubble constant $H_0=70$ km\\,s$^{-1}$\\,Mpc$^{-1}$ and a recession velocity of the host galaxy NGC 2742 $v_{\\mathrm{cor}}=1518$ km\\,s$^{-1}$, corrected for the Local Group infall to the Virgo cluster and taken from the Lyon Extragalactic Data base. There are no observational signatures of the interstellar absorption in the host galaxy (Pastorello \\cite{Pas_03}) so a total extinction is taken to be equal to the Galactic value $A_{B}=0.167$ (Schlegel et al. \\cite{SFD_98}). ",
        "conclusions": "\\label{sec:discon} The goal of this study was to recover the basic parameters of SN 2003Z and to get an idea about progenitor masses of low-luminosity SNe~IIP. We estimated the ejecta mass to be $14.0\\pm1.2~M_{\\sun}$, the explosion energy $(2.45\\pm0.18)\\times10^{50}$ erg, the pre-SN radius $229\\pm39~R_{\\sun}$, and the $^{56}$Ni mass $0.0063\\pm0.0006~M_{\\sun}$. Using the ejecta/wind interaction model, we found the upper limit of the density parameter of the pre-SN wind and, assuming the constant mass-loss rate at the RSG/YSG stage, constrained the progenitor mass by the range of $14.4-17.4~M_{\\sun}$. The estimate of the wind density thus allows us to avoid the uncertainty in the progenitor mass stemming from the weak sensitivity of the hydrodynamic model to the helium core mass.  The only normal SN~IIP, studied in a similar way to SN~2003Z, is SN~1999em. Its pre-SN radius is $\\approx 500~R_{\\sun}$, the ejecta mass is $\\approx 19~M_{\\sun}$, the explosion energy is $\\approx 1.3\\times10^{51}$ erg, and the $^{56}$Ni mass is $\\approx 0.036~M_{\\sun}$ (Utrobin \\cite{Utr_07}), while the progenitor mass is $\\approx 22~M_{\\sun}$ (Chugai et al. \\cite{CCU_07}). A comparison of SN~2003Z with SN~1999em taken together with the fact that the low-luminosity SNe~IIP are very similar indicates that this variety of SNe~IIP originates from less massive progenitors than normal SNe~IIP. Parameters of three SNe~IIP --- SN~1987A (Utrobin \\cite{Utr_05}), SN~1999em (Utrobin \\cite{Utr_07}; Chugai et al. \\cite{CCU_07}), and SN~2003Z (Table \\ref{tab:sumtab}) --- determined on the basis of the similar hydrodynamic modeling, suggest a picture in which ordinary SNe~IIP originate from massive progenitors around $20~M_{\\sun}$, while low-luminosity SNe~IIP originate from less massive progenitors around $\\sim 16~M_{\\sun}$, close to the average mass of the $9-25~M_{\\sun}$ range traditionally associated with SNe~IIP. If our result for SN~2003Z is correct, then the explosion energy and the amount of ejected $^{56}$Ni should significantly increase for the progenitors between $\\sim 16~M_{\\sun}$ and $\\sim 20~M_{\\sun}$ (Fig.~\\ref{fig:nienms}). Interestingly, the empirical correlation between the explosion energy and the $^{56}$Ni mass for normal SNe~IIP has been demonstrated by Nadyozhin (\\cite{Nad_03}). It should be emphasized that these correlations, valid for SNe~IIP, may not be applicable to other SNe~II. At least SN~1994W (Type IIn event) was found to have a low amount of ejected $^{56}$Ni, $\\sim 0.015~M_{\\sun}$, but a normal explosion energy, $\\approx 1.3\\times10^{51}$ erg (Chugai et al. \\cite{CBC_04}).  \\begin{table}[t] \\caption[]{Hydrodynamic models for SN 1987A, SN 1999em, and SN~2003Z.} \\label{tab:sumtab} \\centering \\begin{tabular}{@{ } c @{ } c @{ } c @{ } c @{ } c @{ } c @{ } c @{ } c @{ }} \\hline\\hline \\noalign{\\smallskip} SN & $R_0$ & $M_{env}$ & $E$ & $M_{\\mathrm{Ni}}$ & $v_{\\mathrm{Ni}}^{max}$ & $v_{\\mathrm{H}}^{min}$ & $M_\\mathrm{ms}$ \\\\ & $(R_{\\sun})$ & $(M_{\\sun})$ & ($10^{51}$ erg) & $(10^{-2} M_{\\sun})$ & (km\\,s$^{-1}$) & (km\\,s$^{-1}$) & $(M_{\\sun})$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 87A &  35 & 18 & 1.5   & 7.65 & 3000 & 600 & 21.5 \\\\ 99em & 500 & 19 & 1.3   & 3.60 &  660 & 700 & 22.2 \\\\ 03Z & 229 & 14 & 0.245 & 0.63 &  535 & 360 & 15.9 \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table} Two alternative conjectures about the origin of low-luminosity SNe~IIP have been proposed: (1) progenitors are massive stars, $\\geq 25~M_{\\sun}$ (Turatto et al. \\cite{TMY_98}); (2) these SNe originate from low-mass stars, $\\sim 9~M_{\\sun}$ (Chugai \\& Utrobin \\cite{CU_00}; Kitaura et al. \\cite{KJH_06}). The present estimate of the progenitor mass of SN~2003Z, $14.4-17.4~M_{\\sun}$, is in the disparity with both the high-mass and low-mass scenarios.  Our estimate of the progenitor mass is essentially based on the assumption that the mass-loss rate at the RSG/YSG stage was constant. In fact, the derived wind density refers to the close vicinity of SN~2003Z $r<R_{\\mathrm w}=vt\\sim 10^{16}$ cm (where $v\\approx7\\times10^8$ cm s$^{-1}$ is the SN expansion velocity in the outer layers and $t\\sim10^7$ sec is the characteristic age of the SN). This linear scale with the wind velocity $u>10$ km\\,s$^{-1}$ corresponds to the wind history during the latest $R_{\\mathrm w}/u <300$ yr before the explosion. Most of the RSG/YSG stage, i.e. $4\\times10^5-10^6$ yr, is thus hidden from our sight. One might suggest that the mass-loss rate was substantially more vigorous in the past than immediately before the SN explosion. If this questionable possibility were the case, the $25~M_{\\sun}$ progenitor would lose $\\sim 10~M_{\\sun}$, and the high-mass scenario for SN~2003Z would be saved. Of note, the essentially higher wind density of pre-SN in this case suggests that the dense wind at the large radii could be revealed via radio and X-ray observations at the large age, $t\\geq10$ yr.  A remarkable property of all three SNe is strong mixing between the helium core and the hydrogen envelope indicated by a low minimal velocity of hydrogen matter (Table \\ref{tab:sumtab}). Generally, the model with the unmixed helium core shows a bump in the light curve at the end of the plateau (Utrobin \\cite{Utr_07}); the absence of the bump in the available light curves of SNe~IIP, in addition to the narrow-topped H$\\alpha$ emission at the nebular phase, indicates that substantial mixing between the helium core and the hydrogen envelope is a universal phenomenon for SNe~IIP. This mixing could be induced by the convection during the growth of helium core (Barkat \\& Wheeler \\cite{BW_88}) or/and during the SN explosion (Kifonidis et al. \\cite{KPSJM_06}).  \\begin{figure}[t] \\resizebox{\\hsize}{!}{\\includegraphics{fig9.eps}} \\caption{% Explosion energy (\\textbf{a}) and $^{56}$Ni mass (\\textbf{b}) versus progenitor mass for three core-collapse SNe. } \\label{fig:nienms} \\end{figure} It is noteworthy that the radii of all three pre-SNe~IIP (Table \\ref{tab:sumtab}) are notably lower than a radius of typical massive RSG, e.g., the radius of $\\approx 800~R_{\\sun}$ for $\\alpha$ Ori (Harper et al. \\cite{HBL_01}). We have already shown that the SN~2003Z pre-SN was the YSG rather than the RSG. For SN~1999em the $22~M_{\\sun}$ pre-SN with the luminosity of $\\approx 10^5~L_{\\sun}$ is characterized by the effective temperature of $\\sim 4700$~K, which is intermediate between the RSG and YSG values. Interestingly, the pre-SN of SN~2004et, a normal SN~IIP, was probably a YSG as well (Li et al. \\cite{LVFC_05}). These cases suggest that not only the pre-SN~1987A, but a possibly significant fraction of pre-SNe~IIP, experience an excursion on the Hertzsprung-Russell (HR) diagram from the RSG towards the YSG before the explosion. Note, some pre-SNe~IIP, when becoming the YSG, could get into the instability strip and manifest themselves before the explosion as long period ($P\\sim50-100$ days) Cepheids. If this is the case, the recovery of the Cepheid period of a pre-SN on the basis of a set of pre-explosion observations could provide us with an additional tool for the mass determination from the pulsation period.  The recovered progenitor mass of SN~2003Z raises a crucial question: what happens to the $9-15~M_{\\sun}$ stars which make up $\\approx66$\\% of all the stars from the $9-25~M_{\\sun}$ mass range. Two conceivable suggestions are: (A) all the $9-15~M_{\\sun}$ stars explode as low-luminosity, or even fainter, SNe~IIP; (B) low-luminosity SNe~IIP originate only from a narrow range of progenitors around $\\sim 16~M_{\\sun}$, while the rest of the $9-15~M_{\\sun}$ stars produce different varieties of SNe~II. A large sample of SNe~II which is well observed and studied is certainly needed to distinguish between the options A and B. It should be emphasized that the initial peak, the end of the plateau, and the radioactive tail are of crucial importance for the recovery of reliable SN~II parameters.  An alternative approach to determine the pre-SN mass exploits archival pre-explosion images of host galaxies and stellar evolution tracks on the HR diagram. Interestingly, using the archival HST images Van Dyk et al. (\\cite{VLF_03}) and Maund \\& Smartt (\\cite{MS_05}) derived an upper limit of $\\sim 15~M_{\\sun}$ and $\\sim 12~M_{\\sun}$, respectively, for the progenitor of the low-luminosity SN~1999br. Given a similarity of known low-luminosity SNe~IIP these estimates make the high-mass scenario unlikely.  It should be emphasized that a test of the mass determination method based on the pre-explosion images is needed; for instance, the pre-SN light could be partially absorbed in a dusty circumstellar envelope. In this regard it would be of top priority to determine independently the progenitor mass both from the pre-discovery images and from the hydrodynamic simulations. Among SNe~IIP, except for SN~1987A, only SN~2004et and SN~2005cs were detected in the pre-discovery images (Li et al. \\cite{LVFC_05}; Maund et al. \\cite{MSD_05}; Li et al. \\cite{LVF_06}) and have the light curves and spectra appropriate for the hydrodynamic modeling as well (Sahu et al. \\cite{SASM_06}; Misra et al. \\cite{MPC_07}; Pastorello et al. \\cite{PST_06}; Tsvetkov et al. \\cite{TVS_06}). These two challenging cases require a detailed study."
    },
    "0710/0710.1123_arXiv.txt": {
        "abstract": "Optical and near-infrared spectroscopy of the newly discovered peculiar L dwarf {\\name} are presented. Folkes et al.\\ identified this source as a high proper motion L9$\\pm$1 dwarf based on its strong H$_2$O absorption at 1.4~$\\micron$. We find that the optical spectrum of {\\namesh} is in fact consistent with that of a normal L4.5 dwarf with notably enhanced FeH absorption at 9896~{\\AA}. However, its near-infrared spectrum is unusually blue, with strong H$_2$O and weak CO bands similar in character to several recently identified ``blue L dwarfs''.  Using {\\namesh} as a case study, and guided by trends in the condensate cloud models of Burrows et al.\\ and Marley et al., we find that the observed spectral peculiarities of these sources can be adequately explained by the presence of thin and/or large-grained condensate clouds as compared to normal field L dwarfs. Atypical surface gravities or metallicities alone cannot reproduce the observed peculiarities, although they may be partly responsible for the unusual condensate properties. We also rule out unresolved multiplicity as a cause for the spectral peculiarities of {\\namesh}. Our analysis is supported by examination of {\\em Spitzer} mid-infrared spectral data from Cushing et al.\\ which show that bluer L dwarfs tend to have weaker 10~$\\micron$ absorption, a feature tentatively associated with silicate oxide grains. With their unique spectral properties, blue L dwarfs like {\\namesh} should prove useful in studying the formation and properties of condensates and condensate clouds in low temperature atmospheres. ",
        "introduction": "L dwarfs comprise one of the two latest-type spectral classes of very low mass stars and brown dwarfs, spanning masses at and below the hydrogen burning minimum mass (see \\citealt{kir05} and references therein). They are inexorably linked to the presence and properties of liquid and solid condensates which form in their cool photospheres (e.g., \\citealt{tsu96,bur99,ack01,all01,coo03,tsu05}). These condensates significantly influence the spectral energy distributions and photospheric gas abundances of L dwarfs, by removing gaseous TiO and VO from the photosphere and enabling the retention of atomic alkali species (e.g., \\citealt{feg96,bur99,lod02}). Weakened {\\wat} absorption through backwarming effects (e.g., \\citealt{jon97,all01}) and red near-infrared colors ($J-K$ $\\approx$ 1.5--2.5; \\citealt{kir00}) also result from condensate opacity. In addition, periodic and aperiodic photometric variability observed in several L dwarfs has been associated with surface patchiness in photospheric condensate clouds (e.g., \\citealt{bai99,bai01,gel02,moh02}). Condensate abundances at the photosphere appear to reach their zenith amongst the mid- and late-type L dwarfs \\citep{kir99,cha00,ack01} before disappearing from the photospheres of cooler T dwarfs \\citep{mar96,tsu96b,all01,cus06}. The abundances of photospheric condensates, their grain size distribution, and the radial and surface structure of condensate clouds may vary considerably from source to source, as well as temporally for any one source, and the dependencies of these variations on various physical parameters are only beginning to be explored \\citep{hel01,woi03,kna04}.  With hundreds of L dwarfs now known,\\footnote{A current list is maintained at \\url{http://dwarfarchives.org}.} groupings of peculiar L dwarfs -- sources whose spectral properties diverge consistently from the majority of field objects -- are becoming distinguishable. Examples include young, low surface gravity brown dwarfs \\citep{mcg04,kir06,all07,cru07} and metal-poor L subdwarfs \\citep{me0532,lep1610,giz06,megmos}. There also exists a class of peculiar ``blue'' L dwarfs \\citep{cru03,cru07,kna04}, roughly a dozen sources exhibiting normal optical spectra but unusually blue near-infrared colors and strong near-infrared {\\wat}, FeH and {\\ki} features.  Various studies have attributed these peculiarities to subsolar metallicity, high surface gravity, unresolved multiplicity and peculiar cloud properties \\citep{giz00,cru03,cru07,mcl03,mewide3,kna04,chi06,fol07}. Any one of these characteristics may impact the presence and character of condensates and condensate clouds in low temperature atmospheres.  In an effort to identify new nearby and peculiar L dwarfs, we have been searching for late-type dwarfs using near-infrared imaging data from the Deep Near Infrared Survey of the Southern Sky (DENIS; \\citealt{epc97}). One of the objects identified in this program is DENIS~J112639.9$-$500355, a bright source which was concurrently discovered by \\citet{fol07} in the SuperCOSMOS Sky Survey \\citep[hereafter SSS]{ham01a,ham01b,ham01c} and the Two Micron All Sky Survey (hereafter 2MASS; \\citealt{skr06}). It is designated {\\name} in that study, and we refer to the source hereafter as {\\namesh}. Based on its blue near-infrared colors and deep {\\water} absorption bands, \\citet{fol07} concluded that {\\namesh} is a very late-type L dwarf (L9$\\pm$1) which may have unusually patchy or thin condensate clouds.  In this article, we critically examine the observational properties of {\\namesh} to unravel the origins of its spectral peculiarities, and examine it as a representative of the blue L dwarf subgroup.  Our identification of {\\namesh} and a slightly revised determination of its proper motion using astrometry from the SSS, DENIS and 2MASS catalogs are described in $\\S$~2. Optical and near-infrared spectroscopic observations and their results are described in $\\S$~3, along with  determination of the optical and near-infrared classifications of {\\namesh} and estimates of its distance and space kinematics.  In $\\S$~4 we analyze the properties of {\\namesh} and blue L dwarfs in general, considering metallicity, surface gravity, condensate cloud and unresolved multiplicity effects.  We also introduce a new near-infrared {\\wat} index that eliminates discrepancies between optical and near-infrared types for these sources. Results are discussed and summarized in $\\S$~5. ",
        "conclusions": "Our analysis in $\\S$~4.2 leads us to conclude that the spectral peculiarities of {\\namesh} and other blue L dwarfs have their immediate cause in condensate cloud effects, specifically the presence of thin, patchy or large-grained condensate clouds at the photosphere.  Subsolar metallicities and high surface gravities in of themselves cannot reproduce the observed spectral peculiarities of these sources. However, it is clear that these latter physical properties must play a role in determining the cloud characteristics of blue L dwarfs. Lower metallicities reduce the metal species available to form condensates, resulting in less condensate material overall.  Higher surface gravities may increase the sedimentation rate of condensate grains, potentially resulting in thinner clouds. The large tangential velocities and absence of {\\lii} absorption in the three blue L dwarfs {\\namesh}, 2MASS~J1300+1921 and 2MASS~J1721+3344 support the idea that these sources may be relatively old and possibly slightly metal-poor. However, the influence of other physical parameters on condensate cloud properties must also be considered, including rotation rates, vertical upwelling rates (e.g., \\citealt{sau06}) and possibly magnetic field strengths.  An assessment of how these fundamental physical parameters influence the properties of condensate clouds in low-temperature atmospheres is the subject of ongoing theoretical investigations (e.g., \\citealt{hel01,woi03}; M.\\ Marley, in preparation). Empirical studies are also necessary, particularly those focused on well-characterized samples of blue (and red) L dwarfs. To this end, Table~\\ref{tab_blue} lists all blue L dwarfs currently reported in the literature. We anticipate that this list will grow as near-infrared spectroscopic follow-up of L dwarfs continues."
    },
    "0710/0710.1065_arXiv.txt": {
        "abstract": "We use a semi-analytic circumstellar disk model that considers movement of the snow line through evolution of accretion and the central star to investigate how gas giant frequency changes with stellar mass. The snow line distance changes weakly with stellar mass; thus giant planets form over a wide range of spectral types. The probability that a given star has at least one gas giant increases linearly with stellar mass from 0.4\\,$M_\\odot$ to 3\\,$M_\\odot$. Stars more massive than 3\\,$M_\\odot$ evolve quickly to the main-sequence, which pushes the snow line to 10--15\\,AU before protoplanets form and limits the range of disk masses that form giant planet cores. If the frequency of gas giants around solar-mass stars is 6\\%, we predict occurrence rates of 1\\% for 0.4\\,$M_\\odot$ stars and 10\\% for 1.5\\,$M_\\odot$ stars. This result is largely insensitive to our assumed model parameters. Finally, the movement of the snow line as stars $\\gtrsim$2.5\\,$M_\\odot$ move to the main-sequence may allow the ocean planets suggested by \\citeauthor{2004Icar..169..499L} to form without migration. ",
        "introduction": "\\label{sec:intro}  In the last ten years, the discovery of more than 200 extra-solar planets,\\footnote{http://vo.obspm.fr/exoplanetes/encyclo/encycl.html} and more than 200 debris disks,\\footnote{http://www.roe.ac.uk/ukatc/research/topics/dust/identification.html } suggests that planet formation is a common and robust process. Planet masses inferred from debris disks range from terrestrial to Jovian, at distances as great as tens of AU from the central star \\citep[e.g.][]{2004AJ....127..513K,2005ApJ...619L.187G}. The nature and sensitivity of radial velocity surveys means that most of the planets are $\\sim$Jupiter mass gas giants in close orbits around Sun-like stars. However, recent discoveries as diverse as icy $\\sim$Neptune-mass planets orbiting M dwarfs \\citep[e.g.][]{2005ApJ...634..625R}, and debris disks around A-type stars \\citep[e.g.][]{2005ApJ...620.1010R} show that planet formation occurs over a wide range of spectral types.  Current theory suggests that planets form in similar ways around all stars. Thus, the increasing diversity of stellar hosts and planetary systems provides an opportunity to test these theories. For this reason, the types of planets most likely to form around stars of differing spectral types has become a renewed area of study \\citep[e.g.][]{2005ApJ...626.1045I,2006ApJ...643..501B,2006A&A...458..661K,2006ApJ...650L.139K}, after the idea was first explored by \\citeauthor{1988MNRAS.235..193N} nearly 20 years ago \\citep{1987MNRAS.224..107N,1988MNRAS.230..551N,1988MNRAS.235..193N}.  Theories of Solar System formation generally include the ``snow line,'' where ices condense from the nebular gas. The snow line distance is usually fixed in a disk with a time independent surface density and temperature profile around a main-sequence star \\citep[e.g.][]{2005ApJ...626.1045I}. In a more realistic picture, the disk and stellar properties evolve considerably during the 1--10\\,Myr pre--main-sequence (PMS) lifetime when planets probably form \\citep[e.g.][]{1987Icar...69..249L,1996Icar..124...62P}. As the disk temperature evolves with time, movement of the snow line may therefore influence the properties of theoretical planetary systems \\citep[e.g.][]{2006ApJ...650L.139K,2007ApJ...654..606G}.  Here, we begin to develop a time dependent model for the formation of gas giant cores that considers the PMS evolution of the star and surrounding accretion disk. We introduce a simple semi-analytic disk model, based on the ``minimum mass solar nebula,'' that links movement of the snow line through evolution of disk accretion and stellar luminosity. In contrast to previous studies \\citep[e.g.][]{2005ApJ...626.1045I,2006A&A...458..661K}, our analysis suggests that gas giant formation around stars more massive than the sun is more likely than around less massive stars.  We cover the background important to our story in \\S \\ref{sec:background}, consider the snow line in \\S \\ref{sec:motivation}, and outline our model in \\S \\ref{sec:model}. We present our results in \\S \\ref{sec:results}, and discuss and conclude in \\S \\ref{sec:discussion} and \\S \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  We describe a model for the evolution of the snow line in a planet forming disk, and apply it over a range of stellar masses to derive the probability distribution of gas giants as a function of stellar mass. The two main ingredients for our model are a prescription for movement of the snow line due to accretion and PMS evolution, and rules that determine whether protoplanets are massive enough, and form early enough, to become gas giants.  The snow line distance generally moves inward over time. With our prescription for the accretion rate, accretion dominates over irradiation for stars with $M_\\star \\lesssim 2\\,M_\\odot$. For $\\gtrsim$3\\,$M_\\odot$ stars, irradiation dominates at times $\\gtrsim$1\\,Myr as the star moves up to its main-sequence luminosity. The transition is at a few Myr for $\\sim$2\\,$M_\\odot$ stars. Over the wide range of observed accretion rates for any fixed stellar mass, the snow line in some disks may be set entirely by irradiation.  The snow line generally sets where the innermost gas giant cores form. In relatively massive disks around intermediate mass stars, rocky cores form interior to the snow line. The location of the outermost core is always set by the gas dissipation timescale. The range of disk masses that form cores, and the radial width of the region in the disk where they form, increase with stellar mass. Lower mass disks produce failed icy cores, which are probably similar to Uranus, Neptune, and the observed ``super-Earths.''  The outward movement of the snow line as stars more massive than the Sun reach the main-sequence, and as the disk becomes optically thin, allows the ocean planets suggested by \\citet{2004Icar..169..499L} to form \\emph{in situ}. The change in disk temperature is only large enough for these planets to harbour oceans around stars $\\gtrsim$2.5\\,$M_\\odot$.  Our model includes several poorly determined parameters, which current and future facilities will investigate. While there are current resolved studies of gaseous disks \\citep[e.g.][]{2007arXiv0704.1481B}, the next generation of telescopes such as GMT and ALMA will provide more information on surface density profiles, and how disk properties change with stellar mass and age. These studies will help to constrain input parameters for our model. We have shown that the time dependence of the snow line in part determines where gas giant cores form. This result should motivate future studies of planet formation in disks whose properties change with time.  The subsequent evolution of isolated cores is beyond the scope of this paper, but further work that investigates the growth and dynamical evolution of these objects can investigate the diversity of resulting system structures.  Given an initial distribution of disk masses, the probability that a star has at least one gas giant increases linearly with stellar mass from 0.4\\,$M_\\odot$ to 3\\,$M_\\odot$. If the frequency of gas giants around solar-mass stars is 6\\%, we predict an occurrence rate of 1\\% (10\\%) for 0.4\\,$M_\\odot$ (1.5\\,$M_\\odot$) stars. This result is largely insensitive to changes in our model parameters.  In contrast to the \\citet{2005ApJ...626.1045I} model, where it is hard to form observable gas giants above 1\\,$M_\\odot$, our model predicts a peak at $\\sim$3\\,$M_\\odot$ because we include disk and PMS evolution in our snow line derivation. However, our model does not include migration, so our prediction applies to observable and currently undetectable gas giants.  Though sample numbers are small, it appears that observable gas giant frequency increases with stellar mass across a wide range of host masses \\citep{2007arXiv0707.2409J}. Larger samples of stars that host giant planets, particularly low and intermediate-mass stars, will solidify this result. These studies, and the extension of the results to a wider range of semi-major axes, will provide a basis for comparison with our model predictions."
    },
    "0710/0710.5271_arXiv.txt": {
        "abstract": "The unveiled main-sequence splitting in $\\omega$ Centauri as well as NGC 2808 suggests that matter highly-enriched in He (in terms of its mass fraction $Y\\sim0.4$) was produced and made the color of some main-sequence stars bluer in these globular clusters (GCs). The potential production site for the He-rich matter is generally considered to be massive AGB stars that experience the second dredge-up. However, it is found that massive AGB stars provide the matter with $Y\\sim 0.35$ at most, while the observed blue-shift requires the presence of $Y\\sim 0.4$ matter. Here, we show that extra mixing, which operates in the red giant phase of stars less massive than $\\sim2\\,M_{\\odot}$, could be a mechanism that enhances He content in their envelopes up to $Y\\sim 0.4$. The extra mixing is supposed to be induced by red giant encounters with other stars in a collisional system like GCs.  The $Y\\sim 0.4$ matter released in the AGB phase has alternative fates to (i) escape from a GC or (ii) be captured by kinematically cool stars through encounters. The AGB ejecta in $\\omega$ Cen, which follows the latter case, can supply sufficient He to cause the observed blue-shift. Simultaneously, this scheme generates the extreme horizontal branch, as observed in $\\omega$ Cen in response to the higher mass loss rates, which is also caused by stellar encounters. ",
        "introduction": "The discovery of a split in the main sequence (MS) of \\omcen\\  \\citep{Bedin04} has opened a new window in the study of Galactic globular clusters (GCs). The observed fact that stars of the blue MS (bMS) in \\omcen\\ exhibit slightly stronger absorption of iron lines on average, in comparison with the red MS (rMS) \\citep{Piotto05} which strongly implies that the origin of bMS is attributable to the enhancement of He inside bMS stars. Subsequently, the recent HST observation has provided us with the second sample of MS splitting, NGC 2808 \\citep{Piotto07}.  This GC essentially exhibits no dispersion of [Fe/H] and thus the argument of He enhancement in bMS stars has been reinforced. A comparison between the observed color-magnitude diagrams (CMDs) and synthetic population models suggests that bMS He abundance is enhanced to $Y\\sim 0.4$ for both of two GCs \\citep{Norris04, Lee05, DAntona05, Piotto07}. \\citet{Tsujimoto07} have argued that if the surface of a star on the MS is polluted with  the $Y=0.4$ matter with the mass of $\\sim 0.1 \\msun$, the star can move to a position on the observed bMS in the CMD. This relaxes a severe demand on the amount of He possibly supplied from asymptotic giant branch (AGB) stars, in contrast to the idea that the bMS stars are born from pure AGB ejecta consisting of $Y\\sim 0.4$. Similarly, \\citet{Newsham07} have insisted that surface pollution may explain the high He content of bMS stars.  It should be stressed that $Y\\sim 0.4$ matter is necessary to realize the bMS stars in these GCs. Since massive AGB stars can enhance the He abundance in their envelopes through the second dredge-up in the early AGB phase, they have been considered a possible production site for the $Y\\sim 0.4$ matter \\citep{DAntona05}. However, previous studies indicated that the resultant abundance of He in the envelope of any AGB star does not exceed $Y \\sim 0.35$ \\citep[see e.g.,][]{vandenHoek97}. Therefore, such a He yield from massive AGB stars may not be sufficient to split the MS as observed. In the subsequent section, we will discuss this issue and conclude that it is highly implausible for a massive AGB star to enhance He abundance up to $Y\\sim 0.4$ in its envelope.  Then, we should search for an alternative mechanism that enhances the He abundance more efficiently than the second dredge-up. We believe that this process, which enhances He abundance, should be accompanied by some other elemental signatures. This is reminiscent of the abundance anomaly of red giants in GCs. The abundance anomalies observed for red giants, such as the abundance variations of CNO elements and the O-Na anticorrelation, are common attributes associated with GCs including \\omcen\\  \\citep[see][]{Gratton04}. One of the proposed mechanisms to explain these anomalies is a deep mixing, so-called extra mixing,  on the red giant stage \\citep{Sweigart79, Langer93, Charbonnel98, Fujimoto99,Aikawa01,Chaname05,Suda06}. The extra mixing is assumed to be caused by some kind of rotation-induced mixing, driven by shear instability around the border of the He core \\citep[see][]{Fujimoto88,Zahn92}. Since such a deep mixing naturally draws He from the core, the extra mixing is a promising mechanism of producing He-rich matter. As a mechanism of mixing, most of the previous works assumed a continuous mixing caused by the meridional circulation, while Fujimoto and his coworkers insisted that a flash of H brought into a degenerate core by the rotation-induced mixing drives the intermittent mixing.  In this study, we propose that extra mixing, which occurs in the red giant phase of a star, enhances He abundance in the envelope and ejects He-rich matter when the star evolves to a white dwarf through mass loss. In this scenario, the driving force of the extra mixing is generated through interactions between the red giant and dwarfs in the central region of \\omcen. Our detailed calculations reveal that stars with the masses less than $\\sim$2 \\msun\\ can enhance He abundance up to $Y\\sim 0.4$, provided that the extra mixing operates in the red giant phase.  Based on our scenario, we discuss other conspicuous characteristics of \\omcen, the extremely blue horizontal branch (HB), which is observed in some other Galactic GCs. Although factors that determine HB morphology is still an open question and is known as the second parameter problem, the most effective parameter that affects the location of a HB star in the CMD is the mass $M_{\\rm env}$ of the envelope. A major factor to determine $M_{\\rm env}$ is the mass loss during the red giant phase. We predict that encounters induce a high mass loss rate through a gain of angular momentum and promote the presence of an extremely blue HB population. In the end, two mysteries in \\omcen\\  of MS splitting and HB morphology are discussed in a unified framework of encounters of dwarfs with AGB ejecta and those of red giants with dwarfs. ",
        "conclusions": "We first propose that He-rich matter is synthesized in red giants with the masses of 0.8 - 2 $\\msun$, which experience encounters with other stars in GCs, leading to the extra mixing during the RGB phase in their envelopes. On the other hand, massive AGB stars are unable to produce the He abundance as much as $Y \\sim 0.4$ because (1) their large envelopes work as a buffer against He enrichment and (2) the amount of dredged-up matter is limited by thinning the He-burning shell, due to the growth of C-O core. Therefore, a $Y\\sim 0.4$ matter is a product unique to GCs. In some GCs, this matter is retained in their gravitational potential for a prolonged time and is accreted by kinematically cool stars, although this may not be the case for most GCs. In addition, stellar encounters in GCs induce the extra mixing and increase the mass loss rate during the RGB phase. As a result, such stars are predicted to evolve to significantly blue HB stars. In conclusion, MS splitting is inclined to be accompanied by the existence of an extremely blue HB, as observed in \\omcen\\ and NGC 2808. It should be, however, noted that our models might make a variation in the final He abundance as a result of extra mixing induced by various amounts of angular momentum transfered through encounters. If the variation is too large, then it will erase the MS splitting and lead to a broad MS. Further investigations of the extra mixing model in terms of the MS splitting are surely awaited."
    },
    "0710/0710.0662_arXiv.txt": {
        "abstract": "The study of Wolf-Rayet stars plays an important role in evolutionary theories of massive stars. Among these objects, $\\sim 20\\%$ are known to be in binary systems and can therefore be used for the mass determination of these stars. Most of these systems are not spatially resolved and spectral lines can be used to constrain the orbital parameters. However, part of the emission may originate in the interaction zone between the stellar winds, modifying the line profiles and thus challenging us to use different models to interpret them. In this work, we analyzed the He{\\sevensize II}$\\lambda$4686\\AA \\ + C{\\sevensize IV}$\\lambda$4658\\AA \\ blended lines of WR30a (WO4+O5) assuming that part of the emission originate in the wind-wind interaction zone. In fact, this line presents a quiescent base profile, attributed to the  WO wind, and a superposed excess, which varies with  the orbital phase along the  4.6 day period. Under these assumptions, we were able to fit the excess spectral line profile and central velocity for all phases, except for the longest wavelengths, where a spectral line with constant velocity seems to be present. The fit parameters provide the eccentricity and inclination of the binary orbit, from which it is possible to constrain the stellar masses. ",
        "introduction": "Massive stars are known to drive strong winds, which are responsible for the transfer of a large amount of the stellar mass to the interstellar medium, contributing to the feedback of chemical elements, and to the creation of  cloud cavities in which these objects are found. Typically, O stars present mass-loss rates of $\\dot{M}_{\\rm O} \\sim 10^{-6} - 10^{-5}$ M$_\\odot$ yr$^{-1}$ and wind velocities of $v_{\\rm O} \\sim 2000 - 3500$ km s$^{-1}$, while Wolf-Rayet (WR) stars present $\\dot{M}_{\\rm WR} \\sim 10^{-6} - 10^{-4}$ M$_\\odot$ yr$^{-1}$ and $v_{\\rm WR} \\sim 1000 - 4000$ km s$^{-1}$ (Nugis \\& Lamers 2000, Lamers 2001).  In  massive binary systems in which  both stars present high mass-loss rates and high-velocity winds, the collision of the winds will occur. Therefore, a contact surface is formed where the momenta of the two winds are equal, surrounded by two shocks. The post-shocked gas, cools as it flows along the contact surface and is responsible for strong free-free emission at X-ray and radio wavelengths. X-rays in massive binary systems are orders of magnitude higher than those observed in single massive stars, and are used as an indication of binarity.  UV and optical lines can also indicate the binary nature of a given object, when they present periodic profile variations.  If the lines are of photospheric or atmospheric origin, as the stars move along their orbit, they  suffer Doppler shifts, which   depend on the orbital phase and inclination. However, some massive objects show periodic variable line profiles that cannot be explained under these assumptions.  Seggewiss (1974) noted that the binary system WR79 presented two peaks, superimposed to the C{\\sevensize III} emission line, which changed their position and intensity with time. Typically, in double line spectroscopic binaries, each peak moves in a different direction, indicating opposite velocity components along the line of sight for each star; however in WR79 both peaks moved in the same direction. L$\\rm \\ddot{u}$hrs (1997) presented a model in which the two peaks were not produced by the stellar photosphere, but were generated by the flowing gas at the contact surface between the two strong shocks. This model reproduced well the data for WR 79, but failed to reproduce the line profiles of other WR binary systems, among them WR 30a (Bartzakos, Moffat \\& Niemela 2001). Falceta-Gon\\c calves, Abraham \\& Jatenco-Pereira (2006) improved  L$\\rm \\ddot{u}$hrs' model introducing more realistic parameters, as stream turbulence and  gas opacity, to account for the line broadening and peak displacement.  In the present work, we applied this model to WR 30a, which was classified as a WO4+O5 binary system (Moffat \\& Seggewiss 1984, Crowther et al. 1998); its binary nature was reported by Niemela (1995) based on spectral-line radial  velocities obtained with a high temporal resolution. Later, Gosset et al. (2001) presented a detailed analysis of the spectra of WR30a for several epochs and were able to confirm the binary hypothesis and to determine the period of $P \\sim 4.6$ days. They also noted strong line-profile variations, which made more difficult  the determination of the stellar mass ratio and the orbital inclination from standard methods. They concluded that the C{\\sevensize IV}$\\lambda$4658\\AA \\ line-profile variations were related to wind-wind collision processes, but did not model the lines under such an assumption. The same conclusion was reached by Bartzakos, Moffat \\& Niemela (2001) using the C{\\sevensize IV}$\\lambda$5801\\AA \\ line. Also, Paardekooper et al. (2003) presented photometric measurements at $V$ and $B$ bands; the light curves confirmed the period obtained by Gosset et al. (2001), but also showed higher frequency variability in the $V$ band, with timescale  of hours. They concluded that this could be due to the strong variability of the C{\\sevensize IV}$\\lambda$5801\\AA \\ line, possibly related to the wind-wind interaction.  In the present work, we tested the wind-wind shock emission hypothesis on the C{\\sevensize IV}$\\lambda$4658 \\AA \\ excess line profile variations measured by Gosset et al. (2001) using the model developed by Falceta-Gon\\c calves, Abraham \\& Jatenco-Pereira (2006), which is briefly described in Section 2. In Section 3, we show the results obtained for WR 30a and present a brief discussion, followed by the conclusions in Section 4. ",
        "conclusions": "In this work we presented an application of the wind-wind shock emission model proposed by Falceta-Gon\\c calves et al. (2006) to the line profile variations of WR 30a. In this model, the wind-wind shock structure can be represented by a cone, along which the shocked material flows. This gas will emit radiation, cool down  and eventually reach recombination temperatures. The observed emission lines will then suffer Doppler shifts due to the stream velocity component along the line of sight. During the orbital movement, the cone position will change, as well as the radial velocity, which will cause the line profile variations observed in several massive binary systems.  Gosset et al. (2001) obtained detailed spectra of WR30a during more than 30 days. They determined the orbital period ($P = 4.6$d), and obtained the radial velocity curve for the O-star. Regarding the WR component, they found that the blended He{\\sevensize II}$\\lambda$4686\\AA \\ and C{\\sevensize IV}$\\lambda$4658\\AA \\ lines showed a variable excess emission. In the present paper we modeled this variable emission, being able to reproduce the variations except for the red part of the profiles, which seemed to be unchanged in velocity and were probably generated in the stellar wind of the WR star instead of in the shock region.  The best-fitting result was obtained for  $\\beta = 50^\\circ$, $i = 20^\\circ$, $\\sigma = 0.3$ and $v_{\\rm flow} = 2200$km s$^{-1}$. Also, correlating the orbital phase with the modeled phase angle, it was possible to determine the orbital eccentricity as $e = 0.2$, similar to the value of 0.0 previously assumed {\\it ad hoc} by other authors. Although both  values lead to very  small differences in the orbital shape, its value is important for the determination of the stellar masses and orbital separation between the stars. Using this eccentricity and orbital inclination to model the radial velocity curve for the O-star, we found $K_{\\rm O} = 25$ km s$^{-1}$. If we assume $M_{\\rm O} = 40 - 60$ M$_{\\odot}$, we find $M_{\\rm WR} = 7.5 - 9.7$ M$_{\\odot}$, and the orbital major semi-axis $a_{\\rm O} \\simeq 5.4$ R$_{\\odot}$ and $a_{\\rm WR} \\simeq 30$ R$_{\\odot}$.  The model showed itself a powerful tool for constraining the wind and orbital parameters of massive binary systems.  The fits did not matched the data exactly at all epochs, but considering the difficulties of subtracting the emission excess from the original spectra, the general shape and peak position variations were well reproduced. We must state that this is a very simple approximation, taking into account the complexities found in such systems. Numerical simulations could give a more detailed analysis, and probably more accurate values for the model parameters in future works."
    },
    "0710/0710.0381_arXiv.txt": {
        "abstract": "The evolution of the galaxy stellar mass--star formation rate relationship ($M_*-$SFR) provides key constraints on the stellar mass assembly histories of galaxies.  For star-forming galaxies, $M_*-$SFR is observed to be fairly tight with a slope close to unity from $z\\sim 0\\rightarrow 2$, and it evolves downwards roughly independently of $M_*$.  Simulations of galaxy formation reproduce these trends, broadly independent of modeling details, owing to the generic dominance of smooth and steady cold accretion in these systems.  In contrast, the observed amplitude of the $M_*-$SFR relation evolves markedly differently than in models, indicating either that stellar mass assembly is poorly understood or that observations have been misinterpreted.  Stated in terms of a star formation activity parameter $\\asf\\equiv (M_*/SFR)/(t_{\\rm Hubble}-{\\rm 1 Gyr})$, models predict a constant $\\asf\\sim 1$ out to redshifts $z\\sim 4+$, while the observed $M_*-$SFR relation indicates that $\\asf$ increases by $\\sim\\times 3$ from $z\\sim 2$ until today.  The low $\\asf$ (i.e. rapid star formation) at high-$z$ not only conflicts with models, but is also difficult to reconcile with other observations of high-$z$ galaxies, such as the small scatter in $M_*-$SFR, the slow evolution of star forming galaxies at $z\\sim 2-4$, and the modest passive fractions in mass-selected samples.  Systematic biases could significantly affect measurements of $M_*$ and SFR, but detailed considerations suggest that none are obvious candidates to reconcile the discrepancy. A speculative solution is considered in which the stellar initial mass function (IMF) evolves towards more high-mass star formation at earlier epochs.  Following Larson (1998), a model is investigated in which the characteristic mass $\\hat{M}$ where the IMF turns over increases with redshift.  Population synthesis models are used to show that the observed and predicted $M_*-$SFR evolution may be brought into general agreement if $\\hat{M}=0.5(1+z)^{2} M_\\odot$ out to $z\\sim 2$. Such IMF evolution matches recent observations of cosmic stellar mass growth, and the resulting $z=0$ cumulative IMF is similar to the ``paunchy\" IMF favored by \\citet{far07} to reconcile the observed cosmic star formation history with present-day fossil light measures. ",
        "introduction": "How galaxies build up their stellar mass is a central question in galaxy formation.  Within the broadly successful cold dark matter scenario, gas is believed to accrete gravitationally into growing dark matter halos, and then radiative processes enable the gas to decouple from non-baryonic matter and eventually settle into a star-forming disk~\\citep[e.g.][]{mo98}. However, many complications arise when testing this scenario against observations, as the stellar assembly of galaxies involves a host of other processes including star formation, kinetic and thermal feedback from various sources, and merger-induced activity, all of which are poorly understood in comparison to halo assembly.  A key insight into gas accretion processes is the recent recognition of the importance of ``cold mode\" accretion~\\citep{kat03,bir03,ker05}, where gas infalling from the intergalactic medium (IGM) does not pass through an accretion shock on its way to forming stars.  Simulations and analytic models show that cold mode dominates global accretion at $z\\ga 2$~\\citep{ker05}, and dominates accretion in all halos with masses $\\la 10^{12}$M$_\\odot$~\\citep{bir03,ker05}.  The central features of cold mode accretion are that it is (1) {\\it rapid} (2) {\\it smooth}, and (3) {\\it steady}.  It is rapid because it is limited by the free-fall time and not the cooling time; it is smooth because most of the accretion occurs in small lumps and not through major mergers~\\citep{mur01,ker05,guo07}; and it is steady because it is governed by the gravitational potential of the slowly growing halo. A consequence of cold mode accretion is that the star formation rate is a fairly steady function of time~\\citep[e.g.][]{fin06,fin07}, which results in a tight relationship between the stellar mass $M_*$ and the star formation rate SFR.  Hence simulations of star-forming galaxies generically predict a tight relationship with $M_*\\propto$SFR that evolves slowly with redshift.  In detail, a slope slightly below unity occurs owing to the growth of hot halos around higher-mass galaxies that retards accretion.  The stellar mass--star formation rate ($M_*-$SFR) relation for star forming galaxies has now been observed out to $z\\sim 2$, thanks to improving multiwavelength surveys.  As expected from models, the relationship is seen to be fairly tight from $z\\sim 0-2$, with a slope just below unity and a scatter of $\\la 0.3$~dex~\\citep{noe07a,elb07,dad07}.  It also evolves downwards in amplitude to lower redshift in lock step fashion, suggesting a quiescent global quenching mechanism such as a lowering of the ambient cosmic density, again as expected in models.  The range of data used to quantify this evolution is impressive: \\citet{noe07a} used the AEGIS multi-wavelength survey in the Extended Groth Strip to quantify the $M_*$-SFR relation from $z\\sim 1\\rightarrow 0$; \\citet{elb07} used GOODS at $z\\sim 0.8-1.2$ and SDSS spectra at $z\\sim 0$; and \\citet{dad07} used star-forming BzK-selected galaxies with {\\it Spitzer}, X-ray, and radio follow-up to study the relation at $z\\sim 1.4-2.5$.  In each case, a careful accounting was done of both direct UV and re-radiated infrared photons to measure the total galaxy SFR, paying particular attention to contamination from active galactic nuclei (AGN) emission (discussed in \\S\\ref{sec:syst}).  Multiwavelength data covering the rest-optical were employed to accurately estimate $M_*$.  This broad agreement in $M_*-$SFR slope, scatter, and qualitative evolution between model predictions and these data lends support to the idea that cold mode accretion dominates in these galaxies.  Yet all is not well for theory.  Closer inspection reveals that the {\\it amplitude} of the $M_*-$SFR relation evolves with time in a way that is inconsistent with model expectations.  This disagreement is fairly generic, as shown in \\S\\ref{sec:comparison}, arising in both hydrodynamic simulations and semi-analytic models, and is fairly insensitive to feedback implementation.  The sense of the disagreement is that going to higher redshifts, the observed SFRs are higher, and/or stellar masses lower, than predicted in current models.  The purpose of this paper is to understand the implications of the observed $M_*-$SFR relation for our theoretical view of how galaxies accumulate stellar mass.  In \\S\\ref{sec:tact} it is argued that $M_*-$SFR amplitude evolution implies a typical galaxy star formation history that is difficult to reconcile with not only model expectations, but also other observations of high-$z$ galaxies.  Possible systematic effects that may bias the estimation of SFR and $M_*$ are considered in \\S\\ref{sec:syst}, and it is argued that none of them are obvious candidates to explain the discrepancy.  Finally a speculative avenue for reconciliation is considered, namely that the stellar initial mass function (IMF) has a characteristic mass that evolves with redshift (\\S\\ref{sec:imf}).  This model is constrained based on the $M_*$-SFR relation in \\S\\ref{sec:imfmodel}. The implications for such an evolving IMF are discussed in \\S\\ref{sec:mchar}.  Results are summarized in \\S\\ref{sec:summary}.  A $\\Lambda$CDM cosmology with $\\Omega=0.25$ and $H_0=70$~km/s/Mpc is assumed throughout for computing cosmic timescales. ",
        "conclusions": "\\label{sec:summary}  Implications of the observed stellar mass--star formation rate correlation are investigated in the context of current theories for stellar mass assembly.  The key point, found here and pointed out in previous studies~\\citep{dad07,elb07}, is that the amplitude of the $M_*-$SFR relation evolves much more rapidly since $z\\sim 2$ in observations that in current galaxy formation models.  It is shown here that this is true of both hydrodynamic simulations and semi-analytic models, and is broadly independent of feedback parameters.  In contrast, the slope and scatter of the observed $M_*-$SFR relation are in good agreement with models.  The tight $M_*-$SFR relation with a slope near unity predicted in models is a generic result owing to the dominance of cold mode accretion, particularly in early galaxies, which produces rapid, smooth and relatively steady infall.  The slow amplitude evolution arises because star formation starts at $z\\ga 6$ for moderately massive star-forming galaxies and continues at a fairly constant or mildly declining level to low-$z$.  The large discrepancy in amplitude evolution when compared to observations from $z\\sim 2\\rightarrow 0$ hints at some underlying problem either with the models or with the interpretation of data.  A convenient parameterization of the problem is through the star formation activity parameter, $\\asf\\equiv (M_*/{\\rm SFR})/(t_H-1\\;{\\rm Gyr})$, where $t_H$ is the Hubble time.  A low value corresponds to a starbursting system, a high value to a passive system, and a value near unity to system forming stars constantly for nearly a Hubble time. In models, $\\asf$ remains constant around unity from $z=0-4$, whereas in observations it rises steadily from $\\la 0.2$ at $z\\sim 2$ to close to unity at $z\\sim 0$.  Several ad hoc modifications to the theoretical picture of stellar mass assembly are considered in order to match $\\asf$ evolution, but each one is found to be in conflict with other observations of high-redshift galaxies.  Bursts seem unlikely given the low scatter in $M_*-$SFR. Delaying galaxy formation to match $\\asf$ results in such a low redshift for the onset of star formation, it is in conflict with observations of star forming galaxies at earlier epochs.  Hiding a large population of galaxies as passive runs into difficulty when compared to direct observations of the passive galaxy fraction in mass-selected samples at $z\\sim 2$.  Having an exponentially growing phase of star formation or ``staged\" galaxy formation are perhaps the most plausible solutions from an observational viewpoint, but it is difficult to understand theoretically how such scenarios can arise within hierarchical structure formation.  Hence if observed star-forming galaxies are quiescently forming stars and represent the majority of galaxies at $z\\sim 2$, as other observations seem to suggest, then it is not easy to accomodate the low $\\asf$ inferred from the $M_*-$SFR amplitude.  Systematic uncertainties in $M_*$ or SFR determinations could bias results progressively more at high-$z$ in order to mimic evolution in $\\asf$.  Various currently debated sources are considered, such as the contribution to near-IR light from TP-AGB stars, extinction corrections, assumptions about star formation history, AGN contamination, and PAH emission calibration.  It may be possible to concoct scenarios whereby several of these effects combine to mimic $\\asf$ evolution, but there are no suggestions from local observations that such scenarios are to be expected.  Hence if systematic effects are to explain the low $\\asf$ at high-$z$, it would imply significant and unexpected changes in tracers of star formation and stellar mass between now and high redshifts.  A solution is proposed that the stellar initial mass function becomes increasingly bottom-light to higher redshifts.  Several lines of arguments are presented that vaguely or circumstantially favor such evolution, though no smoking gun signatures are currently known.  A simple model of IMF evolution is constructed, based on the ansatz by \\citet{lar98} that the minimum temperature of molecular clouds is reflected in the characteristic mass of star formation ($\\hat{M}_{\\rm IMF}$) where the mass contribution per logarithmic mass bin is maximized.  The minimum temperature may increase with redshift owing to a hotter cosmic microwave background, more vigorous star formation activity, or lower metallicities within early galactic ISMs.  In order to reconcile the observed $\\asf$ evolution with theoretical expectations of no evolution, an evolving IMF of the form \\begin{eqnarray*} \\frac{dN}{d\\log{M}} &\\propto& M^{-0.3}\\;\\; {\\rm for}\\; M<\\hat{M}_{\\rm IMF},\\\\ &\\propto& M^{-1.3}\\;\\; {\\rm for}\\; M>\\hat{M}_{\\rm IMF};\\\\ \\hat{M}_{\\rm IMF}&=&0.5 (1+z)^{2} M_\\odot \\end{eqnarray*} is proposed.  The exponent of $\\hat{M}_{\\rm IMF}$ evolution is constrained by requiring no $\\asf$ evolution, through careful modeling with the PEGASE.2 population synthesis code.  While the exact form of IMF evolution is not well constrained by present observations, what is required is that the IMF has progressively more high-mass stars compared to low-mass at earlier epochs, and that this IMF applies to the majority of star forming galaxies at any epoch.  It is worth noting that this evolving IMF is only constrained out to $z\\sim 2$ from the $M_*-$SFR relation, though it yields predictions that are consistent with other observations out to $z\\sim 4$.  Extrapolating such evolution to higher redshifts is dangerous, since no observational constraints exist and the precise cause of IMF evolution is not understood.  Implications of such an evolving IMF are investigated.  By leaving the high-mass end of the IMF unchanged, recent successes in understanding the connections between high-mass star formation, feedback, and metal enrichment are broadly preserved.  The cosmic stellar mass accumulation rate would be altered compared to what is inferred from cosmic star formation history measurements using a standard IMF.  It is shown that an evolving IMF is at face value in better agreement with direct measures of cosmic stellar mass assembly~\\citep{per07,els07,wil07}. Furthermore, the evolving IMF goes towards relieving the generic tension between present-day fossil light measures versus observations of the cosmic star formation history.  In particular, the paunchy IMF favored by \\citet{far07} in order to reconcile the observed cosmic star formation history, present-day $K$-band luminosity density, and extragalactic background light constraints, is qualitatively similar to the cumulative IMF of all stars formed by today in the evolving IMF case.  Individually, each argument has sufficient uncertainties to cast doubt on whether a radical solution such as an evolving IMF is necessary.  But taken together, the $M_*-$SFR relation adds to a growing body of circumstantial evidence that the ratio of high-mass to low-mass stars formed is higher at earlier epochs.  It is by no means clear that IMF variations are the only viable solution to the $M_*-$SFR dilemma.  The claim here is only that an evolving IMF is an equally (un)likely solution as invoking unknown systematic effects or carefully crafted star formation histories in order to explain $M_*-$SFR evolution.  It is hoped that this work will spur further efforts, both observational and theoretical, to investigate this important issue.  Future plans include performing a more careful analysis of the evolving IMF in terms of observable quantities, in order to accurately quantify the impact of such IMF evolution on the interpretation of UV, near-IR, and mid-IR light.  Also, it is feasible to incorporate such an evolving IMF directly into simulation runs to properly account for gas recycling in such a scenario.  Finally, including some feedback mechanism to truncate star formation in massive systems may impact the $M_*-$SFR relation in some way, so this will be incorporated into the hydro simulations.  Observationally, pushing SFR and $M_*$ determinations to higher redshifts is key; a continued drop in $\\asf$ to $z\\sim 3$ would rapidly solidify the discrepancies with current models, and would also generate stronger conflicts with other observations of high-$z$ galaxies.  Assessing the AGN contribution and extinction uncertainties from high-$z$ systems is critical for accurately quantifying the light from high-mass star formation, for instance through the use of more direct star formation indicators such as Paschen-$\\alpha$.  Pushing observations further into the mid-IR such as to 70$\\mu$, past the PAH bands at $z\\sim 2$, would mitigate PAH calibration uncertainties in SFR estimates.  Obtaining a large sample of spectra for typical high-$z$ star-forming systems \\citep[like cB58;][]{pet00} would more accurately constrain the SED than broad-band data.  All of these programs push current technological capabilities to their limits and perhaps beyond, but are being planned as facilities continue their rapid improvement.  In summary, owing to the robust form of star formation histories in current galaxy formation models, the $M_*$-SFR relation represents a key test of our understanding of stellar mass assembly.  Current models reproduce the observed slope and scatter remarkably well, but broadly fail this test in terms of amplitude evolution.  Whether this reflects some fundamental lack of physical insight, or else some missing ingredient such as an evolving IMF, is an issue whose resolution will have a significant impact on our understanding of galaxy formation."
    },
    "0710/0710.5047_arXiv.txt": {
        "abstract": "{We present results of an analysis of the optical spectrum of the post-AGB star HD\\,56126 (IRAS\\,07134\\,+\\,1005) based on observations made with the echelle spectrographs of the 6-m telescope with spectral resolutions of R\\,=\\,25000 and 60000 at 4012--8790\\,\\AA. The profiles of strongest lines (HI; FeII, YII, BaII absorptions, etc.) formed in the expanding atmosphere at the base of the stellar wind have complex and variable shapes. To study the kinematics of the atmosphere, the velocities of individual features in these profiles must be measured. Differential line shifts of up to $\\rm V_r$\\,=\\,15$\\div$30\\,km/s have been detected from the lines of metals and molecular features. The star's atmosphere simultaneously contains both expanding layers and layers falling onto the star. A comparison of the data for different times demonstrates that both the radial velocity and the velocity pattern in whole are variable. The position of the molecular spectrum is stable, implying stability of the expansion velocity of the circumstellar envelope around HD\\,56126 detected in observations in the C$_2$ and NaI lines.}  \\authorrunning{Klochkova \\& Chentsov} \\titlerunning{HD\\,56126: structure of the atmosphere and envelope} ",
        "introduction": "We studied the kinematic parameters of the atmosphere of the star HD\\,56126 (SAO\\,96709), which belongs to the asymptotic giant branch (the post-AGB stage). In this short evolution stage, stars are in the process of their transition from the AGB to becoming a planetary nebula; for this reason, they are generally known as proto-planetary nebulae (PPNe). In the Hertzsprung--Russell diagram, post-AGB stars evolve toward the left from the AGB, maintaining nearly constant luminosity while becoming hotter. As descendants of AGB stars, these objects can be used to trace variations in the physical conditions and chemical parameters of the stellar material due to changes in the sources of energy release in the stellar interior, the ejection of the envelope, and mixing.  HD\\,56126 is the optical component of the IR source IRAS\\,07134+1005, which has a double peaked spectral energy distribution (SED), typical for PPNe. In addition to this anomalous SED, associated with the circumstellar dust envelope, the star also displays other characteristics of this type of object [\\cite{Kwok}]: the optical component of the PPN is a F5\\,Iab supergiant outside the galactic plane (b=+9$\\lefteqn{.}^m$99); the central star is surrounded by an extended nebula whose angular size exceeds 4$^{\\rm ''}$, according to Hubble Space Telescope observations [\\cite{Ueta}] (the largest known for this type of PPN); and the optical spectrum displays H$\\alpha$ emission and absorption with a variable line profile [\\cite{Oud1994}]. In addition to these general PPN characteristics noted by Kwok [\\cite{Kwok}], subsequent studies of HD\\,56126 and the associated IR source revealed several important peculiarities expected for this evolution stage: a large excess of carbon and $s$-process elements [\\cite{Kloch1995}] and an emission feature at $\\lambda$\\,=\\,21\\,$\\mu$ in the IR--spectrum.  In the small subgroup of PPN having this emission, the presence of this 21\\,$\\mu$ feature was found to correlate with observational manifestations of the products of nucleosynthesis (the excess of carbon and heavy metals in the outer layers of the atmosphere) [\\cite{Kloch1998, Decin}]. Thus, HD\\,56126 displays all the properties ex pected for a post-AGB object, indicating the importance of detailed studies of its optical spectrum with high spectral resolution in a broad wavelength range. Fortunately, HD\\,56126 is fairly bright (B\\,=\\,9$\\lefteqn{.}^m$11, V = 8.27m), making it the most convenient carbon enriched PPN star for high resolution spectroscopy.  The main moments of our study are spectral features identification and comparison the spectra of HD\\,56126 and the standard $\\alpha$\\,Per (Sp\\,=\\,F5\\,Iab); search for profile variability for spectral features; analyses the radial velocities $\\rm V_r$ to search for differential shifts; and studying the variability of the radial velocity. In Section~2, we briefly describe the techniques used for the observations, processing, and analysis of the spectral data. The main conclusions concerning the spectrum of the star are presented in Section~3.1, while Sections~3.2 and 3.3 present our radial velocity measurements derived using various features of the spectrum, and discuss the temporal behavior of the radial-velocity pattern. Our results are summarized in Section~4.   {\\footnotesize \\begin{table}[t] \\caption{Moments of observations and values of heliocentric radial velocitity $\\rm V_r$ measured. Column~4 contains $\\rm V_r$  values averaged over weak lines (with depths R$_\\lambda$ close to the continuum level, R$_\\lambda \\rightarrow$0). Velocities corresponding to the positions of the strongest components are presented for FeII(42), H$\\alpha$, and D2\\,NaI, with values determined from the weakest components given in parantheses. The two velocities in italics in column~5 are determined from the IR--oxygen triplet OI\\,7773\\,\\AA. Uncertain values are marked by a colon.} \\medskip \\begin{tabular}{ll c| c  c  c  l  l l @{\\quad} | l l l} \\hline \\small Date &\\small  Spectro-- &\\small  $\\Delta\\lambda$  & \\multicolumn{7}{c}{\\small  $V_r$} \\\\ \\cline{4-12} &\\small graph  &\\small  \\AA{} &  R$_\\lambda \\rightarrow $0 & \\small  Fe{\\sc ii} &\\small  H$\\beta$ & \\quad \\small  H$\\alpha$ & \\small  D\\,Na{\\sc i}& \\small C$_2$ &\\multicolumn{3}{c}{\\small  intestellar} \\\\ \\hline \\quad  1& 2& 3&4&5&6&\\quad7& 8&9& 10& 11 & 12\\\\ \\hline 12.01.93&Lynx&5560--8790& 88.8 &{\\it 91} &    &78 (100:)&77        &79: &     &     &    \\\\ 10.03.93&Lynx&5560--8790& 89.0 &{\\it 93} &    &71 (43:) &75:       &76: &     &     &     \\\\ 04.03.99&Lynx&5050--6640& 85.9 & 77      &    &76 (43:) &78        &77.1&     &     &     \\\\ 20.11.02&NES &4560--5995& 89.6 &95 (80:) &89  &         &74.9 (89) &77.2&12.0 &23.5 & 30.8\\\\ 21.02.03&NES &5150--6660& 88.8 & 96:     &    &88 (112:)&75.6 (89) &77.1&12   &24   & 31 \\\\ 12.04.03&NES &5270--6760& 88.4 &         &    &82 (103:)&75.4 (89:)&    &13   &23   & 30.5 \\\\ 14.11.03&NES &4518--6000& 85.3 &96 (86:) &97  &         &75.0 (87:)&76.9&12.5 &     &       \\\\ 10.01.04&NES &5270--6760& 86.7 &         &    &54: (78:)&75.6 (86:)&    &13.0 &23.5 & 31  \\\\ 09.03.04&NES &5275--6767& 89.8 &         &    &58 (74:) &76.1 (89) &    &13   &24   & 31   \\\\ 12.11.05&NES &4010--5460& 82.5 &97 (77:) &98  &         &          &77.5&     &     &      \\\\ \\hline \\end{tabular} \\end{table} } ",
        "conclusions": "Our high resolution spectra of HD\\,56126 have revealed variability of various line profiles: along with the previously known H$\\alpha$ profile variability, the profiles of strongest lines (such as BaII, YII, and FeII) also proved to be variable. The broad spectral range encompassed by our spectra enabled us to measure radial velocities using spectral features that form at various depths in the stellar atmosphere and circumstellar envelope.  We were able to distinguish the C$_2$ absorption molecular bands (and their structure), as well as bands identified with diffuse interstellar bands. We found substantial differential shifts in lines of different intensities within the same spectrum, i.e., appreciable $\\rm V_r$(R) dependences, as well as variation of these shifts with time. This indicates the need for velocity measurements based on a large set of lines in an extended spectral interval.  We conclude that both expanding layers and matter falling onto the star exist simultaneously in the stellar atmosphere. The position of the molecular spectrum is stable, indicating stability of the expansion velocity of the envelope."
    },
    "0710/0710.2451_arXiv.txt": {
        "abstract": "Recently the numerical simulations of the process of reionization of the universe at $z>6$ have made a qualitative leap forward, reaching sufficient sizes and dynamic range to determine the characteristic scales of this process. This allowed making the first realistic predictions for a variety of observational signatures. We discuss recent results from large-scale radiative transfer and structure formation simulations on the observability of high-redshift Ly-$\\alpha$ sources. We also briefly discuss the dependence of the characteristic scales and topology of the ionized and neutral patches on the reionization parameters. ",
        "introduction": "The observations of high-redshift QSO's \\citep{2001AJ....122.2833F,2001AJ....122.2850B} and large-scale CMB polarization \\citep{2007ApJS..170..377S} indicate that the intergalactic medium has been completely ionized by redshift $z\\sim6$ through an extended process. The most probable cause was the ionizing radiation of the First Stars and QSO's. Currently these are the two main direct observational constraints on this epoch. This scarcity of observational data is set to change dramatically in the next few years, however. A number of large observational projects are currently under way, e.g. observations at the redshifted 21-cm line of hydrogen \\citep[e.g.][]{1997ApJ...475..429M,2000ApJ...528..597T, 2002ApJ...572L.123I,2006MNRAS.372..679M,2006PhR...433..181F}, detection of small-scale CMB anisotropies due to the kinetic Sunyaev-Zel'dovich (kSZ) effect \\citep[e.g.][]{2000ApJ...529...12H,2003ApJ...598..756S,kSZ}, and surveys of high-redshift Ly-$\\alpha$ emitters and studies of the IGM absorption \\citep[e.g.][]{2003AJ....125.1006R,2004ApJ...604L..13S, 2006NewAR..50...94B}. The planning and success of these experiments relies critically upon understanding the large-scale geometry of reionization, i.e. the size- and spatial distribution of the ionized and neutral patches. This is best derived by large-scale simulations, although a number of semi-analytical models exist as well \\citep[e.g.][]{2004ApJ...613....1F}. Recently we presented the first large-scale, high-resolution radiative transfer simulations of cosmic reionization \\citep{2006MNRAS.369.1625I,% 2007MNRAS.376..534I} and applied those to derive a range of reionization observables \\citep{2006MNRAS.372..679M,kSZ,pol21,cmbpol,wmap3,2007arXiv0708.3846I}. Here we summarize recent results on the characteristic scales and topology of reionization and implications of our simulations for the observability of high-redshift Ly-$\\alpha$ sources. ",
        "conclusions": ""
    },
    "0710/0710.4936_arXiv.txt": {
        "abstract": "We interpret the recent gravitational lensing observations of Jee et al. \\cite{Jee} as first evidence for a {\\it caustic} ring of dark matter in a galaxy cluster.  A caustic ring unavoidably forms when a cold collisionless flow falls with net overall rotation in and out of a gravitational potential well.  Evidence for caustic rings of dark matter was previously found in the Milky Way and other isolated spiral galaxies.  We argue that galaxy clusters have at least one and possibly two or three caustic rings.  We calculate the column density profile of a caustic ring in a cluster and show that it is consistent with the observations of Jee et al. ",
        "introduction": "Using strong and weak gravitational lensing methods, Jee et al. \\cite{Jee} constructed a column density map of the central region of the galaxy cluster Cl 0024+1654.  The map shows a ring of dark matter of radius $\\simeq$ 400 kpc, width $\\sim$ 150 kpc and maximum column density $\\simeq 58 {M_\\odot \\over {\\rm pc}^2}$.  Remarkably, the overdensity in dark matter is not accompanied by an analogous structure in x-ray emitting gas or luminous matter.  In this regard, Jee et al. discovered a new instance where dark and ordinary matter have dramatically different spatial distributions.  The well-known observations of the ``Bullet Cluster\" 1E0657-56 \\cite{bul} provided an earlier example.  Jee et al. interpret the dark matter ring in Cl 0024+1654 as the product of a near head-on collision, along the line of sight, of two subclusters.  They performed a simulation of the response of the dark matter particles to the time-varying gravitational field and found that, after the collision has occurred, the dark matter particles move outward and form shell-like structures which appear as a ring when projected along the collision axis \\cite{Jee}.  The interpretation of Jee et al. fits with independent lines of evidence that Cl 0024+1654, an apparently relaxed cluster, is a collision of two subclusters.  However, it is shown in ref. \\cite{ZuH} that the observed dark matter ring is reproduced only for highly fine-tuned, and hence unlikely, initial velocity distributions.  The purpose of our paper is to propose an alternative interpretation, to wit that Jee et al. have observed a caustic ring formed by the in and out flow of dark matter particles falling onto the cluster for the first time.  We show that the observed ring is explained assuming only that the dark matter falling onto the cluster has net overall rotation, with angular momentum vector close to the line of sight, and velocity dispersion less than 60 km/s.  When cold collisionless dark matter falls from all directions into a smooth gravitational potential well, the phase space distribution of the dark matter particles is characterized everywhere by a set of discrete flows \\cite{Ips}.  The flows form outer and inner caustics. The outer caustics are formed by outflows where they turn around before falling back in.  Each outer caustic is a fold catastrophe ($A_2$) located on a topological sphere surrounding the potential well.  The inner caustics \\cite{crdm,sing} are formed near where the particles with the most angular momentum in a given inflow reach their closest approach to the center before going back out.  The catastrophe structure of the inner caustics \\cite{inner} depends on the angular momentum distribution of the infalling particles.  If that angular momentum distribution is characterized by net overall rotation, the inner caustics are rings (closed tubes) whose cross-section is a section of the elliptic umbilic catastrophe ($D_{-4}$) \\cite{sing}.  These statements are valid independently of any assumptions of symmetry, self-similarity, or anything else.  It is argued in ref.~\\cite{rob} that discrete flows and caustics are a generic and robust property of {\\it galactic} halos if the dark matter is collisionless and cold. The radii of the outer caustic spheres are predicted by the self-similar infall model \\cite{FG,B} of halo formation.  The radii of the inner caustic rings are predicted \\cite{crdm}, in terms of a single parameter $j_{\\rm max}$, after the model is generalized \\cite{STW} to allow angular momentum for the infalling particles.  Evidence for inner caustic rings distributed according to the predictions of the self-similar infall model has been found in the Milky Way \\cite{crdm,milk,mon} and in other isolated spiral galaxies \\cite{crdm,Kinn}.  The resolution of most present numerical simulations is inadequate to see discrete flows and caustics.  However such features are seen in dedicated simulations which increase the number of particles in the relevant regions of phase space \\cite{Stiff,simca}.  They should also become apparent in fully general simulations of structure formation through the use of special techniques \\cite{White}. ",
        "conclusions": "We conclude that the dark matter ring observed by Jee et al. has properties consistent with a caustic ring of dark matter. The column density agrees both in shape and overall amplitude. At present the data are too imprecise to infer with confidence the nature of the observed ring.  However, our proposal makes a distinct prediction for the column density profile across the ring, as illustrated in Fig.1. This may be tested by future observations.  In this regard, let us emphasize that the ${1 \\over x} I_a({2(x-a) \\over s})$ profile applies to each azimuth.  If the caustic ring interpretation is confirmed, an important corollary is that dark matter falls onto Cl 0024+1654 with net overall rotation.  ~~ We thank Igor Tkachev for making available his numerical codes to solve the equations of the self-similar infall model.  We thank Leanne Duffy for useful discussions. This work was supported in part by the U.S. Department of Energy under contract DE-FG02-97ER41029.  P.S. gratefully acknowledges the hospitality of the Aspen Center of Physics while working on this project."
    },
    "0710/0710.3727_arXiv.txt": {
        "abstract": "The energetics and emission mechanism of GRBs are not well understood. Here we demonstrate that the instantaneous peak flux or equivalent isotropic peak luminosity, $L_{iso}$ ergs s$^{-1}$, rather than the integrated fluence or equivalent isotropic energy, $E_{iso}$ ergs, underpins the known high-energy correlations. Using new spectral/temporal parameters calculated for 101 bursts with redshifts from {\\em BATSE}, {\\em BeppoSAX}, {\\em HETE-II} and {\\em Swift} we describe a parameter space which characterises the apparently diverse properties of the prompt emission. We show that a source frame characteristic-photon-energy/peak luminosity ratio, $K_{z}$, can be constructed which is constant within a factor of 2 for all bursts whatever their duration, spectrum, luminosity and the instrumentation used to detect them. The new parameterization embodies the Amati relation but indicates that some correlation between $E_{peak}$ and $E_{iso}$ follows as a direct mathematical inference from the Band function and that a simple transformation of $E_{iso}$ to $L_{iso}$ yields a universal high energy correlation for GRBs. The existence of $K_{z}$ indicates that the mechanism responsible for the prompt emission from all GRBs is probably predominantly thermal. ",
        "introduction": "The energetics of the central engine which powers the explosion responsible for a GRB are both intriguing and fundamental to our understanding of these cosmic events. The isotropic energy outflow at source, estimated using the integrated gamma-ray fluence, is enormous, up to $E_{iso}\\sim10^{54}$ ergs, and even if the outflow is collimated in jets the total energy involved is still huge, $E_{\\gamma}\\sim10^{51}$ ergs. The possibility that the explosion taps a standard energy resevoir has been pursued by many authors following the initial suggestion from Frail et al. (2001). If this total energy available were, indeed, roughly constant (or predictable through other means) and we could reliably estimate the collimation, then GRBs could be used as a cosmological probe to very high redshifts, Bloom et al. (2003), Ghirlanda et al. (2004).  Early on it was noted that, based on analysis of {\\em BATSE} data, there was a correlation between $E_{p}$, the peak of $E.F(E)$ where $F(E)$ ergs cm$^{-2}$ keV$^{-1}$ is the observed spectrum, and the fluence (Mallozzi et al. 1995, Lloyd et al. 2000). When redshifts became available for long bursts the isotropic energy, $E_{iso}$, could be estimated from the fluence and the peak energy could be transformed into the source frame, $E_{pz}$, the so-called Amati relation, a correlation between $E_{iso}$ and $E_{pz}$ in the sense that more energetic bursts have a higher $E_{pz}$, was discovered using data from {\\em BeppoSAX}, (Amati et al. 2002). This correlation has subsequently been confirmed and extended although there remain many significant outliers, including all short bursts. The physical origin of the correlation may be associated with the emission mechanisms operating in the fireball but the theoretical details are far from settled (see the discussion by Amati (2006) and references therein). More recently a tighter correlation between $E_{iso}$, $E_{pz}$ and the jet break time, $t_{break}$, measured in the optical afterglow has been reported (Ghirlanda et al. 2004). This is explained in terms of a modification to the Amati relation in which $E_{iso}$ is corrected to a true collimated energy, $E_{\\gamma}$, using an estimate of the collimation angle derived from $t_{break}$. The details of the collimation correction depend on the density and density profile of the circumburst medium, Nava et al. (2006) and references therein. Multivariable regression analysis was performed by Liang \\& Zhang (2005) to derive a model-independent relationship, $E_{iso}\\propto E_{pz}^{1.94}t_{zbreak}^{-1.24}$, indicating that the rest-frame break time of the optical afterglow, $t_{zbreak}$ was indeed correlated with the prompt emission parameters.  Other studies have concentrated on the properties of the isotropic peak (maximum) luminosity, $L_{iso}$ ergs s$^{-1}$, measured over some short time scale $\\approx1$ s, rather than the time integrated isotropic energy, $E_{iso}$. Yonetoku et al. (2004) noted a correlation between $L_{iso}$ and $E_{pz}$ for 16 GRBs with firm redshifts. A correlation between $L_{iso}$ and the spectral lag was first identified by Norris et al. (2000) and explained in terms of the evolution of $E_{peak}$ with time. The shocked material responsible for the gamma-ray emission is expected to cool at a rate proportional to the gamma-ray luminosity and it has been suggested that $E_{peak}$ traces the cooling (Schaefer 2004). A similar correlation between $L_{iso}$ and the variability of the GRB ($V$) was described by Reichart et al. (2001). The origin of the $L_{iso}-V$ relation is likely to be related to the physics of the relativistic shocks and the bulk Lorentz factor of the outflow. It could be that high $\\Gamma_{outflow}$ results in high $L_{iso}$ and $V$ while lower luminosity and variability are expected if $\\Gamma_{outflow}$ is low (see, for example, M\\'{e}sz\\'{a}ros et al. 2002). A rather bizzare correlation involving $L_{iso}$, $E_{pz}$ and variability was found by Firmani et al. (2006). They employed the ``high signal'' time, $T_{45}$, as formulated by Reichart et al. (2001) in their study of variability, and showed that $L_{iso}\\propto E_{pz}^{1.62}T_{45}^{-0.49}$ for 19 GRBs with a spread much narrower than that of the Amati relation. There is currently no explanation for such a correlation although it may be connected with the spectral lag and variability correlations and the Amati relation.  The correlation between $E_{iso}$ and $E_{pz}$ supplemented by additional empirical information can be used in pseudo redshift indicators, for example Atteia (2003), Pelangeon \\& Atteia (2006), but the intrinsic spread in the correlation and uncertainty about the underlying physical interpretation introduce errors, typically of a factor $\\sim2$. It may be possible to reduce the errors by simultaneous application of several independent luminosity/energy correlations, and extension of the Hubble Diagram to high redshifts using GRBs has been attempted, see for example Schaefer (2007). However, it is not clear that the correlations briefly described above are truly independent and there may be some underlying principle or mechanism which connects them all together. Recently, and more controversially, Butler et al. (2007) have raised serious doubts about the validity of these correlations suggesting that it is likely that they are introduced by observational/instrumental bias and have nothing to do with the physical properties of the GRBs and hence they conclude that GRBs are probably useless as cosmological probes. Here we take a new look at the source frame spectral and temporal properties of a large number of GRBs for which we have redshifts in order to try and understand what really correlates with what and whether or not this can provide useful intrinsic information about the GRBs and what drives them. In this analysis we include the short-duration GRBs which may share a similar emission mechanism with long bursts despite probably having different progenitors. ",
        "conclusions": "The equivalent isotropic energy, $E_{iso}$ ergs, of a GRB can be expressed as the product of two source frame terms, a characteristic photon energy, $E_{wz}$ keV, calculated from the shape of the spectrum across the range 1-10000 keV and the energy density at the peak of the $E.F_{z}(E)$ spectrum, $Q_{pz}$ ergs keV$^{-1}$. The correlation trend between $E_{wz}$ and $Q_{pz}$ gives rise to the Amati relation. By stacking the samples of a GRB light curve into descending order we can construct a rate profile. The functional form of such rate profiles is common to the vast majority of bursts. Fitting the profile gives us a luminosity time, $T_{Lz}$ s, a measure of the burst duration which can be used to convert the energy density at the peak to a luminosity density at peak, $Q_{pz}/T_{Lz}$ ergs keV$^{-1}$ s$^{-1}$. We can calculate the peak equivalent isotropic luminosity as a product $L_{iso}=E_{wz}Q_{pz}/T_{Lz}=E_{iso}/T_{Lz}$ ergs s$^{-1}$.  $E_{wz}$ is a characteristic photon energy or a measure of the colour or hardness of the burst and $Q_{pz}/T_{Lz}$ is a measure of the instantaneous peak brightness. We have gathered and analysed sufficient spectral and temporal data from 101 bursts to produce the relation between $E_{wz}$ vs. $Q_{pz}/T_{Lz}$ and $E_{wz}$ vs. $L_{iso}$, shown in Figure \\ref{fig11}, which constitutes the closest thing we have to an intrinsic colour-magnitude diagram for the peak emission from GRBs, $E_{wz}\\propto L_{iso}^{0.25}$. All bursts are clustered such that we can construct a intrinsic colour-magnitude quasi constant $K_{z}$, which is a function of the source frame characteristic photon energy/peak luminosity ratio given by Equation \\ref{eq15}. The range of equivalent isotropic energy that drives the expanding fireball is very large, 6 orders of magnitude (Figure \\ref{fig3}), but the instantaneous hardness/brightness of the peak emission covers a very small intrinsic dynamic range, $\\approx4$.  The existence and form of $K_{z}$ indicates that the physical mechanism for the Gamma-ray production at the photosphere of the fireball is common to all bursts and is probably thermal although many other possibilities are not ruled out. If the prompt spectra are dominated by thermal photons the scatter in $K_{z}$ may be attributed to variations in the size and/or Lorentz factor of the fireball. XRFs have low $\\Gamma_{0}$ and/or large radii. Short bursts have high $\\Gamma_{0}$ and/or small radii. The relation between $T_{Lz}$ vs. $Q_{pz}$ clearly separates short from long, but both classes have the same instantaneous peak hardness/brightness."
    },
    "0710/0710.5667_arXiv.txt": {
        "abstract": "We present a review of the standard paradigm for giant planet formation, the core accretion theory. After an overview of the basic concepts of this model, results of the original implementation are discussed. Then, recent improvements and extensions, like the inclusion of planetary migration and the resulting effects are discussed. It is shown that these improvement solve the ``timescale problem''. Finally, it is shown that by means of generating synthetic populations of (extrasolar) planets, core accretion models are able to reproduce in a statistically significant way the actually observed planetary population. ",
        "introduction": "Our current understanding of planet formation is based on several centuries of observations of the planets of our own Solar System, 12 years of extrasolar planets detection,  and several decades of observations of young stellar systems. These studies have let to the general concept that after the collapse of a dense gas cloud, a protostar surrounded by a protoplanetary disk was formed. In this disk, solids started to coagulate from fine dust and grew further by mutual collision to form planetesimals (provided the bottleneck by bodies roughly one meter in size can be overcome), then protoplanets, and ultimately the actual planets. Some of the protoplanets managed to accrete a massive gaseous envelope onto their core. This is the very rough outline of the core accretion model. ",
        "conclusions": "The core accretion paradigm explains in an unified way the formation of giant and terrestrial planets, so that there is no need for a special mechanism for giant planets.  Since the first core accretion models, significant improvement and extensions were made. Such improved and extended core accretion models can form giant planets well within observed disk lifetimes, so that there is no need for a faster formation mechanism.  They have also reached a degree of maturity that allows quantitative tests with observations, of both the giant planets of our own Solar System, and, by means of population synthesis, of the extrasolar planet population. In the latter case, the whole population of detected planets can be used to constrain the models, which excludes model fine tuning for a specific case, and fully exploits the observational investment. As shown by statistical tests,  extended core accretion models can reproduce many observed properties and correlations in the extrasolar planet population in a quantitative significant way with one synthetic population at one time. This means that accretion models can now be used to predict future observations, so that theory can feed back on the design of future instruments"
    },
    "0710/0710.2047_arXiv.txt": {
        "abstract": "We review current state of neutron star cooling theory and discuss the prospects to constrain the equation of state, neutrino emission and superfluid properties of neutron star cores by comparing the cooling theory with observations of thermal radiation from isolated neutron stars. ",
        "introduction": "The equation of state (EOS) of superdense matter in neutron star cores is still a mystery. It is not clear if it is soft, moderate or stiff; if the matter contains nucleons/hyperons, or exotic components. In the absence of good practical theory of supranuclear matter the problem cannot be solved on purely theoretical basis, but it can be solved by comparing theoretical models with observations of neutron stars. The attempts to solve this long-standing problem by different methods are numerous (e.g., Refs.\\ \\cite{nsb1,lp07}). Here we discuss current results obtained from studies of cooling isolated neutron stars.  The first papers on neutron star cooling appeared with an advent of X-ray astronomy, before the discovery of neutron stars. Their authors tried to prove that neutron stars cool not too fast and can be discovered as sources of thermal surface X-ray radiation. The first estimates of thermal emission from cooling neutron stars were most probably done by Stabler \\cite{stabler60} in 1960. Four years later Chiu \\cite{chiu64} made similar estimates and analyzed the possibility to discover neutron stars from their thermal emission. First, simplified calculations of neutron star cooling were done in 1964 and 1965 \\cite{morton64, cs64, bw65b}. The foundation of the strict cooling theory was laid in 1966 by Tsuruta and Cameron \\cite{tc66}, one year before the discovery of pulsars. We review the current state of the cooling theory. More details can be found in recent review papers \\cite{yp04,pgw06}. ",
        "conclusions": "The theory of cooling neutron stars of ages $10^2-10^6$ yr mostly tests the neutrino emission properties of the neutron star core. Its main results are as follows.  (1) Neutrino emission in the outer core (i.e., in the core of a low-mass star) is a factor of 30--100 lower than the modified Urca emission in a nonsuperfluid star.  (2) Neutrino emission in the inner core (of a massive star) is at least a factor of 30--100 higher than the modified Urca emission. It can be enhanced by direct Urca process in nucleon/hyperon inner core or by the presence of pion or kaon condensate, or quark matter.  (3) The scenario with open direct Urca process predicts the existence (Fig.\\ \\ref{exotica}) of massive isolated neutron stars which are much colder than those observed now. In the scenario with pion condensate, the massive stars should be warmer (than those with open direct Urca) but colder than the observed ones. In the scenario with kaon condensate the massive stars should be even warmer but slightly colder than the observed sources. A discovery of cold cooling neutron stars would be crucial to constrain the level of enhanced neutrino emission in the inner core.  (4) Observations of cooling neutron stars can be analyzed together with observations of SXTs in quiescent states. The data on SXTs indicate the existence of very cold neutron stars (first of all, SAX J1808.4--3658) which cool via direct Urca process, but the data and interpretation require additional confirmation.  (5) A transition from slow neutrino emission in the outer core to enhanced emission in the inner core has to be smooth. Current observations of cooling neutron stars and SXTs do not constrain the parameters of this transition. A firm measurement of masses of cooling or accreting stars would help to impose such constraints.  (6) New observations and reliable practical theories of dense matter are vitally important to tune the cooling theory as an instrument for exploring physical properties of neutron star interiors and neutron star parameters. The tuning will imply a careful analysis of many cooling regulators.   \\begin{theacknowledgments} This work was partially supported  by the Russian Foundation for Basic Research (grants 05-02-16245 and 05-02-22003), by FASI-Rosnauka (grant NSh 9879.2006.2), and by the Joint Institute for Nuclear Astrophysics (grant NSF PHY 0216783). \\end{theacknowledgments}"
    },
    "0710/0710.5498_arXiv.txt": {
        "abstract": "We have observed nearly 200~FGK stars at 24~and 70~microns with the {\\em Spitzer} Space Telescope. We identify excess infrared emission, including a number of cases where the observed flux is more than 10~times brighter than the predicted photospheric flux, and interpret these signatures as evidence of debris disks in those systems. We combine this sample of FGK stars with similar published results to produce a sample of more than 350~main sequence AFGKM stars. The incidence of debris disks is 4.2$^{+2.0}_{-1.1}$\\% at 24~microns for a sample of 213~Sun-like (FG) stars and 16.4$^{+2.8}_{-2.9}$\\% at 70~microns for 225~Sun-like (FG) stars. We find that the excess rates for A, F, G, and K stars are statistically indistinguishable, but with a suggestion of decreasing excess rate toward the later spectral types; this may be an age effect. The lack of strong trend among FGK stars of comparable ages is surprising, given the factor of~50 change in stellar luminosity across this spectral range. We also find % that the incidence of debris disks declines very slowly beyond ages of 1~billion years. ",
        "introduction": "Planetary system formation must be studied through various indirect means due to the relative faintness of distant planets and the long timescales over which planetary system evolution takes place. One powerful technique has advanced substantially in the era of the {\\em Spitzer} Space Telescope: investigations of dusty debris disks around mature, main sequence stars. Debris disks arise from populations of planetesimals that remain from the era of planet formation; the analogs in our Solar System are the asteroid belt and the Kuiper Belt. Any system that possesses a debris disk necessarily has progressed toward forming a planetary system to some degree.  The many small bodies that inhabit a debris disk can, on occasion, collide, producing a shower of fragments that grind each other down to dust particles. These dust grains can be heated by the central star to temperatures $\\sim$100~K, where they can be detected at wavelengths of 10--100~microns. Data from the IRAS and ISO satellites were used to identify and characterize debris disks \\citep[e.g.,][]{aumann84,decin00,spangler01,habing01,decin}. The Multiband Imaging Photometer for {\\em Spitzer} (MIPS; \\citet{mips}) offers substantially improved sensitivities at 24~and 70~microns and therefore can be used to advance the study of debris disks and measure the fraction of stars that possess colliding swarms of remnant planetesimals.  A number of these surveys have been carried out using MIPS. \\citet{astars} and \\citet{astars2} observed hundreds of A~stars and found that the number of A~stars showing thermal infrared excess suggestive of collisionally-produced dust decreases as a function of stellar age, from 50\\% or more at ages just past the gas dissipation age of 10~Myr to 30\\% at 500~Myr. The excess rates are lower for older and lower mass stars. \\citet{bryden} found that the excess rate is around 15\\% for stars of mass and age similar to our Sun (for a sample of 69~stars). \\citet{gautier} did not detect any 24~or 70~micron excesses suggestive of debris disks in a sample of 60~field (old) M~stars (with only 13~strongly detected at 70~microns). Finally, multiplicity appears to play a significant role in modulating debris disks, as \\citet{binaries} found that the debris disk rate for A~and F~binaries was higher than that for single stars of similar spectral types.  We present here a survey for excesses around almost 200~F, G, and K~stars, with ages and masses similar to that of our Sun. (Some of these data were published in \\citet{bryden}.) We characterize the excesses that we find and present a few systems of particular interest. We create a larger sample of FGK~stars by adding 75~stars from a similar survey, and derive excess rates as a function of spectral type and as a function of age. Finally, we discuss the implications of our results for planetary system formation. ",
        "conclusions": "\\subsection{Metallicity and excesses}  Four stars in our FGK sample do not have published metallicities. For the 189~stars with known metallicities in Table~\\ref{targetinfo}, the mean metallicity is -0.08$\\pm$0.22, with a median of~-0.06. The mean metallicity of stars with excesses is -0.11$\\pm$0.19 (median is~-0.09). The mean metallicity of stars with no excesses is -0.08$\\pm$0.22 (median is~-0.05). There is no difference between the metallicity of the population of stars with excesses and the population of stars with no excesses. This lack of correlation has been discussed previously (e.g., B06, \\citet{chastpf}).  \\subsection{Excess rates across spectral type \\label{spectral}}  Many parameters affecting infrared excesses change with spectral type, including the importance of grain loss mechanisms (winds, Poynting-Robertson drag, photon pressure); stellar luminosity; and the locations of key temperatures (e.g., the ice line) in the systems. Significant surveys for debris disks across spectral types A--M~have now been published, and we use those data to look for systematic trends across spectral types.  It is well known that excess rates decrease with stellar age \\citep{habing01,spangler01,astars,astars2,siegler}. This dependence must be avoided in testing for changes with spectral type. We take the oldest ($\\ge$600~Myr) A~stars from the \\citet{astars2} sample as our representative sample from that spectral type. For our F, G, and K~samples we take the union of the data presented here and the data in \\citet{chastpf}. The targets and data reduction presented in \\citet{chastpf} are quite similar to the selections and techniques we have employed here, which allows us to merge the two samples relatively seamlessly. We calculate the excess ratios for these F, G, and K samples and take the M~stars excess rates from \\citet{gautier}. This compilation is presented in Table~\\ref{excesssum} and Figure~\\ref{spexcess}.  Within the error bars, the excess rates for the A, F, G, and K~subsamples are essentially indistinguishable. However, there is a suggestion of a trend of decreasing 70~\\micron\\ excess rates with later spectral types (Figure~\\ref{spexcess}). We note, however, that the mean age for the populations increases with later spectral types. It is possible that we are instead detecting a time-related effect, although the decay timescales identified by \\citet{astars}, \\citet{astars2}, and \\citet{siegler} of hundreds of millions of years should long since have diminished all disks at ages of billions of years. We discuss this possibility in Section~\\ref{ages}.  \\citet{chastpf} remarked that the excess rates for K~stars appeared to be lower than that for F and G stars, finding zero excesses among 23~stars later than K2 (and zero excesses among a larger combined sample of 61~K1--M6~stars). We include the \\citet{chastpf} data in our analysis here, and find that, formally, the excess rate for K~stars is not significantly different than the excess rates for earlier (F and G) stars. As \\citet{chastpf} note, and we confirm, none of the 6~K~stars with excesses in our larger sample are later than K2.   Part of the motivation for assembling the F~stars program (PID~30211) was as a control sample for the binary star program presented in \\citet{binaries}. In that study, 69~A3--F8 binary star systems were found, overall, to have relatively high excess rates: 9\\% at 24~microns and 40\\% at 70~microns. It is clear from our results here (see Figure~\\ref{spexcess}) that the (single) F~stars excess rate is equal to or lower than the (single) A~stars excess rate. The binaries excess rates remain significantly high compared to the control sample of A~and F~stars.  Figure~\\ref{fd} shows that, in general, there is no trend of dust distance as a function of stellar effective temperature (spectral type). This may suggest that the processes that drive planetesimal formation do not depend strongly on a single critical temperature, as would be the case in the ``ice line'' model, where protoplanetary disk surface densities increase across certain temperature boundaries. However, our method for calculating dust distances may be too crude to see this effect. We note that all of the (minimum) dust distances given in Table~\\ref{excesstable} would fall within the planetary realm of our Solar System ($<$30~AU). We are not observing disks that are far outside of the potential planetary realm of these systems.  We stated above that there is no particular trend for either fractional luminosity or dust distance with spectral type, but there is an important caveat: no disks with large dust distances or relatively small fractional luminosities were identified among the latest stars in our sample. This dearth may simply allude to the fact that later stars are cooler. Dust at 25~AU around a K1 (5000~K) star would have a temperature around 50~K; this dust would have its peak emission near 70~microns, and cooler (more distant) dust would have its peak emission longward. MIPS 70~micron observations of such a dust population, or a cooler one, would not readily show the presence of this excess. The lack of distant disks around K stars may therefore be an observational bias. Similarly, low fractional luminosity disks would be more difficult to detect around K~stars than around earlier stars, so the lack of low fractional luminosity disks for later stars may also be due to observational bias. These observational biases may be corrected with sufficiently sensitive measurements at $\\sim$100~microns (e.g., with Herschel).  To further address the question of whether there is any detectable trend of excess rate as a function of spectral type using existing published data, we compared the incidence of 70~$\\mu$m excesses across these 5~spectral types (A, F, G, K, M), a total sample size of more than 350~stars. To avoid uncontrolled selection effects, we confined the comparison to stars that would have been detected at 70~microns at a level of at least 2:1 on the photosphere. We then applied a number of tests. First, we computed weighted average values of R70 (observed flux over predicted flux) as in \\citet{gautier}. We find a value of $\\sim$5 for the A stars, but we discount this large average because of the small size of the sample (27~stars). The values for the F, G, K, and M stars are 1.16, 1.23, 1.06, and 1.025, respectively. Errors are difficult to estimate because the excesses are not normally distributed. We also computed straight averages of R70 for these same samples. Here, we calculate 4, 2.6, 1.8, 1.4, and 1.1 for the A, F, G, K, and M stars, respectively. Again, the values need to be interpreted with caution because error estimation is difficult. To circumvent the difficulties in determining errors, we binned the excesses into intervals of 0.5~in excess ratio and used the K-S test to determine if the resulting distributions were likely to have been drawn from the identical parent distribution. For each stellar type, we tested the relevant distribution against the distribution for all the stars, excluding the contribution of the spectral type in question. The result was a set of probabilities of 0.03, 0.8, 0.3, 0.3, and 0.01 that the types A, F, G, K, and M, respectively, were drawn from the same parent distribution as the other types. (For this test, a claim of a significant difference requires a probability of 0.05 or less that the samples are from the same distribution. Values of~0.3 imply that the the G~and K~stars are 1$\\sigma$~different from their respective control samples.) From this suite of tests, we conclude that the incidence of excesses is different for old A stars and for M stars from that of the rest of our sample. We further deduce that there is a possibility of a difference appearing in the K~stars from their lower average excess ratio. The distributions for F and G stars appear to be indistinguishable with our data. We employ this conclusion in the creation of an FG ``supersample'' (Section~\\ref{supersample}).  The higher excesses indicated for the old A stars could be an age effect, since the sample is by necessity significantly younger than the later types (which we selected in general to be $>$ 1 Gyr in age). In fact, \\citet{gorlova} and \\citet{siegler} compare 24$\\mu$m excesses from young A and solar-like stars and find that the incidence is quite similar at a given age. Age effects would presumably cause an apparent decrease of excess incidence with later type among the F, G, and K stars because F stars will evolve off the main sequence quickly enough to bias our sample toward younger objects (Section~\\ref{ages}).  In the end, this discussion may still be suffering from a relatively small number of K~stars sampled. An ongoing {\\em Spitzer} program (PID~30490) to survey nearby stars that were not observed in other programs --- a sample that includes $\\sim$400~K~stars --- should help unravel these statistics. Nonetheless, given our result, it is unlikely that future surveys will find a strong trend among F, G, and K stars of comparable ages --- a range of spectral types that spans more than a factor of~50 in stellar luminosity. This behavior is counter to our expectations and is a challenge to models of debris disk evolution.  \\subsection{Effects of age \\label{ages}}  The lack of strong dependence on spectral type lets us combine data on various types to study the evolution with stellar age. Figure~\\ref{fgkage} shows individual R24 and $\\chi_{70}$ determinations for the stars in our FGK~sample whose ages are known, as a function of system age. There is no correlation apparent for these individual sources, so we look to binned data in a larger combined sample for evidence of trends.  We once again take the union of the data presented here and that presented in \\citet{chastpf}, but this time we include only spectral types F0--K5 from the \\citet{chastpf} sample (that is, we exclude the latest K~stars). This is because the data we present in this paper covers the range F0--K5, and we want the best match to our combined sample. There are 10~stars in the F0--K5 TPF/SIM subsample that have no known ages, and 10~additional stars with ages less than 1~billion years. Of these~20, 2~systems have excesses (10\\%). Since this excess ratio is not significantly different from that of the overall FGK sample, and since the PID~30211 F~stars sample is controlled for age but the PID~41 sample is not specifically controlled for age, we make no attempt to correct the larger sample for age.  Figure~\\ref{ageexcess} shows excess rate as a function of age for this combined sample. To zeroth order, there is no trend as a function of age: a constant excess rate of $\\sim$20\\% adequately fits the data, being consistent at 1$\\sigma$ with the 10~Gyr data point and at 1.5$\\sigma$ with the 8~Gyr data point. Furthermore, the 10~Gyr bin has only 7~targets in it, and the 8~Gyr bin has only 33~targets in it (still a relatively small number).  On the other hand, we note that the data shown in Figure~\\ref{ageexcess} is suggestive of an excess rate that decreases with time through at least 8~Gyr. In this scenario, the 10~Gyr bin would be highly anomalous, although we note that this bin is clearly affected by small number statistics (2~excesses out of 7~stars). In other words, there may be a real trend and a real evolution of planetary systems and debris disks even on the billion-year timescale. However, this may again be a manifestation of the (potential) observational bias shown in Figure~\\ref{fd}, as follows.  The fraction of stars in a given age bin that are K~stars increases for the later age bins. If K~stars truly have fewer excesses (or fewer detectable excesses, according to observational biases) than other spectral types, Figures~\\ref{ageexcess} and~\\ref{agesfd} might indeed be showing a real decrease with increasing age, but caused not by a long-timescale evolution of planetary systems but by the increasing dominance of excess-deficient K~stars at the oldest ages. The data present in our larger sample cannot distinguish between the competing possibilities of K~stars preferentially lacking (detectable) disks, or of old stars increasingly lacking disks. It also remains to be seen whether the high excess rate for the oldest bin in Figure~\\ref{ageexcess} is anything more than a small number statistics anomaly.   The overall high rate of excess incidence in our samples (17\\%) indicates that the 400~Myr decay timescale the drives the evolution of A~star debris disks \\citep{astars2} cannot drive the evolution of debris disks in our $>$1~Gyr Sun-like sample. Instead, Sun-like stars appear to have a relatively constant incidence of 15\\%--20\\% that is not strongly dependent on age, but may be weakly dependent on age through a very long timescale decrease.  \\subsection{Debris disks around Sun-like stars \\label{supersample}}  Because there is no difference in excess rate between F and G stars, we can combine our sample of F0--G9~stars into a single population of ``Sun-like'' stars. To this sample of 169~stars % we add the 56~F~and G~stars from \\citet{chastpf} to create an even larger sample of 225~Sun-like stars (213~stars at 24~microns). We refer to this merged sample of 213~and 225~stars as the ``Sun-like supersample.'' The excess rates for this supersample are 4.2$^{+2.0}_{-1.1}$\\% at 24~microns % and 16.4$^{+2.8}_{-2.9}$\\% at 70~microns % (Table~\\ref{excesssum}). With this large supersample, we can now state the debris disk incidence rate for Sun-like stars with quite good confidence (small error bars).  \\subsection{Implications for planetary system formation}   The majority of the debris disk systems that we present here have excesses at 70~microns only, suggesting temperatures $\\lesssim$100~K and therefore an inner edge to the disks. These dusty debris disks are likely produced by collisions within a swarm of planetesimals akin to the asteroid belt or Kuiper Belt in our Solar System. We interpret the cool temperatures we derive for the dust in these disks as evidence of inner disk holes where the surface density of dust is much smaller than in the planetesimal ring, and potentially zero. This architecture is strongly reminiscent of our Solar System, where planets sculpt the edges of the planetesimal and dust belts. It may be that many of the systems we discuss here similarly have planets sculpting their dust distributions. In several cases, the known planets may indeed be the ones sculpting the inner edges of the disks (Tables~\\ref{colortemps} and~\\ref{excesstable}). Additionally, Lawler et al.\\ (in prep.) have found, using IRS spectra, that the incidence of detectable levels of warm dust may be higher than the 4.2\\% we find here, indicating that dust in a region analogous to our Solar System's asteroid belt may also be somewhat common.  \\citet{bryden07} show that the properties of disks around planet-bearing stars are unlikely to be similar to the properties of disks around stars without known planets. Briefly, the detection rates between these two populations are similar, but the planet-bearing stars are generally farther away and in more confused regions of the sky. They conclude that disks around planet-bearing stars are dustier (i.e., more massive), and probably more common, than disks around stars without known planets. Additionally, \\citet{binaries} recently showed that the excess rate for binary A-F~stars is 9\\% and 40\\% at 24~and 70~\\micron, respectively, significantly higher than our results for FGK stars and for Sun-like stars. We see in Figure~\\ref{spexcess} that these excess rates are relatively high, compared to the large sample of single stars we present here. Combining these two studies, we find strong evidence that the presence of additional massive bodies (whether stellar or planetary companions) in stellar systems appears to promote higher excess rates. This effect may simply be dynamical --- more massive bodies means more stirring, more collisions, and more dust --- but it is not clear that such a process could remain effective for billions of years. Further work is necessary to explain these results.   Figure~\\ref{spexcess} shows that the transition from the high excess rate around A~stars to the more modest excess rate around Sun-like stars is gradual. This gradual decay is likely an age effect (consider the mean ages of the samples given in Figure~\\ref{spexcess}). Our data show that excess rates for FGK stars decline slowly with stellar age, and it is not clear whether we are detecting a billion-year tail of debris disk evolution, a dependence on spectral type, and/or an observational bias.  Young ($<$1~Gyr) debris disks decay on 100--400~million year timescales \\citep{astars,astars2,gorlova,siegler}. The presence of excesses around 16\\% of Sun-like stars at billion year ages indicates that these old debris disks must be driven by a different evolutionary process than debris disks around those younger stars. One interpretation, using our Solar System as an analogy, is that the younger systems are still active in a Late Heavy Bombardment kind of dynamical upheaval. After a billion years, such large scale processes likely have ceased in all but the most unusual systems. Debris disks are then produced from collisions within remaining planetesimal belts (e.g., our asteroid belt) that have been dynamically excited by the previous eon's dynamical stirring. However, there is no trend of fractional luminosity with age (Figure~\\ref{agesfd}), although there is a lack of high fractional luminosity disks at old ages.  The presence of a debris disk indicates that planetary system formation progressed at least to the planetesimal stage in a given system. There is no apparent dependence on spectral type across Sun-like stars for fractional luminosity (i.e., dust mass), dust distance, or even excess rate. This indicates that the processes that give rise to debris disk --- planetesimal formation, dynamical stirring that produces collisions --- must equally be insensitive to stellar parameters (temperature, mass, luminosity). This may argue that planetary system formation is quite robust --- able to occur in many different conditions. While none of the debris disks we observed are very similar to our own Solar System, there can be no question that the process of planetary system formation is quite common."
    },
    "0710/0710.4048.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory % \\abstract % context heading (optional) % {} % %  % aims heading (mandatory) % {For Seyfert galaxies, the AGN unification model provides a simple and well established explanation of the Type~1/Type~2 dichotomy through orientation based effects. The generalization of this unification model to the higher luminosity AGNs that are the quasars remains a key question. The recent detection of Type~2 Radio-Quiet quasars seems to support such an extension. We propose to further test this scenario.} %   {Type~1 but that they are seen at smaller inclination preventing us to see the broad %    emission lines present in the spectrum of Type~1 AGN.} % %  % methods heading (mandatory) % pola + deconvo we search ... {On the basis of a compilation of quasar host galaxy position angles consisting of previously published data and of new measurements performed using HST Archive images, we investigate the possible existence of a correlation between the linear polarization position angle and the host galaxy/extended emission position angle of quasars.} %% %  % results heading (mandatory) {We find that the orientation of the rest-frame UV/blue extended emission is correlated to the direction of the quasar polarization. For Type~1 quasars, the polarization is aligned with the extended UV/blue emission while these two quantities are perpendicular in Type~2 objects. This result is independent of the quasar radio-loudness. We interpret this (anti-)alignment effect in terms of scattering in a two-component polar+equatorial model which applies to both Type~1 and Type~2 objects. Moreover the orientation of the polarization --and then of the UV/blue scattered light-- does not appear correlated to the major axis of the stellar component of the host galaxy measured from near-IR images.}  %  % conclusions heading (optional), leave it empty if necessary {} %   These observations can be simply explained in the framework of an %  extension to quasars of the unification model of AGN where the polarization is produced in two scattering regions.  %% ",
        "introduction": "The study of quasars shows that we can classify them among various categories. Radio-Loud quasars (RLQ) are distinguished from the Radio-Quiet quasars (RQQ) according to their radio power (Kellerman et al. \\cite{ke89}), and the Type~1/Type~2 objects from the presence or absence of broad emission lines in their spectrum (Lawrence \\cite{law87}).  One can then wonder whether there exists a common link between all these objects, i.e. are the physical processes at the origin of all quasars the same? An interesting way to tackle this question consists in the use of the linear optical polarization. Polarimetry, in combination with spectroscopy, has led to major advances in the development of a unified scheme for AGN (see Antonucci \\cite{anto93} for a review). The discovery of hidden broad emission lines in the polarized spectrum of the Type~2 Seyfert galaxy NGC1068 led to consider these objects as intrinsically identical to Type 1 Seyfert, the edge-on orientation of a dusty torus blocking the direct view of the central engine and the broad emission line region (Antonucci \\& Miller \\cite{anto85}).  The question we investigate in this paper relates to the possible existence of a correlation between the optical polarization position angle $\\theta_{Pola}$\\footnote{The polarization position angle $\\theta_{Pola}$ is defined as the position angle of the maximal elongation of the electric vector in the plane of polarization, measured in degrees East of North.} and the orientation of the host galaxy/extended emission $PA_{host}$\\footnote{The orientation of the host galaxy $PA_{host}$ is characterized by the position angle of its major axis projected onto the plane of the sky, measured in degrees East of North.} in the case of RLQs and RQQs. The relation between quasars and their host galaxies may play a fundamental role in our understanding of the AGN phenomenon and in determining the importance of the feedback of AGN on their hosts.  Such a correlation has already been studied in the case of the less powerful AGN that are the Seyfert galaxies. Thompson \\& Martin (\\cite{toma88}) found a tendency for Seyfert 1 to have their polarization angle aligned with the major axis of their host galaxy, an observation that they interpreted as due to dichroic extinction by aligned dust grains in the Seyfert host galaxy. In the case of the more powerful AGN that are quasars, Berriman et al. (\\cite{beri90}) investigated this question. They determined by hand the $PA_{host}$ of 24 PG quasars from ground based images and computed the acute angle $\\Delta \\theta$ between $\\theta_{Pola}$ and $PA_{host}$. They observed that while more objects seem to appear at small values ($\\Delta \\theta \\leq 45\\degr$), this effect was only marginally statistically significant.  Investigating this problem with ground based data for Type~1 quasars is hampered by the inability of separating the faint host galaxy/extended emission\\footnote{In the following, we use indifferently the term host galaxy or extended emission to refer to all the extended emission around the central source including stars, ionized gas or scattered light.} from the blinding light of the quasar nucleus. Because of its high angular resolution and stable Point Spread Function (PSF), the Hubble Space Telescope (HST) permits to properly remove the contribution of the powerful quasar nucleus thus allowing the investigation of the host galaxy parameters. A large number of quasar host observing programs were carried out with the HST leading to a statistically useful sample of quasar host galaxy parameters (e.g. Bahcall et al. \\cite{ba97}; Dunlop et al. \\cite{du03} and many more see Sects.~\\ref{publidata} and \\ref{newosdat}). Our aim is to investigate the $\\theta_{Pola}/PA_{host}$ relation on the basis of high resolution HST quasar images. We will use either position angles given in the literature or $PA_{host}$ we measured ourselves from HST Archive observations (hereafter called ``\\emph{new $PA_{host}$ data}\").   The layout of this paper is as follows. In Sect.~\\ref{publidata} we introduce the samples of quasars used in this study which possess a $PA_{host}$ given in the literature. In Sect.~\\ref{tres}, we outline the samples with good imaging data but for which no $PA_{host}$ were published, and we summarize our data analysis process and the approach followed to model the HST Archive images and to derive the host galaxy parameters. In Sect.\\ref{rapola} we briefly describe the polarization and radio data. Then in Sect.~\\ref{statos} we present the statistical analysis of the sample and the results obtained. In Sect.~\\ref{discu} we discuss the results and compare them to former studies. Finally, our conclusions are summarized in Sect.~\\ref{conclu}.     %__________________________________________________________________ ",
        "conclusions": "\\label{conclu}  Using host galaxy position angles ($PA_{host}$) determined from high resolution optical/near-IR images of quasars and data from the literature, we investigate the possible existence of a correlation between the host morphology and the polarization direction in the case of Radio-Quiet and Radio-Loud quasars. We can summarize our results as follows :  \\begin{enumerate}  \\item We find an alignment between the direction of the linear polarization and the rest-frame UV/blue major axis of the host galaxy of Type~1 quasars. In the case of Type~2 objects, it is well established that the extended UV/blue light is correlated to the observed polarization, as these blue regions are thought to be dominated by scattering. Our results suggest that such an extended UV/blue scattering region is also present in Type~1 quasars.   \\item We do not find such an alignment effect with the near-IR host morphology. This suggests that the morphology of the extended UV/blue emission is not related to the morphology of the stellar component of the host galaxy which dominates in the near-IR.   \\item We observe the same $PA_{host} - \\theta_{Pola}$ behavior for either Radio-Loud or Radio-Quiet objects. This observation supports the idea that the UV/blue continuum is not entirely due to star formation processes triggered by the radio-jet.   \\item The observed correlation fits a unification model where the Type~1/Type~2 dichotomy is essentially determined by orientation effects assuming the two component scattering model of Smith et al. (\\cite{sm04,sm05}). Indeed, depending on the viewing angle to the quasar, the polarization would predominantly arise from either the equatorial or the polar scattering region giving rise to the observed behavior.  \\end{enumerate}  In order to strengthen the conclusions and to further investigate the correlations presented in this paper, new observations of quasars in the rest-frame UV/blue domain are needed, especially of Radio-Quiet objects for which few observations at $\\lambda_{rest} \\le 5000$ \\AA~are available. The detection of the extended blue/UV continuum region and the measurement of its polarization would help to test our interpretation."
    },
    "0710/0710.5721_arXiv.txt": {
        "abstract": "The equation of state for radiation is derived in a canonical formulation of the electromagnetic field. This allows one to include correction terms expected from canonical quantum gravity and to infer implications to the universe evolution in radiation dominated epochs. Corrections implied by quantum geometry can be interpreted in physically appealing ways, relating to the conformal invariance of the classical equations. ",
        "introduction": "\\label{sec:INTRODUCTION}  In theoretical cosmology, many insights can already be gained from spatially isotropic Friedmann--Robertson--Walker models \\begin{equation} \\md s^2= -\\md\\tau^2+a(\\tau)^2\\left(\\frac{\\md r^2}{1-kr^2}+r^2(\\md\\vartheta^2+ \\sin^2\\vartheta\\md\\varphi^2)\\right) \\end{equation} with $k=0$ or $\\pm 1$. The matter content in such a highly symmetric space-time can only be of the form of a perfect fluid with stress-energy tensor $T_{ab}=\\rho u_au_b+P (g_{ab}+u_au_b)$ where $\\rho$ is the energy density of the fluid, $P$ its pressure and $u^a$ the 4-velocity vector field of isotropic co-moving observers. Once an equation of state $P=P(\\rho)$ is specified to characterize the matter ingredients, the continuity equation $\\dot{\\rho}+3H(\\rho+P)=0$ with the Hubble parameter $H=\\dot{a}/a$ allows one to determine the behavior of $\\rho(a)$ in which energy density changes during the expansion or contraction of the universe. This function, in turn, enters the Friedmann equation $H^2+k/a^2=8\\pi G\\rho/3$ and allows one to derive solutions for $a(\\tau)$.  In general, one would expect the equation of state $P=P(\\rho)$ to be non-linear which would make an explicit solution of the continuity and Friedmann equations difficult. It is thus quite fortunate that in many cases linear equations of state $P=w\\rho$ with $w$ constant are sufficient to describe the main matter contributions encountered in cosmology at least phenomenologically. The influence of compact objects on cosmological scales is, for instance, described well by the simple dust equation of state $P(\\rho)=0$. Relativistic matter, mainly electromagnetic radiation, satisfies the linear equation of state $P=\\frac{1}{3}\\rho$. The latter example is an exact equation describing the Maxwell field, rather than an approximation for large scale cosmology. It is thus, at first sight, rather surprising that the dynamics of electromagnetic waves in a universe can be summarized in such a simple equation of state irrespective of details of the field configuration. The result follows in the standard way from the trace-freedom of the electromagnetic stress-energy tensor and is thus related to the conformal symmetry of Maxwell's equations. That the availability of such a simple equation of state is very special for a matter field can be seen by taking the example of a scalar field $\\phi$ with potential $V(\\phi)$. In this case, we have an energy density $\\rho=\\frac{1}{2}\\dot{\\phi}^2+V(\\phi)$ and pressure $P=\\frac{1}{2}\\dot{\\phi}^2-V(\\phi)$. Unless the scalar is free and massless, $V(\\phi)=0$ for which we have a stiff fluid $P=\\rho$, there is no simple relation between pressure and energy density independently of a specific solution.  Any conformal symmetry such as that of elecromagnetism might be broken by quantum effects especially when quantum gravity with its new scale provided by the Planck length is taken into account. The coupling of the electromagnetic field to geometry will then change, and exact conformal symmetries can easily be violated. Accordingly, one expects corrections from quantum gravity to the radiation equation of state and corresponding effects in the universe evolution during radiation dominated epochs. In loop quantum cosmology \\cite{LivRev} equations of state of matter fields are in general modified by perturbative corrections at large scales and non-perturbative ones on small scales \\cite{InvScale}. This has mainly been studied so far for a scalar field for which quantum modifications can be so strong that negative pressure results independently of the chosen potential \\cite{Inflation}. The main reason is the fact that the isotropic scalar field Hamiltonian $H_{\\phi}=\\frac{1}{2}a^{-3}p_{\\phi}^2+ a^3V(\\phi)$, where $p_{\\phi}$ is the momentum of $\\phi$, contains an inverse power of the scale factor $a$. For quantum gravity, this factor has to be quantized, too. Using the methods of \\cite{QSDV}, it turns out that inverse powers receive strong loop quantum corrections at small length scales \\cite{InvScale}. Accordingly, such modifications play a role for effective equations describing the universe after the big bang (or even during the quantum transition through the big bang singularity). During later stages, modifications are expected to decrease in size, but they might still be relevant due to sometimes tight constraints on evolution parameters.  An extension to the usual matter ingredients of cosmology with linear equations of state is, however, difficult since the modification is based on quantizations of the fundamental field Hamiltonians. Equations of state are obtained from fundamental Hamiltonians after an analysis of the matter field equations, which can be difficult in general especially when quantum effects are taken into account. The only exception is the dust case since it implies a constant Hamiltonian (the total mass of dust) which is straightforwardly quantized without any corrections. Thus, although the dust energy density is proportional to $a^{-3}$ and metric dependent in a way which involves the inverse, it does not receive any modification since the Hamiltonian, i.e.\\ total energy $a^3\\rho$, is the essential object to be quantized. For radiation with $\\rho\\propto a^{-4}$ the expectation is not clear since the total energy does behave like an inverse power of $a$, but this follows only after an indirect analysis of the field dynamics. It is not the solution $\\rho(a)\\propto a^{-4}$ of the continuity equation which is quantized but the original field Hamiltonian from which the equation of state has to be derived first. One thus has to go back to the fundamental Maxwell Hamiltonian, derive energy density and pressure and see how quantum effects change the equation of state. If this is completed, one may attempt to solve the continuity equation to obtain corrections to $\\rho(a)$.  We will derive such corrections in this article, using the canonical quantization given by loop quantum gravity \\cite{Rov,ALRev,ThomasRev}. Candidates for Hamiltonian operators of the Maxwell field have been proposed \\cite{QSDV} which show several sources of correction terms. To derive corrections to the equation of state, however, we need to perform the usual calculation in a Hamiltonian formulation.  Thus, we first present the canonical formulation for the free classical Maxwell field to rederive the standard result for the equation of state parameter $w$ without reference to an action or the stress-energy tensor. Appropriate modifications to the matter Hamiltonian $H_{M}$ are then made to derive possible loop quantum gravity corrections to the equation of state $w$. We will show that one case of corrections results again in a linear equation of state, albeit in a corrected way which depends on the basic discreteness scale of quantum gravity. In this case we are able to express, as in the classical case, the full field dynamics in terms of a simple modified $w$, and to solve explicitly for $\\rho(a)$. Our derivation takes into account inhomogeneous field configurations and presents the first modified equation of state obtained for a realistic matter source in loop quantum gravity. ",
        "conclusions": "\\label{sec:DISCUSSIONS}  We have derived here the equation of state of the Maxwell field in a canonical form, including corrections expected from loop quantum gravity. In the canonical derivation, the reason for a linear equation of state, which is trace-freedom in the Lagrangean derivation, is the fact that the same metric dependent factor $q_{ab}/\\sqrt{q}$ multiplies both terms in the Hamiltonian. The Maxwell Hamiltonian is thus simply rescaled if the metric is conformally transformed, which explains the conformal invariance of Maxwell's equations. This is special for the Maxwell field and different from, e.g., a scalar field with a non-vanishing potential.  The same fact allows one to quantize the Hamiltonian in a way which affects both the electric and magnetic term in the same way, at least as far as the metric dependence is concerned. One then obtains a single correction function $\\alpha=\\beta$ which only corrects the metric dependence of the total scale of the Hamiltonian. In this sense, conformal invariance is preserved even after quantization. (But this would not be the case if a quantization is used which results in $\\alpha\\not=\\beta$.)  This preservation of the form of the Hamiltonian explains why we are still able to derive an equation of state independently of the specific field dynamics and that it remains linear. However, the classical value $w=\\frac{1}{3}$ is corrected due to quantum effects in the space-time structure. This modification is also understandable from a Lagrangean perspective, together with basic information from the loop quantization. Employing trace freedom of the stress-energy tensor to derive the equation of state, we have to use the inverse metric in $g^{ab}T_{ab}$. But from loop quantum gravity we know that, when quantized, not all components of the inverse metric agree with inverse operators of the quantization. For the scale factor of an isotropic metric, for instance, we have $\\widehat{a^{-1}}\\not=\\mbox{``$\\hat{a}^{-1}$''}$ since the right hand side is not even defined \\cite{InvScale}. While the left hand side is defined through identities such as (\\ref{poissonbracketofvolume}), it satisfies $\\widehat{a^{-1}}\\hat{a}\\not=1$ and thus shows deviations from the classical expectation $a^{-1}a=1$ on small scales which were captured here in correction functions. As derived in detail, this implies scale dependent modifications to the equation of state parameter $w_{\\rm eff}$.  The result can also be interpreted in more physical terms. The classical behavior $\\rho(a)\\propto a^{-4}$ can be understood as a combination of a dilution factor $a^{-3}$ and an additional redshift factor $a^{-1}$ for radiation in an expanding universe. As we have seen, this is corrected to $\\alpha(a)a^{-4}$ where $\\alpha(a)$ corrects the metric factor $q_{ab}/\\sqrt{q}\\sim a^{-1}\\delta_{ab}$. Since this is only a single inverse power of $a$ for an isotropic solution, we can interpret the result as saying that only redshift receives corrections due to quantum effects on electromagnetic propagation. The dilution factor due to expansion is unmodified, except that the background evolution $a(t)$ itself receives corrections. This agrees with the result for dust, which is only diluted and has an unmodified equation of state even after quantization \\footnote{But it disagrees with \\cite{Metamorph} both for dust and radiation, where a direct quantization of energy densities exclusively for isotropic fields was attempted.}. Unlike dust, for radiation one has to refer to the inhomogeneous field and its quantum Hamiltonian to derive a reliable equation of state, as presented here."
    },
    "0710/0710.5517_arXiv.txt": {
        "abstract": "In a large region of the supersymmetry parameter space, the annihilation cross section for neutralino dark matter is strongly dependent on the relative velocity of the incoming particles. We explore the consequences of this velocity dependence in the context of indirect detection of dark matter from the galactic center. We find that the increase in the annihilation cross section at high velocities leads to a flattening of the halo density profile near the galactic center and an enhancement of the annihilation signal. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.2930_arXiv.txt": {
        "abstract": "Hot Jupiters are new laboratories for the physics of giant planet atmospheres. Subject to unusual forcing conditions, the circulation regime on these planets may be unlike anything known in the Solar System. Characterizing the atmospheric circulation of hot Jupiters is necessary for reliable interpretation of the multifaceted data currently being collected on these planets. We discuss several fundamental concepts of atmospheric dynamics that are likely central to obtaining a solid understanding of these fascinating atmospheres.  A particular effort is made to compare the various modeling approaches employed so far to address this challenging problem. ",
        "introduction": "An exploding body of observations promises to unveil the meteorology of giant planets around other stars.  Because of their high temperatures, short orbital periods, and likelihood of transiting their stars, the hot Jupiters are yielding their secrets most easily, and we now have constraints on radii, composition, albedo, dayside temperature structure, and even day-night temperature distributions for a variety of planets \\citep[e.g.,][] {knutson-etal-2007b, cowan-etal-2007, harrington-etal-2006, harrington-etal-2007, charbonneau-etal-2002, charbonneau-etal-2005, charbonneau-etal-2007, deming-etal-2005, deming-etal-2006}. A knowledge of atmospheric dynamics is crucial for interpreting these observations.  Understanding the atmospheric circulation of {\\it any} planet is a difficult task, and hot Jupiters are no exception.  The difficulty results from the nonlinearity of fluid motion and from the complex interaction between radiation, fluid flow, and cloud microphysics. Turbulence, convection, atmospheric waves, vortices, and jet streams can all interact in a complex manner across a range of temporal and spatial scales.  Furthermore, radiative transfer and dynamics are coupled and cannot be understood in isolation.  For example, the equator-to-pole temperature contrasts on terrestrial planets depend not only on the rate of latitudinal energy transport by the circulation but also on the way the atmospheric radiation field is affected by the advected temperature field.  Cloud microphysics complicates the problem even more, because the large-scale radiative properties of clouds depend on the microphysics of the cloud particles (particle number density, shape, size distribution, and absorption/scattering properties).  We thus have a non-linear, coupled radiation-hydrodynamics problem potentially involving interactions over 14 orders of magnitude, from the cloud-particle scale ($\\sim1\\,\\mu$m) to the global scale ($\\sim10^8\\,$m for hot Jupiters).   It is worth emphasizing here the important difference between the complexity of planetary atmospheres and the relative simplicity of stars, exemplified by the main sequence in the HR diagram.  While specifying a few global parameters---such as mass, composition, and age---is typically sufficient to understand the key observable properties of stars, this is generally not the case for planetary atmospheres. It is possible that a significant observable diversity exists even among a group of extrasolar planets which share similar global attributes.  Understanding such complexity is a challenge for extrasolar planetary science.  A proven method for dealing with the complexity of planetary atmospheres is the concept of a {\\it model hierarchy.} Even a hypothetical computer model that included all relevant processes and perfectly simulated an atmosphere would not, by itself, guarantee an {\\it understanding} of the atmospheric behavior any more than would the observations of the real atmosphere \\citep{held-2005}. This is because, in complex numerical simulations, it is often unclear how and why the simulation produces a specific behavior. To discern which processes cause which outcomes, it is important to compare models with a range of complexities (in which various processes are turned on or off) to build a hierarchical understanding.  For example, east-west jet streams occur in the atmospheres of all the planets in our Solar System.  The study of these jets involves a wide range of models, all of which have important lessons to teach.  The most idealized are pure 2D models that  investigate jet formation in the simplest possible context of horizontal, 2D turbulence interacting with planetary rotation \\citep[e.g.,][]{williams-1978, yoden-yamada-1993, cho-polvani-1996b, huang-robinson-1998, sukoriansky-etal-2007}.    Despite the fact that these simplified models exclude vertical structure, thermodynamics, radiation, and clouds and provide only a crude parameterization of turbulent stirring, they produce jets with several similarities to those on the planets.  Next are one-layer ``shallow-water-type'' models that allow the fluid thickness to vary, which introduces buoyancy waves, alters the vortex interaction lengths, and hence changes the details of jet formation \\citep{cho-polvani-1996a, cho-polvani-1996b, scott-polvani-2007, showman-2007}.  3D models with simplified forcing allow the investigation of jet vertical structure, the interaction of heat transport and jet formation, and the 3D stability of jets to various instabilities \\citep{cho-etal-2001, williams-1979, williams-2003a, schneider-2006, lian-showman-2007}. Such models suggest, for example, that Jupiter and Saturn's superrotating\\footnote{Superrotation is defined simply as a positive (eastward) longitudinally averaged wind velocity at the equator, so that the atmospheric gas rotates faster than the planet.} equatorial jets may require 3D dynamics. Finally are full 3D general-circulation models (GCMs) that include realistic radiative transfer, representations of clouds, and other effects necessary for detailed predictions and comparisons with observations of specific planets. The comparison of simple models with more complex models provides  insights not easily obtainable from one type of model alone.  This lesson applies equally to hot Jupiters: a hierarchy of models ranging from simple to complex will be necessary to build a robust understanding.  Here we review our current understanding of atmospheric circulation on hot Jupiters.  We first describe basic aspects of relevant theory from atmospheric dynamics; this is followed by a detailed comparison of the equations, forcing methods, and results obtained by the different groups attempting to model the atmospheric circulation on these fascinating objects. ",
        "conclusions": "The study of hot Jupiter atmospheres is maturing.  While data of increasing quality are being collected, atmospheric models are also being refined to help build a robust, hierarchical understanding of these unusual atmospheres.  Indeed, one of the main motivations for studying these atmospheres, which is also a main source of difficulty, is the unusual physical regime that characterizes them.  Hence, hot Jupiters represent new laboratories for studying the complex physics of giant planet atmospheres.  In this way, they offer the promise of extending the boundary of comparative planetology well beyond the solar-system planets."
    },
    "0710/0710.4511_arXiv.txt": {
        "abstract": "We present a novel technique to phase-lock two lasers with controllable frequency difference. In our setup, one sideband of a current modulated Vertical-Cavity Surface-Emitting Laser (VCSEL) is phase locked to the master laser by injection seeding, while another sideband of the VCSEL is used to phase lock the slave laser. The slave laser is therefore locked in phase with the master laser, with a frequency difference tunable up to about 35 GHz. The sideband suppression rate of the slave laser is more than 30dB at 30 $\\mu$W seed power. The heterodyne spectrum between master and slave has a linewidth of less than 1 Hz. A narrow linewidth spectrum of coherent population trapping in rubidium is achieved using such beams. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.3041_arXiv.txt": {
        "abstract": "A perturbative analysis is used to investigate the effect of rotation on the instability of a steady accretion shock (SASI) in a simple toy-model, in view of better understanding supernova explosions in which the collapsing core contains angular momentum. A cylindrical geometry is chosen for the sake of simplicity. Even when the centrifugal force is very small, rotation can have a strong effect on the non-axisymmetric modes of SASI by increasing the growth rate of the spiral modes rotating in the same direction as the steady flow. Counter-rotating spiral modes are significantly damped, while axisymmetric modes are hardly affected by rotation. The growth rates of spiral modes have a nearly linear dependence on the specific angular momentum of the flow. The fundamental one-armed spiral mode ($m=1$) is favoured for small rotation rates, whereas stronger rotation rates favour the mode $m=2$. A WKB analysis of higher harmonics indicates that the efficiency of the advective-acoustic cycles associated to spiral modes is strongly affected by rotation in the same manner as low frequency modes, whereas the purely acoustic cycles are stable. These results suggest that the linear phase of SASI in rotating core-collapse supernovae naturally selects a spiral mode rotating in the same direction of the flow, as observed in the 3D numerical simulations of Blondin \\& Mezzacappa (2007). This emphasizes the need for a 3D approach of rotating core-collapse, before conclusions on the explosion mechanisms and pulsar kicks can be drawn. ",
        "introduction": "}  Despite extensive studies, the explosion mechanism of core-collapse supernovae is still elusive. According to the delayed explosion scenario, the shock is first stalled at a distance of a few hundred kilometers, and then revived after neutrinos diffuse out of the proto neutron star. Unfortunately, numerical simulations suggest that neutrino heating may not be efficient enough, at least in spherical symmetry \\citep{lieb05}.  Recent studies have shown that the spherical stalled shock is unstable against non radial perturbations with a low degree $l=1,2$ even if the flow is convectively stable. This result was demonstrated using axisymmetric numerical simulations \\citep{blo03,blo06,sch06,sch08,ohn06} and linear stability analyses \\citep{gal05,fog07,yam07}. Some numerical simulations have shown that this hydrodynamical instability, often called SASI, may assist the revival of the shock and trigger a successful explosion, powered either by neutrino heating \\citep{mar07} or by acoustic waves \\citep{bur06}. Some observed properties of young neutron stars may also be the consequences of SASI, such as their distribution of velocities \\citep{sch04, sch06} or their spin \\citep{blo07a,blo07b}.  Until now, most studies of SASI have assumed that the unperturbed flow is purely radial and not rotating. Since the angular momentum of massive stars is likely to be large \\citep{heg05}, it is desirable to understand how the properties of SASI are affected by rotation. In this Letter, the effect of rotation on the linear stage of SASI is investigated using a perturbative analysis in order to shed light on one of the surprising results observed by \\cite{blo07a} in their 3D numerical simulations: the development of SASI seems to systematically favour a spiral mode rotating in the same direction as the accretion flow. As a consequence of momentum conservation, this mode diminishes and may even reverse the angular momentum acquired by the proto-neutron star from the stationary flow. Incidentally, the present linear study does not address another surprising result of \\cite{blo07a}, that a spiral mode of SASI always dominate the axisymmetric mode even without rotation. Following an approach similar to \\citet{fog07} (hereafter FGSJ07), we first compute the eigenfrequencies by solving accurately a boundary value problem between the shock surface and the accretor surface; in a second step, we use the same WKB method as in FGSJ07 to measure the stability of purely acoustic and advective-acoustic cycles in this region. This approach is different from \\cite{lam07}, which is based on the approximate derivation of a dispersion relation. Rather than the complexity of describing the non-spherical shape of a shock deformed by rotation \\citep{yam05}, we have chosen, as a first step, to solve the much simpler problem of a cylindrical accretion shock. This flow is simple enough to allow for a complete coverage of the parameter space and a physical insight of the main effects of rotation on SASI. Once characterized, these effects can be transposed into the more complex geometry of a rotating stellar core. ",
        "conclusions": "}  \\subsection{Instability mechanism} \\label{mechanism}  As underlined in Sect.~3, the dynamical effect of the centrifugal force on the stationary flow is modest. We anticipated in Sect.~2 that the only linear effect of angular momentum is a Doppler shift of the eigenfrequency $\\omega'=\\omega-m\\Omega(r)$, where $\\Omega(r)$ is the local rotation frequency. This leaves the axisymmetric mode $m=0$ unaffected and explains the relative insensitivity of its growth rate with respect to the rotation rate, at least for moderate angular momentum. The strong effect of rotation on the growth rate of the spiral modes can thus be traced back to this Doppler shifted frequency. What is the mechanism of the instability ? As seen in the previous section, the destabilizing role of rotation does not seem related to the presence or absence of a corotation radius, thus discarding a Papaloizou-Pringle mechanism \\citep{gol85}. Two possibilities have been proposed for the mechanism of SASI without rotation; one is the advective-acoustic mechanism \\citep{fog00,fog01,fog02} and the other is the purely acoustic mechanism \\citep{blo06}. Up to now, there is no satisfactory direct argument for the mechanism of the modes with a long wavelength. FGSJ07 used a WKB approximation to prove that the instability of the modes with a short wavelength is due to an advective-acoustic mechanism and extrapolated this conclusion to the modes with a long wavelength, which are the most unstable. This method, recalled in Appendix~B, is based on the identification of acoustic waves and advected waves at a radius immediately below the shock surface, and the measurement of their coupling coefficients, above this radius due to the shock, and below this radius due to the flow gradients. These coupling processes are responsible for the existence of two cycles, namely a purely acoustic cycle characterized by an efficiency ${\\cal R}$, and an advective-acoustic cycle characterized by an efficiency ${\\cal Q}$.  By using the same method, the present study does not address directly the instability mechanism of long wavelength modes. However, the WKB approximation enables us to describe, in a conclusive manner, the instability mechanism of short wavelength modes affected by rotation. First we checked that when the shock distance is increased ($r_{\\rm sh}=20r_*$), the overtones are also unstable and their growth rate is an oscillatory function of the frequency similar to Fig.~7 of FGSJ07. The effect of rotation on the advective-acoustic cycle is illustrated by Fig~3, for the spiral modes $m=\\pm1$ corresponding to the $10$-th overtone, as a function of the rotation rate. The cycle efficiency ${\\cal Q}$ is strongly amplified by rotation if $m>0$,  while strongly damped if $m<0$. The stabilization of the counter-rotating spiral coincides with a marginally stable cycle ${\\cal Q}\\sim1$. The calculation of the amplification factor ${\\cal R}$ of perturbations during each purely acoustic cycle indicates its stability (${\\cal R}<1$). Contrary to the expectation of \\citet{lam07} (see next subsection), rotation clearly favours the spiral mode of the advective-acoustic cycle.  This consequence of rotation established unambiguously for short wavelength perturbations is identical to the influence of rotation on the fundamental mode of SASI: we consider this a new hint that the advective-acoustic mechanism can be extrapolated to low frequencies. The detailed analysis of the consequences of the Doppler shifted frequency on the increase of the advective-acoustic efficiency ${\\cal Q}$ will be presented elsewhere (\\cite{yam08}, in preparation).  \\subsection{Comments on the Results of Laming (2007)} \\label{comment}  The effect of rotation on the growth rate of SASI, established in Sect. 3 in a cylindrical geometry, is qualitatively similar to the effect conjectured by \\cite{lam07} (hereafter L07). Nevertheless, their investigation about the instability mechanism led them to a different interpretation of the roles of the acoustic and advective-acoustic cycles.  We must point out a fundamental difference between the method of L07 and ours: by using a WKB approximation, we have carefuly defined the range of validity  of our method, namely short wavelength modes. This guarantees that the advective-acoustic interpretation of the instability mechanism is physical and robust, at least in some parameter range. In contrast, the existence of a purely acoustic instability is still a conjecture because the domain of validity of the method used by L07 is ambiguous: their analytical derivation of a dispersion relation when advection is included requires to neglect terms of order $(v_r /\\omega r)$ while terms of order ${\\cal M}$ are retained. This approximation is not supported by the results of their Fig.~2, which indicates that $(v_{r,{\\rm sh}}/\\omega r_{\\rm sh})$ is comparable to or larger than ${\\cal M}_{\\rm sh}$ for the modes $l=0$ and $l=1$. An accurate description of this acoustic mode, even in a simplified set up, would be useful to gain confidence in its possible existence.  In addition to the question of the validity of the approximations used by L07, we find that our results invalidate their reasoning concerning the instability mechanism. They proposed that the advective-acoustic mechanism would be essential if $r_{\\rm sh}/r_*\\ge 10$, whereas a purely acoustic unstable process would be dominant for small shock radii, and they argued that rotation is a key ingredient to discriminate between the two mechanisms. When rotation is included, its effect on SASI has been attributed by L07 to a purely acoustic mechanism, despite the results of their Table 3. However, their view that rotation cannot possibly enhance the growth of the advective-acoustic cycle is clearly incorrect, at least for the short wavelength modes (our Fig.~\\ref{fig4}).  \\subsection{Consequences of rotation on supernova explosions} \\label{consequence}  The perturbative study of a simple cylindrical configuration has enabled us to cover a large parameter space of shock radii and rotation rates, in order to (i) demonstrate the linear selection of non-axisymmetric modes, (ii) establish a correlation between the preferred direction of the spiral SASI and the rotation of the collapsing core, (iii) identify the advective-acoustic mechanism at work for short wavelength spiral perturbations.  The fact that rotation favours a spiral mode $m=1,2$ in a cylindrical flow seems directly connected to the property observed by \\cite{blo07a} in their 3D simulations including rotation. Tracing back the main influence of rotation to the local Doppler shifted frequency $\\omega-m\\Omega$, we may indeed expect a similar destabilization of the spiral modes with a positive value of $m$, a stabilization of the counter-rotating ones, and a comparatively weak influence on the axisymmetric modes.  Even a moderate amount of angular momentum results in a shortening of the growth time of SASI through the destabilization of a non-axisymmetric mode. The promising consequences of SASI on both the explosion mechanisms and the pulsar kick could thus be considerably modified, since they were established on the basis of axisymmetric numerical simulations \\citep{bur06,bur07,mar07,sch04,sch06}. Our study suggests that the effect of rotation on the linear phase of SASI can be safely neglected only for slowly rotating progenitors with a specific angular momentum $L\\ll 2\\pi \\cdot 10^{14}$ cm$^2/$s. Although a fast growth of SASI might be helpful to an early shock revival, the dynamical effects of a spiral mode $m=1$, and even $m=2$, on the possible explosion mechanisms are not known yet.  If the direction of the kick were determined by the geometry of the most unstable $l=1$ SASI mode, our perturbative approach would suggest a kick-spin misalignment. The strength of the equatorial kick may be diminished by the domination of a symmetric mode $m=2$. It is worth noting however that the relationship between the timescale of the most unstable SASI mode and the onset of explosion is not straightforward, and should be evaluated by future 3D numerical simulations. Our linear approach modestly aims at guiding our intuition for the interpretation of these simulations.    \\appendix"
    },
    "0710/0710.3780_arXiv.txt": {
        "abstract": "Imaging data from the Sloan Digital Sky Survey are used to characterize the population of galaxies in groups and clusters detected with the MaxBCG algorithm. We investigate the dependence of Brightest Cluster Galaxy (BCG) luminosity, and the distributions of satellite galaxy luminosity and satellite color, on cluster properties over the redshift range $0.1 \\le z \\le 0.3$. The size of the dataset allows us to make measurements in many bins of cluster richness, radius and redshift. We find that, within \\rtwo\\ of clusters with mass above $3 \\times 10^{13}$\\Msun, the luminosity function of both red and blue satellites is only weakly dependent on richness. We further find that the shape of the satellite luminosity function does not depend on cluster-centric distance for magnitudes brighter than \\Mi\\ $= -19$. However, the mix of faint red and blue galaxies changes dramatically. The satellite red fraction is dependent on cluster-centric distance, galaxy luminosity and cluster mass, and also increases by $\\sim$5\\% between redshifts 0.28 and 0.2, independent of richness. We find that BCG luminosity is tightly correlated with cluster richness, scaling as $L_{BCG} \\sim M_{200}^{0.3}$, and has a Gaussian distribution at fixed richness, with $\\sigma_{log L} \\sim 0.17$ for massive clusters.  The ratios of BCG luminosity to total cluster luminosity and characteristic satellite luminosity scale strongly with cluster richness: in richer systems, BCGs contribute a smaller fraction of the total light, but are brighter compared to typical satellites. This study demonstrates the power of cross-correlation techniques for measuring galaxy populations in purely photometric data. ",
        "introduction": "\\label{sec:intro}  Clusters of galaxies are important systems for studying both galaxy evolution and cosmology. Used as laboratories with well-defined environments, these massive objects are a tool for investigating processes that influence galaxies' physical characteristics. Used as tracers of the underlying mass distribution, they are a tool for investigating the evolution of structure and the nature of dark energy in the Universe.  These objectives are closely linked: cosmological studies require accurate knowledge of cluster selection and redshift- and mass-observable relations, and these facts are directly related the evolutionary properties of the cluster galaxy population.  In the context of galaxy evolution, the high-density environment of galaxy clusters is a particularly interesting place to examine the galaxy population.  Several studies have suggested substantial galaxy transformation in such environments.  Galaxy morphology, star-formation rate, and luminosity have long been known to depend on cluster properties and to depart significantly from the cosmological average \\citep[\\eg,][]{Hubble26, Abell62, Oemler74, Dressler80}. Historically, the cluster galaxy content has been quantified by the luminosity function (LF) and type fraction (such as the late-type or blue fraction).  Measurement of these quantities as a function of cluster mass, redshift, and distance from the cluster center provides insight into the underlying physical mechanisms responsible for these trends.  While the LF is primarily a decreasing function of luminosity, galaxies are bimodal in color and spectral type, with red, early-type galaxies displaying little ongoing star formation, and blue, late-type galaxies exhibiting signs of recent star formation \\citep[\\eg,][]{Strateva01, Baldry04, Bell04, Menanteau06, Blanton05b}.  This bimodality was in place by $z\\sim1$ at the latest and may provide a signpost of galaxy transformation \\citep[\\eg,][]{Faber07}.  Although this bimodality persists in all environments, the fraction of galaxies in each class (a.k.a. the red and blue fractions) changes systematically with local density; this trend is the so-called morphology-density relationship \\citep{Oemler74,Dressler80,Dressler97,Smith05}.  In recent large galaxy surveys this work has been extended to show that a wide range of galaxy properties, including morphology, star-formation rate, and color, depend on local density \\citep{Gomez03,Balogh04a,Balogh04b,Hogg04,Kauffmann04,Tovmassian04,Blanton05b,Christlein05,Croton05,Rojas05,Cooper06,Mandelbaum06,Cooper07}.  The advent of large galaxy surveys has made it possible to place observational constraints on both the type fraction and the LF as a function of cluster mass (the conditional luminosity function, CLF), and to further investigate how these quantities depend on other variables.  In the Sloan Digital Sky Survey \\citep*[][SDSS]{York00}, \\citet{Goto02} and \\citet{Hansen05} examined the LF as a function of cluster richness and cluster-centric distance for systems found in the SDSS Early Data Release, and \\citet{Weinmann06b} measured the CLF measured from a group catalog derived from the spectroscopic sample of SDSS DR2.  Using the 2dF Galaxy Redshift Survey, \\citet{depropris03} and \\citet{Robotham06} compared the LF in high- and low-mass systems; a similar study was performed with a sample of 93 X-ray selected clusters \\citep{LMS04}.  The dependence of the LF on galaxy color has been recently investigated \\citep{PopessoLF} as has the galaxy type fraction \\citep{Goto03,depropris04} for clusters in these large surveys.  The type fraction of galaxies in clusters depends on both cluster richness and redshift: the fraction of star-forming galaxies at fixed local density is larger at higher redshift, an effect known as the Butcher--Oemler effect \\citep{BO78,ButcherOemler}. This effect is now well-documented, if not entirely well-explained, in clusters over a wide range of masses by studies of the blue fraction, or its converse, the red fraction \\citep{Rakos95,Margoniner00,Ellingson01,KodamaBower01,Margoniner01,depropris04,Martinez06,Gerke07}. Other indicators of galaxy state, including galaxy morphology and emission line strength, also show a trend with redshift \\citep{AS93,ODB97,Balogh97,Couch98,vanDokkum00,Fasano00,Lubin02,Goto03,Treu03,Wilman05,Poggianti06,Desai07,vanderwel07}. However, few samples to date have had both large numbers of systems and well-understood mass proxies, so the dynamical range and mass resolution of these previous works has been somewhat limited.  Another characteristic of galaxy clusters is the presence of a highly-luminous galaxy near the cluster center --- the Brightest Cluster Galaxy (BCG).  In addition to being extraordinarily luminous, BCGs differ in a number of ways from other cluster members: they tend to have extended light profiles \\citep{Matthews64, Tonry87, Schombert88, Gonzalez00, Gonzalez03}, larger size at fixed luminosity than other early types \\citep[][and references therein]{Bernardi07} and may contain a larger fraction of dark matter than typical galaxies \\citep[\\eg,][]{Mandelbaum06,Anja07}.  Also, while the traditional fitting function to the LF of \\citet{Schechter76} provides a good fit for satellite galaxies, BCGs follow a different distribution, causing the so-called ``bright end bump'' \\citep[\\eg,][]{Hansen05}. This difference between the BCG and other satellites in a cluster is also manifested in the luminosity gap statistic, the difference in luminosity between the BCG and the next brightest cluster member, which may be indicative of the special accretion history of BCGs \\citep{Ostriker75, Tremaine77, Loh06}. BCGs are also distinct from other galaxies with similar mass that are not at the center of cluster-sized potential wells \\citep{Anja07}. Indeed, the properties of BCGs seem to be closely linked to properties of their host clusters \\citep{Sandage73,Sch83,Schombert88, Edge91, Brough02, Degrandi04,Brough05,Loh06}, including to the masses of their parent halos \\citep{LinMohr04, Mandelbaum06, ZCZ07}.  The outer light profile of BCGs often merges smoothly with the diffuse intra-cluster light (ICL), suggesting again a coupling between the BCG and the host halo \\citep{Gonzalez05}.  Generally the BCG+ICL light is closely linked with cluster mass \\citep{Zibetti05,Conroy07ICL,Purcell07,Gonzalez07}. As BCGs have properties different from those of the rest of the cluster members, BCGs and satellites are often analyzed separately, and we follow this convention here.   It is expected that the properties of cluster galaxies are closely tied to the merging and accretion history of their parent dark matter halos.  Current models based on Cold Dark Matter (CDM) suggest that the stars in BCGs were formed in dense peaks quite early but that the BCGs were assembled in a series of galaxy merging events that continue until relatively recent times \\citep[e.g.][]{AS98,Dubinski98,Gao04a,BK06,DeLucia07}.  Satellite galaxies in clusters are now generally understood to be hosted by smaller dark matter halos that have merged into the parent halo. Several features of the observed cluster galaxy populations can be understood based on the assembly history of the parent halo \\citep[\\eg,][]{Poggianti06, Iro07}.  One such characterization is the core distinction between central galaxies and satellites: the central concentrations of mass and light in massive halos continue to build up while the growth of satellite systems is halted upon accretion.  BCG luminosity, and the correlations between this luminosity and both cluster mass and satellite properties, can thus provide insight into the assembly histories of clusters.  The merger history of dark matter halos alone is not enough to account for the bimodality in galaxy properties and its dependence on redshift and environment.  Several physical processes have been proposed to transform star-forming galaxies into the typical cluster galaxies on the red sequence.  Although the relative strengths of these processes are still hotly debated, a consensus is emerging that the main transformation mechanisms are related to the mass of the host halo and whether (and for how long) the galaxy has been a satellite within a larger system.  Among the suggested transformation mechanisms, some are expected to be most effective in rich clusters, such as ram-pressure stripping \\citep{GunnGott72}, interaction with the cluster potential \\citep{ByrdValtonen90} and high-velocity close encounters \\citep[``harassment;''][]{Moore96}. However, studies of very poor systems have shown that the environmental dependence of galaxy properties is not limited to the richest objects \\citep[\\eg,][]{Zabludoff98,Weinmann06a,Gerke07}. Processes that can operate efficiently in low velocity dispersion systems therefore must also play a role in shaping the galaxy population: \\eg, galaxy mergers \\citep{TT72} and ``strangulation,'' a cutoff to gas accretion onto galaxy disks by stripping or AGN feedback \\citep{Larson80,BNM00,Croton06}. It is likely that some combination of these effects is at work.  Distinguishing their relative significance requires precisely quantifying cluster galaxy properties over a wide range of masses and as a function of cluster-centric distance. These data will provide information on both the assembly histories of clusters and on the physical mechanisms that trigger and quench star formation.  Models of galaxy evolution in a cosmological context most readily predict galaxy properties as a function of halo mass rather than cluster observables.  In order to make these comparisons, a reliable mass--observable relationship is a prerequisite. In addition, there is consensus that the luminosity function and type fraction both depend on a number of variables, complicating detailed comparison between clusters of different mass.  For example, the cluster galaxy LF depends on cluster-centric distance, and the size of the bound regions of clusters scales with mass.  In order to make physically meaningful comparisons between LFs of different mass clusters, an aperture scaled to the bound region is therefore preferable to a fixed metric aperture.  With recent extensive surveys providing well-calibrated cluster catalogs spanning a wide range in mass, it is possible to examine in detail the dependence of the cluster galaxy population on several cluster and galaxy properties simultaneously.  Large, homogeneous photometric surveys such as the SDSS provide rich data with which to characterize the cluster galaxy population.  These data have been used to define large, robust, clean samples of galaxy clusters with accurate photometric redshifts.  These samples are sizable enough to split on several variables allowing detailed statistical exploration of the galaxy populations in clusters. Currently, the largest sample of clusters available is the \\maxbcg\\ catalog from the Sloan Digital Sky Survey \\citep{Koester07a}.  The selection effects of the cluster-finding algorithm are well understood \\citep{Koester07b, Rozo07a}, and there are a number of studies exploring the mass--richness relationship for these objects \\citep{Rozo07b,Becker07,Sheldon07a, Johnston07b,Rykoff07}.  In this work, we convert cluster richness to cluster mass using weak lensing measurements of the \\maxbcg\\ mass--richness relation \\citep{Sheldon07a,Johnston07b}.  The quality and quantity of these data allow for detailed measurements of a variety of cluster galaxy properties as a function of cluster mass and radius.  For satellite galaxies we then measure the luminosity function of all, red, and blue satellites conditional on both mass and cluster radius, and investigate the dependence of the red fraction of satellites on cluster mass, redshift, galaxy luminosity and distance from cluster center; for BCGs we quantify the dependence on cluster mass of both the BCG luminosity and the relationship between the BCG luminosity and satellite galaxy luminosities.  Although this sample of clusters extends only to $z = 0.3$, these objects provide a valuable low-redshift baseline with which higher redshift samples may be compared.  In this paper we use a statistical background-subtraction technique to measure mean galaxy properties over a wide range of color and luminosity in \\maxbcg\\ clusters.  We average the signal from many clusters binned by cluster properties, and statistically subtract the contribution from random galaxies along the line of sight.  This method of cross-correlating clusters with the galaxy population provides very precise statistical measurements, and allows us to study blue and low luminosity galaxies that are indistinguishable from the background in individual clusters.  We test the background-correction algorithm by running the full analysis on realistic mock catalogs, and find that we are able to robustly recover 3D cluster properties using these methods.  The statistical techniques presented here for background correction of photometric data will be directly applicable to future large multi-band imaging programs, including the Dark Energy Survey \\citep[DES\\footnote{\\tt http://www.darkenergysurvey.org},][]{DES} and the Large Synoptic Survey Telescope \\citep[LSST\\footnote{{\\tt http://www.lsst.org}},][]{LSST} and the Panoramic Survey Telescope \\& Rapid Response System \\citep[Pan-STARRS\\footnote{{\\tt http://pan-starrs.ifa.hawaii.edu}},][]{PANSTARRS} and will be essential for leveraging these data to provide the desired insights into cosmology and galaxy evolution.  The paper is organized as follows: in \\S\\ \\ref{sec:data} we describe the SDSS and simulation data used; we present the stacking and background-correction method in \\S\\ \\ref{sec:methods}. Our primary results are given in \\S\\ \\ref{sec:results}: \\S\\ \\ref{sec:sats} presents the luminosity and color characteristics of the satellite population, while \\S\\ \\ref{sec:BCGs} discusses the BCG population. A summary and discussion of the implications of the results is given in \\S\\ \\ref{sec:conclusion}.  The notation used for cluster-related variables in previous \\maxbcg\\ work includes defining \\Ntwo\\ and $L_{200}$ as the counts and $i$-band luminosity of red-sequence galaxies within the measurement aperture of the cluster finder and with L $> 0.4L_*$. We note that as the aperture for cluster finding was determined with a previous definition of richness, it is not strictly the true value of \\rtwo\\ for these systems (in fact, it is larger), and only red galaxies are included in these definitions. In this work we will refer to the total excess luminosity associated with the light from galaxies of {\\em all} colors above a luminosity threshold and within the measured \\rtwo\\ of these systems as $L_{200}$.  In addition, we follow the standard convention of using $R$ to to denote projected, 2D radii and $r$ to refer to deprojected, 3D radii. Where necessary for computing distances, we assume a flat, LCDM cosmology with H$_{0} = 100h$ km s$^{-1}$ Mpc$^{-1}$, $h = 0.7$, and matter density $\\Omega_m = 0.3$. ",
        "conclusions": "\\label{sec:conclusion} In this study, we have examined the properties of cluster-associated galaxies, separating BCGs from satellites and focusing on trends in galaxy color and luminosity as a function of cluster richness, distance from cluster center, and redshift. We use clusters in the SDSS \\maxbcg\\ sample, the largest set of clusters identified to date. Employing photometric data alone, we apply cross-correlation background-correction techniques to characterize the cluster-associated galaxy population, within $5\\times$\\rtwo\\ and brighter than \\Mi\\ $< -19$, around \\nclust\\ systems spanning more than two decades in mass in the redshift range $0.1 \\le z \\le 0.3$.  Our principle results are as follows. \\begin{enumerate}  \\item The luminosity function of satellites within \\rtwo\\ as a function of cluster mass for systems with mass greater than $3 \\times 10^{13}$\\Msun\\ shows remarkable uniformity for \\Mi\\ $< -19$. The characteristic satellite luminosity \\Lstar\\ is only weakly dependent on cluster richness.  \\item The shape of the luminosity function of satellites brighter than \\Mi\\ $< -19$ does not change with cluster-centric radius. However, the color-separated luminosity functions of satellites as a function of \\rad\\ and of \\Ntwo\\ show that the mix of sub-\\Lstar\\ red and blue galaxies changes dramatically as a function of radius. In contrast, the relative number of red and blue galaxies at the bright end is roughly constant with radius.  \\item The average color of satellite galaxies is redder near cluster centers, but this trend is largely a reflection of the changing ratio of red to blue galaxies mentioned above.  This effect is a very weak function of cluster mass over the range investigated here.  \\item The fraction of red galaxies increases with cluster mass over the full range explored, although only weakly for \\Mtwo $\\gae 10^{14}$\\Msun.  This fraction decreases with cluster-centric distance until $2\\times$\\rtwo, decreases with luminosity for $L<$ \\Lstar, and is constant for $r>2\\times$\\rtwo\\ and $L>$ \\Lstar.  \\item The fraction of cluster galaxies within \\rtwo\\ that are red increases by $\\sim$5\\% during the 0.8 Gyr between redshift $z=0.28$ to $z=0.2$; this change is roughly independent of mass over the range investigated.  \\item The luminosity of BCGs and the ratios of BCG luminosity to total cluster luminosity and to characteristic satellite luminosity are all correlated strongly with cluster mass, and we have quantified each of these scalings over the mass range $3 \\times 10^{13}$\\Msun\\ to $9 \\times 10^{14}$\\Msun.  The BCG luminosity has a Gaussian distribution at fixed cluster richness, with dispersion $\\sigma_{logL} \\sim 0.17$ for clusters with \\Mtwo\\ $ > 10^{14}$\\Msun.  \\end{enumerate}  While these results are in general agreement with previous observational work, due to the volume probed we are able to investigate a wider mass range, extending the statistics to higher mass than previous samples.  We are also able to split the sample into finer bins for several variables than has previously been possible.  The \\maxbcg\\ selection function, which is well-understood for most of the richness range used here, is less well quantified for clusters with \\Ntwo\\ $< 10$. This low richness set of clusters may be less complete and pure, which could result in the significant difference in LF shape as compared to higher \\Ntwo\\ systems, or contribute to the drop-off in \\fred\\ for low richness systems. However, not all cluster galaxy properties change significantly at \\Ntwo\\ $= 10$ (\\eg, \\Lbcg/\\Ltwo) and there is no break at a particular \\Ntwo\\ in either the mass--observable relationship \\citep{Johnston07b} or mass-to-light ratios \\citep{Sheldon07b} as measured by lensing. Interestingly, it is the quantities that most closely trace the total mass of the systems (\\ie, \\Lbcg, \\Ltwo and the lensing signal) that are smoothly scaling over the full richness range, while quantities that are related to the mix of galaxies within clusters (\\ie, the LF and \\fred) are the ones that change more dramatically. We hypothesize that, at these low richnesses, \\maxbcg\\ is finding legitimate low-mass systems, but that the selection priors demanding the close proximity of only a few red galaxies result in finding only systems with the observed mix of galaxies and not necessarily {\\em all} systems of this low mass. Further investigation using lower mass threshold mock catalogs is needed to understand in detail the selection of low-richness systems.  Our results can shed light both on the processes that build up the galaxy population in clusters and distinguish central galaxies from satellites, as well as the processes that are responsible for the galaxy transformation from blue to red.  With respect to the former, the results presented here fit well within the basic picture of galaxy formation in CDM: that galaxy properties are likely linked to the formation history of a cluster's dark matter halo and its substructures.  Although detailed comparisons are beyond the scope of this work, our results qualitatively match both HOD constraints from clustering statistics as well as models based on matching the abundance of halos and subhalos to galaxies.  The HOD framework provides a way to examine the galaxy population in both the observed Universe and in models of galaxy formation and evolution. Without needing to identify specific groups or clusters in the data, halo model interpretations of the statistics of luminosity-dependent galaxy clustering result in specific predictions for trends of both central and satellite galaxies as a function of halo mass that are in reasonable agreement with the findings presented here, especially for the relationship between central and satellite luminosities \\citep[\\eg,][]{Berlind03,Skibba07}.  Models which link the properties of galaxies directly to their halos and subhalos also agree broadly with several of the results presented here \\citep[\\eg,][]{ Conroy06, VO07}.  Detailed predictions for several cluster statistics from such a model will be presented in \\citet{Iro07}.  Our results on scatter in the BCG luminosity at fixed cluster richness very likely provide an upper limit on scatter in central galaxy luminosity at fixed mass, as the scatter between halo mass and cluster richness should act to increase this scatter.  The fact that this scatter is already fairly small provides further support for the tight coupling between halo mass and galaxy luminosity that is the basis of these models.  In addition to processes that cause physical changes to satellites resulting from interactions with the cluster gas, cluster potential, or other satellites, presumably some of the satellites are lost due to being accreted onto the BCG. In detail, this process likely results in the stellar component of the disrupted galaxy joining both the BCG and ICL \\citep{Conroy07ICL}, but nonetheless should result in a BCG population closely linked to both halo mass and satellite population, as is observed. Indeed, \\Lbcg/\\Ltwo\\ and \\Lbcg/\\Lstarsat\\ must be intimately related to the processes responsible for BCG growth. That BCGs get brighter as a function of cluster mass faster than do typical satellies may be further evidence that BCGs are different (in their merger history) than typical satellites.   Understanding the timescales and mechanisms for galaxies to transform from star-forming galaxies onto the red sequence is one of the primary current challenges for galaxy formation theories.  There are several processes that can operate within clusters to shape the population of the cluster galaxies, such as ram pressure stripping, harassment and strangulation, that directly influence the galaxies' gas content and thus their subsequent star formation \\citep[for a recent review, see][]{DeLucia06}. Clearly many of the processes responsible for this transformation are related either to the mass of the host halo or how long the galaxies have been satellites.  The relative strengths of various effects are still rather uncertain, however.  What fraction of galaxies become red while they are central galaxies?  Do the processes happen only for satellites in a certain mass range, or do they happen equally for all satellite galaxies?  On what timescales do these processes operate?  Our measurements of how the red fraction scales with cluster mass, radius and redshift will be instrumental in answering these questions.  We address a few of these issues here.  Ram pressure stripping predicts that \\fred\\ will be inversely correlated with cluster-centric distance, and larger for both brighter galaxies and more massive halos. However, our results indicate that \\fred\\ is essentially independent of galaxy luminosity at fixed \\Ntwo\\ for galaxies brighter than \\Lstar\\ and that the radial trend in \\fred\\ is not any more pronounced in high \\Ntwo\\ systems. These results indicate that ram pressure stripping cannot be the dominant process at work to transform the galaxy population.  The harassment scenario predicts that \\fred\\ will be anticorrelated with cluster mass, as is observed; however, if this mechanism were dominant, then at fixed cluster mass \\fred\\ would be expected to be larger for less luminous galaxies \\citep{Weinmann06a}, contradicting the observed trends. Strangulation, where star formation in infalling galaxies is halted because no further gas accretion is allowed, makes several predictions that are in good agreement with our observations.  For example, the model presented by \\citet{Diaferio01} predicts: that the mean satellite color gets bluer as a function of cluster-centric distance until reaching a plateau at the field value around 2--3\\rtwo; that the mean satellite color depends on halo mass only for $M \\lae 5 \\times 10^{13}$\\Msun; and that the incidence of blue galaxies in clusters increases at higher redshift. However, this model also predicts that both bulge- and disk-dominated satellies will get redder toward the central regions of clusters, a trend that we do not see in our red/blue sample split (or course, our color-separated subsamples are not directly comparable to their morphology-separated samples).  Note that the predictions of such a model will in detail depend on its implementation.  Our observation that the satellite red fraction changes in the same way with redshift regardless of cluster richness is significant evidence that the timescale of the physics responsible for quenching is the same in systems over the full mass range examined. The difference in cosmic time between the median redshifts of our two samples is approximately the dynamical timescale, over which we observe a $\\sim$5\\% change in \\fred, and this observation should allow for useful constraints in studies on the details of quenching mechanisms.  A consensus is emerging from a variety of modeling efforts that star formation proceeds most efficiently in a mass range around $L_*$, and is less efficient in more massive halos and for satellite galaxies. Simple models based on this type of assumption can produce many of the rough trends that we have seen here; for example, that the mean color of galaxies will not change significantly as a function of halo mass for the range of masses we investigate here \\citep[\\eg,][]{Diaferio01}, and that the galaxy type fraction should be a weak function of halo mass for $M$ \\gae $10^{13}$\\Msun\\ \\citep{Berlind03,Zheng05,Cooray05}. Furthermore, the general concept of galaxies falling in, quenching, and fading matches well with the observed radial trends of LF and red fraction.  However, most of the detailed semi-analytic modeling efforts have had some trouble matching in detail observations of the color distribution of galaxies and how it changes with environment \\citep[see \\eg,][]{Coil07}.  Unresolved issues include the rates at which star formation is triggered or shut off, and accordingly red galaxies tend to be overproduced.  To understand in detail the physical mechanisms responsible for the quenching of star formation as clusters are assembled, further work is needed to accurately reproduce the observed trends in the cluster galaxy population.  The present findings set a local-universe target for modeling results, and provide some guidance for the relative importance of some of the germane effects.  A substantial effort over the next decade will be devoted to large-scale, multi-band photometric surveys, including DES, Pan-STARRS, and LSST. Although the primary science driver of many of these projects is to investigate the nature of dark energy, the resulting data are likely to also provide strong constraints on the processes of galaxy evolution. The results presented here provide a low-redshift baseline against which current and future high-redshift samples may be compared.  From a technical standpoint, these data are informative to next-generation cluster-finding techniques, and are useful input for creating the mock catalogs necessary for interpreting cluster surveys. Furthermore, the techniques presented here, which use photometric data alone, are directly applicable to these upcoming imaging surveys, and will thus enable detailed studies of the galaxy population at significantly higher redshifts without extensive spectroscopy."
    },
    "0710/0710.1099_arXiv.txt": {
        "abstract": "The reliability of quiet Sun magnetic field diagnostics based on the \\ion{Fe}{1} lines at 6302 \\AA \\, has been questioned by recent work. We present here the results of a thorough study of high-resolution multi-line observations taken with the new spectro-polarimeter SPINOR, comprising the 5250 and 6302 \\AA \\, spectral domains. The observations were analyzed using several inversion algorithms, including Milne-Eddington, LTE with 1 and 2 components, and MISMA codes. We find that the line-ratio technique applied to the 5250~\\AA \\, lines is not sufficiently reliable to provide a direct magnetic diagnostic in the presence of thermal fluctuations and variable line broadening. In general, one needs to resort to inversion algorithms, ideally with realistic magneto-hydrodynamical constrains. When this is done, the 5250~\\AA \\, lines do not seem to provide any significant advantage over those at 6302~\\AA . In fact, our results point towards a better performance with the latter (in the presence of turbulent line broadening). In any case, for very weak flux concentrations, neither spectral region alone provides sufficient constraints to fully disentangle the intrinsic field strengths. Instead, we advocate for a combined analysis of both spectral ranges, which yields a better determination of the quiet Sun magnetic properties. Finally, we propose the use of two other \\ion{Fe}{1} lines (at 4122 and 9000~\\AA ) with identical line opacities that seem to work much better than the others. ",
        "introduction": "\\label{sec:intro}  The empirical investigation of quiet Sun\\footnote{In this work, we use the term ``quiet Sun'' to refer to the solar surface away from sunspots and active regions.} magnetism is a very important but extremely challenging problem. A large (probably dominant) fraction of the solar magnetic flux resides in magnetic accumulations outside active regions, forming network and inter-network patches (e.g., \\citeNP{SNSA02}). It is difficult to obtain conclusive observations of these structures, mainly because of two reasons. First, the size of the magnetic concentrations is much smaller than the spatial resolution capability of modern spectro-polarimetric instrumentation. Estimates obtained with inversion codes yield typical values for the filling factor of the resolution element between $\\sim$1\\% and 30\\% . The interpretation of the polarization signal becomes non-trivial in these conditions and one needs to make use of detailed inversion codes to infer the magnetic field in the atmosphere. Second, the observed signals are extremely weak (typically below $\\sim$1\\% of the average continuum intensity), demanding both high sensitivity and high resolution. Linear polarization is rarely observed in visible lines, so one is usually left with Stokes~$I$ and~$V$ alone.  \\citeN{S73} proposed to use the pair of \\ion{Fe}{1} lines at 5247 and~5250~\\AA \\, which have very similar excitation potentials and oscillator strengths (and, therefore, very similar opacities) but different Land\\'e factors, to determine the intrinsic field strength directly from the Stokes~$V$ line ratio. That work led to the subsequent popularization of this spectral region for further studies of unresolved solar magnetic structures. Later, the pair of \\ion{Fe}{1} lines at 6302~\\AA \\, became the primary target of the Advanced Stokes Polarimeter (ASP, \\citeNP{ELT+92}), mainly due to their lower sensitivity to temperature fluctuations. The success of the ASP has contributed largely to the currently widespread use of the 6302~\\AA \\, lines by the solar community.  Recent advances in infrared spectro-polarimetric instrumentation now permit the routine observation of another very interesting pair of \\ion{Fe}{1} lines, namely those at 15648 and 15653~\\AA \\, (hereafter, the 1.56~$\\mu$m lines). Examples are the works of \\citeN{LR99}; \\citeN{KCS+03}. The large Land\\' e factors of these lines, combined with their very long wavelengths, result in an extraordinary Zeeman sensitivity. Their Stokes~$V$ profiles exhibit patterns where the $\\sigma$-components are completely split for fields stronger than $\\sim$400~G at typical photospheric conditions. They also produce stronger linear polarization. On the downside, this spectral range is accesible to very few polarimeters. Furthermore, the 1.56~$\\mu$m lines are rather weak in comparison with the above-mentioned visible lines.  Unfortunately, the picture revealed by the new infrared data often differs drastically from what was being inferred from the 6302~\\AA \\, observations (e.g., \\citeNP{LR99}; \\citeNP{KCS+03}; \\citeNP{SNSA02}; \\citeNP{SNL04}; \\citeNP{DCSAK03}), particularly in the inter-network. \\citeN{SNSA03} proposed that the discrepancy in the field strengths inferred from the visible and infrared lines may be explained by magnetic inhomogeneities within the resolution element (typically 1\\arcsec). If multiple field strengths coexist in the observed pixel, then the infrared lines will be more sensitive to the weaker fields of the distribution whereas the visible lines will provide information on the stronger fields (see also the discussion about polarimetric signal increase in the 1.56 $\\mu$m lines with weakening fields in \\citeNP{SAL00}). This conjecture has been tested recently by \\citeN{DCSAK06} who modeled simultaenous observations of visible and infrared lines using unresolved magnetic inhomogeneities.  A recent paper describing numerical simulations by \\citeN{MGCRC06} casts some doubts on the results obtained using the 6302~\\AA \\, lines. Our motivation for the present work is to resolve this issue by observing simultaneously the quiet Sun at 5250 and 6302~\\AA . We know that unresolved magnetic structure might result in different field determinations in the visible and the infrared, but the lines analyzed in this work are close enough in wavelengths and Zeeman sensitivities that one would expect to obtain the same results for both spectral regions. ",
        "conclusions": "\\label{sec:conc}  The ratio of Stokes~$V$ amplitudes at 5250 and 5247~\\AA \\, is a very good indicator of the intrinsic field strength in the absence of line broadening, e.g. due to turbulence. However, line broadening tends to smear out spectral features and reduce the Stokes~$V$ amplitudes. This reduction is not the same for both lines, depending on the profile shape. If the broadening could somehow be held constant, one would obtain a line-ratio calibration with very low scatter. However, if the broadening is allowed to fluctuate, even with amplitudes as small as 1~km~s$^{-1}$, the scatter becomes very large.  Fluctuations in the thermal conditions of the atmosphere further complicate the analysis. This paper is not intended to question the historical merits of the line-ratio technique, which led researchers to learn that fields seen in the quiet Sun at low spatial resolution are mostly of kG strength with small filling factors. However, it is important to know its limitations. Otherwise, the interpretation of data such as those in Figure~\\ref{fig:mapratios} could be misleading. Before this work, most of the authors were under the impression that measuring the line ratio of the 5250~\\AA \\, lines would always provide an accurate determination of the intrinsic field strength.  With very high-resolution observations, such as those expected from the Advanced Technology Solar Telescope (ATST, \\citeNP{KRK+03}) or the Hinode satellite, there is some hope that most of the turbulent velocity fields may be resolved. In that case, the turbulent broadening would be negligible and the line-ratio technique would be more robust. However, even with the highest possible spatial resolution, velocity and temperature fluctuations along the line of sight will still produce turbulent broadening.  From the study presented here we conclude that, away from active region flux concentrations, it is not straightforward to measure intrinsic field strengths from either 5250 or 6302~\\AA \\, observations taken separately. Weak-flux internetwork observations would be even more challenging, as demonstrated recently by \\citeN{MG07}. Surprisingly enough, the 6302~\\AA \\, pair of \\ion{Fe}{1} lines is more robust than the 5250~\\AA \\, lines in the sense that it is indeed possible to discriminate between weak and strong field solutions if one is able to rule out a thermal stratification with temperatures that increase outwards. Even so, this is only possible when one employs an inversion code that has sufficient MHD constrains (an example is the MISMA implementation used here) to reduce the space of possible solutions.  The longitudinal flux density obtained from inversions of the 6302~\\AA \\, lines is better determined than those obtained with 5250~\\AA . This happens regardless of the inversion method employed, although using a code like LILIA provides better results than a simpler one such as MELANIE. The best fits to average network profiles correspond to strong kG fields, as one would expect.  An interesting conclusion of this study is that it is possible to obtain reliable results by inverting simultaneous observations at both 5250 and 6302~\\AA . Obviously this would be possible with relatively sophisticated algorithms (e.g., LTE inversions) but not with simple Milne-Eddington inversions.  The combination of two other \\ion{Fe}{1} lines, namely those at 4122 and 9000~\\AA , seems to provide a much more robust determination of the quiet Sun magnetic fields. Unfortunately, these lines are very distant in wavelength and few spectro-polarimeters are capable of observing them simultaneously. Examples of instrument with this capability are the currently operational SPINOR and THEMIS, as well as the planned ATST and GREGOR. Depending on the evolution time scales of the structures analyzed it may be possible for some other instruments to observe the blue and red lines alternatively."
    },
    "0710/0710.3325_arXiv.txt": {
        "abstract": "{} {In this paper we study the possibility of testing $CPT$ symmetry with Cosmic Microwave Background (CMB) measurements.} {Working with an effective lagrangian of the photon with $CPT$ violation ${\\cal L} \\sim p_{\\mu}A_{\\nu}\\tilde F^{\\mu\\nu}$ which causes the polarization vectors of the propagating CMB photons rotated, we determine the rotation angle $\\Delta\\alpha$ using the BOOMERanG 2003 and the WMAP3 angular power spectra.} {In this analysis we have included the newly released $TC$ and $GC$ ($l<450$) information of WMAP3 and found $\\Delta\\alpha=-6.2\\pm3.8$ deg at $68\\%$ confidence level.} {This result increases slightly the significance for the $CPT$ violation obtained in our previous paper (Feng et al. 2006) $\\Delta\\alpha=-6.0 \\pm 4.0$ deg (1$\\sigma$). Furthermore we examine the constraint on the rotation angle from the simulated polarization data with Planck precision. Our results show that the future Planck measurement will be sensitive to $\\Delta \\alpha$ at the level of $0.057$ deg and able to test the $CPT$ symmetry with a higher precision.} ",
        "introduction": "\\label{Introduction}  In the standard model of particle physics $CPT$ is a fundamental symmetry. Probing its violation is an important way to search for the new physics beyond the standard model. Up to now, $CPT$ symmetry has passed a number of high precision experimental tests and no definite signal of its violation has been observed in the laboratory. So, the present $CPT$ violating effects, if exist, should be very small to be amenable to the experimental limits.  The $CPT$ symmetry could be dynamically violated in the expanding universe \\cite{Li:2007}. The cosmological $CPT$ violation mechanism investigated in the literature \\cite{Li:2001st,Li:2002wd,Li:2004hh,Feng:2004mq,Li:2007} has an interesting feature that the $CPT$ violating effects at present time are too small to be detected by the laboratory experiments but large enough in the early universe to account for the generation of matter-antimatter asymmetry \\cite{Li:2001st,Li:2002wd,Li:2004hh,Li:2007}. And more importantly, this type of $CPT$ violating effects could be accumulated to be observable in the cosmological experiments \\cite{Feng:2004mq,Li:2007,Feng:2006dp}. With the accumulation of high quality observational data, especially those from the CMB experiments, cosmological observation becomes a powerful way to test the $CPT$ symmetry.  Here we study the CMB polarizations and $CPT$ violation in the photon sector with an effective lagrangian \\cite{Carroll:1989vb,Carroll:1990zs}: \\begin{equation}\\label{Lagrangian} \\mathcal{L} = -\\frac{1}{4}F_{\\mu\\nu}F^{\\mu\\nu}+\\mathcal{L}_{cs}~, \\end{equation} where $\\mathcal{L}_{cs}\\sim p_{\\mu}A_{\\nu}\\tilde F^{\\mu\\nu}$ is a Chern-Simons term, $p_{\\mu}$ is an external vector and $\\tilde F^{\\mu\\nu}=(1/2)\\epsilon^{\\mu\\nu\\rho\\sigma}F_{\\rho\\sigma}$ is the dual of the electromagnetic tensor. This Lagrangian is not gauge invariant, but the action is gauge independent if $\\partial_{\\nu}p_{\\mu}=\\partial_{\\mu}p_{\\nu}$. This may be possible if $p_{\\mu}$ is constant in spacetime or the gradient of a scalar field in the quintessential baryo-/leptogenesis \\cite{Li:2002wd,Li:2001st,quin_baryogenesis} or the gradient of a function of the Ricci scalar in gravitational baryo-/leptogenesis \\cite{Li:2004hh,R}. The Chern-Simons term violates Lorentz and $CPT$ symmetries when the background value of $p_{\\mu}$ is nonzero.  One of the physical consequences of the Chern-Simons term is the rotation of the polarization direction of electromagnetic waves propagating over large distances \\cite{Carroll:1989vb}. From the Lagrangian (\\ref{Lagrangian}), we can directly obtain the equation of motion for the electromagnetic field: \\begin{equation}\\label{maxwell1} \\nabla_{\\mu}(\\nabla^{\\mu}A^{\\nu}-\\nabla^{\\nu}A^{\\mu})=-p_{\\mu} \\epsilon^{\\mu\\nu\\rho\\sigma}(\\nabla_{\\rho}A_{\\sigma}-\\nabla_{\\sigma}A_{\\rho})~. \\end{equation} After imposing Lorentz gauge condition $\\nabla_{\\mu}A^{\\mu}=0$, it becomes: \\begin{equation}\\label{maxwell2} \\nabla_{\\mu}\\nabla^{\\mu}A^{\\nu}+R^{\\nu}_{\\mu}A^{\\mu}=-p_{\\mu} \\epsilon^{\\mu\\nu\\rho\\sigma}(\\nabla_{\\rho}A_{\\sigma}-\\nabla_{\\sigma}A_{\\rho})~, \\end{equation} where $R^{\\nu}_{\\mu}$ is the Ricci tensor. With the geometric optics approximation, the solution to the equation of motion is expected to be: $A^{\\mu}={\\rm Re}[(a^{\\mu}+\\epsilon b^{\\mu}+\\epsilon^2 c^{\\mu}+...)e^{iS/\\epsilon}]$, where $\\epsilon$ is a small number. With this ansatz, one can easily see that the Lorentz gauge condition implies $k_{\\mu}a^{\\mu}=0$, where the wave vector $k_{\\mu}\\equiv \\nabla_{\\mu}S$ is orthogonal to the surfaces of constant phase and represents the direction which photons travel along with. The vector $a^{\\mu}$ is the product of a scalar amplitude $A$ and a normalized polarization vector $\\varepsilon^{\\mu}$, $a^{\\mu}=A\\varepsilon^{\\mu}$, with $\\varepsilon_{\\mu}\\varepsilon^{\\mu}=1$. Hence in the Lorentz gauge, the wave vector $k_{\\mu}$ is orthogonal to the polarization vector $\\varepsilon^{\\mu}$. Substituting this solution into the modified Maxwell equation (\\ref{maxwell2}) and neglecting the Ricci tensor we have at the leading order of $\\epsilon$ the equation is $k_{\\mu}k^{\\mu}=0$. It indicates that photons still propagate along the null geodesics. The effect of Chern-Simons term appears at the next order, $k^{\\mu}\\nabla_{\\mu}\\varepsilon^{\\nu}=-p_{\\mu}\\epsilon^{\\mu\\nu\\rho\\sigma} k_{\\rho}\\varepsilon_{\\sigma}$. We can see that the Chern-Simons term makes $k^{\\mu}\\nabla_{\\mu}\\varepsilon^{\\nu}$ not vanished. This means that the polarization vector $\\varepsilon^{\\nu}$ is not parallel transported along the light-ray. It rotates as the photon propagates in spacetime.  We consider here the spacetime described by spatially flat Friedmann-Robertson-Walker (FRW) metric. The null geodesics equation is $(k^0)^2-k^ik^i=0$. We assume that photons propagate along the positive direction of $x$ axis, \\emph{i.e.} $k^{\\mu}=(k^0,k^1,0,0)$ and $k^1=k^0$. Gauge invariance guarantees that the polarization vector of the photon has only two independent components which are orthogonal to the propagating direction. So, we are only interested in the changes of the components of the polarization vector, $\\varepsilon^2$ and $\\varepsilon^3$. Assuming $p_{\\mu}=p_0$ to be a non-vanishing constant, we obtain the following equations: $d\\varepsilon^2/d\\lambda+\\mathcal{H}k^0\\varepsilon^2=p_0 k^0\\varepsilon^3$, $d\\varepsilon^3/d\\lambda+\\mathcal{H}k^0\\varepsilon^3= - p_0 k^0\\varepsilon^2$, where we have defined the affine parameter $\\lambda$ which measures the distance along the light-ray, $k^{\\mu}\\equiv dx^{\\mu}/d\\lambda$, and the reduced expansion rate $\\mathcal{H}\\equiv \\dot a/a$. The polarization angle is defined as $\\chi\\equiv\\arctan{(\\varepsilon^3}/{\\varepsilon^2})$. It is easy to find that the rotation angle is \\begin{equation}\\label{kmu1} \\Delta\\chi\\equiv \\chi_0-\\chi_z=-\\int^{\\eta_0}_{\\eta_z} ~p_0 ~d\\eta=\\int^{t_z}_{t_0} p_0 ~\\frac{dt}{a}~, \\end{equation} where the subscript $z$ is the redshift of the source when the light was emitted. For CMB photons, the source is the last scattering surface with $z\\simeq 1100$ \\footnote{Besides the CMB photons which come from the last scattering surface, we might observed the different CMB photons which travelled different distances as well. However, we are only interested in linear perturbations of CMB photons in this paper. CMB polarizations are already linear phenomena. They are not existent at the zeroth order. When calculating the variations of polarizations due to $CPT$ violation in the perturbation theory up to linear order, we may ignore the fluctuations of the travelling distances of CMB photons. Otherwise, these fluctuations combined with polarizations would give higher order result which are beyond the scope of this paper. So, in this paper, we only consider linear perturbations and we can assume that each CMB photon detected by us travelled the same distance.}. The subscript $0$ indicates the present time. As we know, a vector rotated by an angle $\\Delta\\chi$ in a fixed coordinates frame is equivalent to a fixed vector observed in a coordinates frame which is rotated by $-\\Delta\\chi$. So, with the notion of coordinates frame rotation, the rotation angle is \\begin{equation}\\label{kmu2} \\Delta\\alpha=-\\Delta\\chi=\\int^{t_0}_{t_z} ~ p_0 ~ \\frac{dt}{a}~ = p_0 r_z~. \\end{equation} with $r_z$ being the comoving distance of the light source away from us. This phenomena is known as ``cosmological birefringence\". This rotation angle $\\Delta\\alpha$ can be obtained by observing polarized radiation from distant sources such as radio galaxies, quasars and CMB.  The Stokes parameters $Q$ and $U$ of the CMB polarization can be decomposed into a gradient-like ($G$) and a curl-like ($C$) component \\cite{Kamionkowski:1996ks}. For the standard theory of CMB, the $TC$ and $GC$ cross-correlation power spectra vanish. With the existence of cosmological birefringence, the polarization vector of each photon is rotated by an angle $\\Delta\\alpha$, and one would observe nonzero $TC$ and $GC$ correlations, even if they are zero at the last scattering surface. Denoting the rotated quantities with a prime, one gets \\cite{Feng:2004mq,Lue:1998mq}: \\begin{eqnarray}\\label{modify} C_{l}^{'TC} &=& C_{l}^{TG}\\sin(2\\Delta\\alpha)~, \\nonumber\\\\ C_{l}^{'GC} &=& \\frac{1}{2}(C_{l}^{GG}-C_{l}^{CC})\\sin(4\\Delta\\alpha)~,\\nonumber\\\\ C_{l}^{'TG} &=& C_{l}^{TG}\\cos(2\\Delta\\alpha)~,\\nonumber\\\\ C_{l}^{'GG} &=& C_{l}^{GG}\\cos^2(2\\Delta\\alpha) + C_{l}^{CC}\\sin^2(2\\Delta\\alpha)~,\\nonumber\\\\ C_{l}^{'CC} &=& C_{l}^{CC}\\cos^2(2\\Delta\\alpha) + C_{l}^{GG}\\sin^2(2\\Delta\\alpha)~, \\end{eqnarray} while the temperature power spectrum $TT$ remains unchanged. ",
        "conclusions": ""
    },
    "0710/0710.3113_arXiv.txt": {
        "abstract": "{In April $2006$ a $4$-channel acoustic antenna has been put in long-term operation on Lake Baikal. The detector was installed  at a depth of about $100$ m on the instrumentation string of Baikal Neutrino Telescope NT200+. This detector may be regarded as a prototype of subunit for a future underwater acoustic neutrino telescope. We describe the design of acoustic detector and present first results obtained from data analysis.}    \\begin{document} ",
        "introduction": "\\begin{figure*}[t] \\begin{center} \\includegraphics [width=0.98\\textwidth]{icrc0639_fig01.eps} \\end{center} \\caption{Schematic view of  underwater 4-channel digital device for detection of acoustic signals from high energy neutrinos.} \\label{fig1} \\end{figure*}  The large scale neutrino telescopes currently under operation (NT200+ in Lake Baikal, AMANDA/IceCube at the South Pole and ANTARES in the Mediterranean) detect the Cherenkov light emitted in water or ice by relativistic charged particles produced via neutrino interactions with matter. Back in 1957, G.A. Askaryan has shown that a high-energy particle cascade in water should also produce an acoustic signal \\cite{Askarian-1957}. The absorption length for acoustic waves with a frequency about 30 kHz (the peak frequency of acoustic signals from a shower) in sea water is at least an order of magnitude larger than that of Cherenkov radiation, in the fresh Baikal water this ratio is even close to 100 \\cite{Clay}. Therefore acoustic pulses can be detected from considerably larger distances than Cherenkov radiation, and the acoustic method appears to be attractive for the detection of ultra high-energy neutrinos \\cite{Askarian-1977}. However, the technology of acoustic detection in high-energy physics is much worse developed than optical methods. Since several years, however, an increasing number of feasibility studies on acoustic particle detection are performed \\cite{ARENA}.  In order to test the possibility of acoustic detection of high-energy neutrinos in Lake Baikal, the Baikal collaboration started with an in-situ study of acoustic noise which constitutes the background for the acoustic neutrino detection in the lake. For the purpose of noise measurement, an autonomous hydro-acoustic recorder with two input channels has been developed. We have performed a series of hydro-acoustic measurements in Lake Baikal in order to investigate the the background properties  \\cite{Zeuthen-1, Akustika-noise}. It turned out that at stationary and homogeneous meteorological conditions the integral noise power in the frequency range $20$-$50$ kHz can reach levels as low as about $1$ mPa. At the same time, short acoustic pulses with different amplitudes and shapes including bipolar ones have been observed. The latter should be considered as a background for acoustic neutrino detection. However,  the overwhelming majority of the short pulses have probably been generated by quasi-local sources or are due to interference of noise sound waves coming from a layer near the surface.  Taking into account these properties of the noise, we conclude that the most promising way to detect acoustic particle signals is to deploy a net of rather compact acoustic antennas at relatively shallow depths (for example about $100$--$200$ m for Lake Baikal) and monitoring the water volume top-down. It is also necessary to suppress signals from the surface by caps made of a sound-absorbing material and mounted on top of the antennas. ",
        "conclusions": "The results of the experiment have demonstrated the feasibility of the proposed acoustic pulse detection technique in searching signals from cascade showers. Although the Baikal water temperature is close to the temperature of its maximal density, the absence of strong acoustic noise sources in the lake's deep zone, and the very low absorption of sound in freshwater may result in neutrino detection in Lake Baikal with a threshold as low as $10^{18}$ -- $10^{19}$ eV.  This motivates further activities towards a large-scale acoustic neutrino detector in Lake Baikal."
    },
    "0710/0710.5503_arXiv.txt": {
        "abstract": "We present relations of the black hole mass and the optical luminosity with the velocity dispersion and the luminosity of the \\nev\\ and the \\oiv\\ high-ionization lines in the mid-infrared (MIR) for 28 reverberation-mapped active galactic nuclei. We used high-resolution \\spi\\ Infrared Spectrograph and {\\it Infrared Space Observatory} Short Wavelength Spectrometer data to fit the profiles of these MIR emission lines that originate from the narrow-line region of the nucleus. We find that the lines are often resolved and that the velocity dispersion of \\nev\\ and \\oiv\\ follows a relation similar to that between the black hole mass and the bulge stellar velocity dispersion found for local galaxies. The luminosity of the \\nev\\ and the \\oiv\\ lines in these sources is correlated with that of the optical 5100$\\ang$ continuum and with the black hole mass. Our results provide a means to derive black hole properties in various types of active galactic nuclei, including highly obscured systems. ",
        "introduction": "\\label{sec:intro} Following the relation between the black hole (BH) mass, \\mbh, and the bulge stellar velocity dispersion, $\\sigma_*$,  (\\citealt{ferrarese}; \\citealt{gebhardt}; \\citealt{tremaine}), a similar relation was found for the velocity dispersion $\\sigma$ of the \\oiii\\ 5007\\ang\\ line (\\citealt{nelson00}; \\citealt{greene}) that originates from the narrow-line region (NLR) gas of the active galactic nucleus (AGN). The NLR gas is thought to be mostly gravitationally bound to the bulge in high-luminosity AGNs (\\citealt{nelson96}; \\citealt{greene}; \\citealt{laor07}), unlike the broad-line region gas that is virialized due its proximity to the BH (\\citealt{peterson99}). This relation is useful for systems in which the stellar absorption lines cannot be observed, e.g., because of dilution of the stellar light by the AGN continuum. In addition to this relation, \\mbh\\ and the luminosity $L$ of the 5100 \\ang\\ optical continuum, \\lopt , are correlated in AGNs with BH measurements (\\citealt{kaspi00}).  There are two main advantages to expanding such relations in the mid-infrared (MIR). The first is that the \\ion{Ne}{5} and \\ion{O}{4} ions emitting at 14.32 and 25.89 \\micron\\ cannot be easily excited by star-forming regions since they have ionization potentials $\\chi$ of 97.12 and 54.93 eV. The second reason is the low obscuration, which allows for the results to be applied to type 2 AGNs. In this Letter, we investigate for relations between \\mbh\\ and \\lopt\\ with MIR NLR line velocity dispersions and luminosities. ",
        "conclusions": "\\label{sec:origin} The relations between \\snev,\\soiv\\ and \\mbh\\ indicate that the NLR gas kinematics are primarily determined by the potential of the bulge, as it is also believed based on the \\oiii\\ line profiles (\\citealt{whittle92}; \\citealt{nelson96}; \\citealt{greene}). However, \\snev\\ and \\soiv\\ are on average 67 and 62 \\kms\\ higher than $\\sigma_*$, and 51 and 32 \\kms\\ higher than \\soiii . Such discrepancies could be attributed to an increase of the line width with increasing $\\chi$, which is often found for optical NLR lines (e.g. \\citealt{osterbrock}; \\citealt{whittle85b}). It is possible that ions with high ionization potentials have high velocity dispersions because they are located in NLR clouds that are close to the BH sphere of influence. Since the \\ion{O}{4} and, mostly, the \\ion{Ne}{5} ions are predominantly excited by AGNs, the MIR \\msigma\\ relation can be applied even in systems that undergo starbursts. It can also be applied in type 2 AGNs since the lines suffer little from extinction. Moreover, \\snev\\ and \\soiv\\ can be used as surrogates for $\\sigma_*$ in environments where measuring $\\sigma_*$ is hard, such as in bright QSOs. However, the relation does not necessarily hold for AGNs with strong winds and jets such as luminous radio sources (\\citealt{whittle92}), and for AGNs with high Eddington rates, \\eedd\\ (\\citealt{greene}). Such systems can be recognized by the asymmetric wings or the high kurtosis of their line profiles (Whittle 1985a; 1992).  The luminosities of the broad lines (\\citealt{baldwin}) and the narrow lines (Fig.~\\ref{fig:ff}) scale with \\lopt\\ and with the bolometric AGN luminosity, \\lbol , since they depend on the absolute accretion rate onto the central object. Specifically, \\lbol\\ is equal to \\xopt \\lopt, \\xnev \\lnev, and \\xoiv \\loiv , where $x$ is a wavelength-dependent bolometric correction factor. For \\xopt $=$9 (\\citealt{kaspi00}; \\citealt{marconi04}) and for the median values of \\lopt/\\lnev\\ and \\lopt/\\loiv\\ in our sample, we find that \\xnev $=$13000 and \\xoiv $=$4500. That the relation does not have a slope of unity could indicate that \\xnev\\ and \\xoiv\\ depend on $L$ in a manner different than \\xopt\\ does. \\cite{netzer06} found that the equivalent width of \\oiii\\ decreases with increasing \\lopt. This behaviour could be attributed to a different geometric distribution of NLR clouds around the AGN at different luminosities.  Whether the \\nev\\ or \\oiv\\ luminosity of a source can be used to derive its \\mbh\\ partly depends on its Eddington rate. To illustrate how changes of \\eedd\\ affect the relation between the luminosity of a MIR line and \\mbh , we overplot lines of constant \\eedd\\ in Figure~\\ref{fig:ml}. The more quiescent a source is, the more its position is shifted to the top left corner of this diagram. Equation~(\\ref{rel3}) is valid within the \\eedd\\ range 0.003$-$0.6 that our sources span, and implies that the mass of a BH determines its luminosity output. Its advantage is that it can be applied to sources with obscured optical continua or underestimated \\oiii\\ luminosities, such as type 2 AGNs (\\citealt{netzer06}). However its scatter is likely to be larger than that presented in \\S~\\ref{sec:mir_opt} since the Eddington rates of the reverberation-mapped AGNs are not necessarily representative of those of all local AGNs.  We conclude that the MIR NLR gas kinematics trace \\mbh\\ in local reverberation-mapped AGNs in a manner similar to stellar kinematics. The \\nev\\ and \\oiv\\ line widths can be used to estimate \\mbh. The calibration of \\lnev\\ and \\loiv\\ to \\lopt\\ provides a new method to compute bolometric luminosities, black hole masses, and Eddington rates in various types of AGNs, including highly obscured systems, albeit with a large scatter.\\\\ \\\\ This work was based on observations made with the {\\it Spitzer Space Telescope}, % and was supported by NASA through an award issued by JPL/Caltech."
    },
    "0710/0710.2115_arXiv.txt": {
        "abstract": "With a goal toward deriving the physical conditions in external galaxies, we present a survey of the formaldehyde emission in a sample of starburst systems.  By extending a technique used to derive the spatial density in star formation regions in our own Galaxy, we show how the relative intensity of the $1_{10}-1_{11}$ and $2_{11}-2_{12}$ K-doublet transitions of H$_2$CO can provide an accurate densitometer for the active star formation environments found in starburst galaxies.  Relying upon an assumed kinetic temperature and co-spatial emission and absorption from both H$_2$CO transitions, our technique is applied to a sample of nineteen IR-bright galaxies which exhibit various forms of starburst activity.  In the five galaxies of our sample where both H$_2$CO transitions were detected we have derived spatial densities. We also use H$_2$CO to estimate the dense gas mass in our starburst galaxy sample, finding similar mass estimates for the dense gas forming stars in these objects as derived using other dense gas tracers.  A related trend can be seen when one compares $L_{IR}$ to our derived $n(H_2)$ for the five galaxies within which we have derived spatial densities.  Even though our number statistics are small, there appears to be a trend toward higher spatial density for galaxies with higher infrared luminosity.  This is likely another representation of the $L_{IR}$-$M_{dense}$ correlation. ",
        "introduction": "\\label{intro}  Studies of the distribution of Carbon Monoxide (CO) emission in external galaxies (\\cf\\ \\cite{Young1991}) have pointed to the presence of large quantities of molecular material in these systems.  These studies have yielded a detailed picture of the molecular mass in many external galaxies. But, because emission from the abundant CO molecule is generally dominated by radiative transfer effects, such as high optical depth, it is not a reliable monitor of the physical conditions, such as spatial density and kinetic temperature, quantities necessary to assess the possibility of star formation.  Emission from less-abundant, higher-dipole moment molecules are better-suited to the task of deriving the spatial density and kinetic temperature of the dense gas in our and external galaxies.  For this reason, emission line studies from a variety of molecules have been made toward mainly nearby galaxies (see \\cite{Mauersberger1989} (CS), \\cite{Gao2004a} (HCN), \\cite{Nguyen1992} (HCO$^+$), \\cite{Mauersberger1990} and \\cite{Meier2005} (HC$_3$N), \\cite{Mauersberger2003} (NH$_3$), or \\cite{Henkel1991} for a review).  The most extensive sets of measurements of molecular line emission in external galaxies has been done using the J=1-0 transitions of CO \\citep{Helfer2003} and HCN \\citep{Gao2004a}.  Since the J=1-0 transitions of CO and HCN are good tracers of the more generally distributed and the denser gas, respectively, but do not provide comprehensive information about the individual physical conditions of the dense, potentially star-forming gas, another molecule must be observed for this purpose. Formaldehyde (H$_2$CO) has proven to be a reliable density and kinetic temperature probe in Galactic molecular clouds. Existing measurements of the H$_2$CO $1_{10}-1_{11}$ and $2_{11}-2_{12}$ emission in a wide variety of galaxies by \\cite{Baan1986}, \\cite{Baan1990}, \\cite{Baan1993}, and \\cite{Araya2004} have mainly concentrated on measurements of the $1_{10}-1_{11}$ transition.  One of our goals with the present study was to obtain a uniform set of measurements of both K-doublet transitions with which the physical conditions, specifically the spatial density, in the extragalactic context could be derived. Using the unique density selectivity of the K-doublet transitions of H$_2$CO we have measured the spatial density in a sample of galaxies exhibiting starburst phenomena and/or high infrared luminosity.  In \\S\\ref{H2coProbe} we discuss the specific properties of the H$_2$CO molecule which make it a good probe of spatial density. \\S\\ref{Observations} presents our observation summary; \\S\\ref{Results} our H$_2$CO, OH, H111$\\alpha$, and continuum emission measurement results; \\S\\ref{Analysis} analyses of our H$_2$CO, OH, and H111$\\alpha$ measurements, including Large Velocity Gradient (LVG) model fits to and dense gas mass calculations based on our H$_2$CO measurements. ",
        "conclusions": "\\label{Conclusions}  Using measurements of the $1_{10}-1_{11}$ and $2_{11}-2_{12}$ K-doublet transitions of H$_2$CO we have derived accurate \\textit{measurements} of the spatial density (n(H$_2$)) in a sample of starburst galaxies.  The derived densities range from $10^{4.7}$-$10^{5.7}$~cm$^{-3}$, consistent with the suggestion that the high infrared brightness of these galaxies is driven by extreme star formation activity.  We believe that these spatial density measurements are the most accurate measurements of this important physical quantity in starburst galaxies made to-date.  We have also used our H$_2$CO measurements to derive a measure of the dense gas mass which ranges from $0.06-77\\times10^8 M_\\odot$, generally consistent with previous measurements, mainly from studies of the HCN emission in these galaxies.  The linear correlation between the IR luminosity ($L_{IR}$) and the dense gas mass ($M_{dense}$) found in larger samples of the HCN emission in starburst galaxies is also apparent in our H$_2$CO measurements.  This further supports the suggestion that active star formation in IR-bright galaxies is driven by the amount of material available to form stars.  We also note a related trend between $L_{IR}$ and our derived $n(H_2)$ for the five galaxies within which we have derived spatial densities.  The three galaxies with lower IR luminosities have lower ($n(H_2) \\simeq 10^5$~cm$^{-3}$) derived spatial densities, while the two higher luminosity galaxies have higher ($n(H_2) = 10^{5.7}$~cm$^{-3}$) derived spatial densities. This is likely another representation of the $L_{IR}$-$M_{dense}$ correlation."
    },
    "0710/0710.0867_arXiv.txt": {
        "abstract": "The observed cosmic acceleration presents the physics and cosmology communities with amazing opportunities to make exciting, probably even radical advances in these fields.  This topic is highly data driven and many of our opportunities depend on us undertaking an ambitious observational program. Here I outline the case for such a program based on both the exciting science related to the cosmic acceleration and the impressive impact that a strong observational program would have. Along the way, I challenge a number of arguments that skeptics use to question the value of a strong observational commitment to this field. ",
        "introduction": "This is truly remarkable time to be involved in cosmology research.  There are many reasons for this, but none stand out quite as dramatically as the observed cosmic acceleration (attributed in current nomenclature to the ``dark energy'').  In the words of the Dark Energy Task Force (DETF)~\\cite{Albrecht:2006um} ``most experts believe that nothing short of a revolution in our understanding of fundamental physics will be required to achieve a full understanding of the cosmic acceleration''. As things stand, this revolution is being motivated both by remarkable new data sets and by exciting theoretical developments.  Many ambitious researchers in cosmology and related fields have been galvanized by these extraordinary developments.  There have been a number of excellent talks at this conference on the topic of cosmic acceleration which capture some of this climate.  The DETF, charged with charting a way forward with dark energy observations, received 50 thoughtful and thorough whitepapers from leaders in the field (despite the lack of any specific commitment at that time to fund future dark energy experiments).  Most of us look at these developments and are astonished at our good fortune to be part of what will surely be viewed as one of the great moments in the history of science.  Indeed, since the discovery of dark energy every group that has deliberated on future directions for the field has recognized the exciting opportunities and challenges presented by the cosmic acceleration\\footnote{See for example \\cite{Turner:2003pe,Albrecht:2005np,Peacock:2006kj}. Since the 2006 DETF report a number of panels have recommended pursuit of specific ground and space based dark energy projects, and just since PASCOS 07 the US National Research Council's {\\em Committee on NASA's Beyond Einstein Program} named a ``Joint Dark Energy Mission'' the top priority for that program \\url{http://nationalacademies.org/morenews/20070907b}}. But a small number of negative voices continue to be heard alongside the building enthusiasm for further studies of the dark energy.  To some degree this negativity may reflect the natural skepticism of scientists in the face of unbridled enthusiasm (such as the above).  To the extent that that is the explanation, the skepticism is surely a good thing that will lead to a healthy debate and produce more rigorous research.  On the other hand, I suspect some of the negativity is just the result of sloppy thinking and needs to be challenged and simply put to rest.  Regardless of how one might attribute explanations and motivations, one purpose of this paper is to engage in this debate on a number of fronts.  Here are some illustrations of some of the negative views I am talking about. At lunch on the first day of the PASCOS 07 conference there was a lively discussion about the merits of dark energy experiments.  One colleague commented \\begin{quote} {\\em ``Studies of dark energy are unlikely to be interesting because we already have a theory of dark energy.''} \\end{quote} moments later, another  cosmologist declared \\begin{quote} { \\em ``Studies of dark energy are unlikely to be interesting because we have no theory of dark energy.''} \\end{quote} The two were united in their conclusion, and apparently not too bothered by subtle differences in their reasoning  One unavoidable feature of physics research is that at any given time the total amount of data is finite.  That necessarily means there will be more than one theory that fits the data.  In many fields, this universal fact is seen as a source of vitality. Curiosity about further resolving these degeneracies drives exciting new experiments and theoretical work. The momentous progress in physics over the last century can be seen in this light as the fundamental particles, the atomic theory of heat, quantum physics and general relativity emerged from ``under the radar'' of earlier data sets that were fit perfectly by more primitive theories.  As was amply evident in the talks at the PASCOS 07 conference, we have good reason to look forward to similar advances as the LHC data starts coming in.  But for some reason, the fact that a given future dark energy experiment will not remove all uncertainties about the nature of dark energy seems to generate considerable angst among some physicists and astronomers\\footnote{See the question section of my PASCOS 07 talk for an example\\cite{PASCOS07}} and is sometimes given as a reason to be discouraged from even doing more experiments. As I shall quantify below (and as was also shown by others at PASCOS 07), the proposed new experiments will have an impressive impact on our knowledge of dark energy.  This is all one can ever ask of a new experiment.  This paper focuses on two key areas.  In the next section I review some of the thriving theoretical work that has been stimulated by the cosmic acceleration. The striking theoretical issues raised by the cosmic acceleration and the remarkable directions we have been driven in our initial attempts to understand it are the key reasons I find this topic so deeply interesting. I also believe this is why so many excellent researchers are taking risks and changing direction in their careers in order to get involved.  The third section summarizes a number of results which demonstrate the tremendous impact future experiments can have on our understanding of dark energy, both on our understanding of the general properties of dark energy and in terms of constraining specific models of dark energy that are currently of interest.  It is these results that demonstrate that an ``aggressive program of dark energy probes'' is indeed possible.  This paper is based on my talk at the PASCOS 07 meeting at Imperial College.  My slides and a video of the talk are available online\\cite{PASCOS07}. This online material as well as my ``Origins of Dark Energy'' talk\\cite{ODE} and related papers\\cite{Albrecht:2007qy,Abrahamse:2007ip} are a good source for the technical material on which this paper is based.  My goal here is to assemble some key arguments in a concise form. Readers seeking more details should refer to this other material. ",
        "conclusions": "\\label{Sect:CAS}  The case for aggressive pursuit of new data on dark energy is twofold. Firstly, I have outlined how the subject of dark energy has generated very exciting and often radical new theoretical ideas.  These include dramatic proposals that change how we think about equilibrium and initial conditions in cosmology and even how we formulate fundamental theories. Second, impressive new experiments are within reach that could have a tremendous impact on our understanding of dark energy. The best experiments will constrain dark energy properties orders of magnitude better than the current or medium sized future experiments, and will have the ability to strongly discriminate among and even fully eliminate popular dark energy models based on subtle variations in the equation of state.  I have outlined and challenged some of the skeptical perspectives I have heard regarding future dark energy studies. The discovery of the cosmic acceleration has caused a great upheaval in our thinking about fundamental physics and cosmology.  I often sense that the skeptics are hoping this upheaval will end quickly and are grasping for arguments that will allow things to rapidly return to normal.  I feel this outcome is very unlikely, and this is exactly why I find the topic so exciting.  Nature has handed us an amazing opportunity. I hope that the physics and cosmology communities have the strength to face the challenge of the cosmic acceleration head on and give a response that we can be proud of when people write the history of this era.     \\begin{theacknowledgments} I would like to thank A. Abrahamse, M. Barnard, G. Bernstein, B. Bozek, L. Sorbo, M. Yashar and the DETF members who collaborated with me on some of the work reviewed here, and B. Bozek for helpful comments on the manuscript. Also, I thank the organizers, especially Arttu Rajantie, for a really excellent conference.  This work was supported in part by DOE grant DE-FG03-91ER40674 and NSF grant AST-0632901. \\end{theacknowledgments}"
    },
    "0710/0710.5359_arXiv.txt": {
        "abstract": "{% The study of rotation and activity in low-mass stars or brown dwarfs of spectral classes M and L has seen enormous progress during the last years. I summarize the results from different works that measured activity, rotation, and sometimes magnetic fields. The generation of magnetic activity seems to be unchanged at the threshold to completely convective stars, i.e. no change in the efficiency of the magnetic dynamos is observed. On the other hand, a sudden change in the strength of rotational braking appears at the threshold mass to full convection, and strong evidence exists for rotational braking weakening with lower mass. A probable explanation is that the field topology changes from dipolar to small scale structure as the objects become fully convective.  } ",
        "introduction": "Rotation and activity are intimately connected in sun-like stars. Rotation is believed to generate a magnetic field through a dynamo mechanism that scales with rotation. The stellar wind couples to the rotating magnetic field lines carrying away angular momentum so that the star is being braked. Hence stars of spectral type F--K rotate more rapidly and are more active when they are young, but they decelerate and become less active as they age; this is the so-called rotation-activity connection (Noyes et al., 1984; Pizzolato et al., 2003).  The scaling of activity with rotation depends on the type of dynamo inside the star. The fact that the rotation-activity connection is similar in virtually all stars that harbor convective envelopes and at all ages (Pizzolato et al., 2003; Reiners, 2007) indicate that the dynamo mechanism is comparable in all these stars. Around spectral type M3.5, the internal structure of the stars changes; stars later than around M3.5 are believed to be fully convective. They cannot harbor an interface dynamo working at the tachocline, which is believed to be the most important dynamo mechanism in the hotter sun-like stars. If the interface dynamo was the only important mechanism driving a magnetic dynamo, one could expect a sharp break in magnetic field generation around spectral type M3.5. Such a break would imply a sudden change in observable stellar activity and in the braking of stellar rotation.  A break in stellar activity is not observed in X-rays or H$\\alpha$ in the surveys that crossed the M3.5 border (e.g., Delfosse et al., 1998; Mohanty \\& Basri, 2003; West et al., 2004).  Activity rather stays at comparable levels if normalized to bolometric luminosity, and quite surprisingly the fraction of active stars even raises up to $\\sim 80\\,\\%$ in late M-type objects before it goes down again. This clearly shows that the interface dynamo is not the only dynamo operating in stellar interiors and that fully convective stars can have quite efficient dynamos as well.  On the other hand, Delfosse et al., 1998, presented a plot that shows a different behavior in rotational braking in stars later than spectral type M3.5. In their Fig.\\,3, they show that all M dwarfs earlier than M3.5 are slow rotators ($v\\,\\sin{i} \\la 3$\\,km\\,s$^{-1}$) regardless of what disk population they belong to (young or old). M stars later than M3.5, however, show substantial rotation velocities of up to $v\\,\\sin{i} = 50$\\,km\\,s$^{-1}$ in the young disk population, and in the old disk population some late M dwarfs still have velocities around $v\\,\\sin{i} = 10$\\,km\\,s$^{-1}$. This can be interpreted as an indication for a sudden change in the timescales of rotational braking at the mass where stars become fully convective. Delfosse et al.  conclude that spin-down timescales are on the order of a few Gyrs at spectral type M3--M4, and of the order of 10\\,Gyr at spectral type M6.  The investigation of rotation and activity in ultra-cool stars (M7 and later) was put on firm ground by Mohanty \\& Basri, 2003. From high-resolution spectra they determined projected rotation velocities, and from H$\\alpha$ emission they derived the level of activity. Reiners \\& Basri, 2007, added more M stars to this sample. In addition, they measured magnetic fields in low-mass M dwarfs and showed that in M dwarfs the level of activity is still coupled to magnetic flux -- high magnetic flux levels lead to strong H$\\alpha$ emission and stars without H$\\alpha$ emission show no magnetic fields. ",
        "conclusions": "High resolution spectroscopy in M and L dwarfs becomes a technique that allows to investigate the physics and the evolution of low mass objects in great detail. No change in H$\\alpha$ activity is observed at the threshold to complete convection. Activity can be followed to objects as late as mid-M and it might continue to even lower masses but below the current observational threshold. The decline in normalized activity among the ultra-cool dwarfs could be explained by the enhanced electrical resistivity at the low temperatures. Currently, no indications for a mass or temperature dependence of stellar or substellar dynamos can be concluded from the activity measurements.  A sudden change in the behavior of rational braking at spectral class M3.5 was already found by Delfosse et al., 1998. In the young disk, stars earlier than M3.5 rotate slowly while later stars still show significant rotation. The lack of slowly rotating ultra-cool dwarfs and the rise of the minimum rotation rate with spectral class could be explained by rotational braking that is weaker with lower mass or lower temperature. The reason for a weaker braking at later spectral type could be a different magnetic topology.  The geometry of the magnetic field is essential for the strength of magnetic braking as described in Krishnamurti et al.  (1997) and Sills et al.  (2000). These authors find that the description of magnetic braking ($\\omega_\\mathrm{crit}$) needs to be different in mid-M stars and in earlier objects, which they connect to different convective turnover times.  Although the details of magnetic braking in ultra-cool dwarfs are not quite understood, a substantial amount of evidence exists that rotational braking is weaker with lower mass or lower temperature. The limiting factor of rotational braking may be the topology of the magnetic fields which in fully convective stars might be generated on smaller scales so that the topology is different from a dipolar configuration leading to the weaker rotational braking."
    },
    "0710/0710.3396_arXiv.txt": {
        "abstract": "We modify a stellar structure code to estimate the effect upon the main sequence of the accretion of weakly interacting dark matter onto stars and its subsequent annihilation.  The effect upon the stars depends upon whether the energy generation rate from dark matter annihilation is large enough to shut off the nuclear burning in the star.  Main sequence WIMP burners look much like protostars moving on the Hayashi track, although they are in principle completely stable.  We make some brief comments about where such stars could be found, how they might be observed and more detailed simulations which are currently in progress.  Finally we comment on whether or not it is possible to link the paradoxically young OB stars found at the galactic centre with WIMP burners. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.3169_arXiv.txt": {
        "abstract": "We compute the electromagnetic radiative corrections to all leading annihilation processes which may occur in the Galactic dark matter halo, for dark matter in the framework of supersymmetric extensions of the Standard Model (MSSM and mSUGRA), and present the results of scans over the parameter space that is consistent with present observational bounds on the dark matter density of the Universe. Although these processes have previously been considered in some special cases by various authors, our new general analysis shows novel interesting results with large corrections that may be of importance, e.g., for searches at the soon to be launched GLAST gamma-ray space telescope. In particular, it is pointed out that regions of parameter space where there is a near degeneracy between the dark matter neutralino and the tau sleptons, radiative corrections may boost the gamma-ray yield by up to three or four orders of magnitude, even for neutralino masses considerably below the TeV scale, and will enhance the very characteristic signature of dark matter annihilations, namely a sharp step at the mass of the dark matter particle. Since this is a particularly interesting region for more constrained mSUGRA models of supersymmetry, we use an extensive scan over this parameter space to verify the significance of our findings. We also re-visit the direct annihilation of neutralinos into photons and point out that, for a considerable part of the parameter space, internal bremsstrahlung is more important for indirect dark matter searches than line signals. ",
        "introduction": "During the last few years a strong consensus has emerged about the existence of a sizeable dark matter contribution to the total cosmological energy density. The identification of experimental signatures that eventually may determine the nature of the cosmological dark matter is thus becoming ever more important. The present estimates \\cite{wmap} give the fraction of the critical density of cold dark matter particles as $\\Omega_{CDM}h^2\\sim 0.105 \\pm 0.013$, where the Hubble parameter (scaled in units of 100 km/s Mpc$^{-1}$) is $h\\sim 0.70\\pm 0.02$. Also, on the scales of galaxies and smaller, a number of methods including measurements of rotation curves as well as gravitational lensing agree well with the predictions from N-body calculations of gravitational clustering in cold dark matter cosmologies (see e.g. \\cite{vialactea}).  The methods of detection of dark matter (for reviews, see \\cite{reviews}) can be divided into {\\em accelerator} production and detection of missing energy (especially at the LHC at CERN, which will start operating some time in 2008), {\\em direct  detection} (of dark matter particles impinging on a terrestrial detector, with recent impressive upper limits reported by \\cite{direct}), or {\\em indirect detection} of particles generated by the annihilation of dark matter particles in the Galactic halo or in the Sun/Earth. All these methods are indeed complementary -- it is probable that a signal from more than one type of  experiment will be needed to fully identify the particle making up the dark matter. The field is just entering very interesting times, with the LHC soon starting and new detectors of liquid noble gases  being developed for direct detection. For indirect detection, the satellite PAMELA \\cite{pamela} was launched a year ago and will soon reveal its first sets of data for positron and antiproton yields in the cosmic rays \\cite{antimatter}. AMANDA \\cite{amanda} at the South Pole that has searched for detection of neutrinos from the centre of the Earth or the Sun \\cite{neutrinos}, will soon give way to the much larger detector IceCUBE \\cite{icecube}, and for gamma-rays coming from annihilations of dark matter particles in the halo \\cite{gammas} the space satellite GLAST \\cite{glast}, to be launched in 2008, will open up a new window to the high-energy universe, for energies from below a GeV to about 300 GeV.  One problem with all these discovery methods is that the signal searched for may be quite weak, with much larger backgrounds in many cases. For indirect detection through gamma-rays, the situation may in principle be better, due to (i) the direct propagation from the region of production, without significant absorption or scattering; (ii) the dependence of the annihilation rate on the square of the dark matter density which may give \"hot spots\" near density concentrations as those predicted by N-body simulations; (iii) possible characteristic features like gamma-ray lines or steps, given by the fact that no more energy than $m_{\\chi}$ per particle can be released in the annihilation of two non-relativistic dark matter particles (we denote the dark matter particle by $\\chi$).  As an example, it was recently shown \\cite{idm} that in models of an extended Higgs sector, the line signal from the two-body final states $\\gamma\\gamma$ and $Z\\gamma$ could give a spectacular signature in the gamma-ray spectrum between 40 and 80 GeV. On the other hand, in models of universal extra dimensions (UED) \\cite{uedline} or in the theoretically perhaps most favoured, supersymmetric, models of dark matter the line feature is in general not very prominent, except in some particular regions of the large parameter space. However, it was early realised that there could be other important spectral features \\cite{lbe89}, and recently it has been shown that internal bremsstrahlung (IB) from produced charged particles in the annihilations could yield a detectable \"bump\" near the highest energy for heavy gauginos or Higgsinos annihilating into $W$ boson pairs, such as expected in split supersymmetry models \\cite{heavysusy}. In \\cite{birkedal}, it was furthermore pointed out that IB often can be estimated by simple, universal formulas and often gives rise to a very prominent step in the spectrum at photon energies of $E_\\gamma=m_\\chi$ (such as in UED models \\cite{Bergstrom:2004cy}).  Encouraged by these partial results, we have performed a detailed analysis of the importance of IB in the minimal supersymmetric extension to the standard model (MSSM). We have therefore calculated the IB contributions for all two-particle charged final states from general neutralino annihilations. Besides confirming the mentioned partial results for the universal radiative corrections, in particular those relating to soft and collinear bremsstrahlung, we also point out interesting cases of model-dependent ``virtual'' brems\\-strahlung (i.e. photons emitted from charged virtual particles), see Fig.~1. We confirm the suspicion expressed already in \\cite{lbe89} that this type of emission may circumvent the chiral suppression, i.e., the annihilation rate being proportional to $m_f^2$ for annihilation into a fermion pair from an $S$-wave initial state, as is the case in lowest order for non-relativistic dark matter Majorana particles in the Galactic halo (see also \\cite{baltz_bergstrom}). Since this enhancement mechanism is most prominent in cases where the neutralino is close to degenerate with charged sleptons, it is of special importance in the so-called stau coannihilation region in models of minimal supergravity (mSUGRA, as implemented in \\cite{isajet}). We therefore run through an extensive scan over these models (based on \\cite{baltz_peskin}) and find, indeed, remarkable cases of enhancement of the gamma-ray rate in the stau coannihilation region, near the maximal possible photon energy $E_\\gamma=m_\\chi$.  Let us stress  that the radiative corrections to the main annihilation channels, here computed systematically for the first time, may turn out to be of utmost importance when fitting gamma-ray data, e.g. from GLAST, to supersymmetric dark matter templates. Over much of the parameter space we have scanned, these corrections give a large factor of enhancement over the  commonly adopted estimates, especially at the observationally most interesting, highest energies. More importantly, they add a feature, the very sharp step at the dark matter mass, that would distinguish this signal from all other astrophysical background (or foreground) processes. ",
        "conclusions": "As can be seen already from our benchmark points in Table~\\ref{benchmark}, and in more detail from the scatter plots in Fig.~\\ref{FSRcomp1}, the internal bremsstrahlung effects computed in this work can be very significant, changing sometimes by more than an order of magnitude the lowest-order prediction for the high-energy gamma-ray signal from neutralino dark matter annihilation. Although some of these enhancements have been found before \\cite{lbe89,heavysusy,birkedal}, this is the first time the first-order radiative corrections have been computed systematically, for all relevant final states in supersymmetric dark matter models. The resulting enhancements of the expected fluxes are surprisingly large over significant regions in the parameter space of the MSSM, including the more constrained mSUGRA models. Despite the fact that some large corrections apply to absolute rates that are too small to be of practical interest, Fig.~4 shows that the quantity ${\\cal S}$, which is directly proportional to the expected signal in gamma-ray detection experiments, also is significant for the internal bremsstrahlung contribution in large regions of parameter space. For $m_\\chi < 300$ GeV, for example, values of ${\\cal S}_{IB}$  greater than 0.1 are generic, and for masses below 100 GeV, values of 1 or higher are common, which in very many cases is higher than the corresponding values for the line signals $\\gamma\\gamma$ and $Z\\gamma$.  One should also bear in mind that the sensitivity of Air Cherenkov Telescopes increases significantly with energy; detectional prospects for a $m_\\chi\\sim1~$TeV neutralino with ${\\cal S}\\sim0.01$, e.g., correspond very roughly to those for a $m_\\chi\\sim100~$GeV neutralino with ${\\cal S}\\sim0.5$ (see, e.g., \\cite{Morselli:2002nw}). In this light, the situation becomes very interesting even for TeV scale Higgsinos, where IB generically contributes more than 10 times as much as secondary photons.    We note that (as anticipated in \\cite{lbe89}) helicity suppression and also CP selection rules of certain final states may be circumvented by emitting a photon; this is for example the origin of the very substantial enhancements of the signal obtained in the stau annihilation region in mSUGRA models. In this situation, the probability of emitting gamma rays vanishes at zero photon energy but increases rapidly at high energy (see Fig.~\\ref{fig:spec}b and \\ref{fig:spec}c), which gives a photon ``bump'' at $E_\\gamma\\approx m_\\chi$.  We also note that in less constrained versions of the MSSM than considered here, we expect even more situations where large enhancements of the annihilation signal due to internal bremsstrahlung can be found. An example is heavy Wino dark matter \\cite{heavysusy,Chattopadhyay:2006xb}, which becomes possible when relaxing the condition $M_1\\approx \\frac{1}{2}M_2$ (as realized, e.g., in anomaly mediated supersymmetry breaking scenarios \\cite{amsmb}).  Of course, the line signals, in particular $\\gamma\\gamma$, have the virtue of being at the highest possible energy, so in order to make a more accurate comparison between these and the IB signal computed here, one would have to model also the expected spectral shape of possible astrophysical gamma-ray backgrounds and the energy resolution of the detector. This is left for future work \\cite{torstenetal}. We note, however, that in general also the new contributions have a characteristic signature, ({\\em cf.} Fig.~2) which can hardly be mimicked by any known astrophysical gamma-ray source. In fact, in some cases these spectra could even be used by future experiments to distinguish between different dark matter candidates (note that, e.g., the distinction between Kaluza-Klein dark matter and a neutralino in the focus point region like our BM4 point would be possible already with the energy resolution of present Air Cherenkov Telescopes \\cite{Bergstrom:2006hk}).   To conclude, we have shown that the commonly neglected first-order radiative corrections to neutralino dark matter annihilation should definitely be taken into account when predicting rates for  gamma-ray telescopes. In particular, the soon to be launched GLAST space telescope \\cite{glast} will have an enhanced possibility over what has previously been assumed to detect radiation from supersymmetric dark matter annihilation. The routines needed to compute these new processes will be included in the next release of the \\ds\\ package \\cite{ds,joakim}.   \\smallskip"
    },
    "0710/0710.4922_arXiv.txt": {
        "abstract": "Using the high-resolution spectrometer SPI on board the \\emph{International Gamma-Ray Astrophysics Laboratory (INTEGRAL)}, we search for a spectral line produced by a dark matter (DM) particle with a mass in the range $40 keV < M_{DM} < 14 MeV$, decaying in the DM halo of the Milky Way. To distinguish the DM decay line from numerous instrumental lines found in the SPI background spectrum, we study the dependence of the intensity of the line signal on the offset of the SPI pointing from the direction toward the Galactic Centre. After a critical analysis of the uncertainties of the DM density profile in the inner Galaxy, we find that the intensity of the DM decay line should decrease by at least a factor of 3 when the offset from the Galactic Centre increases from $0^\\circ$ to $180^\\circ$. We find that such a pronounced variation of the line flux across the sky is not observed for any line, detected with a significance higher than $3\\sigma$ in the SPI background spectrum. Possible DM decay origin is not ruled out only for the unidentified spectral lines, having low ($\\sim 3\\sigma$) significance or coinciding in position with the instrumental ones. In the energy interval from 20 keV to 7 MeV, we derive restrictions on the DM decay line flux, implied by the (non-)detection of the DM decay line. For a particular DM candidate, the sterile neutrino of mass $M_{DM}$, we derive a bound on the mixing angle. ",
        "introduction": "\\subsubsection*{Dark matter in the Universe}  There is a vast body of evidence, suggesting that the large fraction of matter in the Universe exists in the form of the \\emph{Dark matter} \\emph{(DM)}. However, while the total density of the DM is measured with a very high precision ($\\Omega_{\\rm DM}h^2 = 0.105^{+0.007}_{-0.009}$,~\\citealt{WMAP3}), little is known about its properties apart from this. The possibility that the DM is composed of the Standard Model (SM) particles has been ruled out for a long time already.  Indeed, the DM cannot be made out of baryons, as producing such an amount of baryonic matter would require drastic modifications of the scenario of the Big Bang nucleosynthesis (BBN), which otherwise successfully describes the abundance of light elements~(see for example ~\\citealt{Dar:95}). Recent microlensing experiments rule out the possibility that another type of baryonic DM -- massive compact halo objects (MACHOs) -- constitute dominant fraction of mass in the halo~\\citep{MACHO:00,EROS:00,OGLE:98}. The only non-baryonic DM candidate in the SM candidates -- (left-handed) neutrino -- is ruled out from the large scale structure (LSS) considerations~\\citep[see e.g.][]{Bond:80,Hannestad:03,Crotty:04}.  What are the properties of a successful DM candidate?  First of all, this particle should be massive.  Many extensions of the SM present the DM candidates with the masses ranging from $\\sim 10^{-10}\\ev$~(massive gravitons, \\citealt{Dubovsky:04}) and $\\sim 10^{-6}$~eV (axions) to hundreds of GeV (WIMPs) and even to $10^{13}$~GeV \\citep[WIMPZILLA,][]{Kuzmin:98,Kuzmin:99,Chung:98}. For a review of particle physics DM candidates see e.g.~\\cite{Bergstrom:00,Bertone:04,Carr:06}.  Secondly, there should exist mechanisms of DM production with the correct abundances. The production mechanism in particular determines the velocity distribution of particles in the early Universe. This velocity distribution can, in principle, be probed experimentally.  Namely, if during the structure formation epoch the DM particles have velocities, comparable to the speed of sound in the baryon-photon plasma, they ``erase'' density fluctuations at scales, smaller than the distance, they have traveled (called the \\emph{free-streaming length}).  To differentiate various models in accordance with this property, the DM candidates with the negligible velocity dispersion (and, correspondingly, free-streaming) are called \\emph{cold} DM (CDM), while those with the free-streaming of the order of $\\sim 1\\mpc$ are considered to be \\emph{warm} (WDM).\\footnote{The left-handed neutrino would represent \\emph{hot} DM in this terminology, i.e.  the DM with the free-streaming length $\\gg 1$~Mpc.} It is possible to constrain the free-streaming length of a particular DM candidate by probing the structure of the Universe at galaxy-size scales. This can be done through the analysis of the Lyman-$\\alpha$ forest data~\\citep{Hui:97}.  Lyman-$\\alpha$ analysis puts an upper bound on the free-streaming of the DM particles~\\citep{Hansen:01,Viel:05,Seljak:06,Viel:06,Viel:07}. It should be noted however that currently existing interpretation of the Lyman-$\\alpha$ data is model-dependent, as, apart from a number of astrophysical assumptions (see~\\citealt{Hui:97}) and complicated hydrodynamic simulations, it relies on {\\it a priori} assumptions about the velocity distribution of the DM particles.   A way to differentiate between CDM and WDM models would be to compare the numerical simulations of the DM distribution in the Milky Way-type galaxies with the actual observations.  However, the resolution of the N-body simulations is not yet sufficient to answer the questions about e.g. the DM density profiles in dwarf satellite galaxies. Moreover, most of the simulations include only collisionless DM particles, and do not model the baryons and their feedback on the galaxy structure formation.  These problems are not solved even for the CDM simulations, and WDM simulations have additional serious difficulties.  From an observational point of view, it has been argued for some time already that there is a discrepancy between CDM simulations and observations (see e.g.~\\citealt{Moore:94,Moore:99,Klypin:99,Bode:00,Avila-Reese:01,Goerdt:06}) It has been claimed recently that a number of recent observations of dwarf satellite galaxies of the Milky way and Andromeda galaxy seem to indicate the existence of the smallest scale at which the DM exists~\\citep{Gilmore:06,Gilmore:07a,Gilmore:07b,Koposov:07}. However, this statement and the interpretation of the observations are still subject to debate~\\citep{Klimentowski:06,Penarrubia:07,Strigari:07,Simon:07}.  Therefore it is too early to say what kind of DM models is favoured by comparing simulations and observations.    Usually it is also necessary for the DM candidate to be stable. For the most popular DM candidate -- weakly interacting massive particles (WIMPs), this is related to the fact that the particles of $\\sim$ electroweak mass, having weak strength interaction with SM matter (required to produce the correct amount of DM), would decay too fast and would not be ``dark''.  If, however, the DM particle interacts with the SM more weakly than WIMPs, it could well have a finite (although cosmologically long) life time.  There exist several unstable (decaying) DM candidates e.g. gravitino~\\citep{Borgani:96,Baltz:01,Roszkowski:04,Cerdeno:05,Cembranos:06,Lola:07}. In this paper we will concentrate mainly on one candidate, the sterile neutrino (although our results will be applicable for any type of decaying DM). Constraints on the decaying DM were analyzed in~\\cite{DeRujula:80,Berezhiani:87,Doroshkevich:89,Berezhiani:90a,Berezhiani:90b,Bertone:07,Zhang:07} (see also the book by~\\citealt{Khlopov:97}).   \\subsubsection*{Sterile neutrino DM} \\label{sec:sterile-neutrino-dm}  It was noticed long ago that the right-handed (or as it is often called \\emph{sterile}) neutrino with the mass in the keV range would represent a viable DM candidate~\\citep{Dodelson:93}. Such a neutrino would interact with the rest of the matter only via the quadratic mixing with left-handed (\\emph{active}) neutrinos and therefore (although not stable) could have cosmologically long life-time. At the same time, it could be produced in the early Universe with the correct abundances~\\citep{Dodelson:93,Shi:98,Shaposhnikov:06}.  One of the decay channels of the unstable sterile neutrinos includes emission of photons of the energy equal to half of the sterile neutrino rest energy.  This potentially provides a possibility to observe the decays of DM sterile neutrinos via detection of a characteristic spectral line in the % spectra of astrophysical objects with large DM concentration.  Recently this DM candidate has attracted much attention (see e.g.~\\citet{Shaposhnikov:07a} and references therein).  It was found that a very modest and natural extension of the SM by 3 right-handed neutrinos (making the SM more symmetric as all SM fermions, including neutrino, would have now their left and right handed counterparts) provided a viable extension of the theory, capable of solving several ``beyond the SM'' problems. First of all, such an extension makes neutrinos massive and thus perhaps provides the simplest and the most natural explanation of the phenomenon of ``neutrino oscillations''~(see e.g.~\\cite{Fogli:05,Strumia:06,Giunti:06} for reviews). The smallness of neutrino masses in this model (called \\numsm in~\\citealt{Asaka:05b}) is achieved by the usual see-saw mechanism with Majorana masses of right-handed neutrinos being below electroweak scale.\\footnote{The fact that the \\numsm does not introduce any new scale above the electroweak one, makes this theory especially appealing from the point of view of its experimental verification/falsification.}  Secondly, if two heavier sterile neutrinos ($N_2$ and $N_3$) are almost degenerate in mass and have their masses between $\\mathcal{O}(100)\\mev$ and $\\mathcal{O}(20)\\gev$, the \\numsm provides the mechanism of generating the baryon asymmetry of the Universe. Thirdly, the lightest sterile neutrino $N_1$ can have arbitrary mass and arbitrarily weak coupling with the (active) neutrino sector. At the same time, it can be produced in the early Universe in the correct amounts.  It represents therefore the DM particle in the \\numsm. Thus, altogether the \\numsm represents (arguably) the simplest extension of the SM, capable of explaining three important questions: origin and smallness of neutrino masses, baryon asymmetry in the Universe and the existence of the DM.  \\subsubsection*{Existing restrictions on sterile neutrino DM parameters.}  What are the current restrictions on parameters (mass and \\emph{mixing}) of sterile neutrino DM? First of all sterile neutrino mass should satisfy the universal Tremaine-Gunn lower bound:\\footnote{In its simplest form the Tremaine-Gunn bound comes from the fact that for the fermions there is a maximal density in the phase space~\\citep{Tremaine:79,Dalcanton:00} and therefore the observed phase-space density in various DM dominated systems should be less that this (mass dependent) bound.}  $\\mdm\\gtrsim 300-500\\ev$.\\footnote{A stronger lower bound from Ly-$\\alpha$~\\citep{Seljak:06,Viel:06,Viel:07} can be obtained in the case of the particular production mechanisms -- the Dodelson-Widrow scenario~\\citep{Dodelson:93}. For other possible production mechanisms~\\citep[e.g.][]{Shi:98,Shaposhnikov:06} the Ly-$\\alpha$ constraints should be reanalyzed.}  Next, as the sterile neutrino possesses the (two-body) radiative decay channel: $N_1 \\to \\nu + \\gamma$, the emitted photon would carry the energy $E_\\gamma = \\mdm/2$. A large flux of such photons is expected from the large concentrations of the DM sterile neutrinos, like galaxies or galaxy clusters.  Recently an extensive search of the DM decay line in the region of masses $M_\\dm \\lesssim 20\\kev$ was conducted, using the data of \\emph{Chandra}~\\citep{Riemer:06,Boyarsky:06e,Abazajian:06b} and \\emph{XMM-Newton}~\\citep{Boyarsky:05,Boyarsky:06b,Boyarsky:06c,Watson:06,Boyarsky:06d}. The region of soft X-ray (down to energies $0.2$~keV) was explored by~\\citet{Boyarsky:06f} with the use of the wide field of view spectrometer~\\citep{McCammon:02}.  The non-observation of the DM decay line in X-ray, combined with the first principles calculation of DM production in the early Universe~\\citep{Asaka:06c}, implies that the \\citet{Dodelson:93} (DW) scenario can work only if the sterile neutrino mass is below 4~keV~\\citep{Boyarsky:07a}.  If one takes into account recent \\emph{lower} bound on the mass of sterile neutrino DM in the DW scenario $\\mdm \\ge 5.6\\kev$~\\citep{Viel:07}, it seems that the possibility that all the DM is produced via DW scenario is ruled out~\\citep{Boyarsky:07a}.  The possibility that only fraction of the DM is produced via DW mechanism remains open~\\citep{Palazzo:07}.   There are other viable mechanisms of DM production, including e.g.  resonant oscillation production in the presence of lepton asymmetries~\\citep{Shi:98}. Sterile neutrino DM can be produced by the decay of light inflaton \\citep{Shaposhnikov:06} or in a similar model with the different choice of parameters~\\citep{Kusenko:06a,Petraki:07}. These mechanisms are currently not constrained and remain valid for DM particles with the masses in the keV range and above.  The search for the DM decay line signal produced by sterile neutrinos with masses above $\\sim 20$~keV is complicated by the absence of the focusing optics telescopes (similar to {\\it Chandra} or {\\it XMM-Newton}) in the hard X-ray and $\\gamma$-ray domain of the spectrum.  For example, the existing restrictions in the $20-100$~keV mass range \\citep{Boyarsky:05,Boyarsky:06c} are derived from the observations of diffuse X-ray background, with the help of non-imaging instruments, HEAO-I~\\citep{Gruber:99}. The current status of astrophysical observations in summarized in~\\cite{Ruchayskiy:07}.  In this paper we use the spectrometer SPI on board of INTEGRAL satellite to place restrictions on parameters of decaying DM in the mass range $40\\kev - 14\\mev$.  This range of masses is interesting, for example, the sterile neutrinos, produced in the early Universe in the presence of large lepton asymmetries~\\citep{Shi:98} or through the inflaton decay~\\citep{Shaposhnikov:06}. It is also relevant for the case of gravitino DM~\\citep{Pagels:82,Bond:82}.  When the preparation of this paper was at its final stage,~\\citet[hereafter \\textbf{Y07}]{Yuksel:07} published their work, which used the results of~\\citet[hereafter \\textbf{T06}]{Teegarden:06} to place restrictions on the parameters of sterile neutrino DM in the range $40-700\\keV$. We discuss it in more details in Section~\\ref{sec:discussion}.    \\subsubsection*{SPI spectrometer}  The absence of the focusing optics significantly reduces the sensitivity of the telescopes operating in the hard X-ray/soft \\gr\\ energy band. Most of the instruments operating in this energy band use collimators and/or coded masks to distinguish signals from the sources on the sky from the instrumental background. Contrary to the focusing optics telescopes, both the source and background signals are collected from the entire detector, which significantly increases the irreducible background.  The focusing optics enables to significantly reduce the background only in the studies of point sources. If the source under investigation occupies a large fraction of the sky (e.g. the entire Milky Way galaxy), the performance of the focusing and non-focusing instruments with the same detector collection area are, in fact, comparable.  \\begin{figure} \\centering \\includegraphics[width=\\linewidth]{SENSITIVITY_SPI} % \\caption{Comparison of sensitivity towards the search of the narrow DM decay line for different instruments with the wide FoV. Diagonal straight lines show the improvement of sensitivity (by a factor, marked on the line) as compared with the HEAO-I A4 low energy detector (LED), taken as a reference.} \\label{fig:spi_vs_heao} \\end{figure}  In the case of an extended source, emitting a narrow spectral line, an efficient way of reduction of instrumental background is via the improvement of the spectral resolution of the instrument (in the case of a broad continuum background spectrum, the number of background counts at the energy of the line is proportional to the spectral resolution $\\Delta E$). The best possible sensitivity is achieved when the spectral resolution reaches the intrinsic width of the spectral line (see Fig.\\ref{fig:spi_vs_heao} for the case of wide FoV instruments and~\\cite{Boyarsky:06f} for the case of narrow FoV instruments).  In the case of the line produced by the DM decaying in the Milky Way halo, the line width is determined by the Doppler broadening by the random motion of the DM particles. The velocity dispersion of the DM motion in the halo is about the rotation velocity of the Galactic disk, $v\\sim 200$~km/s.  This means that Doppler broadening of the DM decay line is about \\begin{equation} \\label{eq:7} \\frac{\\Delta E}{E}\\sim \\frac{v}{c}\\simeq 10^{-3}\\;. \\end{equation} Thus, the optimal spectral resolution of an instrument searching for the DM decay line produced by the Milky Way DM halo should be $\\Delta E\\simeq 10^{-3}E$.  Such optimal spectral resolution is almost achieved with the spectrometer SPI on board of \\intgr\\ satellite, which has the maximal spectral resolving power of $E/\\Delta E\\simeq 500$ and works in the energy range 20~keV -- 8~MeV~\\citep{spi}. SPI is a ``coded mask'' type instrument with an array of 19 hexagonal shaped Ge detectors (of which only 17 are operating at the moment).  \\begin{figure} \\begin{center} \\includegraphics[width=\\linewidth]{FoV} \\end{center} \\caption{The geometry of the SPI FoV.} \\label{fig:FoV} \\end{figure} The SPI telescope consists of a coded mask inscribed into a circle of the radius $R_{\\rm mask}=39$~cm, placed at the height $H=171$~cm above the detector plane and of the detector, which has the shape of a hexagon inscribed into a circle of the radius $R_{\\rm det}\\simeq 15.3$~cm (see Fig. \\ref{fig:FoV}). The portion of the sky visible from each point of the SPI detector (the so-called \\emph{fully coded field of view}, FCFOV) has therefore angular diameter \\begin{equation} \\label{eq:12} \\Theta_\\fcfov=2\\arctan\\left[\\frac{R_{\\rm mask}-R_{\\rm det}}H\\right] \\approx 16^\\circ\\;, \\end{equation} while the portion of the sky visible by at least some of the detectors (the \\emph{partially coded field of view}, PCFOV) is \\begin{equation} \\label{eq:14} \\Theta_\\pcfov=2\\arctan\\left[\\frac{R_{\\rm mask}+R_{\\rm det}}H\\right] \\approx 35^\\circ\\;. \\end{equation} The solid angle spanned by the cone with this opening angle is $\\Omega_\\pcfov=2\\pi\\Bigl(1-\\cos(\\Theta_\\pcfov/2)\\Bigr)\\simeq 0.29$ (see Fig.~\\ref{fig:FoV}). Wide field of view makes the SPI telescope suitable for the study of the very extended sources, like the Milky Way DM halo.   \\begin{figure} \\begin{center} \\includegraphics[width=\\linewidth]{Aeff} \\end{center} \\caption{The effective area of the SPI detector for an on-axis source, as a function of the photon energy. The plot is produced by collective the on-axis effective areas of the 17 SPI detectors from the instrumental characteristics files.} \\label{fig:Aeffon} \\end{figure} ",
        "conclusions": "\\label{sec:discussion}  The purpose of this work was to understand how to search for the DM decay line with the SPI spectrometer and to check that none of the strong lines, present in the SPI background, was confused with the DM decay line. Our analysis shows that all the strong lines were, indeed, of instrumental origin and provides the upper bound on the flux of ``weak'' ($3{-}4\\sigma$ above the background) lines, which leads to the corresponding restrictions (see Sec.~\\ref{sec:exclusions}).  To further improve the results, one needs to work with the weak lines (or lines, coinciding in position with instrumental ones). To do this one needs more sophisticated procedures of subtraction of the instrumental background (e.g. imaging).  One of the most interesting cases of the coinciding instrumental and celestial line is the positronium annihilation line at 511 keV. An excess of positron annihilation emission on top of the strong instrumental line (related to positrons annihilating inside the detector) was noticed long ago~\\citep[for an incomplete set of references see e.g.][]{Prantzos:93,Milne:99,Cheng:97,Purcell:97,Knodlseder:05,Weidenspointner:06,Weidenspointner:07}. There exist many attempts of explanation of this excess. In particular, it was attributed to the annihilating or decaying DM~\\citep[see e.g.][]{Boehm:03,Hooper:03,Boehm:06,Frere:06,Picciotto:04,Rasera:05}.  The sterile neutrino DM with the mass $m_s > 1\\mev$ possesses decay channel $N_s\\to e^+e^-\\nu$, with positrons annihilating either in flight or at rest, by forming the positronium atom~\\citep[see e.g.][]{Beacom:06,Sizun:06}. Thus, it is possible that the decay of sterile neutrino DM contributes to such a line.  The detailed analysis of this case will be reported separately.   It should be also mentioned, that the region of masses between $20\\kev \\lesssim m_\\dm \\lesssim 40\\kev$ remains inaccessible for the existing X-ray missions. The strongest restrictions in this region were produced, using the data of HEAO-1 mission \\citep{Boyarsky:06c}.   When the work on this paper was at its final stage, the work of Y07 was published. Y07 obtained the restrictions on parameters of sterile neutrino in the range 40 keV -- 700 keV. To facilitate the comparison, we plot the restrictions of Y07 on Fig.~\\ref{fig:on_limit}, (divided by the factor of 2 to translate them into the restrictions for the Majorana, rather than Dirac sterile neutrino DM, see footnote~\\ref{fn:1}, p~\\pageref{fn:1}).  As the data, used in our work, has about 5 times longer exposure than the \\intgr\\ first years data, on which the results of Y07 are based, we could have expected  results stronger by a factor $\\approx 2$ in our case.  However, the Fig.~\\ref{fig:on_limit} shows the opposite.  The reason for this is as follows. For the SPI, the sensitivity towards the line search from a particular source depends on the shape of the source. In particular, the results of TW06, on which the work of Y07 was based, were obtained under the assumption of a particular diffuse source ($10^\\circ$ Gaussian). As any realistic DM profile is much flatter than the $10^\\circ$ Gaussian, the results of TW06 cannot be applied directly for the case of the DM line search.  They should be rescaled to account for the diffuse nature of the DM source~(c.f. Section~\\ref{sec:exclusions}). Apart from this, the estimated DM signal from the inner part of the Galaxy is about 2 times stronger in Y07 than in our work.  As the DM signal in the direction of the GC is the most uncertain, we have adopted the conservative flat profile everywhere inside the solar radius, to minimize this uncertainty.     \\subsection*"
    },
    "0710/0710.1500_arXiv.txt": {
        "abstract": "Observations of the ELAIS-N1 field taken at 610~MHz with the Giant Metrewave Radio Telescope are presented.  Nineteen pointings were observed, covering a total area of $\\sim9$~deg$^{2}$ with a resolution of 6~$\\times$~5~arcsec$^{2}$, PA $+45\\degr$.  Four of the pointings were deep observations with an rms of $\\sim40$~$\\mu$Jy before primary beam correction, with the remaining fifteen pointings having an rms of $\\sim70~\\mu$Jy.  The techniques used for data reduction and production of a mosaicked image of the region are described, and the final mosaic is presented, along with a catalogue of 2500 sources detected above 6$\\sigma$.  This work complements the large amount of optical and infrared data already available on the region.  We calculate 610-MHz source counts down to 270~$\\mu$Jy, and find further evidence for the turnover in differential number counts below 1~mJy, previously seen at both 610~MHz and 1.4~GHz. ",
        "introduction": "The {\\it Spitzer} Wide-area Infrared Extragalactic \\citep[SWIRE;][]{Lonsdale03} survey has the largest sky coverage of the legacy surveys being performed by the {\\it Spitzer Space Telescope} \\citep{Werner04}.  A total area of $\\sim$49~deg$^{2}$ of sky has been observed with the Infrared Array Camera \\citep[IRAC;][]{Fazio04} and Multiband Imaging Photometer for {\\it Spitzer} \\citep[MIPS;][]{Rieke04} instruments at 3.6, 4.5, 5.8, 8, 24, 70 and 160~$\\mu$m.  The survey is broken down into six fields, three in the northern sky -- ELAIS-N1, ELAIS-N2 and the Lockman Hole -- and three in the south -- ELAIS-S1, {\\it Chandra} Deep Field South and the {\\it XMM}-Large Scale Structure (XMM-LSS) field.  All six regions were selected to be away from the Galactic disk, in order to minimize background cirrus emission.  There is a large amount of multi-wavelength information available on all six SWIRE fields.  The three ELAIS fields were observed as part of the European Large-Area {\\it ISO} Survey \\citep{Oliver00}, which also included another northern (-N3) and southern (-S2) field.  The {\\it Infrared Space Observatory} ({\\it ISO}) observed these regions at 6.7, 15, 90 and 175~$\\mu$m, and a large number of followup observations were carried out in the optical, infrared and radio bands.  A band-merged catalogue, containing the {\\it ISO} data, along with $U$, $g'$, $r'$, $i'$, $Z$, $J$, $H$ and $K$-band detections, and radio observations at 1.4~GHz has been produced -- for more details, see \\citet{RowanRobinson04}, and references therein.  Observations of the ELAIS-N1 region were taken with {\\it Spitzer} in 2004 January, covering $\\sim9$~deg$^{2}$ with the IRAC and MIPS instruments.  The source catalogues have been produced, and are available online \\citep{Surace04}, containing over 280,000 sources. The UK Infrared Deep Sky Survey \\citep[UKIDSS;][]{Lawrence07} intends to cover the ELAIS-N1 region in its Deep Extragalactic Survey plan, observing the full field in the $J$, $H$ and $K$-bands to a depth of $K$ = 21~mag.  This will be a great improvement over the currently available surveys, which have a sensitivity limit of $\\sim$18~mag in the $K$ band.  Data Release 2 \\citep{Warren07} of UKIDSS contains early shallow data on the ELAIS-N1 region.  Further surveys have been carried out in the $R$-band \\citep{Fadda04}, in H$\\alpha$ \\citep{Pascual01}, and with the {\\it Chandra} X-ray telescope \\citep{Manners03,Franceshini05}.  There have been several redshift surveys of the region \\citep{Trichas06,Berta07}, and the ELAIS-N1 region was also partially covered by the Sloan Digital Sky Survey \\citep[SDSS;][]{AdelmanMcCarthy07}.  While there have been a great number of observations of the ELAIS-N1 region at optical and infrared wavelengths, there is comparatively little radio information available.  The existing VLA 1.4~GHz survey of the three northern ELAIS fields \\citep{Ciliegi99}, which has been included into the band-merged catalogue of \\citet{RowanRobinson04}, reaches a 5$\\sigma$ limit of 0.135~mJy over 0.12~deg$^{2}$ but only a 1.15~mJy limit over its full coverage area of 4.22~deg$^{2}$.  The NVSS \\citep{Condon98} and FIRST \\citep{Becker95} surveys both cover the ELAIS-N1 region, but only to relatively shallow $5\\sigma$ limits of 2.25 and 0.75~mJy respectively.  A recent study of polarised compact sources \\citep{Taylor07} at 1420~MHz is underway, using the Dominion Radio Astrophysical Observatory Synthesis Telescope (DRAO ST) centered on $16^{\\rm h}11^{\\rm m}00^{\\rm s}$, $+55\\degr00'00''$ and covering 7.4~deg$^{2}$.  The first 30~per~cent of observations have been completed, with maps in Stokes I, Q and U being produced with a maximum sensitivity of 78~$\\mu$Jy~beam$^{-1}$, although with a resolution of $\\sim$1~arcmin$^{2}$.  In order to extend the information on this region, a much larger deep radio survey is required.  In this paper, we present observations of the ELAIS-N1 survey field taken at 610~MHz with the Giant Metrewave Radio Telescope \\citep[GMRT;][]{Ananthakrishnan05}, covering $\\sim9$~deg$^{2}$ of sky with a resolution of $6\\times5$~arcsec$^{2}$, PA~$+45\\degr$, centred on $16^{\\rm h}11^{\\rm m}00^{\\rm s}$, $+55\\degr00'00''$ (J2000 coordinates, which are used throughout this paper).  This survey, in combination with the deep {\\it Spitzer} data, will be used to study the infrared/radio correlation for star-forming systems \\citep[e.g.][]{Appleton04}, and the link between the triggering of star formation and AGN activity, as well as the properties of the faint radio population at 610~MHz.  In Section~\\ref{sec:observations} we describe the observations and data reduction techniques used in the creation of the survey. Section~\\ref{sec:results} presents the mosaic and a source catalogue containing 2500 sources above 6$\\sigma$, along with a sample of extended sources.  In Section~\\ref{sec:sourcecounts} we construct the 610~MHz differential source counts, and compare them to previous works. ",
        "conclusions": "\\label{sec:results} \\begin{figure} \\includegraphics[width=8cm]{EN1NOISE.PS} \\caption{The rms noise of the final mosaic, calculated using Source Extractor.  The grey-scale ranges between 40 and 350~$\\mu$Jy, and the contours are at 60 and 120~$\\mu$Jy respectively.} \\label{fig:EN1noise} \\end{figure}  Source Extractor \\citep[SExtractor;][]{Bertin96} was used to calculate the rms noise $\\sigma$ across the mosaic.  A grid of $16\\times16$ pixels was used in order to track changes in the local noise level, which varies significantly near the brightest sources. Fig.~\\ref{fig:EN1noise} illustrates the local noise, with the grey-scale varying between 40~$\\mu$Jy (the noise level in the centre of the deep pointings) and 350~$\\mu$Jy (the noise level for the shallow pointings, at the distance where the GMRT primary beam gain was 20~per~cent of its central value).  The 60 and 120~$\\mu$m contours are plotted, which cover the majority of the deep and shallow survey regions respectively.  A sample region of the ELAIS-N1 survey is shown in Fig.~\\ref{fig:EN1sample} to illustrate the quality of the image.  Most of the sources in our survey are unresolved, although there are several with extended structures.  We present a sample of these in Fig.~\\ref{fig:extended}.  \\begin{figure*} \\includegraphics[width=16cm]{EN1area.eps} \\caption{A $70\\times70$~arcmin$^{2}$ region of the 610~MHz image, to illustrate the quality of the survey.  The region is located within the deeper area of the survey, and most sources are unresolved. The grey-scale ranges between $-0.2$ and 1~mJy~beam$^{-1}$, and the noise is relatively uniform and between 40 and 60~$\\mu$Jy, apart from small regions near bright sources where the noise increases.} \\label{fig:EN1sample} \\end{figure*}  \\begin{figure*} \\centerline{\\subfigure[GMRTEN1~J161137.8$+$555955] {\\includegraphics[width=5cm]{IMG12.PS}} \\subfigure[GMRTEN1~J160640.9$+$560136] {\\includegraphics[width=5cm]{IMG13.PS}} \\subfigure[GMRTEN1~J161530.7$+$545231] {\\includegraphics[width=5cm]{IMG02.PS}}} \\centerline{\\subfigure[GMRTEN1~J160929.9$+$552444 and J160931.09$+$552503] {\\includegraphics[width=5cm]{IMG04.PS}} \\subfigure[GMRTEN1~J161858.8$+$545227, J161900.3$+$545305 and J161903.3$+$545240] {\\includegraphics[width=5cm]{IMG05.PS}} \\subfigure[GMRTEN1~J161148.4$+$550049, J161151.5$+$550053 and J161154.4$+$550057] {\\includegraphics[width=5cm]{IMG07.PS}}} \\centerline{\\subfigure[GMRTEN1~J160808.7$+$544723, J160809.5$+$544659, J160810.3$+$544633 and J160810.7$+$544645] {\\includegraphics[width=5cm]{IMG08.PS}} \\subfigure[GMRTEN1~J161027.4$+$541246] {\\includegraphics[width=5cm]{IMG09.PS}} \\subfigure[GMRTEN1~J161757.1$+$545110] {\\includegraphics[width=5cm]{IMG03.PS}}} \\centerline{\\subfigure[GMRTEN1~J161140.5$+$554703 and J161142.8$+$554727] {\\includegraphics[width=5cm]{IMG14.PS}} \\subfigure[GMRTEN1~J160634.8$+$543456] {\\includegraphics[width=5cm]{IMG15.PS}} \\subfigure[GMRTEN1~J160334.6$+$542900 and J160333.1$+$542914] {\\includegraphics[width=5cm]{IMG16.PS}}} \\caption{A selection of extended objects in the ELAIS-N1 610~MHz GMRT survey -- contours are plotted at $\\pm200~\\mu$Jy $\\times$ 1, $\\sqrt{2}$, 2, $2\\sqrt{2}$, 4$\\ldots$, with the exception of object (e), where the contours start at $\\pm$500~$\\mu$Jy.  Negative contours are represented by dashed lines.  The resolution of the beam is shown in the bottom left of each image, and the designations of each source component are given below.} \\label{fig:extended} \\end{figure*}  Our GMRT data suffer from dynamic range problems near the brightest sources, and the final mosaic has increased noise and residual sidelobes in these regions.  We had fewer problems with our survey of the xFLS region, due to the longer time spent on each pointing and the correspondingly better {\\it uv} coverage.  There were also fewer bright sources in the xFLS field, so a much smaller region was affected by residual sidelobes.  Fig.~\\ref{fig:BrightSource} shows an area around one of the bright sources, to illustrate the problems caused by the residual sidelobes.  While the local noise calculated by SExtractor increases due to these residuals, some of them still have an apparent signal-to-noise level that is greater than 6.  We therefore opted for a two-stage selection criteria for our final catalogue.  \\begin{figure} \\includegraphics[width=8cm]{BrightSource.ps} \\caption{Source GMRTEN1~J161212.4$+$552308, and the errors surrounding it.  The grey-scale ranges between $-0.2$ and 1~mJy~beam$^{-1}$, and the source has a peak brightness of 389~mJy.  The region affected by an overdensity of sources is shown by the black circle, of radius 10~arcmin -- see text for more details.} \\label{fig:BrightSource} \\end{figure}  \\subsection{Source fitting} An initial catalogue of 4767 sources was created using SExtractor. The mosaic was cut off at the point where the primary beam correction dropped to 20~per~cent of its central value (a radius of 0\\fdg53 from the outer pointings), however only sources inside the 30~per~cent region (0\\fdg47) were included in the catalogue to avoid the mosaic edges from affecting the estimation of local noise.  The requirements for a source to be included were that it had at least 5 connected pixels with brightness greater than 3$\\sigma$, and a peak brightness greater than 6$\\sigma$.  The image pixel size meant that the beam was reasonably oversampled, so the source peak was taken to be the value of the brightest pixel within a source.  The integrated flux density was calculated using the FLUX\\_AUTO option within SExtractor.  This creates an elliptical aperture around each object \\citep[as described in][]{Kron80}, and integrates the flux contained within the ellipse. Comparisons between the flux density obtained through this method and through the method developed in \\citet{Garn07} -- pixels above a given threshold were summed, then empirically corrected for the elliptical beam shape -- give good agreement between the two techniques.  Sources with Kron flux density above 1~mJy showed no statistical difference between the two flux density measurements, with an uncertainty of 3~per~cent.  Sources below 1~mJy had a Kron flux density that was systematically larger than the \\citet{Garn07} method by 6~per~cent, with an uncertainty of 4~per~cent.  We chose to use the Kron method in order to avoid the empirical correction factor.  In order to estimate the area affected by artefacts near the bright sources, we calculated the number of (potentially spurious) sources in a series of concentric rings centred on a bright source, and converted this to an effective source density as a function of distance.  While the noise does vary across the map, the variation is smooth and relatively slow away from the bright sources.  Using the more distant rings we calculated a mean source sky density $\\mu$, and an estimate of the error in this value, $\\sigma_{\\mu}$, which will be unaffected by the presence of the spurious sources, and will vary slightly between each bright central source due to the changing properties of the image.  Fig.~\\ref{fig:spurioussources} shows the source density (in arbitrary units) away from the source in Fig.~\\ref{fig:BrightSource}.  The clear over-density of sources can be seen within 10~arcmin of the source. If there is a significant peak in the density (greater than $\\mu + 6\\sigma_{\\mu}$), then we define the size of the affected region by finding the first radius at which the source density drops below $\\mu + 3\\sigma_{\\mu}$ -- for this source, the affected radius is 10~arcmin. The affected region has been added to Fig.~\\ref{fig:BrightSource}. Within this region, only sources with a peak brightness greater than 12$\\sigma$ are included in our catalogue -- this value was determined empirically.  The source density plot for a 10~mJy source is shown in Fig.~\\ref{fig:FaintSource}, on the same scale as Fig.~\\ref{fig:spurioussources} -- while there may still be a slight overdensity near the centre, the increased noise level in this region, along with the $6\\sigma$ cut-off reduces the risk of selecting a spurious source near sources with weak over-densities.  \\begin{figure} \\includegraphics[width=8cm]{Bright.eps} \\caption{Density of sources, in concentric rings around the bright source shown in Fig.~\\ref{fig:BrightSource}.  The overdensity of sources near to the bright (389~mJy) central object can be clearly seen.  The mean source density, far from the central source, is given by the large dashed line while the cutoff density defining the affected region is given by short dashes -- see text for more details.} \\label{fig:spurioussources} \\end{figure}  \\begin{figure} \\includegraphics[width=8cm]{Faint.eps} \\caption{Density of sources, in concentric rings around a fainter (10~mJy) object in the ELAIS-N1 survey field.  There is still a slight overdensity near the central source, but the signal-to-noise cutoff of $6\\sigma$ ensures that spurious sources are not included in the catalogue.} \\label{fig:FaintSource} \\end{figure}  This analysis was repeated for all sources with a peak greater than 10~mJy in order to filter spurious sources.  The final catalogue contains 2500 sources -- we have erred on the side of caution in order to produce a catalogue with little contamination from spurious sources.  The size of the affected region is correlated with the peak brightness of a source (with Pearson product-moment correlation coefficient of 0.53), and the number of spurious sources is also correlated with the peak brightness, with correlation coefficient 0.73.  The precise size of the affected region depends on the {\\it uv} coverage for the relevant pointing, the time spent on observations and the local noise levels.  Table $\\ref{tab:catalogue}$ presents a sample of 60 entries in the catalogue, which is sorted by right ascension.  The full table is available via {\\tt http://www.mrao.cam.ac.uk/surveys/}.  Column~1 gives the IAU designation of the source, in the form GMRTEN1~Jhhmmss.s$+$ddmmss, where J represents J2000.0 coordinates, hhmmss.s represents right ascension in hours, minutes and truncated tenths of seconds, and ddmmss represents the declination in degrees, arcminutes and truncated arcseconds.  Columns~2 and 3 give the right ascension and declination of the source, calculated by first moments of the relevant pixel brightnesses to give a centroid position. Column~4 gives the brightness of the peak pixel in each source, in mJy~beam$^{-1}$, and column~5 gives the local rms noise in $\\mu$Jy~beam$^{-1}$.  Columns~6 and 7 give the integrated flux density and error, calculated from the local noise level and source size. Columns 8 and 9 give the $X$, $Y$ pixel coordinates of the source centroid from the mosaic image.  Column 10 is the Source Extractor deblended object flag -- 1 where a nearby bright source may be affecting the calculated flux, 2 where a source has been deblended into two or more components from a single initial island of flux, and 3 when both of the above criteria apply.  There are 232 sources present in our catalogue with non-zero deblend flags; it is necessary to examine the images to distinguish between the case where one extended object has been represented by two or more entries, and where two astronomically distinct objects are present.  \\begin{table*} \\label{tab:catalogue} \\caption{A sample of 60 entries from the 610-MHz ELAIS-N1 catalogue, sorted by right ascension.  The full version of this table is available as Supplementary Material through the online version of this article, and via {\\tt http://www.mrao.cam.ac.uk/surveys/}.} \\begin{tabular}{cccccccccc} \\hline Name & RA & Dec.\\ & Peak & Local~Noise & Int.\\ Flux Density & Error & $X$ & $Y$ & Flags\\\\ & J2000.0 & J2000.0 & mJy~beam$^{-1}$ & $\\mu$Jy~beam$^{-1}$ & mJy & mJy & & & \\\\ (1) & (2) & (3) & (4) & (5) & (6) & (7) & (8) & (9) & (10) \\\\ \\hline GMRTEN1~J160319.2$+$542543 & 16:03:19.22 & $+$54:25:43.6 & 0.496 & 75 & 0.415 & 0.067 & 6599 & 2946 & 0\\\\ GMRTEN1~J160319.2$+$553149 & 16:03:19.24 & $+$55:31:49.1 & 0.469 & 66 & 0.383 & 0.054 & 6527 & 5589 & 0\\\\ GMRTEN1~J160319.3$+$554950 & 16:03:19.38 & $+$55:49:50.2 & 0.514 & 85 & 0.265 & 0.054 & 6506 & 6309 & 0\\\\ GMRTEN1~J160320.8$+$550645 & 16:03:20.89 & $+$55:06:45.9 & 2.617 & 97 & 3.425 & 0.175 & 6545 & 4587 & 0\\\\ GMRTEN1~J160321.1$+$553654 & 16:03:21.13 & $+$55:36:54.2 & 0.418 & 69 & 0.309 & 0.076 & 6510 & 5792 & 0\\\\ GMRTEN1~J160322.6$+$554320 & 16:03:22.67 & $+$55:43:20.8 & 0.560 & 75 & 0.444 & 0.068 & 6495 & 6049 & 0\\\\ GMRTEN1~J160323.1$+$543737 & 16:03:23.13 & $+$54:37:37.3 & 1.920 & 74 & 2.256 & 0.119 & 6564 & 3421 & 0\\\\ GMRTEN1~J160324.0$+$550618 & 16:03:24.05 & $+$55:06:18.2 & 0.635 & 84 & 0.495 & 0.079 & 6527 & 4568 & 0\\\\ GMRTEN1~J160326.0$+$540817 & 16:03:26.00 & $+$54:08:17.1 & 0.595 & 94 & 0.486 & 0.081 & 6579 & 2248 & 0\\\\ GMRTEN1~J160327.7$+$552647 & 16:03:27.78 & $+$55:26:47.3 & 3.240 & 84 & 3.867 & 0.149 & 6484 & 5387 & 0\\\\ GMRTEN1~J160327.9$+$543326 & 16:03:27.99 & $+$54:33:26.0 & 0.452 & 75 & 0.338 & 0.064 & 6540 & 3253 & 0\\\\ GMRTEN1~J160329.5$+$540705 & 16:03:29.55 & $+$54:07:05.7 & 0.869 & 131 & 0.681 & 0.117 & 6559 & 2200 & 0\\\\ GMRTEN1~J160330.8$+$542454 & 16:03:30.83 & $+$54:24:54.3 & 0.483 & 78 & 0.837 & 0.097 & 6533 & 2912 & 0\\\\ GMRTEN1~J160331.1$+$554327 & 16:03:31.14 & $+$55:43:27.2 & 0.473 & 76 & 0.555 & 0.071 & 6447 & 6052 & 0\\\\ GMRTEN1~J160331.3$+$545000 & 16:03:31.39 & $+$54:50:00.7 & 0.408 & 67 & 0.322 & 0.055 & 6503 & 3916 & 0\\\\ GMRTEN1~J160332.3$+$553000 & 16:03:32.38 & $+$55:30:00.2 & 0.508 & 69 & 0.536 & 0.076 & 6454 & 5514 & 0\\\\ GMRTEN1~J160332.6$+$554622 & 16:03:32.66 & $+$55:46:22.6 & 3.212 & 76 & 4.152 & 0.153 & 6435 & 6169 & 0\\\\ GMRTEN1~J160332.9$+$541746 & 16:03:32.90 & $+$54:17:46.8 & 0.474 & 73 & 0.381 & 0.060 & 6528 & 2627 & 0\\\\ GMRTEN1~J160333.1$+$542914 & 16:03:33.11 & $+$54:29:14.2 & 2.377 & 73 & 9.697 & 0.226 & 6515 & 3085 & 3\\\\ GMRTEN1~J160333.6$+$552623 & 16:03:33.69 & $+$55:26:23.7 & 0.773 & 78 & 0.878 & 0.101 & 6451 & 5370 & 0\\\\ GMRTEN1~J160333.7$+$540540 & 16:03:33.70 & $+$54:05:40.6 & 7.541 & 294 & 9.799 & 0.552 & 6536 & 2142 & 0\\\\ GMRTEN1~J160334.6$+$542900 & 16:03:34.68 & $+$54:29:00.2 & 2.326 & 80 & 4.235 & 0.167 & 6506 & 3075 & 3\\\\ GMRTEN1~J160335.1$+$551534 & 16:03:35.19 & $+$55:15:34.1 & 0.462 & 74 & 0.333 & 0.066 & 6454 & 4937 & 0\\\\ GMRTEN1~J160335.2$+$555419 & 16:03:35.24 & $+$55:54:19.5 & 0.762 & 105 & 0.984 & 0.108 & 6412 & 6486 & 0\\\\ GMRTEN1~J160335.2$+$540515 & 16:03:35.25 & $+$54:05:15.3 & 38.494 & 383 & 50.351 & 0.961 & 6528 & 2125 & 0\\\\ GMRTEN1~J160336.5$+$555147 & 16:03:36.55 & $+$55:51:47.5 & 0.500 & 79 & 0.662 & 0.082 & 6408 & 6385 & 0\\\\ GMRTEN1~J160336.9$+$544120 & 16:03:36.90 & $+$54:41:20.7 & 0.532 & 68 & 0.853 & 0.092 & 6480 & 3568 & 0\\\\ GMRTEN1~J160337.6$+$545944 & 16:03:37.64 & $+$54:59:44.2 & 0.547 & 80 & 0.530 & 0.077 & 6457 & 4303 & 0\\\\ GMRTEN1~J160338.7$+$554348 & 16:03:38.77 & $+$55:43:48.1 & 0.771 & 80 & 1.743 & 0.135 & 6404 & 6065 & 0\\\\ GMRTEN1~J160339.0$+$545943 & 16:03:39.05 & $+$54:59:43.8 & 0.500 & 80 & 0.415 & 0.065 & 6449 & 4303 & 0\\\\ GMRTEN1~J160339.3$+$551352 & 16:03:39.37 & $+$55:13:52.0 & 0.516 & 78 & 0.414 & 0.063 & 6432 & 4868 & 0\\\\ GMRTEN1~J160339.6$+$542953 & 16:03:39.68 & $+$54:29:53.2 & 0.701 & 66 & 0.647 & 0.078 & 6476 & 3110 & 0\\\\ GMRTEN1~J160339.7$+$550600 & 16:03:39.75 & $+$55:06:00.4 & 0.584 & 80 & 0.418 & 0.069 & 6438 & 4554 & 0\\\\ GMRTEN1~J160340.1$+$545127 & 16:03:40.10 & $+$54:51:27.4 & 0.711 & 79 & 0.639 & 0.079 & 6451 & 3972 & 0\\\\ GMRTEN1~J160340.1$+$550544 & 16:03:40.15 & $+$55:05:44.9 & 0.495 & 81 & 0.392 & 0.069 & 6436 & 4543 & 0\\\\ GMRTEN1~J160340.8$+$554325 & 16:03:40.83 & $+$55:43:25.8 & 1.447 & 81 & 4.633 & 0.210 & 6392 & 6050 & 0\\\\ GMRTEN1~J160341.2$+$552611 & 16:03:41.27 & $+$55:26:11.8 & 0.475 & 68 & 0.460 & 0.065 & 6408 & 5361 & 0\\\\ GMRTEN1~J160341.5$+$552205 & 16:03:41.59 & $+$55:22:05.4 & 0.638 & 69 & 0.614 & 0.080 & 6411 & 5197 & 0\\\\ GMRTEN1~J160341.6$+$553203 & 16:03:41.66 & $+$55:32:03.4 & 0.462 & 69 & 0.365 & 0.062 & 6400 & 5595 & 0\\\\ GMRTEN1~J160343.1$+$540324 & 16:03:43.13 & $+$54:03:24.6 & 0.908 & 126 & 0.556 & 0.092 & 6483 & 2050 & 0\\\\ GMRTEN1~J160344.8$+$540320 & 16:03:44.84 & $+$54:03:20.4 & 1.168 & 118 & 0.856 & 0.110 & 6473 & 2047 & 0\\\\ GMRTEN1~J160345.8$+$553021 & 16:03:45.84 & $+$55:30:21.4 & 0.381 & 60 & 0.370 & 0.065 & 6378 & 5527 & 0\\\\ GMRTEN1~J160345.8$+$554238 & 16:03:45.88 & $+$55:42:38.3 & 0.653 & 76 & 2.566 & 0.156 & 6365 & 6018 & 0\\\\ GMRTEN1~J160346.5$+$552855 & 16:03:46.56 & $+$55:28:55.8 & 0.577 & 73 & 0.510 & 0.076 & 6375 & 5469 & 0\\\\ GMRTEN1~J160346.6$+$550826 & 16:03:46.62 & $+$55:08:26.0 & 0.525 & 70 & 0.416 & 0.065 & 6396 & 4650 & 0\\\\ GMRTEN1~J160348.3$+$542626 & 16:03:48.37 & $+$54:26:26.4 & 2.328 & 80 & 2.660 & 0.133 & 6429 & 2970 & 0\\\\ GMRTEN1~J160348.6$+$550124 & 16:03:48.61 & $+$55:01:24.3 & 0.477 & 78 & 0.257 & 0.050 & 6392 & 4368 & 0\\\\ GMRTEN1~J160349.2$+$554243 & 16:03:49.21 & $+$55:42:43.9 & 3.349 & 85 & 4.224 & 0.163 & 6346 & 6021 & 0\\\\ GMRTEN1~J160350.0$+$545634 & 16:03:50.06 & $+$54:56:34.6 & 0.435 & 67 & 0.371 & 0.058 & 6389 & 4175 & 0\\\\ GMRTEN1~J160350.4$+$550717 & 16:03:50.43 & $+$55:07:17.2 & 0.597 & 78 & 0.626 & 0.083 & 6376 & 4603 & 0\\\\ GMRTEN1~J160350.5$+$541302 & 16:03:50.51 & $+$54:13:02.7 & 0.770 & 80 & 0.535 & 0.068 & 6430 & 2435 & 0\\\\ GMRTEN1~J160351.1$+$555644 & 16:03:51.13 & $+$55:56:44.4 & 3.996 & 120 & 7.391 & 0.287 & 6321 & 6581 & 0\\\\ GMRTEN1~J160351.2$+$552346 & 16:03:51.24 & $+$55:23:46.1 & 0.519 & 80 & 0.640 & 0.072 & 6354 & 5262 & 0\\\\ GMRTEN1~J160351.3$+$540609 & 16:03:51.38 & $+$54:06:09.9 & 0.951 & 110 & 0.843 & 0.110 & 6432 & 2159 & 0\\\\ GMRTEN1~J160352.6$+$550012 & 16:03:52.62 & $+$55:00:12.8 & 0.510 & 71 & 0.592 & 0.076 & 6370 & 4320 & 0\\\\ GMRTEN1~J160352.8$+$542943 & 16:03:52.80 & $+$54:29:43.5 & 2.258 & 82 & 2.422 & 0.124 & 6400 & 3101 & 0\\\\ GMRTEN1~J160353.2$+$550309 & 16:03:53.29 & $+$55:03:09.9 & 0.570 & 74 & 0.500 & 0.072 & 6363 & 4438 & 0\\\\ GMRTEN1~J160355.5$+$553844 & 16:03:55.52 & $+$55:38:44.9 & 0.463 & 73 & 0.242 & 0.050 & 6314 & 5861 & 0\\\\ GMRTEN1~J160356.2$+$550126 & 16:03:56.20 & $+$55:01:26.7 & 0.472 & 77 & 0.380 & 0.063 & 6348 & 4369 & 0\\\\ GMRTEN1~J160356.2$+$550415 & 16:03:56.25 & $+$55:04:15.3 & 0.506 & 83 & 0.451 & 0.068 & 6345 & 4481 & 0\\\\ \\hline \\end{tabular} \\end{table*}  \\subsection{Comparison with other surveys} In order to test the positional accuracy of our catalogue, we paired it with the 393 objects in the VLA 1.4~GHz source catalogue of \\citet{Ciliegi99}, using a pairing radius of 6~arcsec.  The VLA survey covers only $\\sim25$~per~cent of our 610~MHz observations. Fig.~\\ref{fig:deltaRADEC} shows the position offset of the 263 matched sources compared with their VLA counterparts -- the distribution of offsets is approximately Gaussian, with mean offset in Right Ascension of $-0.9$~arcsec, and $-0.1$~arcsec in Declination.  The standard deviations of the distribution are 0.9 and 0.7~arcsec respectively.  \\begin{figure} \\includegraphics[width=8cm]{deltaRADEC.ps} \\caption{Source positions in the GMRT catalogue relative to the positions found in the VLA catalogue of \\citet{Ciliegi99}, for unique matches within 6~arcsec.  Offsets are distributed in a Gaussian fashion in RA and DEC, and the ellipse corresponding to 1$\\sigma$ for the distribution is shown.} \\label{fig:deltaRADEC} \\end{figure}  We repeated this analysis with data from the FIRST survey \\citep{Becker95}.  The whole of our 610~MHz survey region is covered by FIRST, and even with the reduced sensitivity of FIRST ($\\sim$150~$\\mu$Jy noise), 504 pairs are found within 6~arcsec.  We again find Gaussian-distributed position errors, with RA offset of $-0.1$~arcsec, standard deviation 0.4~arcsec and DEC offset of $-0.1$~arcsec, standard deviation 0.6~arcsec.  We have not corrected the positions given in our catalogue, since it agrees closely with FIRST.  The spectral index distribution of the matched sources from both surveys is shown in Fig.~\\ref{fig:alpha}, using the integrated flux density measurements.  The distribution peaks around $\\alpha=0.7$, where $\\alpha$ is defined such that the flux density $S$ scales with frequency $\\nu$ as $S = S_{0}\\nu^{-\\alpha}$.  \\begin{figure} \\includegraphics[width=8cm]{Alpha.eps} \\caption{Radio spectral index $\\alpha$ between 610~MHz and 1.4~GHz, for sources in the VLA ELAIS-N1 catalogue of Ciliegi et al.\\ (solid lines) and in the FIRST catalogue (dashed lines).} \\label{fig:alpha} \\end{figure}  Fig.~\\ref{fig:AlphaFlux} shows the spectral index distribution for all sources with matches in the GMRT and FIRST catalogues (black diagonal crosses), and sources in the GMRT and \\citet{Ciliegi99} catalogues (red upright crosses).  There are significant biases in Fig.~\\ref{fig:AlphaFlux}, due to the varying sensitivity levels of the three surveys.  In the region covered by \\citet{Ciliegi99}, the 610~MHz completeness level is 360~$\\mu$Jy, shown by the black dotted line.  The limiting spectral indices for sources at the sensitivity levels of the 1.4~GHz surveys are shown (FIRST -- solid black line, Ciliegi -- dashed red line).  In order to look for variations in the source population, we calculate the mean and median spectral indices for sources with detections in the Ciliegi catalogue, with 610~MHz flux density between 500~$\\mu$Jy and 1~mJy -- the point at which the turnover in source counts becomes visible (see Section~\\ref{sec:sourcecounts}) -- and above 1~mJy.  The mean values of $\\alpha$ are $0.22\\pm0.09$ and $0.45\\pm0.04$ respectively, and the median values are $0.36\\pm0.12$ and $0.56\\pm0.04$.  There are 48 and 168 sources in the two flux density bins.  The bias against steep-spectrum sources at low flux densities (which is visible in Fig.~\\ref{fig:AlphaFlux}) means that, for a source with 610~MHz flux density of 500~$\\mu$Jy, the largest value of $\\alpha$ that would be detectable is $\\sim$0.8 and so this apparent flattening at fainter flux densities may simply be due to sample bias.  However, \\citet{Bondi07} also find significantly flatter spectral indices for fainter radio sources, again comparing 610~MHz and 1.4~GHz data, and attribute this to the emergence of a population of low-luminosity AGNs -- see Section~\\ref{sec:sourcecounts} for more details.  \\begin{figure} \\includegraphics[width=8cm]{AlphaFlux.eps} \\caption{The variation in spectral index $\\alpha$ with 610~MHz flux density.  Black diagonal crosses represent sources in the GMRT and FIRST catalogues, with the solid black line showing the limiting spectral index that could be detected, given the respective sensitivity levels.  Red upright crosses represent sources in the corresponding GMRT and Ciliegi et al.\\ catalogues, with the dashed red line showing the limit on $\\alpha$.  The 610~MHz flux density limit is shown by the dotted black line.} \\label{fig:AlphaFlux} \\end{figure}"
    },
    "0710/0710.0506_arXiv.txt": {
        "abstract": "Using a combination of deep MID-IR observations obtained by IRAC, MIPS and IRS on board Spitzer we investigate the MID-IR properties of Lyman Break Galaxies (LBGs) at z$\\sim$3, establish a better understanding of their nature and attempt a complete characterisation of the population. With deep mid-infrared and optical observations of $\\sim$1000 LBGs covered by IRAC/MIPS and from the ground respectively, we extend  the spectral energy distributions (SEDs) of the LBGs to mid-infrared. Spitzer data reveal for the first time that the mid-infrared properties of the population are inhomogeneous ranging from those with marginal IRAC detections to those with bright rest-frame near-infrared colors and those detected at 24$\\mu$m MIPS band revealing the newly discovered population of the Infrared Luminous Lyman Break Galaxies (ILLBGs). To investigate this diversity, we examine the photometric properties of the population and we use stellar population synthesis models to probe the stellar content of these galaxies. We find that a fraction of LBGs have very red colors and large estimated stellar masses $M_{\\ast}$$>$5$\\times$$10^{10}$$M_{\\odot}$. We discuss the link between these LBGs and submm-luminous galaxies and we report the detection of rest frame 6.2 and 7.7 $\\mu$m emission features arising from Polycyclic Aromatic Hydrocarbons (PAH) in the Spitzer/IRS spectrum of an infrared-luminous Lyman break galaxy at z=3.01. ",
        "introduction": "Observation and study of high-redshift galaxies is essential to constrain the history of galaxy evolution and give us a systematic and quantitative picture of galaxies in the early universe, an epoch of rigorous star and galaxy formation. Large samples of high-z galaxies that have recently become available, play a key role to that direction and have revealed a zoo of different galaxy populations at z. There are various techniques for detecting high-z galaxies involving observations in wavelengths that span from optical to far-IR. Among the various methods the Lyman break dropout technique (\\cite{Steidel93}), sensitive to the presence of the 912{\\AA} break, is designed to select z$\\sim$ 3 galaxies. LBGs constitute at the moment the largest galaxy population at z$\\sim$3 (\\cite{Steidel03}). With observations spanning from X-rays (eg. \\cite{Nan02}) to near-infrared (\\cite{Shapley03}) there has been a considerable progress into understanding the nature of population, but to fully characterize their properties (such as stellar mass, dust content, link to other z$\\sim$3 populations) observations of longer wavelengths are required.  With the advent of Spitzer Space Telescope (\\cite{Werner04}) we have access to longer wavelengths. IRAC bands (3.6, 4.5, 5.8, 8.0$\\mu$m) are crucial as they trace the rest-frame near infrared luminosities for galaxies at 0.5$<$z$<$5, (where the bulk of the stellar mass of a galaxy radiates) while MIPS (24, 70, 160 $\\mu$m) and IRS (5.3--40$\\mu$m) provide an insight into the interstellar medium of the population as they are sensitive to PAH features and dust re-radiation.  In this study we use IRAC and MIPS data covering $\\sim$1000 and 244 LBGs respectively, lying on the fields Q1422+2309 (Q1422), DSF2237a,b (DSF), Q2233+1341 (Q2233), SSA22a,b (SSA22), B20902+34 (B0902), QSOHS1700+6416 (Q1700), Extended Groth Strip (EGS) and Hubble Deep Field North (HDFN). Those LBGs have previously beeen identified from their optical colours by \\cite{Steidel03}. In section 2 we search for mid-infrared counterparts of the LBGs, extend their SEDs to mid-infrared and investigate their mid-infrared colours as well as their physical properties such as stellar mass and dust content. In section 3 we examine the possible link between the IRAC/MIPS bright LBGs and the SMGs while in Section 4 using data obtained by IRS, we report the detection of PAH features arising from the mid-inrared spectrum of an ILLBG at z=3.01. In Section 5 we summurize the results of this Spitzer view of LBGs. ",
        "conclusions": "The advent of Spitzer has dramatically improved our understanding of the LBGs. Using data obtained by IRAC/MIPS/IRS on board Spitzer we have reached the following conclusions for the population of the LBGS:  \\begin{itemize} \\item IRAC colors have revealed the diversity of LBGs ranging from those with marginal detection in IRAC bands and R-[3.6]$<$1.5 colors, to those that have bright in IRAC bands and exhibit R-[3.6]$>$1.5 colors. \\item LBGs detected at 8$\\mu$m have redder R-[3.6] colors and on average are more massive, suffer more obscuration and have relatively older stellar populations when compared to the rest of the population. \\item A fraction of about $\\sim$5\\% of the LBGs do have dust as evidenced by MIPS 24$mu$m detections, and are classified as ILLBGs. Those LBGs share many properties in common with the SMGs and preliminary results show that they can be detected at submm bands. It can therefore be suggested that a link between these two populations must exist. \\item Strong PAH features arising from the mid-infrared spectrum of an ILLBG at z=3.01 indicates the existence of dust in the interstellar medium of LBGs and suggest that the emission is dominated by star formation rather than an AGN. \\end{itemize}   \\begin{figure} \\centering \\includegraphics[width=10cm,height=4cm]{magdis_fig2.eps} \\caption {\\small{IRS spectrum of EGS20 J1418+5236. The dashed line is the M82 SED shifted to $z=3.01$. The spectrum has remarkably similar 6.2 and 7.7 $\\mu$m PAH emission-feature strength and shape to those of M82. We cross-correlated the IRS spectrum with that of M82 to derive a redshift of $z =3.01$$\\pm$$0.016$.}} \\end{figure}"
    },
    "0710/0710.5609_arXiv.txt": {
        "abstract": "{} {We present the results of a spectroscopic analysis on three young embedded sources  (HH26 IRS, HH34 IRS and HH46 IRS) belonging to different star-forming regions and displaying well developed jet structures. The aim is to investigate the source accretion and ejection properties and their connection.} {We used VLT-ISAAC near-IR medium resolution ($R\\sim9000$) spectra ($H$ and $K$ bands) to derive, in a self-consistent way,  parameters like the star luminosity, the accretion luminosity and the mass accretion rate. Mass ejection rates have also been estimated from the analysis of different emission features.} {The spectra present several emission lines but no photospheric features in absorption, indicating a large veiling in both $H$ and $K$ bands. In addition to features commonly observed in jet driving sources ([\\ion{Fe}{ii}], H$_2$, \\ion{H}{i}, CO), we detect a number of emission lines due to permitted atomic transitions, such as \\ion{Na}{i} and \\ion{Ti}{i} that are only 2-5 times weaker than the Br$\\gamma$ line. Some of these features remain unidentified. Emission from \\ion{Na}{i} 2.2$\\mu$m doublet is observed along with CO(2-0) band-head emission, indicating a common origin in an inner gaseous disc heated by accretion. We find that accretion provides about 50\\% and 80\\% of the bolometric luminosity in HH26 IRS and HH34 IRS, as expected for accreting young objects. Mass accretion and loss rates spanning $10^{-6}$--$10^{-8}$ M$_{\\sun}$\\,yr$^{-1}$ have been measured. The derived $\\dot{M}_\\mathrm{loss}/\\dot{M}_\\mathrm{acc}$ is $\\sim$0.01 for HH26 IRS and HH34 IRS, and $>$0.1 for HH46 IRS.  These numbers are in the range of values predicted by MHD jet launching models and found in the most active classical T Tauri stars.} {Comparison with other spectroscopic studies performed on Class Is seems to indicate that Class Is actually having accretion-dominated luminosities are a limited number. Although the analysed sample is small, we can tentatively define some criteria to characterise such sources: they have $K$-band veiling larger than 2 and in the majority of the cases present IR features of CO and \\ion{Na}{i} in emission, although these do not directly correlate with the accretion luminosity. Class Is with massive jets have high $L_{\\mathrm{acc}}/L_{\\mathrm{bol}}$ ratios but not all the identified accretion-dominated objects present a jet. As suggested by the SEDs of our three objects, the accretion-dominated objects could be in an evolutionary transition phase between Class 0 and I. Studies of the kind presented here but on larger samples of possible candidates should be performed in order to test and refine these criteria.} ",
        "introduction": "The process of mass accretion accompanying the formation of solar type stars is always associated with mass ejection in form of collimated jets, that extend from few AU up to parsecs distance from the exciting source. According to an established class of models \\citep{koenigl00,casse00}, accretion and ejection are indeed intimately related through the presence of a magnetised accretion disc: the jets carry away the excess angular momentum, so that part of the  disc material can move toward the star. The efficiency of this process is measured by the ratio between the mass ejection and mass accretion rates, and depends on the  jet acceleration mechanism at work. Measurements of such an efficiency have been so far obtained only for classical  T Tauri stars, whose accretion properties are rather well studied through the observation of the excess continuum emission at optical and UV wavelengths Values of $\\dot{M}_\\mathrm{loss}/\\dot{M}_\\mathrm{acc}$ in the range 1-10\\% have been found by different studies \\citep[e.g.][]{hartigan95,woitas05,ferreira06}. It is nevertheless important to test accretion/ejection models in young sources at earlier stages of evolution, when  accretion dominates the energetics of the system and thus the mechanism to extract angular momentum is expected to be more efficient. To this aim, an interesting sample of objects  are Class I sources, i.e. the class of embedded stars characterised by a steeply rising IR spectral energy distribution (SED) between 2 and 10 $\\mu$m and usually considered younger than visible T Tauri stars (the Class II sources). However, the high extinction pertaining to these objects strongly limits the measurement of their stellar and accretion properties, needed to prove that they are indeed in a phase of higher accretion with respect to Class II sources. The general assumption so far applied has been that most of the bolometric luminosity of Class I objects is due to accretion.  Recently, however, thanks to the use of high dispersion sensitive instrumentation, it has become possible to define the stellar properties of small samples of Class I stars through their weak photospheric lines detected in the optical scattered light \\citep{white04} and in the near-IR direct emission \\citep{greene02,nisini05a,doppmann05}. Such studies have shown that the characteristics of Class I objects vary in fact significantly. In particular, the accretion luminosity may span from few percent up to 80\\% of the bolometric luminosity; these findings show that  not all sources defined as Class I are indeed actively accreting objects and suggest that a classification based on different criteria is indeed required.  \\citet{white04}, notably, derived from their analysis of the scattered light spectra of Taurus-Auriga sources, that the average $\\dot{M}_\\mathrm{loss}/\\dot{M}_\\mathrm{acc}$ for Class I of their sample is larger than for Class II and close to unity. They interpret this result as due to an observational bias induced by the effect of disc orientation in their optical spectra; if the Class I are seen prevalently edge-on,  then the extended region emitting the forbidden lines from which they estimate $\\dot{M}_\\mathrm{loss}$ is seen more directly than the obscured stellar photosphere. It is clear that such kind of biases can be minimised by performing spectroscopic observations directly in the IR, where features originating in the photosphere, in the accretion region and in the jet can be simultaneously detected by instrumentation which is sensitive enough.  In this framework, we report here  the results of near-IR spectroscopic observations at medium resolution of three embedded sources (HH34 IRS, HH26 IRS and HH46 IRS)  displaying well developed jet structures and having a spectral index between 2 and 10\\,$\\mu$m, typical of Class I objects. We have derived accretion and ejection parameters of these sources through the analysis of the different features detected on the spectra. The goal is to study how much of their energy is due to accretion and to investigate the efficiency of the ejection mechanism.  We describe the sample and the observations in Sec. 2 and report the results in Sec. 3; in Sec. 4 we present the procedures applied for the analysis of the data and infer the physical properties of the objects and their jets. These results are then discussed in Sec. 5, where a comparison with similar sources analysed in previous studies will be also made. Main conclusions of our work are summarised in Sec. 6.  \\begin{table*}[!t]  \\begin{center} \\caption[]{The observed targets and their main observational properties.} \\begin{tabular}{l c c | c | c c | c c | c c c c} \\hline \\hline Source         & R.A.(2000)\t\t\t&DEC.(2000)\t\t&$m_\\mathrm{J}$& \\multicolumn{2}{|c|}{$m_\\mathrm{H}$}\t \t& \\multicolumn{2}{c|}{$m_\\mathrm{K}$}\t& $D$\t\t& $\\alpha^{(a)}$\t\t& $L_\\mathrm{bol}$\t\t\\\\ &\t\t\t\t\t\t&\t\t\t\t&(mag)\t\t\t\t\t&\\multicolumn{2}{|c|}{(mag)}&\\multicolumn{2}{|c|}{(mag)}&(pc)\t&\t&\t(L$_{\\sun}$)\t\\\\ &\t\t\t\t\t\t&\t\t\t\t&2MASS$^{(1)}$&2MASS$^{(1)}$&This work$^{(1)}$&2MASS$^{(b)}$&This work$^{(b)}$&\t&\t&\t\t\\\\ \\hline HH 26 IRS\t& 05 46 03.9\t& -00 14 52\t\t\t&16.77& 14.07 &14.6  & 11.88 & 12.3 \t\t& 450 \t\t& 2.01 \t\t\t& 4.6--9.2\t\\\\ HH 34 IRS\t& 05 35 29.9\t& -06 26 58\t \t\t&15.06& 13.60 &13.5 \t& 12.38 & 12.4 \t\t& 460 \t\t& 1.14 \t\t\t& 12.4--19.9\t \\\\ HH 46 IRS\t& 08 25 43.9\t& -51 00 36\t \t\t&14.20& 12.88 &14.8 \t& 12.72 & 13.4 \t\t& 450 \t\t& 1.96 \t\t\t& $<$15.0$^{(c)}$ \\\\ \\hline \\end{tabular} \\label{targets} \\end{center} \\vspace{0.2 cm} \\small{Notes. (a) The spectral index $\\alpha=d\\mathrm{Log}(\\lambda F_{\\lambda})/d\\mathrm{Log}(\\lambda)$ is calculated between 2 and 10 $\\mu$m. (b) The magnitude values in the $H$ and $J$ bands have been estimated from the calibrated spectra. (c) Total luminosity: the source is a binary (see text for details). References. (1) from the 2MASS catalogue \\citep{2mass}.}   \\end{table*} ",
        "conclusions": "We have investigated the accretion and ejection properties of three embedded sources (HH26 IRS, HH34 IRS, HH46 IRS) showing prominent jet-like structures. To this aim we have analysed their medium resolution (R $\\sim$ 9000) near IR spectra acquired with VLT-ISAAC. The main results we obtained can be summarised as follows: \\begin{itemize}  \\item The bolometric luminosity and SEDs of the three sources have been revised on the basis of Spitzer and recent sub-mm observations, and theoretical radiative transfer models available in the literature. It turns out that the sources are probably very young, in a transition phase between Class 0 and I.  \\item The spectra of the three sources show important differences in their characteristics: in fact, the number and the absolute and relative intensity of the observed emission features (associated both with the accretion region and the jet environment) vary among the sources. In particular, we point out that there is no clear relationship between the presence of the jet and the spectral accretion signatures detected. Moreover, the spectra show no sign of absorption features, indicating large amount of veiling which in turn suggests the presence of warm dusty envelopes around the sources, heated by the active ongoing accretion. \\item In two of our sources (HH34 IRS and HH26 IRS) we find a  Br$\\gamma$/\\ion{Na}{i} ratio much larger (by a factor of three or more) than that observed in other Class Is or T Tauri stars. Conversely, the ratio between the CO(2-0) 2.3$\\mu$m overtone emission and the \\ion{Na}{i} 2.20$\\mu$m doublet, shows the opposite trend, i.e. is smaller in the objects with jets. This may indicate the presence of massive discs around the jet sources, characterised by large amounts of warm gas in neutral state.  \\item We consistently derive $A_K$, $L_*$ and $L_{\\mathrm{acc}}$ assuming the three sources on the birthline and adopting the relationship between Br$\\gamma$ luminosity and accretion luminosity derived by \\citet{muzerolle98}. In the case of HH34 IRS and HH26 IRS,  accretion largely contributes to the total energy, as expected for young sources in the main accretion phase. In HH46 IRS, we estimate an $L_{\\mathrm{acc}}/L_{\\mathrm{bol}}$ ratio of the order of 0.2, but the real value largely depends on how the bolometric luminosity is shared among the two binary components. Taking into account $K$-band veiling values $r_K$ greater than the lower limits inferred from the spectra would lead in general to higher $A_K$, $L_{\\mathrm{acc}}$ and $\\dot{M}_{acc}$.  \\item Mass accretion and loss rates span (including errors) in the range $10^{-8}$--$10^{-6}$ M$_{\\sun}$\\,yr$^{-1}$. The derived $\\dot{M}_{loss}/\\dot{M}_{acc}$ ratio is $\\sim$0.01-0.03 for HH26 IRS and HH34 IRS, and $>$0.1 for HH46 IRS.  These numbers are in the range of values usually predicted by models and found in the more active classical T Tauri stars.  \\item Comparing the results found in this work with other spectroscopic studies recently performed on Class I sources, we conclude that the number of Class Is actually having accretion-dominated luminosities (Accretion-Dominated Young Objects, ADYOs) could be limited. From the properties inferred in a small sample of objects we can tentatively define some criteria to characterise such sources: ADYOs have all $K$-band veiling larger than 2 and in the majority of the cases present (in addition to \\ion{H}{i}) IR features of CO and \\ion{Na}{i} in emission, although these latter do not directly correlate with the accretion luminosity. Class Is with massive jets have high $L_{\\mathrm{acc}}/L_{\\mathrm{bol}}$ ratios but not all the identified ADYOs present a jet. The SEDs of our small sample of three objects, suggest that accretion-dominated sources could be in an evolutionary phase in transition between Class 0 and I.  Of course, studies of the kind presented here but carried out on larger samples of possible candidates should be performed in order to test and refine the tentative criteria that we have just mentioned.  \\end{itemize}"
    },
    "0710/0710.0440_arXiv.txt": {
        "abstract": "We examine the ability of a future X-ray observatory, with capabilities similar to those planned for the Constellation-X or X-ray Evolving Universe Spectroscopy (XEUS) missions, to constrain dark energy via measurements of the cluster X-ray gas mass fraction, $f_{\\rm gas}$. We find that $f_{\\rm gas}$ measurements for a sample of $\\sim500$ hot ($kT\\gsim5$keV), X-ray bright, dynamically relaxed clusters, to a precision of $\\sim 5$ per cent, can be used to constrain dark energy with a Dark Energy Task Force (DETF; Albrecht et al. 2006) figure of merit of $15-40$, with the possibility of boosting these values by 40 per cent or more by optimizing the redshift distribution of target clusters. Such constraints are comparable to those predicted by the DETF for other leading, planned `Stage IV' dark energy experiments.  A future $f_{\\rm gas}$ experiment will be preceded by a large X-ray or SZ survey that will find hot, X-ray luminous clusters out to high redshifts.  Short `snapshot' observations with the new X-ray observatory should then be able to identify a sample of $\\sim 500$ suitably relaxed systems. The redshift, temperature and X-ray luminosity range of interest has already been partially probed by existing X-ray cluster surveys which allow reasonable estimates of the fraction of clusters that will be suitably relaxed for $f_{\\rm gas}$ work to be made; these surveys also show that X-ray flux contamination from point sources is likely to be small for the majority of the targets of interest. Our analysis uses a Markov Chain Monte Carlo method which fully captures the relevant degeneracies between parameters and facilitates the incorporation of priors and systematic uncertainties in the analysis. We explore the effects of such uncertainties for scenarios ranging from optimistic to pessimistic. We conclude that the $f_{\\rm gas}$ experiment offers a competitive and complementary approach to the best other large, planned dark energy experiments. In particular, the $f_{\\rm gas}$ experiment will provide tight constraints on the mean matter and dark energy densities, with a peak sensitivity for dark energy work at redshifts midway between those of supernovae and baryon acoustic oscillation/weak lensing/cluster number counts experiments. In combination, these experiments should enable a precise measurement of the evolution of dark energy. ",
        "introduction": "\\label{introduction}  In the early 1990s, measurements of the baryonic mass fraction in X-ray luminous galaxy clusters provided compelling evidence that we live in a low density Universe. Under the assumption that large clusters provide approximately fair samples of the matter content of the Universe, X-ray observations require that the mean matter density, $\\Omega_{\\rm m}$, is significantly less than the critical value, with a best-fit value $\\Omega_{\\rm m}\\sim 0.2-0.3$ \\citep[e.g.][]{White:91,Fabian:91, Briel:92, White:93, David:95, White:95, Evrard:97, Mohr:99, Ettori:99, Roussel:00, Grego:01, Allen:02, Allen:04, Allen:07, Ettori:03, Sanderson:03, Lin:03, LaRoque:06}. When combined with the expectation from inflation models, later confirmed by Cosmic Microwave Background (CMB) studies \\citep[][and references therein]{Bennett:03, Spergel:03}, that the Universe should be close to spatially flat, X-ray results on the cluster baryon mass fraction quickly lead to the suggestion that the mass-energy density of the Universe may be dominated by a cosmological constant \\citep[e.g.][]{White:93}.  The first direct evidence for late-time cosmic acceleration, as would be produced by a sizeable cosmological constant, was provided in the late 1990s by \\cite{Riess:98} and \\cite{Perlmutter:98} based on measurements of the light curves of type Ia supernovae (SNIa). Since then, larger SNIa data sets \\citep{knop:03, Riess:04, Astier:06, Riess:07, WoodVasey:07, Davis:07} and an increasingly wide array of other, complementary experiments have confirmed and improved upon this striking measurement. The combination of CMB data from the Wilkinson Microwave Anisotropy Probe (WMAP) \\citep{Spergel:03, Spergel:06, Dunkley:08} with large scale structure (LSS) data from the Sloan Digital Sky Survey (SDSS) \\citep{Eisenstein:05, Percival:07} and/or 2dF Galaxy Redshift Survey (2dFGRS) \\citep{Cole:05} provides powerful evidence for dark energy. The cross-correlation of CMB and LSS fluctuations reveals the effects of dark energy on the Integrated Sachs-Wolfe effect \\citep{Scranton:03, Fosalba:03, Rassat:07}. Measurements of the amplitude and evolution of matter fluctuations using X-ray galaxy clusters \\citep{Borgani:01,Reiprich:02, Allen:03, Schuecker:03, Voevodkin:04, Henry:04, Mantz:07}, optically-selected clusters \\citep{Gladders:06, Rozo:07}, Lyman-$\\alpha$ forest data \\citep{Viel:04, Seljak:04}, and weak lensing \\citep{vanWaerbeke:05, Jarvis:05, Hoekstra:06, Benjamin:07}, also provide important, powerful confirmation of the new, standard cosmological paradigm: a universe in which the main mass and energy components are dark matter and dark energy, and where dark energy drives the current acceleration. The standard model for dark energy remains the cosmological constant, which is mathematically equivalent to vacuum energy. In principle, however, cosmic acceleration could be driven by either dark energy or a modification to the laws of gravity on cosmological scales \\citep[see][for an extensive review]{Copeland:06}.  Building on the early X-ray work, \\cite{Allen:04, Rapetti:05}; and \\cite{Allen:07} showed that measurements of the evolution of the X-ray gas mass fraction, $f_{\\rm gas}$, in the largest, dynamically relaxed galaxy clusters provides a further powerful, complementary approach for studying dark energy. As with SNIa data, $f_{\\rm gas}(z)$ measurements probe the redshift-distance relation; whereas the peak SNIa luminosity varies as the square of the distance, $f_{\\rm gas}$ measurements vary as distance, $d^{1.5}$. \\citep[The distance dependance derives from the way in which $f_{\\rm gas}$ values are determined from the observed X-ray temperature and surface brightness data;][]{Allen:07} In combination with the tight constraint on $\\Omega_{\\rm m}$ provided by the normalization of the $f_{\\rm gas}(z)$ curve, under the assumption of fair matter samples, the $f_{\\rm gas}(z)$ data contain sufficient information to break the degeneracy between $\\Omega_{\\rm m}$ and the dark energy equation of state, $w$, in the distance equations. The additional combination of $f_{\\rm gas}$ and CMB data breaks other important degeneracies between parameters in cosmological analyses \\citep{Rapetti:05, Allen:07}.  \\cite{Allen:07} show that the current constraints on dark energy from the $f_{\\rm gas}$ experiment are of comparable precision to other leading techniques, and are robust under the inclusion of conservative systematic allowances, e.g. relaxing the requirement for exact hydrostatic equilibrium and allowing for moderate redshift evolution in the cluster baryon fraction. These authors also show that intrinsic, systematic scatter remains undetected in the current $f_{\\rm gas}$ data, despite a weighted mean statistical scatter in the individual distance measurements of only $\\sim 5$ per cent; in contrast, SNIa studies \\citep{Riess:07, Jha:07, WoodVasey:07} have established the presence of systematic scatter at the $\\sim 7$ per cent in distance measurements from the best current SNIa data.  The key to determining the nature of dark energy is to obtain precise measurements of its evolution with redshift, $z$, or scale factor, $a=1/(1+z)$.  The Dark Energy Task Force report \\cite[][hereafter DETF]{Albrecht:06} presented estimates of the constraints on dark energy parameters that should be achievable with a number of future proposed or planned dark energy experiments. In particular, the report forecasted the ability of these experiments, in combination with CMB data from the Planck satellite, to constrain a dark energy model of the form $w(a)=w_{\\rm 0}+w_{\\rm a}(1-a)$, and defined a figure of merit (hereafter FoM) to allow for easy comparison of the constraints.  In this paper, we use the same dark energy parameterization and FoM to quantify the constraining power of future $f_{\\rm gas}$ experiments, to be carried out with e.g. the Constellation-X or X-ray Evolving Universe Spectroscopy (XEUS) missions, in combination with CMB data. We show that the $f_{\\rm gas}$ experiment is likely to provide comparable constraining power to the best other, contemporary space and ground-based experiments described by the DETF. When combined, future CMB, SNIa, baryon acoustic oscillation (BAO), weak lensing, cluster number count and $f_{\\rm gas}$ experiments should provide precise, accurate constraints on $w(z)$ and allow significant progress in understanding the origin of cosmic acceleration.  The structure of this paper is as follows: in Section~\\ref{sec:de} we define the dark energy model and the FoM. In Section~\\ref{sec:simdata} we describe the simulated $f_{\\rm gas}$ and CMB data sets.  For the $f_{\\rm gas}$ data, we assume instrument characteristics appropriate for the baseline Constellation-X mission. The CMB data set approximates that expected from two years of Planck data. We also simulate a data set representative of that produced by follow-up observations of the Sunyaev-Zel'dovich effect in the clusters targeted for the $f_{\\rm gas}$ work.  Section~\\ref{analysis} describes the Markov Chain Monte Carlo (MCMC) pipeline and details of the analysis method. Our main results are presented in Section~\\ref{constraints}. Section~\\ref{conclusions} summarizes our conclusions. ",
        "conclusions": "\\label{conclusions}  \\begin{figure} \\includegraphics[width=3.2in]{Prospects_w0waP5_new.ps} \\caption{The 68 and 95 per cent confidence contours in the $w_{\\rm 0}-w_{\\rm a}$ plane determined from the $f_{\\rm gas}$+CMB data (black, dashed contours) using the default dark energy model and $5$ per cent systematic allowances.  The solid red lines show the constraints obtained from the $f_{\\rm gas}$ data alone, using priors on $\\Omega_{\\rm b}h^2$, $\\Omega_{\\rm dm}h^2$, $l_{\\rm a}$ and $h$, as described in the text (Section~\\ref{cmbede}).  The blue, dotted lines show the constraints from the $f_{\\rm gas}$ alone using priors on $\\Omega_{\\rm b}h^2$ and $h$ and assuming flatness. The figure shows how the CMB data contribute in constraining dark energy, especially at early times.} \\label{fig:cmbede} \\end{figure}   We have examined the ability of a future X-ray observatory, with capabilities similar to those planned for Constellation-X, to constrain dark energy via the $f_{\\rm gas}$ experiment. We find that $f_{\\rm gas}$ measurements for a sample of 500 hot ($kT_{2500}\\gsim5$keV), X-ray bright, dynamically relaxed clusters, with a precision of $\\sim 5$ per cent, can be used to constrain dark energy with a FoM of $15-40$. These constraints are comparable to those predicted by the DETF \\citep{Albrecht:06} for other leading, planned (DETF Stage IV) dark energy experiments. We also find that, for the $f_{\\rm gas}$ experiment, the FoM can be boosted up by at least $\\sim 40$ per cent by selecting an optimal redshift distribution of suitable clusters on which to carry out the $f_{\\rm gas}$ observations. Interestingly, the optimal redshift distribution of $f_{\\rm gas}$ measurments appears to be shifted towards low redshifts.  As discussed in the text, a future $f_{\\rm gas}$ experiment will need to be preceded by a large X-ray or SZ cluster survey that will find hot, X-ray luminous clusters out to high redshifts. A survey such as that planned with the Spectrum-RG/eROSITA mission should find several thousand of such clusters. Short `snapshot' follow-up observations of the clusters with a new, large X-ray observatory should be able to identify a sample of $\\sim 500$ suitable systems for $f_{\\rm gas}$ work. Attaining a precision of $\\sim 5$ per cent with individual $f_{\\rm gas}$ measurements should be straightforward for an observatory with characteristics similar to Constellation-X, requiring exposure times of $\\sim 20$ks on average. We note that the population of galaxy clusters in the redshift, temperature and X-ray luminosity range of interest has already been partially probed by the MACS survey \\citep{Ebeling:01}; Chandra observations of MACS clusters are used extensively in current $f_{\\rm gas}$ studies \\citep{Allen:04,LaRoque:06,Allen:07}. The low-level of X-ray flux contamination from point sources observed in MACS clusters also alleviates the requirements on the instrumental PSF for dark energy work via the $f_{\\rm gas}$ method.  In determining the predicted dark energy constraints, we have employed the same MCMC method used to analyze current data. The MCMC method encapsulates all of the relevant degeneracies between parameters and allows one to easily and efficiently incorporate priors and allowances in the analysis. We have included an array of such systematic allowances, with tolerances ranging from optimistic to pessimistic. Our technique differs from the DETF \\citep{Albrecht:06}, who use a simpler Fisher matrix approach in the prediction of dark energy constraints. Despite these differences, we have endeavored to make our calculations of the FoM (Section 2) as comparable as possible.  Benchmarking our results against those of the DETF for other, future `Stage IV' dark energy experiments i.e. large, long-term missions, we find that the $f_{\\rm gas}$ experiment should provide a comparable FoM to future ground-based SNIa (FoM=$8-22$), space-based SNIa (FoM=$19-27$), ground-based BAO (FoM=$5-55$), space-based BAO (FoM=$20-42$) and space-based cluster counting (FoM=$6-39$) experiments. Formally, the predicted FoM for the $f_{\\rm gas}$ experiment is comparable to `pessimistic' scenarios for weak lensing experiments discussed by \\cite{Albrecht:06}, although the value falls short of the most optimistic DETF weak lensing predictions.  The tight constraints on $\\Omega_{\\rm m}$ and $\\Omega_{\\rm de}$ for the $f_{\\rm gas}$ experiment will be of importance when used in combination with other techniques. Interestingly, the `pivot point' for the $f_{\\rm gas}$ experiment lies between those of the SNIa and BAO/weak lensing/cluster number count experiments, offering excellent redshift coverage in attempts to pin down the evolution of dark energy.  We conclude that the $f_{\\rm gas}$ experiment offers a powerful approach for dark energy work, which should be competitive with and complementary to the best other planned dark energy experiments."
    },
    "0710/0710.2997_arXiv.txt": {
        "abstract": "{Gas-rich dwarf galaxies are probably the closest counterparts to primeval objects we can find in the local Universe, therefore it is interesting to study their evolution in different astrophysical contexts.} {We study the effects of interstellar clouds on the dynamical and chemical evolution of gas-rich dwarf galaxies. In particular, we focus on two model galaxies similar to IZw18 and NGC1569 in comparison to models in which a smooth initial distribution of gas is assumed.} {We use a 2-D hydrodynamical code coupled with a series of routines able to trace the chemical products of SNeII, SNeIa and intermediate-mass stars.  Clouds are simulated by adding overdense regions in the computational grid, whose locations are chosen randomly and whose density profiles match observed ones.  We consider both cloud complexes put at the beginning of the simulation and a mechanism for continuous cloud formation.  The clouds are inherently dynamically coupled to the diffuse gas, and they experience heat conduction from a hot surrounding gas.} {Due to dynamical processes and thermal evaporation, the clouds survive only a few tens of Myr. Due to the additional cooling agent, the internal energy of cloudy models is typically reduced by 20 -- 40\\% compared with models of diffuse gas alone.  The clouds delay the development of large-scale outflows by mass loading, therefore helping to retain a larger amount of gas inside the galaxy.  However, especially in models with continuous creation of infalling clouds, their bullet effect can pierce the expanding supershell and create holes through which the superbubble can vent freshly produced metals.  Moreover, assuming a pristine chemical composition for the clouds, their interaction with the superbubble dilutes the gas, reducing the metallicity.  The resulting final metallicity is therefore generally lower (by $\\sim$ 0.2 -- 0.4 dex) than the one attained by diffuse models.  } {} ",
        "introduction": "\\label{intro}  Gas-rich dwarf galaxies are commonly classified into low surface-brightness dwarfs, called dwarf irregulars (dIrrs), and higher surface-brightness objects, usually called blue compact dwarf (BCD) galaxies.  These classes of galaxies tend to have low metallicities, blue colors and complex and chaotic gas phases.  A large fraction of these galaxies shows an ongoing star formation (SF) or at least hints that this process has been quenched in the recent past.  In this case, these objects are commonly referred to as {\\it starburst} galaxies and their gas consumption timescales are much shorter than the Hubble time (Kennicutt \\cite{ken98}), making this a transient phase of their evolution.  Owing to the energy released by stellar winds and supernovae (SNe), intense episodes of SF are also associated to the development of galactic winds or at least of large-scale outflows.  The broad distinction between these two phenomena is the final fate of the outwards-directed flow of gas: galactic winds generally exceed the escape velocity while outflows do not, therefore they tend to recede towards the center of the galaxy.  Clear signatures of outflows are present in NGC1705 (Hensler et al. \\cite{hen98}; Heckman et al. \\cite{hek01}), NGC1569 (Martin, Kobulnicky \\& Heckman \\cite{mkh02}), NGC3079 (Cecil et al. \\cite{cec01}), IZw18 (Martin \\cite{m96}), NGC3628 (Irwin \\& Sofue \\cite{is96}) among others.  Perhaps the best examples of large-scale outflows driven by SN feedback are at large redshifts (Pettini et al. \\cite{pet98}; Pettini et al. \\cite{pet01}).  Although it is not certain, in any of the above-mentioned objects, that the metals will definitely leave the parent galaxy, indirect hints of the ubiquity of galactic winds are given by the mass-metallicity relation (Tremonti et al. \\cite{tre04}; Dave\\'e, Finlator \\& Oppenheimer \\cite{dfo06}) and effective yields (Garnett \\cite{gar02}).  From a theoretical point of view, the study of the evolution of gas-rich dwarf galaxies through numerical simulations has been performed by several authors in the recent past.  The overall picture is that the occurrence of large-scale outflows is initially driven by the thermal pressure of a very hot, high pressurized gas and is favored by a flat distribution of the interstellar medium (ISM), which allows an easy vertical transport of material.  However, since the transport of gas along the disk is very limited, outflows are not able to eject a significant fraction of the ISM, whereas the fraction of ejected metals can be very large (D'Ercole \\& Brighenti \\cite{db99}; MacLow \\& Ferrara \\cite{mf99}; Recchi, Matteucci \\& D'Ercole \\cite{rmd}, hereafter RMD).  For NGC1569, Martin et al. (\\cite{mkh02}) derived a supersolar metal content in the galactic wind from X-ray spectra but also advocated mass-loading of it with the ISM.  Most of these studies, however, have focused on flows in homogeneous media, neglecting the multiphase nature of the ISM, although several attempts to perform multiphase hydrodynamical simulations have been made in the past, particularly using the so-called {\\it chemodynamical} approach (Theis, Burkert \\& Hensler \\cite{tbh92}; Rieschick \\& Hensler \\cite{rh00}; Hensler, Theis \\& Gallagher \\cite{hen04}).  The multiphase nature of the ISM, in particular its clumpiness, is observationally well established in dwarf galaxies (Cecil et al. \\cite{cec01}; Cannon et al. \\cite{cann05}; Leroy et al. \\cite{leroy06}) and it has a solid theoretical background with the seminal work of McKee \\& Ostriker (\\cite{mo77}).  According to this model, the ISM is composed by a cold neutral phase (representing the cores of molecular clouds), confined by a warm medium (with temperatures of the order of 10$^4$ K) and these two phases (which are in pressure equilibrium) are embedded in a hot, diluted intercloud medium (HIM), continuously produced by SN explosions and stellar winds.  Sufficiently dense clouds can pierce the HIM without being swept up, so they can become embedded therein (Vieser \\& Hensler \\cite{vh07b}).  At the interface between clouds and HIM, condensation-evaporation processes establish the final fate of the cloud and its impact on the development of a galactic wind.  In two previous papers, we have studied the dynamical and chemical evolution of model galaxies similar to IZw18 (Recchi et al. \\cite{rec04}, hereafter Paper I) and NGC1569 (Recchi et al. \\cite{rec06}, hereafter Paper II).  The main results can be briefly summarized as follows:  \\begin{itemize}  \\item most of the analyzed models develop large-scale outflows.  These outflows carry out of the galaxy mostly the chemical elements freshly produced during the most recent episodes of SF, with large escape fraction of metals with delayed production (like Fe and N).  \\item Models with very short burst(s) of SF can cool and mix the newly formed metals in a very short timescale, whereas, when the SF is more complex, most of the metals are either directly ejected outside the galaxy through galactic winds or are confined in a too hot medium, therefore cannot contribute to the chemical enrichment of the warm ionized medium observed by emission lines from the \\hii gas.  \\item Models with complex and long-lasting SF episodes reproduce the chemical composition and the abundance ratios of the above-mentioned galaxies much better than models with bursting SF.  \\end{itemize}  In this paper we simulate models with structural parameters similar to IZw18 and NGC1569.  We increase arbitrarily the gas density of some specific regions of the computational grid, in order to create a ``cloudy'' phase, and we address the question how and to which extent a ``cloudy gas phase'' alters the former results.  The clouds possess a specific density profile and can be either added at the beginning of the simulation or continuously created during the evolution of the model.  We then analyze the differences between the dynamical and chemical evolution of these models with the ones presented in Paper I and Paper II.  We point out that, at variance with the above-mentioned works, in this paper we will not specifically look for the best initial setups and the best assumptions in order to reproduce chemical and dynamical features of well-known objects.  We will just stress the main variations produced by a clumpy initial setup.  For this reason, we will also consider models which failed in Paper I and II at reproducing the observations of IZw18 and NGC1569.  The paper is organized as follows: in Sect.~\\ref{cloud} we briefly recall the evolution of a cloud embedded in a hot medium; in Sect.~\\ref{model} we present the model and the adopted assumptions in the simulations.  Results are presented in Sect.~\\ref{results_fix} (models with clouds fixed at the beginning of the simulation) and in Sect.~\\ref{results_inf} (continuous creation of clouds).  Finally, a discussion and some conclusions are drawn in Sect.~\\ref{discussion}. ",
        "conclusions": "\\label{discussion}  In this paper we have computed the chemical and dynamical evolution of model galaxies, with structural parameters similar to IZw18 and NGC1569, but in which a complex of clouds has been added, both perturbing the initial gaseous distribution and creating clouds, at a rate which equals the SF rate, and with infall velocity of 10 km s$^{-1}$ along the polar direction.  The main focus of our work has been the comparison of these models with those presented in previous publications, in which similar setups but a smooth distribution of gas was considered.  We have seen that the clouds are subject to a variety of disruptive phenomena like evaporation (when embedded in a hot medium), formation of shocks, development of thermal instabilities (in particular the Kelvin-Helmholtz instability) and expansion due to the larger pressure compared to the surrounding interstellar medium.  The average lifetime of the clouds is therefore relatively short, depending on the cloud size (which is not constant in our simulations) but being of the order of a few tens of Myr.  In spite of their transient nature, the clouds leave a significant imprint on the dynamical and chemical evolution of dwarf galaxies.  The clouds, when they evaporate inside the superbubble, produce mass loading, increase the mean density of the cavity and, therefore, enhance the radiative losses (which are proportional to the square of the density).  This results in a significant decrease of the total thermal energy (of the order of $\\sim$ 20 -- 40\\% compared to the diffuse models, depending on the assumptions), therefore less energy to drive the development of a large-scale outflow.  On the other hand, the relative motion of supershell and clouds, in particular when the clouds infall motion is considered, can structure, pierce and create holes and fingers in the expanding supershell. These holes destroy the spherical symmetry initially present and favor the rushing out of the highly pressurized gas contained in the cavity. Therefore, in spite of the reduced thermal energy budget, the creation of large-scale outflows is not suppressed but, in most of the explored cases, only slightly delayed.  Complex structures and fingers are indeed relatively common features in galaxies showing large-scale outflows like NGC1800 (Hunter \\cite{hun96}), NGC4214 (MacKenty et al. \\cite{mack00}) or NGC1705 (Heckman et al. \\cite{hek01}).  The pressure inside the cavity is reduced compared to diffuse models, therefore in any case the total amount of ejected pristine gas is very small (smaller than in the models with smooth gas distribution) and, when averaging the size of the supershell in any direction, it turns out to be smaller than in diffuse models.  But the piercing of the supershell can lead to an ejection efficiency of freshly produced metals as high as the one attained by diffuse models.    This has, of course, important consequences on the chemical evolution of these objects.  Since the differential winds are not suppressed, the diminished thermal energy of these models does not imply an increase of metals inside the galactic regions.  On the other hand, the dilution effect of clouds plays a dominant role in determining the final metallicity of our model galaxies.  Since the clouds have primordial chemical composition, their destruction and mixing with the surrounding medium reduces the total chemical composition without altering the abundance ratios.  This produces a final metallicity $\\sim$ 0.2 -- 0.4 dex smaller than the corresponding diffuse models.  We have examined the effect of a different choice of the IMF slope and of the nucleosynthetic set of yields (in particular for what concerns intermediate-mass stars).  Flatter-than-Salpeter IMF slopes lead to an excessive production of energy, able to unbind most of the gas before the end of the simulation.  On the other hand, in models with steeper IMF the development of large-scale outflows is almost completely suppressed.  Different sets of intermediate-mass stars yields affect in particular the log(N/O) ratio.  Renzini \\& Voli (\\cite{rv81}) yields tend to overestimate the primary production of nitrogen.  When compared to the results of models implementing van den Hoek \\& Groenewegen (\\cite{vg97}) yields, the results differ by $\\sim$ 0.3 dex.  Due to the assumption of a metallicity-dependent cooling function, also the dynamics is affected by the choice of the nucleosynthetic prescriptions.  Our main results can be briefly summarized as follows:  \\begin{itemize}  \\item the clouds suffer thermal instabilities, formation of shocks and evaporation, therefore their lifetimes is limited to a few tens of Myr.  \\item In spite of that, they are able to increase the main density of the cavity, provoking a reduction of the total thermal energy by $\\sim$ 20 -- 40\\% compared with a diffuse model.  \\item The interaction clouds-supershell leads to strong structuring and piercing of the shell (in particular for models with continuous creation of infalling clouds), allowing the venting out of metals in spite of the reduced thermal energy.  The development of large-scale outflows is therefore generally delayed but the ejection efficiency of metals remains unchanged.  \\item From a chemical point of view, the effect of the clouds is to significantly reduce the total metallicity of the galaxies, without altering the abundance ratios.  \\end{itemize}"
    },
    "0710/0710.1285.txt": {
        "abstract": "We present results of long-slit spectroscopy in several positions of the Orion nebula. Our goal is to study the spatial distribution of a large number of nebular quantities, including line fluxes, physical conditions and ionic abundances at a spatial resolution of about 1$''$. In particular, we have compared the O$^{++}$ abundance determined from collisionally excited and recombination lines in 671 individual 1D spectra covering different morphological zones of the nebula. We find that protoplanetary disks (proplyds) show prominent spikes of {\\elect}([{\\nii}]) probably produced by collisional deexcitation due to the high electron densities found in these objects. Herbig-Haro objects show also relatively high {\\elect}([{\\nii}]) but probably produced by local heating due to shocks. We also find that the spatial distribution of pure recombination {\\oii} and {\\foiii} lines is fairly similar, in contrast to that observed in planetary nebulae. The abundance discrepancy factor (ADF) of O$^{++}$ remains rather constant along the slit positions, except in some particular small areas of the nebula where this quantity reaches somewhat higher values, in particular at the location of the most conspicuous Herbig-Haro objects: HH 202, HH 203, and HH 204. There is also an apparent slight increase of the ADF in the inner 40$''$ around $\\theta^1$ Ori C. We find a negative radial gradient of {\\elect}([{\\oiii}]) and {\\elect}([{\\nii}]) in the nebula based on the projected distance from $\\theta^1$ Ori C. We  explore the behavior of the ADF of O$^{++}$ with respect to other nebular quantities, finding that it seems to increase very slightly with the electron temperature. Finally, we estimate the value of the mean-square electron temperature fluctuation, the so-called {\\ts} parameter. Our results indicate that the hypothetical thermal inhomogeneities --if they exist-- should be smaller than our spatial resolution element. ",
        "introduction": "\\footnotetext[1]{Based on observations made with the 4.2m William Herschel Telescope (WHT) operated on the island of La Palma by the Isaac Newton Group in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrof\\'\\i sica de Canarias.}  \\label{intro}  The analysis of the spectrum of {\\hii} regions allows to determine the chemical composition of the ionized gas phase of the interstellar medium from the Solar Neighbourhood to the high-redshift Universe. Therefore, it stands as an essential tool for our knowledge of the chemical evolution of the Universe. In photoionized nebulae, the abundance of the elements heavier than He is usually determined from collisional excitation lines (hereinafter CELs), whose  intensity depends exponentially on the electron temperature, {\\elect}, of the gas. It was about 20 years ago when the first determinations of C$^{++}$ abundance from the intensity of the weak recombination line (hereinafter RL) of {\\cii} 4267 \\AA\\ were available for planetary nebulae (PNe). The comparison of the abundance obtained from {\\cii} 4267 \\AA\\ and from the CELs of this ion in the ultraviolet (UV) showed a difference that could be as large as a order of magnitude in some objects \\citep[e.g.][]{french83, rolastasinska94, mathisliu99}. \\citet{peimbertetal93} were the first in determinig the O$^{++}$ abundance from the very weak RLs, obtaining the same qualitative result: the abundances obtained from RLs are higher than those determined making use of CELs. This observational fact is currently known as the ``abundance discrepancy\" (hereinafter AD) problem. In the last years, our group has obtained a large dataset of intermediate and high resolution spectroscopy of Galactic and extragalactic {\\hii} regions using medium and large aperture telescopes \\citep{estebanetal02, estebanetal05, garciarojasetal04, garciarojasetal05, garciarojasetal06, garciarojasetal06b, lopezsanchezetal06}. The general result of these works is that the O$^{++}$/H$^+$ ratio calculated from RLs is between 0.10 and 0.35 dex higher than the value obtained from CELs in the same objects. The value of the AD that we usually find in {\\hii} regions is rather similar for all objects and ions and is much lower than the most extreme values found in PNe. The results for {\\hii} regions obtained by our group are fairly different that those found for PNe, and seem to be consistent with the predictions of the temperature fluctuations paradigm formulated by \\citet{peimbert67}, as it is argued in \\citet{garciarojas06} and \\citet{garciarojasesteban07}. In the presence of temperature fluctuations (parametrized by the mean square of the spatial variations of temperature, the so-called $t^{\\rm 2}$ parameter) the AD can be naturally explained because the different temperature dependence of the intensity of RLs and CELs. The existence and the origin of temperature fluctuations are still controversial problems and a challenge for our understanding of ionized nebulae. Recently, \\citet{tsamispequignot05} and \\citet{stasinskaetal07}  have proposed an hypothesis to the origin of the AD, which is based on the presence of cold high-metallicity clumps of supernova ejecta still not mixed with the ambient gas of the {\\hii} regions. This cold gas would produce most of the emission of the RLs whereas the ambient gas of normal abundances would emit most of the intensity of CELs.  Our group is interested in exploring on what variable or physical process the AD depends from different approaches. One of the most promising is based on the study of the behaviour of this magnitude at small spatial scales, something that has still not been explored in depth in nearby bright Galactic {\\hii} regions. In this paper, we make use of deep intermediate-resolution long-slit spectroscopy of the Orion nebula to study the dependence of the AD with respect to different nebular parameters: electron temperature and density, local ionization state of the gas, presence of high velocity material, and its correlation with different morphological structures (proplyds, ionization fronts, globules, Herbig-Haro objects), in {\\hii} regions.  The spatial distribution of the physical conditions in the Orion nebula has been investigated by several authors. \\citet{baldwinetal91} obtained the density and temperature distribution in 21 and 14 points, respectively, along a 5$'$ line west of $\\theta^1$ Ori C, finding a density gradient that decreases to the outskirts of the nebula and a constant {\\elect}. \\citet{walteretal92} determined electron densities and temperatures and chemical abundances for 22 regions of the Orion nebula. Using also data from the literature, these authors find radial gradients of the physical conditions, but with a positive slope in the case of the temperature determined from {\\foiii} lines. \\citet{poggeetal92} obtained Fabry-Perot images of the inner 6$'$ of the nebula covering several bright CELs and taken with an average seeing of about 1\\farcs8. Those authors present a density map obtained from the ratio of the {\\fsii} doublet confirming the presence of a density gradient that reaches its highest point immediately south-southwest of the Trapezium stars, and some localized density enhancements in the Orion bar and some Herbig-Haro objects. Very recently, \\citet{sanchezetal07} have obtained an integral field spectroscopy mosaic of an area of 5$'\\times 6'$ of the center of the Orion nebula, with a spatial resolution of 2\\farcs7. The electron density map they obtain is consistent with that obtained by \\citet{poggeetal92} but richer in substructures, some of them possibly associated to Herbig-Haro objects. \\citet{sanchezetal07} also obtain a electron temperature map (derived from the line ratio of {\\fnii} lines) that shows clear spatial variations, which rise near the Trapezium and drop to the outer zones of the nebula. However, an important drawback of the temperature map of S\\'anchez et al. is that is based on non-flux calibrated spectra and possible effects due to variations in the dust extinction distribution cannot be disregarded. \\citet{odelletal03} obtained a high spatial resolution map of the electron temperature --derived from the line ratio of {\\foiii} lines-- of a 160$''\\times 160''$ field centered at the southwest of the Trapezium. The data were obtained from narrow-band images taken with the WFPC of the  $HST$. Although they do not find a substantial radial gradient of {\\elect} in the nebula, \\citet{odelletal03} report the existence of small-scale temperature variations down to a few arcseconds compatible with the values of the temperature fluctuations parameter calculated from the AD determinations by \\citet{estebanetal04}. \\citet{rubinetal03} obtained $HST$/STIS long-slit spectroscopy at several slit positions on the Orion nebula analysing the electron temperature and density spatial profiles with resolution elements of 0\\farcs5 $\\times$ 0\\farcs5. These last authors do not find large-scale gradients of the physical conditions along the slits but a relatively large point-to-point variation and some correlation of such variations with several small-scale structures.  The spatial mapping of the AD factor has been performed in few ionized nebulae but largely for PNe. \\citet{liuetal00}, \\citet{garnettdinerstein01}, and \\citet{krabbecopetti06} have found significant differences in the spatial profiles of the O$^{++}$/H$^+$ ratio derived making use of RLs and CELs suggesting the presence of chemical inhomogeneities or additional mechanisms for producing the {\\oii} lines in these objects. \\citet{tsamisetal03} have performed the only available study so far of the spatial distribution of the AD factor in an {\\hii} region: 30 Doradus. However, considering the extragalactic nature of this object and the spatial sampling of 3\\farcs5 used by those authors, their final spatial resolution is very low --about 1pc. In any case, \\citet{tsamisetal03} find a rather constant AD factor along the zone covered with their observations, a quite different behavior than that observed in PNe.  In \\S\\S~\\ref{obsred} and~\\ref{linesel} of this paper we describe the observations, the data reduction procedure and the aperture extraction and measurement of the emission lines. In \\S~\\ref{phiscondabund} we derive the physical conditions and the ionic abundances from both kinds of lines: CELs and RLs. In \\S~\\ref{spat_prof} we present and discuss the spatial profiles of the physical conditions, line fluxes, and the abundance discrepancy factor along the slit positions. In \\S~\\ref{rad_dist} we discuss the large-scale radial distribution of some nebular properties along the nebula. In \\S~\\ref{cor_ADF} we explore possible correlations between the AD and different nebular parameters. In \\S~\\ref{t2} we address and estimate the possible temperature fluctuations inside the nebula. Finally, in \\S~\\ref{conclu} we summarize our main conclusions. ",
        "conclusions": "\\label{conclu} We have studied the spatial distribution of a large number of nebular quantities along five slit positions covering different morphological zones of the Orion nebula. The resolution element of the observations was 1\\farcs2 $\\times$ 1\\farcs03. The studied quantities were c(H$\\beta$), {\\elecd}, {\\elect}([{\\nii}]), {\\elect}([{\\oiii}]), the intensity of several selected lines (H$\\beta$, {\\cii} 4267 \\AA, {\\oii} 4649 \\AA, [{\\oiii}] 4959 \\AA, [{\\feiii}] 4881 \\AA, [{\\nii}] 5755 and 6584 \\AA, [{\\oi}] 6300 \\AA, and [{\\sii}] 6717 + 6731 \\AA), the O$^{++}$/H$^+$ ratio obtained from collisionally excited lines (CELs) and recombination lines (RLs), and the C$^{++}$/H$^+$ ratio obtained from RLs. The total number of apertures or 1D spectra extracted was 730. We have been able to determine the O$^{++}$/H$^+$ ratio from the faint RLs of this ion in a 92\\% of the apertures.  The spatial distribution of {\\elecd} shows a large range of variation --larger than an order of mag\\-ni\\-tu\\-de-- across the nebula, with local maxima associated with the position of protoplanetary disks (proplyds), Herbig-Haro objects, the Orion bar, and the brightest area of the nebula at the southwest of the Trapezium. The proplyds show quite prominent spikes of {\\elect}([{\\nii}]) and much lesser ones of {\\elect}([{\\oiii}]). This fact could be due to collisional deexcitation on the nebular lines of {\\fnii} because of the high densities of these objects. Herbig-Haro objects also show somewhat higher values of {\\elect}([{\\nii}]) but, in this case, the origin could be related to extra heating of the gas due to shock excitation. The spatial distribution of the {\\oii} 4949 \\AA\\ and {\\foiii} 4959 \\AA\\ lines is fairly similar along all the slit positions, a very different behavior to that observed in planetary nebulae. We have found that the abundance discrepancy factor (ADF) of O$^{++}$ --the difference between the O$^{++}$ abundance determined from RLs and CELs-- remains, in general, rather constant along most of the observed areas of the nebula, showing values between 0.15 and 0.20 dex. However, there are some localized enhancements of the ADF, specially at the position of the Herbig-Haro objects HH 202, HH 203, and HH 204.  The combined data of all slit positions indicate a clear decrease of {\\elect}([{\\nii}]) and {\\elect}([{\\oiii}]) with increasing distance from the main ionizing source of the nebula, $\\theta^1$ Ori C. On the other hand, the radial distribution of the ADF shows a rather constant value across the nebula except at the inner 40$''$, where the ADF seems to increase very slightly toward $\\theta^1$ Ori C.  We have explored possible correlations between the ADF of O$^{++}$ and other nebular quantities, finding a possible very weak increase of the ADF for higher electron temperatures. There are not apparent trends between the ADF and c(H$\\beta$), {\\elecd}, the {\\elect}([{\\nii}])/{\\elect}([{\\oiii}]) ratio, O$^{++}$ abundance, and the O$^{++}$/O$^+$ ratio.  Our spatially resolved spectroscopy allows to estimate the value of the mean-square electron temperature fluctuation in the plane of the sky, a lower limit to the traditional {\\ts} parameter. We find very low values in all cases, result that is in contradiction with previous estimates from the literature. Our results indicate that the hypothetical thermal inhomogeneities --if they exist-- should be lower than our spatial resolution limit of about 1$''$.  It is clear that further studies on the  {\\elect}, chemical abundances, and ADF distributions at sub-arcsec spatial scales are necessary in trying to disentangle (a) whether small spatial scale temperature fluctuations and/or metal-rich droplets are really present in the Orion nebula and HII regions in general, and (b) the origin of AD problem and its possible relation with {\\ts} and other nebular properties. The observations needed for this task are very difficult even for ground-based large-aperture telescopes and, by now, unfeasible with the current space telescopes and their available instrumentation.  We thank G. Stasi\\'nska and M. Rodr\\'{\\i}guez for their fruitful comments and help. We are grateful to the referee, Y. Tsamis, for his careful reading of the paper and his comments. This work has been funded by the Spanish Ministerio de Ciencia y Tecnolog\u00eda (MCyT) under project AYA2004-07466."
    },
    "0710/0710.2359_arXiv.txt": {
        "abstract": "{} {We evaluate the generation of magnetosonic waves in differentially rotating magnetized plasma.} {Differential rotation leads to an increase of the azimuthal field by winding up the poloidal field lines into the toroidal field lines. An amplification of weak seed perturbations is considered in this time-dependent background state.} {It is shown that seed perturbations can be amplified by several orders of magnitude in a differentially rotating flow. The only necessary condition for this amplification is the presence of a non-vanishing component of the magnetic field in the direction of the angular velocity gradient.} {} ",
        "introduction": "In astrophysical bodies, differential rotation is often associated with magnetic fields of various strength and geometry. If the poloidal field has a component parallel to the gradient of angular velocity, then differential rotation can stretch toroidal field lines from the poloidal ones. In the presence of the magnetic field, differential rotation can be the reason for various magnetohydrodynamic instabilities, particularly if the field geometry is complex. Some of these instabilities occur in the incompressible limit (Velikhov 1959; Fricke 1969; Acheson 1978; Balbus \\& Hawley 1991) which applies if the magnetic field is weak and the Alfv\\'en velocity is smaller than the sound speed. Other instabilities become important only in sufficiently strong magnetic fields when the effect of compressibility plays a significant role (Pessah \\& Psaltis 2005; Bonanno \\& Urpin 2006, 2007). Note that incompressible instabilities can often be suppressed by a strong magnetic field. For example, the magnetorotational instability (MRI) does not occur if the magnetic field satisfies the condition $B^2 > -8 \\pi \\rho s \\Omega \\Omega' L^2$ where $\\rho$ is the density, $L$ is the lengthscale of disturbances, $\\Omega=\\Omega(s)$ is the angular velocity, and $s$ is the cylindrical radius; $\\Omega'= d \\Omega/ds$ (see, e.g., Urpin 1996, Kitchatinov \\& R\\\"udiger 1997). Moreover, if $\\Omega$ increases with the cylindrical radius MRI cannot arise.  In recent years, many simulations of differentially rotating magnetized bodies have been performed, and much of the dynamics was interpreted as being a direct consequence of the MRI (Brandenburg et al. 1995; Hawley at al. 1995; Matsumoto \\& Tajima 1995; Hawley 2000). Obviously, the MRI cannot be the only instability that operates in a rotating magnetized gas. For example, stratification can lead to a number of strong non-axisymmetric instabilities (Agol et al. 2001; Narayan et al. 2002; Keppens et al. 2002). Blokland et al. (2005) consider the influence of a toroidal field on the growth rate of the MRI and find that it leads to overstability (complex eigenvalue). Van der Swaluw et al. (2005) study the interplay between different instabilities and argue that the growth rate of convection can be essentially increased due to magnetorotational effects. Note, however, that these studies treat the stability of the magnetic field with a vanishing radial component, a condition which is often not met in astrophysical bodies. In fact, the presence of a radial magnetic field can change substantially the stability properties (Bonanno \\& Urpin 2006, 2007).  In this paper, we consider stability of a differentially rotating gas in the presence of a non-vanishing radial magnetic field. Differential rotation causes the azimuthal field to increase with time by winding up the poloidal field lines into the toroidal ones. Therefore, a development of small perturbations occurs in the time-dependent background state. We show that stretching of the azimuthal field leads to the generation of magnetosonic waves in a flow. Magnetohydrodynamic waves and turbulence generated by this instability can play an important role in enhancing transport processes in various astrophysical bodies, such as accretion and protoplanetary Disks, galaxies, stellar radiative zones, etc. ",
        "conclusions": "We have shown that the winding up of toroidal field lines from the poloidal field lines is accompanied by an amplification of seed perturbations of the velocity and magnetic field. A very simplified model has been considered in this paper, but we believe that qualitatively the same results can be obtained for more general background states and perturbations. Differential rotation and compressibility of the gas lead to the generation of magnetosonic waves with the amplitude that grows with time. The physical processes responsible for this amplification are exactly the same that result in the instability considered by Bonanno \\& Urpin (2006, 2007). The only difference is that, in this paper, we consider the development of perturbations on a time-dependent background state and, as a result, the growth of perturbations is not exponential.  The behaviour of seed perturbations depends essentially on various parameters and can generally be rather complex. If the parameter $q=s \\Omega'/\\Omega$ is relatively small ($\\leq 0.1$), then the perturbations of velocity and magnetic field initially grow monotonously and can reach quite high values. For example, the perturbation of the vertical velocity becomes approximately $150-200$ times greater than its initial value after only 15-30 rotation periods (see Figs.~2 and 3). Perturbations of the magnetic field reach even higher values during the initial stage. For instance, the Alfv\\'en velocity corresponding to the perturbation of the radial field component, $C_{As}$, can exceed the initial velocity perturbation by a factor $\\sim (1-3) \\times 10^5$ after the same time, but the perturbations of the toroidal field are even stronger. As a result of such a strong initial amplification, seed perturbations can already reach a non-linear regime  after 15-30 rotation periods if their initial values are sufficiently large. Further evolution of perturbations will then be entirely determined by non-linear effects. However, if the non-linear regime is not reached during this initial stage, the behaviour of perturbations becomes oscillatory with slowly growing amplitude ($\\propto t^{1/2}$). At sufficiently large $t$, the frequency of oscillations grows linearly with time and is given approximately by \\begin{equation} \\omega \\approx \\Omega \\left( \\frac{1}{2} k H \\sqrt{\\beta_s} s \\Omega' t \\right). \\end{equation} Note that perturbations of the magnetic field exhibit more regular behaviour because they can be expressed in terms of the time integrals of a rapidly oscillating velocity. In the case of a strong differential rotation ($q \\sim 1$), perturbations exhibit the oscillatory behaviour from the very beginning and the initial growth of their amplitude is less significant.  The generation of magnetosonic waves occurs even if the magnetic field is very strong and suppresses different MHD-instabilities which can arise in a differentially rotating flow (for example, the MRI). The presence of differential rotation and radial magnetic field is, however, crucially important for the considered process. Since both differential rotation and radial field are quite common in astrophysics, we believe that the considered mechanism can occur in various astrophysical bodies and plays an important role in enhancing transport processes in plasma.    \\vspace{0.5cm} \\noindent {\\it Acknowledgments.} This research project has been supported by a Marie Curie Transfer of Knowledge Fellowship of the European Community's Sixth Framework Programme under contract number MTKD-CT-002995. VU thanks INAF-Ossevatorio Astrofisico di Catania for hospitality."
    },
    "0710/0710.0878_arXiv.txt": {
        "abstract": "Using a population synthesis approach, we compute the total merger rate in the local Universe for double neutron stars, double black holes, and black hole -- neutron star binaries. These compact binaries are the prime source candidates for gravitational-wave detection by LIGO and VIRGO. We account for mergers originating {\\em both\\/} from field populations and from dense stellar clusters, where dynamical interactions can significantly enhance the production of double compact objects. For both populations we use the same treatment of stellar evolution. Our results indicate that the merger rates of double neutron stars and black hole -- neutron star binaries are strongly dominated by field populations, while merging black hole binaries are formed much more effectively in dense stellar clusters. The overall merger rate of double compact objects depends sensitively on the (largely unknown) initial mass fraction contained in dense clusters ($f_{\\rm cl}$). For  $f_{\\rm cl} \\lesssim 0.0001$, the Advanced LIGO detection rate will be dominated by field populations of double neutron star mergers, with a small but significant number of detections $\\sim 20$ yr$^{-1}$. However for a higher mass fraction in clusters, $f_{\\rm cl} \\gtrsim 0.001$, the detection rate will be dominated by numerous mergers of double black holes originating from dense clusters, and it will be considerably higher, $\\sim 25 - 300$ yr$^{-1}$. In addition, we show that, once mergers of double black holes are detected, it is easy to differentiate between systems formed in the field and in dense clusters, since the chirp mass distributions are strikingly different. If significant field populations of double black hole mergers are detected, this will also place very strong constraints on common envelope evolution in massive binaries. Finally, we point out that there may exist a population of merging black hole binaries in intergalactic space. ",
        "introduction": "Gravitational wave astronomy is entering a new era: LIGO (Abramovici et al.\\ 1992) has now taken a full year of data at its design sensitivity; VIRGO (Bradaschia et al.\\ 1990) is nearing completion. Other detectors like GEO or TAMA are also on track. There are well defined plans for improving the LIGO and VIRGO detectors so that, within the next few years, their sensitivity  will increase by a factor of up to 30. The list of potential sources for these high-frequency detectors is long and includes supernova explosions, neutron star oscillations, and persistent radiation from rapidly rotating, nonaxisymmetric neutron stars. However, compact object binaries remain the most promising sources. It has been known for some time that their formation may take place in two very different environments: in the galactic field, where their progenitors are massive binaries that evolve in isolation, and in dense star clusters, where they can form at high rates through dynamical interactions.  The properties of the populations of double compact object binaries have been previously investigated using both observational and theoretical approaches. In the typical observational approach, the known compact objects binaries, i.e., radio pulsars in double neutron star (NS-NS) systems, were analyzed in detail. Based on their observed properties and the radio selection effects the properties of the entire population can be reconstructed and the expected merger rate can be calculated (Kim et al.\\ 2005). This approach, however, cannot be extended to black hole -- neutron star (BH-NS) binaries and double black hole  (BH-BH) binaries since these systems have never been observed. Moreover, we know only one merging NS-NS binary in a globular cluster, and so this method has little power for the population originating in clusters.  The second, theoretical, approach is based on numerical simulations of stellar evolution in field populations, combined with dynamical simulations of star clusters. For field populations some recent examples of theoretical studies include those by Belczynski, Kalogera \\& Bulik (2002, hereinafter BKB02), Voss \\& Tauris (2003), Pfahl et al.\\ (2005), Dewi et al.\\ (2006) and Belczynski et al.\\ (2007a).  The importance of globular cluster evolution for double compact object formation has long been suspected. The fate of BH populations in globular clusters was studied by a number of groups (e.g., Sigurdsson \\& Hernquist 1993; Kulkarni, Hut \\& McMillan 1993; Portegies Zwart \\& McMillan 2000; Merritt et al.\\ 2004). However, previous studies have employed only very simplified treatments of stellar evolution for single stars and binaries in clusters, focusing instead mostly on dynamical interactions. It was pointed out (Phinney et al.\\ 1991; and more recently Grindlay, Portegies Zwart \\& McMillan 2006) that the formation of NS-NS and BH-NS binaries in clusters is not very efficient and the contribution of clusters to their total merger rate is rather small ($\\sim 10-30 \\%$). Although the predicted formation rate of merging double NS has a strong density dependence (Ivanova et al.\\ 2007), even for dense clusters like like 47~Tuc the NS contribution to total cluster merger rates should be overwhelmed by double BH mergers. Such systems are expected to be formed via dynamical interactions in cluster cores (Gultekin, Miller \\& Hamilton 2004; O'Leary et al.\\ 2006, 2007). O'Leary et al.\\ (2006, 2007) used realistic initial conditions for their cluster simulations (which included the stellar evolution of massive stars and binaries in the first few Myr of the cluster life) and assumed immediate creation of a BH subcluster. Their calculations did not consider further stellar or binary evolution and neglected the effects of lower-mass stars on BH populations. Their results provide the most up-to-date estimates of BH-BH merger rates from globular clusters and the corresponding LIGO detection rates assuming rapid BH segregation into an isolated subcluster. Mackey et al.\\ (2007) investigated the effects of BHs on the structural evolution of globular clusters. They used a realistic initial mass function for single stars and observed rapid mass segregation of the BHs into the cluster core (as expected from the Spitzer instability; see Watters et al.\\ 2000 and references therein). However, this study was based on direct $N$-body simulations that did not include {\\em any\\/} primordial binaries, even though binaries could affect the mass segregation and subsequent dynamics very significantly.  In this paper we take the next step and calculate the evolution of representative star clusters from the onset of star formation taking into account the large binary fraction for massive stars. We assume that inelastic interactions of hard binaries in the cluster core are effective enough to prevent BHs from completely separating from the rest of the cluster, and we therefore treat the BHs as always well mixed and in thermal equilibrium with other stars in the cluster core. This simplifying assumption is opposite, and complementary, to the one adopted in the models of O'Leary et al. (2006, 2007). Their results and the merger rates calculated in this work give, respectively, lower and upper bounds on the LIGO detection rates.  In our new models we include both stellar dynamics and full stellar evolution for single and binary stars to predict the merger rate of double compact objects. All stellar populations are evolved and allowed to interact through an entire cluster lifetime ($\\sim 13$ Gyr). Then we combine our estimates for star clusters and the merger rates calculated in O'Leary et al.\\ (2006) with the most recent population synthesis field calculations  (Belczynski et al.\\ 2007a) to deduce the total merger rate. In Section~2 we present a description of the stellar binary evolution treatment used and we introduce our model for globular clusters. In Section~3 we present our results and in Section~4 a brief discussion. ",
        "conclusions": "\\label{summary}  We have presented the ranges of total cosmic merger rate of double compact objects in the local Universe. Our results apply to the distances that will be reached by ground based gravitational-wave detectors such as LIGO or VIRGO ($\\sim 300$ Mpc for a typical NS-NS system). In calculating the rates we have considered formation of close double compact objects in field populations as well as in dense stellar clusters, i.e., globular clusters. The main result of our study shows that the predicted merger rates are too small for detection with the current instruments (i.e., initial LIGO) but are very promising for the upgraded detectors (i.e., advanced LIGO).  We find that field populations dominate the formation of close NS-NS and BH-NS systems. This was already predicted by Phinney (1991) on the basis of observed binary pulsars. Recently Grindlay, Portegies Zwart \\& McMillan (2006) have shown that globular clusters can contribute only up to $\\sim 20 \\%$ to the cosmic NS-NS merger rates. Formation of double compact object systems with NSs is inhibited by the rather large natal kicks NSs receive in supernovae (Hobbs et al.\\ 2005). These lead to {\\em (i)} disruption of binaries hosting NS progenitors and {\\em (ii)} ejection of NSs from the cluster. The remaining systems containing a NS interact with heavier stars, and in particular BHs, and through exchange interactions are removed from binaries and do not form double compact objects. Also, the initial cluster mass that we use in our simulations ($\\sim 5 \\times 10^5\\msun$, which corresponds to the current average globular cluster mass in our Galaxy; e.g., Meylan \\& Heggie 1997, see their Fig.~10.2) is too small to produce a significant number of NS-NS or BH-NS systems. In particular, in our cluster simulations, we do not find any mergers of NS-NS nor BH-NS systems. The results for field populations of NS-NS and BH-NS mergers were discussed in detail by Belczynski et al.\\ (2007a). They find (their model A) that field  Galactic merger rates are $\\sim 15$ Myr$^{-1}$ and $\\sim 0.1$ Myr$^{-1}$ for NS-NS and BH-NS, respectively. In our work we have adopted these rates as the total NS-NS and BH-NS merger rates (field + clusters) since the contribution from clusters is rather small for these systems. As discussed by Belczynski et al. (2007a) these rates are too small for any detection with the initial LIGO, while for advanced LIGO a small but significant number of detections is predicted: $\\sim 15$ yr$^{-1}$ and $\\sim 1$ yr$^{-1}$ for  NS-NS and BH-NS, respectively.  Formation of close BH-BH systems that merge in a Hubble time is expected to be very effective in clusters. This was already noted in studies that use realistic initial conditions for the evolution of BH-BH binaries in clusters (e.g., Miller \\& Hamilton 2002; Gultekin, Miller \\& Hamilton 2004; O'Leary et al.\\ 2006). Here, we have followed the evolution of a realistic average cluster with full stellar evolution and physical treatment of all BH-BH binaries, assuming that binary interactions are able to prevent the BHs from separating into an isolated subcluster. We have found that the production of BH-BH binary mergers under these assumptions is indeed remarkably effective in dense clusters. The most striking and counter-intuitive difference with other studies is that the predicted BH-BH merger rate in clusters is rather constant over long periods of time ($\\sim$ Hubble time). In earlier studies the BH populations were usually introduced in the cluster core and evolved separately from the other stars in the cluster. That leads to a larger merger rate at early times in the cluster evolution followed by a very  rapid drop in the merger rates once BHs are eliminated (through mergers and ejections). However, the assumptions introduced in our work, namely {\\em (i)} stars that form BHs are initially placed throughout the cluster, and {\\em (ii)} BHs do not segregate so strongly as to form a completely decoupled subcluster and are instead allowed to interact with other stars in the cluster, make our results qualitatively different. We find that BHs that form in the core, in fact, produce mergers at the early stages of cluster evolution. But later many massive stars that were in the halo and that formed BHs  will steadily feed the cluster core with BHs, i.e., continued mass segregation in the cluster halo is providing BHs to the cluster core on long timescales. Additionally, exchange interactions of BHs with unevolved (e.g., main sequence stars) in the core are effective over a long periods of time, and usually lead to formation of close BH-BH systems that merge later in the core. A simplified scenario involves two single BHs sinking into the core; each catches a main sequence companion, and then two BH-MS binaries interact, forming a close BH-BH system and two single main sequence stars are released back into the cluster.  Our average cluster merger rate for BH-BH systems is $\\sim 3$ Gyr$^{-1}$ (see Figure~\\ref{f.merg}) for our simulated cluster with mass $M_{\\rm cl}=4.8 \\times 10^5 \\msun$. If we scaled this rate to the mass of the galactic disk ($M_{\\rm MW}=3.5 \\times 10^{10} \\msun$), then the BH-BH merger rate would come up to $\\sim 180$ Myr$^{-1}$ (i.e., this is the expected BH-BH merger rate if the entire mass of our Galactic disk were contained in dense globular clusters). We should compare this rate to the Galactic field BH-BH merger rate, $0.025$ Myr$^{-1}$ (Belczynski et al.\\ 2007a). It is clear that the production of BH-BH mergers is much more efficient (by $\\sim 3-4$ orders of magnitude) in clusters as compared to field evolution\\footnote{ We have used the rate from the same evolutionary model used for field population (model A of Belczynski et al.\\ 2007a) as it was employed in our standard cluster model.}.  To combine our predicted cluster rates with the field rates we need to know the initial stellar mass fraction contained in clusters. If we look at the Galactic globular clusters we find that they contain about 0.001 of the total mass in stars found in the field ($f_{\\rm cl}=0.001$). However, it is reasonable to expect that many clusters may have been completely destroyed and that the clusters we see today were initially more massive and lost significant mass through evaporation (e.g., Vesperini 1998; Joshi et al.\\ 2001). Although we present our predictions for the entire range of plausible $f_{\\rm cl}$ values, we consider that the most reasonable estimate is still $f_{\\rm cl} \\gtrsim 0.001$. While we do not have reliable mass estimates for globular clusters in elliptical galaxies, it has been shown that the specific frequency of GCs per galaxy luminosity in ellipticals is significantly (about an order of magnitude) higher than in spiral galaxies (Kim \\& Fabbiano 2004). Moreover, elliptical galaxies are on average more massive than spiral galaxies. Therefore, an upper limit on the initial mass fraction contained in all clusters, although highly uncertain, can probably be set to $f_{\\rm cl} \\lesssim 0.01$.  The total cosmic merger rate for BH-BH systems is then strongly dependent on the mass contained in globular clusters. For a small fraction ($f_{\\rm cl}=0.001$) we find the merger rate per Milky Way in our model to be $\\sim 0.2$ Myr$^{-1}$, while for a larger fraction ($f_{\\rm cl}=0.01$) the merger rate is $\\sim 2$ Myr$^{-1}$. At this point we can also estimate the detection rate for a given GW detector. We must keep in mind that the average chirp mass of field BH-BH binaries is much smaller ($M_{\\rm c} \\sim 7 \\msun$) than for cluster BH-BH mergers ($M_{\\rm c} \\sim 20 \\msun$). This allows us to observe cluster BH-BH mergers in a much larger volume and their relative contribution to the detection rate is greater than indicated simply by the merger rates (see \\S\\,~\\ref{ligo}). One could also imagine forming of few very massive BHs (up to $\\sim 100\\msun$) in the cluster (e.g., through mergers of massive binaries; Belczynski et al.\\ 2006) which could form BH-BH mergers characterized by extremely high chirp masses. Such mergers would be detectable from much greater distances, making the observed rates even higher.  The predicted ranges of total detection rates, for all types of double compact objects, are presented in Figure~\\ref{f.rates} as a function of cluster contribution. The most likely range of values for this parameter ($f_{\\rm cl} \\sim 0.001\\div 0.01$) is marked with the vertical shaded area in Figure~\\ref{f.rates}. For very low cluster contributions ($f_{\\rm cl}=0.0001$) the detection rates correspond to mergers coming only from field populations and are adopted from the reference model of Belczynski et al.\\ (2007b; their model A). With increasing cluster contribution we see a drastic increase in predicted detection rates. This increase is connected to the very effective production of BH-BH mergers in clusters as discussed above. For advanced LIGO the detection rates could be as high as $\\sim 25-3000$ yr$^{-1}$ and are higher by more than an order of magnitude than the rates just for field populations. The total rates are dominated by dynamically formed BH-BH mergers in dense stellar clusters. If advanced LIGO does not observe this population of BH-BH mergers it will put strong constraints on the initial stellar mass fraction contained in dense stellar clusters.  The production of BH-BH mergers in the field is inhibited by the process identified in Belczynski et al.\\ (2007b): many potential BH-BH progenitors evolve through a common envelope phase while the donor is evolving through the Hertzsprung gap. Such a common envelope leads most likely to a merger and aborts potential formation of a BH-BH system (because, in the Hertzsprung gap, the star has not yet developed a clear core-envelope structure and the inspiral does not stop before complete merger of the two interacting stars). If this current understanding of the common envelope phase is correct, we do not expect detection of more than a few field BH-BH mergers per year. However, if the progenitors somehow survive this phase, we could then expect up to $\\sim100$ detections of field BH-BH mergers per year by Advanced LIGO. Since the chirp mass distribution of the field and cluster populations are so different (see Figure~\\ref{f.mchirp} and Figure 4 of Belczynski et al.\\ 2007b) it would be easy to tell the two populations apart. This would, in turn, allow us to both {\\em (i)} derive the initial mass fraction in clusters, and {\\em (ii)} constrain the fate of massive binary systems going through a common envelope phase. This result highlights the importance of the chirp mass distribution as a diagnostic tool in gravitational wave astronomy (Bulik \\& Belczynski, 2003; Bulik, Belczynski, \\& Rudak, 2004)."
    },
    "0710/0710.2750_arXiv.txt": {
        "abstract": "{The transition from galactic to extragalactic cosmic rays is discussed. One of critical indications for transition is given by the Standard Model of Galactic cosmic rays, according to which the maximum energy of acceleration for iron nuclei is of order of $E_{\\rm Fe}^{\\rm max} \\approx 1\\times 10^{17}$~eV. At $E > E_{\\rm Fe}^{\\rm max}$ the spectrum is predicted to be very steep and thus the Standard Model favours the transition at energy not much higher than  $E_{\\rm Fe}^{\\rm max}$. As observations are concerned there are two signatures of transition: change of energy spectra and elongation rate (depth of shower maximum in the atmosphere $X_{\\rm max}$ as function of energy). Three models of transition are discussed: dip-based model, mixed composition model and ankle model. In the latter model the transition occurs at the observed spectral feature, ankle, which starts at $E_a \\approx 1\\times 10^{19}$~eV and is characterised by change of mass compostion from galactic iron to extragalactic protons. In the dip model the transition occures at the second knee observed at energy $(4 -8)\\times 10^{17}$~eV and is characterised by change of mass composition from galactic iron to extragalactic protons. The mixed composition model describes transition at $E \\sim 3\\times 10^{18}$~eV with mass composition changing from galactic iron to extragactic mixed composition of different nuclei. These models are confronted with observational data on spectra and elongation rates from different experiments, including Auger.} \\begin{document} ",
        "introduction": "The Ultra High Energy Cosmic Ray (UHECR) has two most important problems. One of them is a presence of spectrum features produced by propagation of UHECR particles through Cosmic Microwave Radiation (CMB) and the second is transition from galactic to extragalactic Cosmic Rays (CR).  In the case of extragalactic  protons two spectral signatures caused by interaction with CMB are predicted: Greisen-Zatsepin-Kuzmin (GZK) cutoff \\cite{GZK}  and pair-production dip \\cite{BG88}.  GZK cutoff is most spectacular prediction for UHECR, which status is still uncertain in present observations, though there are the indications to its presence.  The pair-production dip is the spectral feature originated due to electron-positron pair production by extragalactic protons  interacting with CMB: $p+\\gamma_{\\rm CMB} \\rightarrow p+e^++e^-$. Recently this feature has been studied in the works \\cite{Stanev2000,BGGPL,BGG}. The dip has been observed with very good statistical significance $\\chi^2$/d.o.f.$\\sim 1$ by the Fly's Eye, Yakutsk, Akeno-AGASA and HiRes detectors, and with much worse statistical significance by Auger detector.  The pair-production dip and GZK cutoff are signatures of protons. The confirmation of the shape of these features is the evidence for proton-dominated composition of primary CRs. For nuclei as primaries the shape of the dip and GZK cutoff are strongly modified.  The different explanation of the dip has been proposed by Hill and Schramm \\cite{HS85}. They interpreted the dip observed in 1980s in terms of two-component model. The low energy component can be either galactic or produced by Local Supercluster. The similar model has been considered in \\cite{YT}. The Hill-Schramm dip is widely used now for the explanation of the observed dip.  From 1970s in the UHECR spectrum there was observed a flattening, which is called {\\em ankle}. Discovery of this feature at Haverah Park detector was interpreted as transition from the steep galactic component to more flat extragalactic one. The transition at ankle has been recently considered in \\cite{ankle}.  In the dip  model the transition is completed at the beginning of the dip at $E \\approx 1\\times 10^{18}$~eV. The ankle in this model appears as intrinsic part of the dip. Like in ankle model, the transition occurs here also as intersection of flat extragalactic component (this flatness is especially prominent in case of diffusive propagation) with steep galactic spectrum.  In the dip and ankle models the extragalactic component is assumed to be proton dominated, while the galactic component is most probably composed by iron nuclei. In the {\\em intermediate model}, where transition occurs in the middle of the dip, the extragalactic CRs are assumed to have mixed composition \\cite{mixed}.  In this paper all three above-mentioned models of transition are discussed. The logic of our discussion is as follows: we approach first the transition from the high energy end of galactic CRs,  then we discuss the properties of UHECR relevant for transition problem and finally we describe the transition from properties of these two components. ",
        "conclusions": "\\vspace{1.3mm} The region of transition from galactic to extragalactic CRs at energy between $1\\times 10^{17} - 1\\times 10^{19}$~eV is the key energy range for understanding the origin of CRs. At low energy part it includes the high energy end of galactic CRs. The information on maximum energy of acceleration, chemical composition and propagation in Galaxy at these energies will clarify the total picture of origin at lower energies. The low energy part of UHECRs is important for understanding of origin of UHECRs and their propagation in extragalactic magnetic fields. The transition from galactic to extragalactic CRs is the central issue of this energy region.  There are two detectors which cover partially the above-mentioned region: KASCADE-GRANDE \\cite{KASCADE-G} and TALE \\cite{TALE}. There are also the proposals to extend the observations of Auger to energy $E \\sim 1\\times 10^{17}$~eV (see e.g. \\cite{LE-Auger}). The Auger detector has great potential to explore this region, building more dense part of the detector covered with fluorescent, scintillator and muon detectors.  The basic information which can be obtained includes precise measurement  of energy spectra and mass composition (there is little hope to detect anisotropy in this energy region, though in some models the galactic sources can be observed in protons with energy $E \\lesssim 10^{18}$~eV \\cite{BGH}).  At present we have the sufficiently good data on spectra and mass composition at energy range $1\\times 10^{18} - 4\\times 10^{19}$~eV. The spectra are measured with high statistics (especially in case of the Auger detector), but problem is the accuracy of energy determination. From quite disappointing Fig.~\\ref{fig:AgHiYaAu} (left panel) one concludes that scales of energy determination is quite different in all detectors. Energy calibration with help of the pair-production dip suggests that energy measured by scintillator detectors is systematically higher than that by the fluorescence detectors and it gives a reasonable recipe of increasing energies given by fluorescent method and decreasing it for the  scintillation method. In this case the curves 'Yakutsk' and 'Akeno-AGASA' in Fig.~\\ref{fig:AgHiYaAu} go down and 'HiRes' and 'Auger' - up. For HiRes, AGASA and Yakutsk the method of calibration with help of dip works successfully (see Fig.~\\ref{fig:AgHiYa}) with energy shift within the allowed systematic errors, but for Auger it requires the shift by factor 2 greater than systematic error.  The pair-production proton dip in terms of modification factor is an excellent tool to measure {\\em spectrum shape} independently of absolute flux. From Fig.~\\ref{fig:dips} one sees the excellent agreement of the theoretical dip with data of AGASA, HiRes and Yakutsk. By  the standards of cosmic-ray physics the agreement with Auger data is also good, but $\\chi^2$ for comparison with SD data is very large. This is a result of very big statistics in the surface detectors at lowest energies $E \\geq 4.5\\times 10^{18}$~eV. In the lowest energy bin at $E=4.5\\times 10^{18}$~eV there are 4128 events and the error in determination of flux provided mostly by this statistics is $\\delta J/J=0.024$. The theoretical value of modification factor at this energy is only 14\\% higher than experimental value, but owing to very small $\\delta J/J$,the contribution of this bin to $\\chi^2$ is 99.27 ! Most probably the other sources of errors should be included in the bins with small $\\delta J/J$, and a possible source of this error is the energy errors which are changing randomly inside a bin. These could be statistical errors and energy-dependent part of systematic errors. Assuming that number of events are distributed in a bin as $N(E)=K E^{-\\gamma}$ one obtains $\\delta J/J = \\gamma (\\delta E/E)_r$, where $(\\delta E/E)_r$ is the random energy error inside the bin. The estimated value $\\delta J/J$ is much larger than what obtained in Auger analysis for all reasonable values of $(\\delta E/E)_r$ and $\\gamma$. More generally, according to Markus Roth's remark, $\\chi^2$ analysis is not adequate for the cases of small  $\\delta J/J$ and large $(\\delta E/E)$. At this stage of analysis we do not consider Fig.~\\ref{fig:dips} as contradiction with Auger data.\\\\*[4mm] Coming to the transition from galactic to extragalactic CRs, we emphasize that at present there are only two experimental methods to study it: measuring the spectrum and mass composition. The transition will be clearly seen if spectrum of iron nuclei and that of protons are measured separately (see Fig.~\\ref{fig:dip-ankle}), but even without this ideal possibility the total spectrum has signatures of transition in the form of the spectral features - {\\em second knee} in case of the dip model and {\\em ankle} in case of the ankle model. The spectrum can be measured nowadays with high accuracy and its shape contains the information about mass composition, which is the other characteristic of the transition. The pair-production dip with its specific shape is a signature of proton-dominated composition (nuclei contribution should be not more than 10 -15 \\% \\cite{BGGPL}) and its observational confirmation is an argument not weaker than that due to $X_{\\rm max}$ measurement (we remind that only two free parameters are involved in describing about 20  energy bins in each experiment). \\\\*[4.1mm] The mass composition gives another way to test the transition. The best method at present is given by measuring of elongation rate $X_{\\rm max}(E)$. Unfortunately this method has many uncertainties, including those in value of fluorescent yield, absorption of UV light in the atmosphere and uncertainties in the models of interactions, needed to convert the tested mass composition into $X_{\\rm max}$. The systematic errors in measuring $X_{\\rm max}$ can be as large 30~g/cm$^2$ to be compared with difference about 100~g/cm$^2$  between $X_{\\rm max}$  for protons and iron.  The better sensitivity for distinguishing different  nuclei is given by distribution over $X_{\\rm max}$ \\cite{ABBO}.  There are three models of the transition: ankle, dip and mixed-composition model. They differ most notably by the energy of transition (ankle: $E \\sim 1\\times 10^{19}$~eV, dip: $E \\approx 1\\times 10^{18}$~eV and mixed composition model $E \\approx 3\\times 10^{18}$~eV), and by mass composition of extragalactic component (protons -  for the ankle model, proton-dominated -  for the dip model and mixed composition -  for the third model).  The {\\em ankle model} contradicts the Standard Model of Galactic CRs (energy where galactic flux is half of that observed is two orders of magnitude higher than energy of iron knee) and severely disagrees with $X_{\\rm max}$ measured in all experiments at $(1.5 - 5)\\times 10^{18}$~eV.  The {\\em dip model} is based on well confirmed signature of proton interaction with CMB - pair-production dip. The two other models must assume that agreement of pair-production dip with data is accidental and the observed dip is produced by two components, galactic and extragalactic. The dip model assumes the iron-dominated galactic flux below $5\\times 10^{17}$~eV and proton-dominated extragalactic flux above $1\\times 10^{18}$~eV. This mass composition is confirmed by HiRes and HiRes-Mia data for elongation rate. It does not contradict the bulk of all data on $X_{\\rm max}$, but contradicts $X_{\\rm max}$ measured by Auger, especially the highest energy points. The generation spectrum in this model is $E^{-2}$ or $E^{-2.2}$ as needed by shock acceleration with a steepening to $\\gamma_g=2.7$ due to distribution of sources over maximum energy of acceleration of source luminosities. The proton-dominated composition can be produced in some models of injection to the shock acceleration.  The {\\em mixed composition model} assumes mixed composition generation spectrum for extragalactic component with generation index 2.1 - 2.3. It has many free parameters, most notably ones describing the mass composition of the generation spectrum, and thus it can in principle explain any observed mass composition. However, this model has a robust prediction at energy $E \\gtrsim 3\\times 10^{19}$~eV: proton-dominated composition and the GZK feature. As far as Auger elongation rate is concerned, the mixed composition model explains well the break in elongation rate at $2\\times 10^{18}$~eV and contradicts the two Auger points at $E > 2\\times 10^{19}$~eV. The energy where transition to extragalactic CRs is completed in most versions of this model equals $E \\approx 3\\times 10^{18}$~eV. Much better quality of data on $X_{\\rm max}$ is needed to distinguish the dip and mixed-composition models by $X_{\\rm max}$ measurements. Probably it is possible to do using $X_{\\rm max}$ distribution \\cite{ABBO}.\\\\*[1mm] We will comment now on agreement of the transition models with the measured galactic spectrum. For all three models it is reached by the formal subtraction procedure: the galactic spectrum is found as difference between measured total spectrum and calculated extragalactic spectrum. But the galactic spectrum calculated in the Standard Model at $E \\gtrsim 1\\times 10^{17}$~eV is very steep and, as was demonstrated in \\cite{Tanco}, for diffusive model of propagation  all three models contradict the calculated galactic spectrum, the dip model to the less extent. Strictly speaking this contradiction is produced by exponential cutoff in the acceleration spectrum at $E > E_{\\rm max}^{\\rm acc}$. \\\\*[2mm] The most consistent conclusions on nature of observed  UHECRs are obtained at present by HiRes detector: it has confirmed the pair-production dip and thus proton-dominant composition at $1\\times 10^{18} - 4\\times 10^{19}$~eV, the $X_{\\rm max}$ measurements agree with proton-dominant composition at $E > 1\\times 10^{18}$~eV , and $E_{1/2}$ measurement confirms that steepening of the spectrum observed at $E > 4\\times 10^{19}$~eV is really the GZK cutoff. Therefore, according to these data CRs observed at $E \\gtrsim 1\\times 10^{18}$~eV are extragalactic protons exhibiting two signatures of interaction with CMB: pair-production dip and GZK feature."
    },
    "0710/0710.5036_arXiv.txt": {
        "abstract": "We study the 37 brightest radio sources in the Subaru/\\textit{XMM-Newton} Deep Field (SXDF). Using mid-IR (Spitzer MIPS 24 $\\mu \\rm m$) data we expect to trace nuclear accretion activity, even if it is obscured at optical wavelengths, unless the obscuring column is extreme. Our results suggest that above the `FRI/FRII' radio luminosity break most of the radio sources are associated with objects that have excess mid-IR emission, only some of which are broad-line objects, although there is one clear low-accretion-rate FRI. The fraction of objects with mid-IR excess drops dramatically below the FRI/FRII break, although there exists at least one high-accretion-rate QSO. Investigation of mid-IR and blue excesses shows that they are correlated as predicted by a model in which a torus of dust absorbs $\\sim$30\\% of the light, and the dust above and below the torus scatters $\\gtsim$1\\% of the light. ",
        "introduction": "Powerful radio sources are believed to have central super-massive black holes (SMBH) with uniformly high accretion rates at the highest radio luminosities and typically lower accretion rates at lower radio luminosities (\\cite{rs91}). Low-luminosity radio jets can, however, be associated with high-accretion-rate systems, and these so-called `radio quiet' quasars appear to have similar FRI-like radio structures to low-accretion-rate counterparts of similar radio luminosity (e.g. \\cite{hbr07}). At low redshift, the most massive ($\\gtsim$ $10^{8} \\rm M_{\\odot}$) SMBH typically have very low accretion rates with systematically higher average values at $z \\gtsim$ 2, the so-called `quasar epoch' (\\cite{yt02}). These observational results fit in with theoretical ideas that a `quasar mode' of feedback is prevalent in the distant universe, and that a `radio mode' feedback is dominant at low redshift (e.g. \\cite{cro06}).  The central region of an AGN is surrounded by a dusty torus which absorbs light and re-emits it in the infrared. Above and below the plane of the torus, dust scatters light yielding a blue excess. Such mechanisms make it difficult to observe objects viewed through the torus directly in the optical, UV and soft X-rays. The torus creates anisotropic obscuration of the central regions resulting in two different types of observed objects, type 1 that are viewed face-on and type 2 that are viewed edge-on. Here we use mid-IR observations to search for evidence of accretion in a manner which is far less dependent on orientation.  The sample studied here is the 37 brightest radio sources from the VLA survey of the Subaru/\\textit{XMM-Newton} Deep Field (SXDF; \\cite{sim06}) with flux densities greater than 2 mJy at 1.4 GHz. Optical, X-ray and radio observations of the SXDF were made within the 1.3 square degree Subaru/\\textit{XMM-Newton} Deep Field with Subaru, \\textit{XMM-Newton} and the VLA respectively. Thirteen of our objects are not as yet spectroscopically confirmed so we use photometric redshifts in these cases, calculated with the HYPERz code (\\cite{boz00}) and typically nine data points from $B-$band (440 nm) to 4.5 $\\mu \\rm m$ (Vardoulaki et al. in prep). ",
        "conclusions": "\\begin{figure} \\begin{center} \\setlength{\\unitlength}{1mm} \\begin{picture}(140,38) \\put(65,-21){\\special {psfile=\"l1_4_D.ps\" vscale=32 hscale=32 angle=90}} \\put(143,-21){\\special {psfile=\"blue_red.ps\" vscale=32 hscale=32 angle=90}}  \\end{picture} \\end{center} \\vspace{0.5in} {\\caption[junk]{\\label{fig1} a {\\it Left}: Radio Luminosity $\\log_{10}(L_{1.4{\\rm GHz}}/ \\rm [W Hz^{-1} sr^{-1}])$ versus largest projected linear size $D$: symbols indicate optical/IR classification; filled red circles for quasars `Q'; filled blue triangles for obscured quasars `OQ'; filled black squares for possible galaxies `G?'; black squares for secure galaxies `G'; green upside-down triangles for starbursts 'SB'; orange diamonds for weak quasars `WQ'; and light blue stars for BL Lac `BL'. The horizontal lines show the RLF and FRI/FRII breaks calculated from the values in \\cite{fr74} using a typical steep-spectrum spectral index of 0.8 and translated to our assumed cosmology. b {\\it Right}: Blueness versus mid-IR excess. Symbols are the same as in the left figure. The black dotted line corresponds to the best-fit line in the log-linear plane where all objects, were treated as detections; the slope and intercept are 0.26 and -1.27 respectively, giving (see Eqn (2)) $\\rm k_{1} = 0.05$ and $\\rm k_{2} = 0.03$. The blue solid line corresponds to the best-fit line in log-linear plane where objects without detections at 24 $\\mu \\rm m$ were treated as limits; the slope and intercept are 0.10 and -0.41 respectively, giving $\\rm k_{1} = 0.39$ and $\\rm k_{2} = 0.09$. The Buckley-James method in the ASURV statistics package (\\cite{lav92}) was used in these calculations. `SB' and `BL' objects were excluded from the calculations since they have SEDs dominated by different physical processes to those assumed in the model described by Eqns (1) and (2). We adopt a radio spectral index $\\alpha$ = 0.8 ($S_{\\nu} \\propto \\nu^{-\\alpha}$), unless a spectral index could be calculated using 1.4 GHz data from \\cite{sim06} and 325 MHz data from \\cite{tasse} (see Vardoulaki {\\it et al.} in prep.). We assume throughout a low-density, $\\Lambda$-dominated Universe in which $H_{0}=70~ {\\rm km~s^{-1}Mpc^{-1}}$, $\\Omega_{\\rm M}=0.3$ and $\\Omega_{\\Lambda}=0.7$. }} \\end{figure}  \\addtocounter{figure}{0}  We use optical/IR observations to classify a radio source as either Quasar `Q', Obscured Quasar `OQ', Galaxy? `G?', Galaxy `G', Starburst `SB', Weak Quasar `WQ' or BL Lac `BL'. We deem that nuclear accretion is `significant' in objects that obey $\\log_{10}(L_{24 \\mu \\rm m}/ \\rm [W Hz^{-1} sr^{-1}]) > 23.1$ (or $[\\lambda L]_{24 \\mu \\rm m} > 10^{37.3}$ ${\\rm W}$). This value corresponds to $[\\lambda L]_{24 \\mu \\rm m} \\ge 10^{-1.8} L_{Edd}$, a typical lower limit for quasars (\\cite{mcl04}), for a black hole mass $M_{\\rm BH} \\ge 10^{8} M_{\\odot}$, a typical lower limit for radio sources (\\cite{mcl_etal04}); $L_{Edd}$ is the Eddington luminosity. We then define the following categories:\\\\ i) {\\bf Q}: Broad lines in the optical spectrum (3/37 cases). None of these are detected at 24 $\\mu \\rm m$, although their limits are insufficient to rule out significant accretion.\\\\ ii) {\\bf OQ}: Objects with a 24-$\\mu \\rm m$ detection (5/37 cases) and with sufficient $L_{24}$ to represent significant accretion. This class may be incomplete in that some objects in the `G?' class, as described next, have limits above this critical value.\\\\ iii) {\\bf G?}: A galaxy that has a 24-$\\mu \\rm m$ limit consistent with it lying above the $\\log_{10}(L_{24 \\mu \\rm m}/$ $\\rm [W Hz^{-1} sr^{-1}]) = 23.1$ line\\footnote{Because of the 24 $\\mu \\rm m$ flux density limit, these objects are at high redshift, and hence, because of the 1.4-GHz flux density limit, a high-$L_{1.4 \\rm GHz}$ sub-set of the objects lacking Spitzer 24 $\\mu \\rm m$ detections.} (11/37 cases).\\\\ iv) {\\bf G}: All other objects (15/37 cases) without significant accretion, unless they fall into three special categories defined by properties derived from spectroscopy, the SED and the optical structure: {\\bf SB}: evidence from the SED of a starburst component (1/37 cases); {\\bf WQ}: evidence from the SED of a quasar component but no 24 $\\mu \\rm m$ detection (1/37 cases); {\\bf BL}: featureless red continuum and a point source at $K$ (1/37 cases).\\\\  Figure 1a shows the 1.4-GHz radio luminosity at $L_{\\rm 1.4 GHz}$ versus the projected linear size $D$ with the symbols denoting the different optical/IR classes. We see that nearly all `Q', `OQ' and `G?' objects of our sample lie above the `FRI/FRII' luminosity break\\footnote{Although the FRI/FRII classification scheme is on the basis of radio structure, there is a sharp change in radio structure at a characteristic radio luminosity \\cite{fr74}.}, with the exception of the `OQ' sxdf\\_0034 (the `G?' object near sxdf\\_0034 lies very close to the boundary of significant accretion). In previous studies, the quasar fraction has been defined as the number of sources with quasar-like optical features (e.g. broad lines) and has a value of $\\sim 0.1 \\rightarrow 0.4$ over this range of $L_{1.4 \\rm GHz}$ (e.g. \\cite{wil00}). We introduce the `quasar-mode fraction' $f_{\\rm QM}$ to describe the fraction of objects with high accretion rates to the total number of objects. Above the FRI/FRII break $f_{\\rm QM} \\sim 0.5 - 0.9$ (the lower value assumes the 24 $\\mu \\rm m$ limits are much higher than the true 24 $\\mu \\rm m$ values, whereas the higher value assumes the true values lie just below the limits). The one clear exception in this regime is sxdf\\_0001, which has no evidence of a QSO and a clear Twin-Jet (FRI) radio structure.  The quasar-mode fraction drops dramatically below the FRI/FRII break\\footnote{We note that objects in our sample above the FRI/FRII break have median redshift $z_{\\rm med} \\sim 1.6$, whereas those below have $z_{\\rm med} \\sim 0.65$, so evolutionary effects may also be important.}, and whether or not one excludes some of the compact ($D <$ 100 kpc) sources as potentially part of a separate (beamed) population, then $f_{\\rm QM} \\ltsim$ 0.1 because nearly all objects are galaxies `G'. The counter example here are sxdf\\_0034, the only `OQ' below the FRI/FRII break, and potentially an optically-obscured example of unobscured FRI QSOs already studied in this radio luminosity regime (e.g. \\cite{hbr07}).  Inspection of the SEDs (Vardoulaki {\\it et al.} in prep) shows that some of our objects have an excess at 24 $\\mu \\rm m$ above that expected from extrapolation of the stellar populations. This is quantified via a measure of the mid-IR excess, $\\log_{10}([\\nu L]_{10 \\mu \\rm m rest} / [\\nu L]_{1\\mu \\rm m rest})$. A comparison of mid-IR excess and blueness is presented in Fig. 1b where a positive correlation is obvious. The generalised Spearman correlation calculated using survival analysis statistical package ASURV (\\cite{lav92}) is 0.657 with a 99\\% probability for a correlation.  Consider a simple model in which blueness is connected to mid-IR excess through the following equations\\footnote{Equation (2) relies on the $ln(1+x) \\approx x$ approximation which is only accurate around and below the knees of the functions plotted in Fig. 1b.}: \\begin{equation} [\\nu L]_{4000 \\rm \\AA \\rm rest} = \\rm k_{1} \\times [\\nu L]_{1\\mu \\rm m rest}  + \\rm k_{2} \\times [\\nu L]_{10 \\mu \\rm m rest}   \\Rightarrow \\end{equation} \\begin{equation} \\log_{10}\\left(\\frac{[\\nu L]_{4000 \\rm \\AA rest}}{[\\nu L]_{1\\mu \\rm m rest}} \\right) = \\log_{10}(\\rm e) \\times \\frac{\\rm k_{2}}{\\rm k_{1}} \\times \\left(\\frac{[\\nu L]_{10 \\mu \\rm m rest}}{[\\nu L]_{1\\mu \\rm m rest}}\\right) + \\log_{10}(\\rm k_{1}), \\end{equation} where $\\rm k_{1}$ encodes the contribution of the stellar population of a passively evolving galaxy formed at high redshift ($z > 5$), and $\\rm k_{2}$ the mid-IR-excess parameter that we are looking to calculate for this sample of radio sources. This model assumes that light from the nucleus with intrinsic optical luminosity $L_{\\rm opt}$ is i) absorbed by dust and re-emitted in the mid-IR generating luminosity $[\\nu L]_{10 \\mu \\rm m rest}$ and ii) scattered, generating luminosity $[\\nu L]_{4000 \\rm \\AA rest}$. Fig. 1b shows best-fit lines for two scenarios: 1) all objects were treated as detections (black dotted line), and 2) objects are treated as upper limits according to their 24 $\\mu \\rm m$ detection (blue solid line), where in both cases `SB' and `BL' objects were excluded (Fig. 1b). Averaging these results we deduce $\\rm k_{1} \\sim 0.2$ and $\\rm k_{2} \\sim 0.05$, which agrees well with independent evidence. The value deduced for $\\rm k_{1}$ is in line with the expectations of template spectra of galaxies which formed their stars at high redshift. Optical polarisation studies (e.g. \\cite{kis01}) tell us that $[\\nu L]_{4000 \\rm \\AA rest} \\gtsim$ $0.01 [\\nu L]_{\\rm opt}$, which is consistent with our value of $\\rm k_{2}$ given that QSO SED studies suggest $[\\nu L]_{10 \\mu \\rm m rest} \\sim 0.3 [\\nu L]_{\\rm opt}$ (\\cite{rr95}). We conclude that whenever nuclear accretion is significant in our sample of radio sources, dust in the torus absorbs 30\\% of the photons and dust above and below the torus scatters $\\gtsim$1\\% of the photons."
    },
    "0710/0710.0339_arXiv.txt": {
        "abstract": "The observational evidence for central black holes in globular clusters has been argued extensively, and their existence has important consequences for both the formation and evolution of the cluster. Most of the evidence comes from dynamical arguments, but the interpretation is difficult, given the short relaxation times and old ages of the clusters. One of the most robust signatures for the existence of a black hole is radio and/or X-ray emission. We observed three globular clusters, NGC6093 (M80), NGC6266 (M62), and NGC7078 (M15), with the VLA in the A and C configuration with a 3-$\\sigma$ noise of 36, 36 and 25 $\\mu$Jy, respectively. We find no statistically-significant evidence for radio emission from the central region for any of the three clusters. NGC6266 shows a 2-$\\sigma$ detection. It is difficult to infer a mass from these upper limits due to uncertainty about the central gas density, accretion rate, and accretion model. ",
        "introduction": "Although we do not understand how the nuclei of galaxies form or why they have black holes (BH) at their centers, the correlation between BH mass and bulge velocity dispersion does shed light on their formation and evolutionary histories (Gebhardt et al. 2000a, 2000b: Ferrarese and Merritt 2000). A number of different theories (e.g., Silk \\& Rees 1998; Haehnelt \\& Kauffmann 2000; Robertson et al. 2006) predict a BH mass bulge-velocity-dispersion relation, although they predict different slopes and intercepts for this relation. Exploration of the extreme ends of this relationship will help illuminate the underlying physical model, and in this paper we focus on the low mass end. Black holes at the low end of the relations, with masses between 100 and $10^6~\\Msun$, are generally referred to as intermediate-mass black holes (IMBHs). There is significant evidence that black hole masses from $10^5-10^6~\\Msun$ exist from the work of Barth, Greene \\& Ho (2005) and Greene \\& Ho (2006). To go to yet smaller black hole masses, an extrapolation of the correlation between black hole mass and stellar velocity dispersion suggests studying stellar systems with velocity dispersions of 10--20~\\kms. These dispersions are characteristic of globular clusters. Whether the existence of black holes in globular clusters could shed light on the formation and correlations of supermassive black holes is unknown, but clearly it is a possibility. Furthermore, the existence of massive black holes in clusters will have a significant effect on the cluster evolution. Thus, quantifying the demographics of black holes in clusters may be related to how supermassive black holes grow, and will definitely yield useful information about the evolution of clusters.  Theoretical work suggests that we might expect IMBHs at the centers of steller systems (Ebisuzaki et al. 2001; Portegies Zwart \\& McMillian 2002; Miller \\& Hamilton 2002), although it appears to be difficult to make black holes more massive than 100~$\\Msun$.  Gurkan et al. (2004) suggest that IMBHs may be easy to form through runaway collisions with massive stars.  Discoveries of BHs in globular clusters have been claimed --- G1 in M31 (Gebhardt, Rich \\& Ho 2002) and M15 (van der Marel et al. 2002; Gerssen et al. 2002). In fact, the M15 claim has been made for the past 30 years, starting with the result of Newell, da Costa \\& Norris (1976) and subsequently challenged by Illingworth \\& King (1977). The basic issue is being able to distinguish a rise in the central mass-to-light ratio being due to either a black hole or the expected stellar remnants (neutron stars, massive white dwarfs and solar mass black holes). The most recent M15 result has been challenged by Baumgardt et al. (2003a). The result in G1 has also been challenged by Baumgardt et al. (2003b) but Gebhardt, Rich \\& Ho (2005) include additional data and analysis that support the black hole interpretation.  There has been two further observations which strongly support the existence of a black hole in G1. Trudolyubov \\& Priedhorsky (2004) measure X-rays from G1 using the Chandra Observatory, centered to within 2\\arcsec\\ of the center of G1. Subsequently, Pooley \\& Rappaport (2006) suggest the X-ray emmission is from accretion onto a black hole, and Maccarone \\& Koerding (2006) point out that if a black hole is present then a 30 $\\mu$Jy radio source may be expected. The most significant observation comes from Ulvestad, Greene \\& Ho (2007) who find a 28 $\\mu$Jy (4.5$\\sigma$) emission centered on G1. Other interpretations are a pulsar wind or a planetary nebula. The pulsar wind seems unlikely given the age of G1 and the point-like radio source (an old pulsar would have a large size). A planetary nebula would show optical emission lines which are not seen in the HST or Keck spectra of Gebhardt et al. (2003).  Other studies of the existence of black holes in globular clusters have been less compelling. Colpi, Mapelli, \\& Possenti (2003) use indirect dynamical arguments to suggest a few hundred solar mass black hole in NGC~6752. McLaughlin et al. (2006) provide an estimate of black hole in 47Tuc of $900\\pm900~\\Msun$. To date, there are no published upper limits of black hole masses that are significantly below that expected from an extrapolation of the correlation between black hole mass and stellar velocity dispersion.  While the dynamical arguements strongly support the black hole interpretation in at least G1, the radio emission provides a clear and obvious result. Unfortunately, it is difficult to predict the radio emission from a given black hole mass. The next step is to explore other globular clusters with a similar setup and deep exposures. ",
        "conclusions": "Failure to detect radio radiation at 8.6 GHz from the centers of three globular clusters does not prove that no globular clusters have IMBHs at their centers. Besides not having a black hole, other interpretations include 1) accretion by the BH could be episodic and we happened to observe the BHs in an ``off-state\", 2) the gas density could be much lower compared to galaxies, 3) the radiative efficiency may be lower than assumed (although the assumed efficiencies are already quite low), 4) or the accretion model may not be adequate in general. We would predict, using the relation of Merloni et al. (2003) or using standard accretion models and gas density estimates (as done in Maccarone 2004), that we should have detected radio radiation at 8.6 GHz if accretion is steady and the accretion rate times the Bondi rate is 10$^{-4}\\times$ or higher.  We would not have been able to detect the flux density predicted by a rate of 10$^{-5}\\times$ or less. Ulvestad et al. (2007) estimate the fraction of the Bondi rate of just under 1\\% for G1, but it is difficult to interpret due to the unknown radiative efficiency. For galactic black holes, the radiative efficiencies appear to vary greatly with some lower than $10^{-5}$ (Lowenstein et al. 2001), although consistent with rates of around 10\\% of the Bondi rate.  Models which predict 8.6 GHz flux densities from central BHs in globular clusters above about 25 $\\mu \\rm Jy$/beam can be tested with the VLA currently. The EVLA should produce, for continuum observations, a sensitivity improvement of about a factor of 15, making 8.6 GHz flux densities above about 2$\\mu\\rm Jy$/beam detectable."
    },
    "0710/0710.1804_arXiv.txt": {
        "abstract": "Low-mass X-ray binaries, recycled pulsars, cataclysmic variables and magnetically active binaries are observed as X-ray sources in globular clusters. We discuss the classification of these systems, and find that some presumed active binaries are brighter than expected. We discuss a new statistical method to determine from observations how the formation of X-ray sources depends on the number of stellar encounters and/or on the cluster mass. We show that cluster mass is not a proxy for the encounter number, and that optical identifications are essential in proving the presence of primordial binaries among the low-luminosity X-ray sources. ",
        "introduction": "The first celestial maps in X-rays, in the early 1970s, show that globular clusters harbour more X-ray sources than one would expect from their mass. As a solution to this puzzle it was suggested that these bright ($L_x\\gtap10^{36}$\\,erg/s) X-ray sources, binaries in which a neutron star captures mass from a companion star, are formed in close stellar encounters. A neutron star can be caught by a companion in a tidal capture, or it can take the place of a star in a pre-existing binary in an exchange encounter. Verbunt \\&\\ Hut (1987) showed that the probability of a cluster to harbour a bright X-ray source indeed scales with the number of stellar encounters occurring in it; whereas a scaling with mass does not explain the observations.  With the \\textit{Einstein} satellite a dozen less luminous ($L_x\\ltap10^{35}$\\,erg/s) X-ray sources were discovered in the early 1980. \\textit{ROSAT} enlarged this number to some 55, and now thanks to \\textit{Chandra} we know hundreds of dim X-ray sources in globular clusters. The nature and origin of these dim sources is varied. Those containing neutron stars, i.e.\\ the quiescent low-mass X-ray binaries in which a neutron star accretes mass from its companion at a low rate and the recycled or millisecond radio pulsars, have all formed in processes involving close stellar encounters. The magnetically active binaries, on the other hand, are most likely primordial binaries, with stars that are kept in rapid rotation via tidal interaction. Cataclysmic variables are binaries in which a white dwarf accretes matter from a companion. In globular clusters they may arise either via stellar encounters, or from primordial binaries through ordinary binary evolution -- this is expected to depend on the mass and density of the globular cluster.  In this paper we  describe the classification and identification of the dim sources in Section\\,2, and make some remarks on the theory of their formation in Section\\,3. In Section\\,4 we will discuss a new, and in our view more accurate, way to compare the numbers of these sources with theoretical predictions.  \\begin{figure} \\centerline{ \\parbox[b]{0.55\\columnwidth}{\\psfig{figure=verbuntf1a.eps,width=0.54\\columnwidth,clip=t}} \\parbox[b]{0.45\\columnwidth}{\\psfig{figure=verbuntf1b.ps,width=0.44\\columnwidth,clip=t}} }  \\caption{Left: X-ray hardness-luminosity diagram for dim sources in globular clusters. I: quiescent low-mass X-ray binaries, II: cataclysmic variables III: cataclysmic variables and magnetically active binaries. From Pooley \\&\\ Hut (2006). Right: Colour-magnitude diagram of NGC\\,6752 on the basis of HST-WFPC2 data; objects within X-ray position error circles are marked. Left of the main sequence we find cataclysmic variables, above it active binaries. Updated from Pooley et al.\\ (2002a).\\label{xcol}} \\end{figure} ",
        "conclusions": "\\begin{itemize} \\item mass $M$ is {\\em not} a proxy for collision number $\\Gamma$ \\item the number of dim sources scales both with collision number $\\Gamma$ and with mass $M$ \\item scaling with mass only is not acceptable \\item correct treatment of the background is important, esp.\\ for faint sources \\item to prove the mass-dependence optical identifications are essential \\end{itemize}"
    },
    "0710/0710.3030_arXiv.txt": {
        "abstract": "We present the results of the timing analysis of five {\\it Rossi X-ray Timing Explorer} observations of the Black Hole Candidate GRS 1915+105 between 1996 September and 1997 December. The aim was to investigate the possible presence of a type-B quasi-periodic oscillation (QPO). Since in other systems this QPO is found to appear during spectral transitions from {\\it Hard} to {\\it Soft} states, we analyzed observations characterized by a fast and strong variability, in order to have a large number of transitions. In GRS 1915+105, transitions occur on very short time scales ($\\sim$ sec): to single them out we averaged Power Density Spectra following the regular path covered by the source on a 3D Hardness-Hardness-Intensity Diagram. We identified both the type-C and the type-B quasi-periodic oscillations (QPOs): this is the first detection of a type-B QPO in GRS 1915+105.  As the spectral transitions have been associated to the emission and collimation of relativistic radio-jets, their presence in the prototypical galactic jet source strengthens this connection. ",
        "introduction": "\\label{par:intro}  Systematic variations in the energy spectra and intensity of transient Black-Hole Candidates (BHC) have been recently identified in terms of the pattern described in an X-ray Hardness-Intensity diagram (HID, see Homan et al. 2001, Homan et al. 2005b, Belloni et al. 2005). Four main bright states (in addition to the quiescent state) have been found to correspond to different branches/areas of a square-like HID pattern. In this framework much importance is given to the intermediate states (called Hard Intermediate State, HIMS, and Soft Intermediate State, SIMS) and to the transitions between them, identified from the behaviour in several bands of the electromagnetic spectrum (from radio to hard X-rays, see also Fender, Belloni \\& Gallo 2004 and Homan et al. 2005b) and from the timing properties of the X-ray light curve.  Low-frequency Quasi-Periodic Oscillations (LFQPOs) with centroid frequency ranging from mHz to tens of Hz have been observed in the X-ray flux of many galactic BHCs since the '80s (see van der Klis 2006; McClintock \\& Remillard 2006 and references therein). Three main types of LFQPOs, dubbed Type-A, -B and -C respectively, originally identified in the light curve of XTE J1550-564 (Wijnands et al. 1999; Remillard et al. 2002), have been seen in several sources (see Casella et al. 2005 and references therein). We summarize their properties in Table \\ref{ABC_properties}.  In the context of the state classification outlined above, it is possible to ascribe the three LFQPOs to different spectral conditions (see Table \\ref{ABC_properties}, Homan et al. 2001, Homan \\& Belloni 2005, Belloni et al. 2005).  The type-C QPO is associated to the (radio loud) HIMS and to the low/hard state. It is a common QPO seen in almost all BHCs with a variable centroid frequency correlated with the count rate, a high fractional variability and a high coherence ($Q=\\nu/$FWHM$\\sim$10).  The type-B QPO has been seen only in few systems, although it is being seen in a growing number of sources (see Casella et al. 2005 and references therein). It is a transient QPO associated to spectral transitions from the (radio loud) HIMS to the (radio quiet) SIMS. Its features are a $\\sim$ fixed centroid frequency (around $\\sim$6 Hz), lower fractional variability and $Q$ than type-C. Some authors (Fender, Belloni \\& Gallo 2004; Casella et al. 2004) suggested that these spectral transitions are in turn associated to the emission and collimation of transient superluminal relativistic jets visible in radio band. These jets are seen in a number of sources (GRS 1915+105, XTE J1550-564, GX 339-4, XTE J1859+226, GRO J1655-40, etc.). However, not in all of these sources we could resolve the spectral transition to see the transient QPO.  \\begin{table} \\label{tab:ABC} \\centering \\caption{Summary of type-A, -B and -C LFQPOs properties (from Casella et al. 2005).} \\label{ABC_properties} \\scriptsize \\begin{tabular}{lccc} \\hline \\hline Properties & Type-C & Type-B & Type-A \\\\ \\hline Frequency (Hz) & $\\sim0.1-15$ & $\\sim5-6$  & $\\sim8$ \\\\ Q($\\nu /$FWHM) & $\\sim7-12$ & $\\gtrsim6$ & $\\lesssim3$ \\\\ Amplitude (\\% {\\it rms}) & 3-16 & $\\sim2-4$ & $\\lesssim3$ \\\\ Noise & strong flat-top & weak red & weak red \\\\ Phase lag$^{\\mathrm{a}}$ @$\\nu_{QPO}$ & soft$/$hard$^{\\mathrm{b}}$ & hard & soft \\\\ Phase lag @2$\\nu_{QPO}$ & hard & soft & ... \\\\ Phase lag @$\\nu_{QPO}/2$ & soft & soft & ... \\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] With ``hard lag'' we mean that hard variability lags the soft one. \\item[$^{\\mathrm{b}}$] Trend towards soft lags for increasing QPO frequencies \\end{list} \\end{table}  The spectral properties connected to type-A QPO are similar to those introduced for the type-B. This QPO has been seen in few systems (Casella et al. 2005). It is broader, weaker and less coherent than the type-B QPO.  GRS 1915+105 is a transient BHC discovered on August 15 1992 with the WATCH instrument on board GRANAT (Castro-Tirado et al. 1992, 1994). It is the first galactic source observed to have apparently superluminal transient relativistic radio jets (Mirabel \\& Rodriguez 1994), commonly interpreted as ejection of ultra-relativistic plasma, with a speed close to the speed of light (up to $\\sim98$\\%). Its radio variability was discovered to correlate with the hard X-ray flux (Mirabel et al. 1994). Thanks to VLA-radio observations of these jets, Rodriguez et al. (1995) estimated a distance of 12.5 kpc; more recent estimates attest a distance of $6.5 \\pm 1.6$ kpc (Kaiser et al. 2005). The mass of the compact object was estimated through IR spectroscopic studies (Greiner, Cuby \\& McCaughrean 2001, Harlaftis \\& Greiner 2004) to be $14.0 \\pm 4 M_{\\odot}$ which unambiguously makes GRS 1915+105 a BHC.  The various and rich phenomenology of this source was classified by Belloni et al. (2000): they analyzed 163 RXTE observations, showing that the complex behaviour of GRS 1915+105 can be described in terms of spectral transitions between three basic states, A, B and C (not to be confused with the name of the LFQPOs introduced before), that give rise to 12 variability classes. The non standard behaviour of GRS 1915+105 (it is a very bright transient source continuing the same outburst started in 1992) was interpreted as that of a source that spends all its time in Intermediate States (both in its hard and soft flavors), never reaching the LS or the quiescence (see e.g. Fender \\& Belloni, 2004).  GRS 1915+105 also shows strong time variability on time scales of fractions of second, revealing low- and high-frequency QPOs (see Morgan et al. 1997) whose properties (frequency and fractional variability) are tightly correlated with the spectral parameters (Morgan et al. 1997; Muno et al. 1999; Markwardt et al. 1999; Rodriguez et al. 2002a, 2002b; Vignarca et al. 2003). In particular, all LFQPOs observed from this system can be classified as type-C QPOs. Although GRS 1915+105 makes a large number of fast state transitions, which have been positively associated to radio activity and jet ejections, no type-B QPO has been observed to date.  \\begin{table*} \\centering \\caption{Log of the 5 RXTE/PCA observations analyzed in this work} \\label{log_obs} \\begin{tabular}{c c c c c c} \\hline \\hline N$^{\\circ}$ & Obs. Id.& Date & Starting MJD & Exp. (s) & Classification\\\\ \\hline 1 & 10408-01-35-00 & 1996 Sep 09 & 50348.271 & 9448& $\\mu$\\\\ 2 & 20402-01-45-03 (\\#1, 2, 3) & 1997 Sep 09 & 57700.250 & 10038& $\\beta$\\\\ 3 & 20402-01-53-00 & 1997 Oct 31 & 50752.013 & 9656& $\\beta$\\\\ 4 & 20402-01-53-01 \\& 20402-01-53-02(\\#1) & 1997 Nov 04-05 & 50756.412 & 6322&$\\mu$\\\\ 5 & 20402-01-59-00 & 1997 Dec 17 & 50799.091 & 9784& $\\beta$\\\\ \\hline \\end{tabular} \\flushleft Obs. 4 is composed of two orbits from two separate observations\\\\ \\# indicates the number of the RXTE orbit within the observation\\\\ Classification is from Belloni et al. (2000). \\end{table*}  In this paper we present the discovery with RXTE of the type-B QPO in the X-ray light curve of GRS 1915+105. The QPO was present during fast spectral transitions that we identify with the HIMS to SIMS transition observed in other BHCs. ",
        "conclusions": "\\label{par:discussion}   We analyzed 5 RXTE/PCA observations collected during the first two years of the mission. In the power density spectra of all five observations we detected several peaks which we identify with two different types of QPO already seen in many other BHCs: the type-C and type-B QPOs. This is the first identification of a type-B QPO in GRS 1915+105. To detect it, we looked in detail at spectral transitions in observations characterized by a fast and intense variability. Spectral transitions in GRS 1915+105 are usually very fast, often occurring on timescales of $\\sim$ seconds. This is at variance with most of other black-hole binaries in which spectral transitions are observed to last hours or days. In order to study the spectral transitions in GRS 1915+105 we performed an energy-dependent timing analysis by averaging power spectra on the pattern the source recursively tracks in the 3-dimensional hardness-hardness-intensity diagram.  Applying this method, we found a type-B QPO in all five observations. In all of them, we detected the type-B QPO together with the type-C, in region \\# 3 of the HHID (see Tab.  \\ref{tab:zone}). In three of them, we also detected a type-B QPO alone, in region \\# 4 of the HHID. Two of these observations belong to class $\\mu$ and one to class $\\beta$, which excludes any relation between the class of variability and the presence of the type-B QPO. This means that the presence of the long hard intervals (which differentiate the $\\beta$-class from the $\\mu$-class light curves) does not influence the fast timing properties of the source outside these intervals.  We could not find any property (as e.g. hardness, rate) correlated with the presence of the type-B QPO alone in region {\\em 4} of the HHID. In three observations we also found a type-B bump in region {\\em 2}. No correlations were found neither between the presence of the type-B bump and the type-B QPO in region {\\em 4} nor with the hardness and the count rate.  \\begin{figure} \\begin{tabular}{c} \\resizebox{7cm}{!}{\\includegraphics{LAG_ALL_DEFINITIVO.ps}} \\end{tabular} \\vspace{-2.0cm} \\caption{Phase lags of the detected QPOs in all observations. For each QPO we extracted the phase lag in a range centered at the QPO peak frequency and corresponding to the width itself ($\\nu_p \\, \\pm \\, FWHM/2$). Errors bar on the X-axis are not shown for clarity.} \\label{fig:lag_all} \\end{figure}  \\subsection{QPOs identification} \\label{subpar:disc_identif}  In order to identify the two types of QPO that we found in our data sets, we compare them with the known LFQPOs in BHCs. In particular we first analyze their position and behaviour in the Hardness-Intensity diagram (right panel of Figure \\ref{fig:cd_hid_arrows}). When the source moves through regions \\#{\\em 1} and \\#{\\em 2} up to region \\#{\\em 3} the first QPO shows a behaviour very similar to that of type-C QPOs: its frequency is correlated with the count rate and is inversely correlated with the hardness. At a certain hardness, this QPO disappears. A second type of QPO is also detected: as in the case of type-B QPOs, this second QPO appears in a narrow frequency range (often around $\\sim$6 Hz) and in a limited range in hardness.  Its frequency and quality factor appear to be slightly correlated with the count rate, particularly when at its lowest frequencies (2.44-2.84 Hz). We interpret this as the first evidence of an increase of the coherence of the type-B QPO from $Q<2$ when at low frequencies to $Q>2$ when reaching frequencies around $\\sim$6 Hz, possibly suggesting the presence of a resonance at this frequency (see Casella et al. 2004).  The combined evolution of the QPOs in GRS 1915+105 is strongly reminiscent of the known behaviour of type-C and type-B QPOs in BHCs (see Casella et al. 2005 and references therein).  To verify this identification, we plot in Figure \\ref{fig:rms_1} the {\\it rms} fractional variability of the detected QPOs as a function of their frequency for three energy bands. The two QPO types have a somewhat similar energy dependency (being stronger at high energy) but they clearly show different behaviours in these diagrams. In each of the three panels, two well-identified groups of points are evident. A comparison of these two groups with Figure \\#3 of Casella et al. 2004 helps to classify the observed QPO in GRS 1915+105: the first group is diagonally spread across the plots, covering the whole frequency range between $\\sim2$ and $\\sim15$ Hz and a large range in {\\it rms} (particularly at high energies, see the right panel of Figure \\ref{fig:rms_1}). The second group is clustered both in frequency (between $\\sim2.5$ and $\\sim7$ Hz) and in fractional {\\it rms}. The observed behaviour is clearly consistent with that known to be typical of type-C and type-B QPOs in BHCs (see Casella et al. 2004, 2005).  \\subsection{Phase lags} \\label{subpar:disc_lag}  The association between QPO-type and phase lag in literature is not conclusive: although an average behaviour can be identified (see Tab. \\ref{tab:ABC}, Casella et al. 2005 and reference therein) there are a number of exceptions (see e.g. Belloni et al. 2005 and Homan et al. 2005a). Nevertheless we performed a phase-lag analysis in order to have a comprehensive view of the behaviour of the QPOs in GRS 1915+105.  On the basis of the analyzed data it is not possible to characterize unambiguously the phase lag behaviour of any of the two types of QPO we observe: lags appear to vary between different observations, although both types show in average negative values of phase lag (see Fig. \\ref{fig:lag_all}).  This can possibly due to the fast variability of the analyzed light curves: we extracted power spectra over time intervals 2 seconds long, without applying any procedure to detrend the variability on longer time scales. However, this variability is very strong (see Fig. \\ref{fig:mu_beta_class}), which results in a leakage at higher frequencies as strong as to actually dominate the phase-lag continuum.  \\subsection{GRS 1915+105 as a ``normal'' source} \\label{subpar:disc_normal}  The type-B QPO was detected in a few BHCs (see Casella et al. 2005), and associated to spectral transitions from HIMS to SIMS (Homan \\& Belloni 2005, Belloni et al. 2005 and reference therein). Some authors (Gallo et al 2003; Fender, Belloni \\& Gallo 2004) suggested a relation between these spectral transitions and the emission and collimation of transient relativistic radio jets: if we consider the type-B QPO as the signature of radio jet emission, we have this ``signature'' also in the prototypical galactic jet source.  As already pointed out by Belloni et al. (2005), the non-detection of a type-B QPO in GRS 1915+105 was rather interesting, especially if you interpret the X-ray$/$radio correlation in this source and in other transients in the framework of the same model (Fender, Belloni \\& Gallo 2004). In GRS 1915+105 the association between X-ray and radio activity is well known (Pooley \\& Fender 1997, Mirabel et al. 1998, Fender \\& Belloni 2004 and references therein). Belloni et al. 2005 suggested that the elusiveness of this QPO (in GRS 1915+105) could be due to high velocity of the movement of the source through the HID. The result presented in this work thus strengthens the interpretation of GRS 1915+105 in the framework of the same model of other BHCs: a type-B QPO appears in correspondence of spectral transition from the B-state to the A-state (Transition, Region {\\em 3}, Stars in Figure \\ref{licu_symbols}) and when the source is in the A-state (Region {\\em 4}, Upward Triangles in Figure \\ref{licu_symbols}). In the light curve in Figure \\ref{licu_symbols} we see state oscillations CTAB TAB TAB that we can identify as fast passages between the HIMS and the SIMS of the other BHCs (Casella et al. 2004; Fender, Belloni \\& Gallo 2004). The type-B QPO appears also in correspondence of the transition from the B-state to the C-state (Squares-Stars-Circles in Figure \\ref{licu_symbols}), therefore not in an oscillation event but in a transition from a soft to a hard state.  Furthermore, GRS 1915+105 is at present the heaviest black hole (although uncertainties on BH mass values are rather large, see McClintock \\& Remillard 2006 and references therein) in which a type-B QPO has been observed at 6 Hz (when its quality factor $Q>2$). This strengthens the already known independence of this type of QPO on the mass of the compact object (see Casella et al. 2005).  \\subsection{X-ray$/$radio association} \\label{subpar:disc_radio}  Klein-Wolt et al. (2002) made a detailed study of simultaneous radio$/$X-ray observations of GRS 1915+105, focusing in particular on radio oscillation events, and found that $\\beta$ class observations are usually associated to strong radio oscillations, while $\\mu$ class observations are usually associated to a weak steady radio activity. These authors point out that the long hard intervals (which are the main macroscopic difference between the two variability classes) appear to be directly linked to the production of radio oscillations. We found the type-B QPO in both classes $\\beta$ and $\\mu$. We could not find any timing property (on time scales shorter than 2 seconds) clearly differentiating the two classes.  This apparently weakens the link between the presence of the type-B QPO and radio activity. Unfortunately, the lack of radio coverage during the analyzed RXTE observations does not allow any conclusive result. In order to clarify this issue it will be fundamental to extend our analysis to RXTE observations for which a radio coverage is available."
    },
    "0710/0710.3340_arXiv.txt": {
        "abstract": "We report the first determination of a distance bracket for the high-velocity cloud (HVC) complex~C. Combined with previous measurements showing that this cloud has a metallicity of 0.15 times solar, these results provide ample evidence that complex~C traces the continuing accretion of intergalactic gas falling onto the Milky Way. Accounting for both neutral and ionized hydrogen as well as He, the distance bracket implies a mass of 3--14\\tdex6~\\Msun, and the complex represents a mass inflow of 0.1--0.25~\\Msunpyr. We base our distance bracket on the detection of \\CaII\\ absorption in the spectrum of the blue horizontal branch star SDSS\\,J120404.78+623345.6, in combination with a significant non-detection toward the BHB star BS\\,16034-0114. These results set a strong distance bracket of 3.7--11.2~kpc on the distance to complex~C. A more weakly supported lower limit of 6.7~kpc may be derived from the spectrum of the BHB star BS\\,16079-0017. ",
        "introduction": "\\par The evolution of galaxies is strongly driven by the gas in the interstellar medium. There is strong evidence for the infall of new material that provides fuel for galaxy growth. This gas may originate in accreted satellite galaxies, as gas tidally pulled out of passing galaxies, or from pristine intergalactic gas. The cool, infalling clouds appear to be embedded in an extended (100--200~kpc radius) hot Corona (Sembach et al.\\ 2003). Indirect evidence for infalling gas is provided by two arguments: (a) At the current rate of star formation, all of the ISM will be turned into stars within about a Gyr. (b) The narrowness of the distribution of metallicities of long-lived stars implies that the metallicity of the ISM remains more or less constant over a Hubble time, which can happen if there is a continuing inflow of low-metallicity material with a present-day rate of about 1~\\Msunpyr. Item (b) is known as the ``G-dwarf problem'' (van den Bergh 1961). Using the infall hypothesis to solve it has been the subject of much theoretical work (see e.g.\\ Pagel 1997 for a good summary). Continuing infall is essential in detailed numerical modeling of the chemical evolution of the Galaxy and the development of abundance gradients (e.g.\\ Chiappini et al.\\ 2001 and references therein). Infall of low-metallicity gas also seems necessary to reproduce the relatively high abundance of deuterium measured in the local interstellar medium (Linksy et al.\\ 2006). \\par Direct observational evidence for infalling low-metallicity gas is provided by the high-velocity clouds (HVCs; see reviews by Wakker \\& van Woerden 1997; Richter 2006). Subsolar metallicities have now been determined for eleven clouds (see van Woerden \\& Wakker 2004 for a summary). In particular, the metallicity of complex~C is well established as 0.15 times solar (see summary by Fox et al.\\ 2004). Complex~C also has a high deuterium abundance (Sembach et al.\\ 2004). Distance brackets have been more elusive, with just one known before 2006 (8--10~kpc for complex~A -- van Woerden et al.\\ 1999a; Wakker et al.\\ 2003). Thom et al.\\ (2006) derive an 8.8~kpc upper limit for cloud WW\\,35, while in a separate paper (paper~I, Wakker et al.\\ 2007), we present new results for two HVCs (9.8--15.1~kpc for complex~GCP and 5.0--11.7~kpc for the Cohen Stream). In this letter we report a distance bracket for the HVC covering the largest sky area -- complex~C. We summarize our method in Sect.~2. The data are described in Sect.~3, the results in Sect.~4, while in Sect.~5 we summarize the implications. ",
        "conclusions": "\\par Forty years after the first attempt (Prata \\& Wallerstein 1967), we report the first successful detection of interstellar \\CaII\\ H and K absorption from HVC complex~C. This sets an upper limit on the distance of core CIII (left side in Fig.~1) of 11.2~kpc. For core CI (right side in Fig.~1) we find a lower limit of 3.7~kpc, possibly 6.7~kpc. Although the stars are 27\\deg\\ apart on the sky, it is still safe to conclude that complex~C is located at Galactocentric radius $<$14~kpc, and lies high above the Galactic plane ($z$=3--9~kpc). A more precise determination requires a lower limit for core CIII and an upper limit for CI. \\par Integrating $N$(\\HI) across the cloud, we estimate $M$(\\HI) as 0.7--6\\tdex{6}~\\Msun. \\Ha\\ emission has also been detected (Tufte et al.\\ 1998). We can assume either that the H$^+$ and \\HI\\ are thorougly mixed or that the H$^+$ originates in a photoionized skin around the cloud. In either case, the observed \\Ha\\ intensity suggests that there is roughly as much ionized as neutral gas. \\par We can also estimate the mass inflow associated with complex~C, using a method described in paper~I. Including the neutral and ionized hydrogen, as well as a 40\\% contribution from helium, we derive that complex~C represents about 0.1--0.25~\\Msunpyr\\ of infalling gas. This is a substantial fraction of the theoretically required amount of 1~\\Msunpyr. Other HVCs may contribute the rest, but we have not yet determined distances and metallicities for the most likely candidates. \\par From our results, we conclude that the mystery of the distances to the HVCs is beginning to be solved. The evidence shows that several HVCs are located in the upper reaches of the gaseous Galactic Halo and that they contribute significantly to the inflow of metal-poor gas onto the Galaxy. Once the mass inflow rate is constrained from observations of a sufficient number of HVCs, the next step will be to determine their three-dimensional structure, so that we can use their velocities and galactic location to derive orbits and solve the outstanding mystery of their ultimate origins.  \\bigskip Acknowledgements \\par B.P.W., D.G.Y, R.W. and T.C.B. acknowledge support from grant AST-06-07154 awarded by the US National Science Foundation. T.C.B. also acknowledges NSF grants AST 04-06784 and PHY-02-16783; Physics Frontier Center/Joint Institute for Nuclear Astrophysics (JINA). \\par Some of the data presented 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 the National Aeronautics and Space Administration. The Observatory was made possible by the generous financial support of the W.M. Keck Foundation. \\par The William Herschel telescope is operated on the Island of La Palma by the Isaac Newton group in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrophysica de Canarias. \\par 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/. The SDSS is managed by the Astrophysical Research Consortium for the Participating Institutions. The Participating Institutions are the American Museum of Natural History, Astrophysical Institute Potsdam, University of Basel, University of Cambridge, Case Western Reserve University, University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington."
    },
    "0710/0710.1345_arXiv.txt": {
        "abstract": "The results obtained from a study of the mass distribution of 36 spiral galaxies are presented. The galaxies were observed using Fabry-Perot interferometry as part of the GHASP survey. The main aim of obtaining high resolution H$\\alpha$ 2D velocity fields is to define more accurately the rising part of the rotation curves which should allow to better constrain the parameters of the mass distribution. The H$\\alpha$ velocities were combined with low resolution HI data from the literature, when available. Combining the kinematical data with photometric data, mass models were derived from these rotation curves using two different functional forms for the halo: an isothermal sphere and an NFW profile. For the galaxies already modeled by other authors, the results tend to agree. Our results point at the existence of a constant density core in the center of the dark matter halos rather than a cuspy core, whatever the type of the galaxy from Sab to Im. This extends to all types the result already obtained by other authors studying dwarf and LSB galaxies but would necessitate a larger sample of galaxies to conclude more strongly. Whatever model is used (ISO or NFW), small core radius halos have higher central densities, again for all morphological types. We confirm different halo scaling laws, such as the correlations between the core radius and the central density of the halo with the absolute magnitude of a galaxy: low luminosity galaxies have small core radius and high central density. We find that the product of the central density with the core radius of the dark matter halo is nearly constant, whatever the model and whatever the absolute magnitude of the galaxy. This suggests that the halo surface density is independent from the galaxy type. ",
        "introduction": "Rotation curves are a fundamental tool for studying the dynamics and mass distribution in galaxies. The distribution of the total mass can then be compared with the distribution of visible light, assuming a certain mass-to-light (M/L) ratio. Observations have clearly established that dark matter is needed for explaining the rotation velocities observed in the outer parts of spirals. The dark matter is most often considered as being distributed in a spherical dark halo but its density profile, especially at small radii, is still a matter of debate (Blais-Ouellette et al. 2001, de Blok \\& Bosma 2002, Swaters et al. 2003, Navarro 2004, Graham et al. 2006, Kusio de Naray 2006, Hayashi et al. 2007).  However, mass models of spiral and dwarf galaxies have well-known degeneracies (Barnes et al. 2004) that prevent a unique mass decomposition, the most important being the unknown value of the stellar mass-to-light ratio (Dutton et al. 2005). Unfortunately, stellar population models (e.g. de Jong \\& Bell 2006) still cannot predict accurately the (M/L) values of the stellar disks based on their colors. The main problem comes from the fact that the normalization of this relation (color vs M/L) depends critically on the shape of the stellar IMF at the low mass end. It is well known that the faint stars contribute significantly to the mass but not to the luminosity and color of the stellar disks. This means that it will prove difficult to tighten this relation and lessen the effect of the disk-halo degeneracy.  Cosmological numerical simulations favor cuspy dark halos, although the value of the inner slope $\\gamma$ of the radial density profile, where $\\rho$ $\\propto$ r $^{\\gamma}$ (see equation 6 of section 4) depends on the authors, with $\\gamma$ = -1.5 for Moore et al. 1999 as well as Fukushige \\& Makino 2001 or -1.0 for Navarro, Frenk \\& White 1997 (hereafter NFW). Recent simulations (Graham et al. 2006) have extrapolated inner logarithmic profile slope ranging from -0.2 to -1.5, with a typical value at 0.1 kpc around -0.7.  Most observers conclude however that $\\gamma$ is closer to 0.0 than to -1.0 and that mass models give better results with an isothermal (or pseudo-isothermal) sphere halo rather than a NFW profile (e.g. Blais-Ouellette et al. 2001, de Blok \\& Bosma 2002, Swaters et al. 2003, de Blok et al. 2003, de Blok 2005, Kassin et al. 2006a). Most of these observations are based on the rotation curves of dwarf and Low Surface Brightness (LSB) galaxies (Kuzio de Naray et al. 2006). It is still unclear if this is true also for High Surface Brightness (HSB) systems and for all morphological types.  One problem in this cusp-core debate is that numerical simulations mainly predict the halo density profile shape at the time of formation while galaxies are observed after many Gyrs of evolution. The problem is that internal dynamics and interaction between the dark halo and the luminous disk (e.g. adiabatic contraction: Dutton et al. 2005) and interaction with the environment (Maccio et al. 2007) may have altered the shape of the halo density profile.  Moreover, the shape of the gravitational potential in CDM halos may explain the core-like halo density profiles seen in LSB systems. Hayashi et al. 2007 suggest that galactic disks may be forming in elliptical gravitational potential. This could result in significant non-circular motions in systems such as LSBs which would mimic constant density cores. Thus, taking into account the 3D shape of the dark mass distribution could reconcile the constant density cores observed in LSBs with the predicted cuspy mass profiles of CDM halos.  Finally, it is unclear if CDM simulations, mainly obtained to compare with the large scale structure, have sufficient resolution to reliably probe the kpc scales necessary to compare with the observationally derived dark matter halo density profiles on the scale of a few halo core radii. As discussed by Navarro (2004), unfortunately rotation curves constraints are strongest where numerical simulations are the least reliable. In fact, rotation curves are usually compared with extrapolations of the simulation data that rely heavily on the applicability and accuracy of fitting formulae such as the NFW profile to regions that may be compromised by numerical artifact. This is especially true for LSB and dwarf galaxies (for instance Dutton et al. 2005).  Blais-Ouellette et al. 1999 and Barnes et al. 2004 have shown the necessity of optical integral field spectroscopy to accurately determine the rotation curves in the inner parts of spiral and dwarf galaxies, for which the HI data are affected by beam smearing (Swaters et al. 2000, van den Bosch et al. 2000). While the use of two-dimensional data does not necessarily alter the halo parameters derived from optical long-slit data it however decreases the uncertainties by roughly a factor of 2 (de Naray et al. 2006). Blais-Ouellette et al. 2001 and Blais-Ouellette et al. 2004 also pointed out the great sensitivity of the mass distribution parameters to the inner rotation curve. The ideal rotation curve will therefore combine high resolution optical data, for the inner part, with radio data, for the outer part extending well beyond the luminous disk. ",
        "conclusions": "Our study, based on 36 galaxies of different morphological types, confirms the result already claimed by other authors about the shape of the dark matter halo in the center of spiral galaxies, namely that its density profile is probably closer to an isothermal sphere profile than to an NFW profile: DM halos are rather flat than cuspy. Our mass models (ISO and NFW) are compared when fitting the rotation curve as a whole, however a comparison of the chi2 values limited to the central regions would be still clearer since the inner slope of the RC found with the NFW model is systematically steeper and above the data points, compared with the ISO model.  Thus far, most of the results found in the literature concerned dwarf or Low Surface Brightness galaxies. 9 galaxies of our sample are LSB (UGC 2304, 5272, 5279, 5789, 7524, 7699, 9465, 11707 and 12060, with types going from Scd to Im) and, for all of them, we systematically find better results with the ISO model (smaller $\\chi^{2}$) than with NFW, except UGC 2304 for which both are equivalent. Interestingly, our study suggests that this holds for most spirals since the 36 galaxies of this study have morphological types ranging from Sab to Im. For almost all of them, the best fit is obtained with the ISO model (the reverse is found only for 4 galaxies out of 36). Also, no significant difference can be seen when comparing the quality of the fits obtained with the NFW and the ISO model as a function of the morphological type.  However, a difference can be seen in the way the rotation curves are decomposed into several components, with the halo being the main component for late types and the disk (or disk + bulge) the main component for early types. This result, although needing to be confirmed with a larger sample, merely reflects the well known fact that later type galaxies are more dark matter dominated.  Finally, we confirm different halo scaling laws seen previously by other authors such as C\\^ot\\'e et al (2000), Kormendy \\& Freeman 2004 and Barnes et al. 2004. Among those, it appears clearly that low luminosity galaxies have small core radius and high central density, the product of the two parameters being nearly constant with absolute magnitude. This means that the galaxy halo surface density is independent of galaxy type or luminosity. Trends are also seen for the core radius as a function of luminosity and color but should be confirmed with a larger sample."
    },
    "0710/0710.4560_arXiv.txt": {
        "abstract": "We study the effect of primordial nongaussianity on large-scale structure, focusing upon the most massive virialized objects.  Using analytic arguments and N-body simulations, we calculate the mass function and clustering of dark matter halos across a range of redshifts and levels of nongaussianity.  We propose a simple fitting function for the mass function valid across the entire range of our simulations.  We find pronounced effects of nongaussianity on the clustering of dark matter halos, leading to strongly scale-dependent bias.  This suggests that the large-scale clustering of rare objects may provide a sensitive probe of primordial nongaussianity.  We very roughly estimate that upcoming surveys can constrain nongaussianity at the level $|\\fnl|\\lesssim 10$, competitive with forecasted constraints from the microwave background. ",
        "introduction": "One of the fundamental predictions of standard (single-field, slow-roll) inflationary cosmology is that the density fluctuations in the early universe that seeded large-scale structure formation were nearly gaussian random (e.g.\\ \\cite{maldacena,Acquaviva:2002ud,Creminelli:2003iq,Lyth_Rodriguez,Seery_Lidsey}). Constraining or detecting non-gaussianity (NG) is therefore an important and basic test of the cosmological model.  To the extent that it can be measured, gaussianity has so far been confirmed; the tightest existing constraints have been obtained from observations of the cosmic microwave background \\cite{wmap3,Creminelli_wmap}. Recently, several inflationary models have been proposed which predict a potentially observable level of nongaussianity, see \\eg \\cite{ArkaniHamed:2003uz,Bartolo:2003jx,Lyth:2005du,Rigopoulos:2005ae, Allen:2005ye,Chen_DBI,Barnaby:2006cq,Barnaby:2006km,Barnaby:2007yb, Sasaki:2006kq,Chen:2006nt,Chen_Easther_Lim,Battefeld_Easther, Assadullahi:2007uw,Battefeld:2007en,Shandera} and \\cite{Bartolo:2004if} for a review. Improved limits on NG would rule out some of these models; conversely, a robust detection of primordial nongaussianity would dramatically overturn standard inflationary cosmology and provide invaluable information about the nature of physical processes in the early universe. In this regard, there has been a resurgence in studying increasingly more sophisticated methods and algorithms to  constrain (or, if we are lucky, detect) nongaussianity \\cite{Babich:2005en,Babich_shape,Creminelli_estimators,Smith_Zaldarriaga,Fergusson_Shellard}.  Nongaussianity manifests itself not only in the cosmic microwave background \\cite{Falk_Ran_Sre,Luo_Schramm,Gangui_etal,Wang_Kam}, but also in the late-time evolution of large-scale structure.  For example, detailed measurements of higher order correlations like the bispectrum or trispectrum of galaxy clustering could provide a handle on primordial nongaussianity \\cite{Verde:2000vr,Scoccimarro:2003wn,Sefusatti:2007ih}.  The abundance of galaxy clusters, the largest virialized objects in the universe, has also long been recognized as a sensitive probe of primordial NG \\cite{Lucchin:1987yv,Robinson:1999se,Benson:2001hc,Matarrese:2000iz,verde01,Scoccimarro:2003wn,Komatsu:2003fd}.  Because clusters are rare objects which form from the largest fluctuations on the tails of the density probability distribution, their abundance is keenly sensitive to changes in the shape of the PDF such as those caused by nongaussianity.  Large statistical samples of massive clusters have already been compiled from wide-area optical imaging and spectroscopic surveys such as the Sloan Digital Sky Survey \\cite{maxBCG,Koester:2007bg}, the Two-Degree Survey \\cite{Eke:2004ve}, and from the Red Sequence Survey \\cite{Yee:2007if} and from X-ray surveys using the Chandra and XMM-Newton observatories \\cite{Willis:2005ag,Valtchanov:2003it}.  Future missions, such as the Dark Energy Survey, Supernova/Acceleration Probe and Large Synoptic Survey Telescope, will detect and study tens of thousands of clusters, revolutionizing our understanding of cluster physics as well as providing important constraints on cosmology \\cite{Haiman_Mohr_Holder,Majumdar_Mohr,Wang_Haiman, Battye_Weller,Lima_Hu_05, Marian_Bernstein,Takada_Bridle}.  To exploit the potential of these upcoming surveys as probes of primordial nongaussianity, it is important to calibrate the effects of NG on the abundance and clustering of virialized objects.  While no previous work has attempted to quantify the effects of NG on halo clustering, several groups over the past decade have constructed fitting formulae for the halo mass function \\cite{Robinson_Baker,Robinson_Gawiser_Silk,MVJ}.  All of this work, however, was analytic and relied on the validity of the Press-Schechter \\cite{press-schechter} formalism, plus various further approximations. The resulting analytic estimates are, in general, rather cumbersome to compute and have questionable accuracy.  As discussed below, the Press-Schechter model provides only a qualitative description of halo abundance, and fails to reproduce the halo mass function to within an order of magnitude over the mass and redshift range accessible to current and future cluster surveys.  Therefore, analytic models for NG cluster abundance based on the Press-Schechter ansatz may not be sufficiently accurate.  Given the high-quality data soon to be available, a much more precise calculation of cluster statistics will be required.  Quite recently, two groups have attempted to quantify the mass function of clusters in NG models using N-body simulations \\cite{Kang,Grossi}, reaching contradictory conclusions.  In this paper, we use analytic arguments and numerical simulations to estimate the effect of NG on the abundance and clustering of virialized objects. Because N-body simulations can be expensive and there is a wide NG parameter space, we also strive to make our results useful to a cosmologist who is not necessarily equipped with the machinery or patience to run simulations or evaluate difficult analytic expressions. To this end, we provide a simple, physically motivated fitting formula for the halo mass function and halo bias, which we calibrate to our N-body simulations.  Our main results are that the mass function and correlation function of massive halos can be significantly modified by primordial nongaussianity.  We find a somewhat weaker effect of NG on the mass function than previous analytic estimates.  We also show analytically and numerically that NG strongly affects the clustering of rare objects on large scales, implying that measurements of the large-scale power spectrum can place stringent bounds on NG.  The plan of the paper is as follows.  In Section \\ref{anal} we derive analytic expressions for the abundance and clustering of rare peaks. In Section \\ref{sec:method} we describe our N-body simulations, followed in Section \\ref{sec:mf} by a discussion of our measured halo mass function, and our fitting formula for the mass function.  In Section \\ref{sec:pk} we present measurements of halo clustering within our simulations, and in Section \\ref{sec:cosmo} we discuss cosmological implications of our findings. ",
        "conclusions": "We have quantified the effects of primordial nongaussianity on the abundance and power spectra of massive halos.  Our two principal results are as follows.  First, we have provided a new fitting formula for the halo mass function. The formula is based on matching halos in Gaussian and non-Gaussian simulations: for $\\fnl>0$ the corresponding halos are more massive than in the Gaussian case, and vice versa. The formula is consistent with the measured mass function from our simulations to within $\\sim10\\%$ over the entire range of masses and redshifts that we consider. Being essentially a convolution of the Gaussian mass function and a Gaussian kernel (Eqs.~(\\ref{eq:mf_conv})-(\\ref{eq:rms_Mf_def})), the formula is also easy to use and does not require estimating the extreme tails of the nongaussian PDF of the density field. Our results also indicate that previous work based on Extended Press-Schechter type formulae overestimated the effects of nongaussianity on the abundance of halos by a factor of $\\sim 2$ over the relevant mass scales.  Secondly, we showed both analytically and numerically that nongaussianity (in the $\\fnl$ model) leads to strong scale dependence of the bias of dark matter halos.  We find remarkably good agreement between our analytic expression and our numerical results. Measurement of the power spectrum of biased objects therefore provides a new avenue to detect and measure nongaussianity.  While cluster counts can constrain NG at a level comparable to existing CMB constraints, $|\\fnl|\\lesssim 100$, we found that future large-scale redshift surveys can potentially do much better, roughly $|\\fnl|\\lesssim 10$.  We do not find significant degeneracies between $\\fnl$ and dark energy parameters in our Fisher matrix calculations, either for mass function measurements or power spectrum measurements.  More precise estimates will require considerably more sophisticated treatments than we have attempted in our illustrative examples above.  We close this paper by considering, in light of our findings, the optimal methods for constraining NG of the $\\fnl$ form.  Measurements of the power spectrum would appear the most promising; observations of high redshift, highly clustered objects on large scales would allow the strongest constraints on the scale-dependent bias signature of $\\fnl$.  Fortunately, upcoming BAO surveys will likely provide the necessary observations of, e.g. luminous red galaxies (LRGs).  Photometric surveys may also be useful in this regard.  Since the effects of NG are most pronounced on large scales, rather than small scales, precise spectroscopic redshifts may not be necessary.  Photometric redshifts with errors of order $\\Delta z\\approx 0.03$ have already been achieved for LRGs and for optically selected groups and clusters with prominent red sequences \\cite{Padmanabhan:2004ic,Ilbert:2006dp,Yee:2007if}.  At $z=0.5$, this corresponds to roughly 100 $h^{-1}$Mpc comoving, fairly small compared to the $\\sim$Gpc scales where NG becomes most important.  Since photometric surveys can cover wider areas more deeply than spectroscopic surveys, they may turn out to provide tighter bounds.  Besides their abundance and clustering, the internal properties of massive halos may also be sensitive to nongaussianity. For instance, the concentrations and substructure content of massive halos have been found to depend upon primordial NG \\cite{avila03}. Our simulations lacked sufficient force resolution to explore this in detail, but we note in passing that multiple groups find a tension between observations of massive lensing clusters and theoretical predictions for Gaussian perturbations \\cite{rcs,arcs04,hennawi07,broadhurst08}.  Another intriguing possibility for probing primordial NG is to use statistics of the largest voids in the universe.  Just as the abundance and clustering of high density peaks are affected by nongaussianity, so are the same properties for deep voids (albeit with an opposite sign, c.f.\\ Fig.~\\ref{fig:slice_sims}).  In a sense, because voids are not as nonlinear as overdense regions, their properties are more easily related to the initial Lagrangian underdensities whose statistics are straightforward to compute.  Voids may be detected at high redshift as a deficit of Lyman-$\\alpha$ forest absorption features in QSO spectra. The Sloan Digital Sky Survey (SDSS) has already measured spectra for high redshift QSO's over a roughly $\\sim$8000 deg$^2$ area, corresponding to a volume of $\\gtrsim 30 ({\\rm Gpc}/h)^3$ \\cite{sdss_lyaf}.  Each QSO spectrum typically probes $\\sim 400 h^{-1}$ Mpc, and the typical transverse separation between QSO sightlines in SDSS is $\\sim 100 h^{-1}$ Mpc, (P.\\ McDonald, priv.\\ comm.)  so measurements of the clustering of $\\sim 10$ Mpc-sized voids on $\\sim$ Gpc scales may already be feasible.   Finally, we note that our conclusions are based on simulations implementing a very specific type of local primordial nongaussianity quantified by the $\\fnl$ parameter. The validity of our conclusions in the context of other type of primordial nongaussianity is the subject of ongoing studies."
    },
    "0710/0710.2493_arXiv.txt": {
        "abstract": "{The annihilations of WIMPs produce high energy gamma-rays in the final state. These high energy gamma-rays may be detected by IACTs such as the H.E.S.S. array of Imaging Atmospheric Cherenkov telescopes. Besides the popular targets such as the Galactic Center or galaxy clusters such as VIRGO, dwarf spheroidal galaxies are privileged targets for Dark Matter annihilation signal searches. H.E.S.S. observations on the Sagittarius dwarf galaxy are presented. The modelling of the Dark Matter halo profile of Sagittarius dwarf is discussed. Constraints on the velocity-weighted cross section of Dark Matter particles are derived in the framework of Supersymmetric and Kaluza-Klein models. The future of H.E.S.S. will be briefly discussed. \\PACS{ {98.70.Rz}{Gamma-rays : observations}   \\and {95.35.+d}{Dark Matter} } % } % ",
        "introduction": "\\label{intro} H.E.S.S. (High Energy Stereoscopic System) is an array of four Imaging Atmospheric Cherenkov Telescopes (IACTs) located in the Khomas Highlands of Namibia at an altitude of 1800 m above sea level. Each telescope has an optical reflector of 107 m$^2$ composed of 382 round mirrors \\cite{bernlohr}. The Cherenkov light is emitted by charged particles from electromagnetic showers initiated by the interaction of the primary gamma-rays in the Earth's upper atmosphere. The light is collected by the reflector which focuses it on a camera made of 960 fast photomultiplier tubes (PMTs) with individual field of view of 0.16$^{\\circ}$ in diameter \\cite{vincent}. Each PMT is equipped with a Winston cone to maximize the light collection and limit the background light. The total field of view of the H.E.S.S. instrument is 5$^{\\circ}$ in diameter. The stereoscopic technique allows for an accurate reconstruction of the direction and energy of the primary gamma-rays as well as for an efficient rejection of the background induced by cosmic ray interactions. The energy threshold is about 100 GeV at zenith and the angular resolution is better than 0.1$^{\\circ}$ per gamma-ray. The point source sensitivity is $\\rm 2 \\times 10^{-13} cm^{-2}s^{-1}$ above 1 TeV for a 5$\\sigma$ detection in 25 hours \\cite{crabe}.\\\\ The H.E.S.S. instrument is designed to detect very high energy (VHE) gamma-rays in the 100 GeV - 100 TeV energy regime and investigate their possible origin. Scenarios beyond the Standard Model of particle physics predict plausible WIMP (weakly interacting massive particle) candidates to account for the Cold Dark Matter. Among these are the minimal supersymmetric extension of the Standard Model (MSSM) or universal extra dimension (UED) theories. With R-parity conservation, SUSY models predict the lightest SUSY particle (LSP) to be stable. In various SUSY breaking scenarios, the LSP is the lightest neutralino $\\rm \\tilde{\\chi}$. In Kaluza-Klein (KK) models with KK-parity conservation, the lightest KK particle (LKP) is stable \\cite{KK}, the most promising being the first KK mode of the hypercharge gauge boson, $\\rm \\tilde{B}^{(1)}$. The annihilation of WIMPs in galactic halos may produce gamma-ray signals detectable with Cherenkov telescopes (for a review, see \\cite{bertone}). Generally, their annihilations will lead to a continuum of gamma-rays with energy up to the WIMP mass resulting from the hadronization and decay of the cascading products, mainly from $\\pi^0$'s generated in the quark jets. The gamma-ray flux from DM particle annihilations of mass m$_{DM}$ in a spherical halo can be expressed as: \\begin{equation} \\frac{d\\Phi(\\Delta\\Omega,E_{\\gamma})}{dE_{\\gamma}} = \\frac{1}{4\\pi}\\,\\frac{\\langle \\sigma v \\rangle}{m^2_{DM}}\\frac{dN_{\\gamma}}{dE_{\\gamma}}\\,\\times\\,\\bar{J}\\Delta\\Omega \\end{equation} which is a product of a particle physics term containing the velocity-weighted cross section $\\langle \\sigma v \\rangle$ and the differential gamma-ray spectrum $dN_{\\gamma}/dE_{\\gamma}$, and an astrophysics term $\\bar{J}$ given by: \\begin{equation} \\bar{J} = \\frac{1}{\\Delta\\Omega}\\int_{\\Delta\\Omega}\\int_{l.o.s}\\rho^2 ds \\end{equation} This term corresponds to the line-of-sight-integrated squared density of the DM distribution which is averaged over the solid angle integration region spanned by the H.E.S.S. point spread function (PSF).\\\\ Plausible astrophysical targets have been proposed to search for DM, from local galactic objects to extragalactic objects such as galaxy clusters. This paper reports on three recent results on targets relevant to indirect dark matter search: the Galactic Center, the center of the VIRGO cluster (M87) and the Sagittarius dwarf spheroidal galaxy. ",
        "conclusions": "\\label{conclusion} The H.E.S.S. collaboration has studied potential targets for DM annihilations. The TeV $\\gamma$-ray energy spectrum measured by H.E.S.S. in the Galactic Center region is unlikely to be interpreted in terms of WIMP annihilations. Constraints on the velocity-weighted annihilation cross-section have been derived in the case of a NFW DM halo profile. Neither pMSSM nor KK models can be ruled out. The detection of a temporal variability of the VHE signal from the M87 nucleus in the VIRGO galaxy cluster excludes the bulk of the gamma-ray signal to come from DM annihilations. H.E.S.S. observed the Sagittarius dwarf spheroidal galaxy and no significant gamma-ray excess has been detected. Two DM halo modellings of Sgr have been investigated to derive contraints on the velocity-weighted annihilation cross-section. Some models can be ruled out in the case of a cored profile.\\\\ Dark matter searches will continue and searches with H.E.S.S. 2 will start in 2009.  The phase 2 will consist of a new large 28 m diameter telescope located at the center of the existing array. With the availability of the large central telescope H.E.S.S. 2, the analysis energy threshold will be lowered down to less than 80 GeV and will allow to explore more supersymmetric models."
    },
    "0710/0710.2941_arXiv.txt": {
        "abstract": "Although the occurrence of steady supercritical disk accretion onto a black hole has been speculated about since the 1970s, it has not been accurately verified so far. For the first time, we previously demonstrated it through two-dimensional, long-term radiation-hydrodynamic simulations. To clarify why this accretion is possible, we quantitatively investigate the dynamics of a simulated supercritical accretion flow with a mass accretion rate of $\\sim 10^2 L_{\\rm E}/c^2$ (with $L_{\\rm E}$ and $c$ being, respectively, the Eddington luminosity and the speed of light). We confirm two important mechanisms underlying supercritical disk accretion flow, as previously claimed, one of which is the radiation anisotropy arising from the anisotropic density distribution of very optically thick material. We qualitatively show that despite a very large radiation energy density, $E_0\\gsim 10^2L_{\\rm E}/4\\pi r^2 c$ (with $r$ being the distance from the black hole), the radiative flux $F_0\\sim cE_0/\\tau$ could be small due to a large optical depth, typically $\\tau\\sim 10^3$, in the disk. Another mechanism is photon trapping, quantified by ${\\bm v}E_0$, where ${\\bm v}$ is the flow velocity. With a large $|{\\bm v}|$ and $E_0$, this term significantly reduces the radiative flux and even makes it negative (inward) at $r < 70 r_S$,  where $r_S$ is the Schwarzschild radius. Due to the combination of these effects, the radiative force in the direction along the disk plane is largely attenuated so that the gravitational force barely exceeds the sum of the radiative force and the centrifugal force. As a result, matter can slowly fall onto the central black hole mainly along the disk plane with velocity much less than the free-fall velocity, even though the disk luminosity exceeds the Eddington luminosity. Along the disk rotation axis, in contrast, the strong radiative force drives strong gas outflows. ",
        "introduction": "Recently, very bright objects which may be undergoing supercritical (or super-Eddington) accretion flows have successively been found. Good examples are ultraluminous X-ray sources (ULXs; \\citeauthor{WMM01} \\citeyear{WMM01}; \\citeauthor{Ebisawa03} \\citeyear{Ebisawa03}; \\citeauthor{Okajima06} \\citeyear{Okajima06}). These are pointlike off-center X-ray sources whose X-ray luminosity significantly exceeds the Eddington luminosity of a neutron star \\citep{Fabbiano89}. Because of substantial variations, it is reasonable to assume that the ULXs are single compact objects powered by accretion flows \\citep{Makishima00}. If so, there are two possibilities to account for large luminosities exceeding the Eddington luminosity for a mass of 100 $M_\\odot$: sub-critical accretion onto an intermediate-mass black hole (IMBH) and supercritical accretion onto a stellar-mass black hole. We support the latter possibility, since through the fitting to several {\\it XMM-Newton} EPIC data of ULXs, which have been claimed as good IMBH candidates, we have found evidence of supercritical flows \\citep{Vierdayanti06}. Another interesting group is narrow-line Seyfert 1 galaxies \\citep[see][for a review]{Boller04}. Because of their relatively small black hole masses, they have in general large Eddington ratios ($L/L_{\\rm E}$ with $L$ and $L_{\\rm E}$ being the luminosity and the Eddington luminosity, respectively), and some of them seem to fall in the slim-disk regimes (\\citeauthor{Mineshige00} \\citeyear{Mineshige00}; \\citeauthor{Kawaguchi03} \\citeyear{Kawaguchi03}; see also \\citeauthor{Wang99} \\citeyear{Wang99})  Despite growing evidence indicating the existence of supercritical accretion flows in the universe, theoretical understanding is far from being complete. It is well known that any spherically accreting object, irrespective of the nature of the central source, cannot emit above the Eddington luminosity, since otherwise significant radiative force will prevent accretion of the gas. If we examine detailed radiation-matter interactions in the interior, however, we notice that the situation is not so simple.   Actually, radiation produced at the very center cannot immediately reach the surface (i.e., photosphere), since photons generated at the center should suffer numerous scatterings with accreting material and thus take a long time to reach the surface. If the matter continuously falls, and if the mass accretion timescale is shorter than the mean travel time for photons to reach the surface (the diffusion timescale), photons at the core may  not be able to go out. This is the so-called photon-trapping effect \\citep[e.g.,][]{Begelman78,HC91}. Here we define the trapping radius, inside which radiation-matter interaction is so frequent that photons are trapped within the accretion flow. Inside this trapping radius, therefore, radiative flux can be negative (i.e., inward). Thus, the apparent luminosity is reduced as compared with the case without photon trapping. Nevertheless, the concept of the Eddington luminosity is still valid, since far outside the trapping radius radiative flux should be outward and its absolute value should be less than $L_{\\rm E}/4\\pi r^2$, where $r$ is the distance from the central black hole, in the quasi-steady state. Radiation-hydrodynamic (RHD) simulations of spherically symmetric supercritical accretion flows have been performed by \\citet{BK83}.  The situation may differ in the case of disk accretion, since the radiation field is not isotropic due to inhomogeneous matter distribution. That is, matter can fall predominantly along the disk plane, whereas radiation can go out along its rotating axis, where matter density is low. In other words, the main directions of the inward matter flow and the outward radiative flux are not parallel to each other in disk accretion, leading to a situation in which the radiative force does not completely counteract the gravitational force. There is room for the possibility of a supercritical accretion flow with super-Eddington luminosity.  Based on such an argument, many researchers have speculated about the occurrence of supercritical flows in disk accretion systems. \\citet{SS73} discussed the possibility of supercritical disk accretion based on a one-dimensional steady model \\citep[see also][]{MRT76}. They mentioned that the mass accretion rate would not be steady but oscillate if the mass accretion rate exceeded the critical rate; otherwise, a part of the accreting matter might be ejected from the disk as a disk wind. Radiatively driven outflows from supercritical disks were investigated by \\citet{Meier79}, \\citet{JR79}, and \\citet{Icke80}.  Despite a long history in the study of supercritical disk accretion flows, the occurrence of steady supercritical disk accretion has not yet been accurately verified. Similar simulations have been performed since the 1980s, but all of them calculated only the initial transient phase. Their conclusions then were not general, since the back-reactions (i.e., enhanced radiation pressure), which may inhibit steady flow, were not accurately evaluated. Although the radiative force predominantly has an effect in the vertical direction, it should also have an effect on the material within the disk. Hence, in the direction parallel to the disk plane, the situation may be similar to or more severe than the case of spherical accretion. This is because the radiative force, together with the centrifugal force, may possibly overcome the gravitational force. To study the possibility of supercritical disk accretion flows, precise and quantitative research treating both supercritical disks and outflows, and which also takes into consideration multi-dimensional effects, is needed.   \\citeauthor{O05} (\\citeyear[][hereafter Paper I]{O05}) have confirmed the occurrence of quasi-steady supercritical disk accretion onto a black hole by two-dimensional RHD simulations. The motivation of the present study is to investigate the physical mechanisms which make supercritical disk accretion possible. For this purpose we examine quantitatively the flow motion and force fields via the radiative flux, rotation, and gravity of a supercritical disk accretion flow onto a central black hole, based on the two-dimensional RHD simulation data from Paper I. Through detailed inspection of the results, it will be possible to clarify the physics behind supercritical disk accretion flows. In \\S 2 we plot the spatial distributions of several key quantities which control flow dynamics. A discussion is given in \\S 3. ",
        "conclusions": "\\subsection{Structure of the Supercritical Disk Accretion Flow} We summarize our simulation results in a schematic picture (see Fig. \\ref{pic}). \\begin{figure*} \\epsscale{0.87} \\plotone{Fig07_astroph.eps} \\caption{Schematic picture of the supercritical disk accretion flow around a black hole. The gas motion is shown on the right: the high-velocity outflow ({\\it solid arrows}) and the slow accretion flow ({\\it dashed arrows}). The left side indicates the radiative fluxes in the comoving frame ({\\it black arrows}) and the rest frame ({\\it white arrows}). \\label{pic} } \\end{figure*} In this figure the radiative flux in the comoving frame, $F_0^r$, is positive (outward) except in the vicinity of the black hole. However, the radiative flux in the rest frame, $F^r (\\sim F_0^r+4v_r E_0/3)$, is negative (inward) via photon trapping, $v_r E_0<0$, in the trapping region. In the vicinity of the black hole, both $F_0^r$ and $F^r$ are negative. Matter slowly accretes inside the disk, since the sum of the radiative and centrifugal forces is nearly balanced with the gravitational force. Here we stress again that the radiative force counteracts the gravitational force in spite of the trapping region because $F_0^r>0$. Radiatively driven high-velocity outflows appear above and below the disk. In the very vicinity of the black hole, the gas is accelerated inward by the radiative force and the gravitational force.  As far as the physical quantities around the equatorial plane are concerned, the simulated profiles of the density, temperature, and radial as well as rotational velocities roughly agree with the prediction of the slim-disk model \\citep{Abramowicz88}. Such features have already been shown in Figure 11 in Paper I. However, only about $10\\%$ of the injected mass can accrete onto the black hole, and an almost equal amount of matter is ejected as high-velocity outflows. The mass accretion rate is not constant in the radial direction and decreases near the black hole (see Fig. 6 in Paper I). Thus, we conclude that the simulated flows do not perfectly agree with the slim-disk model with regard to the whole structure of the flow.    \\citet{HD07} have investigated local maximum values of the accretion rate in the supercritical disk accretion flows. In their paper they focused only on the force balance in the vertical and radial directions around the equatorial plane. Multi-dimensional effects were not taken into consideration. They revealed that the vertical radiative force limits the maximum accretion rate at the inner disk region, leading to a decrease of the accretion rate with a decrease of the radius. Their results imply that the disk loses mass via the radiatively driven outflows and the mass accretion rate decreases at the inner region. Such tendencies agree with our results. As shown in Figure \\ref{vr}, our simulations show that radiatively driven outflows form above and below the disk. The mass accretion rate decreases with a decrease of the radius (see Fig. 6 in Paper I).   \\subsection{Dependency of the Mass Accretion Rate} In the present study, focusing on numerical simulations in which the mass input rate at the outer boundary is set to be $\\dot{M}_{\\rm input}=10^3L_{\\rm E}/c^2$, we show that the radiative force is attenuated in the disk region via large optical thickness, which makes supercritical disk accretion possible. Such dilution of the radiative force would effectively operate even if the mass input rate (mass accretion rate) varied. In fact, simulations with mass input rates of $3\\times 10^2L_{\\rm E}/c^2$ and $3\\times 10^3L_{\\rm E}/c^2$ show that $E_0/\\rho$ is almost independent of the mass input rate, although the radiation energy density goes up as the mass input rate increases. That is, the dynamics is not sensitive to the precise value of $\\dot{M}_{\\rm input}$, as long as it largely exceeds the critical value.  So far, we have studied steady accretion flows. Although a highly supercritical disk ($\\dot{M}_{\\rm input}>3\\times 10^2L_{\\rm E}/c^2$) is quasi-steady, it has been revealed that a moderately supercritical disk [$\\dot{M}_{\\rm input}=(10-10^2)L_{\\rm E}/c^2$] is unstable and exhibits limit-cycle behavior (\\citeauthor{O06} \\citeyear{O06}, \\citeyear{O07}; see also \\citeauthor{SH75} \\citeyear{SH75}; \\citeauthor{SS76} \\citeyear{SS76}). The luminosity goes up and down around the Eddington luminosity.  \\citet{O07} has reported that the time-averaged mass, momentum, and kinetic energy output rates via the outflow, the mass accretion rate, and the disk luminosity increase as the mass input rate increases, $\\propto \\dot{M}_{\\rm input}^{0.7} -\\dot{M}_{\\rm input}^{1.0}$ for $\\alpha=0.5$ and $\\propto \\dot{M}_{\\rm input}^{0.4} -\\dot{M}_{\\rm input}^{0.6}$ for $\\alpha=0.1$.      \\subsection{Future Work} As we have already mentioned in \\S 3.1, the sum of the accreting matter and the matter ejected as high-velocity outflows is $20\\%$ of the injected mass, and $80\\%$ of the injected matter is ejected from the computational domain as low-velocity outflows, whose velocities do not exceed the escape velocity at the outer boundary. Since such outflowing matter tends to be accelerated by the radiative force even at the outside of the computational domain, it would be blown away from the system. However, a part of the outflowing matter might return to the vicinity of the black hole through the disk region, since the radiative force does not exceed the gravity near the equatorial plane. Numerical simulations with larger computational domains would make this point clear.  Whereas the resulting mass accretion rate onto the black hole is around $10^2 L_{\\rm E}/c^2$ in the present simulations, \\citet{HD07} have indicated that the mass accretion rate can increase up to $10^4 L_{\\rm E}/c^2$. They have reported that the vertical force balance breaks down via a strong radiative force if the mass accretion rate exceeds this limit. However, even in such a case, the matter might accrete onto the black hole, although the strong radiative force would produce powerful outflows. To investigate the maximum value of the accretion rate is an outstanding issue. We should perform numerical simulations with larger computational domains, since the trapping region is expected to expand with the increase of the mass accretion rate.  We reveal in the present paper that photons generated deep inside the disk are effectively trapped in the flow, leading to supercritical disk accretion. Although the magnetic fields are not solved in our simulations, magnetic buoyancy might play an important role in the transportation of matter, as well as photons, toward the disk surface (\\citeauthor{Parker75} \\citeyear{Parker75}; \\citeauthor{SR84} \\citeyear{SR84}; \\citeauthor{SC89} \\citeyear{SC89}). Magnetic buoyancy might lead to photon generation near the disk surface if the magnetic fields rise quickly without dissipation deep inside the disk. Thus, magnetic buoyancy would dilute the photon-trapping effect. A photon bubble instability, which is induced in the magnetized, radiation-pressure-dominated region, might also suppress the photon trapping \\citep{Begelman02,Turner05}. In these cases the enhanced radiative force would more effectively accelerate the matter around the disk surface, working to decrease the mass accretion rate. However, the magnetic fields might prevent such acceleration if they strongly tie the matter near the disk surface with the disk matter. In the disk region the matter might easily accrete toward the black hole, since the radiation energy density decreases. Global radiation-magnetohydrodynamic (RMHD) simulations would make these points clear. Local RMHD simulations of accretion flows have been performed by \\citet{Turner03} and \\citet{HKS06}. In addition, it is thought that disk viscosity has magnetic origins (\\citeauthor{HBS01} \\citeyear{HBS01}; \\citeauthor{MMM01} \\citeyear{MMM01}; for a review see \\citeauthor{Balbus03} \\citeyear{Balbus03}). Hence, we should stress again that RMHD simulations are very important to more realistically investigate viscous accretion flows, although an $\\alpha$-viscosity model is employed in the present study."
    },
    "0710/0710.2807_arXiv.txt": {
        "abstract": "We have identified two moderately hot ($\\sim$18000--22000\\,K) white dwarfs, SDSS\\,J1228+1040 and SDSS\\,J1043+0855, which exhibit double-peaked emission lines in the CaII\\,$\\lambda\\lambda$\\,8600 triplet. These line profiles are unambiguous signatures of gaseous discs with outer radii of $\\sim1R_\\odot$ orbiting the two white dwarfs. Both stars accrete from the circumstellar material, resulting in large photospheric Mg abundances. The absence of hydrogen emission from the discs, and helium absorption in the white dwarf photospheres demonstrates that the circumstellar material is depleted in volatile elements, and the most likely origin of these gaseous rings are tidally disrupted rocky asteroids. The relatively high mass of SDSS\\,J1228+1040 implies that planetary systems  can not only form around $4-5\\,M_\\odot$ stars, but may also survive their post main-sequence evolution. ",
        "introduction": "While more than 250 extra-solar planets orbiting main-sequence stars have been discovered, the destiny of planetary systems in the late stages of the evolution of their host stars is very uncertain, and so far no planet has been found around a white dwarf. Infrared excess detected around a number of white dwarfs has been interpreted as the signature of dust discs \\citep[e.g.][]{zuckerman+becklin87-1, becklinetal05-1, kilicetal06-1, vonhippeletal07-1}. The photospheres of these white dwarfs are rich in metals \\citep{zuckermanetal07-1}, indicating ongoing accretion from the circumstellar material.  The likely origin of these debris discs are tidally disrupted asteroids \\citep{jura03-1}, and hence they represent a close association with the planetary systems that the white dwarf progenitor stars may have had. However, while the infrared excess detected around these white dwarfs can be explained in terms of a dusty debris disc, the observations actually do not provide any strong constraint on the geometry of the source of the infrared light \\citep[e.g.][]{reachetal05-1}.  We summarise here our recent discovery of two white dwarfs in the SDSS spectroscopic data base which exhibit double-peaked emission lines of Ca\\,II$\\lambda\\lambda$8600, unambiguously confirming a circumstellar disc-like structure \\citep{gaensickeetal06-3, gaensickeetal07-1} ",
        "conclusions": "It has been suggested that planetary systems may survive the post-main sequence evolution of their host stars \\citep{burleighetal02-1, villaver+livio07-1}, however, no planet has yet been discovered around a white dwarf. The detection of debris discs from rocky asteroids around white dwarfs, such as SDSS\\,J1228+1040 and SDSS\\,J1043+0855 and the cooler white dwarfs with dusty debris discs lends strong support to the survival hypothesis. It appears also entirely possible that these white dwarfs may still have planetesimal objects or planets. SDSS\\,J1228+1040 is particular interesting, as its relatively high mass implies that its progenitor must have had a mass of $\\sim4\\,M_\\odot$ \\citep{dobbieetal06-1}, suggesting that also short-lived massive stars may be host to planetary discs. This is in accordance with the detection of a relatively massive debris disc around the young A2e star MWC\\,480 \\citep{manningsetal97-1}."
    },
    "0710/0710.2813_arXiv.txt": {
        "abstract": "{ We report on a sensitive search for mid-infrared molecular hydrogen emission from protoplanetary disks. We observed the Herbig Ae/Be stars UX Ori, HD 34282, HD 100453, HD 101412, HD 104237 and HD 142666, and the T Tauri star HD 319139, and searched for H$\\,_2~0-0~S(2)~(J=4-2)$ emission at 12.278 micron and H$\\,_2~0-0~S(1)~(J=3-1)$ emission at 17.035 micron with VISIR, ESO-VLT's high-resolution mid-infrared spectrograph. None of the sources present evidence for molecular hydrogen emission at the wavelengths observed. Stringent 3$\\sigma$ upper limits to the integrated line fluxes and the mass of optically thin warm gas ($T=$ 150, 300 and 1000 K) in the disks are derived. The disks contain less than a few tenths of Jupiter mass of optically thin H$_2$ gas at 150 K, and less than a few Earth masses of optically thin H$_2$ gas at 300 K and higher temperatures. We compare our results to a Chiang and Goldreich (1997, CG97) two-layer disk model of masses 0.02 M$_{\\odot}$ and 0.11 M$_{\\odot}$. The upper limits to the disk's optically thin warm gas mass are smaller than the amount of warm gas in the interior layer of the disk, but they are much larger than the amount of molecular gas expected to be in the surface layer. If the two-layer approximation to the structure of the disk is correct, our non-detections are consistent with the low flux levels expected from the small amount of H$_2$ gas in the surface layer. We present a calculation of the expected thermal H$_2$ emission from optically thick disks, assuming a CG97 disk structure, a gas-to-dust ratio of 100 and T$_{\\rm gas}$ = T$_{\\rm dust}$. We show that the expected H$_2$ thermal emission fluxes from typical disks around Herbig Ae/Be stars are of the order of 10$^{-16}$ to 10$^{-17}$ erg s$^{-1}$ cm$^{-2}$ for a distance of 140 pc. This is much lower than the detection limits of our observations (5 $\\times$ $10^{-15}$ erg s$^{-1}$ cm$^{-2}$). H$_2$ emission levels are very sensitive to departures from the thermal coupling between the molecular gas and dust in the surface layer. Additional sources of heating of gas in the disk's surface layer could have a major impact on the expected H$_2$ disk emission. Our results suggest that in the observed sources the molecular gas and dust in the surface layer have not significantly departed from thermal coupling (T$_{\\rm gas}$/T$_{\\rm dust}<$ 2) and that the gas-to-dust ratio in the surface layer is very likely lower than 1000. ",
        "introduction": "Circumstellar disks surrounding low- and intermediate-mass stars in their pre-main sequence phase are the locations where planets presumably form. Such protoplanetary disks are composed of gas and dust. Their mass is initially dominated by gas (99\\%), specifically by molecular hydrogen (H$_2$), which is the most abundant gas species. The dust constitutes only  a minor fraction of the total disk mass, however, it is the main source of opacity. Consequently, most of what we know observationally about protoplanetary disks has been inferred from studies of dust emission and scattering (for recent reviews see Henning et al. 2006; Natta et al. 2007; Dullemond et al. 2007; Watson et al. 2007). In order to understand the structure and evolution of protoplanetary disks, it is necessary to study their gaseous content independently from the dust. For example, a basic physical quantity such as the disk mass is conventionally deduced from dust continuum emission at millimeter wavelengths assuming an interstellar gas-to-dust ratio of 100 (e.g., Beckwith et al. 1990; Henning et al. 1994). If dust is physically processed in the disk, as should be the case in order to form planets, the gas-to-dust ratio must change with time. The disk dissipation time scale - another fundamental quantity required to disentangle proposed giant planet formation scenarios (Pollack et al. 1996; Boss et al. 1998) - is deduced from observations of thermal infrared excess emission produced by dust grains (Haisch et al. 2001, 2005). Although recent studies (e.g., Sicilia-Aguilar et al. 2006) suggest a parallel evolution of the dusty and gaseous components, it still remains to be demonstrated that the gaseous disks disappear over the same time scale as the infrared excess.  A variety of spectral diagnostics of the gas disk have been observed from the UV to the millimeter (see reviews by Najita et al. 2007; Dutrey et al. 2007). However, the only diagnostic that is potentially able to probe the warm gas in the regions where giant planets are thought to form is the mid-infrared (mid-IR) emission lines of H$_2$. UV and near-infrared diagnostics only probe the innermost regions of the disk (R $<$ few AU), and mm and sub-mm diagnostics are limited to probe the cold outermost regions of the disk (R $>$10 AU).  \\begin{table*} \\caption{Summary of the observations.}             % \\label{table:observations}      % \\centering                          % \\begin{tabular}{@{}l c l c c c c l c c c c c c c @{}}        % \\hline\\hline                 % Star    & $\\lambda$ &  Date            &  U.T.    & $t_{exp}$ & Airmass $^a$ & V$_{\\oplus\\,{\\rm rad}}\\,^b$ & \\multicolumn{2}{c}{Calibrator $^c$} & $t_{exp}$ & \\multicolumn{2}{c}{Airmass} \\\\ & [$\\mu$m]  &                  &  [hh:mm] &   [s]        &              &  [km s$^{-1}$] & \\multicolumn{2}{c}{ }          & [s]       & \\multicolumn{2}{c}{       } \\\\  \\hline                        % UX Ori    & 12.278 & 11 January 2006  & 02:28 & 3600 & 1.0 - 1.2  & 2.06 & HD 36167  & (P)(F)  & 1000 & 1.1(P)  & 1.2(F) \\\\ & 17.035  & 4 January 2007 & 01:54 & 3600 & 1.0 - 1.0  & 4.50 & HD 25025 & (P)    & 600 &   1.0 (P)  & ...\\\\ HD 34282  & 12.278 & 10 January 2006 & 04:01 & 3600 & 1.0 - 1.5  & 2.17 & HD 36167  & (P)(F)  & 1000 & 1.1(P)  & 1.7(F) \\\\ HD 100453 & 12.278  & 22 March 2006 & 07:27 & 3600 & 1.2 - 1.5  & 24.45 & HD 89388  & (P)(F)  & 600   & 1.5 (P) & 2.0 (F) \\\\ & 17.035  & 27 March 2006 & 06:26 & 3600 & 1.3 - 1.7  & 22.96 & HD 89388  & (P)     & 1000 & 1.5 (P) & ... \\\\ HD 101412 & 12.278  & 30 March 2006 & 04:40 & 3600 & 1.2 - 1.4  & 23.82 & HD 91056  & (P)(F)  & 600  & 1.3 (P) & 1.7 (F) \\\\ & 17.035  & 30 March 2006 & 01:27 & 3600 & 1.3 - 1.2  & 24.02 & HD 91056  & (F)     & 1000 & ...     & 1.3 (F) \\\\ HD 104237 & 12.278  & 10 February 2006 & 05:43 & 3600 & 1.7 - 1.7  & 26.09 & HD 92305  & (P)(F)  & 600  & 1.7 (P) & 1.8 (F) \\\\ & 17.035  & 12 February 2006 & 06:36 & 3600 & 1.6 - 1.7  & 26.15 & HD 92305  & (F)     & 1000 & ...     & 1.8 (F) \\\\ HD 142666 & 17.035  & 28 February 2006 & 06:28 & 3600 & 1.1 - 1.0  & 27.06 & HD 169916 & (P)    & 1000 & 1.1 (P)  & ...\\\\ HD 319139 & 17.035  & 30 March 2006 & 08:30 & 3600 & 1.1 - 1.0  & 22.47 & HD 169916 & (P)    & 1000 & 1.2 (P)  & ...\\\\  \\hline \\end{tabular} \\flushleft $^a$ The airmass interval is given from the beginning to the end of the observations. $^b$ V$_{\\oplus\\,{\\rm rad}}$ is the expected velocity shift of the spectra due to the reflex motion of the Earth around the Sun and the radial velocity of the star. $^c$ The standard stars were observed prior (P) and/or immediately following (F) the science observations. \\end{table*}  Molecular hydrogen is by far the most abundant molecular species in protoplanetary disks. Unfortunately, H$_2$ is one of the most challenging molecules to detect. Since H$_2$ is a homonuclear molecule, it lacks a permanent dipole moment and its transitions are thus electric quadrupole in nature. The small Einstein coefficients, characteristic of the quadrupole transitions, imply that H$_2$ emission features are very weak. In addition, in the case of protoplanetary disks, the H$_2$ lines are not sensitive to the warm gas in the optically thick regions where the dust and gas are at equal temperature. Practical observational challenges also have to be faced. The mid-IR H$_2$ emission from the disk needs to be detected on the top of a strong mid-IR continuum. From the ground, the mid-infrared windows are strongly affected by sky and instrument background emission, and the H$_2$ transitions at 12 and 17 $\\mu$m lie close to atmospheric absorption lines highly dependent on atmospheric conditions. The advent of high spectral resolution spectrographs mounted on larger aperture telescopes, allows for the first time the study of H$_2$ emission from the ground, but the search is still limited to bright targets. From space, the problems of atmosphere absorption are alleviated and the $J = 2-0$ feature at 28 $\\mu$m is visible. However, the beam sizes are large and the spectral resolution of space mid-infrared spectrographs are usually low when compared to ground-based facilities, therefore, they are not very appropriate for small line-to-continuum ratios.  H$_2$ emission from protoplanetary disks in the mid-IR has been reported from ISO observations (Thi et al. 2001). However, subsequent ground-based efforts (Richter et al. 2002; Sheret et al. 2003; Sako et al. 2005) did not confirm the ISO detections. H$_2$ emission in the mid-IR has been searched towards debris disks using Spitzer (Hollenbach et al. 2005, Pascucci et al. 2006, Chen et al. 2006) with no detection reported. Most recently, Bitner et al. (2007) and Martin-Za\\\"idi et al. (2007) reported the detection of mid-IR H$_2$ emission in two Herbig Ae/Be stars (AB Aur and HD 97048) from the ground, and Lahuis et al. (2007) reported the detection of mid-IR H$_2$ emission in 6 T Tauri stars with Spitzer.  Here, we report on a sensitive search for molecular hydrogen emission from protoplanetary disks. We observed six southern nearby (d$<$400 pc) Herbig Ae/Be   stars and one T Tauri star, employing the Very Large Telescope (VLT) imager and spectrometer for the  mid-infrared (VISIR; Lagage et al. 2004)\\footnote{http://www.eso.org/instruments/visir}, ESO's VLT mid-infrared high-resolution spectrograph, and searched for H$_2~0-0~S(1)~(J=3-1)$~emission at 17.035~$\\mu$m and H$_2~0-0~S(2)~(J=4-2)$ emission at 12.278~$\\mu$m. The paper is organized as follows: in Sect. 2 we present the sample studied, and the details of how the observations were performed. In Sect. 3 we discuss the data reduction. In Sect. 4 we deduce upper limits to the H$_2$ fluxes, and using the optically thin approximation, we derive upper limits to the mass of warm (150 - 1000 K) gas in the disks. In Sect. 5 we discuss our results in the context of the Chiang and Goldreich (1997) two-layer disk model. Finally, we summarize our results and conclusions in Sect. 6. \\begin{figure*} \\centering \\includegraphics[angle=0,width=0.95\\textwidth]{7846fig1.eps} \\caption{Spectra obtained for the H$_2$ 0--0 S(2) (J=4--2) line at 12.278 $\\mu$m. The left panel shows a zoom to the -100 to 100 km s$^{-1}$ interval of the atmospheric corrected spectra. A Gaussian of {\\it FWHM} = 15 km s$^{-1}$ and integrated line flux equal to the line-flux upper limits obtained is overplotted at the expected velocity shifted location (vertical dashed line, see Table 2). The central panel shows the full corrected spectra. Dotted lines show spectral regions strongly affected by telluric or standard star absorption features. The right panel shows the continuum normalized spectra of the standard star and the target before telluric correction. The uppermost right panel displays the sky spectrum from a half-chop cycle. The spectra are not corrected for the radial velocity of the targets.       } \\end{figure*} \\begin{figure*} \\centering \\includegraphics[angle=0,width=0.95\\textwidth]{7846fig2.eps} \\caption{Spectra obtained for the H$_2$ 0--0 S(1) (J=3--1) line at 17.035 $\\mu$m. The left panel shows a zoom to the -100 to 100 km s$^{-1}$ interval of the atmospheric corrected spectra. A Gaussian of {\\it FWHM} = 15 km s$^{-1}$ and integrated line flux equal to the line-flux upper limits obtained is overplotted at the expected velocity shifted location (vertical dashed line, see Table 2). The central panel shows the full corrected spectra. Dotted lines show spectral regions strongly affected by telluric or standard star absorption features. The right panel shows the continuum normalized spectra of the standard star and the target before telluric correction. The uppermost right panel displays the sky spectrum from a half-chop cycle. The spectra are not corrected for the radial velocity of the targets. } \\end{figure*} ",
        "conclusions": "We observed a sample of nearby pre-main sequence stars with evidence for cold ($T < 50$ K) gas disk reservoirs and searched for emission of the warm gas with $T > 150$ K, which is expected to be present in the inner region of these disks. None of the targets show any evidence for H$_2$ emission at 17.035 $\\mu$m or at 12.278 $\\mu$m. From the 3$\\sigma$ upper limits of the H$_2$ line fluxes, we found stringent upper limits to the mass of optically thin warm H$_2$. The disks contain less than a few tenths of Jupiter mass of optically thin H$_2$ at 150 K, and less than a few Earth masses of optically thin H$_2$ at 300 K and higher temperatures.  Assuming that T$_{\\rm gas}$=T$_{\\rm dust}$ and a gas-to-dust ratio of 100, we compared our results to models of disks employing a Chiang and Goldreich (1997) optically thick two-layer disk model of masses 0.02 M$_{\\odot}$ and 0.11 M$_{\\odot}$. The upper limits to the disk optically thin warm gas mass are smaller than the warm gas mass in the interior layer of the disk, but they are much larger than the amount of molecular gas expected to be in the surface layer. The amount of mass in the surface layer is very small ($<$ 10$^{-2}$ M$_J$) and almost independent of the total disk mass. We calculated the expected H$_2$ S(1) and H$_2$ S(2) line fluxes emitted from a two-layer disk for the low-mass and the high-mass cases assuming a distance of 140 pc, and LTE thermal emission. The predicted line fluxes of the two-layer disk model are of the order of $\\sim$10$^{-16}-10^{-17}$ erg s$^{-1}$ cm$^{-2}$, much smaller than the detection limits of our observations (5 $\\times$ $10^{-15}$ erg s$^{-1}$ cm$^{-2}$).  {\\it If the two-layer approximation to the structure of the disk is correct, we are essentially ``blind\" to most of the warm H$_2$ in the disk because it is located in the optically thick interior layer of the disk. Our non-detections are explained because of the small flux levels expected from the little mass of H$_2$ present in the optically thin surface layer.} Naturally, the two-layer disk model is only an approximation of the real structure of a protoplanetary disk. The puffed-up inner rim, which could be an important contributor to the H$_2$ emission, is not included in our models, and in reality there will be a smooth transition zone between the disk hot surface layer and the cool disk interior, which again could contribute significantly to the H$_2$ emission. In addition, the surface layers of the disk are likely to have gas temperatures hotter than the dust temperatures or a gas-to-dust ratio higher than the canonical value of 100. Both effects could potentially increase the H$_2$ emission by significant amounts. We presented additional calculations in which the gas-to-dust ratio and the temperature of the molecular gas in the surface layer of the disk were increased. We showed that detectable S(1) and S(2) H$_2$ line flux levels can be achieved if T$_{\\rm gas\\,surf}$/T$_{\\rm dust\\,surf}>$ 2 and if the  gas-to-dust ratio in the surface layer is greater than 1000. H$_2$ emission levels are very sensitive to departures from the thermal coupling between the molecular gas and dust in the surface layer. Our results suggest that in the observed sources the molecular gas and the dust in the surface layer have not significantly departed from thermal coupling and that the gas-to-dust ratio in the surface layer is very likely lower than 1000. A definitive interpretation of our results awaits the development of future, more sophisticated models."
    },
    "0710/0710.2955_arXiv.txt": {
        "abstract": "We report on the discovery of three new pulsars in the first blind survey of the north Galactic plane (45$^{\\circ}$ < l  < 135$^{\\circ}$ ; |b| < 1$^{\\circ}$) with the Giant Meterwave Radio telescope (GMRT)  at an intermediate frequency of  610 MHz. The timing parameters, obtained in follow up observations with the Lovell Telescope at Jodrell Bank Observatory and the GMRT, are  presented. ",
        "introduction": "Most pulsar surveys have been carried out with single dish telescopes where there is a trade-off between  the collecting area and the beam-width, and consequently the rate of the survey. In a multi-element telescope such as the Giant Meterwave Radio Telescope (GMRT), a large number of smaller antennas can be combined to provide a high sensitivity and yet retain a relatively large beam-width. In this paper, we report on the discovery of three new pulsars in the first blind survey of the north Galactic plane (45$^\\circ$ < l  < 135$^\\circ$ ; |b| < 1$^\\circ$) with the GMRT at an intermediate frequency of  610 MHz, which represents the best trade-off between the increased flux density at low frequency for pulsars, interstellar scattering and dispersion and beam-width. The GMRT's multi-element nature was also exploited to determine the positions of the pulsars to an accuracy of 5 arcminutes and this technique is also described. ",
        "conclusions": ""
    },
    "0710/0710.0999_arXiv.txt": {
        "abstract": "We investigate the fate of particle production in an expanding universe dominated by a perfect fluid with equation of state $p = \\alpha\\rho$. The rate of particle production, using the Bogolioubov coefficients, are determined exactly for any value of $\\alpha$ in the case of a flat universe. When the strong energy condition is satisfied, the rate of particle production decreases as time goes on, in agreement to the fact that the four-dimensional curvature decreases with the expansion; the opposite occurs when the strong energy condition is violated. In the phantomic case, the rate of particle production diverges in a finite time. This may lead to a backreaction effect, leading to the avoidance of the big rip singularity, specially if $- 1 > \\alpha > - \\frac{5}{3}$. ",
        "introduction": "The phenomenon of particle production in an expanding universe is due, essentially, to the fact that in curved space-time the vacuum is not unique \\cite{birrel,jacobson}. As the universe evolves, and the curvature changes, the vacuum state also changes: the initial vacuum state, representing the state with no particle, becomes later a multiparticle state. The particle production is directly connected with the curvature of the universe. If the universe is spatially flat, the cosmic evolution leads asymptotically to a Minkowski space-time, where the phenomenon of particle production does not occur anymore. However, this is true only if the strong energy condition is verified: the energy density $\\rho$ and the pressure $p$ must satisfy the relation $\\rho + 3p > 0$. If the strong energy condition is violated the particle production should be zero initially, increasing as the universe evolves, in connection with the increase of the four-dimensional curvature with time when $\\rho + 3p < 0$. \\par In order to compute the particle production in a given cosmological model, it is necessary to fix an initial vacuum state. If the universe is always decelerating, there is no natural and unique choice for this, since all modes are initially outside the Hubble radius feeling strongly the curvature of the space-time. However, a natural choice for the initial vacuum state can be done for the case of an accelerating universe, since all physical modes are initially well inside the Hubble radius, and they behave as in a Minkowski space-time. This is one of the reasons why the primordial inflationary scenario is so successful theoretically: it encodes naturally a mechanism for quantum fluctuations, which will be later the seeds for the large scale structure of the universe \\cite{jerome}. \\par Observations indicate that we live today in an inflationary phase, since the universe is accelerating \\cite{spergel,uzan}. Hence, the mechanism of particle production may play again a very important role today. Suppose we define a quantum state for a given field today. The question we want to address concerns the rate of particle production. In particular, we would like to know if such particle production is so important that it can alters, by back-reaction, the later evolution of the universe. The back-reaction due to quantum effects has already studied in the case of the cosmological constant \\cite{ford2}. The back-reaction may be particularly relevant for the case the universe is dominated by a phantom fluid since, classically, in such a situation the universe will inevitably evolve towards a new singular state, called the big rip. We remember that the observational data favors somehow a phantom scenario today \\cite{caldwell}. A phantom fluid has many special features. Local configurations of a phantom fluid leads to regular black holes \\cite{kirill}. In cosmology, a universe dominated by a phantom fluid implies a future singularity that will occur in a future finite proper time \\cite{frampton}. \\par We will solve the rate of the particle production when the universe is flat and dominated by an equation of state $p = \\alpha\\rho$, for an arbitrary value of $\\alpha$. A simplified scenario, not considering the different phase transitions that occur in the real universe, will be adopted. The calculations will be performed for a massless scalar field, for which the corresponding Klein-Gordon equation is solved. Through the quantization of the field, the rate of particle may be determined using the technique of Bogolioubov's coefficients. The general solution for this problem is expressed in terms of Bessel functions. The rate of particle production is determined exactly for any value of $\\alpha$. The connection of this rate with the strong energy condition is established. It is shown that for a phantom fluid the rate of particle production diverges as the big rip is approached. This may lead to a back-reaction effect changing the effective equation of state of the universe, implying perhaps the avoidance of the big rip itself: the particle produced does not necessarily obeys the phantom equation of state and may become the dominant component of the universe if $ - 1 > \\alpha > - \\frac{5}{3}$. Hence, the back-reaction of the particle production may be an ineluctable mechanism to the avoidance of the big rip if the pressure is not excessively negative. The limiting case, $\\alpha = - \\frac{5}{3}$, has already been found before in a complete different context, that is, in the evolution of classical scalar perturbations in phantom cosmological models \\cite{finelli,fabris}. \\par This paper is organized as follows. In the next section we review the formalism of particle creation in a curved space-time. In section $3$ we apply the formalism to a perfect fluid cosmological model, and we determine the rate of particle creation for any value of $\\alpha$. In section $4$ we present our conclusions. Even if many steps of the formalism used and of the calculation performed are quite well known, we will present them with some details in order to trace the main features of the final results. ",
        "conclusions": "We have evaluate the rate of creation of massless scalar particle in a universe dominated by a perfect fluid whose equation of state is given by $p = \\alpha\\rho$. An analytic expression has been found in terms of Hankel's functions. Since we have not considered a complete cosmological scenario, with a sequence of different phases, with an initial inflationary phase (as in standard cosmological model), the calculation performed here makes sense, strictly speaking, only when the strong energy condition is violated, that is, $\\alpha < - 1/3$. For such a case, there is a natural initial vacuum state, from which the particle occupation number can be determined. However, we formally extended the calculation for any value of $\\alpha$. It has been found that for those \"dark energy\" scenarios the rate of particle creation diverges as $t \\rightarrow \\infty$, that is, in the future infinity. \\par We have payed special attention to the phantom scenario. The reason is that a universe dominated by a phantom fluid may develop a singularity in a finite future time, the so-called big rip. In this case, we have found that it is possible that the energy density associated to the particles created (which obey in present case an equation of state of the type $p_s = \\rho_s$) becomes dominant over the phantom fluid if $- 1 > \\alpha > - 5/3$. Hence, the big rip can be avoided if the pressure is not deeply negative. \\par It is interesting to remark that a similar critical point, $\\alpha = - 5/3$, has been found in the case of classical scalar perturbation \\cite{fabris}. In that case, however, the evolution of scalar perturbation may destroy the conditions of homogeneity (necessary for the big rip) if $\\alpha < - 5/3$, that is, if the pressure is negative enough. The result found here is exactly the opposite: quantum effects can be operative in the sense of destroying the conditions for the big rip if the pressure is not negative enough. It must be stressed, however, that the evaluation made in the present work must be complemented by a study of more general quantum fields and by a deeper thermodynamical analysis of the energy balance between the phantom fluid and the created particles as the big rip is approached. \\newline \\vspace{0.5cm} \\newline {\\bf Acknowledgements}:\\\\ We thank CNPq (Brazil) and the french/brazilian scientific cooperation CAPES/COFECUB (project number 506/05) for partial financial support. We thank specially J\\'er\\^ome Martin for his criticisms on the text."
    },
    "0710/0710.1167_arXiv.txt": {
        "abstract": "A symplectic integrator algorithm suitable for hierarchical triple systems is formulated and tested. The positions of the stars are followed in hierarchical Jacobi coordinates, whilst the planets are referenced purely to their primary. The algorithm is fast, accurate and easily generalised to incorporate collisions. There are five distinct cases -- circumtriple orbits, circumbinary orbits and circumstellar orbits around each of the stars in the hierarchical triple -- which require a different formulation of the symplectic integration algorithm.  As an application, a survey of the stability zones for planets in hierarchical triples is presented, with the case of a single planet orbiting the inner binary considered in detail.  Fits to the inner and outer edges of the stability zone are computed. Considering the hierarchical triple as two decoupled binary systems, the earlier work of Holman \\& Wiegert on binaries is shown to be applicable to triples, except in the cases of high eccentricities and close or massive stars. Application to triple stars with good data in the multiple star catalogue suggests that more than 50 \\% are unable to support circumbinary planets, as the stable zone is non-existent or very narrow. ",
        "introduction": "\\label{sec:intro}  It is well known that many of the stars in the solar neighbourhood exist in multiple systems. As the number of planetary surveys increases, planets are regularly being found not only in single star systems, but binaries and triples as well. For example, recently a hot Jupiter has been claimed in the triple system HD 188753 (Konacki 2005; but see also Eggenberger et al. 2007), and \\citet{DB07} lists 33 binary systems and 2 other hierarchical triples known to harbour exoplanets. As the majority of work on planetary dynamics has been for single star systems, the dynamics of bodies in these multiple stellar systems is of great interest. At first sight, it might not seem likely to expect long term stable planetary systems to exist in binary star systems, let alone triples, but numerical and observational work is starting to show otherwise.  In recent years, much study has been devoted to the stability of planets and planetesimal discs in binary star systems. There are several investigations of individual systems (e.g. work on $\\gamma$ Cephei by \\citealt{Dv03}, \\citealt{Ha06} and \\citealt{VE06}) as well as substantial amounts of work on more general limits for stability of smaller bodies in these systems.  Notably, \\citet{HW99} approach this problem by using numerical simulation data to empirical fit critical semimajor axes for test particle stability as functions of binary mass ratio and eccentricity.  These general studies are of great use in the investigation of observed systems and their stability properties, giving an effective and fast method of placing limits on stability in the large parameter space created by observational uncertainties.  The aim of this work is to investigate test particle stability in hierarchical triple star systems, and to see if any similar boundaries can be defined for them.  To do this, a new method for numerically integrating planets in such systems is presented, following the ideas for binary systems presented by \\citet{CQDL02}.  Although there have been empirical studies of the stability of hierarchical star systems themselves, there do not appear to have been studies of small bodies in such systems. There is a great deal of literature concerning periodic orbits in the general three and four body problems (see e.g. \\citealt{Sz67}), but these are almost entirely devoted to considerations of planetary satellites in single star systems, for example satellite transfer in the Sun-Earth-Moon systems. Also, while periodic orbits prove that stable solutions exist in these problems, they are of little practical use in determining general stability limits.  The problem of planetary orbits in triple systems is more complex than for those in binary systems, as there are many different orbital configurations possible relative to the three stars. These are considered in Section~\\ref{sec:orbits}, and a method for classifying them is described in order to simplify the discussion in this paper. The derivation of a method to numerically integrate such planets is then given in Section~\\ref{sec:maths}.  In Section~\\ref{sec:stats} is a brief overview of the statistical properties of known triple systems, as a basis for deciding the parameters of the systems used in the numerical simulations. The results of numerical investigations into stability properties are then presented in Section~\\ref{sec:sp} for one of the possible orbital types. ",
        "conclusions": "\\label{sec:conclusion}  The main achievement of this paper is the formulation of a symplectic integrator algorithm suitable for hierarchical triple systems. This extends the algorithm for binary systems presented by \\citet{CQDL02}. The positions of the stars are followed in hierarchical Jacobi coordinates, whilst the planets are referenced purely to their primary. Each of the five distinct cases, namely circumtriple orbits, circumbinary orbits and circumstellar orbits around each of the stars in the hierarchical triple, requires a different splitting of the Hamiltonian and hence a different formulation of the symplectic integration algorithm. Here, we have given the mathematical details for each of the five cases, and presented a working code that implements the algorithm.  As an application, a survey of the stability zones for circumbinary planets in hierarchical triples is presented. Here, the planet orbits an inner binary, with a more distant companion star completing the stellar triple. Using a set of numerical simulations, we found the extent of the stable zone which can support long-lived planetary orbits and provided fits to the inner and outer edges. The effect of low inclination on this boundary is minimal. A reasonable first approximation to a behaviour of a hierarchical triple is to regard it as a superposition of the dynamics of the inner binary and a pseudo-binary consisting of the outer star and a point mass approximation to the inner binary. If it is considered as two decoupled binary systems, then the earlier work of Holman \\& Wiegert (1999) on binaries is applicable to triples, except in the cases of high eccentricities and close or massive stars.  The implication here is that the addition of a stable third star does not distort the original binary stability boundaries. As mentioned, \\citet{MW06} have shown that overlapping sub-resonances are the cause of the boundary in the binary case. It is reasonable to expect that in triples the same process is responsible, and the similarities between the binary and triple results support this theory. It is also expected that there is a regime in which the resonances from each sub-binary start to overlap as well, further destabilising the test particles. Evidence of this is the deviation from the binary results when the stars are close, massive and very eccentric, when resonances would be both stronger and wider. Since the parameter space investigated was chosen to reflect the observed systems it would seem to be a reflection on the characteristics of known triple stars that most lie in the decoupled regime. The relatively constant nature of the stellar orbits in the simulations is however a consequence of the test particle orbits being destabilised long before the stars are close enough to interact. By extension of all these arguments, it is expected that the binary criteria can be used to fairly accurately predict the stability zones in any hierarchical stellar system, no matter the number of stars.  The results presented here can be used to estimate the number of known hierarchical triple systems that could harbour S(AB)-P planets. \\citet{To97} lists 54 systems with semimajor axis, eccentricity and masses for both the inner and outer components. The mutual inclinations of most are not well known, but there are nine systems listed for which this angle can be determined.  For five of these it is less than $15^\\circ$, two are around $40^\\circ$ and two are retrograde. Although a small sample it suggests that there are systems that fall within the low inclination regime investigated here.  Using the criteria of \\cite{HW99} and those found here for the position of the inner and outer critical semimajor axis the size of the coplanar stable region for each of these triples can be calculated.  This can be considered an upper limit, since it is likely that very non-coplanar systems and those with significant eccentricities for both binary components will further destabilise planets.  Of the 54 systems 13 are completely unstable to circumbinary planets according to the four parameter fits (compared to 11 using \\citet{HW99}'s criteria). Figure \\ref{fig:zone} shows a plot of the width of the stable region for the remaining systems.  Interestingly, the majority seem to have very small stable zones, with 16 smaller than an au. Whether this is a feature of triple systems, or an observation bias is not apparent.  It does indicate though that circumbinary planets are unlikely to exist in at least 50 \\% of observable systems.    \\begin{figure} \\centering \\includegraphics[width=\\textwidth]{fig14.eps} \\caption{ Widths of the circumbinary stability zones for known triple systems in au and as a percentage of the area between the two sub-binaries, calculated using the four parameter fits derived in Section~\\ref{sec:sp}. Note that using the criteria of \\citet{HW99} gives almost identical results.} \\label{fig:zone} \\end{figure}"
    },
    "0710/0710.1021_arXiv.txt": {
        "abstract": "\\noindent We test statistically the hypothesis that radio pulsar glitches result from an avalanche process, in which angular momentum is transferred erratically from the flywheel-like superfluid in the star to the slowly decelerating, solid crust via spatially connected chains of local, impulsive, threshold-activated events, so that the system fluctuates around a self-organised critical state. Analysis of the glitch population (currently 285 events from 101 pulsars) demonstrates that the size distribution in individual pulsars is consistent with being scale invariant, as expected for an avalanche process. The measured power-law exponents fall in the range $-0.13\\leq a \\leq 2.4$, with $a\\approx 1.2$ for the youngest pulsars. The waiting-time distribution is consistent with being exponential in seven out of nine pulsars where it can be measured reliably, after adjusting for observational limits on the minimum waiting time, as for a constant-rate Poisson process. PSR J0537$-$6910 and PSR J0835$-$4510 are the exceptions; their waiting-time distributions show evidence of quasiperiodicity. In each object, stationarity requires that the rate $\\lambda$ equals $- \\epsilon \\dot{\\nu} / \\langle\\Delta\\nu\\rangle$, where $\\dot{\\nu}$ is the angular acceleration of the crust, $\\langle\\Delta\\nu\\rangle$ is the mean glitch size, and $\\epsilon\\dot{\\nu}$ is the relative angular acceleration of the crust and superfluid. Measurements yield $\\epsilon \\leq 7 \\times 10^{-5}$ for PSR J0358$+$5413 and $\\epsilon \\leq 1$ (trivially) for the other eight objects, which have $a < 2$. There is no evidence that $\\lambda$ changes monotonically with spin-down age. The rate distribution itself is fitted reasonably well by an exponential for $\\lambda \\geq 0.25\\,{\\rm yr^{-1}}$, with $\\langle \\lambda \\rangle = 1.3^{+0.7}_{-0.6}\\,{\\rm yr^{-1}}$. For $\\lambda < 0.25\\,{\\rm yr^{-1}}$, its exact form is unknown; the exponential overestimates the number of glitching pulsars observed at low $\\lambda$, where the limited total observation time exercises a selection bias. In order to reproduce the aggregate waiting-time distribution of the glitch population as a whole, the fraction of pulsars with $\\lambda > 0.25\\,{\\rm yr^{-1}}$ must exceed $\\sim 70$ per cent. ",
        "introduction": "} Glitches are tiny, impulsive, randomly timed increases in the spin frequency $\\nu$ of a rotation-powered pulsar, sometimes accompanied by an impulsive change in the frequency derivative $\\dot{\\nu}$. They are to be distinguished from timing noise, a type of rotational irregularity where pulse arrival times wander continuously, although there is evidence that timing noise is the cumulative result of frequent microglitches in certain pulsars \\citep{cor85,dal95}.  At the time of writing, 285 glitches in total have been detected in 101 objects ($\\sim 6\\%$ of the known radio pulsar population), the majority in the last four years, facilitated by the Parkes Multibeam Survey, refined multifrequency ephemerides, and better interference rejection algorithms \\citep{hob02,kra03,kra05,lew05,jan06}. Efforts to analyse the data statistically have focused on the correlation of glitch activity with age \\citep{mck90,she96,ura99,lyn00,wan00} and Reynolds number \\citep{peralta06,mel07}, the post-glitch relaxation time-scale \\citep{wan00,won01}, the size distribution \\citep{mor93a,mor93b,peralta06}, and the correlation between glitch sizes and waiting times \\citep{wan00,won01,mid06,peralta06}. \\citet{hob02} reviewed the role of observational selection effects.  Most glitching pulsars ($65\\%$) have been seen to glitch once, but a minority glitch repeatedly; the current record holder is PSR J1740$-$3015, with 33 glitches. Of those objects which glitch repeatedly, most do so at unpredictable intervals, but two (PSR J0537$-$6910 and Vela) are quasiperiodic; Vela, in particular, has been likened to a relaxation oscillator \\citep{lyn96}. The fractional increase in $\\nu$ spans seven decades ($3\\times 10^{-11} \\leq \\Delta\\nu \\leq 2\\times 10^{-4}$) across the glitch population and as many as four decades in a single object (e.g.\\ $7\\times 10^{-10} \\leq \\Delta\\nu \\leq 2\\times 10^{-6}$ in PSR J1740$-$3015). The spin-down age $\\tau_{\\rm c}= -\\nu/(2\\dot{\\nu})$ of glitching pulsars spans four decades, from $1\\times 10^3\\,{\\rm yr}$ to $3\\times 10^7\\,{\\rm yr}$. In many respects, therefore, the glitch phenomenon is {\\em scale invariant}. This striking property invites physical interpretation.  Theories of pulsar glitches have focused mainly on the local microphysics of the superfluid in the stellar interior and its coupling to the solid crust, for example the strength of vortex pinning \\citep{and75,jon98}, the rate of vortex creep \\citep{lin96}, or the conditions for exciting superfluid turbulence \\citep{per05,per06,mel07,and07}. Ultimately, however, the local microphysics must be synthesized with the global, {\\em collective} dynamics in order to make full contact with observational data. (Likewise, a practical model of earthquakes must synthesize the microphysics of rock fracture with the macrodynamics of interacting tectonic plates.) For example, if approximately $10^{16} (\\Delta\\nu/1\\,{\\rm Hz})$ vortices unpin from crustal lattice sites in sympathy during a glitch, they must communicate rapidly across distances much greater than their separation. How? And why does the number that unpin fluctuate so dramatically (by up to four orders of magnitude) from glitch to glitch in a single pulsar, while always amounting to a small fraction ($\\Delta\\nu/\\nu$) of the total?  Such collective, scale invariant behavior is a generic feature of a class of natural and synthetic far-from-equilibrium systems, called self-organized critical systems, that are discrete, interaction dominated, and slowly driven, and that adjust internally via erratic, spatially connected {\\em avalanches} of local, impulsive, threshold-activated, relaxation events \\citep{jen98}. Such systems fluctuate around a stationary state towards which they evolve spontaneously, in which global driving balances local relaxation on average over the long term. The archetype of a self-organized critical system is the sandpile \\citep{bak87}.  In this paper, we study pulsar glitches as an avalanche process, as first proposed by \\citet{mor93a}. After reviewing self-organised criticality in \\S\\ref{sec:gli2}, we define the statistical sample on which our study is based (\\S\\ref{sec:gli3}) and analyze the observed distribution of glitch sizes (\\S\\ref{sec:gli4}) and waiting times (\\S\\ref{sec:gli5}). Some implications for glitch physics are explored in \\S\\ref{sec:gli6}. We only include radio pulsars in the sample, to preserve its homogeneity, even though glitches have now been observed in anomalous X-ray pulsars (magnetars) as well \\citep{dal03,kas03}. ",
        "conclusions": "} In this paper, we analyze the size and waiting-time distributions of pulsar glitches, taking advantage of the enlarged data set produced by the Parkes Multibeam Survey. We conclude that the data are consistent with the hypothesis that pulsar glitches arise from an avalanche process. In each of seven pulsars with $N_{\\rm g} > 5$, the size distribution is consistent with being scale invariant across the observed range of $\\Delta\\nu$ (up to four decades), and the waiting-time distribution is consistent with being Poissonian. These features are natural if the system is driven globally at a constant rate (as the pulsar spins down), and each glitch corresponds to a locally collective, threshold activated relaxation of one of the many spatially independent, metastable stress reservoirs in the system (e.g. via a vortex unpinning or crust cracking avalanche). In two pulsars, PSR J0537$-$6910 and PSR J0835$-$4510, the dynamics may include a second, quasiperiodic component, comprising $\\sim 20\\%$ of the events. The size and waiting-time distributions of the quasiperiodic component are narrowly peaked, as expected for rare, system-spanning avalanches, which relax a large fraction of the total stress accumulated in the system. This two-component behavior is observed widely in self-organised critical systems, including experiments on magnetic flux vortices in type II superconductors, which are closely analogous to neutron star superfluids \\citep{fie95}.  The power-law exponent of the size probability density function differs from pulsar to pulsar, spanning the range $-0.13 \\leq a \\leq 2.4$. Calculating $a$ theoretically from first principles is a deep problem which has not yet been solved for other self-organised critical systems, let alone glitching pulsars, although some progress has been made on two-dimensional sandpiles using renormalization group techniques \\citep{pie94,jen98}. In the mean-field approximation, which is exact in four or more dimensions, theoretical calculations on sandpiles (and other systems in their universality class) yield $a = 1.5$, whereas three-dimensional cellular automata output $a=1.3$ \\citep{jen98}.  The size distribution transmits two important lessons concerning the microphysics of glitches. First, the fact that $a$ differs from pulsar to pulsar implies that the strength and level of conservation of the local (e.g.\\ pinning and intervortex) forces also differs \\citep{ola92,fie95}. By contrast, in equilibrium critical systems like ferromagnets, $a$ depends only on the dimensionality of the system and its order parameter and is therefore universal. Second, except for the two pulsars which show quasiperiodicity, $a$ appears to vary smoothly with spin-down age, with $a\\approx 1.2$ for the youngest pulsars (e.g.\\ the Crab). Figure \\ref{fig:gli9} depicts the trend between $a$ and $\\tau_{\\rm c}$. It is suggestive; after all, local pinning forces do depend on temperature and hence $\\tau_{\\rm c}$. Interestingly, however, there is no clear trend between $a$ and $\\nu$, even though the mean vortex spacing (and hence intervortex force) is proportional to $\\nu^{1/2}$. It will pay to study these trends more thoroughly as more glitch data is collected.  An avalanche process predicts a specific relation between the distributions of glitch sizes $\\Delta\\nu$ and lifetimes $T$ (as opposed to waiting times $\\Delta t$). Specifically, in a self-organized critical state, the lifetime probability density function is also a power law, $p(T) \\propto T^{-b}$, with \\begin{equation} b = 1 + (a-1)\\gamma_2 / \\gamma_3~. \\label{eq:gli16} \\end{equation} The constants $\\gamma_2$ and $\\gamma_3$ are defined such that the cardinality of an avalanche scales with its linear extent ($L$) as $L^{\\gamma_2}$ and its lifetime (i.e.\\ duration) scales as $L^{\\gamma_3}$ \\citep{jen98}. Both $\\gamma_2$ and $\\gamma_3$ depend on the effective dimensionality of the local forces and can be calculated numerically using a cellular automaton. In two dimensions, avalanches are compact, not fractal, and one has $\\gamma_2=2$; in three dimensions, one has $2 < \\gamma_2 < 3$. At present, radio timing experiments cannot measure $T$; most glitches are detected as unresolved, discontinuous, spin-up events with $T < 120\\,{\\rm s}$ \\citep{mcc90}. \\footnote{In the Crab, some spin-up events seem to be resolved, e.g.\\ at epochs MJD 50260 ($T \\approx 0.5\\,{\\rm d}$) and MJD 50489 (secondary spin up, $T \\approx 2\\,{\\rm d}$) \\citep{won01}. If these are rare but otherwise standard glitches originating from the long-$T$ tail of the lifetime distribution, it is puzzling that other, shorter, but still resolved (and presumably more common) spin-up events are not observed, e.g.\\ with $T\\sim 0.1\\,{\\rm d}$ or $0.01\\,{\\rm d}$. Alternatively, the events at MJD 50620 and MJD 50489 may have been triggered by a different physical mechanism.} In the future, however, single- and/or giant-pulse timing experiments with more sensitive instruments (e.g.\\ the Square Kilometer Array) will test this prediction. If confirmed, it will independently corroborate the avalanche hypothesis.  The mean glitching rates of the nine pulsars studied here are fairly narrowly distributed, spanning the range $0.35\\,{\\rm yr^{-1}} \\leq \\lambda \\leq 2.6\\,{\\rm yr^{-1}}$. The probability density function for $\\lambda$ is adequately fitted by an exponential, as for solar flare avalanches \\citep{whe00}, with $\\langle\\lambda\\rangle = 1.3_{-0.6}^{+0.7}$ {\\rm yr}$^{-1}$, or by an exponential with a lower cutoff, at $\\lambda_{\\rm min} \\approx 0.25$ {\\rm yr}$^{-1}$. A theoretical derivation of $\\langle\\lambda\\rangle$ from first principles is currently lacking, although estimates of how long it takes to crack the crust locally predict reasonable rates, if the critical strain angle approaches that of imperfect terrestrial metals \\citep{alp96,mid06}.  Figure \\ref{fig:gli10} plots $\\lambda$ versus $\\tau_c$ for the nine pulsars examined individually in this paper. There is no significant trend. The data are consistent with the notion that old pulsars glitch less frequently than young pulsars \\citep{she96}, but they are equally consistent with the notion that the glitching rate is independent of age.  Many authors have searched for a correlation between waiting time and the size of the next glitch. Such a correlation appears to be absent from the data, e.g.\\ Figure 17 in \\citet{wan00} and Figure 10 in \\citet{won01}. At first blush, this is surprising: the vortex unpinning and crust fracture paradigms, which are driven by the accumulation of differential rotation and mechanical stress respectively, seem to be natural candidates for a `reservoir effect'. Avalanche dynamics resolves this apparent paradox. The reservoir effect does operate locally, but the star contains many reservoirs, insulated from each other by relaxed zones, whose storage capacities evolve stochastically in response to the slow driver and avalanche history. During a glitch, a single reservoir (often small but sometimes large) relaxes at random via an avalanche, releasing its stored $\\Delta\\nu$ (and destabilizing neighboring reservoirs in preparation for the next glitch). Some of the $\\Delta\\nu$ is accumulated since the previous glitch, but the remainder is `borrowed' from earlier epochs, when some other reservoir relaxed instead. All self-organized critical systems share these dynamics; the waiting time is uncorrelated with the size of the next avalanche \\citep{jen98}. The only exceptions are large, system-spanning avalanches, which always have roughly the same sizes and waiting times, and which account for $\\sim 20\\%$ of the glitches in PSR J0537$-$6910 and PSR J0835$-$4510.  A corollary of the previous paragraph is that the total $\\Delta\\nu$ released in glitches up to some epoch is less than the total crust-superfluid differential rotation accumulated since that epoch, viz. \\begin{equation} \\sum_{i=1}^{N_{\\rm g}} \\Delta\\nu_i \\leq \\epsilon |\\dot{\\nu}|  \\sum_{i=1}^{N_{\\rm g}} \\Delta t_{i}, \\label{eq:gli17} \\end{equation} where $\\epsilon\\dot{\\nu}$ is the relative angular acceleration of the crust and superfluid due to electromagnetic spin down. The `staircase' described by (\\ref{eq:gli17}) has been noted previously \\citep{she96,lyn00}, both in quasiperiodic glitchers like PSR J0537$-$6910 [e.g.\\ Figure 8 in \\citet{mid06}], where the reservoir effect is obvious, and in Poisson glitchers like PSR J0534$+$2200, [e.g.\\ Figure 12 in \\citet{won01}], where the trend is more subtle because it reverts to the mean over long times, not after every glitch. Upon dividing (\\ref{eq:gli17}) by $N_{\\rm g}$, and averaging over long times, the inequality becomes an equality (provided there is no secular accumulation of differential rotation in the system) and we recover (\\ref{eq:gli5}).  It is fundamentally impossible to measure $\\epsilon$ for individual pulsars with current data, because $\\langle \\Delta \\nu \\rangle$ is dominated by large (and therefore rare) glitches for $a < 2$. It is therefore wrong to assume stationarity over a typical, $40$-yr observation interval. Consequently, we are prompted to reassess the familiar correlation between activity and spin-down age \\citep{she96}. Our definition of $\\epsilon\\dot{\\nu}$ is identical to $\\dot{\\nu}_{\\rm glitch}$ in \\citet{lyn00} (but for individual objects, not in aggregate) and $A_{\\rm g}$ in \\citet{won01}. It is closely related to the original activity parameter defined by \\citet{mck90}, which equals $N_{\\rm g} \\nu^{-1} \\epsilon |\\dot{\\nu}|$. For PSR J0358$+$5413, we measure $\\epsilon\\leq 7 \\times 10^{-5}$, lower than the {\\it aggregate} value $0.017\\pm0.002$ measured by \\citet{lyn00} for objects with $\\tau_{\\rm c} > 10\\,{\\rm kyr}$ (binned by semi-decades in $\\dot{\\nu}$). \\footnote{The aggregate value $\\dot{\\nu}_{\\rm glitch}$ \\citep{lyn00}, binned over semi-decades in $\\dot{\\nu}$, effectively averages together different pulsars. While this approach reduces the formal error bar on $\\dot{\\nu}_{\\rm glitch}$, its physical interpretation is less straightforward, given the likelihood that $\\epsilon$ is different in different pulsars.} Interpreted in terms of the vortex unpinning model, this result suggests that $0.007$--$2$ \\% of the angular momentum outflow during spin down may be stored in metastable reservoirs on average over time. On the other hand, five other objects have $0.04 \\leq \\epsilon_{\\rm max} \\leq 0.8$, under the questionable assumption that the maximum physical size is $\\Delta \\nu_{\\rm upper} = 2 \\times 10^{-4}$ in all pulsars. Our data are therefore inadequate to update usefully the value $A_{\\rm g}/|\\dot{\\nu}| = 1\\times 10^{-5}$ measured by \\citet{won01} for PSR J0534$+$2200.  In the context of vortex unpinning, it has been argued that the aggregated $\\epsilon$ measured by \\citet{lyn00} partly corroborates the hypothesis that younger pulsars are still in the process of forming their capacitive elements, e.g.\\ by creating pinning centers through crust fracture, while older pulsars have mostly completed the task \\citep{alp96,won01}. However, the full picture is more complicated. Vela's quasiperiodic avalanches point to a richly connected network of reservoirs \\citep{alp96}, yet its aggregated value $\\dot{\\nu}_{\\rm glitch}$ is relatively low. On the other hand, the other quasiperiodic glitcher, PSR J0537$-$6910, is relatively young ($\\tau_{\\rm c} = 4.9\\,{\\rm kyr}$); how did it form a richly connected reservoir network so quickly? And, if its network is so richly connected, why is its aggregated $\\dot{\\nu}_{\\rm glitch}$ value so low? Likewise, PSR J0358$+$5413 is the oldest object in the sample ($\\tau_{\\rm c} = 560 \\,{\\rm kyr}$), yet its $\\epsilon$ value arguably points to a dearth of capacitive elements, characteristic of a young object. There are no obvious grounds (e.g.\\ quasiperiodicity) on which to treat PSR J0358$+$5413 as exceptional.  Do all pulsars glitch eventually? It has been speculated in the past that there is something special physically about the minority of pulsars that do glitch. While it is impossible to reject this hypothesis unequivocally with the data at hand, the results presented here suggest that all pulsars are capable of glitching. However, most do so infrequently (low $\\lambda$) and hence have not been detected during the last four decades of timing experiments. We find that up to $\\sim 30 \\%$ of the pulsar population can glitch at rates lower than $\\lambda_{\\rm min}= 0.25$ {\\rm yr}$^{-1}$ and still conform with the measured aggregate waiting-time distribution.  Once verified, the claimed Poissonian nature of the glitch mechanism can be invoked to exclude broad classes of glitch theories, e.g.\\ those that rely on `defects' or `turbulence' at special locations (like the pole), or that involve a pair of dependent events (A.\\ Martin, private communication). It is important to interpret aftershocks carefully in this light \\citep{won01}. In self-organized critical systems, the excess number of avalanches following a large avalanche (over and above the Poissonian baseline following a small avalanche) scales inversely with the time elapsed, a property known as Omori's law for earthquakes \\citep{jen98}.  In this paper, we do not analyze post-glitch relaxation times and glitch-activated changes in $\\dot{\\nu}$ in the context of avalanche processes, e.g.\\ the correlation between $\\Delta\\dot{\\nu}$ and the transient component of $\\Delta\\nu$ \\citep{won01}. We also assume implicitly that the quantized superfluid vortices in the vortex unpinning model are organized in a rectilinear array, even though recent work suggests that meridional circulation destabilizes the array and converts it into a turbulent tangle \\citep{per05,per06}. Further study of these matters is deferred to future work.  We acknowledge the computer time and system support supplied by the Australian Partnership for Advanced Computation (APAC) and the Victorian Partnership for Advanced Computation (VPAC). We thank Andre Trosky for illuminating conversations on self-organized criticality and cellular automata. This research was supported by a postgraduate scholarship from the University of Melbourne. It makes use of the Australia Telescope National Facility Pulsar Catalogue \\citep{man05}, which can be accessed on-line at {\\tt http://www.atnf.csiro.au/research/pulsar/psrcat}."
    },
    "0710/0710.3815_arXiv.txt": {
        "abstract": "The Newtonian solid-mechanical theory of non-compressional spheroidal and torsional nodeless elastic vibrations in the homogenous crust model of a quaking neutron star is developed and applied to the modal classification of the quasi-periodic oscillations (QPOs) of X-ray luminosity in the aftermath of giant flares in SGR 1806-20 and SGR 1900+14. A brief outline is given of Rayleigh's energy method which is particular efficient when computing the frequency of nodeless elastic spheroidal and torsional shear modes as a function of multipole degree of nodeless vibrations and two input parameters -- the natural frequency unit of shear vibrations carrying information about equation of state and fractional depth of peripheral seismogenic layer. In so doing we discover that the dipole overtones of both spheroidal and torsional nodeless vibrations possess the properties of Goldstone soft modes. It is shown that obtained spectral formulae reproduce the early suggested identification of the low-frequency QPOs from the range $30\\leq \\nu \\leq 200$ Hz as torsional nodeless vibration modes $\\nu(_0t_\\ell)$ of multipole degree $\\ell$ in the interval $2\\leq \\ell \\leq 12$. Based on this identification, which is used to fix the above mentioned input parameters of derived spectral formulae, we compute the frequency spectrum of nodeless spheroidal elastic vibrations $\\nu(_0s_\\ell)$. Particular attention is given to the low-frequency QPOs in the data for SGR 1806-20 whose physical origin has been called into question. Our calculations suggest that these unspecified QPOs are due to nodeless dipole torsional and dipole spheroidal elastic shear vibrations: $\\nu(_0t_1)=18$ Hz and $\\nu(_0s_1)=26$ Hz. ",
        "introduction": "The discovery of QPOs of X-ray luminosity in the aftermath of giant flares SGR 1806-20 and SGR 1900+14 (Israel et al 2005; Strohmayer \\& Watts 2006; Israel 2007; Watts \\& Strohmayer 2007) with concomitant interpretation of QPOs as caused by quake-induced differentially rotational seismic vibrations has stimulated remarkable developments in the magnetar asteroseismology (e.g. Piro 2005, Samuelsson \\& Andersson 2007a, 2007b; Lee 2007; Levin 2007; Watts \\& Reddy 2007; Sotani et al 2007; Bastrukov et al 2007a and references therein). Following the above interpretation and presuming the dominant role of the elastic restoring force, the focus of  most theoretical works is on computing the frequency spectra of odd-parity torsional mode of shear vibrations and less attention is paid to the even parity spheroidal elastic mode. However, from the viewpoint of modern global seismology (Lay \\& Wallace 1995; Aki \\& Richards 2003), the spheroidal vibrational mode in a solid star and planet has the same physical significance as the toroidal one in the sense that these two fundamental modes owe their existence to one and the same restoring force (e.g. McDermott, Van Horn, Hansen 1988; Bastrukov, Weber, Podgainy 1999, Bastrukov et al 2007b). In this light there is a possibility that, by not considering both these modes on an equal footing, we may miss discovering certain essential novelties which are consequences of solid mechanical laws governing seismic vibrations of superdense matter of neutron stars. Adhering to this attitude and continuing the investigations recently reported by Bastrukov et al (2007a), we derive here spectral equations for the frequency of both spheroidal and torsional elastic nodeless vibrations in the solid crust of quaking neutron star and examine what conclusions can be drawn regarding low-frequency QPOs whose physical nature still remain unclear. In Sec.2 by use of the energy variational method we derive spectral formulae for the frequencies of the nodeless spheroidal and torsional elastic vibrations locked in the finite-depth seismogenic layer. Particular attention is given to the dipole spheroidal and torsional vibrations possessing properties of Goldstone's soft modes. In sec.3, the obtained spectral formulae are applied to a modal analysis of available data on the above mentioned QPOs. The obtained results are highlighted in Sec.4. ",
        "conclusions": "The exact spectral formulae, which has been obtained here within the framework of Newtonian, non-relativistic, solid-mechanical theory of seismic vibration for the first time, are interesting in its own right from the viewpoint of general theoretical seismology (e.g. Lay, Wallace 1995) because they can be utilized in the study of seismic vibrations of more wider class of solid celestial objects such as Earth-like planets. One of the remarkable findings of our investigation is that the dipole overtones of nodeless elastic shear vibrations trapped in the finite-depth crust of seismically active neutron star possess properties of Goldstone soft modes. It is shown that obtained spectral equations are consistent with the existence treatment of low-frequency QPOs in the X-ray luminosity of flares SGR 1900+14 and  SGR 1806-20 as caused by quake-induced torsional nodeless vibrations (Samuelsson \\& Andersson 2007a; 2007b). What is newly disclosed here is that previously non-identified low-frequency QPOs in data for SGR 1806-20 can be attributed to nodeless dipole torsional and spheroidal vibrations, namely, $\\nu(_0t_1)=18\\,\\mbox{Hz}$ and $\\nu(_0s_1)=26\\,\\mbox{Hz}$."
    },
    "0710/0710.5508_arXiv.txt": {
        "abstract": "We analyse a sample of 33 extensive air showers (EAS) with estimated primary energies above $2\\cdot 10^{19}$~eV and high-quality muon data recorded by the Yakutsk EAS array.  We compare, event-by-event, the observed muon density to that expected from CORSIKA simulations for primary protons and iron, using SIBYLL and EPOS hadronic interaction models. The study suggests the presence of two distinct hadronic components, ``light'' and ``heavy''. Simulations with EPOS are in a good agreement with the expected composition in which the light component corresponds to protons and the heavy component to iron-like nuclei. With SYBILL, simulated muon densities for iron primaries are a factor of $\\sim 1.5$ less than those observed for the heavy component, for the same electromagnetic signal. Assuming two-component proton-iron composition and the EPOS model, the fraction of protons with energies $E>10^{19}$~eV is $0.52 ^{+0.19}_{-0.20}$ at 95\\% confidence level. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.4162_arXiv.txt": {
        "abstract": "Estimates of inflationary parameters from the CMB $B$-mode polarization spectrum on the largest scales depend on knowledge of the reionization history, especially at low tensor-to-scalar ratio. Assuming an incorrect reionization history in the analysis of such polarization data can strongly bias the inflationary parameters. One consequence is that the single-field slow-roll consistency relation between the tensor-to-scalar ratio and tensor tilt might be excluded with high significance even if this relation holds in reality. We explain the origin of the bias and present case studies with various tensor amplitudes and noise characteristics. A more model-independent approach can account for uncertainties about reionization, and we show that parametrizing the reionization history by a set of its principal components with respect to $E$-mode polarization removes the bias in inflationary parameter measurement with little degradation in precision. ",
        "introduction": "\\label{sec:intro}   Temperature and polarization power spectra of the cosmic microwave background (CMB) are consistent with predictions of the simplest inflationary models~\\cite{Gut81,AlbSte82,Lin82,Sat81}: a nearly flat geometry, superhorizon correlations as probed by the spectrum of acoustic peaks, and primordial scalar perturbations that are adiabatic, Gaussian, and close to scale-invariant~\\cite{HuWhi96c,SpeZal97,HuSpeWhi97,Peietal03,Speetal07}. One of the key remaining signatures of inflation, tensor perturbations (i.e.\\ gravitational waves)~\\cite{KamKosSte97,SelZal97}, has yet to be detected. Depending on the amplitude of the tensor perturbations, which is not well constrained theoretically, it may be possible to measure the angular power spectrum of the inflationary gravitational waves in the $B$-mode component of the CMB polarization on large scales. Non-detection of the tensor spectrum does not necessarily rule out inflation, but upper limits on $r$ can be used to exclude particular models of inflation and limit its energy scale. Many experiments have been proposed to search for this signal~\\cite{Planck,CMBTaskForce,Oxletal04,LawGaiSei04,Ruhl:2004kv,Mafetal05,Yooetal06,Kogut:2006nd,Macetal07,Polenta07}.  Measurement of  tensor perturbations in the $B$-mode polarization power spectrum would test models of inflation by constraining inflationary parameters. These parameters include the tensor-to-scalar ratio, $r$, and the tensor spectral index, $n_t$, which are related by a consistency relation under the simplest single-field slow-roll inflationary scenarios. If the tensor spectrum can be detected, precise measurements over a wide range of scales could test the consistency relation.   CMB constraints on $r$ and $n_t$ depend on the ability to accurately determine the large-scale power in $B$-modes due to tensor perturbations, independent of the effects of other cosmological parameters. On the largest scales, the tensor $B$-mode spectrum depends not only on inflationary parameters but also on the reionization history of the universe~\\cite{Zal97}. The main impact of reionization on the spectrum is through the total optical depth, $\\tau$.  The 3-year \\wmap\\ measurements of $E$-mode polarization determine $\\tau$ to an accuracy of about $30\\%$~\\cite{Pagetal07,Speetal07}, and future CMB experiments should constrain $\\tau$ at the~$5-10\\%$~level~\\cite{Holetal03,KeaMil06,Planck,MorHu07b}.  However, the detailed evolution of the reionization history also significantly affects the large-scale polarization spectra.  Uncertainty in this history leads to added uncertainty in inflationary parameters.  Moreover, incorrect inferences due to an oversimplified treatment of reionization may bias estimates of inflationary parameters. For unbiased estimation of the optical depth from the $E$-mode reionization peak, the solution is to use a complete, principal-component-based description of reionization effects  when estimating parameters from CMB polarization data~\\cite{HuHol03,MorHu07b}. In this paper, we extend this approach to tensor $B$-mode polarization and show that it is equally if not more effective in ensuring accurate measurements with little loss in precision.   The outline of the paper is as follows. We discuss the effects of reionization and inflationary parameters on the polarization power spectra and the large-scale degeneracy between these parameters in \\S~\\ref{sec:param}. A brief overview of the principal component parametrization of the reionization history follows in \\S~\\ref{sec:pcs}. In \\S~\\ref{sec:mcmc}, we describe our Markov Chain Monte Carlo analysis of simulated polarization data and give the resulting constraints on $\\tau$, $r$, and $n_t$, which we discuss further in \\S~\\ref{sec:discuss}.       \\begin{figure} \\centerline{\\psfig{file=f1.eps, width=3.0in}} \\caption{$B$-mode tensor spectra illustrating the degeneracy between $\\tau$ and $n_t$ for large-scale measurements, with angular power spectra plotted in the upper panel and fractional deviations from the base model in the lower panel. For the base model (\\emph{solid}), $r=0.03$, $\\tau=0.1$, and $n_t=-0.00375$. The other two models have $\\{\\tau,n_t\\}=\\{0.12,-0.00375\\}$ (\\emph{long dashed}) and $\\{0.12,0.13\\}$ (\\emph{short dashed}), with a pivot scale of $k_{\\rm pivot}=0.01~{\\rm Mpc}^{-1}$. The reionization history here is assumed to be instantaneous. Cosmic variance of $\\cltens$ for the base model, which excludes the variance from lensing, is shown by the shaded band in the lower panel. } \\label{fig:degen} \\end{figure} ",
        "conclusions": "\\label{sec:discuss}   The value of the optical depth to reionization estimated from the CMB $E$-mode polarization spectrum on large scales can be biased by adopting a model that has insufficient freedom to describe the true reionization history.  Likewise, the use of simple reionization models can bias inflationary parameters such as the tensor-to-scalar ratio and tensor tilt that depend on the large-scale amplitude of the $B$-mode spectrum of primordial gravitational waves. In each case, the problem can be solved by using a more general parametrization of the reionization history. We have shown that using a small but complete set of the principal components of the reionization history effectively yields unbiased constraints on both reionization and inflationary parameters.   Measurements of $r$ and $n_t$ are only affected by the assumed form of the reionization history if the reionization peak of the tensor $B$-mode spectrum at the very largest scales is needed to precisely constrain the parameters.  If, instead, good constraints can be obtained using only the $B$-mode recombination peak at intermediate scales, then assumptions about reionization do not affect tests of the consistency relation between $r$ and $n_t$.  They would instead appear as false evidence for running of the tensor tilt in violation of slow-roll expectations. Measurement of the recombination peak however is inhibited by experimental noise and contamination from $E$-mode power converted to $B$-mode power by gravitational lensing, both of which become more important at smaller scales.   To study the potential impact of reionization on parameter constraints from $B$-mode polarization, we have employed a Markov Chain Monte Carlo analysis of simulated CMB polarization power spectra and compared results for two descriptions of reionization: a simple, one-parameter, instantaneous reionization model, and a parametrization using principal components of the reionization history with respect to the $E$-mode polarization power spectrum.  By varying the properties of the simulated polarization power spectra, including the fiducial tensor-to-scalar ratio and the noise spectrum, we have determined over what range of scales CMB polarization data is most important for constraining inflationary parameters in various scenarios. In particular, the question of whether the large-scale reionization peak of $\\cltens$ or the smaller-scale recombination peak is more important determines the severity of bias in inflationary parameters when reionization is modeled incorrectly.  If the tensor-to-scalar ratio is near the current upper limit of $r\\sim 0.3$ and measurements of $B$-mode polarization are limited only by cosmic variance, then the spectrum on scales $20 \\lesssim \\ell \\lesssim 500$ dominates constraints on $r$ and $n_t$ and incorrect assumptions about reionization do not strongly bias the results. If the true tensor-to-scalar ratio is more than a factor of a few smaller than this upper bound, however, then lensing $B$-modes limit the information that can be extracted from the recombination peak of the tensor spectrum alone. Furthermore, all-sky experiments in the foreseeable future are likely to have noise that exceeds the lensing signal, making tests of inflation even more reliant on the reionization peak of the tensor $B$-modes on large scales. In all of these cases, a general parametrization of reionization such as that provided by principal components allows the use of the $B$-mode reionization peak for inflationary parameter constraints without significantly worsening the errors on those parameters.   \\vfill"
    },
    "0710/0710.4448_arXiv.txt": {
        "abstract": "The first high resolution \\textit{Spitzer} IRS 9-37$\\mu$m spectra of 29 Seyfert galaxies (about one quarter) of the 12$\\mu$m Active Galaxy Sample are presented and discussed. The high resolution spectroscopy was obtained with corresponding off-source observations. This allows excellent background subtraction, so that the continuum levels and strengths of weak emission lines are accurately measured. The result is several new combinations of emission line ratios, line/continuum and continuum/continuum ratios that turn out to be effective diagnostics of the strength of the AGN component in the IR emission of these galaxies. The line ratios [NeV]/[NeII], [OIV]/[NeII], already known, but also  [NeIII]/[NeII] and [NeV]/[SiII] can all be effectively used to measure the dominance  of the AGN. We extend the analysis, already done using the 6.2$\\mu$m PAH emission feature, to the equivalent width of the 11.25$\\mu$m PAH feature, which also anti-correlates with the dominance of the AGN. We measure that the 11.25$\\mu$m PAH feature has a constant ratio with the H$_2$ S(1) irrespective of Seyfert type, approximately 10 to 1. Using the ratio of accurate flux measurements at about 19$\\mu$m with the two spectrometer channels, having aperture areas differing by a factor 4, we measured the source extendness and correlated it with the emission line and PAH feature equivalent widths. The extendness of the source gives another measure of the AGN dominance and correlates both with the EWs of [NeII] and PAH emission. Using the rotational transitions of H$_2$ we were able to estimate temperatures (200-300K) and masses (1-10 $\\times$ 10$^{6}$ M$_{\\sun}$), or significant limits on them, for the warm molecular component in the galaxies observed. Finally we used the ratios of the doublets of [NeV] and of [SIII] to estimate the gas electron density, which appears to be of the order of n$_e$ $\\sim$ 10$^{3-4}$ cm$^{-3}$ for the highly ionized component and a factor 10 lower for the intermediate ionization gas. ",
        "introduction": "Mid-infrared (mid-IR) spectroscopy provides a powerful tool to investigate the nature and physical processes in active galactic nuclei (AGNs) and in the starburst dominated regions frequently associated to them. Because of the large variety of fine structure lines present in the mid-IR, covering a wide range in ionization/excitation conditions and gas density \\citep[e.g.][]{sm92} mid-IR spectroscopy of the Narrow Line Regions (NLR) in AGNs can add information not available from classical optical spectroscopy, especially when dust extinction is high. Furthermore, the electronic transitions responsible for the infrared emission lines of various elements are less sensitive to uncertainties in temperature than the corresponding optical lines. Moreover the brightest H$_2$ rotational lines, that can be used to quantify the presence and excitation of warm molecular gas, as well as a prominent Polycyclic Aromatic Hydrocarbons \\citep{pule89}, hereafter PAH, feature at 11.25 $\\mu$m are in the mid-IR. The first detailed mid-IR spectroscopic studies of AGNs and Starburst galaxies using multiple ionic transitions of various elements were performed by various authors \\citep[][for a review]{stu02, spi05, ver03, ver05} with the \\textit{Short Wavelength Spectrometer} (SWS) \\citep{deg96} onboard of the \\textit{Infrared Space Observatory} (ISO) \\citep{kes96}. However, the improved sensitivity of the Infrared Spectrometer (IRS) \\citep{hou04} on the \\textit{Spitzer Space Telescope} now enables a detailed investigation of the nature and physical processes in large samples of galactic nuclei from nonstellar (AGN) and stellar (starburst) power sources.  Even at low resolution, the IRS spectra of the first few classical AGNs already showed the diversity of the mid-IR spectral features: silicate absorption and emission, PAH emission and strong fine structure lines \\citep{wed05}.   \\citet{buc06} examined 51 low resolution IRS spectra of 12$\\mu$m selected Seyfert galaxies \\citep{rms93}, exactly the sample we are considering in our study. They report a few major findings: (1) the sample contains a very wide range of continuum types, with no more than about 3 galaxies being closely similar to one another, however principal component analysis applied to their data suggests that the relative contribution of starburst emission may be the dominant cause of variance in the observed spectra; (2) the starburst component in the sample objects does not contribute more than 40\\% of the total IR flux density; (3) Seyfert 1's have higher ratios of infrared to radio emission \\citep[see also][]{rme96}; (4)  the Seyfert 2 galaxies typically show stronger starburst contributions than Seyfert 1's.  \\citet{Gor07} found a strong correlation between the [NeV]14.3$\\mu$m and the [NeIII]15.5$\\mu$m lines in Narrow Line Regions (NLR) of AGNs, spanning 4 orders of magnitude in luminosity. This would imply a very narrow range in ionization parameter (-2.8$<$logU$<$-2.5) for simple constant density photoionization models. \\citet{dud07} discussed the ratio of the doublets of [NeV] and [SIII] in a heterogeneous sample of active galaxies. Finally ULIRG galaxies, more than Seyfert galaxies, have been so far the object of systematic spectroscopic studies with Spitzer IRS \\citep{hig06,arm07,des07,far07}.  In this article we present the first  high resolution IRS spectra of 29 galaxies  from the original list of Seyfert galaxies of the \"12$\\mu$m Galaxy Sample\" (12MGS) \\citep{rms93}, an IRAS-selected all-sky survey flux-limited to 0.22 Jy at 12$\\mu$m.  The sample selection is briefly described in \\S 2, the observations are described in \\S 3, the results are presented in \\S 4, the diagnostics diagrams using line ratios, equivalent widths and other observed quantities are presented and discussed in \\S 5 and a comparison between Seyfert 1's and Seyfert 2's is given in \\S 6. The conclusions are given in \\S  7. ",
        "conclusions": "We have identified a family of IRS observables which quantify the proportion of the total IR emission coming from a Seyfert nucleus, all of which are intercorrelated with each other. The ratios of ionic fine structure lines [NeV]/[NeII] and [OIV]/[NeII] were already proposed to measure the importance of the AGN component. We also see that [NeIII]/[NeII] and OIV or [NeV]/[SiII]35$\\mu$m can be used to quantify the AGN dominance.  It was also known from ISO that the equivalent width of the 6.2$\\mu$m (and 7.7$\\mu$m) PAH features is inversely related to the AGN dominance; we find that the same holds for the equivalent width of the 11.25$\\mu$m PAH feature. We also discovered two additional IRS observables: the equivalent width of [NeII]12.8$\\mu$m and the extendness of the 19$\\mu$m continuum, which also quantify the dominance of the AGN component compared with the emission from the underlying spiral galaxy. All of these observables are correlated with each other, since they are measuring this same astrophysical quantity, and they are all correlated with the hardness of the far-IR continuum, since a more dominant AGN component is already known to be correlated with hotter dust.   There is no clear indication that recent star formation is much more important on average in the Seyfert 2's in our sample, compared with the Seyfert 1's. Although the Seyfert 1's generally tend to have more dominant AGN than the Seyfert 2's, there is a strong overlap between these two classes.  The relatively small difference between the averages of these observables for Seyfert 1's and Seyfert 2's indicates that those two AGN categories are not extremely different in the mid-IR range. Thus for example, the AGN in Seyfert 2's are clearly observable in their 10$\\mu$m continuum emission, and in [NeIII], [OIV] and [NeV], to almost the same degree as in Seyfert 1's, which is not consistent with some proposed torus models for Seyfert 1/2 unification.   Finally, we do not fully confirm the observational claims of \\citet{dud07} : 75 - 80\\%  of our Seyfert 1's and Seyfert 2's show [SIII] or [NeV] densities larger than the low-density limit. In fact, after applying plausible aperture corrections to the [SIII] line ratio, only three Seyfert 2's and one Seyfert 1 have a [SIII] line ratio below the density limit. A similar result is also found for the [NeV] ratio, for which we have two Seyfert 1's (that became four including the upper limits) and 2 Seyfert 2's below the low density limit. These few cases do not in our view require enormous dust extinction values."
    },
    "0710/0710.2617.txt": {
        "abstract": "We consider the propagation of cosmic rays in turbulent magnetic fields. We use the models of magnetohydrodynamic turbulence that were tested in numerical simulations, in which the turbulence is injected on large scale and cascades to small scales. Our attention is focused on the models of the strong turbulence, but we also briefly discuss the effects that the weak turbulence and the slab Alfv\\'enic perturbations can have. The latter are likely to emerge as a result of instabilities with in the cosmic ray fluid itself, e.g., beaming and gyroresonance instabilities of cosmic rays. To describe the interaction of cosmic rays with magnetic perturbations we develop a non-linear formalism that extends the ordinary Quasi-Linear Theory (QLT) that is routinely used for the purpose. This allows us to avoid the usual problem of 90 degree scattering and enable our computation of the mean free path of cosmic rays. We apply the formalism to the cosmic ray propagation in the galactic halo and in the Warm Ionized medium (WIM). In addition, we address the issue of the transport of cosmic rays perpendicular to the mean magnetic field and show that the issue of cosmic ray subdiffusion (i.e., propagation with retracing the trajectories backwards, which slows down the diffusion) is only important for restricted cases when the ambient turbulence is far from what numerical simulations suggest to us. As a result, this work provides formalism that can be applied for calculating cosmic ray propagation in a wide variety of circumstances. ",
        "introduction": "The propagation and acceleration of cosmic rays (CRs) is governed by their interactions with magnetic fields. Astrophysical magnetic fields are turbulent and, therefore, the resonant and non-resonant (e.g., transient time damping, or TTD) interaction of cosmic rays with MHD turbulence is the accepted principal mechanism to scatter and isotropize cosmic rays (see Schlickeiser 2002). In addition, efficient scattering is essential for the acceleration of cosmic rays. For instance, scattering of cosmic rays back into the shock is a vital component of the first order Fermi acceleration (see Longair 1997). At the same time, stochastic acceleration by turbulence is entirely based on scattering.    It is generally accepted that properties of turbulence are vital for the correct description of CR propagation. Historically, the most widely used model is the model composed of slab perturbations and 2D MHD perturbations (see Bieber, Smith, \\& Matthaeus 1988). The advantage of this empirical model is its simplicity and the ability to account for the propagation of CRs in magnetosphere given a proper partition of the energy between the two types of modes.  Numerical simulations (see Cho \\& Vishniac 2000, Maron \\& Goldreich 2001, M\\\"uller \\& Biskamp 2000, Cho, Lazarian \\& Vishniac 2002, Cho \\& Lazarian 2002, 2003, see also book of Biskamp 2003, as well as, Cho, Lazarian \\& Vishniac (2003) and Elmegreen \\& Scalo 2004 for reviews), however, do not show 2D modes, but instead show Alfv\\'enic modes that exhibit scale-dependent anisotropy consistent with predictions in Goldreich \\& Sridhar (1995, henceforth GS95). The approach in the latter work makes productive use of the earlier advances in understanding of MHD turbulence, that can be traced back to Iroshnikov (1963) and Kraichnan (1965) work and the classical work that followed (see Montgometry \\& Turner 1981, Higdon 1984, Montgomery, Brown \\& Matthaeus 1987).  A careful analysis shows that there is no big gap between the Reduced MHD and the GS95 model. In fact, it was shown in Lazarian \\& Vishniac (1999) that the numerical results in Matthaeus et al. (1998) are consistent with GS95 predictions. While particular aspects of the GS95 model, e.g., the particular value of the spectral index, are the subject of controversies\\footnote{To address quantitatively these controversies, we need much better numerical resolution. For instance, hydrodynamic turbulence simulations in Kritsuk et al. (2007) showed that only starting with the 1028$^3$ numerical cubes the bottleneck effects stop dominating the measured spectral slope. While the simulations in Kowal \\& Lazarian (2007) show that the bottleneck is less important for their MHD code, the exact value of the spectral slope is still uncertain. At the same time, the particular theoretically-predicted features of turbulence, for instance, the existence of the scale-dependent anisotropy, can be reliably established, while the exact scaling of this dependence, e.g., like $k_{\\|}\\sim k_{\\bot}^{2/3}$ as in GS95 model or $k_{\\|}\\sim k_{\\bot}^{\\alpha*}$, where $\\alpha*<2/3$ (see Beresnyak \\& Lazarian 2006), also require higher resolution simulations.} (see M\\\"uller \\& Biskamp 2000, Boldyrev 2005, 2006, Beresnyak \\& Lazarian 2006, Gogoberidze 2006, Mason et al. 2007), we think that, at present, GS95 model provides a good starting for developing models of CR scattering. as was done in Chandran (2000), Yan \\& Lazarian (2002, 2004, henceforth YL02, Paper I, respectively), Brunetti \\& Lazarian (2007) etc. In particular, the latter three papers used the decomposition of MHD turbulence over Alfv\\'en, slow and fast modes as in Cho \\& Lazarian (2003) and identify the fast modes as the major source of CR scattering in interstellar and intracluster medium.  However, while the turbulence injected on large scales may correspond to GS95 model and its extensions to compressible medium (Lithwick \\& Goldreich 2001, Cho \\& Lazarian 2002, 2003), one should not disregard the possibilities of generation of additional perturbations by CR themselves. Indeed,  the slab Alfv\\'enic perturbation can be created, e.g., via streaming instability (see Wentzel 1974, Cesarsky 1980) or kinetic gyroresonance instability (see its application for CR transport in Lazarian \\& Beresnyak 2006). These perturbations, that are present for a range of CRs energies (e.g., $\\lesssim 100$GeV for the instabilities above in ISM) owing to non-linear damping arising from ambient turbulence (YL02, Paper I, Farmer \\& Goldreich 2004, Lazarian \\& Beresnyak 2006), should also be incorporated into the comprehensive models of CR propagation and acceleration.  At present, the propagation of the CRs is an advanced theory, which makes use both of analytical studies and numerical simulations. However, these advances have been done within the turbulence paradigm which is being changed by the current research in the field. As we discussed above, instead of the empirical 2D+slab model of turbulence, numerical simulations suggest anisotropic Alfv\\'enic modes (an analog of 2D, but not an exact one, as the anisotropy changes with the scale involved) + fast modes or/and slab modes. This calls for important revisions of the CR propagation, which is the subject of the current paper.    The perturbations of turbulent magnetic field are usually accounted for by direct numerical scattering simulations (Giacalone \\& Jokipii 1999, Mace et al 2000, Qin at al. 2002) or by quasi-linear theory, QLT (see Jokipii 1966, Schlickeiser 2002). The problem with direct numerical simulations of scattering is that the present-day MHD simulations have rather limited inertial range. At the same time, creating synthetic turbulence data which would correspond to scale-dependent anisotropy in respect to the local magnetic field (which corresponds, e.g., to GS95 model) is challenging and has not been practically realized, as far as we know.  While QLT allows easily to treat the CR dynamics in a local magnetic field system of reference, a key assumption in QLT, that the particle's orbit is unperturbed, makes one wonder about the limitations of the approximation. Indeed, while QLT provides simple physical insights into scattering, it is known to have problems. For instance, it fails in treating $90^o$ scattering  (see Jones, Birmingham \\& Kaiser 1973, 1978; V\\\"olk 1973, 1975; Owens 1974; Goldstein 1976; Felice \\& Kulsrud 2001) and perpendicular transport (see K\\'ota \\& Jokipii 2000, Matthaeus et al. 2003).     Indeed, many attempts have been made to improve the QLT and various non-linear theories have been attempted (see Dupree 1966, V\\\"olk 1973, 1975, Jones, Kaiser \\& Birmingham 1973, Goldstein 1976). Currently we observe a surge of interest in finding way to go beyond QLT. Those include recently developed nonlinear guiding center theory (see Matthaeus et al. 2003), weakly nonlinear theory (Shalchi et al. 2004), second-order quasilinear theory (Shalchi 2005a) (see also Shalchi 2006, Webb et al. 2006, Qin 2007, Le Roux \\& Webb 2007). At the same time, most of the analysis so far has been confined to traditional 2D+slab models of MHD turbulence. Following the reasoning above, we think that it is important to extend the work to the non-linear treatment of CR scattering to models MHD turbulence that are supported by numerical simulations.  Propagation of CRs perpendicular to the mean magnetic field is another important problem in which QLT encounters serious difficulties. Compound diffusion, resulting from the convolution of diffusion along the magnetic field line and diffusion of field line perpendicular to mean field direction, has been invoked to discuss transport of cosmic rays in the Milky Way (Getmantsev 1963; Lingenfelter, Ramary \\& Fisk 1971; Allan 1972). The role of compound diffusion in the acceleration of CRs at quasi-perpendicular shocks were investigated by Duffy et al. (1995) and Kirk et al. (1996).  Indeed, the idea of CR transport in the direction perpendicular to the mean magnetic field being dominated by the field line random walk (FLRW, Jokipii 1966, Jokipii \\& Parker 1969, Forman et al. 1974) can be easily justified only in a restricted situation where the turbulence perturbations are small and CRs do not scatter backwards to retrace their trajectories. If the latter is not true, the particle motions are subdiffusive, i.e., the squared distance diffused growing as not as $t$ but as $t^{\\alpha}$, $\\alpha<1$, e.g., $\\alpha=1/2$ (K\\'ota \\& Jokipii 2000, Mace et al 2000, Qin at al. 2002, Shalchi 2005b). If true, this could indicate a substantial shift in the paradigm of CR transport, a shift that surely dwarfs a modification of magnetic turbulence model from the 2D+slab to a more simulation-motivated model that we deal here.   It was also proposed that with substantial transverse structure, {\\it i.e.}, transverse displacement of field lines, perpendicular diffusion is recovered (Qin et al 2002). Is it the case of the MHD turbulence models we deal with?   How realistic is the subdiffusion in the presence of turbulence? The answer for this question apparently depends on the models of turbulence chosen. In this paper we again seek the answer for this question within domain of numerically tested models of MHD turbulence.  There are three major thrusts of the paper:\\\\ I. Extend QLT by taking into account magnetic mirroring effect on large scales.\\\\ II. Describe CR propagation in Milky Way (e.g., calculate CR mean free path for different phases of ISM).\\\\ III. Address the problem of perpendicular transport of CR.    In what follows, we discuss the cosmic ray transport in incompressible turbulence in \\S2. We shall describe the \\S2.1 dispersion of guiding center of CRs and introduce the broadened resonance function to replace the $\\delta$ function in QLT, following which we shall discuss the scattering in strong and weak incompressible turbulence respectively in \\S2.2 and \\S2.3. Then we shall consider the scattering by fast modes in \\S3 and apply the analysis to ISM and get mean free path for different phases of ISM (\\S4). In \\S5, we shall study the perpendicular transport of cosmic rays on both large and small scales. We shall also discuss the applicability of the subdiffusion. Discussion and summary are provided in \\S6 and \\S7 respectively. ",
        "conclusions": "The present paper extends our study in Paper I. As in Paper I we mostly deal with the magnetic perturbations that are part of the large scale turbulent cascade, which is consistent with the Big Power Law in the sky observed via radio-scattering and scintillation technique (Armstrong et al. 1995). In both papers we use the description of the MHD turbulence that follows from numerical simulations.  In Paper I we have the CR scattering calculated in the selected interstellar environments making use of Quasi-Linear Theory (QLT). Because of the limitations of the QLT, we could not provide calculations of the mean free path in Paper I, which limited the utility of the study. In this paper we extended the non-linear approach suggested in V\\\"olk (1975) to treat the scattering, which allows us to calculate the mean free paths that arise from CR interactions with the fast modes. In doing so, similar to Paper I, we take into account damping of the fast modes in the presence of the field wondering induced by the Alfv\\'enic modes.  Our results show that in WIM and halo of our Galaxy, confinement of bulk CRs are mostly due to the compressible modes. We obtain CR mean free paths about a few parsec, consistent with what observations indicate. The major difference with earlier picture is the dependencies of CR transport parameters on the medium properties. The dependence appears as a result of damping of the fast modes. For low energy CRs ($\\lesssim 100$GeV), if dominated by viscous damping, the mean free path of CRs would decrease with energy; with collisionless damping, however, CRs' mean free path stays almost a constant. Field line wandering in general increases the damping of the fast modes and reduces the scattering efficiency of CRs. For higher energy CRs, the influence of damping is limited, and their mean free path increases with energy.  The dependencies on the turbulence damping and therefore the phase properties should have various implications from ratio of secondary to primary elements, diffuse Galactic $\\gamma$ ray emission, to the CMB synchrotron foreground. With precise measurements, the understanding of CMB is now constrained by our understanding of the foreground. The variation of CR index over the Galaxy may paralyze the synchrotron templates. Such variations can be addressed on the basis of the more elaborate CR propagation theory.  The importance of this study goes beyond the interstellar medium. For instance, Brunetti \\& Lazarian (2007), treated acceleration of CRs for plasma in clusters of galaxies appealing to the fast modes, which is the approach to CRs similar to that in Paper I. We believe that the non-linear treatment may be useful for such cases as well. In addition, stochastic acceleration by the MHD turbulence is a promising mechanism for generating high energy particles during solar flares (see, e.g., Petrosian \\& Liu 2004, and references therein). An application to the acceleration of CRs in solar flares will be given in Yan, Lazarian \\& Petrosian (2007, in preparation).   In our treatment we attempted to use the scalings that (a) are consistent with numerical calculations and (b) whose amplitudes we can estimate with a sufficient degree of precision. Therefore our present study does not deal with scattering of CRs by the fast modes on the scales $l>LM_A^2$, $M_A<1$, i.e., on the scales where the Alfv\\'enic turbulence in the weak regime. It was suggested by Chandran (2005) that the weak fast modes at small pitch angles tend to steepen due to the coupling with the Alfv\\'en modes. When the resulting scaling of the fast modes becomes clearer, our approach will be applicable to them.  We have not quantitatively dealt in the present paper with the case of the slab Alfv\\'en modes created by instabilities\\footnote{Streaming instability (see Cesarsky 1980) is an example of such instability. However, the instability is suppressed by both ion-neutral damping (Kulsrud \\& Pierce 1969) and the ambient turbulence (YL02, Farmer \\& Goldreich 2004, Paper I, Lazarian \\& Beresnyak 2006). Another example is the gyroresonance instability discussed in the context of CRs in Lazarian \\& Beresnyak (2006).}. The CR scattering by the perturbations created by those modes may dominate over the gyroresonance with the fast modes, especially for CRs of low energies, i.e., whose gyroresonance with the fast modes is inefficient due to the fast modes damping (see estimates in Lazarian \\& Beresnyak 2006). Progress in quantitative description of the non-linear stages of the instabilities that can create slab modes should enable comprehensive models that include both the fast modes and the slab modes.  In addition, we addressed the issue of perpendicular diffusion, the issue that we have not dealt with in Paper I. We found, that similar to the case of thermal diffusion discussed in Lazarian (2006), the diffusion of CRs depends on the Alfv\\'enic Mach number $M_A$. We found that the suppression of the perpendicular diffusion compared to the parallel one scales as $M_A^4$ for $M_A<1$. Approaching the issue of subdiffusion, we found that it is negligible for CRs in the Alfv\\'enic turbulence."
    },
    "0710/0710.4354_arXiv.txt": {
        "abstract": "We present deep WIYN H$\\alpha$ imaging of the dwarf irregular starburst galaxy NGC 1569, together with WIYN SparsePak spatially-resolved optical spectroscopy of the galactic outflow. This leads on from our previous detailed analyses of the state of the ISM in the central regions of this galaxy. Our deep imaging reveals previously undetected ionized filaments in the outer halo. Through combining these results with our spectroscopy we have been able to re-define the spatial extent of the previously catalogued superbubbles, and derive estimates for their expansion velocities, which we find to be in the range 50--100~\\kms. The implied dynamical ages of $\\lesssim$\\,25~Myr are consistent with the recent star- and cluster-formation histories of the galaxy. Detailed decomposition of the multi-component H$\\alpha$ line has shown that within a distinct region $\\sim$$700\\times 500$~pc in size, roughly centred on the bright super star cluster A, the profile is composed of a bright, narrow (FWHM $\\lesssim$ 70~\\kms) feature with an underlying, broad component (FWHM $\\sim$ 150~\\kms). Applying the conclusions found in our previous work regarding the mechanism through which the broad component is produced, we associate the faint, broad emission with the interaction of the hot, fast-flowing winds from the young star clusters with cool clumps of ISM material. This interaction generates turbulent mixing layers on the surface of the clouds and the evaporation and/or ablation of material into the outflow. Under this interpretation, the extent of the broad component region may indicate an important transition point in the outflow, where ordered expansion begins to dominate over turbulent motion. In this context, we present a multi-wavelength discussion of the evolutionary state of the outflow. ",
        "introduction": "Outflows powered by the collective injection of kinetic energy and momentum from massive stars and supernovae (SNe) can drastically affect the structure and subsequent evolution of galaxies. Thus, a good understanding of the feedback mechanisms between massive stars, star clusters and the ISM is fundamental. In particular, dwarf galaxies are thought to be strongly affected by the effects of feedback since their smaller gravitational potentials mean that supernova-heated gas can escape more easily \\citep{larson74}. Although the ejection of the ISM potentially could have significant consequences on the star-formation rate within these low-mass systems \\citep{dekel86}, more recent work suggests that ejection of hot gas through bubble blow-out may not be as efficient as first thought \\citep{deyoung94, martin98}. It is therefore important to study such systems to understand how gas is removed and what effects this has in the evolution of the galaxy.  NGC 1569 (UGC 3056, Arp 210, VII Zw 16, IRAS 4260+6444) is a nearby \\citep[2.2~Mpc;][]{israel88}, low metalicity \\citep*[0.25~\\Zsun;][]{devost97, kobulnicky97} dwarf irregular galaxy that has recently undergone a period of starburst activity. The most recent burst is thought to have peaked between $\\sim$10--100~Myr ago with an average star-formation rate of $\\sim$0.5~\\Msun~yr$^{-1}$ \\citep{greggio98}. At some point near the end of this event, the two well-known super-star clusters (SSCs) A and B were formed (\\citealt{arp85}; \\citealt*{oconnell94}; \\citealt{demarchi97}), and together with the slightly older cluster 30 \\citep{hunter00, origlia01}, appear to dominate the energetics of the central regions.  H\\one\\ observations of NGC 1569 \\citep{stil98, stil02, muhle05} clearly show morphological and kinematical signatures caused by the starburst. Firstly, the H\\one\\ kinematics are ``strongly disrupted'' within the central 900~pc, with little or no evidence for the disc rotation seen at lager radii. Furthermore, large seemingly tidal structures are seen, including a so-called bridge connecting the galaxy to a low-mass H\\one\\ cloud \\citep[the companion;][]{stil98}, a large H\\one\\ arm extending to the west of the disc, and a very faint filamentary H\\one\\ stream wrapping around the south of the disc \\citep{muhle05}. A `hot spot' (a region of velocity crowding) on the western edge of the disk was found by \\citet{muhle05}, who interpreted it as the impact location of infalling gas from the companion. This provides a compelling explanation of how the starburst event was triggered.  H$\\alpha$ images of this galaxy show an equally chaotic, complex structure to the warm ionized component, exhibiting many filaments, bubbles and loops. Many of these were identified by \\citet*{hunter93} from deep H$\\alpha$ imaging. Later kinematical studies found that these filaments form part of a cellular outflow structure, comprising large-scale expanding superbubbles on the northern and southern sides of the disc (\\citealt*{tomita94}; \\citealt{heckman95, martin98}). By analysing spectra from two long-slits placed parallel to the major and minor axes, \\citet{heckman95} found evidence of shocks in the outer regions of the ionized halo. The ratios of [O\\one]/H$\\alpha$, [S\\two]/H$\\alpha$ and [N\\two]/H$\\alpha$ \\citep[used often to diagnose and trace the conditions within ionized gas;][]{veilleux87, dopita00} were all found to be high in the outer filaments, indicating either an increase in the importance of shocks, or that photoionization becomes less important with distance as the ionizing radiation from the central starburst becomes diluted.  The existence of shocked gas is supported by high-resolution X-ray observations of NGC 1569. \\citet*{martin02} examined the X-ray properties of the outflow with \\textit{Chandra}, and found significant soft (0.3--0.7~keV), diffuse emission coincident with the H$\\alpha$ morphology. Although they found the X-ray colour variations to be inconsistent with a free-streaming wind, they concluded that the X-ray emission probably originates in the halo shock generated by the outflowing gas, possibly from the mixing layers between the shock and the bubble interior.  In order to properly characterise the structure of the outflow, it is essential to study its kinematics. From optical long-slit spectroscopy, \\citet{martin98} found the outflow to be composed of numerous superbubbles. In general, she finds the redshifted component of the split-line profile to be stronger in the south and weaker in the north, suggesting an inclined outflow aspect. This is consistent with more recent X-ray absorption measurements \\citep{martin02} and H\\one\\ observations \\citep[from which an inclination angle of $\\sim$60$^{\\circ}$ was derived][]{stil02}. Although it is unclear whether the two SSCs, A and B, alone are sufficiently powerful to provide enough mechanical energy and ionizing radiation to drive the whole outflow and power the galaxy's diffuse ionized medium \\citep{martin97, martin02}, what is clear from the H$\\alpha$ morphology is that energy is being injected throughout the central starburst zone of the disc from multiple sources.  A detailed look at the spectral line profiles, however, shows that near SSC A, ``distinctly non-Gaussian'' H$\\alpha$ emission can be found, exhibiting weak but very broad wings \\citep{heckman95}. Broad emission line wings have been detected in other nearby starburst galaxies (e.g.~\\citealt{izotov96, homeier99, marlowe95, mendez97}; \\citealt*{sidoli06}; \\citealt{westm07c}). Due to mismatches in spectral and spatial resolution and in the specific environments observed, the nature of the energy source for these broad lines has been contested, and has resulted in the proposal of number of possible explanations. However, a detailed IFU (integral field unit) study of the ionized ISM conditions in four regions within the 200~pc region surrounding SSC A by \\citet[][hereafter \\citetalias{westm07a}]{westm07a} and \\citet[][hereafter \\citetalias{westm07b}]{westm07b} have shed a considerable amount of light on this problem.  By accurately decomposing the emission line profiles across each of the IFU fields, we found the line shape to be, in general, composed of a bright narrow feature (intrinsic FWHM $\\sim$ 50~\\kms) superimposed on a fainter broad component (FWHM $\\sim$ 200--400~\\kms). By mapping out their individual properties, we identified a number of correlations between the line components that allowed us to investigate in detail the state of the ionized ISM. We concluded that the broad underlying component is most likely produced in a turbulent mixing layer \\citep{slavin93, binette99} on the surface of cool gas knots, set up by the impact of the fast-flowing cluster winds \\citep{pittard05}. Our analysis revealed a very complex environment with many overlapping and superimposed components, but surprisingly no evidence for organised bulk motions. We concluded that the four regions sampled are all located well within the wind energy injection zone \\citep{shopbell98} at the very roots of the outflow, and that the collimation processes required to transform the turbulent motions into an organised net outflow forming the large-scale superbubbles must occur between 100--200~pc from the central star clusters.  With this in mind, we have obtained new deep H$\\alpha$ imaging of NGC 1569, together with spatially-resolved spectra of the outer halo, to investigate in detail the morphology and kinematics (including the line profile shapes) of the warm ionized component of the outflow at these large radii. In Section~\\ref{sect:data} we present the observations and in Section~\\ref{sect:maps} we map out the properties of the line components and discuss the results of the spectroscopy. Since a number of SparsePak fibres are coincident with or lie adjacent to the Gemini GMOS/IFU fields presented in Papers I and II, we compare the results obtained with the two instruments in Section~\\ref{sect:comparison}. In Section~\\ref{sect:disc} we discuss the state of the outflow, including the conditions in the inner and outer halo, and we summarise our findings and conclusions in Section~\\ref{sect:concs}. ",
        "conclusions": "\\label{sect:concs} We have presented WIYN MiniMo deep H$\\alpha$ imaging of NGC 1569 covering a field-of-view of $9.5\\times 10.7$~arcmins. The depth and large dynamic range of the observations have enabled the identification of previously undetected faint ionized filaments in the halo and the study of the gas morphology right down into the bright, central regions of the disc. We have also presented WIYN SparsePak ``formatted field unit'' observations covering the outer galactic wind flow of NGC 1569 in four pointings with integration times of 4--4.5 hours per field. The large diameter of the SparsePak fibres makes this instrument ideal for probing the faint ionized gas found in the halos of galaxies. This light-collecting power allowed us to choose a high-resolution spectrograph set-up, enabling us to characterise the line profile shapes of the important nebular diagnostic lines of H$\\alpha$ and [S\\two] to an accuracy limited only by the S/N achieved. We now summarise our main findings.  \\begin{itemize} \\item We find H$\\alpha$ emission out to radii of $\\sim$1.5~kpc from the disc, and detect emission in almost every SparsePak fibre over the combined FoV. The presence of such an extensive system of ionized filaments results from the ongoing starburst that is supplying both mechanical energy to eject material and the Lyman continuum luminosity to keep it photoionized. \\item Through detailed Gaussian line fitting, we find that within a distinct region $\\sim$$700\\times 500$~pc in size, roughly centred on the location of SSC A, the nebular emission line profiles are composed of a bright, narrow (FWHM $\\lesssim$ 70~\\kms) component with an underlying, broad component (FWHM $\\sim$ 150~\\kms). At larger radii, we find two narrow components to the H$\\alpha$ line, each representing one half of a split-line profile. \\item By comparing our results to observations of the central regions directly surrounding SSC A \\citepalias{westm07a, westm07b}, we conclude that the physical mechanisms that give rise to the underlying broad emission seen within this zone must be the same as within the regions directly surrounding SSC A sampled by Gemini GMOS/IFU observations. The broad emission is most likely to result from turbulent mixing layers on the surface of cool gas clumps set up by the impact of the hot, fast-flowing cluster winds, and from evaporation and/or ablation of material from the clumps. \\item The extent of this broad component region, coincident with the point at which we start observing signatures of large-scale bubble expansion, may indicate a transition point where ordered expansion begins to dominate over turbulent motion. Further observations are needed to investigate this in more detail. \\item By combining our deep H$\\alpha$ imaging and spectroscopy, we redefine the spatial extents of superbubbles A and B, and confirm their published \\citep{martin98} expansion velocities (90 and 85~\\kms, respectively). We estimate the dynamical ages of these bubbles to be $\\sim$10--15~Myr. Contrary to what has been previously suggested \\citep{martin98}, we find no kinematic or morphological evidence to suggest that either of these two superbubbles have ruptured and are venting their interiors into the galactic halo. \\item Our data indicate that the halo of NGC 1569 contains only 4 superbubbles. Following the terminology introduced by \\citet{martin98}, we propose that her superbubble F should should encompass the whole south-eastern bubble complex, where the velocity ellipses identified by \\citet{martin98} and used to define shells G and F, are simply one level of a `hierarchy of structure'. Furthermore, we propose that superbubble E should encompass what were previously referred to as shells D and E and the large north-eastern X-ray spur \\citep{martin02}. \\item We derive new measurements of the expansion velocity, $v_{\\rm exp}$ (calculated from the difference in the radial velocities between the two split-line components), for the superbubble complexes E and F of $v_{\\rm exp}\\sim50$~\\kms{} and $\\lesssim$100~\\kms, respectively. Assuming a diameter of $\\sim$1.3~kpc for these two structures implies dynamical ages of $\\lesssim$25~Myr. \\item The derived ages of the supershells are consistent with the recent cluster formation history of NGC 1569 \\citep{anders04}, implying that each shell is associated with a specific star-forming event, such as a young massive star cluster that can provide a large mechanical luminosity from its many type II supernovae. \\item The consistent reversal of strengths between the blue and red components in the northern and southern outflows provides evidence of preferred outflow directions approximately perpendicular to the inclined and flattened H\\one\\ disc \\citep{stil02}. \\item In addition to characterising the H$\\alpha$ line profile, we have also measured [S\\two] derived electron densities and [S\\two]$\\lambda$6717+$\\lambda$6731/H$\\alpha$ line ratios. We find that much of the ionized gas in the broad component region and in the outer-wind regions is at or below the low density limit, as is expected for a rarefied outflow. We find the highest [S\\two]/H$\\alpha$ ratios are associated with the faintest H$\\alpha$ fluxes and the largest galactocentric distances. log([S\\two]/H$\\alpha$) ratios $>$0 are found in the superbubbles E, B and A, indicating that in these regions the gas emission may be significantly shock-excited. \\end{itemize}  In summary, the outflow in NGC~1569 appears to consist of several superbubbles in various phases of development, as noted by \\citet{martin98}. This situation is further complicated by the disturbed state of the H\\one\\ that includes features well out of the main plane of the galaxy (Fig.~\\ref{fig:xray_HI}b). Our data confirm this model and thus indicate that the evolution of the outflow will be determined by the development of the superbubbles. In superbubbles A, B, and F ionized gas arcs seen in H$\\alpha$ and the X-ray morphology suggest that much of the hot gas still is confined, albeit moving at velocities that are comparable to those needed to escape. Thus the NGC~1569 outflow does not currently appear to be in the form of an approximately steady state galactic wind, even though it may eventually lead to mass loss from the system. It therefore differs from the well known M82 outflow, whose outer regions can be modelled by a supersonic galactic wind \\citep[e.g.][]{suchkov94, shopbell98, zirakashvili06}."
    },
    "0710/0710.2003_arXiv.txt": {
        "abstract": "The apparent alignment of the cosmic microwave background multipoles on large scales challenges the standard cosmological model. Scalar field inflation is isotropic and cannot account for the observed alignment. We explore the imprints, a non-standard spinor driven inflation would leave on the cosmic microwave background anisotropies. We show it is natural to expect an anisotropic inflationary expansion of the Universe which has the effect of suppressing the low multipole amplitude of the primordial power spectrum, while at the same time to provide the usual inflationary features. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.0837_arXiv.txt": {
        "abstract": "\\citet{chang06} reported millisecond duration dips in the X-ray intensity of Sco X-1 and attributed them to occultations of the source by small trans-Neptunian objects (TNOs). We have found multiple lines of evidence that these dips are not astronomical in origin, but rather the result of high-energy charged particle events in the {\\it RXTE} PCA detectors. Our analysis of the {\\it RXTE} data indicates that at most 10\\% of the observed dips in Sco X-1 could be due to occultations by TNOs, and, furthermore, we find no positive or supporting evidence for any of them being due to TNOs.  We therefore believe that it is a mistake to conclude that any TNOs have been detected via occultation of Sco X-1. ",
        "introduction": "\\label{sec:intro}  \\citet{chang06} found statistically significant 1-2 millisecond duration dips in the count rate during X-ray observations of the bright X-ray source Sco X-1 carried out with the Proportional Counter Array (PCA) on the {\\it Rossi X-ray Timing Explorer} ({\\it RXTE}) and attributed them to occultations of the source by small objects orbiting the Sun beyond the orbit of Neptune, i.e., trans-Neptunian objects (TNOs).  In all, \\citet{chang06} found some 58 dips in approximately 322 ks of Sco X-1 observations.  Given that the {\\it RXTE} spacecraft moves through the diffraction-widened shadows of any TNOs at a velocity of $\\sim30$ km s$^{-1}$, dips of $\\sim$2 ms duration should correspond to a TNO size of $\\sim$60 m.  If the identification of these dips with occultations by TNOs is correct, the dips would provide extremely valuable information on the number and distribution of solar system objects of $\\sim20$-100 m in size.  We have found evidence that these dips are produced by electronic dead-time as a result of high-energy charged particle events in the RXTE PCA detectors.  Preliminary reports of our results were given by \\citet{tajatel, tajast}; herein we give a more detailed and complete report. ",
        "conclusions": "\\label{sec:disc}  The observed dips have widths and depths that are approximately what one might expect to be produced by occultations by TNOs, even though much wider dips would be detectable in principle (given appropriate depths).  Thus we are obligated to seriously consider the hypothesis that some or all of the observed dips are the product of TNO occultations.  However, close examination of the {\\em RXTE} PCA data reveals six signatures that independently indicate that few and possibly none of the observed dips are due to occultations by TNOs. The signatures are (1) the numbers of SM1-other events during the dips; (2) the numbers of VLE events during the dips; (3) the absence of the expected diffraction sidelobes; (4) the temporal asymmetry of the dips; (5) the almost total lack of dips longer than $\\sim$1 ms; and (6) the lack of correlation between dip duration and depth.  We discuss each of these in turn.  (1) {\\bf SM1-Other Events:} Fig.~\\ref{fig:xraylc} shows that there is, on average for 201 dips, a large excess of SM1-other events at or near the times of the dips. On average, individual dips should show an excess of SM1-other events at the $2.7\\, \\sigma$ level.  In Fig.~\\ref{fig:otherhist} we show a histogram of the number of SM1-other events in the 1/8-s time bins corresponding to the dips expressed in standard deviations from the mean.  The mean value obtained from the histogram is $\\sim2.7\\, \\sigma$, as expected.  If one makes the reasonable assumption that the numbers of SM1-other events should not be affected by true occultations (other than by negligible increases due to reductions in the electronic dead time), then one may estimate the maximum fraction, $f_{\\rm occ}$, of the observed dips that represent genuine occultation events that is consistent with this distribution of SM1-other events.  We constructed the following simple function with which to fit the histogram, and thereby constrain $f_{\\rm occ}$: \\begin{equation} {\\rm probability} = \\frac{1}{\\sqrt{2\\pi}}f_{\\rm occ} e^{-C^2/2} + \\frac{1}{\\sqrt{2\\pi \\sigma_{cr}^2}}(1-f_{\\rm occ})e^{-(C- \\overline{C}_{cr})^2/2\\sigma_{cr}^2} \\end{equation} where $C$ is the number of excess SM1-other counts (in units of standard deviations of the counts per bin in each PCA light curve), and $\\overline{C}_{cr}$ is the mean of $C$ for those dips which are not the results of occultation events and which we take to be $\\approx2.7/(1-f_{\\rm occ})$.  The distribution of the numbers of excess SM1-other counts is wider than what would be expected from a Poisson distribution with a mean equal to the slightly increased (on average) number of SM1-other events per bin; the width of this component of the fitting function is adjusted by means of the parameter $\\sigma_{cr}$. Fits of the function to the histogram in Fig.~\\ref{fig:otherhist} were carried out with $\\chi^2$ fits using both Gaussian and Cash (1979) statistics.  If we neglect the tail of the distribution at high numbers of excess SM1-other events, i.e., at $>9\\sigma$, we obtain formally acceptable fits with values of $f_{\\rm occ}$ in the range 0.0 to 0.12 and values of $\\sigma_{cr}$ in the range 1.85 to 2.35 (based on Gaussian statistics; the limits represent the formal joint 95\\% confidence range). Using Cash statistics, we obtain formally acceptable fits with values of $f_{\\rm occ}$ in the range 0.0 to 0.11 and values of $\\sigma_{cr}$ in the range 1.65 to 2.15 (95\\% confidence). These results indicate that fewer than 11\\% of the 203 dips might be the product of TNO occultations.  \\begin{figure*} \\epsscale{0.85} \\plotone{f7.eps} \\caption{Histogram of the numbers of excess SM1-other events in 1/8-s time bins corresponding to the times of 201 of the 203 dips.  The number of excess SM1-other events is given in terms of the square root of the mean number of SM1-other events per bin away from the time of the dip.  The solid smooth curve represents the best fit with $f_{\\rm occ} = 0$, $\\sigma_{cr} = 2.1$ and the dashed curve represents a formally {\\em un}acceptable fit with $f_{\\rm occ} = 0.24$, $\\sigma_{cr} = 1.8$ (see text). } \\label{fig:otherhist} \\end{figure*}  (2) {\\bf VLE events:} Figure~\\ref{fig:xraylc} also shows that there is an excess of VLE events around the times of the dips.  The difference between the background rate and that in the 1/8-s bin containing the dips is very close to 4 (actually $3.9 \\pm 0.5$) extra VLE events per dip.  The peak in Fig.~\\ref{fig:xraylc} is significant at the $7 \\sigma$ level.  If there is precisely one VLE per operating PCU for each non-TNO dip, then we would expect on average an excess of 4.67 VLEs per non-TNO dip. If only the non-TNO dips contribute to the excess VLE events then there is an upper limit to $f_{\\rm occ}$ that is consistent with the observations.  If we further allow that the statistical mean excess number of VLE events per dip may have been as small as 2.9, then a simple calculation gives the limit $f_{\\rm occ} \\lesssim 0.38$ (95\\% confidence). This limit is weaker than for the SM1-other events, and, furthermore, is compromised by the possibility that more than one VLE event could be produced in a operating PCU in a single cosmic-ray induced dip.  (3) {\\bf Lack of diffraction sidelobes:} In Fig.~\\ref{simprof} we showed an average profile of simulated dips that had been inserted into actual PCA data.  It should be compared to averages of the actual measured dip profiles in Figs.~\\ref{fig_superpos_202} and \\ref{fig_superpos_109}.  The average model dip profile shows a clear bump of $\\sim8$\\% amplitude on either side of the dip due to diffraction, whereas the averages of the actual profiles show no significant evidence for diffraction sidelobes.  Thus, we conclude that the fraction $f_{\\rm occ}$ of legitimate TNO occultations can be no larger than $\\sim$30\\%, otherwise diffraction sidelobes likely would have been detected.  Again, while this is a clear strike against the dips being due to TNOs, the limiting statistically significant constraint that can be set due to the lack of diffraction sidelobes is not as significant as for the SM1-other events.  (4) {\\bf Asymmetry:} A comparison of the simulated with the actual dip profiles (as in [3]) above, clearly shows a marked asymmetry for the real dip events.  This is physically implausible if the dips are the product of occultation events and therefore testifies against a TNO origin for most of the dip events.  We estimate that the statistical significance of the asymmetry is $\\sim 6\\,\\sigma$.  Unfortunately, there is no direct way to use this information to constrain the fraction of legitimate TNO occultations.  The problem is that we do not know, a priori, how large the asymmetry is, on average, for non-TNO dips.  Therefore, we can not tell how `diluted' the non-TNO events are by potentially real ones.  Nonethless, this marked asymmetry is another solid indication that few of the dips are the product of TNO occultations.  (5) {\\bf Lack of dips longer than $\\sim$1 ms:} From Fig.~\\ref{fig_realwd} we can see that all of the dips, except for a single event, have RMS widths $\\sigma < 1.1$ ms.  In Section~\\ref{sec:data} we described a computer simulation of the production, detection, and analysis of dips caused by TNO occultations.  For a relative speed between the {\\it RXTE} satellite and the shadows of the putative TNOs of $v_{\\rm rel} = 25$ km s$^{-1}$ we find that the fraction of recovered simulated dips with $\\sigma > 1.1$ ms is $\\sim$27\\%.  For $v_{\\rm rel} = 35$ km s$^{-1}$, $9$\\% of the dips have $\\sigma > 1.1$ ms.  We estimate that the average relative velocity between {\\it RXTE} and the shadows of any TNOs was not higher than $v_{\\rm rel} \\sim 30$ km s$^{-1}$.  For this speed, $16$\\% of the dips are characterized by $\\sigma > 1.1$ ms.  Therefore, if {\\em all} of the dips are the result of TNO occultations the number of longer-duration dips should be $\\sim$30, whereas the observed number is actually 1.  On the other hand, if only 15\\% of the dips are due to TNO occultations, then we would expect only $\\sim 5$ dips with $\\sigma > 1.1$ ms.  This expected number is marginally statistically consistent, i.e., at $\\sim$5\\% confidence, with the detection of one dip with $\\sigma > 1.1$ ms.  Therefore, we conclude that the lack of longer dips allows an upper limit of 15\\% to be set on the fraction, $f_{\\rm occ}$, of potentially real TNO occultations.  (6) {\\bf Lack of correlation between width and depth:} If the dips were due to TNO occultations of Sco X-1, we would expect a strong correlation between the widths of the dips and their depths.  This results from the fact that diffraction produces shallow occultations for the smaller size occulters, while it produces deeper more geometric-shadowing-like occultations for the larger occulters.  As can be seen from the distribution of dip widths vs. depths in Fig.~\\ref{fig_realwd}, there is no such correlation, with almost all of the dips confined to a narrow range of widths (between 0.4 and 0.8 ms) and depths that range all the way from 45\\% to nearly 100\\%. Thus, the fact that the dips we detect include a significant number, i.e., $\\sim$20\\%, that are both narrow ($\\sigma < 0.7$ ms) and deep (minimum normalized count rate below 0.2) whereas only $\\sim2$\\% of the `detected' simulated dips (for $v_{\\rm rel} \\sim 30$ km s$^{-1}$) are this narrow and deep, indicates that $\\lesssim10$\\% of the dips might be due to TNO occultations.  Given the effects of statistical fluctuations on the observed number of narrow deep dips and the fact that the simulation is based upon somewhat uncertain parameters, it is more reasonable to use these numbers to set an upper limit of $\\sim20$\\%.  Summarizing the results from approaches (1) through (6) above, we find limits on the fraction of valid TNOs to be $f_{\\rm occ} <11\\%$, $< 38\\%$, $< 30\\%$, $< Q\\%$, $<15\\%$, $<20\\%$, respectively, where ``Q'' denotes that a formal limit could not be set, but the approach provides an important independent indication that the dips are, for the most part, not the result of TNO occultations.  We believe that the combined upper limit on $f_{\\rm occ}$ due to the joint application of all six approaches is simply the minimum value achieved by the most sensitive of these, i.e., the constraints cannot be combined.  The reason, in short, is that the effects we explore serve only to statistically limit the number of events which could be due to TNOs rather than to identify specific qualifying events.  Therefore, our final limit is simply $f_{\\rm occ} \\lesssim 10\\%$.  One might argue, as did \\citet{chang07}, that since $\\sim10$\\% of the observed dips cannot be formally eliminated as being due to TNOs, they serve as viable potential candidates for TNO detections.  However, we argue that if 90\\% of the dips can be securely eliminated as TNO occultations, and there are six different and independent indicators that point in the direction of a common cause due to cosmic ray interactions in the detector, then it is most plausible that {\\em all} of the dips have this common origin.  While our results cast serious doubt on whether any true occultation events have been detected, one cannot yet conclude with a high degree of confidence that no such events have been detected. Further investigations of the dip phenomenon and its possible causes would be of interest.  We are working to obtain a new measurement of, or upper limit on, the rate of occurrence of occultations of Sco X-1 by analyzing the data that are being obtained in a new series of {\\it RXTE} observations of Sco X-1 with high-time-resolution information on VLE events."
    },
    "0710/0710.3252.txt": {
        "abstract": "% The role of optical Fe\\,{\\sc iii} absorption lines in B-type stars as iron abundance diagnostics is considered. To date, ultraviolet Fe lines have been widely used in B-type stars, although line blending can severely hinder their diagnostic power. Using optical spectra, covering a wavelength range $\\sim$ 3560 -- 9200 \\AA, a sample of Galactic B-type main-sequence and supergiant stars of spectral types B0.5 to B7 are investigated. A comparison of the observed Fe\\,{\\sc iii} spectra of supergiants, and those predicted from the model atmosphere codes {\\sc tlusty} (plane-parallel, non-LTE), with spectra generated using {\\sc synspec} (LTE), and {\\sc cmfgen} (spherical, non-LTE), reveal that non-LTE effects appear small. In addition, a sample of main-sequence and supergiant objects, observed with FEROS, reveal LTE abundance estimates consistent with the Galactic environment and previous optical studies. Based on the present study, we list a number of Fe\\,{\\sc iii} transitions which we recommend for estimating the iron abundance from early B-type stellar spectra. ",
        "introduction": "%-----------------------  Iron lines dominate the spectra of many astrophysical objects, such as novae \\citep{{mck97},{hat07}}, photoionized H II regions \\citep{{rub97},{rod02},{est02}} and active galactic nuclei \\citep*{{sig03},{sig04},{zha06}}. The atomic processes for iron and other iron-group ions have been the subject of numerous investigations, for example the IRON Project \\citep{hum93} which considers applications in astrophysical and laboratory plasmas \\citep{pra96a}. Absorption lines of iron provide important metallicity diagnostics for both stars and galaxies, and also play a key role in investigating star formation histories through element ratios, such as [$\\alpha$/Fe] \\citep{gil98}. However, there is a lack of robust Fe abundance determinations in external galaxies,  e.g. the Magellanic Clouds (see for example, \\citealt*{{rol02},{tru02},{tru07}} and \\citealt{mok07}), due to the complex Grotian diagrams for Fe, the low metallicity environment of the Magellanic Clouds and the reliability of the currently available atomic data.  Early B-type stars are important for studying the chemical composition of our own and other galaxies \\citep{{kil92},{duf98}}. In the optical spectra of B-type stars, iron lines due to a number of ionization stages are observed (see for example, \\citealt*{{gie92},{len93},{sma97},{mor06}}). Absorption features arising from Fe\\,{\\sc iii} are primarily detected \\citep{har70}, with Fe\\,{\\sc ii} also found in later B-type stars \\citep{pin93}. Fe\\,{\\sc iv} lines are not expected in the optical spectra of B-type stars due to their intrinsic weakness in this temperature range.   The optical Fe\\,{\\sc iii} line spectrum has not been widely employed to determine abundances, as it has often been believed to be too weak to provide reliable measurements \\citep{ken94a}. However, it has been used for chemical composition studies of several bright, narrow-lined main-sequence B-type stars, such as $\\zeta$\\,Cas, $\\gamma$\\,Peg, $\\iota$\\,Her, $\\tau$\\,Sco and $\\lambda$\\,Lep \\citep{{sni69},{pet70},{har70},{pet76},{pet85}}. These Galactic objects have sufficiently high Fe content, coupled with narrow metal absorption lines (due to their low projected rotational velocities), so that, even with relatively poor quality optical spectra, Fe\\,{\\sc iii} features can be detected. However, the quality of the available spectra did not allow the study of Fe\\,{\\sc iii} lines in objects of lower metallicity. Therefore, more recently the optical wavelength range has been largely overlooked in favour of the ultraviolet domain \\citep{{swi76},{pet85},{dix98}}, and in particular the very rich spectral region around 1900 \\AA\\ \\citep*{{tho74},{heb83},{ken94a},{gri96},{moe98}}. On the other hand, due to the high density of absorption features and resultant blending, continuum placement in the ultraviolet is difficult, and hence significant errors may be present in the derived abundances \\citep{moe98}.  Due to instrumental advances in more recent years, it has become routine to obtain high resolution and signal-to-noise (S/N) spectra, and the use of the optical Fe\\,{\\sc iii} lines as a diagnostic has been revisited (e.g. \\citealt{{cun94},{kil94}}). A number of main-sequence objects have subsequently been re-examined, for example $\\gamma$\\,Peg and $\\iota$\\,Her \\citep*{{pin93},{zon98}}, producing results consistent with the earlier optical analyses. In addition, the higher S/N ratios of the available spectra have made it possible to detect the weak optical Fe\\,{\\sc iii} lines in a number of lower metallicity objects, such as AV\\,304 in the Small Magellanic Cloud \\citep{rol03}, and the globular cluster post-AGB stars ZNG-1 in M\\,10 \\citep{moo04}, Barnard\\,29 in M\\,13 and ROA\\,5701 in $\\omega$\\,Cen \\citep{tho07}.   A number of B-type stars have been studied using both ultraviolet and optical spectra, but the corresponding abundance estimates are generally in poor agreement. Estimates from the ultraviolet spectra are consistently lower than those from the optical, for example in the post-AGB stars Barnard\\,29, ROA\\,5701 \\citep{tho07}, BD\\,+33${^\\circ}$2642 \\citep*{nap94} and HD\\,177566 \\citep{{ken94a},{ken94b}}. Young Galactic objects  including $\\gamma$\\,Peg \\citep{moe98} and $\\iota$\\,Her \\citep{gri96}, plus stars in known metallicity environments such as the Magellanic Clouds \\citep{duf07}, also display similar discrepancies.   Here, high resolution spectra for a number of narrow-lined Galactic B-type main-sequence and supergiant stars, covering a range of spectral types, are analysed using LTE and non-LTE model atmosphere techniques. Details on the observations, models and Fe\\,{\\sc iii} line selection can be found in Sections \\ref{sec_obs} and \\ref{sec_data}. An LTE approximation is considered, and the reliability of such an assumption, along with an assessment of the individual Fe\\,{\\sc iii} lines, is discussed in Section \\ref{sec_discuss}.   %----------------------- ",
        "conclusions": "\\label{sec_conc} %-----------------------  Due to their intrinsic weakness, Fe\\,{\\sc iii} absorption lines have not been widely considered for use in chemical composition studies. Instead, the very rich ultraviolet spectral region has been favored. However, the results found in this paper suggest that the optical region can provide consistent results. The stars here display abundance estimates that agree with the Galactic metallicity, and are consistent with previous studies using optical spectra, where available. By contrast, previous determinations from ultraviolet spectra have followed the trends observed in other studies (for example \\citealt{{tho07},{duf07}}), providing lower abundance estimates, in these cases by approximately 0.5 dex.   Although our study has concentrated on B-type stars found in the Milky Way, the optical Fe\\,{\\sc iii} lines examined here can be applied to studies of B-type stars in other galaxies, provided suitable S/N spectra are employed. For example, \\citet{tru02} analysed a sample of B-type supergiants in M\\,31, obtaining an Fe abundance from the Fe\\,{\\sc iii} line at 4419 \\AA. \\citet{rol02} analysed a sample of OB-type main-sequence stars from the LMC, finding an Fe abundance for one object (PS\\,34-16), while \\citet{rol03} investigated a B-type dwarf from the SMC (AV\\,304), finding agreement with other giant stars. More recently, \\citet{tru07} used similar methods to those detailed here and obtained Fe abundances for B-type stars in the Galaxy, LMC and SMC, using the two Fe\\,{\\sc iii} lines at 4419 and 4430 \\AA. The results found were consistent with the present accepted metallicities of these systems. These studies indicate that the optical Fe\\,{\\sc iii} lines can provide reliable abundance indicators in different galaxies.  Our comparison of stars analysed using the model atmosphere codes {\\sc cmfgen} and {\\sc tlusty} generally shows little difference in the abundance estimates, indicating that the different physical assumptions, in particular non-LTE effects, are small for this species. Therefore, the results suggest that the optical Fe\\,{\\sc iii} absorption line spectrum may be used with confidence in chemical composition studies, and an LTE analysis provides reliable results.  The atomic data of \\citet{nah96}, employed both in this paper and by \\citet{cro06}, appear to provide appropriate abundances, although there are some features, such as those at 4005 and 4273 \\AA, whose log\\,{\\it gf} values may be incorrect. Comparing the values in Table \\ref{tab_atdata} shows that, for some features, e.g. the 4166.88, 4419.60 and 5272.90 \\AA\\ lines, there are large differences between the atomic data from the Kurucz database and \\citet{nah96}. Further work is required to refine the atomic data for this species.   \\begin{table} \\caption{Recommended Fe\\,{\\sc iii} lines for use as abundance diagnostics.} \\label{tab_touse} \\begin{tabular}{@{}lcccc} \\hline Line\t& Spectral Type\t\t&Line\t& Spectral Type\t\\\\ (\\AA)\t& Range\t\t\t&(\\AA)\t& Range\t\t\\\\ \\hline 4419\t& B0.5--B7\t\t&5156\t& B0.5--B7 \t\\\\ 4431\t& B1--B7\t\t&5272\t& B1.5--B5 \t\\\\ 5063\t& B1.5--B7\t\t&5282 \t& B2--B7 \t\\\\ 5086\t& B0.5--B7\t\t&5299 \t& B1.5--B4\t\\\\ 5127\t& B0.5--B7\t\t&5302\t& B1.5--B4\t\\\\ \\hline \\end{tabular} \\end{table}   In Table \\ref{tab_touse} we list recommended Fe\\,{\\sc iii} lines which we believe, based on the present study, will provide reliable diagnostics for the iron abundance in early B-type stars. Lines have been selected based on the following criteria:  \\begin{itemize}  \\item Relatively strong, isolated feature, free from known blends.  \\item Yields an abundance estimate within $\\pm$ 0.2 dex of the mean iron abundance for all well observed lines.  \\end{itemize}  \\begin{figure*} \\includegraphics[angle=0,width=0.8\\textwidth]{figure3_a.eps} \\caption{Observed Fe\\,{\\sc iii} lines in HD\\,79447 as listed in Table \\ref{tab_touse}, including the 5193 \\AA\\ line. Overplotted are theoretical fits (smooth line) to the Fe\\,{\\sc iii} lines, calculated using the average abundance estimate for the star of 7.63 dex.} \\label{fig_felines} \\end{figure*}   The range of spectral types over which it is advisable to use the lines as an abundance diagnostic is also listed in Table \\ref{tab_touse}. This spectral type range is not the same for all lines, generally due to the fact that some of the weaker transitions are only detected over a restricted span of spectral types. Observations for all of the lines in this Table are shown in Fig. \\ref{fig_felines} for HD\\,79447. The line at 5193 \\AA\\ has been included in the Figure, but not the Table, because its derived abundance is more than 0.2 dex larger than the mean value for three of the stars studied here, namely HD\\,108002 (B1), HD\\,142768 (B1.5) and HD\\,53138 (B3). However, the feature, along with others observed (see Tables \\ref{tab_tcab}, \\ref{tab_sab} and \\ref{tab_mab}) may be suitable for diagnostic use, depending on the quality of the spectra used. We note that the stars included in this study have been selected due to being narrow lined and are typical of their individual luminosity classes. However, if objects with larger $\\upsilon$\\,sin{\\it i} values were employed, blending between Fe lines and other metal features may occur, thus, care should be taken when using objects with larger projected rotational velocities."
    },
    "0710/0710.2145_arXiv.txt": {
        "abstract": "\\footnote{To appear in proceedings of the Puerto Vallarta Conference on ``New Quests in Stellar Astrophysics II: Ultraviolet Properties of Evolved Stellar Populations'' eds. M. Chavez, E. Bertone, D. Rosa-Gonzalez \\& L. H. Rodriguez-Merino, Springer, ASSP series.} Recent HST/ACS images of M82 covering the entire galaxy have been used to detect star clusters. The galaxy is known to contain a young population (age $<10^7$~yr) in its starburst nucleus, surrounded by a post-starburst disk of age $<10^9$~yr. We detect more than 650 star clusters in this galaxy, nearly 400 of them in the post-starburst disk. These data have been used to derive the luminosity, mass and size functions separately for the young nuclear, and intermediate-age disk clusters. In this contribution, we discuss the evolutionary status of these clusters, especially, on the chances of some of these clusters surviving to become old globular clusters. ",
        "introduction": "\\label{sec:1}  Super star clusters (SSCs) and globular clusters (GCs) represent the youngest and the oldest stellar aggregates known in the Universe. The environments in which these two kinds of clusters are found are vastly different --- SSCs are found in violent star-forming regions, whereas GCs are found in the halos of galaxies. Yet, the similarity in their compactness and mass, is a reason compelling enough to think of an evolutionary connection between them. The growing popularity of the hierarchical model of galaxy formation in the years following the discovery of SSCs, and the possibility of observing the epochs of galaxy (and GC) formation at high redshifts, have also generated interest in looking for a common origin for these two seemingly different classes of clusters.  In order to investigate the relation between the two types of clusters, it is important to analyze the survival of SSCs for a Hubble time. Star clusters are vulnerable to a variety of disruption processes that operate on three different timescales \\citep[see][for more details] {Fal01, Mai04, deG07}. On short timescales ($t\\sim10^7$~yr), the exploding supernovae and the resulting superwinds are responsible for cluster expansion and disruption, a process popularly dubbed as infant mortality. On intermediate timescales ($10^7<t<{\\rm few} \\times 10^8$~yr), the mass-loss from evolving stars leads to the disruption of the clusters. On even longer timescales ($t> {\\rm few} \\times 10^8$~yr), stellar dynamical processes, especially evaporation due to two-body scattering, and tidal effects on a cluster as it orbits around the galaxy, known as gravitational shocks, come into play in the removal of stellar mass from clusters. The GCs represent those objects that have survived all these processes, whereas young SSCs are just experiencing them. Intermediate age SSCs are the ideal objects to investigate the influence of disruption processes on the survival of star clusters. Almost all the star formation in the disk of M82 took place in a violent disk-wide burst around 100--500~Myr ago, following the interaction of M82 with the members of M81 group \\citep{May06}. Cluster formation is known to be efficient during the burst phase of star formation \\citep{Bas05}, and hence we expect large number of clusters of intermediate age ($\\sim100$~Myr) in its disk. Hence, M82 offers an excellent opportunity to assess the evolutionary effects on the survival of star clusters, and to look for a possible evolutionary connection between the SSCs and GCs. ",
        "conclusions": "Luminosity and Mass functions of star clusters in M82 follow power-law functions, with the power law index showing a tendency for flattening of the profile with age. In other words, there is a deficiency of low-mass clusters among the older clusters. We also find the mean size of the older clusters to be smaller as compared to the younger clusters for masses $<10^5$~\\msun. These two results together imply the selective destruction of loose clusters. The tidal forces experienced by the clusters as they orbit around the galaxy lead to exactly such a destruction process. If this process continues in M82, the LF of surviving clusters can mimic the presently observed LF of the Galactic GCs, provided the clusters move around the galaxy in highly elliptical orbits, with perigalactic distance as small as 350~pc. The resulting LF contains 85 clusters with the function peaking at the same luminosity as for the Galactic GCs at 5~Gyr age, and fainter by $\\sim$0.5~mag at 10~Gyr. On the other hand, if the clusters move in nearly circular orbits, the LF will retain the power-law form, with the number of surviving clusters even higher.  \\vspace*{0.3cm} This work is partly supported by CONACyT (Mexico) research grants 42609-F and 49942-F. We would like to thank the Hubble Heritage Team at the Space Telescope Science Institute for making the reduced fits files available to us."
    },
    "0710/0710.0189_arXiv.txt": {
        "abstract": "I review the evolutionary connection between low-mass X-ray binaries (LMXBs) and pulsars with binary companions (bPSRs) from a stellar binary evolution perspective.  I focus on the evolution of stellar binaries with end-states consisting of a pulsar with a low-mass ($<1.0 \\msun$) companion, starting at the point the companion's progenitor first initiates mass transfer onto the neutron star.  Whether this mass transfer is stable and the physics driving ongoing mass transfer partitions the phase space of the companions's initial mass and initial orbital period into five regions. The qualitative nature of the mass-transfer process and the binary's final end-state differ between systems in each region; four of these regions each produce a particular class of LMXBs.  I compare the theoretical expectations to the populations of galactic field LMXBs with companion-mass constraints and field bPSRs.  I show that the population of accreting millisecond pulsars are all identified with only two of the four LMXB classes and that these systems do not have readily identifiable progeny in the bPSR population.  I discuss which sub-populations of bPSRs can be explained by binary evolution theory and those that currently are not.  Finally I discuss some outstanding questions in this field. ",
        "introduction": "Since the discovery of the class prototype \\citep{backer82}, there has been a posited evolutionary connection between millisecond radio pulsars (MSPs) and low-mass X-ray binaries (LMXBs)---mass transferring binaries with a neutron star (NS) accretor and donor companion with a mass  $M_2 \\lesssim 1 \\msun$ \\citep{alpar82}.  The central idea behind this connection is that LMXBs can provide the long-lived phase  ($\\sim \\mathrm{Gyr}$) of moderate mass transfer rates ($\\Mdot \\lesssim \\Mdot_{\\mathrm{Edd}} \\approx 10^{-8} \\msun\\,\\mathrm{yr}^{-1}$, where $ \\Mdot_{\\mathrm{Edd}}$ is the Eddington mass-transfer rate) thought necessary to spin-up the NS to spin periods of $\\Pspin < 10 \\mathrm{ms}$ as observed in the MSP population.  \\begin{figure} \\includegraphics[height=.425\\textheight]{proc_eps_sys_bPSRs_lmxbs.eps} \\caption{The galactic field population of radio pulsars with binary companions (filled circles) in the $M_2$-$\\Porb$ plane. The data are from the ATNF catalog \\citep{manchester05}. Plotted are each system's minimum $M_2$, with horizontal lines extending to the system's median $M_2$ (assuming $i$ is randomly distributed). Filled  circle size indicates each pulsar's $\\Pspin$ (see plot legend). Symbols circumscribing a filled circle indicates the binary's eccentricity (again see plot legend).  For comparison, LMXBs with independent $M_2$ estimates are plotted with open stars and filled triangles, the latter indicating the minimum $M_2$ for \\emph{accreting MSP} systems.  For the open stars, the dash-dotted lines indicate the \\emph{total} estimated $M_2$ range in each system. For the accreting MSPs, the dash-dotted lines extend to the median $M_2$.  This sample of LMXB systems represents a union of systems with $M_2$ estimates in the \\citet{ritter98} catalog and a targeted literature search on LMXB systems where $\\Pspin$ has been determined.  Thus, this plot likely does not present in total the current census of LXMBs with $M_2$ estimates. } \\label{fig:bPSRs_sys} \\end{figure}  Observational evidence for millisecond variability in LMXBs has been growing steadily in the form of detections of kilohertz quasi-periodic oscillations (kHz QPOs), X-ray burst oscillations, and accretion-powered oscillations \\citep[see][]{chakrabarty05}. In terms of support for the LMXB-MSP connection, pride of place has been given to the accretion-powered pulsations systems (also known as accreting MSPs) since the pulsation period in these systems are identifiable directly with $\\Pspin$ \\citep{chakrabarty05}.   However, recent work has also shown that $\\Pspin$ is of order the oscillation periods in the kHz QPO and X-ray burst oscillation sources \\citep{chakrabarty03,strohmayer03,wijnands03,linares05}, establishing that these systems also harbor a rapidly rotating NS.  Thus, it is now well established that NSs in LMXBs can be spinning rapidly enough to produce radio MSPs once the LMXB phase ends.  However, understanding fully the evolutionary connection between LMXBs and MSPs requires not only explaining the NS's spin evolution but also accounting for other properties seen in the radio MSP population.  In particular, this includes the distribution in orbital period, $\\Porb$ of MSPs that retain a remnant binary companion, how this remnant's $M_2$ correlates with $\\Porb$, the distribution of binary eccentricity, $e$, and the production of \\emph{isolated} MSPs.  Indeed, the ideal test of accretion-torque theory would be accomplished by understanding the evolution in the LMXB-phase well enough to correlate final $\\Pspin$ with these other quantities.  This is a rather ambitious goal since the $\\Pspin$ evolution depends not only on the secular $\\Mdot$ evolution but also on the efficiency with which matter accretes onto the NS, whether accretion onto the NS occurs sporadically due to disk instabilities, how $\\Pspin$ evolves during mass-transfer outbursts and when unstable disks are in quiescent phases, and how each LMXB transitions into an MSP system.  Turning from where one would like to be to where we are now, my goal for this contribution is to approach the LMXB-MSP connection from the vantage point of stellar binary evolution theory.  To do so, I will expand the view somewhat and review our understanding of NS-main sequence (MS) binaries whose evolutionary endpoints are NSs with a low-mass ($M_2 \\lesssim 1.0 \\msun$) binary companion. In doing so, the focus of the discussion will shift from solely the MSP population to making connections between NS-MS binaries and various populations of radio pulsars in binaries, bPSRs (one would like to simply say ``binary pulsars'' here, but the discovery of \\emph{the} Binary Pulsar \\citep{burgay03} has lead to this term often causing confusion; I'll leave it to the reader to decide whether the scientific windfall from this system compensates sufficiently for necessitating such unwieldy terminology).  The starting point for this discussion is Figure \\ref{fig:bPSRs_sys}, which shows the location of galactic field bPSRs in the $\\Porb$-$M_2$ plane by the filled circles \\citep{manchester05}. Only field sources are included so as to compare theory to a sample of bPSRs whose properties have not been influenced by dynamic interactions.  The filled circles indicate each bPSR's minimum-$M_2$; the horizontal lines extend to each system's median $M_2$ (corresponding to a binary inclination of $i = 60^\\circ$).  Spin period and $e$ are encoded via filled circle size and circumscribed symbols, respectively. Filled triangles show the same information for accreting MSP systems \\citep{chakrabarty98,galloway02,galloway05,kaaret06,krimm07,markwardt02,markwardt03}.  Open stars indicate other field LMXB systems with $M_2$ determinations \\citep{bhattachar06,casares06,cominsky89,cornelisse07,finger96,heinz01,hinkle06,jonker01,parmer86,pearson06,reynolds97}; for these latter systems, the dashed-dotted horizontal lines show the estimated range of $M_2$ (not its median value).  In the following, I'll compare the predictions of binary evolution theory for how systems evolve in this $\\Porb$-$M_2$ plane to the location of systems in Fig. \\ref{fig:bPSRs_sys}. While this will allow tentative positive identifications of the evolutionary connections discussed above,  it will also serve to highlight sub-populations of bPSRs whose formation is \\emph{not} currently explained by binary evolution theory. I'll discuss the evolution of NS-MS binaries, focusing on how initial conditions lead to four different classes of X-ray binaries.  I'll compare between this theory and the observations, pointing out where agreement between the two is better and worse. Finally, I'll close by discussing several open questions related to the evolution of LMXBs and bPSR formation.  For other discussions and reviews of the bPSR population see \\citep{lorimer05} and \\citep{phinney94}. Also see \\citep{flamb05} for a conference proceeding that reviews NS spin evolution under accretion. ",
        "conclusions": ""
    },
    "0710/0710.5881_arXiv.txt": {
        "abstract": "We report observations of a radio burst that occurred on the flare star AD~Leonis over a frequency range of 1120-1620~MHz ($\\lambda\\approx$18--27~cm). These observations, made by the 305~m telescope of the Arecibo Observatory, are unique in providing the highest time resolution (1~ms) and broadest spectral coverage ($\\Delta \\nu/\\nu=0.36$) of a stellar radio burst yet obtained. The burst was observed on 2005 April 9. It produced a peak flux density of $\\sim 500$ mJy and it was essentially 100\\% right-circularly polarized. The dynamic spectrum shows a rich variety of structure: patchy emission, diffuse bands, and narrowband, fast-drift striae. Focusing our attention on the fast-drift striae, we consider the possible role of dispersion and find that it requires rather special conditions in the source to be a significant factor. We suggest that the emission may be due to the cyclotron maser instability, a mechanism known to occur in planetary magnetospheres. We briefly explore possible implications of this possibility. ",
        "introduction": "The use of radio dynamic spectra has played a central role in identifying and clarifying the physical mechanisms at work in the solar corona \\citep[see][for reviews]{1985srph}.  The application of similar techniques to active stars has long been an important goal, but it has been hampered by limitations in available instrumentation. Past studies of radio emission from M dwarf flare stars led to the discovery of extreme stellar radio bursts, characterized by close to 100\\% circularly polarized emission with brightness temperatures in excess of 10$^{14}$K and durations less than a few tens of milliseconds \\citep{gudel1989,bastian1990}. However, these spectroscopic investigations of the coherent radio bursts on flare stars have typically been limited by relatively long integration times \\citep{bb1987,gudel1989} and/or limited frequency bandwidth ratios $\\Delta\\nu/\\nu$  \\citep[usually just a few percent; e.g.,][]{bastian1990, abada1997a}. The necessary combination of high time resolution and a large frequency bandwidth ratio has only been available infrequently \\citep{stepanov2001, zaitsev2004}, precluding measurements of key parameters such as the intrinsic frequency bandwidth or frequency drift rate of the radio bursts, making the interpretation of these puzzling events difficult.  It is only with the recent advent of radio spectrometers capable of supporting both a large bandwidth ratio and high time resolution simultaneously that progress in understanding the physics of radio bursts in the coronas of other stars becomes possible.  AD~Leonis, a young disk star at a distance of 4.9 pc from the Sun, is one of the most active flare stars known, producing intense, quasi-steady chromospheric and coronal emissions \\citep{hawleyetal2003, hunschetal1999, jackson1989} seen at UV, X-ray, and radio wavelengths. The star is also highly variable, producing flares from radio to X-ray wavelengths \\citep[e.g.,][]{bastian1990, hawleypettersen1991, hawleyetal2003, hawleyetal1995, favataetal2000}. Its propensity for frequent and extreme radio bursts \\citep[with intensities peaking at $>$ 500 times the quiescent radio luminosity of 5.5$\\times$10$^{13}$ erg s$^{-1}$ Hz$^{-1}$;][]{jackson1989} makes it a frequent target for radio investigations of stellar flares. In a previous paper \\citep[][hereafter, Paper I]{ob2006} we described the initiation of a pilot program to observe active M dwarfs with the Arecibo Observatory's Wideband Arecibo Pulsar Processor (WAPP), and first results from that program. Here we describe the next phase, which increased the time resolution by a factor of 10 to 1~ms. ",
        "conclusions": "We have described observations of a unique set of stellar radio bursts, which take advantage of the wide bandwidth and high time resolution capabilities of the Wideband Arecibo Pulsar Processor at the Arecibo Observatory. These ultra-high time resolution observations reveal phenomena that differ from those previously described using a similar observational setup, pointing out the complexity and diversity of processes likely occurring in stellar coronal plasmas.  Whereas in Paper I we concluded that a plasma emission process appeared to be producing the two types of radio bursts observed in June 2003, in the current paper we prefer a different explanation, a cyclotron maser instability, for the fast-drift striae observed in April 2005.  While all sets of phenomena show drifting structures of highly circularly polarized radiation, key discriminants between them are the durations and bandwidths of spectral features, as well as the magnitude and sign of the drift rates.  In Paper I and here, we have demonstrated that the analysis of dynamic spectra of stellar radio bursts provide observational constraints which can be used as a measure to gauge the likelihood that a particular emission process is operative. Extensions of the current observational setup can look for dynamics at even higher time resolution, search for harmonic emissions over larger frequency bandwidths, expand the observational program to other dMe flare stars, and search for high time resolution behavior on other classes of active stars. Given the complexity of solar radio emissions at meter wavelengths compared with the already rich variety of decimetric phenomena, the observational results presented here for the dMe flare star AD Leo suggest that the next generation of radio instrumentation, particularly at metric wavelengths, promises to reveal a wealth of new phenomena which can diagnose plasma processes occurring in stellar coronae. As highly circularly polarized radio emission appears to be a common phenomenon on active stars, these spectacular radio bursts on M dwarf flare stars apparently represent the tip of the iceberg of stellar coronal plasma physics soon to be available for study."
    },
    "0710/0710.5286_arXiv.txt": {
        "abstract": "We present new observations of the strongly-barred galaxy NGC~1365, including new photometric images and Fabry-Perot spectroscopy, as well as a detailed re-analysis of the neutral hydrogen observations from the VLA archive. We find the galaxy to be at once remarkably bi-symmetric in its I-band light distribution and strongly asymmetric in the distribution of dust and in the kinematics of the gas in the bar region. The velocity field mapped in the \\ha\\ line reveals bright HII regions with velocities that differ by 60 to $80\\;$\\kms\\ from that of the surrounding gas, which may be due to remnants of infalling material.  We have attempted hydrodynamic simulations of the bar flow to estimate the separate disk and halo masses, using two different dark matter halo models and covering a wide range of mass-to-light ratios ($\\Upsilon$) and bar pattern speeds ($\\Omega_p$).  None of our models provides a compelling fit to the data, but they seem most nearly consistent with a fast bar, corotation at $\\sim1.2r_B$, and $\\Upsilon_ I \\simeq 2.0 \\pm 1.0$, implying a massive, but not fully maximal, disk.  The fitted dark halos are unusually concentrated, a requirement driven by the declining outer rotation curve. ",
        "introduction": "The centrifugal balance of the circular flow pattern in a near-axisymmetric spiral galaxy yields a direct estimate of the central gravitational attraction as a function of radius.  However, the division of the mass giving rise to that central attraction into separate dark and luminous parts continues to prove challenging.  The radial variation of the circular speed simply does not contain enough information to allow a unique decomposition between the baryonic mass, which has an uncertain mass-to-light ratio, $\\Upsilon$, and the dark halo, whose density profile is generally described by some adopted parametric function \\citep{albada3198, LF89, BSK04}.  Predictions for $\\Upsilon$ from stellar population synthesis models that match broad-band colors \\citep[e.g.][]{Bell03} are useful, but not precise.  Despite intense effort, they are still sufficiently uncertain to be consistent with both maximum and half-maximum disk, which is the range of disagreement \\citep[e.g.][]{Sackett97, Bottema97, S99Rutgers}.  \\citet{McGaugh05} argues that the values can be refined by minimizing the scatter in the Tully-Fisher and/or mass discrepancy-acceleration relation.  A number of dynamical methods have been employed to break the disk-halo degeneracy.  \\citet{Casertano83}, \\citet{bosma98}, and others have suggested that the slight decrease in orbital speed near the edge of the optical disk of a bright galaxy -- the ``truncation signature'' -- could be used as an indicator of disk $\\Upsilon$, but in practice it does not provide a tight constraint.  \\citet{ABP87} and \\citet{Fuchs03} attempt to constrain the disk mass using spiral structure theory.  \\citet{Bottema97} and \\citet{verheijen04} measure the vertical velocity dispersion of disk stars in a near face-on galaxy, which they assume has the same mean thickness of similar galaxies seen edge-on \\citep{KvdKdG}, to constrain the disk mass.  A similar approach is reported by \\citet{CD04} using velocity measurements of individual planetary nebulae.  One of the most powerful, although laborious, methods for barred galaxies was pioneered by Weiner, Sellwood \\& Williams (2001), who made use of the additional information in the driven non-circular motions caused by the bar.  By modeling the observed non-axisymmetric flow pattern of the gas in a 2-D velocity map, they were able to determine the mass-to-light ratio of the visible disk material.  They found that the luminous disk and bar contributed almost all the central attraction in NGC~4123 inside $\\sim 10\\;$kpc, requiring the dark halo to have a very low central density.  \\citet{Ben3095} reports a similar result for a second case, NGC 3095.  The method has also been applied by \\citet{perez04} for several barred galaxies and by \\citet{kranz03} who modeled motions caused by spiral arms. \\citet{BEG03} present a similar study for the Milky Way.  Earlier studies \\citep[\\eg][]{DA83} did not attempt to separate the disk from the dark matter halo \\citep[see][for a review]{SW93}.  Here, we apply the \\citep{wein2} method to the more luminous barred galaxy NGC~1365 in the Fornax cluster.  As one of the most apparently regular, nearby barred spiral galaxies in the Southern sky, NGC~1365 was selected by the Stockholm group for an in-depth study \\citep[see \\eg][]{Lindblad1365}.  Hydrodynamic models of the bar flow pattern were already presented by \\citep{Linlinatha}, based mainly on the velocities of emission-line measurements from many separate long-slit observations.  J\\\"ors\\\"ater \\& van Moorsel (1995, hereafter JvM95) present a kinematic study using the 21 cm line, which suggests that the galaxy is somewhat asymmetric in the outer parts, where the shape of the rotation curve is hard to determine. \\citet{sandqvist} find substantial amounts of molecular gas, but only within 2~kpc of the nucleus, which is resolved in interferometric observations \\citep{Sakamoto07} into a molecular ring in the center plus a number of CO hot spots. \\citet{Galliano05} have found previously unknown MIR sources in the inner 10\\arcsec\\ around the AGN. They are able to correlate some of these MIR sources with radio sources, which they interpret in terms of embedded star clusters because of the lack of strong optical counterparts. \\citet{JCA97} present H-band photometry of the bright inner disk, finding an elongated component in the central region suggesting that the NGC~1365 is a double-barred galaxy, although they note that the light in this component is not as smooth as in their other nuclear bar cases. \\citet{LSKP} also classify it as a double barred galaxy. However, \\citet{Emsellem01} and \\citet{Erwin04} argue against a nuclear bar, citing an HST NICMOS image which resolves the feature into a nuclear spiral. \\citet{Emsellem01} also present stellar kinematics from slit spectra using the $^{12}$CO bandhead.  They propose a model for the inner 2.5 kpc of NGC 1365 consisting of a decoupled nuclear disk surrounded by spiral arms within the inner Lindblad resonance (ILR) of the primary bar. \\citet{Beck05} observed NGC 1365 in radio continuum at 9\\arcsec-25\\arcsec\\ resolution and find radio ridges roughly overlapping with the dust lanes in the bar region. They propose that magnetic forces can control the flow of gas at kiloparsec scales.  Here we present new photometric images, a full 2-D velocity map of the \\ha\\ emission, and a reanalysis of the neutral hydrogen data from JvM95.  We also compare many hydrodynamic models to these new data in an effort to determine the separate disk and dark halo masses in this galaxy. ",
        "conclusions": "We have presented a detailed study of the strongly-barred galaxy NGC~1365, including new photometric images and Fabry-Perot spectroscopy, as well as a detailed re-analysis of the neutral hydrogen observations by J\\\"ors\\\"ater \\& van Moorsel (1995).  We find the galaxy to be at once both remarkably bi-symmetric in its I-band light distribution and strongly asymmetric in the distribution of dust and gas, and in the kinematics of the gas.  These asymmetries extend throughout the galaxy, affecting the bar region, the distribution of gas in the spiral arms and the neutral hydrogen beyond the edge of the bright disk.  The velocity field mapped in the \\ha\\ line showed bright HII regions with velocities that differed by up to $\\sim 80\\;$\\kms\\ from that of the surrounding gas.  Our sparsely-sampled line profiles in these anomalous velocity regions hint at unresolved substructure, suggesting a possible double line profile.  The strong bar and spiral arms complicate the determination of the projection geometry of the disk, assuming it can be characterized as flat in the inner parts.  The inclination of the plane we derive from the kinematic data is smaller by about 10$^\\circ$ from that determined from the photometry.  The strong spiral arms that cross the projected major axis far out in the disk seem likely to bias the photometric inclination and we therefore adopt, in common with other workers, the inclination derived from the gas kinematics.  This preference is supported by the much poorer fits to the observed kinematics obtained when we adopt the photometric inclination (\\S~\\ref{comparsect}).  Our attempts to derive the rotation curve of NGC~1365 were complicated by the fact that neither the \\ha\\ nor the HI velocity maps are consistent with a simple circular flow pattern over a significant radial range.  The bar and spirals clearly distort the gas flow in the luminous disk.  The neutral hydrogen extends somewhat beyond the visible disk but unfortunately has neither a uniform distribution nor regular kinematics.  JvM95 attempted to fit a warp to the outer HI layer that extends into the visible disk, and derived a strongly declining rotation curve.  We chose instead to assume a coplanar flow out to a deprojected radius of 255\\arcsec\\ and to neglect the asymmetric velocities in the neutral hydrogen beyond.  The velocities derived separately from the \\ha\\ and HI data are in good agreement.  Our resulting rotation curve shows a gentle decrease beyond a radius of $\\sim 10\\;$kpc, similar to those observed in other massive galaxies \\citep{CvG91, noordermeer07}.  We used our deprojected I-band image to estimate the gravitational field of the luminous matter, which can be scaled by a single mass-to-light ratio, $\\Upsilon_I$.  We also employed a gradual increase to $\\Upsilon_I$ in the central few kpc to allow for an older bulge-like stellar population, although the light distribution does not appear to have a substantially greater thickness near the center. We combined the central attraction of the axially-symmetrized disk for various values of $\\Upsilon_I$ with two different halo models to fit the observed rotation curve in the region outside the bar -- finding, as always, no significant preference for any $\\Upsilon_I$.  We attempt to fit hydrodynamic simulations of the gas flow pattern in the bar region, in order to constrain $\\Upsilon_I$.  For each type of halo adopted, we run a grid of simulations covering a range of both $\\Upsilon_I$ and $\\Omega_p$, the pattern speed of the bar.  We then project each simulation to our adopted orientation of the galaxy and compare the gas flow velocities in the model with those observed. Since the light distribution in NGC~1365 is highly symmetric, our simulations were constrained to be bi-symmetric, yet the observed gas flow has strong asymmetries.  None of our simulations is capable, therefore, of fitting both sides of the bar simultaneously.  The anomalous position of the dust lane in the western part of the bar suggests that side is the more likely to be disturbed, and we therefore fit our models to the eastern half of the bar only.  After smoothing the model to match the resolution of the kinematic data and masking out five blobs of gas with strongly anomalous velocities, we are able to obtain moderately satisfactory fits to the remaining velocities.  The best fit pattern speed is $\\Omega_p = 24\\;$\\kms$\\rm{kpc}^{-1}$ for both types of halo which places corotation for the bar at $r_L \\simeq 1.23r_B$, in excellent agreement with the value found by \\citet{Linlinatha} and consistent with most determinations of this ratio for other galaxies \\citep[\\eg][]{ADC03}.  Our estimated mass-to-light ratio values are $\\Upsilon_I \\simeq 2.50 \\pm 1$ for the isothermal halo models and $\\Upsilon_I=1.75 \\pm 1$ for the NFW halo.  While the constraints are disappointingly loose, the preferred mass-to-light ratio in both our halo models, $\\Upsilon_I \\simeq 2.0 \\pm 1$, is consistent with that obtained by \\citet{wein2} and \\citet{Ben3095} in two other cases.  For NGC~1365, however, this value implies a massive, but not fully maximal, disk, and we do not find support for the disk-only model with no halo that was suggested by JvM95.  Although such a model can reproduce the declining rotation curve (see their Fig. 24), the simulated gas flow produced by such a model (which in the I band is $\\Upsilon\\sim 3.75$) is quite strongly excluded.  The preferred value of $\\Upsilon_I$ is nicely consistent with those obtained in the two previous studies using this method \\citep{wein2, Ben3095}, but suggest somewhat more massive disks than predicted by population synthesis models \\citep{BdJ,Bell03} for galaxies of these colors.  The halos of our two models required to fit the declining rotation curve in the outer disk are distinctly non-standard, however.  The circular speed in the pseudo-isothermal model declines steadily outside the large core, while the NFW halo has a very high concentration and small scale radius.  Even allowing for compression of the halo as the massive disk forms within it, the original NFW halo has $c \\simeq 22$, $V_{200} \\simeq 123\\;$kms.  The total dark matter mass out to $r_{200}$ in this case is less than three times our estimated disk mass, and the halo is quite unlike those predicted by LCDM models for a galaxy of this mass.  The disturbed distribution and kinematics of the gas in this galaxy clearly complicates our attempt to identify a preferred mass model. Its projected position near the center of the Fornax cluster, together with its velocity within 200~\\kms\\ of the cluster mean, suggest it is a cluster member.  The disturbed nature of the outer HI distribution should not therefore be regarded as surprising.  But the high central density of this massive galaxy should ensure that tidal forces have little influence on the inner part, where we indeed see that the bar and inner spirals are very pleasingly bi-symmetric in the I-band light.  The existence of such strong asymmetries in the inner parts of the gas and dust is rather surprising, therefore.  The asymmetry in the dust distribution and the kinematic map, combined with the existence of a number of patches of \\ha\\ emission with anomalous velocities all suggest that the agent that caused the disturbance was an infalling gas cloud.  We cannot say whether the gas was an isolated intergalactic cloud not associated with a galaxy, or whether it could be a stream of debris from a gas-rich dwarf galaxy that had been tidally disrupted.  The anomalous velocities clearly suggest that the infalling gas has yet to be assimilated in the disk of NGC~1365."
    },
    "0710/0710.5765_arXiv.txt": {
        "abstract": "We examine halo gas cross sections and covering fractions, $f_c$, of intermediate redshift {\\MgII} absorption selected galaxies.  We computed statistical absorber halo radii, $R_{\\rm x}$, using current values of $dN/dz$ and Schechter luminosity function parameters, and have compared these values to the distribution of impact parameters and luminosities from a sample of 37 galaxies.  For equivalent widths $W_r(2796) \\geq 0.3$~{\\AA}, we find $43 \\leq R_{\\rm x} \\leq 88$~kpc, depending on the lower luminosity cutoff and the slope, $\\beta$, of the Holmberg--like luminosity scaling, $R \\propto L^{\\beta}$.  The observed distribution of impact parameters, $D$, are such that several absorbing galaxies lie at $D>R_{\\rm x}$ and several non--absorbing galaxies lie at $D < R_{\\rm x}$.  We deduced $f_c$ must be less than unity and obtain a mean of $\\left< f_c \\right> \\sim 0.5$ for our sample. Moreover, the data suggest halo radii of {\\MgII} absorbing galaxies do not follow a luminosity scaling with $\\beta$ in the range of $0.2-0.28$, if $f_c= 1$ as previously reported. However, provided $f_c \\sim 0.5$, we find that halo radii can remain consistent with a Holmberg--like luminosity relation with $\\beta \\simeq 0.2$ and $R_{\\ast} = R_{\\rm x}/\\sqrt{f_c} \\sim 110$~kpc.  No luminosity scaling ($\\beta=0$) is also consistent with the observed distribution of impact parameters if $f_c \\leq 0.37$. The data support a scenario in which gaseous halos are patchy and likely have non--symmetric geometric distributions about the galaxies.  We suggest halo gas distributions may not be govern primarily by galaxy mass/luminosity but also by stochastic processes local to the galaxy. ",
        "introduction": "Understanding galaxy formation and evolution is one of the most important topics of modern astronomy. The extended distribution of baryonic gas surrounding galaxies holds great potential for constraining theories of their formation. However, the sizes of gaseous galaxy halos along with the distribution of gas within are not well understood. Numerical models have been able to synthesize the formation and evolution of large scale structures, however, there are unresolved issues regarding the evolution of individual galaxies and halos. The halo baryon--fraction problem \\citep[e.g.,][]{mo02} and the rapid cooling of gas \\citep[e.g.,][]{white78} result in galaxy halos which have little or no gas soon after they form. These effects are not seen in the observable universe since there is an abundance of galaxies where gas has been detected in halos via quasar absorption lines.  From an observational standpoint, quasar absorption lines provides a unique means of probing the extent and abundance of halo gas. Although, quasar absorption line observations to date are sufficient the recognize the aforementioned problems, they are lacking the detail required to statistically constrain the distribution of the baryonic gas in the halos of simulated galaxies. Cross--correlations between absorbers and galaxies hold the promise to yield useful information on cloud sizes and halo gas covering fractions. First steps towards incorporating multi--phase gas in semi--analytical models and numerical simulations suggest that warm gas in halos extends out to galactocentric distances of $\\sim 150$~kpc with cloud covering fractions of $\\sim 0.25-0.6$ \\citep{maller04,kaufmann06}.  The association of {\\MgIIdblt} doublet absorption in quasar spectra with normal, bright, field galaxies has been firmly established \\citep[e.g.,][]{bb91,sdp94,cwc-china}.  In an effort to understand halo sizes and gas distributions, \\citet[][hereafter S95]{steidel95} searched for foreground galaxies associated with {\\MgII} absorption within $\\sim10''$ ($\\sim65$~kpc for $z=0.5$) of quasars\\footnote{Throughout we adopt a $h=0.70$, $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$ cosmology. All quoted physical quantities from previously published works have been converted to this cosmology.}. The sample consisted of 53 absorbing and 14 non--absorbing galaxies with a {\\MgII} $\\lambda 2796$ equivalent width sensitivity limit of $W_{r}(2796) > 0.3$~{\\AA}. S95 directly fitted the data by assuming a Holmberg--like luminosity scaling, \\begin{equation} R(L) = R_{\\ast} \\left( \\frac{L}{L^{\\ast}} \\right) ^{\\beta} \\quad {\\rm kpc}, \\label{eq:rl} \\end{equation} and minimizing the number of non--absorbing and absorbing galaxies above and below the $R(L)$ relation. The best fit obtained clearly showed that absorbing and non--absorbing galaxies could be separated and that the halo radii $R(L_K)$ and $R(L_B)$ scale with luminosity with $\\beta = 0.15$ and $\\beta = 0.2$, respectively, where an $L_B^{\\ast}$ galaxy has a gas halo cross section of $R_{\\ast} = 55$~kpc. Furthermore, since almost none of the absorbing galaxies were observed above the $R(L)$ boundary and that almost none of the non--absorbing galaxies were observed below the $R(L)$ boundary, S95 inferred that {\\it all\\/} $L> 0.05L^{\\ast}$ galaxies are hosts to {\\MgII} absorbing gas halos characterized by a covering fraction of unity and a spherical geometry which truncates at $R(L)$. Examination of this now ``standard model'' has been the subject of several theoretical studies \\citep[e.g.,][]{cc96,mo96,lin01}.  \\citet{gb97} determined a steeper value of $\\beta = 0.28$ for the B--band luminosity obtained from a best fit to the upper envelope of the distribution of impact parameters of 26 absorbing galaxies. They found $R_{\\ast} = 67$~kpc.  Using a reverse approach of establishing foreground galaxy redshifts and then searching for {\\MgII} absorption in the spectra of background quasars yields results inconsistent with a covering fraction of unity. For example, \\citet{bowen95} identified 17 low--redshift galaxies with background quasar probing an impact parameter range between $3-162$~kpc. Galaxies that were probed at impact parameters greater than 13~kpc had no absorption in the halo ($W_{r}(2796) \\geq 0.40-0.9$~{\\AA}), however, four of the six galaxies within 13~kpc of the halo produced {\\MgII} absorption. For intermediate redshift galaxies, \\citet{bechtold92} reported a covering fraction $f_c \\simeq 0.25$ for $W_{r}(2796) \\geq 0.26$~{\\AA} for eight galaxies with $D \\leq 85$ kpc. Also, \\citet{tripp-china} reported $f_c \\sim 0.5$ for $W_{r}(2796) \\geq 0.15$~{\\AA} for $\\sim 20$ galaxies with $D \\leq 50$ kpc.  These results are also consistent with the findings of \\citet{cwc-china} who reported very weak {\\MgII} absorption, $W_{r}(2796) < 0.3$~{\\AA}, well inside the $R(L)$ boundary of bright galaxies; these galaxies would be classified as ``non--absorbers'' in previous surveys.  They also report $W_r(2796) > 1$~{\\AA} absorption out to $\\simeq 2 R(L)$.  All these results suggest that there are departures from the standard model, that the covering fraction of {\\MgII} absorbing gas is less than unity, and that the halo sizes and the distribution of the gas appear to diverge from the $R(L)$ relation with spherical geometry.  Another approach to understanding halo sizes and gas distributions is to determine the statistical properties of {\\MgII} absorbing gas and then compute the statistical cross section from the redshift path density, $dN/dz$ \\citep[see][]{lanzetta95}.  The downfall of this method is that a galaxy luminosity function must be adopted in order to estimate $R_{\\ast}$.  \\citet{nestor05} acquired a sample of over 1300 {\\MgII} absorption systems, with $W_{r}(2796) \\geq 0.3$~{\\AA} from the Sloan Digital Sky Survey (SDSS). Using the $K$--band Holmberg--like luminosity scaling and luminosity function of MUNICS \\citep{drory03}, Nestor {\\etal} computed $R_{\\ast} = 60-100$~kpc for adopted minimum luminosity cutoffs of $L_{min}=0.001-0.25L^{\\ast}$. They found no redshift evolution of $R_{\\ast}$ over the explored range of $0.3\\leq z \\leq 1.2$.  \\citet{zibetti06} studied the statistical photometric properties of $\\sim2800$ {\\MgII} absorbers in quasar fields imaged with SDSS.  Using the method of image stacking, they detected low--level surface brightness (SB) azimuthally about the quasar. The SB profiles follow a decreasing power law with projected distance away from the quasar out to $100-200$~kpc. These results imply that absorption selected galaxies may reside out to projected distances of 200~kpc.  However, it is worth noting that the extended light profiles may be an artifact of clustering of galaxies. Cluster companions of the {\\MgII} absorbing galaxies could extend the observed light profile over hundreds of stacked images. Thus, one would infer that {\\MgII} absorbing galaxies are present at a larger impact parameters than would be found in direct observation of individual galaxies.   Motivated by recent expectations from simulations that halo gas is dynamically complex and sensitive to the physics of galaxy formation, we investigate the standard halo model of {\\MgII} absorbers. We also aim to provide updated constraints on $f_c$ and $\\beta$ for galaxy formation simulations. In this paper, we demonstrate that $f_c < 1$ and question the validity of the Holmberg--like luminosity scaling (Eq.~\\ref{eq:rl}). Using high resolution quasar spectra, we explore {\\MgII} absorption strengths to an order of magnitude more sensitive than previous surveys which allow us to re--identify non--absorbing galaxies as ``weak'' absorbing galaxies.  In \\S~\\ref{sec:data} we describe our sample and analysis. In \\S~\\ref{sec:results}, we present new calculations of the statistical absorber radius computed using the statistically measured absorption path density $dN/dz$ and the Schechter luminosity function.  We then compare these values to the empirical results of S95 and to a sample of known {\\MgII} absorption selected galaxies with measured luminosities and impact parameters. We also examine how individual halos behave with respect to the statistical halo. In \\S~\\ref{sec:dis}, we discuss the properties and distribution of gas in halos.  Our concluding remarks are in \\S~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  In conclusion, the gas covering fraction must be less than unity since the observed impact parameter distribution of absorbing galaxies does not fall exclusively within the statistical absorber halo radius in the range of $43 \\leq R_{\\rm x}\\leq 88$~kpc.  The fact that some absorbing galaxies are found at $D>R_{\\rm x}$ and some non--absorbing galaxies are found at $D<R_{\\rm x}$ implies $f_c < 1$ and that the standard halo model cannot describe halos on a case by case basis. This highlights the power of using the statistics of absorption line surveys to constrain the properties of halos in relation to the measured distributions in absorption selected galaxy surveys.  By quantifying how individual galaxy halos deviate from a ``standard'' halo, we have obtained an average gas covering fraction of $\\left< f_c \\right> \\sim 0.5$. It is possible that $f_c$ exhibits both a radial and an equivalent width dependence, though we cannot address this with our sample.  Values of $f_c$ are likely to depend on galaxy star formation rates, and galaxy--galaxy mergers and harassment histories; processes that give rise to patchy and geometrically asymmetric gas distributions. Alternatively, the absorption properties of intermediate redshift halos may be governed by the dark matter over density, $\\Delta \\rho /\\rho$, and redshifts at which the galaxies formed \\citep{cwc06}.  Our results also show that, if $f_c<1$, the sizes of {\\MgII} absorbing halos can still follow a Holmberg--like luminosity relation with $\\beta$ in the range of $0.2-0.28$ \\citep[S95;][]{gb97}, which corresponds to $R_{\\ast}\\sim 110$~kpc. If $\\beta=0$ is assumed, then $f_c \\leq 0.37$ for our sample to be consistent with no luminosity scaling. In semi--analytical models in which {\\MgII} absorbing gas is infalling and is pressure confined within the cooling radius of hot halos \\citep[e.g.,][]{mo96,burkert00,lin00,maller04}, a Holmberg--like luminosity relation in quasar absorption line systems naturally arises \\citep{mo96}. However, these models have great difficulty explaining {\\MgII} absorption at impact parameters greater than $\\sim 70$~kpc. If on the other hand halo gas spatial distributions are governed by stochastic mechanical processes, as suggested by \\citet{kacprzak07}, then there is no {\\it a priori} reason to expect a clean halo--size luminosity scaling. It is likely that some combination of these scenarios contribute to the statistical values of $f_c$ and $\\beta$. Thus, it is reasonable to suggest that {\\MgII} halos sizes may not be strictly coupled to the host galaxy luminosity.  Further work on the cross--correlations between absorbers and galaxies would provide better estimates of $f_c$ and $\\beta$, two quantities that provide direct constraints of galaxy formation simulations. Also needed are additional constrains on the relative kinematics of the absorbing halo gas and galaxies \\citep[e.g.,][]{s02,ellison03,kacprzak07b}. What is required is the development of techniques to quantitatively compare observational data with mock quasar absorption line analysis of simulated galaxy halos \\citep{cwcaas06}."
    },
    "0710/0710.5179.txt": {
        "abstract": "We propose a scenario for the formation of the Main Belt in which asteroids incorporated icy particles formed in the outer Solar Nebula. We calculate the composition of icy planetesimals formed beyond a heliocentric distance of 5 AU in the nebula by assuming that the abundances of all elements, in particular that of oxygen, are solar. As a result, we show that ices formed in the outer Solar Nebula are composed of a mix of clathrate hydrates, hydrates formed above 50 K and pure condensates produced at lower temperatures. We then consider the inward migration of solids initially produced in the outer Solar Nebula and show that a significant fraction may have drifted to the current position of the Main Belt without encountering temperature and pressure conditions high enough to vaporize the ices they contain. We propose that, through the detection and identification of initially buried ices revealed by recent impacts on the surfaces of asteroids, it could be possible to infer the thermodynamic conditions that were present within the Solar Nebula during the accretion of these bodies, and during the inward migration of icy planetesimals. We also investigate the potential influence that the incorporation of ices in asteroids may have on their porosities and densities. In particular, we show how the presence of ices reduces the value of the bulk density of a given body, and consequently modifies its macro-porosity from that which would be expected from a given taxonomic type. ",
        "introduction": "In recent years, some objects within the Main Belt of asteroids have been found to display cometary characteristics (Hsieh \\& Jewitt 2006). Objects such as 133P/Elst-Pizarro, P/2005 U1 and 118401 (1999 RE$_{70}$) occupy orbits that are entirely decoupled from Jupiter within the Main Belt, and are probably bodies that have undergone a recent collision, revealing previously buried volatile material, and leading to the observed dusty outgassing.  In addition, present-day surface water ice and possible water sublimation have been reported on Ceres (Lebofsky et al. 1981; A'Hearn \\& Feldman 1992; Vernazza et al. 2005). This is consistent with recent Hubble Space Telescope (HST) observations which suggest that Ceres' shape is the result of the dwarf planet consisting of a rocky core surrounded by an ice-rich mantle (Thomas et al. 2005) - an idea in agreement with several thermal evolution scenarios (McCord \\& Sotin 2005) that suggest that the ice content of the asteroid is between 17\\% and 27\\%, by mass. These observations are supported by the evidence of hydrated minerals in meteorites which provide samples of rock from asteroids in the Main Belt. Most of these minerals formed as a result of water ice accreting with the chondritic meteorite parent bodies, melting, and driving aqueous alteration reactions (Clayton \\& Mayeda 1996; Jewitt et al. 2007).  It seems likely, then, that some objects in the asteroid belt have incorporated significant amounts of water ice (and possibly other volatiles) during their formation in the early stages of the solar System. These bodies would have incorporated icy particles\\footnote{By icy particles is meant planetesimals composed of a mix of ices and rocks.} coming from the outer nebula that survived their inward drift due to gas-drag through the disk (Mousis \\& Alibert 2005 -- hereafter MA05). The volatile fraction incorporated in this manner could vary depending on the inward flux of icy planetesimals from the external region and the heliocentric location of the asteroid, together with the density of the proto-solar nebula. Given that the current asteroid belt lies closer to the Sun than the ``snow-line'', postulated to lie at around 5 AU in the Solar Nebula, these results are a little unexpected. In this context, understanding how volatiles were incorporated into the asteroids is therefore important, not only for the study of the asteroids themselves, but also for our understanding of the processes by which the solar system came into being.  To this end, MA05 studied the possibility of determining the nature and composition of the ices which were incorporated into Ceres. They used a time dependant model of the Solar Nebula and showed that icy particles of sizes between 0.1 and 10 metres could drift from heliocentric distances greater than 5 AU to the present location of Ceres without encountering temperatures or pressures high enough to vaporise the ices within. The authors then suggested that ices produced in the outer Solar Nebula were transported inwards to become incorporated in the solids which accreted to form Ceres.  The present work aims to improve upon the calculation detailed in MA05, along with expanding the results to involve the entire asteroid belt, rather than just its largest member. In particular, MA05 postulated that all volatiles, except CO$_2$\\footnote{CO$_2$ is the only major volatile species which does not form a clathrate hydrate in the Solar Nebula because it condenses as a pure ice prior to being trapped by water.}, were trapped by water in the form of hydrates or clathrate hydrates in the outer solar nebula. This assumption was supported by the work of Hersant et al. (2001) who estimated that Jupiter was formed at temperatures higher than $\\sim$40 -- 50 K. The accretion of ices in the form of hydrates and clathrate hydrates was thus required during the formation of the planet in order to explain the volatile enrichments observed in its atmosphere\\footnote{The abundances of volatile species in Jupiter's atmosphere have been measured using the mass spectrometer on board the {\\it Galileo} probe. These measurements reveal that the giant planet's atmosphere is enriched by a factor of $\\sim$3 in Ar, Kr, Xe, C, N, and S compared to the solar abundances (Owen et al. 1999).} (Gautier et al. 2001a,b). Indeed, since these ices are usually formed at temperatures higher than that reached by the nebula at the time of Jupiter's completion, as defined by Hersant et al. (2001), they can be incorporated in the planetesimals accreted by the giant planet during its growth. However, the amount of water that would be required in the nebula to trap all these volatiles as hydrates and clathrate hydrates exceeds that derived from the solar oxygen abundance. Therefore, MA05 made the {\\it ad hoc} hypothesis that oxygen was ``oversolar'' in the gas-phase in order to provide enough available water in the Solar Nebula\\footnote{ The oxygen abundance required for to allow the trapping of all volatile species in the form of hydrates or clathrate hydrates is $\\sim$1.9 times the solar abundance, with CO$_2$:CO:CH$_4$ = 1:1:1 and N$_2$:NH$_3$ = 1:1 (the nominal nebula gas phase ratios used in this work).}. Additionally, Hersant et al. (2001) only used an evolutionary Solar Nebula model to derive the disk's temperature at the time when the mass of Jupiter's feeding zone was equal to that of the gas in its current envelope. They thus neglected many important effects such as the influence of protoplanet formation on the structure of the disk (e.g. Fig. 2 of Alibert et al. 2004). However, recent giant planet core-accretion formation models that include migration, disk evolution, such as that proposed by Alibert et al. (2004), have shown that the disk's temperature can be as low as $\\sim$10 -- 20 K at the end of Jovian formation. This implies that Jupiter itself can accrete ices during its formation that were produced at temperatures lower than those required for clathration. As a result, no extra water is required in the nebula to allow all the volatile species to be trapped in clathrate hydrates, and the oversolar oxygen abundance condition in the nebula can be relaxed.  In Section 2, we calculate the composition of ices produced in the outer Solar Nebula under the assumption that the abundances of all elements, in particular that of oxygen, are solar. In Section 3, we consider the inward migration of particles produced at various locations in the nebula, and at different times. This allows us to examine whether some planetesimals formed in the outer Solar Nebula may have drifted to the current position of the Main Belt without encountering temperature and pressure conditions high enough to vaporize the ices they contain. In Section 4, we examine the uncertainties in the determination of the physical properties of asteroids. We also investigate the potential influence that the incorporation of ices in these objects may have on their porosities and densities. Section 5 is devoted to summary and discussion. ",
        "conclusions": "In order to explain the presence of hydratation and cometary features in the Main Belt, we have proposed that asteroids incorporated during their formation icy particles formed in the outer Solar Nebula. We have then calculated the composition of the ices trapped in these planetesimals formed beyond a heliocentric distance of 5 AU in the nebula, in a manner consistent with the formation of Jupiter, by assuming that the gas-phase abundances of all elements, in particular that of oxygen, are solar. As a result, we have found that the ices being formed in the outer Solar Nebula are composed of a mix of clathrate hydrates, hydrates formed above 50 K, and pure condensates produced at temperatures between $\\sim$20 K and $\\sim$50 K. We have noted that, whatever the input parameters adopted in the modelling of the disk, or the formation location considered for icy planetesimals at heliocentric distances beyond 5 AU, their composition remains almost constant, provided that the gas-phase abundances are homogeneous in the nebula. We have argued in this work that gas-drag is responsible for the inward drift of icy particles formed in the outer nebula towards the forming Main Belt. To support this hypothesis, we have showed that, at some epochs of the disk's evolution, some particles produced in the outer nebula may drift to the current position of the Main Belt without encountering temperature and pressure conditions high enough to vaporize the ices they contain.  The current distribution of ices potentially existing in asteroids has probably been deeply altered after their formation. The effect of solar insolation may have vaporized the ice within nearer asteroids (semi-major axes of $\\sim$2 AU), melted the ice of mid-range asteroids situated at $\\sim$3 AU, but should not have affected the ice in asteroids located at greater heliocentric distances. Inner and outer asteroids would therefore display no detectable hydratation features, either because the ice was vaporized and dissipated, or because the ice never melted and thus did not react with the surface minerals to a sufficient extent as to allow detection (Cyr et al. 1998). In this context, we have proposed that, from the detection and identification of initially buried ices revealed by recent impacts on the surfaces of asteroids, it could be possible to infer the thermodynamic conditions that occurred within the Solar Nebula during the accretion of these bodies, as well as during the inwards migration of the icy planetesimals which they incorporated. However, this statement requires that either no parent body processing or modification took place during and after the formation of asteroids. For example, we have noted that subsequent alteration of the volatile phases in asteroids may occur due to catalytic reactions in their interiors.  We have also investigated the potential influence that the incorporation of ices in asteroids may have on their porosities and densities. In particular, we have showed that the presence of ices can considerably reduce the value of the bulk density of the body, and consequently its macro-porosity, that would be expected from a given taxonomic type.  That volatiles were delivered to areas within the ice line is clearly beyond doubt. In addition to the gas-drag mechanism described in this work, it is also likely that a significant amount of volatile material was dynamically driven inwards in the latter stages of planet formation. We still see the tail of this dynamical, chaotic volatile movement today -- the comets we observe passing through the inner solar system are the bearers of ices formed far beyond the snow line, and held in deep freeze since the early days. During the latter stages of planetary migration, the flux of such objects passing through the inner solar system, and hence encountering the asteroids, was significantly higher. Of particular interest, when one considers veneers of volatile material near the surface of the asteroids, is the Late Heavy Bombardment. In the Nice model, (see e.g. Gomes et al, 2005), vast amounts of volatile-rich material is flung inwards from the outer solar system approximately 700 Million years after its birth. This event, caused by the resonant destabilisation of the outer solar system, would have coincided with a simultaneous stirring of the asteroid belt, leading to an impact flux upon the Earth containing approximately even proportions of asteroidal and cometary material. It is clear, though, that the Earth would not be the only object to encounter volatiles injected in this way, and the possibility of a late-veneer of ice arriving in the asteroid belt is surely something which must be acknowledged in future work. In addition to this aggressive and chaotic injection of material, there is also a\u00de gentler mechanism by which volatiles can be driven inwards as a result of planetary migration. As planets migrate, material can be trapped in the locations of mean-motion resonances (MMR), which sweep in front of the planet through it's motion. Evidence of material being swept outwards in the resonances of Neptune is clear for all to see -- the Plutino family of objects are locked in the 2:3 MMR with the planet, and have an inclination distribution which can tell us a great deal about the distance over which the planet migrated, sweeping them along. Inward migration can have the same effect -- the interior resonances of a planet can collect material as it moves inwards, and sweep it along -- giving a mechanism by which volatile material can be eased inwards, with the migration of a giant such as Jupiter. Work such as Fogg \\& Nelson (2006) has shown that such resonant forcing can operate with a resonable efficiency, even for significantly faster migration than expected in our solar system, and so the effects of this behaviour should not be ignored in future work.  In spite of the growing pool of evidence pointing towards the existence of water ice in the Main Belt, its detection on asteroids is a challenging observational problem. Large bodies such as Ceres are suspected to have retained a large amount of water since their formation, perhaps even including an internal liquid ocean, throughout the age of the solar system (McCord \\& Sotin 2005). This could particularly be the case if this internal water was originally mixed with some ammonia, in agreement with our composition calculations in section \\ref{icy}, which would have the effect of lowering the melting point of the water (ammonium bearing minerals have been suggested by King et al. (1992) as an alternate explanation for the origin of the 2.07 $\\mu$m band seen in the spectrum of Ceres). Nevertheless, internal water can only be indirectly probed, either by measuring the hydrostatic shape of the object, as was done in the case of Ceres, or by inferring its density from its size and mass, when they are known or, more evidently, from outgasing activities. The case of Ceres is particularly interesting since, in spite of several possible pieces of evidence which support it being a highly hydrated body, the only report of water detection on the dwarf planet is the observation of OH escaping from its northern pole\\footnote{An OH atmosphere was indeed observed around Ceres after perihelion by A'\u00d5Hearn and Feldman (1992) by performing IUE long exposure spectra, with column densities of the order of $10^{11}$~cm$^{-2}$.}, is still not confirmed. Nevertheless, this detection could be explained in the context of the accumulation of ice during winter on the surface or within the subsurface layer, which would then dissipate during summer, when the surface temperature rises. Similar transient events have been suggested as possible mechanisms to trigger the geyser-like activity taking place near the south pole of Enceladus and reported by Cassini (Porco \\& Team 2006).  It is interesting to note that, considering the gravity and the day-side temperature of Ceres, any outgassed atmosphere would be rapidly lost. The mean thermal velocity $v_0$ of H$_2$O, for instance, would be close to the escape velocity  ($v_{\\infty}=516$~km~s$^{-1}$). Assuming a subsolar temperature of 215~K (Dotto et al. 2000), $v_0$ would vary between 450 to 350~km~s$^{-1}$ from the subsolar point to a zenith angle of 80$^{\\circ}$. As a consequence, hydrodynamical escape would occur ($v_{esc}^{2}/v_{0}^{2} \\le 2$). The photolysis of H$_2$O by solar EUV makes this atmospheric escape even more efficient by giving the photodissociation products OH and H some additional kinetic energy. Considering the short lifetime of H$_2$O at $\\sim$3~AU ($<$9~days), and the fact that the mean thermal velocity of H atoms exceeds $v_{\\infty}$, a tenuous atmosphere of OH is expected if water is outgassed by the asteroid at a sufficient rate. Due to the transcient nature of the atmosphere, the loss of water to space is limited by the flux of water from the interior to the surface. At low latitude, where ice is not stable, the continuous flux of water from the interior to space is too low to be detected.  Only an accumulation of water ice at high latitude before perihelion, followed by an outgassing of H$_2$O  associated with post-perihelion warming seems to result in an observable column density of OH. These results were found to be consistent with an earlier work done by Fanale and Salvail (1989), who estimated the mean loss rate of H$_2$O to be in the range 30-300~g~s$^{-1}$. Even if one assumes that the atmospheric loss observed by  A\u00d5'Hearn and Feldman (1992) occurs continuously at the same rate and at all latitudes (which is obviously wrong as this maximum loss requires high latitudes and post-perihelion conditions), the water loss remains below 4 kg~s$^{-1}$, which, integrated over 4.5~Gyr, corresponds to only 0.07\\% of the mass of Ceres. If the loss rate of H$_2$O in Ceres remained constant throughout its thermal history, the initial water reservoir is thus likely to be integraly preserved. Moreover, since the other volatile species are expected to be trapped as hydrates, clathrate hydrates and pure condensates in this reservoir, we can conclude that they have also been preserved from outgassing throughout the thermal history of the asteroid.  Ceres being the largest and, due to its size, probably the wettest Main Belt asteroid, it is an ideal target for carrying out observations aiming at constraining its water regime. The experiment searching for water being vaporised within the polar regions of Ceres should be repeated with the state-of-the-art instrumentation available today on large telescopes. Such a detection would confirm unambiguously the presence of a large amount of water near the surface of Ceres. Direct observation of water ice, or of the effects of hydration, on the surface of Ceres can also be attempted for lower latitudes on the asteroid using a combination of high-angular resolution and spectroscopic instruments permitting the full resolution of its surface to the ~30-40 km level. Due to its low spectral resolution, imaging of the surface of Ceres, even when it is spatially resolved using HST or adaptive optics, is not sensitive to the presence of ice, while the detection of such is within the reach of low resolution spectroscopic observations (e.g. the detection of absorption features in the 1.0-3.5 $\\mu$m region). A spatially resolved spectroscopic mapping of the surface of Ceres in the near-infrared can be done with today's ground-based telescopes and would permit the mapping of the strength of the 3 $\\mu$m band, and allow the search for regions on the surface where interstitial water ice, or hydration features could be present, for instance at the location of cracks within the surface of Ceres, or the locations of deep impact craters. Indeed, recent HST (Thomas et al. 2005) and adaptive optics (Carry et al. 2007) imaging observations of Ceres revealed the presence of large impact craters across its surface which have likely disrupted the outer crust of the asteroid enough to directly expose the sub-surface mantle of wetter material. Finally, a spectroscopic study of the surface of Ceres, in order to search for the spectral signature of water and maybe those of other volatiles, should not be limited to one wavelength region (although the near-infrared range offers many diagnostic bands) but should, instead, encompass a wider range, from the near-UV to infrared wavelengths, in order to improve the identification of the chemicals species responsible for these spectral features.  Finally, the NASA Discovery mission {\\it Dawn}, which has been launched in September 2007 and whose arrival at Ceres is scheduled for 2015, will certainly bring new constraints on the presence of volatiles in the Main Belt. In particular, the {\\it Dawn} mapping spectrometer (MS) covers the spectral range from the near UV (0.25 \\micron) through the near IR (5 \\micron) and has moderate to high spectral resolution and imaging capabilities (Russell et al. 2004). These characteristics make it an appropriate instrument for determining the asteroid's global surface composition. Near-infrared mapping of the surface of Ceres at small spatial scales will be very sensitive to volatile concentrations and may reveal ice spots on fresh impact-crater ridges. Moreover, the gravitation investigation of Ceres will allow the determination of its gravity field up to the 12th harmonic degree (Russell et al. 2004). Such a measurement will enable the shape and gravity models to characterize crustal and mantle density variations and, consequently, the amount of volatiles trapped therein."
    },
    "0710/0710.4280_arXiv.txt": {
        "abstract": "Long Gamma Ray Bursts (GRBs) constitute an important tool to study the Universe near and beyond the epoch of reionization. We delineate here the characteristics of an 'ideal' instrument for the search of GRBs at $z\\ge 6-10$. We find that the detection of these objects requires soft band detectors with a high sensitivity and moderately large FOV. In the light of these results, we compare available and planned GRB missions, deriving conservative predictions on the number of high-$z$ GRBs detectable by these instruments along with the maximum accessible redshift. We show that the {\\it Swift} satellite will be able to detect various GRBs at $z\\ge 6$, and likely at $z\\ge 10$ if the trigger threshold is decreased by a factor of $\\sim 2$. Furthermore, we find that INTEGRAL and GLAST are not the best tool to detect bursts at $z\\ge 6$: the former being limited by the small FOV, and the latter by its hard energy band and relatively low sensitivity. Finally, future missions (SVOM, EDGE, but in particular EXIST) will provide a good sample of GRBs at $z\\ge 6$ in a few years of operation. ",
        "introduction": "The study of the Universe at the epoch of reionization is one of the main goal of available and future space missions. In the last few years, our knowledge of the early Universe has been enormously increased mainly owing to the observation of Quasars by the SDSS survey (Fan 2006). Long gamma ray bursts (GRB) may constitute a complementary way to study the cosmos and the early evolution of stars avoiding the proximity effects and possibly probing even larger redshifts up to $z\\sim 10$. The five GRBs detected at $z\\gsim 5$, over a sample of about 200 objects observed with the {\\it Swift} satellite (Gehrels et al. 2004), show that a large percentage of GRBs is detected at high-$z$. The current record holder is $z = 6.29$ (Tagliaferri et al. 2005, Kawai et al. 2006). The identification of a large number of GRBs at $z\\ge 6$ will open a new window in the study of the early Universe. Just to give some example, GRBs can be used to constrain the reionization history (Totani et al. 2006, Gallerani et al. 2007), to study the metallicity and dust content of normal galaxies at high-$z$ (Campana et al. 2007), to  probe the small-scale power spectrum of density fluctuations (Mesinger, Perna \\&  Haiman 2005). Moreover, available and future GRB missions might be the first observatories to detect individual Population III stars, provided that massive metal-free stars were able to trigger GRBs (see Bromm \\& Loeb 2007 for a review). Finally, the study of GRBs at high redshift is interesting by itself. In particular, thanks to cosmological time dilation, the study of the early phases of the afterglow is easier and can provide fundamental highlight on the central engine and burst physics.   In this paper, we delineate the main characteristics of an 'ideal' instrument for the search of GRBs at $z\\ge 6-10$. In particular, we explore different observational bands, deriving the best combination of sensitivity and field of view (FOV) in order to detect bursts at $z\\gsim 10$. In the light of these results, we compare available and planned X-- and Gamma--ray missions, deriving conservative predictions on the number of high-$z$ GRBs detectable by these instruments along with the maximum accessible redshift.  The paper is organized as follows. In Sect.~2 we briefly describe the different models here adopted. In Sect.~3, we derive the main characteristic of an 'ideal' instrument for exploring the high-$z$ GRB population, whereas predictions for available (planned) GRB missions are given in  Sect.~4 (Sect.~5). Finally, we summarize our results in Sect.~6. ",
        "conclusions": "We have explored the characteristics of an ideal mission for the search of GRBs near and beyond the epoch of reionization, i.e. $z> 6-10$. In particular, we considered different observational bands, deriving the best combination of sensitivity and field of view in order to detect bursts at $z\\gsim 10$. We found that such an experiment requires soft band detectors and high sensitivity, whereas large FOV or wide energy coverage are less important. Assuming 3 sr FOV, an observational band of 8-200 keV, and a sensitivity as low as 0.1 ph s$^{-1}$ cm$^{-2}$, this instrument would be able to detect $\\sim 40$ (3) GRBs at $z\\ge 6$ ($z\\ge 10$) in one year of mission.  In the light of these results, we compared available and planned GRB missions, deriving conservative predictions on the observable number of GRBs at $z\\ge 6$ and $z\\ge 10$ along with the maximum accessible redshift. We have shown that {\\it Swift} is a viable tool to detect GRBs at $z\\sim 6$. At the actual trigger threshold, $1.3-4$ GRBs per year should be identified above this redshift. We discuss also the possibility of increasing the number of high-$z$ detections by lowering the {\\it Swift} trigger threshold. We found that the number of detectable GRBs doubles by lowering this by about 50\\%. Assuming to be able to further lowering the trigger threshold down to 0.1 ph s$^{-1}$ cm$^{-2}$, {\\it Swift} should detect $\\sim 1$ GRB per year at $z\\ge 10$.  The INTEGRAL and GLAST satellites do not appear the best tools to search for GRBs at very high redshift. The former is limited by the very small FOV whereas the GLAST hard energy observational band and relatively low sensitivity do not allow to detect more than 1 GRB per year at $z\\ge 6$. No GRB at $z\\ge 10$ is expected during the entire mission of both instruments.  Finally, we show that future missions, like SVOM, EDGE, and in particular EXIST, will be able to collect a good number of GRBs at $z\\ge 6$ in a few years of operations. This sample can be use to study the early Universe, possible providing strong constrain on the reionization process (Gallerani et al. 2007), and deriving estimate on the star formation and metallicity/dust content in normal high-$z$ galaxies."
    },
    "0710/0710.5529_arXiv.txt": {
        "abstract": "The strong spectral features near 2.2 $\\mu$m in early-type galaxies remain relatively unexplored. Yet, they open a tightly focused window on the coolest giant stars in these galaxies -- a window that can be used to explore both age and metallicity effects. Here, new measurements of K-band spectral features are presented for eleven early-type galaxies in the nearby Fornax galaxy cluster. Based on these measurements, the following conclusions have been reached: (1) in galaxies with no signatures of a young stellar component, the K-band \\naK\\ index is highly correlated with both the optical metallicity indicator \\mgfe\\/ and the central velocity dispersion $\\sigma$; (2) in the same galaxies, the K-band Fe features saturate in galaxies with $\\sigma > 150$ \\kms\\/ while \\naK\\/ (and \\mgfe) continues to increase; (3) [Si/Fe] (and possibly [Na/Fe]) is larger in all observed Fornax galaxies than in Galactic open clusters with near-solar metallicity; (4) in various near-IR diagnostic diagrams, galaxies with signatures of a young stellar component (strong \\hb, weak \\mgfe) are clearly separated from galaxies with purely old stellar populations; furthermore, this separation is consistent with the presence of an increased number of M-giant stars (most likely to be thermally pulsating AGB stars); (5) the near-IR \\naK\\/ vs.~$\\sigma$ or \\fe\\/ vs.~$\\sigma$ diagrams discussed here seem as efficient for detecting putatively young stellar components in early-type galaxies as the more commonly used age/metallicity diagnostic plots using optical indices (e.g H$\\beta$ vs.~\\mgfe). The combination of these spectral indices near 2.2 $\\mu$m with high {\\it spatial} resolution spectroscopy from ground-based or space-based observatories promises to provide new insights into the nature of stellar populations in the central regions of distant early-type galaxies. ",
        "introduction": "\\label{sec:intro}  Understanding the stellar content of early-type galaxies is fundamental to understanding their star formation and chemical evolution history.  Most early-type galaxies are too distant to resolve their individual stars with current technology, rendering the direct study of their stellar populations impossible.  Thus, their stellar populations must be studied using indirect methods.  In recent decades, significant effort has gone into trying to better constrain the stellar contents for early-type galaxies using optical spectroscopic data. The most commonly studied features have been Ca I H and K 0.38 $\\mu$m, H$\\beta$, Mgb 0.52 $\\mu$m, Fe $\\mu$m 0.53, Na 0.82 $\\mu$m, and CaT 0.86 $\\mu$m. Interpretation of all such spectral features is intrinsically complicated by their blended nature -- each feature is really the super-position of many spectral lines, usually from several different elements, blurred together by the line-of-sight velocity dispersion within each galaxy. There is no way to overcome this problem -- it must simply be taken into account during analysis. As population synthesis models have become more sophisticated and digital stellar libraries more complete, this problem has become more tractable over time.  Another challenge arises from the composite nature of galaxies: each observed feature is the luminosity-weighted integrated sum of that feature from all stars in the observed line-of-sight.  Naturally, luminosity-weighted does not imply mass-weighted. A relatively small fraction of the mass can dominate the observed luminosity and mask the underlying stellar population (e.g. as happens during a starburst event within a pre-existing galaxy).  Even in relatively quiescent galaxies, light from stars at several important evolutionary stages contribute roughly equally to the observed spectral features between 0.4 -- 1 $\\mu$m range. Hence, a feature depth change could be due to (e.g.) a change near the (mostly) age-driven main-sequence turnoff or the (mostly) metallicity-driven red giant branch.  The details can become quite complicated, as illustrated by the long standing controversy about whether observed changes in Balmer line strength arise from the presence of younger main sequence stars, more metal-poor main sequence stars, or an extended horizontal giant branch (for recent discussions of this debate, see Maraston \\& Thomas 2000 and Trager et al. 2005). A similar controversy surrounds Na 0.82 $\\mu$m feature: is it driven by metallicity-driven red giant branch changes, initial mass function related differences in the relative number of cool dwarf and giant stars or both \\citep[e.g.][]{car86, all89, del92}?  However, the properties of the RGB component can be isolated by observing in the K-band (centered near 2.2 $\\mu$m).  At those wavelengths, cool giants near the tip of the first-ascent red giant branch (RGB) dominate the integrated light in old ($\\geq$ 3 Gyr) stellar populations. In combination with optical observations, K-band observations should facilitate the separation of MSTO and RGB light contributions. There are two possible complications to this scenario. First, a very young stellar population containing red supergiants will contribute a significant fraction K-band light. Fortunately, such a population is obvious from the presence of H~II region emission lines at shorter wavelengths. Second, a somewhat older population (1 -- 2 Gyr, i.e. an intermediate-age population) may contain bolometrically bright carbon stars that can contribute a detectable amount of K-band light (see discussions in Silva \\& Bothun 1998a,b). Such a population may or may not be connected to increased H$\\beta$ strength.  Initial development of these ideas can be found in Silva et al. (1994), Mobasher \\& James (1996), James \\& Mobasher (1999), Mobasher \\& James (2000), all of whom focused on the CO 2.36 $\\mu$m feature. Origlia et al. (1997) observed a Si dominated feature at 1.59 $\\mu$m as well as CO dominated features. These observational studies were limited by small-format detectors to relatively low resolving powers and/or small wavelength ranges per observation. In the cases of Silva et al. and Origlia et al., only small, heterogeneous samples of galaxies were observed. A general conclusion of the James \\& Mobasher studies was that changes in CO strength between early-type galaxies in high-density and low-density regions were statistically consistent with different fraction contributions of intermediate-age AGB light and hence galaxies in low-density regions had younger luminosity-weighted mean ages. Origilia et al. argued that [Si/Fe] was super-solar in the four elliptical galaxies they observed.  To further develop these ideas and investigate the usefulness of other K-band spectral indices in the study of early-type galaxies, new data have been obtained for eleven E/S0 galaxies in the nearby Fornax cluster. Only measurements in the central regions of these galaxies are discussed here.  In Section~\\ref{sec:data}, the galaxy sample and its observations are discussed, while in Section~\\ref{sec:proc} the data processing methodology is described. The measurement of spectral feature strength is explained in Section~\\ref{sec:lines} while basic observation results are presented in Section~\\ref{sec:results}. The broader astrophysical implications of our observational results are discussed in Section~\\ref{sec:disc}. A summary is provided at the end. ",
        "conclusions": "\\label{sec:disc}  No self-consistent theoretical spectral synthesis models for the interpretation of integrated K-band spectra of early-type galaxies are widely available. Development of such models is on-going (e.g. Marmol Queralt, 2007, in preparation). Until such models exist, we must rely on basic (and perhaps imperfect) astrophysical intuition and available tools, such as the high-resolution near-IR spectral atlas of Wallace \\& Hinkle (1996) (hereafter WK96).  \\subsection{Galaxies with young populations}  Three galaxies in our sample (NGC 1316, NGC 1344, and NGC 1375) have stronger \\hb\\/ and weaker \\mgfe\\/ features than the rest of our sample (see Figure~\\ref{fig:fornax_pops}). This suggests a simple, two component stellar population model. One stellar component is cold -- in the luminosity-weighted mean, it is presumably old ($>$ 8 Gyr) and metal-rich ([Fe/H] $>$ --0.3) with spectral properties consistent with the observed central velocity dispersion. In other words, it has properties similar to the purely old galaxies in our sample (see Figure~\\ref{fig:fornax_pops}). The other component is warm -- it contains a significant number of A/F dwarf or sub-giant stars. This component could have several origins but only two are considered here. Each has different consequences for K-band spectral index behavior.  First, the warm component could be associated with a young population with a warm ($\\sim$2 \\msun) main sequence turnoff (MSTO). The MSTO stars are tied to thermally pulsating asymptotic giant branch (TP-AGB) stars with bolometric magnitudes that place them above the tip of the first-ascent red giant branch (TRGB). The TP-AGB stars are cooler than stars on the first-ascent RGB and hence have an M spectral type. In the underlying luminosity function (spectral type vs. number of stars), the M-star bins are relatively more populated than in a purely old galaxy. The net effect is that integrated 2.2 $\\mu$m spectra should become more M-like.  Second, the warm component could be associated with a metal-poor population with a warm MSTO. These MSTO stars are tied to RGB stars with similar luminosity as the corresponding metal-rich RGB stars but with warmer \\teff. In the underlying luminosity function, the K-star bins will be relatively more populated than in a purely old population. The net effect is to make the integrated 2.2 $\\mu$m spectra more K-like.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=0]{f16.ps}} \\caption{\\label{fig:galModel} Optical-NIR corrections for young populations.  See Figure~\\ref{fig:galLineSigma} for a symbol explanation. The observed index strengths for the galaxies with young components are connected to the predicted index strengths ({\\it solid squares}) after the effects of a young component has been removed. In the lower panel, the {\\it slanted dashed line} indicates the unweighted linear regression fit for the purely old galaxies with $\\sigma >$ 150 \\kms. In the middle panels, the {\\it horizontal dashed line} indicates the unweighted mean value for purely old galaxies with $\\sigma >$ 150 \\kms.} \\end{figure}  Can the observations presented here distinguish between these two scenarios?  Consider Figure~\\ref{fig:galModel} -- the galaxies with relatively strong \\hb\\/ and weak \\mgfe\\/ have relatively stronger (i.e. more M-like) near-IR features.  Qualitatively, this is consistent with the first scenario. Can this conclusion be better quantified? What is the relative ratio of young to old stellar mass? Are the changes in optical and near-IR spectral features {\\it quantitatively} consistent?  To begin to answer these questions, the Thomas et al. (2003) models were used to create very simple two-component models that produced optical index values that matched the observed values in the three galaxies with strong \\hb. One component was old (11 Gyr) with metallicity set to [Z/H] $=$ 0.35 for NGC 1316 and NGC 1344 and 0.0 for NGC 1375. The other component had the same metallicity but a young (1 Gyr) age. The relative mass fractions of these components were varied until the model \\hb\\/ strength matched the observed \\hb\\/ strength. The resultant young mass fractions (where total mass $=$ 1) were 0.135, 0.075, and 0.140 for NGC 1316, NGC 1344, and NGC 1375, respectively. The implied differential correction between the purely old model and the two-component model was then applied to the observed data (see Figure~\\ref{fig:galModel}).  Next, observed \\fe\\/ strength was adjusted manually until all three galaxies lay within the locus of purely old galaxies with similar $\\sigma$ in the \\fe\\/ vs. \\naK\\/ panel. As part of this adjustment, $\\Delta$(\\fe)/$\\Delta$(\\naK) was forced to agree with the value (0.62) determined for the observed Galactic cluster giant stars (see Figure~\\ref{fig:starLineLine}). This is a trend in effective temperature that is relatively insensitive to metallicity. Cooler giant stars (more late K and early M like) have strong K-band spectral features.  In addition to the $\\Delta$(\\fe)/$\\Delta$(\\naK) measured from the Galactic open cluster stars, two other key numerical trends can be computed:  \\begin{eqnarray} \\Delta(\\fe)/\\Delta(\\naK)   & = & 0.62 \\\\ \\Delta(\\naK)/\\Delta(\\mgfe)  & = & 0.70 (1.10) \\\\ \\Delta(\\fe)/\\Delta(\\mgfe) & = & 0.28 (0.44) \\end{eqnarray}  NGC 1316 and NGC 1344 could be forced to have consistent trajectories. However, NGC 1375 (values shown in parenthesis) appears to follow somewhat different trajectories. Obviously, this kind of cartoon model is illustrative only and surely hides a plethora of details. Indeed, the real trajectories are unlikely to be linear. Nevertheless, a clear astrophysical conclusion emerges: for the \\hb-strong galaxies, the observed optical and near-IR features all change in concert, consistent with a warm MSTO tied to an extended RGB, which is in turn consistent with the presence of a young stellar component, not a metal-poor stellar component.  \\subsection{Galaxies with bright K-band SBF}  Liu, Graham, \\& Charlot (2002) found that relative I-band and K-band surface brightness fluctuation strength (expressed as an SBF color, $\\bar{I} - \\bar{K}$) was essentially constant for the majority of the early-type galaxies they studied in Fornax. However, two galaxies in common with our sample (NGC 1419 and 1427) were found to have brighter (larger) K-band SBF relative to the mean relationship (see also Mei et al. 2001). A larger $\\bar{K}$ is thought theoretically to arise from the presence of an extended giant branch, i.e. more cool, bright, M giants.  In the spectroscopic data discussed here, there are no indications of such extended giant branches in these galaxies -- their central line indices are consistent with an integrated stellar populations dominated by an old, metal-rich population. Of course, the SBF measurements were made in the outer regions of these galaxies, as opposed to the optical and near-IR observations of the central regions discussed here. It may be that the integrated luminosity of such a component (if present at all) is not high enough in the central region for detection by our method.  \\subsection{Purely old galaxies}  In purely old galaxies, we have seen that: \\begin{itemize}  \\item{} \\naK\\ is stronger than in Galactic open clusters and is highly correlated with $\\sigma$ and \\mgfe.  \\item{} \\fe\\/ saturates for $\\sigma \\gtrsim 150$ \\kms.  \\item{} \\co\\ is somewhat correlated with $\\sigma$ and \\mgfe, while \\caK\\/ is not (unless NGC 1399 is not included in the regression fits, see footnote 2 above).  \\end{itemize}  The observed K-band spectral features are named for the dominant elemental species in the {\\it solar} spectrum. However, at the effective temperatures of interest here, the on-band index definition and their companion off-band continuum bands contain lines from other elemental species as well (as Ramirez et al.~1997 discuss comprehensively). By referring to the high resolution spectra of WK96, the contribution from these other absorbers can be investigated and used to explain the overall spectral feature behavior. In turn, underlying astrophysical parameters related to galaxy formation and evolution are revealed.  \\subsubsection{\\naK\\/ index}  In the current galaxy sample, \\naK\\/ is significantly stronger than observed in the Galactic open cluster stars.  The off-bands defined for the \\naK\\/ feature are relatively line-free, so we can focus on the on-band spectral region (see the $\\lambda$ Dra, M0 III spectrum in WK96, p. 352). In addition to sodium, Sc, Si, and (to a lesser extend) V are significant absorbers in this region. Origlia et al. (1997) argued that [Si/Fe] is super-solar in a small sample of early-type galaxies based in measurements of the Si I 1.59 $\\mu$m feature. Trager et al. (2000) have argued that both [Na/Fe] and [Si/Fe] are super-solar in early-type galaxies based on models of optical spectral feature behavior. The K-band \\naK\\/ measurements presented here are consistent with enhanced silicon in the observed Fornax galaxies and hence indirectly confirm the conclusions of Origlia et al. and Trager et al.  \\subsubsection{\\fe\\/ index}  The \\fea\\/ feature contains approximately equal absorption contributions from Fe, Sc and Ti while \\feb\\/ is dominated by Fe absorption features with some additional Sc absorption (in WK96, see the $\\lambda$ Dra, M0 III spectra on p. 346 and 348). Hence, the \\fe\\/ index (a combination of \\fea\\/ and \\feb) is dominated by absorption lines from Fe-peak elements). For purely old galaxies with $\\sigma >$ 150 \\kms, \\fe\\/ is observed to remain constant while \\naK\\/ (an index dominated by absoprtion lines from $\\alpha$-elements) and \\mgfe\\/ (a total metallicity [Z/H] indicator) become stronger as $\\sigma$ increases. It is tempting to conclude that above $\\sigma >$ 150 \\kms\\/ luminosity-weighted mean total metallicity [Z/H] continues to increase as a function of central velocity dispersion, driven by a relative increase of [$\\alpha$/H], while [Fe/H] remains constant.  However, within the narrow [Fe/H] range of observed Galactic open cluster stars, \\fe\\/ is more strongly correlated with effective temperature than cluster [Fe/H] (see Figure~\\ref{fig:starLineColor}). Increasing mean total metallicity [Z/H] in the galaxies should correspond to cooler mean RGB effective temperature and hence increased \\fe. Recall that the relationship between \\fe\\ and mean RGB effective temperature has already been exploited above to explain the observed behavior of \\fe\\/ (and \\naK) in galaxies with young stellar components. The dilemma is clear: how can a temperature-sensitive index like \\fe\\/ remain constant while mean RGB effective temperature decreases due to increasing [Z/H]?  Without appropriate population synthesis models or stellar spectra, only a few speculative thoughts can be offered. In their detailed comparison of non-solar aboundance population synthesis models with measurements of spectral indices in the optical Lick system, Trager et al. (2000) concluded that their best-fit models included enhancements in C, N, Na, and Si (among others). As [$\\alpha$/H] increases with $\\sigma$ (as suggested by increasing \\naK), stronger CN bands in the \\fea\\/ and \\feb\\/ on-band and off-band windows (see the WK96 spectra referenced above and the band definitions in Table~\\ref{tab:features}) may have the effect of depressing the local continuum and hence decreasing the value of \\fe\\/ $=$ (\\fea\\/ $+$ \\feb)/2. For example, Trager et al. noted a similar effect from C$_{\\rm 2}$ bands in the optical Mg {\\it b} index. Therefore, it may be that above some critical mean [$\\alpha$/H] and below some critical mean effective temperature, \\fe\\/ stays within a narrow range and provides no useful information about mean [Z/H], [Fe/H], or effective temperature.  A conclusive astrophysical explanation of \\fe\\/ in early-type galaxies with $\\sigma >$ 150 \\kms\\/ awaits larger galaxy samples, observations of appropriate stars (e.g. in the Galactic bulge), and detailed population synthesis models.   \\subsubsection{\\co\\ index}  The \\co\\/ feature is dominated by the $^{\\rm 12}$CO(2,0) bandhead and has no significant absorption components from other elements. Like \\fe\\/ above, \\co\\/ is highly correlated with $T_{eff}$ in the observed Galactic open cluster stars and reaches larger values than observed in the galaxies. As a function of central velocity dispersion $\\sigma$ in the observed galaxies,  \\co\\/ does not appear to saturate (or at least not as definitively as \\fe). This adds credence to the suspicion that \\fe\\ is not tracing mean RGB effective temperature in galaxies with $\\sigma >$ 150 \\kms.  Based on the current measurements, no conclusions can be reached about relative carbon or oxygen abundance in these galaxies.  \\subsubsection{\\caK\\ index}  No obvious explanation presents itself for the behavior of \\caK\\/ in this galaxy sample.  The \\caK\\/ feature is intrinsically complex, consisting of contributions from many absorbers: Ca, S, Si, and Ti (all $\\alpha$-elements) as well as Sc and Fe (see $\\lambda$ Dra, MO III spectrum in WK96, p. 342). As a function of \\teff\\/ in the range of interest, some of the contributing absorption lines get stronger (Sc and Ti, the former faster than the latter), some features get weaker (Si), and some remain roughly constant (Ca, S, and Fe) (cf. Ramirez et al. 1997). Relative to solar abundance ratios, Trager et al. (2000) have argued that Si and S are over-abundant and Ca is under-abundant in early-type galaxies while [Ti/Fe] and [Sc/Fe] have their solar values. In the Galactic open cluster stars observed here, the \\caK\\/ feature is stronger at a given $J-K$ ($T_eff$) in the solar-metallicity stars than in the more metal-poor stars.  Within the galaxies observed here, \\caK\\/ appears to have significant scatter at any given $\\sigma$ or \\mgfe.  No obvious explanation for this scatter (or the global behavior of \\caK\\ with $\\sigma$ and \\mgfe) presents itself at this time.    Using new, moderate resolution ($R \\sim 2500$) K-band spectra, spectral indices have been measured in the central regions of eleven early-type galaxies in the nearby Fornax cluster. Based on these measurements, the following conclusions were reached:  \\begin{enumerate}  \\item{} The \\naK\\/ feature is much stronger in these Fornax early-type galaxies than observed in solar-metallicity Galactic open cluster stars. This is attributed to relative [Si/Fe] (and possible [Na/Fe]) differences between the open cluster stars and the Fornax galaxies, i.e. both are larger in the Fornax galaxies than in the cluster stars.  \\item{} In various near-IR diagnostic diagrams, galaxies with optical indices indicative of a warm stellar component are clearly separated from galaxies dominated by colder, presumably old ($\\geq$ 8 Gyr) stellar populations. Changes in the near-IR spectra features are consistent with the presence of an cool component dominated by late K and/or early M giants stars. In combination, the optical and near-IR observations are consistent with the presence of a young stellar component with a warm MSTO and a significant extended giant branch consisting of TP-AGB stars.  \\item{} For detecting a young stellar component, the \\naK\\/ vs.~$\\sigma$ or \\fe\\/ vs.~$\\sigma$ diagnostic diagram seems as efficient as using H$\\beta$ vs. \\mgfe\\/ (or other similar combinations of optical indices).  The near-IR features have the additional advantage that no emission-line correction is needed (as it is necessary for the stronger optical Balmer lines).  \\item{} The \\fe\\/ index saturates in galaxies with central velocity dispersion $\\sigma$ $>$ 150 \\kms\\/ dominated by old ($\\geq$ 8 Gyr) stellar populations. For $\\sigma >$ 150 \\kms, these Fe features are unlikely to be useful for investigating stellar population differences between early-type galaxies. Above $\\sigma >$ 150 \\kms, the continued increase in \\naK\\ (and \\mgfe) strength with $\\sigma$ coupled with constant \\fe\\ presents an astrophysical challenge. Although it is tempting to conclude that [Fe/H] reaches a maximum value, while [$\\alpha$/H] (and hence [Z/H]) continues to increase, more observational and population synthesis work is needed to understand conclusively the behavior of \\fe\\/ in these high-mass early-type galaxies.  \\end{enumerate}  Adding these near-IR indices to the standard diagnostic toolkit for analyzing the integrated light of early-type galaxies clearly has great potential. To develop this potential, three obvious steps are needed: observe more galaxies over a larger range of central velocity dispersion (including field galaxies with existing optical data) and extend current population synthesis models to the near-IR. We know that various groups are working on both steps.  It would also be useful to study the radial behavior of these indices within individual galaxies to compare and contrast index behavior between galaxies.  For the foreseeable future, the study of the central populations in early-type galaxies will remain the study of integrated light. As distance increases, our ability to study the central regions of early-type galaxies within metric apertures equivalent to nearby galaxies relies on the high {\\it spatial} resolution spectroscopy achievable from space or with adaptive optics on the ground. In the former case, the James Webb Space Telescope represents the frontier -- yet, no optical spectrograph that works below 0.8 $\\mu$m is currently planned. In the latter case, the frontier will be shaped by further development of adaptive optics systems -- and these systems will achieve their best performance beyond 1 $\\mu$m. Hence, understanding how to interpret near-IR spectral indices in nearby galaxies is key to facilitating the kind of investigation and characterization of more distant early-type galaxies that will not be possible at optical wavelengths."
    },
    "0710/0710.4059.txt": {
        "abstract": "The turbulent magnetic diffusivity tensor is determined in the presence of rotation or shear. The question is addressed whether dynamo action from the shear--current effect can explain large-scale magnetic field generation found in simulations with shear. For this purpose a set of evolution equations for the response to imposed test fields is solved with turbulent and mean motions calculated from the momentum and continuity equations. The corresponding results for the electromotive force are used to calculate turbulent transport coefficients. The diagonal components of the turbulent magnetic diffusivity tensor are found to be very close together, but their values increase slightly with increasing shear and decrease with increasing rotation rate. In the presence of shear, the sign of the two off-diagonal components of the turbulent magnetic diffusion tensor is the same and opposite to the sign of the shear. This implies that dynamo action from the shear--current effect is impossible, except perhaps for high magnetic Reynolds numbers. However, even though there is no alpha effect on the average, the components of the $\\alpha$ tensor display Gaussian fluctuations around zero. These fluctuations are strong enough to drive an incoherent alpha--shear dynamo. The incoherent shear--current effect, on the other hand, is found to be subdominant. ",
        "introduction": "Many of the stellar and planetary magnetic fields are believed to be the result of a dynamo process that converts kinetic energy from turbulent motions and shear into magnetic energy. A particular challenge consists in explaining the field on length scales that exceed the scale of the turbulence. This topic has traditionally been addressed within the framework of mean--field electrodynamics (Krause \\& R\\\"adler 1980).  Over the decades the applicability of this theory has repeatedly been questioned (e.g., Piddington 1981, Vainshtein \\& Cattaneo 1992). Meanwhile, direct simulations of hydromagnetic turbulence have begun to show dynamo action (Meneguzzi et al.\\ 1981, Meneguzzi \\& Pouquet 1989, Nordlund et al.\\ 1992, Brandenburg et al.\\ 1996, Cattaneo 1999). In some particular cases, large-scale fields are being generated (Glatzmaier \\& Roberts 1995, Brandenburg et al.\\ 1995, Brandenburg 2001) which raises the question about the mechanism responsible for this phenomenon. In cases where the flow is systematically non-mirror symmetric the association with an $\\alpha$ effect is obvious. However, there are now also examples of nonhelical large-scale dynamos owing to turbulence under the influence of shear alone (Brandenburg 2005a, Yousef et al.\\ 2007). Their interpretation is not straightforward, because several possible mechanisms have been proposed that might produce dynamo action from turbulence and shear alone, i.e.\\ {\\it without} rotation and stratification that otherwise would have been the main ingredients of an $\\alpha$ effect. The most detailed investigations have been carried out in connection with the so-called shear--current effect (Rogachevskii \\& Kleeorin 2003, 2004, R\\\"adler \\& Stepanov 2006, R\\\"udiger \\& Kitchatinov 2006). Another possibility is a magnetic $\\alpha$ effect that is driven by a current helicity flux, as was suggested by Vishniac \\& Cho (2001; see also Brandenburg \\& Subramanian 2005c). A third possibility might be an incoherent (random) $\\alpha$ effect with zero mean and finite variance, suggested by Vishniac \\& Brandenburg (1997) in connection with accretion discs (see also Sokolov 1997, Silant'ev 2000, Fedotov et al.\\ 2006, Proctor 2007). The only reliable way to determine what is the dominant effect is to calculate all relevant components of the $\\alpha$ and turbulent magnetic diffusivity tensors in a general expansion of the electromotive force in terms of the mean magnetic field.  The case considered in Brandenburg (2005a) is unnecessarily complicated because the shear employed there depends on two Cartesian coordinates. A simpler possibility is to consider a shear flow depending linearly on only one coordinate and we shall pursue this idea in the present paper. The shear--current effect and the incoherent $\\alpha$ effect could then still operate. Because we will use periodic boundary conditions there can be no magnetic helicity flux, so the Vishniac \\& Cho (2001) effect is then ruled out, even though it could still, at least in principle, explain the generation of a mean magnetic field in the simulations of Brandenburg (2005a), which do possess a helicity flux.  In this paper we calculate all relevant components of $\\alpha_{ij}$ and $\\eta_{ijk}$ using the so-called test field method. This method was introduced by Schrinner et al.\\ (2005, 2007) in connection with convection in a spherical shell and used later by Brandenburg (2005b), Sur et al.\\ (2007) and Brandenburg et al.\\ (2008) in connection with forced turbulence in Cartesian boxes. The essence of this method consists in solving evolution equations for the fluctuations of the magnetic field around suitably defined test fields such that all relevant coefficients can be computed. ",
        "conclusions": "The present work has demonstrated that the test field method provides a robust means of determining all components of the turbulent magnetic diffusivity tensor that are relevant for mean fields depending only on $z$ and $t$. Both rotating and  weakly shearing turbulence are studied. In either case the diagonal components of the turbulent diffusivity tensor are about equal to each other. Shear slightly enhances the turbulent magnetic diffusivity while rotation quenches it. In the presence of rotation, the $\\OO\\times\\meanJJ$ effect occurs, which is described by the off-diagonal components of the turbulent magnetic diffusivity tensor. Shear leads to the shear--current effect, again described by off-diagonal components of this tensor. In both cases the results are consistent with those found in the framework of the second-order correlation approximation.  The possibility of the so-called shear-current dynamo has been scrutinized. It depends crucially on the sign of the component $\\eta_{21}$ of the magnetic diffusivity tensor. It turns out that, within the ranges of parameters considered, its sign is in general not suited for driving a dynamo based on this effect, with a possible exception at large magnetic Reynolds numbers. In this way the analytic results found in the second--order correlation approximation for incompressible fluids (R\\\"udiger \\& Kitchatinov 2006; R\\\"adler \\& Stepanov 2006) are confirmed and generalized.  Direct numerical simulations are presented which exhibit growing mean magnetic fields in shear flow turbulence. An interpretation as a (coherent) shear-current dynamo is hardly possible. Instead, it is argued that it can be explained by an incoherent alpha--shear dynamo. The incoherent shear--current effect has also been determined, but it is found to be less important."
    },
    "0710/0710.0760_arXiv.txt": {
        "abstract": "{ This series of papers is dedicated to a new technique to select galaxies that can act as standard rods and standard candles in order to perform geometrical tests on large samples of high redshift galaxies to constrain different cosmological parameter. The goals of this paper are (1) to compare different rotation indicators in order to understand the relation between rotation velocities extracted from observations of the H${\\alpha} \\lambda$6563\\AA\\ line and the [OII]$\\lambda 3727$\\AA\\ line, and (2) to determine the scaling relations between physical size, surface brightness and magnitude of galaxies and their rotation velocity using the SFI++, a large catalog of nearby galaxies observed at I-band. A good correlation is observed between the rotation curve-derived velocities of the H${\\alpha}$ and [OII] observations, as well as between those calculated from velocity histograms, justifying the direct comparison of velocities measured from H${\\alpha}$ rotation curves in nearby galaxies and from [OII] line widths at higher redshifts. To provide calibration for the geometrical tests, we give expressions for the different scaling relations between properties of galaxies (size, surface brightness, magnitude) and their rotation speeds.  Apart from the Tully-Fisher relation, we derive the size-rotation velocity and surface brightness-rotation velocity relations with unprecedentedly small scatters. We show how the best size-rotation velocity relation is derived when size is estimated not from disc scale lengths but from the isophotal diameter $r_{23.5}$, once these have been corrected for inclination and extinction effects.}   \\titlerunning{Geometrical Tests of Cosmological Models II.} \\authorrunning{Saintonge et al.} ",
        "introduction": "With the new generation of telescopes and instruments, such as multi-object spectrograph and large field of view cameras, it has become possible to perform large, deep redshift surveys.  Using these data, the structure and geometry of the Universe can be studied through new channels, namely geometrical tests such as the angular diameter test (angular size - redshift relation), the Hubble diagram (magnitude - redshift relation) and the Hubble test (count - redshift relation).  These tests are generally difficult, requiring differentiation between the natural evolution of galaxies and the effects of geometry.  Large redshift surveys ease the task in that respect.  To perform the geometrical tests mentioned above, a population of objects that can be tracked through redshift needs to be defined.  More specifically, the angular diameter test requires a standard rod to be identified, the angular size of which is then being measured over the redshift range of interest.  The objects taken to serve that purpose could be anything from galaxies \\citep[e.g.][]{sandage72}, clusters of galaxies \\citep[e.g.][]{sereno,cooray98}, or dark matter halos \\citep[]{cooray01}.  Regardless of its nature, a standard rod should satisfy simple criteria: its size should be measured reliably up to high redshifts, and it should be observable in the local Universe to provide calibration.  In this study, the relation between the physical size of a spiral galaxy and its rotation velocity (or global profile width) is used.  By virtue of this relation, selecting a population of objects with a given rotation velocity is equivalent to selecting them based on their physical size.  In other papers of this series \\citep[][thereafter Paper I and Paper III]{paperI,paperIII}, the angular diameter test is performed on a dataset of precursor Vimos/VLT Deep Survey (VVDS) \\citep[]{vvds} galaxies, using the linewidth-diameter relation to select standard rods.  The test is used to isolate the effects of disk evolution with redshift, and to constrain the values of cosmological parameters.  Though the VVDS provides excellent data at high redshift, the angular diameter test highlights one of its shortcomings: limited volume coverage at low redshift and therefore the lack of a comparison sample in the local Universe.  In this context, the goal of this paper is twofold.  First, using a large sample of nearby galaxies \\citep[SFI++;][]{spr07} the physical properties of the standard rods and standard candles used to perform the geometrical tests are established, free from any evolutionary biases.  This is especially useful to calibrating the angular diameter relation and to separating the effects of galaxy evolution from those of geometrical variations (see Paper I).  By establishing clearly the scaling relations between the size, surface brightness and magnitude of these nearby galaxies observed at I-band and their rotation velocity, we establish the calibration that will allow us to select standard rods/candles simply by tracing through redshift galaxies with a given rotation speed. Having a strong handle on the structural parameters of disc galaxies in the local universe gives the unique opportunity of tracing over time the evolution of these properties for disc galaxies hosted in dark matter halos of the same mass, if rotation velocity is used as a proxy for halo mass, and compare them with predictions made under the hierarchical scenario for the growth of structures \\citep[e.g.][]{mo98}.  The second goal of this paper is to compare the rotation velocity indicator used for the high redshift data to those used in the local Universe, cross-calibrating rotation indicators used at different redshifts. This last point is of relevance since rotation information for spiral galaxies can be obtained through the observation of various spectral lines, the choice of which varies with the redshift of the sample.  For galaxies with $z \\lesssim 0.1$, the HI 21 cm line is an excellent candidate.  The H${\\alpha} \\lambda$6563\\AA\\ line is also frequently used for galaxies with low to moderate redshifts.  However, the H${\\alpha}$ line is quickly redshifted into the near-infrared and becomes unavailable to ground observers using optical telescoptes for galaxies with $z \\gtrsim 0.4$, even though with new near-infrared spectrographs such as SINFONI it is now possible to obtain H$\\alpha$ rotation curves for high redshift galaxies \\citep[e.g.][]{forster06}.  Even with the advent of such instruments, most large studies of galaxies at high redshift rely on the [OII]$\\lambda 3727$ \\AA\\ line, including the VVDS and the DEEP2 Redshift Survey \\citep[]{davis01}.  In order to compare sets of local and distant galaxies, it is therefore necessary to understand how rotation velocities extracted from these different lines relate.  It has already been shown that velocity widths derived from HI 21cm and H${\\alpha}$ observations are in excellent agreement \\citep[see for example][]{courteau97,vogt}.  The correlation between the [OII] and HI line widths has also been investigated \\citep[]{kobulnicky}.  Using a sample of 22 nearby late-type spirals they find [OII] widths to be accurate to within $10 \\%$ for galaxies with a roation width of 200 km s$^{-1}$, which is comparable to the overall scatter in the local Tully-Fisher relation.  However, they conclude that the uncertainties go up to about 50\\% for galaxies with widths $<150$ km s$^{-1}$.  Here, a sample of 32 spiral galaxies with $0.155<z<0.25$ is used to compare velocity widths from the H${\\alpha} \\lambda$6563\\AA\\ and [OII]$\\lambda 3727$ \\AA\\ lines (hereafter, H${\\alpha}$ and [OII]).  Even though these galaxies are more distant than in the \\citet[]{kobulnicky} sample, the scatter in the rotation curves data points is significantly smaller, resulting in a better estimation of the rotation velocities.   Our data also allows for a direct comparison between H${\\alpha}$ and [OII], therefore bridging nicely the gap between low and high redshift samples.  Contrary to the \\citet[]{kobulnicky} study, the galaxy sample used here is restricted to late spiral types, therefore we will not address the question of possible biasses due to varying morphologies.  In \\S \\ref{data} a description of the galaxies observed and of the data reduction process is given, while in \\S \\ref{widths} we describe how velocity widths were extracted from the data and how they compare for both sets of emission lines.  In \\S \\ref{localtf} we presente a sample of nearby galaxies used as calibrators for the purpose of the angular diameter test (Paper III) and discussion of our results are given in \\S \\ref{discussion}.  Calculations were done assuming $H_0=70$\\ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M=0.3$, $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "}  The main goal of this paper was to provide calibration in two different ways for the angular diameter test that we performed on data from the VVDS: (1) by directly comparing velocity indicators commonly used in the local and distant Universe, and (2) by using a large sample of nearby galaxies to provide zero points for the radius - rotation velocity, magnitude - rotation velocity and surface brightness - rotation velocity relations.  Previous studies have focused on comparing the H${\\alpha}$ or [OII] rotation velocities with HI data for nearby galaxies.  \\citet[]{kobulnicky} have used a sample of 22 nearby galaxies to compare [OII] and HI linewidths.  Their galaxy sample was smaller and less homogeneous that the one used in this study.  They found the two sets of velocity estimates to be consistent within $10\\%$ overall, with the [OII] underestimating the HI rotation velocity by up to $50\\%$ in a few extreme cases.  In the local Universe, the consistency of the HI and H${\\alpha}$ rotation velocities was also established \\citep[see for example][]{courteau97,vogt}.  Most recently, \\citet[]{catinella} has shown that on average HI widths are larger than H${\\alpha}$ ones by $\\sim10\\%$ for a galaxy with $w=100$ km s$^{-1}$ and that optical rotational widths measured from velocity histograms are affected by systematic biases and therefore are less reliable that those derived by using the full spatial information through the fitting of the rotation curves.  For this study, we first analysed the rotation curves for all the galaxies in our sample and concluded that the rotation velocities derived from the [OII]$\\lambda 3727$ \\AA\\ line translate directly into the velocities obtained from H${\\alpha}$ rotation curves, given that the sampling of the [OII] rotation curve extends far enough to reach the region of constant (or slowly increasing) velocity.  Even under this condition and as shown in Figure \\ref{vrot_plot}, there is scatter in the relation of about 30 km s$^{-1}$.  The next step was to compare the techniques used to determine rotation velocities in the local Universe (H${\\alpha}$ rotation curves) and at high redshift ([OII] velocity histograms).  Once again, there is a direct relation between the two but significant scatter on the order of 50 km s$^{-1}$ (see Figure \\ref{sigma_plot}).  We conclude that all the rotation velocity indicators studied here give comparable results, but before combining samples analysed with different methods one should be aware of the important scatter in the different relations.  This scatter is due mostly to the poorer spatial extent of the [OII] rotation profiles compared to the H${\\alpha}$ ones obtained with the same integration time, and can be as large as $25\\%$ for a galaxy with $w=200$ km s$^{-1}$ or $100\\%$ if $w=50$ km s$^{-1}$ in the case of the $\\sigma_{rot}$([OII])-$v_{rot}$(H${\\alpha}$) relation (Fig.\\ref{sigma_plot}).  Even considering this large scatter, the fact that there is no systematic deviation from a 1:1 relation allows us to use both methods in their respective context to divide galaxies in classes based on rotation velocity, as required to perform the geometrical tests with the high redshift sample (see Paper III).  In Paper I it was showed that this is a viable way of selecting standard rods for the purpose of the angular diameter test.  Since one cannot perform the test only with high redshift data, it was of great importance to establish that a set of local calibrators can be reliably compared to it, even if the rotation measure is different.  Once it is established that the various rotation indicators can be reliably compared, the next critical step is to derive the scaling relations between luminosity, size, surface brightness and rotation velocity for a set of nearby galaxies, free from any evolution. We used the SFI++ catalog, a very large homogeneous sample of nearby spiral galaxies with I-band photometry and rotation information obtained either through optical long-slit spectroscopy or HI measurements. We have particularly focussed our attention on the $r-v_{rot}$ relation which is central to the cosmological tests performed in the other papers of this series. We have shown that by carefully correcting sizes for inclination and extinction effects as well as obtaining accurate distances to the galaxies, the $r_{23.5}-v_{rot}$ relation can be almost as tight as the Tully-Fisher relation, as judged by their Pearson correlation coefficients.  The relation derived using $r_{23.5}$ as the size measurement is considerably tighter than the ones obtained using the disc scale length $r_d$ or a radius defined to contain a given fraction of the light of the galaxy (in our case $r_{83}$ or $r_{50}$).  The next phase for this study on scaling relations will be to see what constraints our $r-v_{rot}$ relation imposes on models for disc galaxy formation, for example the range of spin parameters its tight scatter allows. These scaling relations also offer exciting possibilities for the study of disc galaxy evolution. Since they provide accurate values for luminosity, size and surface brightness as a function of rotation velocity, measurements of the properties for galaxies with the same rotation velocity (i.e. hosted in a dark halo of the same mass) over a large redshift range can provide in a new and unique way information about the evolution of galaxies and a mean of testing predictions from models under the hierarchical scenario for the growth of structures. A pilot study is presented in the next paper of this series."
    },
    "0710/0710.3763_arXiv.txt": {
        "abstract": "{} {To measure the supernova (SN) rates at intermediate redshift we performed the Southern inTermediate Redshift ESO Supernova Search (STRESS). Unlike most of the current high redshift SN searches, this survey was specifically designed to estimate the rate for both type Ia and core collapse (CC) SNe.} {We counted the SNe discovered in a selected galaxy sample measuring SN rate per unit blue band luminosity. Our analysis is based on a sample of $\\sim 43000$ galaxies and on 25 spectroscopically confirmed SNe plus 64 selected SN candidates. Our approach is aimed at obtaining a direct comparison of the high redshift and local rates and at investigating  the dependence of the rates on specific galaxy properties, most notably their colour.} {The  type Ia SN rate, at mean redshift $z=0.3$, amounts to 0.22$^{+0.10 +0.16}_{-0.08 -0.14}$ $h_{70}^2$ SNu, while the CC SN rate, at $z=0.21$, is 0.82$^{+0.31 +0.30}_{-0.24 -0.26}$ $h_{70}^2$ SNu. The quoted errors are the statistical and systematic uncertainties.  } {With respect to local value, the CC SN rate at $z=0.2$ is higher by a factor of $\\sim 2$ already at redshift , whereas the type Ia SN rate remains almost constant.  This implies that a significant fraction of SN~Ia progenitors has a lifetime longer than $2-3$ Gyr. We also measured the SN rates in the red and blue galaxies and found that the SN~Ia rate seems to be constant in galaxies of different colour, whereas the CC SN rate seems to peak in blue galaxies, as in the local Universe.  SN rates per unit volume were found to be consistent with other measurements showing a steeper evolution with redshift for CC SNe with respect to SNe~Ia.  Finally we have exploited the link between SFH and SN rates to predict the evolutionary behaviour of the SN rates and compare it with the path indicated by observations. We conclude that in order to constrain the mass range of CC SN progenitors and SN Ia progenitor models it is necessary to reduce the uncertainties in the cosmic SFH. In addition it is important to apply a consistent dust extinction correction both to SF and to CC SN rate and to measure SN~Ia rate in star forming and in passive evolving galaxies in a wide redshift range. } ",
        "introduction": "We now detail how all the ingredients described in the previous sections have been combined to compute the control time of our galaxy sample. First, for a given SN type and filter $F$, the control time of the $i$-observation of the $j$-galaxy was computed as:  \\begin{equation} CT^{\\rm SN,F}_{j,i} = \\int{\\tau^{\\rm SN,F}_{j}(m) \\,\\epsilon^{\\rm F}_{i} (m)\\,{\\rm d}m} \\end{equation} where $\\tau^{\\rm SN,F}_{j}(m)$ is the time spent by the SN in the magnitude range $m$ and $m+{\\rm d}m$ and $\\epsilon^{\\rm F}_{i} (m)$ is the detection efficiency. More specifically, we convolved the distribution of the absolute magnitude at maximum ($M^{\\rm SN}_{B}(0)$) and of the absorption due to host galaxy extinction ($A_B^h$), so as to determine the distribution of the quantity $M^{\\rm SN}_{B}(0)+A_B^h$ appearing in Eq.~\\ref{lc}, and computed $\\tau^{\\rm SN,F}_{j}(m)$ as a weighted average of the individual times over this combined distribution.  Then, the total control time $CT^{\\rm SN,F}_j$ of the $j$-galaxy was computed by summing the contribution of individual observations. If the temporal interval elapsed since the previous observation is longer than the control time, that contribution is equal to the control time of the observation, otherwise it is equal to the interval of time between the two observations \\citep{Cappellaro99}.  The total control time of the galaxy sample is obtained as the $B$ band luminosity weighted ($\\overline {CT}^{\\rm SN,F}_j$) average of the individual galaxies.  Since we merged all CC subtypes, including type Ib/c IIP and IIL, the control time for CC SNe is computed as follows:  \\begin{equation}\\label{eqctcc} \\overline{CT}^{\\rm CC}_j = f_{\\rm Ib/c}\\, \\overline{CT}^{\\rm Ib/c}_j +  f_{\\rm IIL}\\, \\overline{CT}^{\\rm IIL}_j +  f_{\\rm IIP}\\, \\overline{CT}^{\\rm IIP}_j \\end{equation}  where the relative fractions of the different CC subtypes is assumed to be constant with redshift and equal to that observed in the local Universe, namely 20$\\%$ of Ib/c and 80$\\%$ of II \\citep{Cappellaro99}, out of which 35$\\%$ are IIL and 65$\\%$ are IIP events \\citep{Richardson02}.  In order to illustrate the role of the host galaxy extinction, we calculated the CC SN rate by adopting the control time both for a {\\em standard} ($\\tau(0)=1.0$), and for a {\\em high} ($\\tau(0)=5.0$) extinction scenario. ",
        "conclusions": "\\label{Discussion}  In this section we discuss the evolution of the SN rates and investigate on the link between SF and SN rates. First we verify the consistency of our estimate of the redshift evolution of SN rates with other measurements from the literature  (Sec.~\\ref{Complit}). Then we compare the observed evolution of the CC SN rate with those expected for different SFHs (Sec.~\\ref{CompSFR}). Finally, we convolve the SFH of \\citet{Hopkins06} with the delay time distribution (DTD) for  different progenitor scenarios, and compare the results to the measurements of the SN~Ia rate evolution (Sec.~\\ref{Iamod}).  \\subsection{Comparison with other estimates}\\label{Complit}  Published measurements of the SN rates at intermediate and high redshift are expressed in units of co-moving volume.  To convert our rates from SNu to volumetric rates we assumed that the SN rates are proportional to the galaxy $B$ band luminosity. If this assumption is true, by multiplyng the rates in SNu by the total $B$ band luminosity density in a given volume we derived the total rates in that volume, even if we did not sample the faint end of the galaxy luminosity function. We accounted for the evolution of the $B$ band luminosity density with redshift. A compilation of recent measurements of the $B$ band luminosity density is plotted as function of redshift in Fig \\ref{lumdens}, where we also show a linear least-square fit to the data in the range $z=0-1$: $j_B(z) = (1.03+1.76\\times z)$ ~$10^8 L_{\\odot}^{B}  {\\rm Mpc}^{-3}$.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{lumden.eps}} \\caption{Measurements of the galaxy luminosity density at different redshifts. The DEEP2 and COMBO-17 data are taken from Table 2 in \\citet{Faber}. The line represents the linear least-square fit in the redshift interval $z=0-1$.} \\label{lumdens} \\end{figure}  Multiplying our measurements by the value of $j_B$ at the average redshifts of the Ia and CC SN samples, the rates per unit of co-moving volume result: \\begin{itemize} \\item $r^{\\rm Ia}(z=0.30) = 0.34^{+0.16 +0.21}_{-0.15 -0.22}$ $10^{-4} h_{70}^3 \\mbox{yr}^{-1} \\mbox{Mpc}^{-3}$\\\\ \\item$r^{\\rm CC}(z=0.21) = 1.15^{+0.43 +0.42}_{-0.33 -0.36}$ $10^{-4} h_{70}^3 \\mbox{yr}^{-1} \\mbox{Mpc}^{-3}$ \\end{itemize} where both statistical and systematic errors are indicated. Also local rates are converted into volumetric rates: $r^{\\rm Ia}(z=0.01) =0.18\\pm0.05$ $10^{-4} h_{70}^3 \\mbox{yr}^{-1} \\mbox{Mpc}^{-3}$ and $r^{\\rm CC}(z=0.01) = 0.43\\pm0.17$ $10^{-4} h_{70}^3 \\mbox{yr}^{-1} \\mbox{Mpc}^{-3}$. With respect to the rates in SNu's, the volumetric rates evolve more rapidly with redshift, due to the increase of the $B$ band luminosity density. We find an increase of a factor of $\\sim 2$ at $z=0.3$ for SNe~Ia, and a factor of $\\sim 3$ at $z=0.21$ for CC SNe.  Measurements of Ia and CC SN rate as function of redshift are shown in Fig.~\\ref{obsrate}, where those originally given in SNu \\citep{Hardin,Blanc,Cappellaro05} were converted into measurements per unit volume as above.  \\begin{figure*} \\resizebox{\\hsize}{!}{\\includegraphics{obs.eps}} \\caption{Observed SN rates as function of redshift from different authors as indicated in the legend. The black (gray) symbols indicate SN~Ia (SN~CC) rate measurements.  The shaded area represents the 1 $\\sigma$ confidence level of our rate evolution estimate as deduced from the MLE fit.} \\label{obsrate} \\end{figure*}  As it can be seen, the few measurements of the CC SN rate appear to be fully consistent, while those of the SN~Ia rate show a significant dispersion which increases with redshift, in particular in the range $0.5<z<0.7$ where the values of \\citet{Barris} and \\citet{Dahlen04} are 2-3 times higher that those of \\citet{Pain02} and \\citet{Neill07}. Our estimate of the SN~Ia rate  is consistent  with all other measurements in the redshift range we explored; our result does not help to discriminate between the steep trend suggested by the \\citet{Barris} and \\citet{Dahlen04} measurements and the slow evolution indicated by the \\citet{Neill07} measurement. The robust indication from the current data appears that the SN~Ia rate per unit volume at redshift 0.3 is a factor of $\\sim 2$ higher that in the local Universe, while in the same redshift range the CC SN rate increases by a factor of $\\sim 5$.   \\subsection{Comparison with the predicted evolution of the CC SN rate}\\label{CompSFR}  The stellar evolution theory predicts that all stars more massive than 8-10~M$_\\odot$  complete the eso-energetic nuclear burnings, up to the development of an iron core that cannot be supported by any further nuclear fusion reactions or by electron degenerate pressure. The subsequent collapse of the iron core results into the formation of a compact object, a neutron star or a black hole, accompanied by the high-velocity ejection of a large fraction of the progenitor mass. TypeII SNe originate from the core collapse of stars that, at the time of explosion, still retain their H envelopes, whereas the progenitors of type Ib$/$Ic SNe are thought to be massive stars which have lost their H (and He) envelope \\citep{Heger}. Given the short lifetime of their progenitors ($<30$~Myr), there is a simple, direct relation between the CC SN and the current SF rate:  \\begin{equation} r^{\\rm CC}(z) = K^{\\rm CC} \\times \\psi(z) \\end{equation}  where $\\psi(z)$ is the SFR and $K^{CC}$ is the number of CC SN progenitors from a 1 $M_\\odot$ stellar population:  \\begin{equation} K^{\\rm CC} = \\frac{\\int_{m_l^{\\rm CC}}^{m^{\\rm CC}_u} \\phi(m) dm}{\\int_{m_L}^{m_U}m\\phi(m) dm} \\end{equation} where $\\phi(m)$ is the IMF, $m_L-m_U$ is the total stellar mass range, and $m^{\\rm CC}_l-m^{\\rm CC}_u$ is the mass range of CC SN progenitors.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{rate_cc.eps}} \\caption{Comparison between the SN~CC and SF rate evolution. Symbols are as in Fig.~\\ref{obsrate} with additional open symbols (measurements not corrected for extinction) and filled black symbols (estimates for the high extinction correction). Lines are selected SFR evolutions from the literature. All SFHs have been scaled to the SalpeterA IMF.} \\label{ccrate} \\end{figure}  In principle, if accurate measurements of the CC SN and SF rates are available,  it is possible to probe the possible evolution with redshift of either the CC SN progenitor scenarios or the IMF by determining the value of $K^{CC}$.  Here however we assume that $K^{CC}$ does not evolve significantly in the redshift range of interest and compare the observed evolution of the CC SN rate with that predicted by the SFH.  The estimate of the cosmic SFH is based on different SF indicators, depending on redshift range, and requires a suitable parameterization and accurate normalization. Since there is a large scatter between the measurements obtained by different SFR indicators (a factor 2-3  at $z=0.3-0.5$ and increasing with redshift) it is  hard to obtain a consistent picture of the SFH \\citep{Hopkins06}.  Indeed observations made at different wavelengths, from X-rays to radio, sample different facets of the SF activity and are sensitive to different time scales over which the SFR is averaged. Thus, different assumptions are required  to convert the observed luminosities at the various wavelengths to the SFR, and different systematic uncertainties affect the SFR estimates. In particular the significant difference between the SFR inferred from the UV and $H\\alpha$ luminosity and that inferred from the far-IR (FIR) luminosity may be related to the effect of dust extinction (expecially for the UV) and/or to the contribution to the light from old stars and AGNs (for the FIR).  To illustrate this point, we select three representative prescriptions for the SFH, namely:  \\begin{itemize} \\item the piecewise linear fit of selected SFR measurements in the range $0<z<6$ by \\citet{Hopkins06}, \\item the fit to the SFR measurements from the $H\\alpha$ emission line, by \\citet{Hippelein}, with an exponential increase from $z=0$ to $z=1.2$, \\item the prescriptions by \\citet{Hernquist}, i.e. a double exponential function that peaks at redshift $z\\sim 5.5$, obtained from an analytical model and hydro-dynamic simulations. \\end{itemize}  All these SFHs were converted to the same IMF, a modified Salpeter IMF (SalA) with a turnover below 0.5 $M_{\\odot}$ and defined in the mass range $m_L=0.1$$M_{\\odot}$ to $m_U=120$ $M_{\\odot}$   \\citep{Baldry}. We assumed a  mass range of $8-50 M_{\\odot}$ for CC SN progenitors, which gives a scale factor $K^{\\rm CC}=0.009$.  The measurements of the CC SN rate per unit volume and the predicted evolutionary behaviours are shown in Fig~\\ref{ccrate}.   The observations confirm the steep increase with redshift expected by the SFH from \\citet{Hopkins06} and \\citet{Hippelein}. For a look-back time of 3~Gyr ($z=0.25$) both the SFR and the CC SN rate increase by a factor of $\\sim3$ compared with the local values. A flat evolution, as that proposed by \\citet{Hernquist}, appears inconsistent with the observed CC SN rates in the overall range of redshift.  With the adopted $K^{\\rm CC}$, the level of the CC SN rate predicted by the SFH of \\citet{Hippelein} and, in general, by the $UV$ and $H\\alpha$ based SFHs fits well the data. Instead, for the SFH of \\citet{Hopkins06} and in general the SFHs inferred through FIR luminosity, the predicted CC SN rate is higher than observed over the entire redshift range \\citep[see also][]{Dahlen04,Hopkins06,Mannucci07}.  If we correct the observed CC SN rate according to the high extinction scenario, we obtain an acceptable agreement between the data and the predictions with \\citet{Hopkins06} SFH, as shown in Fig.~\\ref{ccrate}.  However, this correction would require an  extremely high dust content in galaxies which is not favored by present measurements. Indeed, \\citet{Mannucci07} derived an estimate of the fraction of SNe which are likely to be missed in optical SN searches because they occur in the nucleus of starburst galaxies or, in general, in regions of very high extinction and found that, at the average redshift of our search, the fraction of missing CC SNe is only $\\sim 10\\%$, far too small to fill the gap between observed and predicted rates.  Alternatively we may consider the possibility of a narrower range for the CC SN progenitor masses: in particular, a lower limit of $10-12$ $M_{\\odot}$ would bring the observed CC SN rates in agreement with those predicted from FIR based SFHs. In this respect we notice that, from a theoretical point of view, there is the possibility that a fraction of stars between 7-8 M$_\\odot$ and 10-12 M$_\\odot$ avoids the collapse of the core and ends up as ONeMg White Dwarfs \\citep{Ritossa,Poelarends}. On the other hand, estimates of the progenitor mass from the detection in pre-explosion (HST)  images has been possible for a few SNe IIP (e.g. SN 2003gd \\citep{Hendry}, SN 2005cs \\citep{Pastorello}): their absolute magnitudes and colours seem indicate a moderate mass ($8-12$ $M_{\\odot}$ )\\citep{VanDyk,Smartt,Maund,Li06}.  Given these controversies, we conclude that in order to constrain the mass range of CC SN progenitors it is necessary to reduce the uncertainties in the cosmic SFH. In addition it is important to apply a consistent dust extinction correction both to SFH and to CC SN rate.  \\subsection{Comparison with the predicted evolution of SN Ia rate}\\label{Iamod}    According to the standard scenario SNe~Ia originate from the thermonuclear explosion of a Carbon and Oxygen White Dwarf (C-O WD) in a binary system.  In the first phase of the evolution the primary component, a star less massive than $8 M_{\\odot}$, evolves into a C-O WD. When the secondary expands and fills its Roche Lobe, two different paths are possible, depending on whether a common envelope forms around the two stars (double degenerate scenario, DD) or not (single degenerate scenario, SD). In the SD scenario, the WD remains confined within its Roche Lobe, grows in mass until it reaches the Chandrasekhar limit and explodes, while in the DD scenario the binary system evolves into a close double WD system, that merges after orbital shrinking due to the emission of gravitational wave radiation.  In both scenarios two basic ingredients are required to model the evolution of the SN~Ia rate: the fraction of the binary systems that end up in a SN~Ia, and the distribution of the time elapsed from star formation to explosion (delay-time). In the SD scenario, the delay time is the evolutionary lifetime of the secondary; in the DD, the gravitational radiation timescale has to be added. In both cases, the distribution of the delay times depends on the distribution of the binary parameters \\citep{Greggio05}. In principle, the SD and DD scenarios correspond to different realization probabilities and different shape of the delay time distribution (DTD) functions, hence rather different evolutionary behavior of the SN~Ia rate. As mentioned previously, the observations indicate that the distribution of the delay times of SN~Ia progenitors is rather wide.  The cosmic SFH is the other critical ingredient that modulates the evolution of the SN~Ia rate. Following \\citet{Greggio05} the SN~Ia rate is given by:  \\begin{equation}\\label{sniaeq} r^{\\rm Ia}(t) = k_{\\alpha}A^{\\rm Ia}\\int_{\\tau_{i}}^{min(t,\\tau_{x})} f^{\\rm Ia}(\\tau) \\psi(t- \\tau)d{\\tau} \\end{equation}  where $k_{\\alpha}$ is the number of stars per unit mass of the stellar generation, $A^{\\rm Ia}$ is the realization probability of the SN~Ia scenario (the number fraction of stars from each stellar generation that end up as SN~Ia), $f^{\\rm Ia}(\\tau)$ is the distribution function of the delay times and $\\psi(t-\\tau)$ is the star formation rate at the epoch $t-\\tau$. The integration is extended over all values of the delay time $\\tau$ in the range $\\tau_{i}$ and $min(t,\\tau_{x})$, with $\\tau_{i}$ and $\\tau_{x}$ being the minimum and maximum possible delay times for a given progenitor scenario. Here we assumed that both $k_{\\alpha}$ and $A^{\\rm Ia}$ do not vary with cosmic time.  A detailed analysis of the predicted evolution of the SN~Ia rate for different SFHs and DTDs is presented elsewhere \\citep{Forster,Blanc07}. Here we consider only one SFH, the piecewise interpolation of \\citet{Hopkins06},  since this is conveniently defined also at high redshift, and limit our analysis to few DTDs representative of different approaches to model the SN~Ia rate evolution: three models from \\citet{Greggio05}, and two different parametrizations \\citep{Mannucci06,Strolger} designed to address some specific observational constraints, regardless the  correspondence to a specific progenitor scenario.  Specifically, among the \\citet{Greggio05} models we select one SD and two DD-models, one of the ``close'' and the other of the ``wide'' variety, the latter being an example of a relatively flat DTD. This choice is meant to represent the full range of plausible DTDs. The minimum delay time for the SD model is the nuclear lifetime of the most massive secondary stars in the SN Ia progenitor's system, i.e. $8 M_{\\odot}$. In principle, for the DD model the minimum delay time could be appreciably larger than this because of the additional gravitational waves radiation delay. In practice, also for the DDs the minimum delay is of a few $10^{7}$ yrs. The maximum delay time is quite sensitive to the model for the SN~Ia progenitors. The reader should refer to \\citet{Greggio05} for a more detailed description of these models.  The DTD parametrization proposed by \\citet{Mannucci06} is the sum of two distinct functions: a Gaussian centered at 5$\\times10^{7}$ yr, representative of a ``prompt'' progenitor population which traces the more recent SFR, and an exponentially declining function with characteristic time of 3 Gyr, a ``tardy'' progenitor population proportional to the total stellar mass. This DTD was introduced to explain, at the same time, the dependence of the SN~Ia rate per unit mass on the galaxy morphological type, the cosmic evolution of the SN~Ia rate and, in particular, the high SN~Ia rate observed in radio-loud galaxies.  Finally, \\citet{Strolger} showed that the best fit of the apparent decline of SN~Ia rate at $z> 1$ is achieved for a DTD with a Gaussian distribution centered at about 3~Gyr. Nevertheless this DTD fails to reproduce the dependence of the SN rate on galaxy colours which is observed in the local Universe \\citep{Mannucci06}.   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{delay.eps}} \\caption{Delay time distributions as derived by \\citet{Greggio05} for SD and DD models, \\citet{Mannucci06} and \\citet{Strolger}.} \\label{delay} \\end{figure}  The selected DTD functions are plotted in Fig.~\\ref{delay} while the predicted evolutionary behaviours of the SN~Ia rate are compared with all published measurements in Fig.~\\ref{rateia}.  In all cases, the value of $k_{\\alpha}A^{\\rm Ia}$ was fixed to match the value of the local rate; depending on the model it ranges between 3.4-7.6$\\times 10^{-4}$. This normalization implies that, for the adopted SalA IMF, and assuming a mass range for the progenitors of $3-8$ $M_\\odot$,  the probability that a star with suitable mass becomes a SN~Ia, is $\\sim$ 0.01-0.03. \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{rate_ia.eps}} \\caption{SN~Ia rate measurements fitted with different DTD functions and the SFH by \\citet{Hopkins06}. Symbols for measurements are as in Fig.~\\ref{obsrate}.} \\label{rateia} \\end{figure}   The models obtained with the different DTDs are all consistent with the observations with the exception of the \"wide\" DD model, whose redshift evolution is definitely too flat. On the other hand none of the DTD functions, with the adopted SFH, is able to reproduce at the same time the very rapid increase from redshift 0 to 0.5 suggested by some measurements \\citep{Barris,Dahlen04} and the decline at redshift $>1$ \\citep{Dahlen04}. We note that a new measurement of \\citet{Poznanski} suggests that the SN~Ia rate decline at high redshift may be not as steep as estimated by \\citet{Dahlen04}.   Given the current uncertainties of both SN~Ia rate and SFH it is difficult to discriminate between the different DTD functions and hence between the different SN~Ia progenitor models. To improve on this point, more measurements of SN~Ia rate at high redshifts are required  to better trace the rate evolution. At the same time measurements in star forming and in passive evolving galaxies in a wide redshift range can provide important evidence about the SN~Ia progenitor models. In addition, it is essential to estimate the cosmic SFH more accurately because the position of the peak of the SFH was found to be the crucial parameter for the recovered delay time \\citep{Forster}."
    },
    "0710/0710.1550_arXiv.txt": {
        "abstract": "We present VLT/FORS2 spectroscopic observations of globular clusters (GCs) in five low surface brightness (LSB) dwarf galaxies: KK211 and KK221, which are both dwarf spheroidal satellites (dSph) of NGC~5128, dSph KK84 located close to the isolated S0 galaxy NGC~3115, and two isolated dwarf irregular (dIrr) galaxies UGC~3755 and ESO~490-17.~Our sample is selected from the Sharina et al. (2005) database of Hubble Space Telescope WFPC2 photometry of GC candidates in dwarf galaxies. For objects with accurate radial velocity measurements we confirm 26 as genuine GCs out of the 27 selected candidates from our WFPC2 survey.~One candidate appears to be a distant galaxy.~Our measurements of the Lick absorption line indices in the spectra of confirmed GCs and the subsequent comparison with SSP model predictions show that all confirmed GCs in dSphs are old, except GC KK211-3-149 ($6 \\pm$2 Gyr), which we consider to be the nucleus of KK211. GCs in UGC~3755 and ESO~490-17 show a large spread in ages ranging from old objects ($t>10$ Gyr) to clusters with ages around 1 Gyr. Most of our sample GCs have low metallicities $\\zh \\le -1$. Two relatively metal-rich clusters with $\\zh \\approx -0.3$ are likely to be associated with NGC~3115. Our sample GCs show in general a complex distribution of $\\alpha$-element enhancement with a mean $\\langle$[$\\alpha$/Fe]$\\rangle=0.19\\pm0.04$ derived with the $\\chi^{2}$ minimization technique and $0.18\\pm0.12$ dex computed with the iterative approach. These values are slightly lower than the mean $\\langle$[$\\alpha$/Fe]$\\rangle=0.29\\pm0.01$ for typical Milky Way GCs. We compare other abundance ratios with those of Local Group GCs and find indications for systematic differences in N and Ca abundance. The specific frequencies, $S_N$, of our sample galaxies are in line with the predictions of a simple mass-loss model for dwarf galaxies and compare well with $S_N$ values found for dwarf galaxies in nearby galaxy clusters. ",
        "introduction": "\\label{intro} The hierarchical structure formation scenario predicts that dwarf galaxies are the first systems to form in the Universe \\citep{peebles68}, and that more massive galaxies form through dissipative processes from these smaller sub-units. The involved physical mechanisms of this sequence depend on the density and mass of the parent dark matter halo, in the sense that more massive halos initiate star formation at earlier epochs and form their stars at a faster rate \\citep[e.g.][]{peebles02, renzini06, ellis07}. Because of this environmental gradient, we expect that dwarf galaxies in the field formed the first stellar population relatively late and at a lower pace compared to their counterparts in dense galaxy clusters. In other words, the difference in age and chemical composition between the oldest stellar populations in cluster and field dwarf galaxies should reflect the delay in the onset of structure formation in these two environments.  The task of measuring the age and chemical composition of the oldest stellar populations in distant dwarf galaxies from their integrated light is very challenging. An alternative approach is to investigate the oldest globular clusters (GCs) that are found in dwarf galaxies. Several photometric surveys of extragalactic GCs in dwarf galaxies outside the Local Group have been performed in the past decade \\citep[see review by][]{miller06}. However, only a handful of those were followed up with 8--10m-class telescopes to derive spectroscopic ages and chemical composition. Observations of galaxies in groups and clusters provide more and more evidence that environment is a major factor influencing the process of GC formation \\citep[e.g.][]{west93, tully02, grebel03, miller98}.~Recent progress in modeling the assembly history of GC systems in massive elliptical galaxies suggests that a significant fraction of metal-poor GCs were accreted from dwarf satellites at later times compared to the number of GCs initially formed in the parent galaxy halo \\citep{ppm07}. These results underline the ideas put forward in the work of \\cite{forte82} and \\cite{muzzio87} as well as the models of \\cite{cote98, cote02, hilker99}, who suggested that the rich GC systems of massive galaxies may be the result of significant GC accretion through tidal stripping of less massive systems.  Spectroscopic studies of a few GCs in cluster and field dwarf galaxies showed that most of these systems host at least some old GCs with ages $t\\ga10$ Gyr \\citep{puzia00, sharina03, strader03a, strader03b, beasley06, conselice06}.~Although today's accuracy of relative spectroscopic age determinations ($\\Delta t/t \\approx 0.2-0.3$) is not sufficient to resolve the expected delay of $\\sim\\!0.5-4$ Gyr in the onset of star-formation between cluster and field environment \\citep[depending on cosmology, ionizing source population, ionization feedback efficiency, etc., see][]{kauffmann96, treu05, thomas05, delucia06, clemens06}, the old ages combined with information on abundance ratios can provide a powerful tool to decide whether stellar populations in field dwarf galaxies followed the same early enrichment history as their analogs in denser environments. Furthermore, any difference in GC chemical composition between dwarf and more massive galaxies opens an attractive way of chemically tagging accreted sub-populations in massive halos and, therefore, enables us to quantify the mass accretion history of galaxies, a task that for old galaxies is infeasible from studies of the diffuse galaxy light alone. Similar ideas have been formulated for stellar populations that make up the diffuse component of the most nearby galaxies, which are close enough for high-resolution spectroscopy of individual stars \\citep{fh02, geisler07}. The obvious advantage of GC systems is that their spectra can be observed out to about 10 times greater distances.  In this work we analyze spectroscopic observations of GCs in dwarf irregular (dIrr) and dwarf spheroidal galaxies (dSph) in the field/group environment. Our sample consists of systems that are representatives of the lowest-mass bin of the Local Volume (LV) galaxy population, limited to distances $D<10$ Mpc. We refer to \\cite{karachentseva85} and \\cite{grebel99} for a morphological type definition of our sample galaxies. The paper is organised as follows. In Section~\\ref{observations} we describe our observations and data reduction as well as the methods of measuring radial velocities. Section~\\ref{analysis} summarizes the measurement of Lick line indices, their calibration, and the determination of spectroscopic ages, metallicities, and abundance ratios. Section~\\ref{discussion} is devoted to the discussion of our results. Conclusions are presented in section~\\ref{conclusion}.  \\begin{deluxetable*}{lccccccclcr}[!ht] \\tabletypesize{\\scriptsize} \\tablecaption{Properties of sample dwarf galaxies \\label{dwgprop}} \\tablewidth{0pt} \\tablehead{ \\colhead{Galaxy} & \\colhead{RA (J2000)}& \\colhead{DEC (J2000)} & \\colhead{$ \\mu_0$} & \\colhead{$D_{\\rm MD}$} & \\colhead{$A_{B}$} & \\colhead{$B$} & \\colhead{$B-I$} & \\colhead{$M_V$} & \\colhead{$N_{\\rm GC}$} & \\colhead{$S_{N}$} } \\startdata KK211, AM1339-445   & 13 42 06 & $-$45 13 18& 27.77     & 0.25       & 0.477   & 16.3$\\pm$0.2 & 1.8$\\pm$0.2 & $-$12.58&  2  & 18.6\\\\ KK221               & 13 48 46 & $-$46 59 49& 28.00     & 0.50       & 0.596   & 17.3$\\pm$0.4 & 2.0$\\pm$0.4 & $-$11.96&  6  & 95.1\\\\ KK084, UA200, KDG65 & 10 05 34 & $-$07 44 57& 29.93     & 0.03       & 0.205   & 16.4$\\pm$0.2 & 1.7$\\pm$0.2 & $-$14.40&  7  & 10.4\\\\ UGC3755, PGC020445  & 07 13 52 & $+$10 31 19& 29.35     &$\\sim5$     & 0.384   & 14.1$\\pm$0.2 & 1.1$\\pm$0.2 & $-$16.12&  32 & 11.4\\\\\\smallskip E490-017, PGC019337 & 06 37 57 & $-$25 59 59& 28.13     & 4.50       & 0.377   & 14.1$\\pm$0.2 & 1.1$\\pm$0.2 & $-$14.90&  5  &  5.4 \\\\ \\enddata \\tablecomments{Columns contain the following data: (1) galaxy name, (2) equatorial coordinates, (3) and (4) are the distance modulus and distance from the nearest bright galaxy in Mpc from \\cite{kara04}, and from \\cite{tully05} for UGC3755, (5) reddening from \\cite{sch98} in the B-band, (6) is the integrated $B$ magnitude, and (7) the integrated $B-I$ color, derived from surface photometry on the VLT-FORS2 images in this work (see Appendix~\\ref{sbprof}), (8) absolute $V$ magnitude from SPM05, (9) number of GCs according to SPM05 and this paper, (10) specific frequency, $S_N=N_{\\rm GC}10^{0.4(M_V+15)}$ \\citep{harris81}.} \\end{deluxetable*} ",
        "conclusions": "\\label{conclusion}  Numerous photometric and spectroscopic studies of globular clusters in Virgo and Fornax cluster dwarf galaxies have been undertaken in the last years, which targeted bright dwarf galaxies down to $M_V\\approx-15$ mag \\citep[see][]{miller06}. Due to observational selection effects dwarf galaxies fainter than this are missed at distances of $D \\approx 17$ Mpc. Faint LSB dwarf galaxies down to $M_V \\approx -12$ mag have long been thought to be free of globular clusters, because they have insufficient mass. Our HST/WFPC2 survey of low-mass dwarf galaxies (SPM05), situated at distances $2-6$ Mpc in the Local Volume, revealed a rich population of globular cluster candidates (GCCs). In this work, we observed five of these galaxies with the VLT/FORS2 spectrograph in MXU mode and found that all targeted GCCs except one are genuine globular clusters. We could also confirm five additional globular clusters in our sample galaxies. Two clusters appear to be the nuclei of KK84 and KK211. The confirmed globular clusters are in general old and metal-poor, and show a range of \\afe\\ ratios. The mean $\\langle$[$\\alpha$/Fe]$\\rangle=0.19\\pm0.04$ that was determined with the $\\chi^{2}$ minimization technique and $0.18\\pm0.12$ dex which was computed using the iterative approach appears slightly lower than the mean $\\langle$[$\\alpha$/Fe]$\\rangle=0.29\\pm0.01$ for typical Milky Way clusters. Globular clusters in the two isolated, relatively bright dwarf galaxies UGC~3755 and ESO~490-17 show a wide range of ages from 1 to 9 Gyr, and imply extended star formation histories in these galaxies. This goes in hand with the measured low \\afe\\ ratios for the younger clusters and is consistent with low intensity star bursts. The oldest clusters with the highest \\afe\\ are found in KK84, a companion of NGC~3115. Other chemical abundances indicate potentially interesting differences between globular clusters in dwarf and more massive galaxies and, if confirmed, would facilitate the quantification of the accreted mass in rich GC systems of massive early-type galaxies."
    },
    "0710/0710.4143_arXiv.txt": {
        "abstract": "The  gravitational  magnification   and  demagnification  of  Type  Ia supernovae  (SNe)  modify  their  positions  on  the  Hubble  diagram, shifting    the    distance     estimates    from    the    underlying luminosity-distance  relation.    This  can  introduce   a  systematic uncertainty in the dark energy  equation of state (EOS) estimated from SNe,  although  this  systematic  is  expected  to  average  away  for sufficiently large data sets.   Using mock SN samples over the redshift range $0 < z \\leq 1.7$ we quantify the lensing bias.  We find that the bias on the dark  energy EOS is less than half a percent for  large  datasets  ($\\gtrsim$   2,000  SNe).   However,  if  highly magnified  events  (SNe  deviating   by  more  than  2.5$\\sigma$)  are systematically removed from the analysis, the bias increases to $\\sim$ 0.8\\%.  Given that the EOS parameters measured from such a sample have a  1$\\sigma$  uncertainty of  10\\%,  the  systematic  bias related  to lensing in  SN data out to $z \\sim 1.7$  can  be safely ignored  in future  cosmological measurements. ",
        "introduction": "\\label{sec:introduction}  Since   the   discovery  of   the   accelerating   expansion  of   the universe~\\citep{Riess:98,  Perl:99, Knop:03,  Riess:04}, the  quest to understand the physics responsible  for this acceleration has been one of  the  major  challenges  of  cosmology.  At  present  the  dominant explanation  entails  an additional  energy  density  to the  universe called dark energy. The physics  of dark energy is generally described in terms of its equation of  state (EOS), the ratio of its pressure to density. In some  models this quantity can vary  with redshift.  While there exist a variety of probes  to explore the nature of dark energy, one of  the most compelling entails  the use of type  Ia supernovae to map the  Hubble diagram, and thereby directly  determine the expansion history of  the universe.  With increasing sample  sizes, SN distances can potentially provide multiple independent estimates of the EOS when binned   in   redshift   \\citep{Huterer:05,Sullivan:07a,Sullivan:07b}. Several present and  future SN surveys, such as  SNLS \\citep{snls} and the Joint  Dark Energy Mission  (JDEM), are aimed at  constraining the value of the dark energy EOS to better than 10\\%.  Although SNe  have been shown  to be good standardizable  candles, the distance  estimate to  a given  SN  is degraded  due to  gravitational lensing   of  its   flux  \\citep{Frieman:97,  Wambsganss:97, HolzWald:98}.  The  lensing becomes more prominent as  we observe SNe out to higher  redshift, with the extra dispersion  induced by lensing becoming  comparable   to  the  intrinsic   dispersion  (of  $\\sim0.1$ magnitudes) at  $z \\gtrsim 1.2$ \\citep{Holz:05}.  In  addition to this dispersion, which  leads to an  increase in the error  associated with distance   estimate  to  each   individual  supernova,   lensing  also correlates  distance estimates  of nearby  SNe on  the sky,  since the lines-of-sight   pass   through   correlated  foreground   large-scale structure  \\citep{Cooray:05,Hui:06}.  Although this  correlation error cannot be statistically eliminated by  increasing the number of SNe in the  Hubble  diagram,  the  errors  can be  controlled  by  conducting sufficiently wide-area ($> 5\\mbox{  deg}^2$) searches for SNe (in lieu of small-area pencil-beam surveys).  In addition  to the statistical  covariance of SN  distance estimates, gravitational lensing also  introduces systematic uncertainties in the Hubble  diagram  by  introducing  a  non-Gaussian  dispersion  in  the observed  luminosities of  distant SNe.   Since lensing  conserves the total  number  of  photons,  this  systematic bias  averages  away  if sufficiently large  numbers of SNe  per redshift bin are  observed. In this case the  average flux of the many  magnified and demagnified SNe converges  on the unlensed  value \\citep{Holz:05}.   Nonetheless, even with thousands  of SNe  in the  total sample it  is possible  that the averaging remains  insufficient, given  that one may  need to  bin the Hubble diagram at very small redshift intervals to improve sensitivity to the EOS. Furthermore, SNe at higher redshifts are more likely to be significantly lensed.  If ``obvious''  outliers to the  Hubble diagram are  removed from  the sample,  this introduces  an important  bias in cosmological  parameter  determination,  and  can lead  to  systematic errors in the determination of the dark energy EOS.   In this paper we  quantify the bias introduced in the estimation of the  dark energy  EOS due  to weak lensing  of supernova  flux.  We consider the  effects due  to the non-Gaussian  nature of  the lensing magnification  distributions~\\citep{Wang:02},  performing  Monte-Carlo simulations by  creating mock  datasets for future  JDEM-like surveys. The paper is organized as  follows: In $\\S$2.1  we discuss our parameterization   of   the  dark   energy   EOS,  $\\S$2.2   discusses gravitational  lensing, and $\\S$3  is an  in-depth description  of our methodology. We present our results in $\\S$4. ",
        "conclusions": "\\label{sec:discussion}  We first present  our results for the magnitude  case, as described in \\S\\ref{sec:mag}.    The  left   panel   of  Figure~\\ref{fig:2}   shows histograms  of  the  best-fit  values  of $w_0$  from  the  likelihood analysis,  after   marginalizing  over  $w_a$  and   $h$.   The  empty histogram, which  peaks at -1.009  (marked with a  vertical dot-dashed line),  is  for  the  model  with  300  SNe.   The  shaded  histogram, representing  the 2,000  SN  case, peaks  at  -1.003 (vertical  dashed line), while the hatched histogram representing the 10,000 SN case has its peak  at -1.002 (vertical  solid line).  These  distributions have 1$\\sigma$  widths  of 0.016,  0.006,  and  0.003, respectively.   This scatter is primarily due  to the intrinsic uncertainty associated with absolute calibration,  and is not  dominated by lensing.   Without the inclusion of  lensing, however, the distributions peak  at exactly -1, and show no bias.  The shifted mode  gives us a rough idea of the bias to be expected, on average, due  to lensing.  We find that 68\\% of the time a random sample of 300 SNe will have an estimated value for $w_0$ within  3\\% of  its fiducial  value, and  this drops  to 0.5\\%  when a sample size of 10,000 SNe is considered.  The right panel of  Figure~\\ref{fig:2} shows the same distributions as the  left panel,  but  this time  using  the flux-averaging  technique instead of averaging over  magnitudes.  The empty, shaded, and hatched histograms peak  at -1.007, -1.003, and  -1.001, respectively, showing the mean  bias for  the 300,  2,000 and 10,000  SN cases  (marked with dot-dashed, dashed,  and solid vertical  lines).  With flux-averaging, we expect that 68\\% of the time  a random sample of 300 SNe will yield a value of $w_0$ within 2.5\\%  of the fiducial value, and within 0.5\\% for a sample of 10,000 SNe.  The  1$\\sigma$ parameter  uncertainty on  $w_0$ ranges  from  the 20\\% level (for  300 SNe) to less  than 5\\% (for 10,000  SNe), dwarfing the bias due  to lensing.   Thus, we need  not be concerned  about lensing degradation  of  dark  energy  parameter estimation  for  future  {\\it JDEM}-like surveys.  We note, however,  that our estimated bias on the EOS is larger than the lensing  bias of $w<0.001$ quoted in Table 7 of \\citet{Wood:07}.   This is  not  surprising, given  their  use of  the simple Gaussian  approximation to lensing  from \\citet{Holz:05}, which is less effective for low statistics. Nonetheless, we agree with their conclusion that lensing  is negligible.  A similar conclusion was also reached by \\citet{Martel:07} who used  a compilation of 230 Type Ia  SNe \\citep{Tonry:03} in the redshift range $0  < z < 1.8$ to show  that the lensing errors  are small compared  to the intrinsic SNe errors.  \\begin{figure}[!tb] \\begin{center} \\includegraphics[scale=0.5]{f3.eps} \\caption{Histograms  showing the distribution  of the  values obtained for $w_0$ after marginalizing over $w_a$ and $h$ for the 2,000 SN case (using  the flux-averaging technique).   The shaded  histogram assumes that   the  full   sample  of   2,000  SNe   is  used   for  parameter estimation. The  hatched histogram shows the shift  when outliers (SNe that are shifted  above or below the Hubble diagram  by more than 25\\% on  either  side)  are  removed  from  the sample.  The  bias  in  the distribution is due to  the removal of highly-magnified lensing events from the sample.} \\label{fig:3} \\end{center} \\end{figure}  We now discuss the bias which arises if anomalous SNe are removed from the  sample.   Gravitational lensing  causes  some  SNe  to be  highly magnified, and  it is conceivable that these  ``obvious'' outliers are subsequently removed from  the analysis. In this case  the mean of the sample will be  shifted away from the true  underlying Hubble diagram, and a bias will be  introduced in the best-fit parameters. To quantify this  effect, we  remove  SNe  which deviate  from  the expected  mean luminosity-distance relation  in the Hubble diagram by  more than 25\\% (corresponding roughly to a $2.5\\sigma$  outlier). The SN scatter is a result of the convolution of  the intrinsic error (Gaussian in flux of width 0.1)  and the  lensing PDF,  and the outlier  cutoff leads  to a removal of  $\\sim$ 50 SNe  out of the  2,000 SNe.  These  outliers are preferentially  magnified,  due to  the  strong  lensing  tail of  the magnification distributions.   The demagnification tail is  cut off by the empty-beam  lensing limit, and  therefore isn't as  prominent. The hatched  histogram in Figure~\\ref{fig:3}  shows the  distribution when events with convolved error greater than 2.5$\\sigma$ are removed.  The vertical dot-dashed line  at -1.0075 shows the average  value of $w_0$ obtained in  this case, representing a  bias in the  estimate of $w_0$ roughly three times  larger than when the full  2,000 SNe are analyzed (shown by shaded histogram).  This bias is a result of cutting off the high  magnification tail of  the distribution,  and thus  shifting the data towards a net dimming of observed SNe, leading to a more negative value of $w_0$.  We also  apply a cutoff at  $3\\sigma$, in addition  to the $2.5\\sigma$ discussed above. This results in a removal of $\\sim20$ SNe on average, for each  2,000 SN  sample, and  leads to a  bias of  $\\sim0.6$\\%. Any arbitrary cut  on the  (non-Gaussian) convolved (lensing  + intrinsic) sample leads  to a  net bias  in the distance  relation, and  even for large outliers  and large SN  samples, this can lead  to percent-level bias in the best-fit values for $w_0$.  To  summarize, we  have quantified  the effect  of  weak gravitational lensing on  the estimation of dark  energy EOS from  type Ia supernova observations.  With  generated mock  samples of 2,000  SNe distributed uniformly in redshift up to  z$\\sim$1.7 (as expected in future surveys like {\\it JDEM}), we have  shown that the bias in parameter estimation due to lensing is less than  1$\\%$ (which is well within the 1$\\sigma$ uncertainty expected  for these missions). Analyzing the  data in flux or magnitude  does not alter  this result.  If lensed  supernovae that are highly magnified (such that  the convolved error is more than 25\\% from the  underlying Hubble  diagram) are systematically  removed from the sample,  we find  that the  bias increases by  a factor  of almost three. Thus, so long as all  observed SNe are used  in the Hubble diagram, including  ones that  are highly magnified,  the bias  due to lensing in the  estimate of the dark energy  EOS will be significantly less  than the  1$\\sigma$  uncertainty.  Even  for  a post-{\\it  JDEM} program with 10,000 SNe, lensing bias can be safely ignored."
    },
    "0710/0710.1634_arXiv.txt": {
        "abstract": "We present 107 new epochs of optical monitoring data for the four brightest images of the gravitational lens SDSS J1004+4112 observed between October 2006 and June 2007.  Combining this data with the previously obtained light curves, we determine the time delays between images A, B and C. We confirm our previous measurement finding that A leads B by $\\Delta t_{BA}=40.6\\pm1.8$~days, and find that image C leads image A by $\\Delta\\tau_{CA}=821.6\\pm2.1$ days. The lower limit on the remaining delay is that image D lags image A by $\\Delta\\tau_{AD}>1250$ days.  Based on the microlensing of images A and B we estimate that the accretion disk size at a rest wavelength of 2300\\AA\\ is $10^{14.8\\pm0.3}$~cm for a disk inclination of $\\cos{i}=1/2$, which is consistent with the microlensing disk size-black hole mass correlation function given our estimate of the black hole mass from the MgII line width of $\\log M_{BH}/M_\\odot=8.44\\pm0.14$. The long delays allow us to fill in the seasonal gaps and assemble a continuous, densely sampled light curve spanning 5.7 years whose variability implies a structure function with a logarithmic slope of $\\gamma = 0.35\\pm0.02$. As C is the leading image, sharp features in the C light curve can be intensively studied 2.3 years later in the A/B pair, potentially allowing detailed reverberation mapping studies of a quasar at minimal cost. ",
        "introduction": "The quasar SDSS~J1004+4112 at $z_s=1.734$ is split into five images by an intervening galaxy cluster at $z_l=0.68$ \\cite{inada,inada2,oguri}. With a maximum image separation of $14\\farcs62$, it is a rare example of a quasar gravitationally lensed by a cluster \\cite{wambsganss,inada3}. One of the most interesting applications of this system is to use the time delays between the lensed images to study the structure of the cluster.  If we assume the Hubble constant is known, then the delays break the primary model degeneracy of lensing studies (the ``mass sheet degeneracy''), and the delay ratios constrain the structure even if the Hubble constant is unknown.  After its discovery, several groups modeled the expected time delays in SDSS J1004+4112 and their dependence on the mean mass profile of the cluster \\cite{kawano,oguri,williams}.   When we measured the shortest delay in the system, between images A and B, we found a longer delay than predicted by the models (Fohlmeister et al. 2007, hereafter Paper I) where the discrepancy probably arose because the models included the cD galaxy and the cluster halo but neglected the significant perturbations from the member galaxies.   As we measure the longer delays, where the cluster potential should be relatively more important than for the merging A/B image pair, we would not expect cluster substructures to play as important a role.  We also expect this lens to have a fairly short time scale for microlensing variability created by stars either in the intracluster medium or in galaxies near the images. The internal velocities of a cluster are much higher than in a galaxy (700~km/s versus 200~km/s), and  SDSS~J1004+4112's position on the sky is almost orthogonal to the CMB dipole (Kogut et al. 1993), giving the observer a projected motion on the lens plane of almost 300~km/s.  In Paper I, we detected microlensing of the continuum emission of the A/B images in Paper I and there is also evidence for microlensing of the CIV broad line \\cite{richards,lamer,gomez}.  Once we have measured the time delays we can remove the intrinsic quasar variability and use the microlensing variability to estimate the mean stellar mass and stellar surface density, the transverse velocities, and the structure of the quasar source \\cite{gilmerino,mortonson,poin,morgan}.  Finally, we note that SDSS~J1004+4112 could be an ideal laboratory for studying correlations in the intrinsic variability of quasars.  With, image C leading images A and B by 2.3 years, sharp variations in image C can be used to plan intensive monitoring of images A and B to measure the response times as a function of wavelength (e.g. Kaspi et al. 2007), with the additional advantage that the delay between A and B provides redundancies that protect against weather, the Moon and the Sun.  The long delays between the images also mean that seasonal gaps are completely filled, and we can examine the structure function of the variability with a densely-sampled, gap-free light curve (modulo corrections for microlensing). Such data generally do not exist, since most time variability data for quasars (other than nearby reverberation mapping targets, e.g. Peterson et al. 2004) have very sparse sampling (e.g. Hawkins 2007 on long time scales for a small number of objects or Vanden Berk et al. 2004 on shorter time scales for many objects).  In Paper I \\cite{fohli} we presented three years of optical monitoring data for the four brightest images of SDSS J1004+4112 spanning 1000 days from December 2003 to June 2006. The fifth quasar image, E, is too faint to be detected in our observations. We measured the time delay between the A and B image pair to be $\\Delta\\tau_{BA}=38.4\\pm2.0$ days. While larger separation lenses tend to have longer time delays, for these two images the propagation time difference is small, because they form a close image pair ($3\\farcs8$) from the source lying close to a fold caustic.  For the more widely separated C and D images we could only estimate lower limits on the delays of 560 and 800 days relative to image B and A.  In this paper we present the 107 new optical monitoring epochs for the 2006/2007 season in \\S2.  When combined with our previous data we have light curves spanning 1250 days that allow us to measure the AC delay in \\S3.  In \\S4 we use the microlensing variability of the A/B images to measure the size of the quasar accretion disk, and in \\S5 we measure the structure function of the intrinsic variability.  We discuss the future prospects for exploiting this system in \\S6.  \\begin{figure*}[t] \\centering \\includegraphics[bb= 30 100 520 700, width=10cm,angle=0,clip]{f1.ps} \\caption{ Light curves of the A, B, C and D images of the quasar SDSS J1004+4112 from December 2003 to June 2007. Images C and D have been offset by 0.3 and 1.0 mag, respectively, in order to avoid overlap. We present a running average of one point every 5 days averaged over $\\pm7$ days to emphasize trends and to avoid confusion by noise. \\label{lcurve}} \\end{figure*} ",
        "conclusions": "We present a fourth season of monitoring data for the four bright images of the five image gravitational lens system SDSS J1004+4112.  We confirm our previous estimate for the time delay between the merging A/B pair, finding that B leads A by $40.6\\pm1.8$ days.  We measure the delay for image C for the first time, finding that it leads image A by $821.6\\pm2.1$~days.  We note that this is nearly twice the longest previously measured delay (the 417 day delay in Q0957+561 \\cite{schild, kundic}).  We find a lower bound that D lags A by more than approximately 1250~days.  Our current mass model predicts that D lags A by approximately 2000 days, which is consistent with the present limit. The fractional uncertainties in the AB delay are still dominated by sampling and microlensing, while the fractional uncertainties in the AC delay are dominated by cosmic variance due to density fluctuations along the line of sight rather than our measurement uncertainties of 0.3\\% (e.g. Barkana 1996).  A detailed model of this system, including the constraints from the multiply imaged, higher redshift arcs (Sharon et al. 2005), the X-ray measurements \\cite{ota,lamer} and a detailed understanding of the uncertainties will be a challenge.   We lack a completely satisfactory model for the system at present, in the sense that the modeling process is extraordinarily slow due to the ability of the gravitational potentials associated with the cluster member galaxies to generate additional but undetected images of the quasar, making it impossible to carry out a reliable model survey.  The record of models for this system is discouraging.  As we noted in Paper I, all three model studies (Oguri et al. 2004; Williams \\& Saha 2004; Kawano \\& Oguri 2006) generically predicted shorter AB delays than the observed 40 days, and that this could be plausibly explained by the absence of substructure (i.e. galaxies) in the potential models.  The longer AB-C and AB-D delays should be less sensitive to substructure.  Oguri et al. (2004) do not include an estimate of the AB-C delays and have A-D delays consistent with our present limits.  The range of B-C delays in Williams \\& Saha (2004) is consistent with our measurement of 820 days, but they predict AD delays shorter than our current lower bound of 1250~days.  Kawano \\& Oguri (2006) predict a range for the longer delays over a broad range of mass distributions, none of which match our delays in detail.  However, models with sufficiently long C-B delays generally have C-D delays long enough to agree with our present limits.  Based on our present mass model we used the microlensing between the A and B images to make an estimate of the size of the quasar accretion disk at 2300\\AA\\ in the quasar rest frame.  If we convert this to the expected size at 2500\\AA\\ assuming the $R_\\lambda \\propto \\lambda^{4/3}$ scaling for a thin disk and assume the mean disk inclination $\\cos (i)=1/2$ the scale on which the disk temperature matches the photon energy is $R_{2500\\AA}=10^{15.0\\pm0.3}$~cm. Comparisons to other disk models should use the half-light radius which is $2.44$ times larger. Based on the quasar MgII emission line width we estimate that the black hole mass is $10^{8.4\\pm0.2} M_\\odot$. For this mass, the microlensing accretion disk size-black hole mass correlation found by Morgan et al. (2007) predicts that $R_{2500\\AA}=10^{15.3}$~cm, which is in broad agreement with the measurement.  Further observations, the inclusion of additional images, and monitoring in multiple bands should improve these measurements and potentially allow us to determine the mean surface density in stars near the images $\\kappa_*$ and their average mass $\\langle M\\rangle$.  Similarly, the ability to construct continuous light curves of the intrinsic variability and to use image C to provide early warning of sharp flux changes that can then be intensively monitored in images A and B may make this system a good candidate for applying reverberation mapping techniques to a massive, luminous quasar.  At present, we already see that the system has a structure function typical of quasars."
    },
    "0710/0710.3631_arXiv.txt": {
        "abstract": "We use high-resolution, three-dimensional hydrodynamic simulations to study the hydrodynamic and gravitational interaction between stellar companions embedded within a differentially rotating common envelope.  We evaluate the contributions of the nonaxisymmetric gravitational tides and ram pressure forces to the drag force and, hence, to the dissipation rate and the mass accumulated onto the stellar companion.  We find that the gravitational drag dominates the hydrodynamic drag during the inspiral phase, implying that a simple prescription based on a gravitational capture radius significantly underestimates the dissipation rate and overestimates the inspiral decay timescale.  Although the mass accretion rate fluctuates significantly, we observe a secular trend leading to an effective rate that is significantly less than the rate based on a gravitational capture radius. We discuss the implications of these results within the context of accretion by compact objects in the common-envelope phase. ",
        "introduction": "To understand the evolution of binary star systems, it is essential to analyze the interactions between their stellar components. Examples of such influences include the spin-orbit tidal interaction and mass transfer, as well as interactions that result in the loss of mass and angular momentum. Equally important are the interactions of stars orbiting about their common center of mass within a differentially rotating common envelope. It is generally accepted that such an evolutionary stage is essential for the formation of short-period binary systems containing compact objects (see, e.g., Iben \\& Livio 1993 and Taam \\& Sandquist 2000). In this case, the interaction determines the orbital evolution of the system and the conditions under which the common envelope is ejected, leading to the survival of a remnant binary system or to a merger that forms a rapidly rotating single star.  The amount of mass and angular momentum accreted by the inspiralling components during this phase also has direct implications for the properties of the compact object population in binary systems.  Lacking multidimensional hydrodynamical simulations of the common-envelope phase, the initial numerical and semi-analytical studies of the problem used simple prescriptions for the stellar interactions based on the pioneering work by Hoyle \\& Lyttleton (1939) and Bondi \\& Hoyle (1944), as generalized by Bondi (1952).  These seminal studies focused on the idealized problem of the capture of matter by a gravitating point object moving supersonically with respect to a uniform medium.  In this framework, a gravitational capture radius, $\\rcap$, plays an important role in determining the rates of mass accretion and energy dissipation.  $\\rcap$ is given by \\begin{equation} \\label{Eqn:rcap} \\rcap = {2 G M \\over \\vrel^2 + \\cs^2}\\ , \\end{equation} where $M$ is the mass of the gravitating object, $\\vrel$ is the velocity of the object with respect to the medium, and $\\cs$ is the local speed of sound.  When a density gradient with scale height $H$ is present, the effective accretion radius $\\racc$ is (Dodd \\& McCrea 1952) \\begin{equation} \\label{Eqn:racc} \\racc = {\\rcap \\over 1 + (\\rcap/2H)^2}\\ . \\end{equation} The energy dissipation rate is then $L_{\\rm d} \\approx \\pi \\racc^2 \\rho \\vrel^3$, where $\\rho$ is the upstream density.  To improve on these estimates, hydrodynamic effects were approximated analytically by Ruderman \\& Spiegel (1971), Wolfson (1977), and Bisnovatyi-Kogan et al.\\ (1979) as well as numerically by Hunt (1971, 1979), Shara \\& Shaviv (1980), and Shima et al.\\ (1985).  These early multidimensional simulations considered axisymmetric flow, and their results have been used to calibrate the energy loss rate.  In particular, the drag coefficients obtained from such simulations (see, e.g., Shima et al.\\ 1985) have been used to estimate the rate of energy dissipation in the common envelope.  However, many of the simplifying assumptions underlying these studies are inadequate for direct application to common-envelope interactions. The flow is nonaxisymmetric and distinctly nonuniform, reflecting the existence of velocity or density gradients (the density may span several scale heights within $\\racc$). The effect of relaxing these assumptions has been studied in two dimensions by Fryxell \\& Taam (1988) and Taam \\& Fryxell (1989) and in three dimensions by Sawada et al.\\ (1989) and Ruffert (1999).  These studies could not encompass the full complexity of common-envelope interactions, since the envelope's self-gravity was ignored. Furthermore, because the companions move in elliptical orbits, their cores interact with matter that has already been affected in previous orbital phases. Thus, the state of the gas and its environment in these calculations must be regarded as highly idealized.  Within the past decade, three-dimensional numerical studies of the common-envelope phase that have relaxed earlier geometrical assumptions have been carried out by Sandquist et al.\\ (1998, 2000), DeMarco et al.\\ (2003a,b), and Taam \\& Ricker (2006).  Recently, we have carried out high-resolution adaptive mesh refinement (AMR) simulations of common-envelope evolution with effective resolutions of $2048^3$ (R. E. Taam \\& P. M. Ricker, in preparation), allowing the interaction of the stars within the common envelope to be examined and quantified.  In this Letter we report on some results of our numerical studies.  We focus on analyzing a single high-resolution simulation to determine the hydrodynamic and gravitational contributions to the drag forces affecting the orbital motion of the stellar components during the early inspiral phase.  We also analyze the accumulation of mass by the stellar components within the common envelope to compare their magnitudes to estimates based on an accretion radius formalism.  In \\S~2, we briefly describe the numerical method and our assumed model for the binary system.  Descriptions of the method of analysis and the numerical results are presented in \\S~3.  Finally, we summarize our results and comment on their possible implications for applications involving compact objects in short-period binary systems. ",
        "conclusions": "We have quantitatively described the interaction of stars within a common envelope on the basis of the analysis of the early inspiral stage of a $1.05 \\msun$ red giant and a $0.6 \\msun$ binary companion.  The orbital decay is dominated by the nonaxisymmetric gravitational drag associated with the self-gravitating matter in the common envelope. On the basis of high-resolution three-dimensional hydrodynamic simulations, this drag is 1-2 orders of magnitude greater than the hydrodynamic drag.  As a consequence, the orbital decay timescale is much shorter than that derived from analyses based on the Hoyle-Lyttleton-Bondi picture of accretion from a uniform medium by a gravitating point mass moving at supersonic speeds. In this latter description the gravitational drag associated with an accretion wake is not significantly larger than the hydrodynamic drag. The effect of long-range gravitational interactions is critical for reliable estimates of the orbital decay timescale and energy dissipation rate.  In this picture the drag and the mass accretion rate are not as directly related as they are in the Hoyle-Lyttleton-Bondi-type description, because the dominant drag term is gravitational, rather than hydrodynamic, in origin.  This decoupling is reflected in the much larger ratio $L_{\\rm d} / \\dot{M} \\sim 10^{15}-10^{16}$~cm$^2$~s$^{-2}$ measured in the simulation than expected from the gravitational capture radius formalism ($\\sim 10^{14}$~cm$^2$~s$^{-2}$).  Although the mass accretion rates estimated from our simulation should be regarded as only indicative due to the lack of a detailed inner boundary treatment for the companion, the structure of the flow (dominated as it is by tidal effects) strongly suggests that the true mass accretion rate should be much smaller than the rate expected in the gravitational capture radius formalism. The effective capture radius, on the basis of the observed ambient density and relative velocity, is almost an order of magnitude smaller than the expected value. In any case, the expected value would lead to an unrealistic level of accretion over the common-envelope period. Furthermore, the inspiral time suggested by the early evolution of our three-dimensional simulation is much shorter than that found in earlier one- and two-dimensional calculations that are based on the Bondi-Hoyle-Lyttleton picture. Assuming that the discrepancy in mass accretion rate continues into the deep inspiral phase, the total accumulation of matter onto the companion should be much smaller than previously expected.  These results would have little effect on the mass of an embedded main-sequence star because of its tendency to expand as a result of the high entropy within the common envelope (see Hjellming \\& Taam 1991); however, it can significantly affect the outcome for neutron stars within a common envelope.  In particular, the estimated accretion rates exceed $10^{-3} \\mpy$, for which steady state accretion flows with neutrino losses are possible (Chevalier 1989; Houck \\& Chevalier 1991). At these hypercritical mass accretion rates, photons are trapped in the flow and the Eddington limit is not applicable.  On the basis of this hypercritical accretion flow regime, Chevalier (1993) and Brown (1995) suggest that neutron stars embedded in the common envelope would accrete sufficient mass to form low-mass black holes (although see Chevalier 1996), and, hence, the formation of binary radio pulsars would require an evolutionary scenario involving progenitor stars of nearly equal mass (Brown 1995). However, the population synthesis of binary black holes and neutron stars by Belczynski et al.\\ (2002), including hypercritical accretion, resulted in an average accretion of $0.4 \\msun$.  Such a high rate of mass accretion is inconsistent with the observed masses of binary radio pulsars ($\\sim 1.35 \\msun$; Thorsett \\& Chakrabarty 1999) and indicates the need for a reduction in accreted matter during the common-envelope phase (see also Belczynski et al.\\ 2007). Our calculations show that the necessary reduction may arise naturally as a result of a more realistic treatment of the common-envelope phase. Consequently, this reduction could also lead to a reduction in the number of low-mass black holes, depending on the maximum mass of neutron stars, resulting from the accretion induced collapse of massive accreting neutron stars in the common-envelope phase.  Similarly, the mass accretion, which was found to be as large as several solar masses for black hole accretors, would also be reduced, thereby affecting the masses and spins of double black holes emerging from the common-envelope phase.  Further investigations are planned to examine the generality of these results regarding mass accretion and to quantify the importance of these processes for determining the properties (mass and spin) and ultimate fate of the compact components in short-period binary system populations. Such studies are not only important for determining the masses of binary neutron star and black hole systems resulting from the common-envelope phase (Belczynski et al.\\ 2007), but also their orbital periods, which directly influence the expected merger rates of such binary populations as sources for gravitational wave detection in the advanced LIGO experiment."
    },
    "0710/0710.3588_arXiv.txt": {
        "abstract": "We assess the relative merits of weak lensing surveys, using overlapping imaging data from the ground-based Subaru telescope and the Hubble Space Telescope (HST). Our tests complement similar studies undertaken with simulated data. From observations of 230,000 matched objects in the 2 square degree COSMOS field, we identify the limit at which faint galaxy shapes can be reliably measured from the ground. Our ground-based shear catalog achieves sub-percent calibration bias compared to high resolution space-based data, for galaxies brighter than $i^{\\prime}\\simeq$24.5 and with half-light radii larger than $1.8\\arcsec$. This selection corresponds to a surface density of 15 galaxies arcmin$^{-2}$ compared to $\\sim 71$ arcmin$^{-2}$ from space. On the other hand the survey speed of current ground-based facilities is much faster than that  of HST, although this gain is mitigated by the increased depth of space-based imaging desirable for tomographic (3D) analyses. As an independent experiment, we also reconstruct the projected mass distribution in the COSMOS field using both data sets, and compare the derived cluster catalogs with those from $X$-ray observations. The ground-based catalog achieves a reasonable degree of completeness, with minimal contamination and no detected bias, for massive clusters at redshifts $0.2<z<0.5$. The space-based data provide improved precision and a greater sensitivity to clusters of lower mass or at higher redshift. ",
        "introduction": "\\label{sec:Introduction}  Dark matter dominates the gravitational component of the cosmic energy density and thus provides the framework for structure formation in the Universe. However, by its nature, the distribution and cosmic growth are challenging to observe. The most promising probe is weak gravitational lensing: analysis of the distorted shapes of ordinary galaxies behind foreground mass concentrations. Several numerical techniques are now available to recover the projected mass distribution from these distortions, and tests on simulated datasets are underway to verify their precision \\citep{step1,step2}. There is great optimism in the weak lensing community that such methods will enable both the tomographic mapping of dark matter structures in time and space. This will also provide a robust statistical measure of the nature of dark energy over redshifts 0$<z<$1 \\citep{mellier99, refregier03}  Observational progress has been particularly dramatic. The first detections of statistical ``cosmic shear'' were only published in 2000 \\citep{bacon00, kaiser00, wittman00, vanwaerbeke00}. In the subsequent \\com{seven} years, weak lensing surveys have measured the dark matter power spectrum \\citep{brown03, heymans05, hoekstra06, sembolini06}, traced the evolution of structure \\citep{bacon05, kitching06, massey07a}, enabled the construction of lensing-selected cluster catalogs \\citep{miyazaki02a, wittman06, schirmer07, miyazaki07}, and non-parametrically reconstructed the total mass distribution both in clusters \\citep{kneib03, clowe06, jee07} and on larger scales \\citep{massey07b}. As a result, weak lensing has been identified as the most promising route to understanding the nature of dark energy by the ESA-ESO Working Group on Fundamental Cosmology\\footnote{\\tt http://www.stecf.org/coordination/esa\\_eso/cosmology.php}, joint NSF-NASA-DOE Astronomy and Astrophysics Advisory Committee\\footnote{\\tt http://www.nsf.gov/mps/ast/aaac.jsp}, and NSF-DOE High Energy Physics Advisory Panel Dark Energy Task Force\\footnote{\\tt http://www.nsf.gov/mps/ast/detf.jsp}.  The primary signal of any weak lensing analysis is the statistically coherent distortion of background galaxies along adjacent lines of sight. The main sources of statistical noise are the finite density of galaxies that can be sufficiently well-detected and resolved for accurate shape measurement, plus their intrinsic morphologies. The density of resolved galaxies also governs the angular resolution and fidelity of a reconstructed mass map which, in turn, determines the limiting halo mass that can be detected. On the other hand, statistical analyses of the dark matter power spectrum are less concerned with individual halos but require panoramic fields to counter the effects of cosmic (sample) variance. Minimizing statistical errors in such an analysis, within a finite survey lifetime, requires an optimal balance between area and depth.  A key debate in the development of future weak lensing experiments concerns the relative merits of ground- versus space-based platforms. Ambitious surveys now being planned with dedicated, ground-based facilities (eg VST-KIDS, DES, Pan-STARRS, LSST). These are driven by technological progress including panoramic cameras with small  optical distortions, highly sensitive imaging detectors, and (in the case of Pan-STARRS) on-chip active correction to reduce the width of the point spread function (PSF). Future surveys spanning significant fractions of the celestial sphere are envisaged, promising tight constraints on the  cosmological parameters.  However, measurements with current ground-based facilities are limited by the size and temporal variations of the PSF. There is concern in many quarters that wide-field facilities operating in space (e.g.\\ DUNE, SNAP, JDEM) will ultimately be required to achieve the precision required (particularly) to distinguish between various models of dark energy. Space-based facilities will be more costly but will likely offer increased depth, better photometric performance and a stable PSF. The key issue in gauging their merits is not statistical error, but the extent to which potential biases in ground-based data may act as a ``systematic floor'' to prevent complete exploitation.  Some valuable answers can be obtained by comparing simulated ground and space-based images, \\citep{wittman05,lampton06} and the Shear TEsting Programme \\citep[STEP:][]{step1,step2}. However, the input parameters used to generate the simulated data may not be realistic or address all the instrumental idiosyncrasies. Of particular concern are the stability and vagaries of the PSF. No simulations have  yet adequately addressed this point -- which may, ultimately, be the limiting problem for ground-based data. It is often argued that future facilities will be carefully designed to mitigate any limitations realized with current observational facilities. While progress can no doubt be expected, both on the ground and in space, we believe many lessons can be learned from extant data and hardware with proven engineering pedigree.  In this paper, we present the first direct comparison of weak lensing analysis  {\\it for the same sky field} using ground and space-based data. Deep, panoramic imaging has been obtained for the 1.64 deg$^2$ COSMOS field \\citep{scoville07a} by both the Advanced Camera for  Surveys ({\\it ACS}) on board the Hubble Space Telescope (HST) \\citep{scoville07b} and the {\\it Suprime-Cam} imager at the prime focus of the Subaru 8.2m telescope \\citep{taniguchi07}. In both cases, the entire field was covered by mosaicing many independent exposures.  The SuPrimeCam instrument was constructed with weak lensing analysis particularly in mind, and currently provides the best image performance available from any ground-based telescope, in terms of optical distortions over a large field. A comparison of these datasets should therefore provide a realistic and valuable assessment of the relative performance of state-of-the-art imagers on the ground and in space.  The paper is organized as follows. In \\S\\ref{sec:theory}, we briefly review the relevant theory. In \\S\\ref{sec:data}, we describe the two data sets, data reduction pipelines and weak lensing analyses. We then present the results. In \\S\\ref{sec:analyses}, we compare shear measures on a galaxy-by-galaxy  basis to determine the optimum depth at which the ground-based data matches the performance of the (deeper) space  based data. This permits us to determine the relative survey speeds of Subaru and HST for high precision cosmic shear experiments. In \\S\\ref{sec:Maps}, we construct maps of the mass distribution, treating the Subaru and HST maps as independent probes of the same field, and contrast these against X-ray data. This permits us to evaluate the completeness and reliability of a lensing-selected halo catalog, and evaluate the precision of their inferred masses as a function of redshift. In \\S\\ref{sec:conc}, we summarize our results and discuss their wider implications for future missions. ",
        "conclusions": "\\label{sec:conc}  We have performed parallel weak lensing analyses of Subaru and Hubble Space Telescope imaging in the COSMOS field. Our comparisons of the observed shear and convergence signals have revealed a number of issues, and suggest that such a study with real data usefully complements the independent approach based on blind analyses of simulated data \\citep{step1, step2}.  For statistical ``cosmic shear'' analyses, shear measurement with an existing ground based telescope, using existing measurement techniques, can be achieved with less than 1\\% bias relative to higher resolution space-based data, for a galaxy surface density of 15 arcmin$^{-2}$. One limitation of our approach is that we cannot check the performance of our space-based analysis on the additional, small galaxies. At first sight the factor of $\\sim$3 shortfall in surface density seems inconsequential given the lower cost and improved areal mapping speed of existing ground-based cameras such as SuPrime-Cam. However, accompanying the brighter Subaru limit is a reduction in survey depth and hence the redshift distribution of background sources. More distant sources contain a larger signal, and a narrower range of redshifts also hinders tomographical tests \\citep{bacon05,massey07b}, which tighten cosmological parameter constraints significantly.  A key issue is whether this limiting depth is a fundamental one for all future ground-based cameras. PanSTARRS, VST, and even LSST each have significantly smaller primary mirrors than Subaru, so achieving even the $S/N$ discussed here would require formidable exposure times. Most importantly, the deep infrared imaging that is required for photometric redshifts to enable tomographic analyses is likely to be difficult over large survey fields from the ground, because of increased sky background. Recent weak lensing analyses are limited at roughly the same level by uncertainty in galaxy shape measurement and photometric redshift estimation.  Statistical measurements from the ground are also hindered by variable atmospheric seeing. Past experience has taught the authors that data collected in seeing worse than $0.8\\arcsec$ is of little use for weak lensing analysis. The apparently rapid speed of data collection for our Subaru data belies the time spent waiting for better seeing, even with the excellent atmospheric conditions above Mauna Kea and the well-controlled dome seeing of Subaru. For this small-scale survey, we obtained exceptional quality imaging during a fortuitous observing window. The relevant quantity for larger-scale surveys in the future will be the time-averaged seeing quality, and the fraction of time spent with seeing better than $0.8\\arcsec$. This is particularly true for surveys like Pan-STARRS and LSST, that plan to adopt a strategy of co-adding many shorter exposures. Their advantage is that the stacked images will achieve a near-uniform image quality, by virtue of the independent PSFs in each short exposure. This can then be tuned to the required image quality by rejecting a certain fraction of exposures.  Variable seeing conditions is also of concern for the reconstruction of mass maps (c.f.\\ Green et al.\\ {\\it in prep.}). Difficulties in the analysis of one pointing in the SW corner of the Subaru map result in a patchy recovery of large-scale structure; with \\com{more noise and a lower range of probed redshift} in certain regions.   Most importantly, \\com{however, the four most massive clusters out of eight detected from space are also detected from the ground -- with one intriguing additional signal and two more confirmed clusters just ouside the field of view observed from space. Reassuringly, the three clusters conservatively deemed ``secure'' by the independent analysis of \\citep{miyazaki07} have now been confirmed via space-based weak lensing and $X$-ray observations. The physical properties of the four massive halos in common (A, B and C) are remarkably consistent whether derived from from ground- or space-based weak lensing. The measured masses and radial profiles of these clusters are consistent and, for the lower redshift clusters, the error bars are comparable.} A Dark Energy Task Force ``Stage 3'' survey from the ground appears eminently feasible.   A wide-field space-based platform would open up many new applications. Very important for statistical applications is the increased redshift range of resolved background galaxies. These not only contain a larger shear signal, but also more readily split into redshift bins for tomographic analysis. Three-dimensional analysis techniques will tighten constraints on cosmological parameters by factors of $3-5$, and directly measure quantities that depend upon the properties of dark energy, like the growth of structure over cosmic time and the redshift-distance relation. They will also eliminate sources of error due to the intrinsic correlations of galaxy shapes. Sufficiently good photometric redshifts require deep, wide-field near-IR imaging, and these are also realistically possible over large surveys only from space. Recent weak lensing analyses with relatively shallow near-IR coverage like \\citep{massey07a} are limited to roughly the same degree by uncertainty in galaxy shape measurement and photometric redshift estimation. Full implementations of cross-correlation cosmography will almost certainly require deep near-IR imaging from space. Such advanced techniques will become particularly important as ground-based surveys expand to encompass the entire observable sky.  The increased surface density of galaxies resolved from space also improves maps of the mass distribution. \\com{As shown in figures~\\ref{emode} and \\ref{bmode}, the statistical noise and systematic contamination in the $B$-mode are significantly reduced. Eight clusters are detected in the COSMOS $E$-mode signal without any contamination from spurious peaks. With the increased mass and spatial resolution of mass reconstructions from space, it becomes possible to detect halos the size of galaxy groups, as well as clusters over a wide range of redshifts -- thus tracing their formation,} which is governed by the properties of dark matter and the nature of gravity. Space-based data also crosses the threshold to mapping even filamentary large-scale structure in three dimensions \\citep{massey07b}. Obtaining the detailed, 3D distribution of mass will be particularly important near regions of interest like the Bullet cluster \\citep{clowe06}, where the small differences between the location of mass and baryons in a small patch of sky may yield the best possible information about the properties of dark matter. In this and other astrophysical phenomena, knowledge of the local mass environment and nearby large-scale structure is critical.  Overall, we conclude that ground-based weak lensing surveys can perform several tasks remarkably well, with sufficiently small amount of systematic bias to easily justify the next generation of dedicated  ground-based surveys. Two dimensional statistical analyses will be able to produce order-of-magnitude improvements in weak lensing constraints, using proven hardware technology and software pipelines. On the other hand, a wide-field space-based imager would provide important control over some systematic effects, and open up many new applications that are, at least currently, unachievable from the ground. For several of the most exciting techniques that will directly probe the nature of dark matter and dark energy, eventual space-based imaging is likely to be essential."
    },
    "0710/0710.3541_arXiv.txt": {
        "abstract": "The massive OB-type binary $\\sigma$~Ori~AB is in the centre of the very young $\\sigma$~Orionis cluster. I~have computed the most probable distances and masses of the binary for several ages using a dynamical parallax-like method. It incorporates the $BVRIH$-band apparent magnitudes of both components, precise orbital parameters, interstellar extinction and a widely used grid of stellar models from the literature, the Kepler's third law and a $\\chi^2$ minimisation. The derived distance is 334$^{+25}_{-22}$\\,pc for an age of 3$\\pm$2\\,Ma; larger ages and distances are unlikely. The masses of the primary and the secondary lie on the approximate intervals 16--20 and 10--12\\,$M_\\odot$, respectively. I~also discuss the possibility of $\\sigma$~Ori~AB being a triple system at $\\sim$385\\,pc. These results will help to constrain the properties of young stars and substellar objects in the $\\sigma$~Orionis~cluster. ",
        "introduction": "\\label{introduction}  The Trapezium-like system {$\\sigma$~Ori}, that illuminates the encolure of the {Horsehead Nebula}, is the fourth brightest star in the young Ori~OB~1~b association. The multiple system is composed of at least five early-type stars (Burnham 1892; Greenstein \\& Wallerstein 1958; van Loon \\& Oliveira 2003; Caballero 2007b). The two hottest components, $\\sigma$~Ori~A and B (O9.5V and B0.5V), are separated by only $\\sim$0.25\\,arcsec and were for a long time ``the most massive visual binary known'' (${\\mathcal M}_{\\rm A} + {\\mathcal M}_{\\rm B}\\sim 25+15~M_\\odot$; Heintz 1974). Although the binary has not yet completed a whole revolution, the orbital parameters are relatively well determined (Hartkopf, Mason \\& McCalister 1996; Heintz 1997; Horch et~al. 2002). It has been suggested that $\\sigma$~Ori~AB is a hyerarchical triple, being the primary a short-period, double-line spectroscopic binary (Frost \\& Adams 1904; Henroteau 1921; Miczaika 1950; Bolton 1974; Morrell \\& Levato 1991). However, a large amount of accurate, comprehensive spectroscopic investigations have failed to confirm this hypothesis (Heard 1949; Conti \\& Leep 1974; Humphreys 1978; Bohannan \\& Garmany 1978; Garmany, Conti \\& Massey 1980; Sim\\'on-D\\'{\\i}az \\& Lennon, priv.~comm.).  The $\\sigma$~Ori system is located in the centre of the well-known $\\sigma$~Orionis open cluster. The proper motions, radial velocities and spacial distribution of stars in this cluster strongly suggests a physical association between $\\sigma$~Ori itself and the young cluster (Zapatero Osorio et~al. 2002a; Caballero 2007a, 2007c -- see also Jeffries et al. [2006], who discovered a second older and kinematically and spacially distinct population). Because of its youth, comparative nearness and low extinction, the cluster has become the richest hunting ground for brown dwarfs and planetary-mass objects in the whole sky (B\\'ejar et~al. 1999; Zapatero Osorio et~al. 2000, 2002b, 2007). There is plentiful material in the literature about this cluster, covering topics like the initial mass function down to a few Jupiter masses (B\\'ejar et~al. 2001; Gonz\\'alez-Garc\\'{\\i}a et~al. 2006; Caballero et~al. 2007), jets and Herbig-Haro objects (Reipurth et~al. 1998; Andrews et~al. 2004), the frequency of accretors and discs (Zapatero Osorio et~al. 2002a; Oliveira, Jeffries \\& van Loon 2004;  Kenyon et~al. 2005; Oliveira et~al. 2006; Hern\\'andez et~al. 2007; Caballero et~al. 2007), the X-ray emission from young objects (Walter et~al. 1997; Sanz-Forcada et~al. 2004; Franciosini, Pallavicini \\& Sanz-Forcada 2006) or their photometric variability (Caballero et~al. 2004; Scholz \\& Eisl\\\"offel 2004). The most used values of heliocentric distance and age of the $\\sigma$~Orionis cluster are $d \\sim$ 360\\,pc and $\\sim$3\\,Ma (Brown, de Geus \\& de Zeeuw 1994; Perryman et~al. 1997; Zapatero Osorio et~al. 2002a; Oliveira et~al. 2002). There is a consensus in the literature that the cluster is younger than 8\\,Ma and older than 1\\,Ma. There is, however, a strong divergence of opinion on the heliocentric distance. Caballero (2007a) compiled determinations in the literature of the distance to the $\\sigma$~Orionis cluster from the 352$^{+166}_{-168}$\\,pc from {\\em Hipparcos} parallax to almost 500\\,pc from colour-magnitude diagrams.  Apart from the uncertainties of theoretical isochrones at very young ages, the derivation of the initial mass function of the cluster is strongly affected by the uncertainty in the actual age and heliocentric distance (Jeffries et~al. 2006; Caballero et~al. 2007). There are other investigations that require a precise age determination, such as the evolution of the angular momentum due to discs and stellar winds (Eisl\\\"offel \\& Scholz 2007), disc dissipation (Hern\\'andez et~al. 2007) and evolution of hot massive stars. % In this Letter, I~revisit a well known method for distance determination: the dynamical parallax (e.g. Russell 1928). I~apply it to the $\\sigma$~Ori~AB binary using state-of-the-art data and tools to determine its mass, age and heliocentric~distance. ",
        "conclusions": "The theoretical effective temperatures that correspond to the optimal fits lie on the intervals $T_{\\rm eff}$ = 30.4--34.6\\,kK for the primary and $T_{\\rm eff}$ = 25.2--27.5\\,kK for the secondary. The hottest temperatures are for the youngest ages. These values are consistent with the expected $T_{\\rm eff}$ for O9.5V and B0.5V stars, respectively (e.g. O9.5V: 30--35\\,kK -- Popper 1980; Gulati, Malagnini \\& Morossi 1989; Castelli 1991; Vacca, Garmany \\& Shull 1996; Martins, Schaerer \\& Hillier 2005), and with previous measurements of the $T_{\\rm eff}$ of $\\sigma$~Ori~A (30.0--33.0\\,kK -- Morrison 1975; Underhill et~al. 1979; Morossi \\& Crivellari 1980; Repolust et~al. 2005). The corresponding theoretical gravities ($\\log{g} \\sim$ 4.00--4.22) are also normal for class V at such temperatures.  The minima of $\\chi^2$ in Table~\\ref{bestfits} are very sensitive to the variations of heliocentric distance and of mass: on the one hand, at fixed mass and age, a fluctuation of $d$ of barely 30\\,pc results in a change of three orders of magnitude in $\\chi^2$; on the other hand, at fixed distance and age, a fluctuation of ${\\mathcal M}_{\\rm A}$ of barely 5\\,$M_\\odot$ results in a change of almost two orders of magnitude in $\\chi^2$. The minima of $\\chi^2$ are, however, quite unresponsive to the variations of the age between 1.0 and 4.9\\,Ma (see last column in Table~\\ref{bestfits}). A younger age gives a slightly better fit results and that favours a slightly larger distance ($d \\sim$ 350\\,pc). The results do not strongly suggest a younger age of 1\\,Ma, but they are useful in excluding an older age. The absence of a solution at 10\\,Ma agrees with previous upper limits on the ages of the {Ori~OB~1~b} association from the presence of very early-type stars in the main sequence (Blaauw 1964) and of the $\\sigma$~Orionis cluster from spectral synthesis surrounding the Li~{\\sc i} $\\lambda$670.78\\,nm line (Zapatero Osorio et~al.~2002a).  The derived distance interval for 3$\\pm$2\\,Ma, $d$ = 334$^{+25}_{-22}$\\,pc, is consistent with the canonical distance to the $\\sigma$~Orionis cluster of $d \\sim$ 360\\,pc, but is difficult to conciliate with the distance of 440\\,pc for 2.5\\,Ma that Sherry, Walter \\& Wolk (2004) used. The derived distance interval also deviates from very recent determinations of the distance to some elements in the {Ori~OB~1} complex. In particular, Terrell, Munari \\& Siviero (2007), using the eclipsing spectroscopic binary {VV~Ori} close to {Alnilam} ($\\epsilon$~Ori) in Ori~OB~1~b, and Sandstrom et~al. (2007), employing the Very Large Baseline Array in the {Orion Nebula Cluster} in {Ori~OB~1~a}, have determined very accurate heliocentric distances of 388--389\\,pc (see also Menten et~al. 2007). These values are also lower than the classical distance to the {Ori~OB~1} complex of $\\sim$440\\,pc from average {\\em Hipparcos} parallax (Brown, Walter \\& Blaauw 1999; de Zeeuw et~al. 1999). Because of projection effects and the large physical size of Ori~OB~1 (Reynolds \\& Ogden 1979), $\\sigma$~Ori could be easily contained within the complex. Given their kinematic and spacial association, the $\\sigma$~Ori system and the young $\\sigma$~Orionis cluster are likely at the same heliocentric distance and also age, if one assumes that massive and low mass star formation in a cluster is coeval (Prosser et~al. 1994; Massey, Johnson \\& Degioia-Eastwood 1995; Stauffer et~al. 1997; see, however, Sacco et~al. 2007). Recently, it has been suggested that the $\\sigma$~Orionis cluster is actually kinematically distinct from the {Ori~OB~1~b} association (Jeffries et~al. 2006), just as {25~Ori} is distinct from Ori~OB~1~a (Brice\\~no et~al.~2007).  If $\\sigma$~Ori~AB were a hyerarchical triple, as described in Section~\\ref{introduction}, it would be located at a larger heliocentric distance. The hypothetical companion to $\\sigma$~Ori~A, to which I~tentatively call {$\\sigma$~Ori~F}, would be 0.5\\,mag fainter than the primary in the 370--493\\,nm interval according to Bolton (1974). This wavelength interval corresponds to the $U,~B$ Johnson bands. Taking into account $m_{\\rm A} = m_{\\rm A+F} + 2.5 \\log{\\left( 1 + 10^{-\\frac{m_{\\rm A}-m_{\\rm F}}{2.5}} \\right)}$ and the difference $B_{\\rm A+F}-B_{\\rm B}$ in Table~\\ref{photometry}, then the apparent magnitudes in the blue band of the three components would be related to the combined magnitude $B_{\\rm A+F}$ through: $B_{\\rm A} \\approx B_{\\rm A+F} + 0.53$\\,mag, $B_{\\rm B} \\approx B_{\\rm A+F} + 1.33$\\,mag and $B_{\\rm F} \\approx B_{\\rm A+F} + 1.03$\\,mag. I~have looked for the distances and theoretical masses whose corresponding apparent magnitudes match the $B_{\\rm [A,B,F]}$ relations and the Kepler's third law, $({\\mathcal M}_{\\rm A} + {\\mathcal M}_{\\rm F}) + {\\mathcal M}_{\\rm B} = 7.45~10^{-7} d^3$. There are only a few solutions that simultaneously verify $T_{\\rm eff,A} \\approx$ 30.0--33.0\\,kK and $T_{\\rm eff,B} \\approx$ 26.0--30.0\\,kK, as expected from the spectral types of $\\sigma$~Ori~A+F and $\\sigma$~Ori~B. In the triple scenario, the F component would have an intermediate temperature between the A and B stars (roughly B0.0V) and would orbit very close to $\\sigma$~Ori~A. The valid solutions are found for the narrow distance interval 370--400\\,pc and only for ages between 1.0 and 4.9\\,Ma. Although distances less than 290\\,pc and larger than 450\\,pc are ruled out again, the most probable distance to $\\sigma$~Ori under the triple hypothesis, $d \\sim$ 385\\,pc, agrees very well with those of VV~Ori and the Orion Nebula Cluster. Further high-resolution spectroscopic studies are needed to ascertain the existence and characteristics of $\\sigma$~Ori~F."
    },
    "0710/0710.1896_arXiv.txt": {
        "abstract": "We present an analysis of five X-ray Multi-Mirror Mission (\\xmm) observations of the anomalous X-ray pulsar (AXP) 1E 2259+586 taken in 2004 and 2005 during its relaxation following its 2002 outburst. We compare these data with those of five previous \\xmm observations taken in 2002 and 2003, and find the observed flux decay is well described by a power law of index $-0.69\\pm0.03$. As of mid-2005, the source may still have been brighter than preoutburst, and was certainly hotter. We find a strong correlation between hardness and flux, as seen in other AXPs. We discuss the implications of these results for the magnetar model. ",
        "introduction": "\\label{sec:intr} It is now commonly believed that soft gamma-ray repeaters (SGRs) and anomalous X-ray pulsars (AXPs) are neutron stars with ultra-strong magnetic fields, i.e. magnetars \\citep{dt92a}.  Their common nature was conclusively demonstrated when AXP \\tfe was observed to emit SGR-like bursts in 2001 \\citep{gkw02} and \\ttfn, in the supernova remnant (SNR) CTB 109, was seen to undergo a major SGR-like outburst in 2002 \\citep{kgw+03,wkt+04}. Subsequently, a variety of different types of activity in AXPs have been seen, including short- and long-term flux variations \\citep{gk04,dkg07} and slow and rapid pulse profile changes \\citep{ikh92,kgw+03,wkt+04,icd+07,dkg07,dkg07b}, in addition to bursts and outbursts (\\citealp{gkw06,wkg+05,dkg07}; see \\citealt{kas07} for a recent review).  During \\ttfn's 2002 outburst, the pulsed and persistent fluxes rose suddenly by a factor of $\\ge$20 and decayed on a timescale of months. Coincident with the X-ray brightening, the pulsar suffered a large glitch of fractional frequency change $4\\times10^{-6}$ \\citep{kgw+03,wkt+04}. In the first few hours of the outburst, the pulsar's pulse profile changed significantly, its pulsed fraction decreased, and its spectrum hardened dramatically. Over 80 short SGR-like bursts from the pulsar were observed at the same time \\citep{gkw04}. A near-infrared (K$_s$) enhancement was also observed during the epoch of the outburst \\citep{kgw+03}.  Combining {\\it Rossi X-Ray Timing Explorer } (\\xte) observations and \\xmm observations of \\ttfn\\, taken during and after the outburst, \\cite{wkt+04} found that the decay of \\ttfn's unabsorbed flux (mostly inferred from \\xte pulsed fluxes) after the outburst was well characterized by two power law components: a rapid steep decay visible only during the first several hours ($<1$ day) of the outburst, and a slower decay of index $-0.22$ for the next several months. \\cite{tkvd04} found that the near-infrared enhancement at late times decayed at the same rate as the slow X-ray decay, although there were no IR observations during the first few hours of the outburst.  Other AXPs have also exhibited transient behavior that could be explained by SGR-like outbursts. AXP \\etten is called transient because it was only discovered in 2003 when it suddenly became brighter by a factor of 100 \\citep{ims+04,ghbb04}. \\cite{gh07} found that the flux of \\etten after 2003 followed an exponential decay of timescale 233.5 days. Similarly, the AXP CXOU~J164710.2$-$455216 was found to have brightened by a factor of $\\sim$300 between two \\xmm observations taken 5 days apart in 2006 September \\citep{icd+07}. Candidate AXP \\etff, was discovered in an observation made in 1993 by {\\it ASCA} \\citep{gv98,tkk+98}. Follow-up observations in 1999 showed that the source's flux was smaller by a factor of $\\sim$10 \\citep{vgtg00}. \\cite{tkgg06} found that \\etff\\, remains undetected in \\chandra observations taken in 2003, with its flux $\\sim $260-430 times fainter than observed in 1993.  The transient AXP phenomena summarized above are qualitatively similar to the 1998 August 27 flare of \\ninet, in which the X-ray flux decayed with a power law of index $\\sim -0.9$ \\citep{fmw+03}, and the flux decay of SGR 1627$-$41 since 1998, which followed a power law of index $\\sim -0.47 $ and lasted for $\\sim$800 days \\citep{kew+03}. However, thus far,  the AXP outbursts have been much less energetic than most SGR outbursts. Also, most of the burst energy was released during the afterglows of the AXP outbursts, while for SGR outbursts, the X-ray afterglows have less integrated energy than the burst itself.  With now a handful of AXP and SGR outbursts and subsequent relaxations observed, we can begin to look for correlations between different outburst and relaxation properties in the hope of constraining magnetar physics. For example, SGR outburst recoveries have been modeled as crustal cooling following impulsive heat injection, and in principle can yield constraints on the nature of the crustal matter \\citep{let02}. Alternatively, the AXP events have been interpreted in terms of magnetospheric twisting \\citep{tlk02,bt07}, whose recovery depends on electrodynamics in the region of the magnetosphere immediately outside the stellar surface. On the other hand, \\citet{gogk07} suggest that AXP recoveries can be modeled with a stationary magnetosphere, with only the surface temperature changing.  They argue that their model, which includes the stellar atmosphere, can be used to quantitatively determine the source's magnetic field.  In this paper we present a spectral and pulsed flux analysis of 10 \\xmm\\, observations of AXP 1E~2259+586\\ taken between 2002 and 2005, as the source relaxed back toward quiescence following its 2002 outburst. We compare the X-ray flux and spectral evolution of \\ttfn\\, with those of other magnetars, and interpret these results in terms of the magnetar model. ",
        "conclusions": "\\label{sec:disc}  In this paper, we have presented a comprehensive study of the X-ray recovery of AXP 1E~2259+586 following its 2002 outburst.  Here we discuss the properties of this recovery, compare them with those of other magnetar outbursts, and consider how they constrain the magnetar model.  \\subsection{Return to ``Quiescence''}  In our 2004 and 2005 {\\it XMM} observations, the source's temperature and unabsorbed fluxes were still higher than the preoutburst value (Fig. \\ref{fig:all}).  This suggests that the source was not fully back to the preoutburst flux level.  Our power law fit to the flux decay shows that the after-outburst quiescent flux level is $(1.75 \\pm 0.02)\\times 10^{-11}\\, \\rm ergs,s^{-1} cm^{-2}$, which is significantly higher than the preoutburst value [$(1.59\\pm0.01) \\times 10^{-11}\\,\\rm ergs\\,s^{-1} cm^{-2} $; Table \\ref{tab:fit}].  Either the 2005 flux had still not returned to its quiescent level, or perhaps it had returned to quiescence but the flux just before the event was unusually low.  Also possible is that this (and other) AXPs do not have well-defined constant quiescent fluxes, but have long-term flux variations.  Indeed, there is evidence for some X-ray flux variability in \\ttfn\\, over the years since its discovery in 1981 \\citep{bs96}. Other AXPs also show variability on a variety of timescales (see \\citealp{kas07} for a review).  \\subsection{Comparison with other Magnetar Recoveries}  It is useful to compare the behavior observed from \\ttfn\\, with that of other magnetars.  \\ninet's flux was found to follow a power law of index $-0.713\\pm0.025 $ after its 1998 August 27 flare \\citep{wkg+01}.\\footnote{Later the afterglow of the \\ninet\\, August 27 flare was fitted with a power law plus constant model instead of the single power law model used by \\cite{wkg+01}, and a decay index of $\\sim0.9$ was obtained \\citep{fmw+03}.}  This has been interpreted as the cooling of the magnetar outer crust following a sudden release of magnetic energy \\citep{let02}.  This model predicts a power law decay of index $\\sim -2/3$.  The flux of SGR 1627$-$41 was found to decay following a power law of index $\\sim-0.47$ since its 1998 source activation.  Approximately 800 days after the source activation, SGR 1627$-$41's flux suddenly declined by a factor of 10.  This behavior is also well fitted by the crust cooling model, although with some fine tuning \\citep{kew+03}. We fit the \\xmm 2--10 keV unabsorbed fluxes of \\ttfn\\, with a power law plus constant model, and found the best-fit power law index to be $-0.69 \\pm 0.03 $, close to that of \\ninet, and that predicted by the model. This suggests that the \\ttfn\\, outburst afterglow may also be explained by the diffusion of heat in the outer crust.  The transient AXP \\etten exhibited an outburst in 2003. \\cite{ims+04} found that the afterglow of the \\etten\\, outburst as observed by \\xte could be described by a power law decay model ($F\\propto t^{-\\beta} $) with $\\beta=0.45-0.73 $.  This is similar to the behavior of \\ttfn\\, and other SGRs.  However, \\cite{gh07} found that, with more observations taken by \\chandra\\, from 2003 to 2006, the afterglow of the \\etten\\, outburst actually followed an exponential decay of timescale 233.5 days. As we have shown in this paper, the pulsed and unabsorbed X-ray flux decay of \\ttfn\\, favors the power law decay model over the exponential decay. Perhaps the physical processes involved in the 2003 outburst of \\etten were different from those in 2002 outburst of \\ttfn.  \\subsection{Twisted Magnetosphere Model}  \\cite{tlk02} reported that, if there exists a global twist of the magnetosphere, the decay timescale $\\tau$ of this twist would be \\begin{equation} \\label{eq:twist} \\tau = 40\\Delta \\phi^2 (\\frac{L_X}{10^{35}{\\rm ergs\\,s^{-1}}})^{-1}(\\frac{B_{\\rm pole}}{10^{14} \\rm G})^2(\\frac{R_{\\rm NS}}{10 \\rm km})^3 \\rm yr. \\end{equation} \\cite{wkt+04} argued that, for \\ttfn, the twist angle $\\Delta \\phi$ should be $\\sim 10^{-2}$ rad. Thus, the predicted twist relaxation timescale of \\ttfn\\, is several hours, which is coincidently the timescale of the steeper flux decay observed at the beginning of the afterglow.  However, \\citet{bt07} have shown more recently that this decay timescale is actually expected to be much larger than equation~(\\ref{eq:twist}) suggests.  This is because, in their model, the self-induction of the twisted portion of the magnetosphere accelerates particles from the stellar surface and initiates avalanches of pair creation which forms the corona.  This corona persists in dynamic equilibrium, maintaining the electric current, as long as dissipation permits.  The relevant timescale in this picture for the decay of a sudden twist is given by \\begin{equation} \\tau \\simeq 0.3 (\\frac{L_X}{10^{35}{\\rm ergs \\, s^{-1}}}) (\\frac{e\\Phi_e}{\\rm GeV})^{-2}(\\frac{R_{\\rm NS}}{10 \\rm km}) \\,\\, {\\rm yr}, \\end{equation} where $L_X$ is the peak X-ray luminosity and $e\\Phi_e$ is the voltage along the twisted magnetic field lines and should nearly universally be $\\sim$1~GeV \\citep[see][]{bt07}. For 1E~2259+586, we find $\\tau \\simeq 1.2$~yr, given the peak luminosity $L_X \\sim 4 \\times 10^{35}(d/3 \\; {\\rm kpc})$~ergs~s$^{-1}$.  Thus, the longer observed decay after the initial steep decline may indeed correspond to the untwisting of a coronal flux tube in the \\citet{bt07} picture, although the predicted timescale is somewhat smaller than the observed time to return to quiescence.  We note that the \\citet{bt07} model predicts a linear flux decline, in contrast to what we have observed for 1E~2259+586 and what has been observed for XTE~J1810$-$197 \\citep{gh07}.  Moreover, in the $\\sim$5 yr of {\\it RXTE} monitoring of 1E~2259+586 prior to its 2002 outburst \\citep{gk02}, its pulsed X-ray luminosity in the 2--10~keV band was roughly constant at $\\sim 2 \\times 10^{34}$~ergs~s$^{-1}$.  This also is puzzling given the \\citet{bt07} prediction that if the time between large-scale events is longer than the decay time from the previous event, the magnetar should enter a quiescent state in which the observed luminosity is dominated by the surface blackbody emission.  Why should the ``quiescent'' blackbody emission from 1E~2259+586 be a full order of magnitude larger than that from XTE~J1810$-$197, especially given the latter's much larger inferred magnetic field ($1.7 \\times 10^{14}$ versus $6 \\times 10^{13}$~G)? This disparity in ``quiescent,'' steady luminosities is even larger when considering AXP 1E~1841$-$045, which has an apparently steady 2--10~keV luminosity of $1.4 \\times 10^{35}\\rm ~ergs~s^{-1}$, and comparing with probable AXP AX~1845$-$0258, which has quiescent luminosity approximately 2 orders of magnitude smaller \\citep{tkgg06}.  Distance uncertainties may contribute but not on a scale that can significantly alleviate this problem. This remains an interesting puzzle in magnetar physics.  The twisted magnetosphere or flux tube models generically predict that the flux and spectral hardness of magnetars in outburst should be roughly correlated due to increased scattering optical depth when the twist is larger.  However, a similar prediction for a flux/hardness correlation was made by \\cite{og07} in their thermally emitting magnetar model, using a simple prescription for the magnetosphere and scattering geometry, with the latter stationary, i.e. invoking no variable magnetospheric twists. \\citet{gogk07} found that their model could reproduce the existing data for XTE~J1810$-$197.  We note that hardness-intensity correlations have now been observed for \\seven\\, \\citep{cri+07}, \\tfe\\, \\citep{tgd+07}, and as we report, in our \\ttfn\\, \\xmm observations.  It would be interesting to apply analysis of \\cite{og07} to these data, but it is outside the scope of this paper.  \\subsection{Other Observed Recovery Properties}  The fact that the rms and area pulsed fractions remained largely constant while the blackbody radius (in the blackbody plus power law model) changed by a factor of $\\sim$2 (Fig. \\ref{fig:all}) is worth considering, if the empirical blackbody plus power law spectrum model somehow resembles the real radiation mechanism.  % Pulsed fraction should generally decrease when the thermally radiating region on the star grows, provided that this region is not very small compared to the entire surface.  Any realistic spectral model which takes radiative transfer in the atmosphere and scattering through the magnetosphere into account should be able to reproduce the observation in this regard as well.  A clear anti-correlation between  \\tfe's pulsed fraction and unabsorbed flux has been observed \\citep{tmt+05,gkw06,tgd+07}. However, we found no such correlation in the 2--10 keV band for \\ttfn. On the contrary, its 0.1--2 keV area pulsed fractions seem to be correlated with both 0.1--2 and 2--10 keV unabsorbed fluxes (see Fig. \\ref{fig:flxpf}). \\cite{gh07} found that \\etten's pulsed fraction measured between 2003 and 2006 after its outburst decreased with the decay of its flux, i.e. \\etten's pulsed fraction is also correlated with flux.  Thus, the striking anti-correlation between pulsed fraction and flux observed from \\tfe\\, is clearly not universal.  Finally, we note that the near-infrared flux decay of \\ttfn\\, was found to follow a power law of index $-0.75^{+0.22}_{-0.33}$ when fitted to a power law plus constant model \\citep{tkvd04}.  This decay index is close to what we found for the X-ray flux decay, thus confirming the reported correlation between near-IR and X-ray fluxes postoutburst.\\footnote{The $-$0.22 X-ray decay index reported by \\cite{wkt+04} and the $-$0.21 near-infrared flux decay index reported by \\cite{tkvd04} were obtained from a simple power law fitting, i.e. with no quiescent level included in the fit.}  \\citet{tgd+08} and \\citet{wbk+08} showed that the near-IR flux of 1E~1048.1$-$5937 do show correlation with X-rays at times of outbursts.  However, \\citet{crp+07} show that the near-IR flux variation of XTE~J1810$-$197 is not simply correlated with X-ray flux nor even monotonic postoutburst.  Thus, the AXP picture with regard to near-IR variability is not yet fully clear."
    },
    "0710/0710.3893_arXiv.txt": {
        "abstract": "{Neutrinos heavier than $M_Z/2\\sim 45$ GeV are not excluded by particle physics data. Stable neutrinos heavier than this might contribute to the cosmic gamma ray background through annihilation in distant galaxies as well as to the dark matter content of the universe.} { We calculate the evolution of the heavy neutrino density in the universe as a function of its mass, $M_N$, and then the subsequent gamma ray spectrum from annihilation of distant $N\\bar{N}$ (from $0<z<5$). } { The evolution of the heavy neutrino density in the universe is calculated numerically. In order to obtain the enhancement due to structure formation in the universe, we approximate the distribution of $N$ to be proportional to that of dark matter in the GalICS model. The calculated gamma ray spectrum is compared to the measured EGRET data.} { A conservative exclusion region for the heavy neutrino mass is 100 to 200~GeV, both from EGRET data and our re-evalutation of the Kamiokande data. The heavy neutrino contribution to dark matter is found to be at most 15\\%.  } {} ",
        "introduction": "The motivation for a fourth generation neutrino comes from the standard model of particle physics. In fact, there is nothing in the standard model stating that there should be exactly three generations of leptons (or of quarks for that matter).   The present limits on the mass of a fourth generation of neutrinos are only conclusive for $M_N\\lesssim M_Z/2 \\approx 46$ GeV \\citep[p.~35]{2006JPhG...33....1Y}. This limit is obtained from the measurement of the invisible width of the $Z^0$-peak in LEP, which gives the number of light neutrino species, as $N_\\nu = 2.9841 \\pm 0.0083$ \\citep{2001hep.ex...12021T}.  In \\cite{2000PhLB..476..107M}, a fourth generation of fermions is found to be possible for $M_{N}\\sim 50$ GeV, while heavier fermions are shown to be unlikely. However, this constraint is only valid when there is a mixing between the generations \\citep{Novikov:2001md} and since this is not necessarily true, we will not take it for certain.  In the context of cosmology and astrophysics there are other contraints. Light neutrinos, with $M_N\\lesssim 1$ MeV, are relativistic when they decouple, whereas heavier neutrinos are not. The light neutrinos must have $\\sum m_\\nu\\lesssim 46$ eV in order for $\\Omega_\\nu h^2<1$ to be valid \\citep{2006PrPNP..57..309H}. For the dark matter (DM) content calculated by \\cite{2003ApJS..148..175S}, the bound is $\\sum m_\\nu\\lesssim 12$~eV. The number of light neutrino species are also constrained to $N_\\nu = 4.2^{+1.2}_{-1.7}$ by the cosmic microwave background (CMB), large scale structure (LSS), and type Ia supernova (SNI-a) observations at 95\\% confidence \\citep{2006JCAP...01..001H}.  Neutrinos heavier than about 1 MeV, however, leave thermal equilibirum before decoupling and therefore their number density drops dramatically, see for example \\cite{RevModPhys.53.1}. This will be discussed in more detail in Sect.~\\ref{se:Evolution}.  The most important astrophysical bound on heavy neutrinos comes from Kamiokande \\citep{1992PhLB..289..463M} and this will be considered separately in the end.  In \\cite{PhysRevD.52.1828}, it is found that the mass range $60\\lesssim M_N \\lesssim 115$ GeV is excluded by heavy neutrino annihilation in the galactic halo. However, according to \\citet[p.~57]{2002PhR...370..333D} this constraint is based on an exaggerated value of the density enhancement in our galaxy.  Other works constraining the heavy neutrino mass include \\cite{1998JETPL..68..685F,1999astro.ph..2327F} and \\cite{2004hep.ph...11093B}. There has also been a study of the gamma ray spectrum of dark matter (DM) in general \\citep{2007PhRvD..75f3519A}.  For an exhaustive review of modern neutrino cosmology, including current constraints on heavy neutrinos, see \\citet{2002PhR...370..333D}. It is concluded that there are no convincing limits on neutrinos in the mass range $50\\lesssim M_N \\lesssim 1000$ GeV. A review of some cosmological implications of neutrino masses and mixing angles can be found in \\cite{2003nema.conf...53K}.    In this paper we consider a stable fourth generation heavy neatrino with mass $M_N \\gtrsim 50$ GeV possessing the standard weak interaction. We assume that other particles of a fourth generation are heavier and thus do not influence the calculations.  We assume a $\\Lambda$CDM universe with $\\Omega_{tot} = \\Omega_m + \\Omega_\\Lambda = 1$, where $\\Omega_m = \\Omega_b + \\Omega_{DM} = 0.135/h^2$, $\\Omega_b = 0.0226/h^2$ and $h = 0.71$ \\citep{2003ApJS..148..175S}, using WMAP data in combination with other CMB datasets and large-scale structure observations % (2dFGRS + Lyman $\\alpha$). %  Throughout the article we use natural units, such that the speed of light, Planck's reduced constant and Boltzmann's constant equal unity, $c = \\hbar = k_B = 1$.  If heavy neutrinos ($M_N\\gtrsim 50$ GeV) exist, they were created in the early universe. They were in thermal equilibrium in the early stages of the hot big bang, but froze out relatively early. After freeze-out, the annihilation of $N\\bar N$ continued at an ever decreasing rate until today. Since those photons that were produced before the decoupling of photons are lost in the CMB, only the subsequent $N\\bar N$ annihilations contribute to the photon background as measured on earth.  The intensity of the photons from $N\\bar N$-annihilation is affected by the number density of heavy neutrinos, $n_N$, whose mean density decreases as $R^{-3}$, where $R$ is the expansion factor of the universe. However, in structures such as galaxies the mean density will not change dramatically, and since the number of such structures are growing with time, this will compensate for the lower mean density. Note that the photons are also redshifted with a factor $R$ due to their passage through space-time. This also means that the closer annihilations will give photons with higher energy than the farther ones. ",
        "conclusions": "\\label{sec:Discussion}  The numerical calculation of the evolution of the heavy neutrino number density indicates that in the mass region $100\\lesssim M_N \\lesssim 200$, the cosmological neutrinos would give a cosmic ray signal that exceeds the measurements by the EGRET telescope \\citep{1998ApJ...494..523S}. Note that the clumping factor for these limits is rather conservative. In \\citet{2002PhRvD..66l3502U}, this factor is much larger, which would also produce a stronger limit on the heavy neutrino mass.  We can also compare our neutrino density with the results from the Kamiokande collaboration \\citep{1992PhLB..289..463M}. We scale the neutrino signal in their Fig.~2 to $\\Omega_N/\\Omega_{DM}$, where we use $h_0=0.71$, $\\Omega_{m}=0.2678$ and $\\Omega_b=0.044$. This is shown in Fig.~\\ref{fig:compKamiokande}, where we compare our numerical results for the relic neutrino density to the observed muon flux in the Kamiokande detector. \\begin{figure}[here!] \\resizebox{\\hsize}{!}{\\includegraphics{compKamiokande.pdf}} \\caption{ Predicted signal from enhanced $N\\bar N$ annihilation in the earth and the sun compared to the measured signal in the Kamiokande. On the y-axis: the number of muons (per 100~m$^2$year) produced by muon neutrinos resulting from heavy neutrino collisions in the sun and the earth, as evaluated by \\cite{1992PhLB..289..463M}, but scaled to our $\\Omega_N(M_N)$. On the x-axis: the heavy neutrino mass in GeV. } \\label{fig:compKamiokande} \\end{figure} This gives an exclusion region of $80\\lesssim M_N \\lesssim 400$ GeV. Our analytical results, which are comparable to the traditional relic neutrino densities, is about a factor two lower, giving an exclusion region of $90\\lesssim M_N \\lesssim 300$~GeV. The model that gives these limits \\citep{1987ApJ...321..571G} is rather complicated and not verified experimentally, so these results cannot be taken strictly. Note also that in the three-year WMAP analysis \\citep{2007ApJS..170..377S}, the value of $\\Omega_{DM}$ depends on which other data the WMAP data are combined with. For WMAP+CFHTLS $\\Omega_{DM}$ can be as high as 0.279 and for WMAP+CBI+VSA it can be as low as 0.155. The higher of these possibilities would give an exclusion region of $85\\lesssim M_N \\lesssim 350$ GeV. The lower boundary value would give an exclusion region of $75\\lesssim M_N \\lesssim 500$ GeV. A conservative limit based on the Kamiokande data gives the exclusion region $100\\lesssim M_N \\lesssim 200$ GeV.   If a heavy neutrino exists with a mass $M_N\\sim100$ GeV or $M_N\\sim200$ GeV it would give a small bump in the data at $E_\\gamma\\sim1$ GeV. Currently the data points are too far apart and the error bars too large to neither exclude nor confirm the eventual existence of such a heavy neutrino. Most of this part of the gamma ray spectrum is usually attributed to blazars, which have the right spectral index, $\\sim 2$ \\citep{1997ApJ...490..116M}.    We note that there could be an enhancement in the signal due to the higher DM densities within galaxies compared to the mean density in the halos. On the other hand, from within galaxies there will also be an attenuation due to neutral hydrogen, thus reducing the enhancement. There will also be a certain degree of extinction of the signal due to neutral hydrogen along the line of sight, but even if we assume complete extinction above $z=4$ the resulting spectrum decreases with only about 20\\%.    We are also aware of the ongoing debate concerning the antiprotons -- whether or not the DM interpretation of the EGRET gamma excess is compatible with antiproton measurements \\citep{2006JCAP...05..006B, 2006astro.ph.12462D}. We note the argument by de Boer that antiprotons are sensitive to electromagnetic fields, and hence their flux need not be directly related to that of the photons, even if they too were produced by $N\\bar N$ annihilation.  In the advent of the Large Hadron Collider, we also point out that there may be a possibility to detect the existence of a heavy neutrino indirectly through the invisible Higgs boson decay into heavy neutrinos \\citep{2003PhRvD..68e4027B}.   It will of course be interesting to see the results of the gamma ray large area space telescope (GLAST). It has a field of view about twice as wide (more than 2.5 steradians), and sensitivity about 50 times that of EGRET at 100~MeV and even more at higher energies. Its two-year limit for source detection in an all-sky survey is $1.6 \\times 10^{-9}$ photons cm$^{-2}$~s$^{-1}$ (at energies $>$ 100~MeV). It will be able to locate sources to positional accuracies of 30 arc seconds to 5 arc minutes. The precision of this instrument could well be enough to detect a heavy neutrino signal in the form of a small bump at $E\\sim 1$ GeV in the gamma spectrum, if a heavy neutrino with mass $\\sim$100 or 200~GeV would exist.  There are also some other possible consequences of heavy neutrinos that may be worth investigating. The DM simulations could be used to estimate the spatial correlations that the gamma rays would have and to calculate a power spectrum for the heavy neutrinos. This could be interesting at least for masses $M_N\\sim 100$~GeV and $M_N\\sim 200$~GeV. The annihilation of the heavy neutrinos could also help to explain the reionization of the universe. Another possible interesting application of heavy neutrinos would be the large look-back time they provide \\citep{2006PhLB..639...14S}, with a decoupling temperature of $\\gtrsim 10^{13}$ K \\citep{1989NuPhB.317..647E}."
    },
    "0710/0710.4966_arXiv.txt": {
        "abstract": "{ The EGRET excess of diffuse Galactic gamma rays  shows all the features expected from dark matter annihilation (DMA): a spectral shape given by the fragmentation of mono-energetic quarks, which is the same in all sky directions and an intensity distribution of the excess expected from a standard dark matter halo, predicted by the rotation curve. From the EGRET excess one can predict the flux of antiprotons from DMA. However, how many antiprotons arrive at the detector strongly depends on the pro\\-pagation model. The conventional isotropic propagation models trap the antiprotons in the Galaxy leading to a local antiproton flux far above the observed flux. According to Bergstr\\\"{o}m et. al. this excludes the DMA interpretation of the EGRET excess. Here it is shown that more realistic anisotropic propagation models, in which most antiprotons escape by fast transport in the z-direction, are consistent with the B/C ratio, the antiproton flux and the EGRET excess from DMA. } % ",
        "introduction": "\\label{intro} The interpretation of the observed EGRET excess of diffuse Galactic gamma rays as Dark Matter annihilation (DMA) (see \\cite{us} or contributions by W. de Boer, C. Sander and M. Weber, this volume) could be a first hint at the nature of dark matter. The excess was observed in all sky directions. From the spectral shape of the excess the WIMP mass was constrained to be between 50 and 100 GeV and from the distribution of the excess in the sky the Dark Matter (DM) halo profile was obtained. One of the most important criticisms of this analysis was a paper by Bergstr\\\"{o}m et. al. \\cite{bergstrom1} claiming that the antiproton flux from DMA would be an order of magnitude higher than the observed antiproton flux. They used a conventional propagation model assuming the propagation of charged particles to be the same in the halo and the disk. However, the propagation in the halo (perpendicular to the disk) can be much faster than the propagation in the disk \\cite{breit}.  In this paper we show that the local antiproton flux from DMA can be strongly reduced in an anisotropic propagation model and that the DMA interpretation of the EGRET excess can by no means be excluded by Galactic antiprotons. In section \\ref{CM} we discuss the pro\\-blems of the isotropic model for cosmic ray transport leading to the fact that our galaxy can work as a large storage box for antiprotons. An anisotropic pro\\-pagation model, which simultaneously describes the \\\\ EGRET excess and and the observed local fluxes of charged cosmic rays is introduced in section \\ref{CP}. Section \\ref{conclusion} summarizes the results. ",
        "conclusions": "Tracing of charged particles in realistic models of the regular Galactic magnetic fields with a turbulent (small-scale) component has shown that CRs remember the regular field lines, even if the irregular component is of the same order of magnitude as the regular, thus leading to enhanced diffusion in $\\phi$ and $z$ (see Fig. A1 in \\cite{blasi}).  With such an anisotropic  propagation model the amount of antiprotons expected from DMA can be reduced by one to two orders of magnitude. Therefore the claim by \\cite{bergstrom1} that the DMA interpretation of the EGRET excess of diffuse Galactic gamma rays is excluded is strongly pro\\-pagation model dependent. It only applies to a pro\\-pagation model with isotropic diffusion. An anisotropic pro\\-pagation model with different pro\\-pagation in the halo and the disk can reconcile the EGRET excess with the antiproton flux and the ratios of secondary/primary and unstable/stable nuclei. Clearly the DMA search for light DM particles is propagation model dependend.  Taking these uncertainties into account shows that DMA is a viable explanation of the EGRET excess of diffuse Galactic gamma rays, as shown in \\cite{us} and and can by no means be excluded by the antiproton flux predicted by a specific model.  \\vspace*{5mm}"
    },
    "0710/0710.0774_arXiv.txt": {
        "abstract": "In the third part of our photometric study of the star-forming region NGC~346/N~66 and its surrounding field in the Small Magellanic Cloud (SMC), we focus on the large number of low-mass pre-main sequence (PMS) stars revealed by the Hubble Space Telescope Observations with the Advanced Camera for Surveys. We investigate the origin of the observed broadening of the pre-main sequence population in the $V-I$, $V$ CMD. The most likely explanations are either the presence of differential reddening or an age spread among the young stars.  Assuming the latter, simulations indicate that we cannot exclude the possibility that stars in NGC 346 might have formed in two distinct events occurring about 10 and 5 Myr ago, respectively. We find that the PMS stars are not homogeneously distributed across NGC 346, but instead are grouped in at least five different clusters. On spatial scales from 0.8$''$ to 8$''$ (0.24 to 2.4\\,pc at the distance of the SMC) the clustering of the PMS stars as computed by a two-point angular correlation function is self-similar with a power law slope $\\gamma \\approx -0.3$. The clustering properties are quite similar to Milky Way star forming regions like Orion OB or $\\rho$\\,Oph.  Thus molecular cloud fragmentation in the SMC seems to proceed on the same spatial scales as in the Milky Way. This is remarkable given the differences in metallicity and hence dust content between SMC and Milky Way star forming regions. ",
        "introduction": "It is well known that Galactic OB associations also host large numbers of fainter, low-mass pre-main sequence (PMS) stars (e.g.\\ Preibisch \\& Zinnecker 1999; Sherry et al.\\ 2004; Brice\\~{n}o et al.\\ 2005). Low-mass PMS stars in stellar associations provide a longer-lasting record of the most recent star formation events than the short-lived high-mass stars.  Large-scale surveys have identified hundreds of PMS stars in nearby OB associations (Brice\\~{n}o et al. 2007). In order to understand the star formation triggering and history, or the Initial Mass Function (IMF), one has to study both high- and low-mass stars in star forming regions.  In many cases the low-mass populations of galactic OB associations cannot be easily distinguished from foreground main-sequence stars (e.g. Brice\\~{n}o et al. 2001).  Recently PMS stars have been discovered for the first time in an extragalactic stellar association, which suffers much less from contamination by the galactic disk.  {\\em Hubble Space Telescope} (HST) observations of the association LH\\,52 in the Large Magellanic Cloud (LMC) revealed $\\approx$500 PMS stars with masses down to 0.3\\,M$_\\odot$ (Gouliermis et al. 2006a). The investigation of individual extragalactic pre-main sequence stars opens a new field of study. As both high angular resolution and wide field-of-view are required, HST is the ideal observatory to carry out these studies. While HST observations of extragalactic associations cannot reach the detection limits achieved for local associations, they provide a unique opportunity for studying low-mass star formation in other galaxies.  The OB association NGC~346 in the Small Magellanic Cloud (SMC) is located in the central part of the brightest {\\sc H ii} region in the SMC, named LHA~115-N~66 or in short N~66 (Henize 1956). With 33 spectroscopically confirmed O and B stars, NGC~346 hosts the largest sample of young, massive stars in the SMC (Walborn 1978; Walborn \\& Blades 1986; Niemela et al. 1986; Massey et al. 1989; Walborn et al. 2000; Evans et al.  2006).  Photoionization models by Rela\\~{n}o et al. (2002) imply that these stars are a major source of ionizing flux for the surrounding diffuse interstellar medium (ISM). Recent imaging from the Wide-Field Channel (WFC) of the {\\em Advanced Camera for Surveys} (ACS) on-board HST (GO Program 10248; PI: A. Nota) revealed the PMS stellar content of the general region of NGC~346/N~66 down to the sub-solar mass regime (Nota et al.\\ 2006; Gouliermis et al.\\ 2006b -- hereafter Paper I -- ; Sabbi et al.\\ 2007).  Nota et al.\\ (2006) suggest that all PMS stars in the association are the product of a single star formation event, taking place 3 to 5 Myr ago.  The PMS distribution in the $V-I$, $V$ color-magnitude diagram (CMD), however, shows a prominent widening, which may be explained by an age spread of $\\approx$10 Myr. This raises the question {\\em if the PMS stars in NGC\\,346 are indeed the result of a single star formation event.}  Sabbi et al.\\ suggest that the PMS population is mainly concentrated in a number of subclusters (three of them at the central part, where the association is located), which formed at the same time from the turbulence-driven density variations, and not following a sequential process.  Star counts, however, reveal only a few compact PMS clusters with a significant number of stars in the area around the association. The main body of the association cannot be divided into separate subclusters (Paper I). {\\em Could the remote clusters be the product of a star formation event occurring before or after the event which triggered the formation of the association?}   In the first part of our study of this extraordinary star-forming region (Paper I), we compiled our ACS photometry of almost 100,000 stars, and presented preliminary results on stellar types and their distinct spatial distributions. We confirmed the co-existence of massive OB stars and low-mass PMS stars in NGC 346. In the present paper, we explore in greater detail the properties of the low-mass PMS stars with a focus on clustered star formation in NGC~346/N~66.  The paper is organized as follows: In \\S~2 we describe the datasets from HST/ACS, present the color-magnitude diagram, carry out a comparison of our photometry of the brightest stars with previous photometric studies, and discuss the reddening in the region. The broadening of the PMS distribution in the CMD, a thorough discussion of possible causes and explanations of this phenomenon as well as the spatial variations in the CMD are presented in \\S~3.  In \\S~4 we present the clustering properties of the PMS stars and discuss them in terms of hierarchical star formation. Our analysis of H\\alp\\ observations with HST/ACS as well as previous {\\em Spitzer} Space Telescope results on Young Stellar Objects (YSOs) are presented in \\S~5 and \\S~6, respectively. In \\S~7 we summarize the results. ",
        "conclusions": "In this paper we present the results from our photometric study on the recent star formation history of N~66 the brightest {\\sc H~ii} region in the SMC, related to the OB association NGC~346, as it is recorded in the observed PMS population of the region. Our deep photometry revealed an extremely large number of low-mass PMS stars in the association and the surrounding region, easily distinguishable in the $V-I$, $V$ CMD. We show that factors such as reddening, binarity and variability can cause a broadening of the positions of these stars in the CMD of the association, wide enough so that they can be misinterpreted as the result of multi-epoch star formation, and not as the product of a single star formation event. We found that a modest reddening like the one we found for the general area of NGC~346 ($E(B-V)\\simeq 0.08$ mag) can make the PMS stars appear younger than what they actually are, and therefore an age of 10 Myr better fits the observed sequence of PMS stars. However, our results do not exclude the possibility of multi-epoch star formation in the area of the association if the reddening is even lower, as suggested from the OB stars of this region ($E(B-V)\\simeq 0.04$ mag). In this case, two star formation events 10 and 5 Myr ago can explain the observed broadening of the PMS stars due to age spread and factors such as reddening and binarity. No specific dependence of the estimated ages of the PMS stars to their loci within the association, as a signature of sequential star formation, was found.  It has been previously suggested that three different generations of stars occurred through sequential star formation in the region of NGC~346/N~66 (Rubio et al. 2000) within the last 3 Myr (based on the age of the youngest OB stars) and that the PMS stars of NGC~346 represent a star formation event that took place 3 - 5 Myr ago (Nota et al. 2006). However, the study by Massey et al. (1989) based on the presence of evolved 15~M{\\solar} stars suggests that there might have been earlier star formation events in the region making NGC~346 as old as $\\sim$~12 Myr. In addition, there are recent indications that NGC~346 might host classical Be-type stars (Evans et al. 2006), and if so, the age of the association should be at least 10~Myr, a threshold given by Fabregat \\& Torrej{\\'o}n (2000) for classical Be-type stars to form. These results fit very well to the hypothesis that star formation in NGC~346 did already occur about 10~Myr ago, as our observations and simulations of the PMS stars suggest.  From star counts based on our ACS photometry we identify at least five PMS clusters across the region, covering a range of ages. On spatial scales from 0.8$''$ to 8$''$ (0.24 to 2.4\\,pc at the distance of the SMC) the clustering of the PMS stars as computed by a two-point angular correlation function is self-similar with a power law slope $\\gamma \\approx -0.3$. The clustering properties are quite similar to Milky Way star-forming regions like the Orion OB association or $\\rho$\\,Oph. Thus molecular cloud fragmentation in the SMC seems to proceed on the same spatial scales as in the Milky Way. This is remarkable given the differences in metallicity and hence dust content between SMC and Milky Way star-forming regions.  The youngest PMS stars are located mostly to the north of the bar of N~66, where three PMS clusters are identified. This area is also characterized by a high concentration of candidate YSOs (Simon et al. 2007), H\\alp-excess stars (found with our photometry), and IR-emission peaks (Rubio et al.  2000). This indicates that star formation probably still takes place in an arc-like feature, as it is outlined by the spatial distribution of these sources. In an accompanying letter, we combine these results with previous multi-wavelength studies of the region to investigate the star formation history, which helped to shape NGC~346/N~66 (Gouliermis et al. 2007b)."
    },
    "0710/0710.0890_arXiv.txt": {
        "abstract": "A next-generation lunar laser ranging apparatus using the 3.5~m telescope at the Apache Point Observatory in southern New Mexico has begun science operation. APOLLO (the Apache Point Observatory Lunar Laser-ranging Operation) has achieved \\emph{one-millimeter} range precision to the moon which should lead to approximately one-order-of-magnitude improvements in several tests of fundamental properties of gravity. We briefly motivate the scientific goals, and then give a detailed discussion of the APOLLO instrumentation. ",
        "introduction": "\\subsection{Scientific Motivation\\label{sub:Scientific-Motivation}}  A variety of observations and theoretical explorations---including the apparent acceleration of the expansion of the universe \\citep{sn1,sn2}, the possible existence of extra dimensions \\citep{extra_dim}, and attempts to reconcile quantum mechanics and gravity---provide motivation for improved tests of the fundamental aspects of gravity.  Lunar Laser Ranging (LLR) currently provides the best tests of a number of gravitational phenomena \\citep{jgw-96,jgw-latest} such as:  \\begin{itemize} \\item the strong equivalence principle (SEP): $\\eta\\approx5\\times10^{-4}$ sensitivity \\item time-rate-of-change of the gravitational constant: $\\dot{G}/G<10^{-12}$ yr$^{-1}$ \\item geodetic precession: 0.6\\% precision confirmation \\item deviations from the $1/r^{2}$ force law: $\\sim10^{-10}$ times the strength of gravity at $10^{8}$ meter scales \\end{itemize} LLR also tests other gravitational and mechanical phenomena, including for example gravitomagnetism \\citep{gravmag}, preferred frame effects \\citep{alpha-1,alpha-2}, and Newton's third law \\citep{newtonsthird}. LLR may also provide a window into the possible existence of extra-dimensions via cosmological dilution of gravity \\citep{lue,dvalimoon}. Besides the SEP, LLR tests the weak equivalence principle (WEP) at the level of $\\Delta a/a<1.3\\times10^{-13}$, but the LLR constraint is not competitive with laboratory tests. In addition, LLR is used to define coordinate systems, probe the lunar interior, and study geodynamics \\citep{dickey}.  These constraints on gravity are based on about 35 years of LLR data, although the precision is dominated by the last $\\sim$15 years of data at 1--3~cm precision.  APOLLO aims to improve tests of fundamental gravity by approximately an order-of-magnitude by producing range points accurate at the one-millimeter level.   \\subsection{A Brief History of LLR\\label{sub:A-Brief-History}}  The first accurate laser ranges to the moon followed the landing of the first retroreflector array on the Apollo 11 mission by less than two weeks (August 1, 1969). These were performed on the 3.0~meter telescope at the Lick Observatory. One month later, a second station using the 2.7~meter telescope at the McDonald Observatory began ranging to the moon \\citep{bender}. The operation at the Lick Observatory was designed for demonstration of initial acquisition, so that the scientifically relevant observations over the next decade came from the McDonald station, which used a ruby laser with 4~ns pulse width, firing at a repetition rate of about 0.3 Hz and $\\sim3$ J/pulse. This station routinely achieved 20~cm range precision, with a photon return rate as high as 0.2 photons per pulse, or 0.06 photons per second. A typical ``normal point''---a representative measurement for a run typically lasting tens of minutes---was constructed from approximately 20 photon returns.  In the mid 1980's, the McDonald operation was transferred to a dedicated 0.76~m telescope (also used for satellite laser ranging) with a 200~ps Nd:YAG laser operating at 10~Hz and 150 mJ/pulse. This station is referred to as the McDonald Laser Ranging System: MLRS \\citep{mlrs}. At about the same time, a new station began operating in France at the Observatoire de la C\\^ote d'Azur (OCA) \\citep{oca}. Using a 1.5~meter telescope, a 70~ps Nd:YAG laser firing at 10~Hz and 75 mJ/pulse, this became the premier lunar ranging station in the world. In recent years, the MLRS and OCA stations have been the only contributors to lunar range data with typical return rates of 0.002 and 0.01 photons per pulse, respectively. Typical normal points from the two stations consist of 15 and 40 photons, respectively.  Other efforts in LLR are described in \\citet{ep-llr}, and more detailed histories may be found in the preceding reference as well as in \\citet{bender,dickey}.   \\subsection{Millimeter Requirements\\label{sub:Millimeter-Requirements}}  The dominant source of random uncertainty in modern laser ranging systems has little to do with the system components, but rather comes from the varying orientation of the lunar retroreflector arrays. Although the arrays are nominally pointed within a degree of the mean earth position, variations in the lunar orientation---called libration---produce misalignments as large as 10 degrees, and typically around 7 degrees. This means the ranges between the earth and the individual array elements typically have a root-mean-square (RMS) spread of 15--36~mm, corresponding to about 100--240~ps of round-trip travel time. This dominates over uncertainties associated with the laser pulse width, and with jitter in the detector and timing electronics. A typical normal point containing 16 photons will therefore be limited to 4--9~mm range precision by the array orientation alone, though range residuals reported by analysis at the Jet Propulsion Laboratory tend to be larger than this.  Reaching the one-millimeter precision goal demands at a minimum the collection of enough photons to achieve the appropriate statistical reduction. Assuming an ability to identify the centroid of $N$ measurements---each with uncertainty $\\sigma$---to a level of $\\sigma_{\\mathrm{net}}=\\sigma/\\sqrt{N}$, the uncertainty stemming from the retroreflector array orientation typically demands 225--1300 photons in the normal point to reach the one millimeter mark. Worst-case orientations push the individual photon uncertainty to 50~mm, demanding 2500 photons. This is far outside of the capabilities of the aforementioned LLR stations. We point out that any constant range bias is accommodated in the analysis, so that only \\emph{variations} in the range are important to the experiment.  While adequate photon number is sufficient to reduce statistical uncertainty to the one-millimeter level, other sources of error could potentially limit the ultimate scientific capacity of LLR.  Most importantly, the gravitational physics is sensitive to the center-of-mass separations of Earth and Moon, while one measures the distance between a telescope and reflectors that are confined to the body surfaces.  The earth's surface in particular has a rich dynamic---experiencing diurnal solid-earth tides of 350~mm peak-to-peak amplitude, plus crustal loading from oceans, atmosphere, and ground water that can be several millimeters in amplitude.  Moreover, the earth atmosphere imposes a propagation delay on the laser pulse, amounting to $\\sim$1.5~m of zenith delay at high-altitude sites.  Satellite laser ranging, very long baseline interferometry, and other geodetic efforts must collectively contend with these same issues, for which accurate models have been produced.  A good summary of these models is published by the International Earth Rotation and Reference Systems Service \\citep[IERS:][]{iers}.  As an example of the state of these models, the long-standing atmospheric model by \\citet{marini-murray} has recently been replaced by a more accurate model \\citep{atmo1,atmo2}. The model differences for a high-altitude site are no more than 2~mm for sky elevation angles greater than 40 degrees---providing an indicative scale for the model accuracy.  The primary input for this model is the atmospheric pressure at the site, as this represents a vertical integration of atmospheric density, which in turn is proportional to the deviation of the refractive index, $n$, from unity.  Thus the zenith path delay, being an integration of $n-1$ along the path, is proportional to surface pressure under conditions of hydrostatic equilibrium.  A mapping function translates zenith delay to delays for other sky angles.  Measuring pressure to a part in 2000 (0.5~mbar) should therefore be sufficient to characterize the 1.5~m zenith delay at the one-millimeter level.  Our experiment records atmospheric pressure to an accuracy of 0.1~mbar.  The principal science signals from LLR appear at well-defined frequencies. For example, the equivalence principle signal is at the synodic period of 29.53 days, and even secular effects ($\\dot{G}$, precession) are seen via the comparative phases between periodic (monthly) components in the lunar orbit.  Because many of the effects discussed in the preceding paragraphs are aperiodic, they will not mimic new physics.  To the extent that these effects are not adequately modeled, they contribute either broadband noise or discrete ``signals\" at separable frequencies.  The science output from APOLLO may be initially limited by model deficiencies.  But APOLLO's  substantial improvement in LLR precision, together with a high data rate that facilitates deliberate tests of the models, is likely to expose the nature of these deficiencies and therefore propel model development---as has been historically true for the LLR enterprise.  Ultimately, we plan to supplement our LLR measurement with site displacement measurements from a superconducting gravimeter (not yet installed), in conjunction with a precision global positioning system installation as part of the EarthScope Plate Boundary Observatory (installed February 2007 as station P027).  \\subsection{The APOLLO Contribution}  APOLLO---operating at the Apache Point Observatory (APO)---provides a major improvement in lunar ranging capability. The combination of a 3.5~meter aperture and 1.1~arcsecond median image quality near zenith translates to a high photon return rate. Using a 90~ps FWHM (full-width at half-maximum) Nd:YAG laser operating at 20~Hz and 115~mJ/pulse, APOLLO obtains photon return rates approaching one photon per pulse, so that the requisite number of photons for one-millimeter normal points may be collected on few-minute timescales. To date, the best performance has been approximately 2500 return photons from the Apollo 15 array in a period of 8 minutes. The average photon return rate for this period is about 0.25 photons per shot, with peak rates of 0.6 photons per pulse. Approximately half of these photons arrived in multi-photon bundles, the largest containing eight photons. APOLLO brings LLR solidly into the multi-photon regime for the first time.  This paper describes the physical implementation of the APOLLO apparatus, including descriptions of the optical and mechanical design, the electronics implementation, and system-level design. For early reports on APOLLO, see \\citet{iwlr12,spacepart,iwlr13,iwlr14}. For an analysis of our expected photon return rate, see \\citet{weak_LLR}. A list of acronyms commonly-used in this paper appear in Appendix~\\ref{app:acronyms}. ",
        "conclusions": ""
    },
    "0710/0710.0632_arXiv.txt": {
        "abstract": "We present environmental dependence of the build-up of the colour-magnitude relation (CMR) at $z \\sim 0.8$. It is well established that massive early-type galaxies exhibit a tight CMR in clusters up to at least $z \\sim 1$. The faint end of the relation, however, has been much less explored especially at high redshifts primarily due to limited depths of the data. Some recent papers have reported a deficit of the faint red galaxies on the CMR at $0.8 \\lsim z \\lsim 1$, but this has not been well confirmed yet and is still controversial. Using a deep, multi-colour, panoramic imaging data set of the distant cluster RXJ1716.4+6708 at $z=0.81$, newly taken with the Prime Focus Camera (Suprime-Cam) on the Subaru Telescope, we carry out an analysis of faint red galaxies with a care for incompleteness. We find that there is a sharp decline in the number of red galaxies toward the faint end of the CMR below $M^*+2$. We compare our result with those for other clusters at $z \\sim 0.8$ taken from the literature, which show or do not show the deficit. We suggest that the \"deficit\" of faint red galaxies is dependent on the richness or mass of the clusters, in the sense that poorer systems show stronger deficits.  This indicates that the evolutionary stage of less massive galaxies depends critically on environment. ",
        "introduction": "\\label{sec:intro}  It is well known that red early-type galaxies exhibit a tight sequence on colour-magnitude diagrams, which is called the colour-magnitude relation (CMR) (e.g., Visvanathan \\& Sandage 1977; \\citealp{bow92}). In nearby clusters, the CMRs extend down to at least $5-6$ magnitude fainter than the brightest cluster galaxies (e.g., Terlevich, Caldwell \\& Bower 2001). The small colour scatter around the CMR is indicative of the homogeneity of early-type galaxies in clusters (e.g., Bower et al.\\ 1992, 1998). At high redshifts, the CMR has already been well established in clusters at least out to $z\\sim1$ as far as the bright end is concerned (e.g., \\citealt{ell97}; \\citealt{kod98}; \\citealt{sta98}; \\citealp{van98}; 2001; Blakeslee et al.\\ 2003; Stanford et al.\\ 2006; Mei et al.\\ 2006a,b). The faint end of the CMR, however, has been much less explored and still highly uncertain.  Some recent deep studies of distant galaxy clusters have shown a relatively small number of galaxies at the faint end of the CMR compared to local clusters. De Lucia et al.\\ (2004, 2007) showed such a deficit of faint red galaxies in $z=0.6-0.8$ clusters observed by the ESO Distant Cluster Survey (EDisCS; \\citealt{whi05}). A similar result was shown in Stott et al.\\ (2007). They compared the faint end of the luminosity functions of red galaxies in $z\\sim 0.5$ clusters from Massive Cluster Survey (MACS; Ebeling et al. 2001) with those of $z \\sim 0.1$ clusters from Las Campanas/AAT Rich Cluster Survey (LARCS; Pimbblet et al. 2006). \\cite{tan05} analysed the RXJ0152.7--1357 cluster (hereafter RXJ0152) at $z=0.83$ based on wide field data taken with the Subaru Prime Focus Camera on the Subaru Telescope (Suprime-Cam; Miyazaki et al.\\ 2002), and they also showed a deficit of faint red galaxies on the CMR. Based on these results, De Lucia et al. (2004, 2007) and \\cite{tan05} discussed that the faint end of the CMR well visible in the present-day universe was established at relatively later epochs as faint blue galaxies stopped their star formation after $z\\sim0.8$ in contrast to much earlier ($z\\gg1$) termination of star formation in massive galaxies. \\cite{tan05} classified galaxy environment into ``cluster'', ``group'' and ``field'', and examined the environmental dependence of the faint end of the CMR as well. They suggest that the build-up of the CMR depends also on environment in the sense that it is more delayed in lower-density environment.  The deficit of the faint end of the CMR is often discussed in the line of a currently favoured observational phenomenon called ``down-sizing''. This trend was first noted for field galaxies by Cowie et al.\\ (1996). They showed in their Hawaii Deep Field that most massive galaxies tend to show low star formation rates while less massive galaxies still show on-going star formation activity at $z\\lsim1$. Such a trend has been extended in both redshift space and magnitude range. Kauffmann et al.\\ (2003) showed in the local SDSS data that massive galaxies are dominated by red old galaxies. By contrast, less massive galaxies show bluer colours due to some on-going star formation, and galaxies below a few times $10^{10}$ \\msun {} in stellar mass are predominantly blue. A very similar trend was reported at $z\\sim1$ by Kodama et al.\\ (2004). They looked in the Subaru/XMM Deep Field and showed the distribution of galaxies at $z\\sim1$ on the colour-magnitude diagram. A clear bi-modality on the colour-magnitude diagram was observed again. Since then, a large number of papers have discussed this interesting issue. One of the most convincing cases is based on $\\sim$8,000 galaxies with spectroscopic redshifts within $0.7<z<1.4$ in the DEEP2 survey (Bundy et al.\\ 2006). They derived stellar mass functions of red and blue galaxies and showed that the mass where the dominant contribution is switched from red to blue galaxies shifts to smaller masses as time goes on. This down-sizing trend is found also in clusters as already mentioned above.  Recently, however, Andreon (2006) claimed that the faint end of the CMR is fully in place in the rich cluster MS1054--0321 (hereafter MS1054) at $z=0.83$ and questioned the universality of the deficiency of faint red galaxies at $z \\sim 0.8$. A critical problem is that the number of galaxy clusters having deep enough imaging data so that we can discuss the faint end of the CMR is still very limited at high redshifts. In fact, so far only a few clusters at $z \\sim 0.8$ (of which some are optically-selected clusters) have been studied in this respect (i.e., MS1054 by \\citealt{and06}, RXJ0152 by \\citealt{tan05}, and some optically-selected clusters by EDisCS in \\citealt{del04}, 2007). Therefore, it is crucial to increase the number of clusters and see if the deficit of faint red galaxies is universally observed or not and see what determines the degree of the deficit. In this paper, we examine another cluster at $z=0.81$, RXJ1716.4+6708 (hereafter RXJ1716), in order to obtain a more general picture of $z\\sim0.8$ clusters. We will also discuss a possible origin of the cluster-to-cluster variations.  The structure of this paper is the following. In Section 2, we introduce our PISCES project, and also summarize the properties of the RXJ1716 cluster shown by some previous works. We present a deficit of faint red galaxies in Section 3, and we discuss the environmental dependence of the nature of faint galaxies in Section 4. Finally, we summarize our results in Section 5. Throughout this paper we use $\\Omega_M =0.3$, $\\Omega_{\\Lambda} =0.7$, and $H_0 =70$ km s$^{-1}$Mpc$^{-1}$. Magnitudes are all given in the AB system, unless otherwise stated. ",
        "conclusions": "\\label{sec:summary} Using a deep, multi-colour, panoramic imaging data set of the distant cluster RXJ1716.4+6708 at $z=0.81$, newly taken with the Prime Focus Camera (Suprime-Cam) on the Subaru Telescope, we have carried out an analysis of red-sequence galaxies with a care for incompleteness. We have found that there is a sharp decline in the number of the red galaxies toward the faint end of the CMR below $M^*+2$. We compared our results with those for other clusters at $z \\sim 0.8$ taken from the literature, by calculating the luminous-to-faint ratio to quantify the degree of the ``deficit'' and by combining the information on richness of the individual clusters from X-ray properties.  We suggest that the deficit of faint red galaxies is dependent on the richness or mass of the clusters in the sense that poorer systems show stronger deficits.  This indicates that the evolutionary stage of less massive galaxies depends critically on environment. In order to confirm this interesting trend, we need a much larger sample of galaxy clusters over a wide range in richness, and not only at similar redshifts but also at other redshifts."
    },
    "0710/0710.4407_arXiv.txt": {
        "abstract": "Numerical simulations of planetesimal accretion in circumprimary and circumbinary orbits are described. The secular perturbations by the companion star and gas drag are included in our models. We derive limits on the parameters of the binary system for which accretion and then planetary formation are possible. In the circumbinary case we also outline the radial distance from the baricenter of the stars beyond which accumulation always occurs. Hydrodynamical simulations are also presented to validate our N--body approach based on the axisymmetric approximation for the gas of the disk. ",
        "introduction": "The formation of terrestrial planets and cores of giant planets within circumstellar disks involves the accumulation of a large number of planetesimals, solid bodies with initial sizes of roughly several kilometers (Lissauer 1993; Wetherill \\& Stewart 1993). The initial growth of the planetesimals can follow different paths depending on the their mutual velocities. If runaway growth occurs, a limited number of large planetary embryos form on a short timescale (about $10^{4} - 10^5$ years) followed by a period of violent mutual collisions until the planets reach their final mass. If the encounter velocities exceeds the planetesimal's escape velocities, the size distribution of the entire population exhibits an orderly growth until larger bodies are formed on a much longer timescale.  Most observed extrasolar planets are believed to have formed from planetesimals. The core--accretion model (Pollack et al., 1996) seems to explain a large fraction of the observed physical and dynamical properties of extrasolar gaseous giants in particular after the inclusion of migration by interaction with an evolving disk and gap formation (Alibert et al. (2005)). Neptune--size extrasolar planets possibly formed directly by planetesimal accumulation without reaching the critical mass to accrete a massive gaseous envelope. Around single stars the efficiency of planetesimal accumulation is very high, leading easily to planet formation. The influence of collective perturbations like stirring by mutual gravitational perturbations and damping by collisions and gas drag effects has been studied in detail in order to understand the conditions favoring runaway growth.  Radial velocity surveys have shown that exoplanets are found also in binary or higher multiplicity stellar systems (Raghavan et al. 2006, Desidera \\& Barbieri 2007). Planetesimal accumulation and then planet formation in binary (or multiple) stellar systems appears to be a more complex process than around single stars. The gravitational secular perturbations by the companion star may overcome the mutual planetesimal interactions and significantly affect the initial stage of accretion by exciting high eccentricities and then affecting significantly the relative velocity distribution. In this paper we explore the velocity evolution of planetesimals in S or C--type orbits under both, the perturbing effects of the companion star and gas drag. We recall that planetesimals revolving just about one star in a binary pair are on so-called \"S-type\" pr \"circumprimary\" orbits, whereas those that revolve about both stars have \"P-type\" or \"circumbinary\" orbits. ",
        "conclusions": "Numerical simulations of planetesimal evolution support the scenario in which planet formation may undergo even in binary star systems. Planetesimals in both S--type and P--type orbits keep their relative velocities low enough to allow accumulation rather than fragmentation for a wide range of binary orbital and physical parameters even if orderly growth or the so--called Type II runaway growth \\citep{kort01} are possibly more common than the conventional fast runaway growth presumed to occur around single stars. Both terrestrial planets and giant planets are supposed to form in binary systems unless extreme orbital conditions for the two stars are met like large binary eccentricity or very short separation (in the case of circumprimary disks). The potential lower rate of planet discovery around double stars may be ascribed to these cases rather than to a general effect related to the presence of a companion star."
    },
    "0710/0710.4893_arXiv.txt": {
        "abstract": "There is a large body of work that has used the excellent Chandra observations of nearby galaxies with neglible low mass X-ray binary (LMXB) populations. This has culminated in a ``Universal'' X-ray luminosity function (XLF) for high mass X-ray binaries  (HMXBs). However,  a number of methods have been used to convert from source intensities to luminosities when creating these XLFs. We have taken advantage of the XMM-Newton observations of the nearby starbursting spiral galaxy NGC 253 to test some of these methods. We find the luminosities derived from these various methods to vary by a factor of $\\sim$3. We also find the most influential factor in the conversion from intensity to luminosity to be the absorption. We therefore conclude that a more consistent approach is required for determining the true Universal XLF for HMXBs. Ideally, this would involve individual spectral fitting of each X-ray source. Certainly, the line-of-sight absorption should be determined from the observations rather than assuming Galactic absorption. We find the best approach for obtaining an XLF from low-count data to be the splitting of the X-ray sources into two or more intensity intervals, and obtaining a conversion from intensity to flux for each group from spectral modelling of the summed spectrum of that group. ",
        "introduction": "The X-ray populations of external galaxies have been well studied for the last $\\sim$20 years. Historically, studies of the individual sources have been severely limited by low count rates and signal to noise, and several methods have been used to derive the X-ray luminosity of a source from its intensity. Grimm et al. (2003, hereafter known as G03) used Chandra and ASCA surveys of nearby starburst galaxies, along with ASCA, MIR-KVANT/TTM and RXTE/ASM observations of HMXBs in our Galaxy and the Magellanic Clouds to obtain a correlation between the X-ray properties of HMXBs and the star formation rate (SFR) of their host galaxies. They chose their sample of galaxies to have sufficiently high SFR to total mass ratios so that their X-ray populations would be dominated by HMXBs, with negligible LMXB contributions. G03 used published Chandra X-ray luminosity functions (XLFs), scaled  assuming the Hubble constant to be 70 km s$^{-1}$ Mpc$^{-1}$. They found the XLFs of these galaxies to be strikingly similar, when normalised by the SFR of the galaxy, and proposed a universal HMXB XLF. Additionally,  they found that the number of sources with 2--10 keV luminosities $>$2$\\times$10$^{38}$ erg s$^{-1}$ to be proportional to SFR$^{1.06\\pm0.07}$. Furthermore they discovered a linear relation between the total HMXB X-ray flux of a galaxy and its SFR, for SFRs $\\ga$4 M$_{\\odot}$ yr$^{-1}$.   Several different methods were used to convert from X-ray intensity to luminosity when creating the XLFs for galaxies in the G03 sample. Some XLFs were derived  assuming  a standard X-ray binary emission model (a power law with spectral index, $\\Gamma$, $\\sim$1.7 or a 5 keV bremsstrahlung) with Galactic line-of sight absorption (e.g. Zezas et al., 2002; Soria \\& Kong, 2002). Others used the X-ray colours to estimate the emission spectrum, using Galactic absorption (e.g. Eracleous et al., 2002), or deriving the absorption from the colours also (e.g. Lira et al., 2002). Also, some XLFs were derived using best fit spectra to individual bright sources (Smith \\& Wilson, 2001), or to the stacked X-ray population (Roberts et al., 2002).   XMM-Newton (XMM) is the most sensitive X-ray imaging telescope in the 0.3--10 keV band. We can therefore use deep XMM observations of nearby galaxies to glean the accuracy of these methods by comparing their resulting XLFs with the XLF derived from freely fitting each X-ray source. To do this, we chose XMM-Newton observations of NGC 253. NGC 253 is a star-bursting spiral  galaxy in the Sculptor  Group at a distance of $\\sim$4 Mpc; it is almost edge on, with a D$_{25}$ region of $\\sim$25$'$$\\times$7$'$. We fully describe our analysis of the observations in Barnard et al. (2007, MNRAS submitted,  hereafter known as B07), and concentrate here on the 2003, 110 ks XMM observation of NGC253. ",
        "conclusions": "Grimm et al. (2003) report a Universal HMXB XLF derived from published XLFs of several nearby galaxies. They also derive relations between the star formation rate and (i) the total luminosity of the point X-ray sources in the galaxies and (ii) the number of X-ray sources in a galaxy with 2--10 keV luminosity $>$2$\\times$10$^{38}$ erg s$^{-1}$.  However, the published XLFs were produced using a number of methods, in most cases assuming Galactic line-of-sight absorption.  We have tested several of these models using a deep XMM-Newton observation of the nearby galaxy NGC253, included in a secondary sample of Grimm et al. (2003). We obtained freely modelled luminosities for the 140 brightest sources in the field and also obtained the conversion factors from intensity to flux for some of these different models. We  found them to vary by a factor of $\\sim$3. We found the biggest influence on the conversion factor to be the absorption; Model I assumed Galactic line-of-sight absorption, while absorptions 20--50 times higher than this were obtained for the other methods. Since the universal XLF and relations between SFR and X-ray properties were obtained using a mixture of methods, we find them to be inconclusive.  We have also found from freely fitting the spectra of 140 bright sources that the corresponence between count-rate and flux is non-linear; this is largely due to the systematic softening of the spectra of more luminous sources. Hence it is unwise to employ a single emission model when describing the X-ray populations of nearby galaxies. Ideally, one would construct XLFs only from luminosities derived from free spectral modelling. To get the most out of the low-photon data, we recommend the stacking method of e.g. Roberts et al. (2002)."
    },
    "0710/0710.4588_arXiv.txt": {
        "abstract": "{Stability properties of  magnetic-field configurations containing the toroidal and axial field are considered. The stability is treated by making use of  linear analysis. It is shown that the conditions required for the onset of instability are essentially different from those given by the necessary condition $d (s B_{\\varphi})/ds > 0$, where $s$ is the cylindrical radius. The growth rate of instability is calculated for a wide range of the parameters. We argue that the instability can operate in two different regimes depending on the strength of the axial field and the profile of the toroidal field. ",
        "introduction": "Turbulence generated by MHD instabilities can play an important role in enhancing transport processes in various astrophysical bodies, such as accretion and protoplanetary disks, galaxies, stellar radiative zones, etc. The anomalous turbulent transport can be particularly important in magnetized gas where a wide variety of MHD instabilities  can occur (see, e.g., Barnes et al. 1999). In this case, the onset of instability can be caused both by hydrodynamic motions (for instance, differential rotation; see, e.g., Velikhov 1959; Chandrasekhar 1960) or unstable magnetic configurations. { Which field strength and topology can sustain a stable magnetic configuration is still rather uncertain despite extensive work (see Borra et al. 1982; Mestel 1999 for review).}  Most likely, the best-studied magnetic configuration is one with a purely toroidal field. Ever since the paper by Tayler (1973), it has been known that toroidal fields can be unstable close to the axis of symmetry, if there is a non-zero electric current density on the axis. The growth rate of this instability is expected to be of the order of the time taken for an Alfv\\'en wave to travel around the star on a toroidal field line. However, even a purely toroidal field is stable if it decreases rapidly with the cylindrical radius $s$. For instance, Tayler (1973; see also Chanmugam 1979) argued that the toroidal field $B_{\\varphi}$ is stable against axisymmetric perturbations if it satisfies the condition $d ( B_{\\varphi}/s)/ds < 0$ and to non-axisymmetric perturbations if $d (s B_{\\varphi}^2)/ds <0$. Note that a purely toroidal field can also be subject to the magnetic buoyancy instability (Parker 1955; Gilman 1970; Acheson 1978) but the Tayler instability likely appears first as the strength of the toroidal field increases (Spruit 1999).  The stability of a purely toroidal field in the radiative zones of stars and accretion disks has been studied by a number of authors. Numerical modeling by Braithwaite (2006) confirms that the toroidal field with $B_{\\varphi} \\propto s$ or $\\propto s^2$ is unstable to the $m=1$ mode ($m$ is the azimuthal wave number) as predicted by Tayler (1973). The linear stability of the toroidal field in rotating stellar interiors has been considered by  Kitchatinov \\& R\\\"udiger (2007), who conclude that the magnetic instability is essentially three-dimensional and that the finite thermal conductivity has a strong destabilizing effect. Terquem \\& Papaloizou (1996) and Papaloizou \\& Terquem (1997) considered the stability of an accretion disk with an embedded toroidal magnetic field. These authors find that the disks containing a purely toroidal field are always unstable and obtained spectra of unstable modes in the local approximation. They argue that one type of modes is driven primarily by buoyancy, while the other is driven by shear independently of the magnetic configuration.  { Stability properties of purely poloidal magnetic configurations  have also been well-studied. Since the papers by Wright (1973) and Markey \\& Tayler (1973, 1974), it is known that a poloidal field is subject to dynamical instabilities in the neighbourhood of neutral points/lines if the field lines are closed inside the star. These authors recognizes that the magnetic field in the neighbourhood of a neutral line resembles that of a toroidal pinched discharge that is known to be unstable. Although instabilities involving significant displacements in the direction of gravity were strongy inhibited, other instabilities were not affected. The instability of poloidal configurations is rather fast: its growth time can reach a few Alfv\\'en crossing time (Van Assche et al. 1982; Braithwaite \\& Spruit 2006). With numerical simulations, the stability of poloidal magnetic configurations has been studied by Braithwaite \\& Spruit (2006), who apply the results to the internal magnetic configuration of neutron stars. Note, however, that a toroidal field might exert a stabilising influence on the instabilities of a poloidal field in the neighbourhood of neutral points (Tayler 1980).}  { On the contrary, the addition of even a relatively weak poloidal field alters the stability of the toroidal field substantially. For example, as first shown by Howard \\& Gupta (1962; see also Knobloch 1992; Dubrulle \\& Knobloch 1993), a necessary (but not sufficient) condition for the instability of a toroidal field in the presence of the axial field reads} \\begin{equation}\\label{1} \\frac{d}{ds} (s B_{\\varphi}) > 0. \\end{equation} Howard \\& Gupta (1962) argue that, for a fixed value of $m$, the growth rate of instability caused by condition (1) must vanish in the limit of a vanishing axial magnetic field, thereby providing a connection with the stability criterion obtained by Tayler (1973). Note that the presence of a radial field is also crucial for stability properties of rotating magnetic configurations (Bonanno \\& Urpin 2006). { It turns out that configurations containing both toroidal and poloidal fields are more stable than purely toroidal or purely poloidal ones (Prendergast 1956; Tayler 1980). With numerical simulations, Braithwaite \\& Nordlund (2006) studied the stability of a random initial field in the stellar radiative zone. The star was modeled on a Cartesian grid, and the authors found that the stable magnetic configurations generally have the form of tori with comparable poloidal and toroidal field strengths.}  In the present paper, we address the stability properties of magnetic configurations by considering the stability of the toroidal magnetic field with respect to axisymmetric perturbations in the presence of axial fields of a various strength. We show that the instability may occur in such magnetic field configurations under the conditions that differ substantially from those imposed by the Tayler criterion or the necessary condition (\\ref{1}). We argue that, in some cases, the instability is caused by the new type of MHD waves with the growth rate proportional to $\\sqrt{B_z B_{\\varphi}}$ where $B_z$ is the axial magnetic field. Depending on the profile $B_{\\varphi}(s)$ and the ratio $B_z/B_{\\varphi}$, the instability can occur in two regimes that have substantially different growth rates. We also show that the range of unstable wavelength in the $z$-direction can be essentially different, depending on the $B_z/B_{\\varphi}$ ratio. ",
        "conclusions": "{ We have considered the hydromagnetic stability of cylindrical configurations containing the toroidal and axial magnetic fields. Dissipative effects were neglected in our study. We treated a linear stability assuming that the behaviour of small perturbations is governed by equations of incompressible hydrodynamics. This approximation is  justified if the magnetic field is subthermal and the Alfv\\'en velocity is low compared to the sound speed.} The stability of the magnetic configurations is a key issue for understanding the properties of various astrophysical bodies, such as peculiar A and B stars, magnetic white dwarfs, neutron stars, etc. { Even though  various dynamo models predict that the toroidal field should be typically stronger than the poloidal one, the effect of a poloidal field on the stability usually cannot be neglected.}  To demonstrate this, we treated the simplest model of a highly conducting fluid between two cylindrical surfaces. We assumed that the toroidal and axial fields depends on the cylindrical radius alone. In a short-wavelength approximation, we  derived the growth rate and a sufficient criterion of instability analytically (Eqs.~(30)-(31)). For large-scale perturbations, the condition of instability and its growth rate were calculated numerically. The analytical and numerical results are in  good qualitative agreement. The obtained conditions of instability differ substantially from what is predicted by the necessary condition (1). For instance, according to Eq.~(1), if the instability occurs in the magnetic configuration, then the toroidal field profile satisfies the condition $\\alpha > -1$. In fact, the instability occurs only if the toroidal field decreases with $s$ much slower (or even increases): the critical value of $\\alpha$ is $\\approx -0.1$ if $B_z/ B_{\\varphi 0} =0.1$ and $\\approx 1$ if $B_z/ B_{\\varphi 0} =0.5$. { If $B_{z}$ depends on $s$, then the critical values of $\\alpha$ should be even higher.}  Depending on the profile of the toroidal field and the strength of the axial field, the instability can arise in two essentially different regimes. In the case of a weak axial field, $B_{\\varphi 0} \\gg B_z$, the value of $\\alpha$ that distinguishes between the regimes is $\\approx 1$. If $\\alpha > 1$, then the instability grows on the Alfv\\'en timescale determined by the toroidal field and is rather fast. If $\\alpha < 1$, then the instability is much slower and grows on the timescale determined by the axial field. The transition between these two regimes occurs at larger $\\alpha$ if the axial field increases. The efficiency of the considered instability turns out to be rather low if $\\alpha < 1$ and $B_z$ is weak.  It is worth noticing the very particular properties of instability in the case $\\alpha \\approx 1$. In such a configuration, the particular type of MHD waves is given by the dispersion equation (19). The growth rate of these waves (or the frequency, if the waves are stable) is proportional to the product of $B_z$ and $B_{\\varphi}$. These waves cannot exist in purely toroidal or purely poloidal fields. The instability of the configuration with $\\alpha \\approx 1$ is caused by the generation of this particular type of wave. { These waves can probably determine the instability of magnetic configurations near the axis of symmetry where $B_{\\varphi} \\rightarrow 0$.}  A sufficiently strong axial field always suppresses the instability. For more or less plausible values of $\\alpha \\leq 1$, the strength of the axial field stabilising the configuration is $\\sim 0.1-1 B_{\\varphi}$. A much stronger field is required, however, to stabilise the configuration with larger $\\alpha$.   \\vspace{0.5cm}  \\noindent {\\it Acknowledgments.} This research project was  supported by a Marie Curie Transfer of Knowledge Fellowship of the European Community's Sixth Framework Programme under contract number MTKD-CT-002995. VU thanks also INAF-Ossevatorio Astrofisico di Catania for hospitality."
    },
    "0710/0710.3682_arXiv.txt": {
        "abstract": "We report on the latest (2007 Jan) observations of supernova remnant (SNR) 1987A from the {\\it XMM-Newton} mission.  Since the 2003 May observations of Haberl et al. (2006), 11 emission lines have experienced increases in flux by factors $\\sim 3$ to 10 ($6 \\pm 0.6$ on average), with the 775 eV line of O~{\\sc viii} showing the greatest increase.  Overall, we are able to make Gaussian fits to 17 emission lines in the {\\it RGS} spectra and obtain line fluxes; we have observed 6 lines of Fe~{\\sc xvii} and Fe~{\\sc xviii} previously unreported by {\\it XMM-Newton}.  A two-shock model representing plasmas in non-equilibrium ionization is fitted to the {\\it EPIC-pn} spectra, yielding temperatures of $\\sim 0.4$ and $\\sim 3$ keV, as well as elemental abundances for N, O, Ne, Mg, Si, S and Fe.  We demonstrate that the abundance ratio of N and O can be constrained to $\\lesssim 20\\%$ accuracy $\\left(\\mbox{N/O} = 1.17 \\pm 0.20 \\right)$.  Within the same confidence interval, the same analysis suggests that the C+N+O abundance varies from $\\sim 1.1$ to $1.4 \\times 10^{-4}$, verifying the {\\it Chandra} finding by Zhekov et al. (2006) that the C+N+O abundance is lower by a factor $\\sim 2$ compared to the value obtained in the optical/ultraviolet study by Lundqvist \\& Fransson (1996).  Normalizing our obtained abundances by the Large Magellanic Cloud (LMC) values of Hughes, Hayashi \\& Koyama (1998), we find that O, Ne, Mg and Fe are under-abundant, while Si and S are over-abundant, consistent with the findings of Aschenbach (2007).  Such a result has implications for both the single-star and binary accretion/merger models for the progenitor of SNR 1987A.  In the context of the binary merger scenario proposed by Morris \\& Podsiadlowski (2006, 2007), material forming the inner, equatorial ring was expelled after the merger, implying that either our derived Fe abundance is inconsistent with typical LMC values or that iron is under-abundant at the site of Sanduleak -69$^\\circ$202. ",
        "introduction": "\\label{sect:intro}  SNR 1987A is the Rosetta Stone (Allen 1960) of Type II supernova remnants, resolved and well-studied in multiple wavebands, including the infrared (Bouchet et al. 2006; Kjaer et al. 2007), optical/ultraviolet (Gr\\\"{o}ningsson et al. 2007; Heng 2007) and radio (Gaensler et al. 2007).  The physical mechanism partially powering optical/ultraviolet emission from the reverse shock is the same as the one at work in Balmer-dominated supernova remnants (Heng \\& McCray 2007; Heng et al. 2007).  The detection of a neutrino burst confirmed the core collapse nature of the progenitor (Koshiba et al. 1987; Svoboda et al. 1987), though a pulsar has yet to be detected (Manchester 2007).  A system of three rings may be the result of a binary merger between two massive stars about 20,000 years prior to the supernova explosion (Morris \\& Podsiadlowski 2006, 2007; hereafter MP0607).  Reviews of the multi-wavelength studies of SNR 1987A can be found in McCray (1993, 2005, 2007).  Mixing of the stellar envelope and core by Rayleigh-Taylor instabilities within the progenitor star, Sanduleak -69$^\\circ$202, has been invoked to explain the early emergence of the 847 keV $\\gamma$-ray line from SNR 1987A, which was predicted by Shibazaki \\& Ebisuzaki (1988) to reach its peak around 1.1 years after the explosion, if one assumes a mixed mass of about 5$M_\\sun$ (Ebisuzaki \\& Shibazaki 1988).  Instead, Matz et al. (1988) observed the 847 keV line $\\sim$ 6 to 8 months post-explosion, suggesting even more extensive mixing of $^{56}$Co than assumed.  A similar explanation (Ebisuzaki \\& Shibazaki 1988) was given for the early emergence of 16 to 28 keV X-rays (Sunyaev et al. 1990; Inoue et al. 1991).  The $\\gamma$-rays originate from the radioactive decay of $^{56}$Co, while the X-rays are from the Compton degradation of the $\\gamma$-rays (McCray, Shull \\& Sutherland 1987).  In the soft X-ray, SNR 1987A was first observed by Beuermann, Brandt \\& Pietsch (1994).  Subsequently, Hasinger, Aschenbach \\& Tr\\\"{u}mper (1996) tracked a steady increase of the soft X-ray flux over 4 years with {\\it ROSAT}.  Extensive work has since been done by the {\\it Chandra} (Michael et al. 2002; Park et al. 2002, 2004, 2005 [hereafter P05], 2006, 2007 [hereafter P07]; Zhekov et al. 2005, 2006 [hereafter Z06]) and {\\it XMM-Newton} groups (Haberl et al. 2006, hereafter H06; Aschenbach 2007, hereafter A07). The general picture gleaned from these studies is of a bimodal plasma distribution present in the region between the forward and reverse shocks (Fig. \\ref{fig:87a}).  The soft X-rays ($\\sim$ 0.3 to 0.5 keV) are from the decelerated shock front interacting with dense protrusions (``fingers'') on the inner, equatorial ring, while the ``hard'' X-rays ($\\sim$ 2 to 3 keV) are from a fast shock propagating into more tenuous material.  The soft X-rays appear to be correlated with optical ``hot spots'', believed to be emission from the shocked fingers, appearing around the equatorial ring, while the hard X-ray and radio images exhibit structures that coincide.  Between Days 6000 and 6200, the soft X-ray light curve experienced an upturn and departure from an exponentially increasing profile, which P05 interpreted as evidence that the blast wave had reached the main body of the dense circumstellar material of the equatorial ring.  Future studies on the nature of the soft X-ray light curve are relevant to the issue of pre-ionization of the supernova ejecta, which can potentially extinguish H$\\alpha$ and Ly$\\alpha$ emission from the reverse shock (Smith et al. 2005; Heng et al. 2006).  In the studies described, relatively little attention has been paid to the subject of inferring elemental abundances from X-ray analyses. Fits to the X-ray spectra yield N, O, Ne, Mg, S, Si and Fe abundances; such results have been tabulated and studied by Z06 and H06, using {\\it Chandra} and {\\it XMM}, respectively.  Here, we examine such an approach using the latest {\\it XMM} data set of SNR 1987A, taken in early 2007.  In \\S\\ref{sect:data}, we describe our observations and data reduction techniques.  Our results are presented in \\S\\ref{sect:results}.  In \\S\\ref{sect:discussion}, we perform a detailed error analysis of individual abundances and their ratios. We find that the N/O ratio as well as the individual N and O abundances can be constrained to $\\sim 20\\%$ accuracy.  Our derived elemental abundances for O, Ne, Mg and Fe are under-abundant, while Si and S are over-abundant, relative to typical values for the Large Magellanic Cloud (LMC; Hughes, Hayashi \\& Koyama 1998).  Such a result has implications for modeling the progenitor of SNR 1987A, which we discuss. ",
        "conclusions": "\\label{sect:discussion}  \\subsection{INDIVIDUAL ABUNDANCES \\& THEIR RATIOS} \\label{subsect:abund}  We compare the derived elemental abundances to those of Z06 and H06 and list them {\\it relative to hydrogen} in Table \\ref{tab:abundances}; uncertainties in the AG89 and W00 abundance tables are not propagated.  We emphasize that care must be taken to specify the abundance table used, as this may lead to widely differing values of the derived abundances (relative to hydrogen).  In modeling the LMC absorption, Z06 used the elemental abundance table of AG89, while H06 chose the table of W00 because the lower oxygen abundance fitted the K absorption edge in the {\\it EPIC} data better.  Fransson et al. (1989) found N/O = 1.6 $\\pm$ 0.8, about 12 times higher\\footnote{The AG89 and W00 values for solar nitrogen-to-oxygen abundance are N/O = 0.132 and 0.110 by number, respectively.} than the AG89 solar value, which they interpret as evidence of substantial CNO processing.  LF96 found N/O = 1.1 $\\pm$ 0.4, while Sonneborn et al. (1997) found N/O = 1.7 $\\pm$ 0.5.  All three of these values were derived from optical/ultraviolet data.  Next, we turn our attention to the N/O ratio derived from X-ray studies.  Linearly propagating the errors listed by Z06, we find that their results yield N/O = $1.10^{+0.47}_{-0.45}$; they remark that their derived C, N and O abundance, C+N+O $\\approx 1.98 \\times 10^{-4}$, is lower by about a factor of 2 compared to the $3.72 \\times 10^{-4}$ value of LF96.  Model A ({\\tt VNEI+VRAYMOND}) in H06 yields C+N+O $\\approx 1.67 \\times 10^{-4}$ and a rather wide range in the nitrogen-to-oxygen ratio, N/O = $1.33^{+1.47}_{-0.71}$.  Our {\\tt VPSHOCK+VPSHOCK} fit to the {\\it EPIC-pn} data yields C+N+O $\\approx 1.29 \\times 10^{-4}$ and N/O = $1.17^{+0.37}_{-0.34}$ (using the W00 table); we call this combination of N and O the ``best fit point''. Note that the C abundance is held fixed at 0.09 relative to solar ($\\sim 3 \\times 10^{-5}$ relative to hydrogen) for the C+N+O values derived from the X-ray studies.  We next perform a more careful analysis of the N/O ratio.  We first generate a $\\chi^2$ map quantifying the inter-dependence of the fits to the N and O abundance.  Contour lines in the $\\chi^2$ map form ``error ellipses'', which are shown for different $\\Delta \\chi^2$ values from the best fit point (Fig. \\ref{fig:NO}).  At the $\\Delta \\chi^2 = 2.706$ level, N/O $= 1.17 \\pm 0.20$ for the {\\it EPIC-pn} data.  In linearly propagating the errors in the individual abundances, one is in essence adopting the largest possible range of ratios, which can be visualized as the edges of a rectangle in the contour map.  Our error analysis improves the uncertainties because it considers only values of the abundance ratio within the specified contour.  Within the same confidence interval considered, the corresponding C+N+O value is from $\\sim$ 1.1 to $1.4 \\times 10^{-4}$. We see that the errors in the N/O ratio and the individual N and O abundances can be constrained at the $\\sim 20\\%$ level.  We thus confirm the C+N+O under-abundance noted by the {\\it Chandra} studies of Z06, who suggest a couple of physical reasons for such a result: the sub-LMC abundance of C+N+O within the progenitor star, Sanduleak -69$^\\circ$202, and/or an extra source of possibly non-thermal, X-ray continuum.  We perform the same analysis for N/S (Fig. \\ref{fig:othercontours}).  Again using the W00 table, linear propagation of the errors in the abundances obtained from the {\\tt VPSHOCK+VPSHOCK} fit yields  N/S = $4.59^{+1.51}_{-1.36}$, while the error ellipse analysis gives N/S = $4.59^{+1.59}_{-1.33}$.  The iron, nitrogen and oxygen lines are predominantly from the lower energy part ($\\lesssim 1$ keV) of the spectrum, while the sulphur lines are situated between $\\sim 2.2$ and 3.1 keV.  (The silicon lines are located between $\\sim 1.8$ and 2.2 keV.)  Abundance ratios based on lines of widely differing energies lead to rounder error ellipses --- ``error circles''.  In such cases, a more thorough error analysis will not constrain the abundance ratio better, as in the case of N/S.  This is partially an instrumental effect --- when the energy resolution is comparable to the spacing of the lines, the line complexes overlap and are only partially resolved.  There is also the issue of choosing a coordinate system in which the fitting parameters are ``orthogonal''.  Hydrogen and helium have no X-ray lines --- their abundances are derived from the strength of the X-ray continuum.  When the pair of elements considered are situated close to each other in energy, the hydrogen abundance has to first order a linear dependence on the continuum strength.  Thus, the ratio of the considered abundances is tightly constrained as the individual abundances track each other closely.  By contrast, when the pair of abundances considered are located far apart in energy, this linear dependence of hydrogen abundance on the continuum is broken as the dominant uncertainties are in temperature rather than in flux.  Orthogonality is now absent.  To attain orthogonality for such pairs of lines, one has to construct models that directly fit to the abundance ratio considered, an approach which is not explored in this paper.  The error in the abundance ratios increases as one moves a given $\\Delta \\chi^2$ away from the minimum point.  In Fig. \\ref{fig:ratioerror}, we compute the mean error sustained by the various ratios as a function of $\\Delta \\chi^2$.  According to Avni (1976), if three interesting parameters are considered (see \\S\\ref{subsect:spectral}), the 90\\% confidence interval is situated at $\\Delta \\chi^2 = 6.25$.  In this case, the N/O abundance ratio suffers from errors $\\sim 25\\%$.   \\subsection{THE PROGENITOR OF SNR 1987A: SINGLE-STAR OR BINARY MODEL?} \\label{subsect:metallicity}  A more revealing approach to analyze the elemental abundances is to normalize the results listed in Table \\ref{tab:abundances} by the ``canonical'' values of the LMC abundances (Hughes, Hayashi \\& Koyama 1998).  These were derived using a sample of 7 middle-aged SNRs in the LMC (N23, N49, N63A, DEM 71, N132D, 0453-68.5 and N49B).  This approach was first explored by A07, who showed that the normalized abundances appear to cluster in two groups: N, O, Ne, Mg and Fe are slightly more than half their LMC values, while Si, S and Ni exceed their LMC values.  We generalize the A07 approach by considering both sets of abundances derived from using the AG89 and W00 tables.  The normalized abundance, relative to its respective LMC value, is $R_{\\rm{87A/LMC}}$; it is plotted as a function of the elemental mass number in Fig. \\ref{fig:abundratios}.  The error bars for $R_{\\rm{87A/LMC}}$ are computed by linearly propagating the errors listed in Table \\ref{tab:abundances} and in Hughes, Hayashi \\& Koyama (1998).  We caution that additional systematic errors may be present that are not taken into account.  For example, the derived abundance for Fe is a sensitive function of temperature and may vary substantially when small changes are made to $T_{\\rm{low}}$ \\footnote{The Fe abundance obtained from the {\\tt VPSHOCK+VPSHOCK} fit (W00 table) varies by $\\sim 16\\%$ when $T_{\\rm{low}}$ is changed by $\\sim 10\\%$.}.  We see that the elements O, Ne, Mg and Fe are under-abundant, while Si and S are over-abundant, consistent with the findings of A07.  With the exception of Fe, there is a tendency for $R_{\\rm{87A/LMC}}$ to increase with larger elemental mass number, a trend that is independent of the AG89 or W00 tables, though we note that it is more pronounced with the latter.  The Fe abundance derived is essentially independent of the AG89 or W00 tables, and is under-abundant by about 70\\% relative to the LMC \\footnote{The O abundance relative to H for the AG89 and W00 tables are the same.}.  The under-abundance of Fe and O was previously noted by Hasinger et al. (2006), who argued for the existence of iron-oxygen ``rust grains''.  The reduced abundance of Fe alone suggests that the iron is locked up in dust grains.  However, Dwek \\& Arendt (2007) showed from an analysis of the infrared-to-X-ray flux ratio --- ${\\cal R}_{\\rm{IRX}} < 1$ versus the theoretically expected value of $\\sim 10^2$ to $10^3$ --- that the dust in SNR 1987A is severely depleted compared to standard dust-to-gas mass ratios in the LMC, suggesting low dust condensation efficiency or dust destruction in the hot X-ray gas.  In fact, ${\\cal R}_{\\rm{IRX}}$ was shown to {\\it decrease} with time, which is direct evidence for dust destruction.  Our derived plasma temperatures are consistent with this scenario --- even if dust could form, it would be destroyed at these temperatures.  In light of Fig. \\ref{fig:abundratios}, the central question to ask is whether the progenitor of SNR 1987A arose from a single star or a binary system?  Sanduleak -69$^\\circ$202 was known to be a blue supergiant (BSG) at the time of the supernova explosion, contrary to the expectation that massive stars end their lives as red supergiants (RSGs).  Observations of low-velocity, nitrogen-rich circumstellar material are interpreted as the progenitor star being a RSG until about $\\sim 20,000$ years before its death (Fransson et al. 1989).  Such a time scale has in turn been interpreted as the Kelvin-Helmholtz time of the helium core (Woosley et al. 1997).  This BSG-RSG-BSG evolution remains one of the greatest challenges for the single-star model (Woosley et al. 1997; Woosley, Heger \\& Weaver 2002), as is the observed system of three rings produced $\\sim 20,000$ years before the explosion.  The favored single-star models require a combination of reduced metallicity and ``restricted semi-convection'' (Woosley 1988), the former of which is supported by our derived abundance ratios.  An additional challenge for the single-star model is to reproduce the over-abundance of Si and S (Fig. \\ref{fig:abundratios}), which may require the invoking of some non-standard mixing process (H.-T. Janka 2007, private communication).  Binary solutions to Sanduleak -69$^\\circ$202 are sub-divided into accretion and merger models (see Woosley, Heger \\& Weaver [2002] and references therein).  The binary accretion models allow helium- and nitrogen-rich material to be added to the progenitor star and require the disappearance of the mass donor in an earlier supernova event.  In the binary merger scenario proposed by MP0607, two stars with masses $\\sim 5 M_\\sun$ and $\\sim 15 M_\\sun$ are initially orbiting each other with a period $\\sim 10$ years.  The more massive companion transfers mass to the less massive star only after the former has completed helium burning in the core.  A common envelope (Paczy\\'{n}ski 1976) is formed, during which core material from the primary star is dredged up to the surface.  The merger process takes a few hundred years, culminating in an initially over-sized RSG, which loses its excess thermal energy over a few thousand years to become a BSG.  The spun-up, rapidly rotating BSG produces a fast stellar wind, sweeping up ejecta associated with the merger, producing the triple ring nebula we now see in projection.  The nearly axi-symmetric but highly non-spherical nature of the rings suggests that rotation played a role in their formation and is consistent with the proposed scenario.  The beauty of the model lies in the fact that it requires no physically ad hoc assumptions --- apart from a small kick of $\\sim 2$ km s$^{-1}$ given to the ejecta to displace the center of the outer rings from their symmetry axis --- and makes a number of predictions.  In their favored model, MP0607 assert that the outer rings are ejected before the stellar core material is dredged up, while the inner, equatorial ring --- the site of the observed X-ray emission --- is ejected {\\it afterwards} \\footnote{As noted by MP0607, this hypothesis is verifiable/refutable, as the inner ring should exhibit helium enhancement and more CNO processing relative to the outer rings.}.  A relevant consequence of this model is that the dredged-up heavy elements will manifest themselves in the form of X-ray emission lines.  This may explain the trend we see in Fig. \\ref{fig:abundratios} --- the challenge for binary merger models is to reproduce the derived $R_{\\rm{87A/LMC}}$ values.  Stellar nucleosynthesis can add and {\\it not} subtract iron --- in the context of the MP0607 binary merger model, we expect $R_{\\rm{87A/LMC}}$(Fe) $\\ge 1$.  We are again led to the question: where is the iron?  If the Fe abundance derived is an upper limit on the iron used to form Sanduleak -69$^\\circ$202, then it is clearly inconsistent with ``standard'' Fe abundances in the LMC.  An alternative interpretation is that there is a strong spatial variation in the Fe abundance throughout the LMC, such that the iron is sub-LMC at the site of SN 1987A and equal to its LMC abundance elsewhere.  The C+N+O under-abundance suggested in \\S\\ref{subsect:abund} supports such a conclusion.  An improvement over using the Hughes, Hayashi \\& Koyama (1998) {\\it ASCA} abundance values is to re-analyze and expand upon their SNR sample using {\\it Chandra} and {\\it XMM}.  Existing studies (e.g., Hughes et al. 2006) tend to pick out regions of interest that may include ejecta enrichment; instead, X-ray emission should be extracted from the {\\it entire} blast wave, from which average abundances can be inferred.  Future studies will be invaluable towards resolving these issues.   \\scriptsize K.H. acknowledges the kind hospitality and financial support of: the Max Planck Institutes for Astrophysics (MPA) and Extraterrestrial Physics (MPE) during June to October 2007, where he was a visiting postdoctoral scientist; and the Lorentz Center (Leiden) during the August 2007 workshop, ``From Massive Stars to Supernova Remnants''.  He thanks Dick McCray, Sangwook Park, Svet Zhekov, Jack Hughes, John Raymond, Philipp Podsiadlowski, Rashid Sunyaev, Peter Lundqvist, Claes Fransson, Dmitrijs Docenko, Thomas Janka, Carlos Badenes, Roger Chevalier, Nathan Smith and Jacco Vink for engaging and helpful discussions.  Special mentions go out to: Dick and John, who provided a crash course on non-equilibrium ionization plasmas during a sunny bicycle ride in Leiden; Svet, who pointed out relevant material on X-ray physics, as well as for critical comments following his meticulous scrutiny of the manuscript. The authors are collectively grateful to Kazik Borkowski for providing the updated {\\tt VPSHOCK} model files for use in {\\tt XSPEC}.  The {\\it XMM-Newton} project is supported by the {\\it Bundesministerium f\\\"{u}r Wirtschaft und Technologie/Deutsches Zentrum f\\\"{u}r Luft- und Raumfahrt} (BMWI/DLR, FKZ 50 OX 001) and the Max Planck Society. \\normalsize"
    },
    "0710/0710.4294_arXiv.txt": {
        "abstract": "Infrared surveys indicate that the dust content in debris disks gradually declines with stellar age. We simulated the long-term collisional depletion of debris disks around solar-type (G2~V) stars with our collisional code. The numerical results were supplemented by, and interpreted through, a new analytic model. General scaling rules for the disk evolution are suggested. The timescale of the collisional evolution is inversely proportional to the initial disk mass and scales with radial distance as $r^{4.3}$ and with eccentricities of planetesimals as $e^{-2.3}$. Further, we show that at actual ages of debris disks between 10~Myr and 10~Gyr, the decay laws of the dust mass and the total disk mass are different. The reason is that the collisional lifetime of planetesimals is size-dependent. At any moment, there exists a transitional size, which separates larger objects that still retain the ``primordial'' size distribution set in the growth phase from smaller objects whose size distribution is already set by disruptive collisions. The dust mass and its decay rate evolve as that transition affects objects of ever-larger sizes. Under standard assumptions, the dust mass, fractional luminosity, and thermal fluxes all decrease as $t^\\xi$ with $\\xi = -0.3$...$-0.4$. Specific decay laws of the total disk mass and the dust mass, including the value of $\\xi$, largely depend on a few model parameters, such as the critical fragmentation energy as a function of size, the primordial size distribution of largest planetesimals, as well as the characteristic eccentricity and inclination of their orbits. With standard material prescriptions and a distribution of disk masses and extents, a synthetic population of disks generated with our analytic model agrees quite well with the observed Spitzer/MIPS statistics of 24 and 70 \\micron\\ fluxes and colors versus age. ",
        "introduction": "Since the IRAS discovery of the excess infrared emission around Vega by \\citet{aumann-et-al-1984}, subsequent infrared surveys with ISO, Spitzer and other instruments have shown the Vega phenomenon to be common for main-sequence stars. The observed excess is attributed to second-generation circumstellar dust, produced in a collisional cascade from planetesimals and comets down to smallest grains that are blown away by the stellar radiation. While the bulk of such a debris disk's mass is hidden in invisible parent bodies, the observed luminosity is dominated by small particles at dust sizes. Hence the studies of dust emission offer a natural tool to gain insight into the properties of planetesimal populations as well as planets that may shape them and, ultimately, into the evolutionary history of circumstellar planetary systems.  In recent years, various photometric surveys of hundreds of nearby stars have been conducted with the Spitzer Space Telescope. These are the GTO survey of FGK stars \\citep{beichman-et-al-2005,bryden-et-al-2006,beichman-et-al-2006b}, the FEPS Legacy project  \\citep{meyer-et-al-2004,kim-et-al-2005}, the A star GTO programs \\citep{rieke-et-al-2005,su-et-al-2006}, the young cluster programs \\citep{gorlova-et-al-2006}, and others. These observations were done mostly at 24 and 70~\\textmu{m} with the MIPS photometer, but also between 5 and 40~\\textmu{m} with the IRS spectrometer \\citep{jura-et-al-2004,chen-et-al-2006}. Based on these studies, about 15\\% of mature solar-type (F0--K0) stars have been found to harbor cold debris disks at 70~\\textmu{m}. For cooler stars, the fraction drops to 0\\%--4\\% \\citep{beichman-et-al-2006b}. For earlier spectral types, the proportion increases to about 33\\% \\citep{su-et-al-2006}. At 24~\\textmu{m}, the fraction of systems with detected excess stays similar for A~stars, but appreciably decreases for FGK ones. Similar results in the sub-millimeter range are expected to become available soon from a survey with SCUBA and SCUBA2 on JCMT \\citep{matthews-et-al-2007}. Preliminary SCUBA results for M dwarfs suggest, in particular, that the proportion of debris disks might actually be higher than suggested by Spitzer \\citep{lestrade-et-al-2006}.  All authors point out a decay of the observed infrared excesses with systems' age. However, the values reported for the slope of the decay, assuming a power-law dependence $t^{-\\alpha}$, span a wide range. \\citet{greaves-wyatt-2003} suggest $\\alpha \\la 0.5$, \\citet{liu-et-al-2004} give $0.5 < \\alpha < 1.0$, \\citet{spangler-et-al-2001} report $\\alpha \\approx 1.8$, and \\citet{greaves-2005} and \\citet{moor-et-al-2006} derive $\\alpha \\approx 1.0$. Fits of the upper envelope of the distribution of luminosities over the age yield $\\alpha \\approx 1.0$ as well \\citep{rieke-et-al-2005}. Besides, the dust fractional luminosity exhibits a large dispersion at any given age.  In an attempt to gain theoretical understanding of the observed evolution, \\citet{dominik-decin-2003} assumed that equally-sized ``comets'' produce dust through a cascade of subsequent collisions among ever-smaller objects. If this dust is removed by the same mechanism, the steady-state amount of dust in such a system is proportional to the number of comets. This results in an $M/M_0 \\approx \\tau/t$ dependence for the amount of dust and for the number of comets or the total mass of the disk. Under the assumption of a steady state, this result is valid even for more complex systems with continuous size distributions from planetesimals to dust. Tenuous disks, where the lifetime of dust grains is not limited by collisions but by transport processes like the Poynting-Robertson drag \\citep{artymowicz-1997,krivov-et-al-2000,wyatt-2005}, follow $M \\propto t^{-2}$ rather than $M \\propto t^{-1}$.  More recently, \\citet{wyatt-et-al-2007a} lifted the most severe simplifying assumption of the Dominik-Decin model, that of equal-sized parent bodies, and included them into the collisional cascade. A debris disk they consider is no longer a two-component system ``comets + dust''. Instead, it is a population of solids with a continuous size distribution, from planetesimals down to dust. A key parameter of the description by \\citet{dominik-decin-2003} is the collisional lifetime of comets, $\\tau$. \\citet{wyatt-et-al-2007a} replaced it with the lifetime of the largest planetesimals and worked out the dependencies on this parameter in great detail. Since the collisional timescale is inversely proportional to the amount of material, $\\tau \\propto 1/M_0$, the asymptotic disk mass becomes independent of its initial mass. Only dynamical quantities, i.e. the disk's radial position and extent, the orbiting objects' eccentricities and inclinations, and material properties, i.e. the critical specific energy and the disruption threshold, as well as the type of the central star determine the very-long-term evolution.  Still, there are two important simplifications made in the model by \\citet{wyatt-et-al-2007a}: (i) the disk is assumed to be in collisional equilibrium at all sizes, from dust up to the largest planetesimals and (ii) the minimum specific energy needed to disrupt colliding objects is independent of their size. As a consequence of (i) and (ii), the size distribution of solids is a single power-law. To check how reasonable these assumptions are, realistic simulations of the disks with collisional codes are necessary \\citep[e.g.,][]{thebault-et-al-2003,krivov-et-al-2005,krivov-et-al-2006,thebault-augereau-2007}.  The aim of this paper is two-fold. First, we follow the evolution of debris disks with our elaborate numerical code  \\citep{krivov-et-al-2005,krivov-et-al-2006} to check the existing analytic models and the assumptions (i) and (ii) they are based upon. Second, in order to make these numerical results easier to use, we develop a new analytic model for the evolution of disk mass and dust mass that relaxes both assumptions (i) and (ii) above.  Section~\\ref{secNumerics} summarizes the basic ideas and assumptions and describes our numerical model and the runs of the collisional code. In Section~\\ref{secScalings} the numerical results are presented and dependences of the collisional timescale on the disk mass, distance to the star, and mean eccentricity of parent bodies are derived. In section~\\ref{secAnalytics}, the analytic model for the evolution of disk mass and dust mass is developed. Section~\\ref{secEvolLuminosity} analyzes the evolution of dust luminosities. In Section~\\ref{secObservations}, we use the analytic model to synthesize representative populations of debris disks and compare them with statistics of debris disks derived from the Spitzer surveys. A summary is given and conclusions are drawn in Section~\\ref{secConclusions}. \\pagebreak ",
        "conclusions": " \\begin{enumerate} \\item The timescale of the collisional evolution is inversely proportional to the initial disk mass. For example, halving the total mass doubles all collisional timescales. This rule is valid for systems where collisions are the only loss mechanism of particles and only as long as $\\beta$-meteoroids are unimportant for the collisional budget. \\item Numerics and analytics consistently yield a $\\tau\\propto  r^{4.3}$ dependence of the timescale of the collisional evolution on the radial distance. \\item Numerical simulations show that the collisional timescale varies with the average eccentricity of dust parent bodies as $\\tau \\propto e^{-2.3}$. The analytic approach suggests a somewhat weaker dependence, $\\tau\\propto e^{-5/3}$. \\item An evolving three-slope size distribution is proposed to approximate the numerical results. The biggest objects are still distributed primordially, with a slope $q\\sbs{p}$. The objects below a certain transitional size are already reprocessed by collisions and thus have a quasi-steady-state size distribution, determined by their self-gravity (for intermediate-sized objects, slope $q\\sbs{g}$) or by material strength (for smallest objects, slope $q\\sbs{s}$). That transitional size corresponds to the largest objects for which the collisional lifetime is still shorter than the age of the system. The transitional size increases with time, meaning that ever-larger planetesimals get involved into the collisional cascade. \\item At actual ages of debris disks, $\\sim$10~Myr to $\\sim$10~Gyr, the decay of the dust mass and the total disk mass follow {\\em different} laws. The reason is that, in all conceivable debris disks, the largest planetesimals have longer collisional lifetimes than the system's age, and therefore did not have enough time to reach collisional equilibrium. If the system were let to evolve for sufficiently long time, both dust mass and disk mass would start to follow $t^{-1}$. However, this requires time spans of much longer than 10~Gyr. \\item The loss rate of the dust mass, and the decay rate of fractional luminosity, primarily depend on the difference between the slope $q\\sbs{p}$ of the ``primordial'' size distribution of largest planetesimals and the slope $q\\sbs{g}$ of the size distribution of somewhat smaller, yet gravity-dominated, planetesimals that already underwent sufficient collisional evolution. With ``standard'' values of $q\\sbs{p}$ and $q\\sbs{g}$, the dust mass and the thermal fluxes follow approximately $t^\\xi$ with $\\xi = -0.3\\ldots -0.4$. \\item Specific decay laws of the total disk mass and the dust mass largely depend on a few model parameters. Most important are: the critical fragmentation energy $Q\\sbs{D}^*$ as a function of size, the slope of the ``primordial'' size distribution of planetesimals $q\\sbs{p}$ and their maximum size $s\\sbs{max}$, and the characteristic eccentricity $e$ and inclination $I$ of planetesimals. \\item The property that the maximum possible dust luminosity for a given age does not depend on the initial disk mass, established by \\citet{wyatt-et-al-2007a}, is only valid in cases of very rapid collisional evolution, i.e. in closer-in or dynamically very hot disks. For most of the systems at ages $<10$~Gyr, an increase of the initial disk mass leads to an increase of the dust luminosity, unless that initial mass is assigned extreme values, incompatible with our understanding of planetesimal disks. \\item Assuming standard material prescriptions and disk masses and extents, a synthetic population of disks generated with our analytic model generally agrees with the observed statistics of 24 and 70~\\textmu{m} fluxes versus age. Similarly, the synthetic [24]-[70] colors are consistent with the observed disk colors. \\end{enumerate}  As every model, our numerical model makes a number of general simplifying assumptions; the analytic one imposes further simplifications: \\begin{itemize} \\item The collisional evolution is assumed to be smooth and unperturbed. Singular episodes like the aftermath of giant break-ups or special periods of the dynamical evolution such as the late heavy bombardment are not included. \\item Effects of possible perturbing planets are taken into account only indirectly: through the eccentricities of planetesimals (dynamical excitation) and confinement of planetesimal belts (truncation of disks). Further effects such as resonant trapping or ejection of material by planets are neglected. \\item We only consider disruptive collisions. This is a reasonable approximation for disks that are sufficiently ``hot'' dynamically. However, cratering collisions become important when the relative velocities are insufficient for disruption to occur. \\item Neither dilute disks under the regime of Poynting-Robertson drag nor very dense disks with collisional timescales shorter than orbital timescales and with avalanches \\citep{grigorieva-et-al-2007} are covered by the present work. \\item Explaining the initial conditions or deriving them from the dynamical history of the systems at early stages of planetesimal and planetary accretion was out of the scope of this paper. Correlations between disk masses, disk radii, and the presence of planets, for example, were not considered, although they might alter the scalings we found here. \\end{itemize}  Despite these limitations, our models reproduce, in essential part, the observed evolution of dust in debris disks. We hope that they may serve as a starting point for in-depth studies that will certainly be undertaken in the future, motivated by questions that remain unanswered, as well as by new data expected from ongoing and planned observational programs."
    },
    "0710/0710.2367_arXiv.txt": {
        "abstract": "We present measurements of the neutron-capture elements Rb and Pb for bright giants in the globular clusters M4 and M5. The clusters are of similar metallicity ([Fe/H] $\\simeq -1.2)$ but M4 is decidedly $s$-process enriched relative to M5: [Ba/Fe] = +0.6 for M4 but 0.0 for M5. The Rb and Pb abundances were derived by comparing synthetic spectra with high-resolution, high signal-to-noise ratio spectra obtained with MIKE on the Magellan telescope. Abundances of Y, Zr, La, and Eu were also obtained. In M4, the mean abundances from 12 giants are [Rb/Fe] = 0.39 $\\pm$ 0.02 ($\\sigma$ = 0.07), [Rb/Zr] = 0.17 $\\pm$ 0.03 ($\\sigma$ = 0.08), and [Pb/Fe] = 0.30 $\\pm$ 0.02 ($\\sigma$ = 0.07). In M5, the mean abundances from two giants are [Rb/Fe] = 0.00 $\\pm$ 0.05 ($\\sigma$ = 0.06), [Rb/Zr] = 0.08 $\\pm$ 0.08 ($\\sigma$ = 0.11), and [Pb/Fe] = $-$0.35 $\\pm$ 0.02 ($\\sigma$ = 0.04). Within the measurement uncertainties, the abundance ratios [Rb/Fe], [Pb/Fe] and [Rb/X] for X = Y, Zr, La are constant from star-to-star in each cluster and none of these ratios are correlated with O or Na abundances. While M4 has a higher Rb abundance than M5, the ratios [Rb/X] are similar in both clusters indicating that the nature of the $s$-products are very similar for each cluster but the gas from which M4's stars formed had a higher concentration of these products. ",
        "introduction": "\\label{sec:intro}  Globular clusters continue to provide a source of fascination and frustration to both theorists and observers. Two notable accomplishments include the use of globular clusters to (a) check the age of the Universe (e.g., \\citealt{gratton03c}) and to (b) test and refine our understanding of stellar evolution (e.g., \\citealt{renzini88}). Despite these successes, globular clusters continue to present bewildering puzzles. The most persistent puzzle relates to chemical composition.  For many years, globular clusters have been known to exhibit star-to-star abundance variations for the light elements C, N, O, Na, Mg, and Al (e.g., see reviews by \\citealt{smith87}, \\citealt{kraft94}, and \\citealt{gratton04}). While the amplitude of the star-to-star abundance dispersion can vary from cluster to cluster, the now familiar anticorrelations between C and N, O and Na, and Mg and Al reveal that the abundance variations are likely produced during hydrogen burning at high temperatures via the CNO, Ne-Na, and Mg-Al cycles. (The O-Na and Mg-Al anticorrelations are not seen in field stars.) However, the stars responsible for the nucleosynthesis and the nature of the pollution mechanism(s) remain poorly understood (see \\citealt{lattanzio06} for a recent summary).  One possible explanation for the observed abundance anomalies is internal mixing and nucleosynthesis (e.g., \\citealt{sm79,charbonnel95}) within the present cluster members, the so-called evolutionary scenario. The systematic variation of the C and N \\citep{ss91} and Li \\citep{grundahl02} abundances with luminosity along the red giant branch demand an evolutionary component to the star-to-star abundance variations. Dredge-up of CN-cycled material accounts for the C and N variations. Development of a giant's convective envelope leading to mixing with highly Li-depleted gas accounts for the decline of the Li abundance with increasing luminosity. The proton-capture reactions causing the O, Na, Mg, and Al variations demand much higher temperatures and much deeper mixing than those required for CN-cycling. Such mixing is not predicted by standard theoretical models of red giants and the discovery of the O, Na, Mg, and Al anomalies in main sequence stars (e.g., \\citealt{briley96}, \\citealt{gratton01}) eliminates deep mixing as a viable explanation for the O-Al variations. The interiors of main sequence stars are too cool to process Ne to Na or Mg to Al. Therefore, the cluster gas must have been inhomogeneous when the present stars were formed. This alternative explanation for the abundance anomalies is the so-called primordial scenario.  In the primordial scenario, intermediate-mass asymptotic giant branch stars (IM-AGBs) from the generation to which the observed stars belong have long been considered candidates for synthesizing the abundance variations \\citep{cottrell81}. In IM-AGBs, the convective envelope can reach the top of the hydrogen-burning shell, a process called hot-bottom burning. For sufficiently massive and metal-poor AGBs, the temperatures at the base of the convective envelope can exceed 100 million degrees thereby allowing the efficient operation of the CNO, Ne-Na, and Mg-Al cycles (e.g., \\citealt{karakas03}). That IM-AGBs do not alter the abundances of the alpha or iron-peak elements (as required by observations) adds to their qualitative appeal. However, quantitative tests reveal problems with the IM-AGB primordial scenario. Theoretical yields from IM-AGBs combined with a chemical evolution model \\citep{fenner04} suggest that O is not sufficiently depleted, Na is overproduced, Mg is produced rather than destroyed, the isotope ratios of Mg do not match the observations, and the sum of C+N+O increases substantially in contrast to the observations. \\citet{ventura05a,ventura05b,ventura05c} find that many of the flaws noted above can be alleviated when IM-AGB yields are calculated using a revised treatment for convection and mass-loss. However, \\citeauthor{ventura05a} note that problems persist, namely with the Mg isotope ratios, and warn that the predictive power of the current AGB models is limited. Recently, \\citet{prantzos06}, \\citet{smith06}, and \\citet{decressin06} suggest that the winds from massive stars may be more promising candidates than IM-AGBs. There is no satisfactory explanation for the complex patterns for the light element abundances exhibited by every well studied Galactic globular cluster. Therefore, our present understanding of globular cluster chemical evolution and/or stellar nucleosynthesis is incomplete.  Determinations of the stellar abundances of the trans-iron or heavy elements offer clues to the history behind the chemical evolution of globular clusters. Here, we provide novel information -- the Rb and Pb abundances -- for giants in M4 and M5, a pair of clusters of similar metallicity but with distinctly different levels of $s$-process products. The quintessential $r$-process element Eu has similar abundances in the two clusters and, indeed, across the collection of Galactic globular clusters. In sharp contrast, the $s$-process products are more evident in M4 than in M5 and other clusters of similar metallicity: [Ba/Fe] is about +0.6 in M4 but 0.0 in M5. The questions - Are there differences in the Rb and Pb abundances between this pair of clusters? and Are the star-to-star variations in the abundances of light elements (O, Na, Mg, and Al) reflected in variations among the abundances of Rb and Pb? -- seem likely to probe the origins of the $s$- and $r$-process products for globular clusters.  Due to a critical branching point in the $s$-process path at $^{85}$Kr, the abundance of Rb relative to Sr, Y, or Zr can differ by a factor of 10 depending upon the neutron density at the $s$-process site. In the case of AGB stars, the neutron density in the He-shell is dependent on the stellar mass (e.g., see \\citealt{tomkin83}, \\citealt{lambert95}, \\citealt{busso99}, and \\citealt{abia01} for further details). Since the isotopes of Pb and Bi are the last stable nuclei on the $s$-process path, the $s$-process terminates at these elements and overabundances of Pb and Bi will arise if seed nuclei are shuffled by neutron captures down the entire $s$-process path. In particular, metal-poor AGB stars may produce large overabundances of Pb and Bi if the neutron supply per seed exceeds a critical value (e.g., see \\citealt{goriely01}, \\citealt{travaglio01}, and \\citealt{busso01} for further details). The suspicion is that the star-to-star abundance variations for light elements are due to contamination by IM-AGBs. Some contend that IM-AGBs also synthesize $s$-process nuclides and then one might expect to see star-to-star variations in the Rb and Pb abundances as well as correlations with light element abundances.  To further examine the possible role of IM-AGBs in the chemical evolution of globular clusters, \\citet{rbpbsubaru} measured Rb and Pb in NGC 6752 and M13, the two clusters that exhibit the largest amplitude for Al variations. It was found that the abundance ratios [Rb/Zr] and [Pb/Fe] were constant from star-to-star within the measurement uncertainties. If IM-AGBs do synthesize Rb and Pb, then they may not be responsible for the abundance variations. On the other hand, if IM-AGBs are responsible for the abundance variations, they cannot synthesize Rb or Pb.  In this paper, we extend the measurements of Rb and Pb to the globular clusters M4 and M5. While these clusters are more metal-rich than NGC 6752 or M13, both M4 and M5 are known to exhibit large dispersions and correlations for the light element abundances [see pioneering studies on CN bimodality by \\citet{norris81a} and \\citet{smith83} as well as recent high-resolution spectroscopic studies by \\citet{M4,M5}, \\citet{ramirez03}, and references therein]. In particular, as noted above, M4 is remarkably, perhaps uniquely among globular clusters, enriched in $s$-process products. ",
        "conclusions": "\\label{sec:summary}  In this paper we present measurements of the neutron-capture elements Rb and Pb in the globular clusters M4 and M5. While both clusters exhibit star-to-star abundance variations for the light elements, we find that the abundances of Rb and Pb are constant. None of the abundance ratios [Rb/Fe], [Rb/Zr], and [Pb/Fe] are correlated with O or Na abundances. In the primordial scenario, the abundance variations for the light elements are attributed to different levels of accretion of ejecta from IM-AGBs or massive stars. The fact that the heavy elements including Rb and Pb do not show abundance variations implies that the accreted material has the same composition as the ambient material for the heavy elements (i.e., the accreted material cannot be highly underabundant or overabundant in these elements). That the ratios [Rb/X] for X = Y, Zr, La are similar for M4 and M5 suggests that the source of the $s$-process elements are similar and that M4 had a greater concentration of these products.  There remains a need to pursue additional observational tests of the primordial scenario. In particular, present data on the Rb and Pb abundances in field and cluster stars are sparse. The indication that the Mg isotopic ratios of unpolluted or normal cluster stars differ from those of field stars of the same metallicity deserves closer scrutiny by, in particular, extending the measurement of these isotopic ratios to additional clusters."
    },
    "0710/0710.5637_arXiv.txt": {
        "abstract": "{ A new method for the determination of open cluster membership based on a cumulative effect is proposed. In the field of a plate the relative $x$ and $y$ coordinate positions of each star with respect to all the other stars are added. The procedure is carried out for two epochs $t_1$ and $t_2$ separately, then one sum is subtracted from another. For a field star the differences in its relative coordinate positions of two epochs will be accumulated. For a cluster star, on the contrary, the changes in relative positions of cluster members at $t_1$ and $t_2$ will be very small. On the histogram of sums the cluster stars will gather to the left of the diagram, while the field stars will form a tail to the right. The procedure allows us to efficiently discriminate one group from another. The greater the distance between $t_1$ and $t_2$ and the more cluster stars present, the greater is the effect. The accumulation method does not require reference stars, determination of centroids and modelling the distribution of field stars, necessary in traditional methods. By the proposed method 240 open clusters have been processed, including stars up to $m<13$. The membership probabilities have been calculated and compared to those obtained by the most commonly used Vasilevskis-Sanders method. The similarity of the results acquired the two different approaches is satisfactory for the majority of clusters. ",
        "introduction": "Various methods based on the analysis of positions, proper motions, radial velocities, magnitudes and their combinations have been proposed to determine the members of open clusters.  The first mathematically rigorous procedure for determination of open cluster membership was developed by Sanders (\\cite{Sanders}) with a statistical analysis of proper motions. It is also the most widely used method. Sanders's approach is based on the model of overlapping distributions of field and cluster stars in the neighborhood and within the region of visible grouping of stars, introduced by Vasilevskis et al. (\\cite{Vasilevskis}).  Vasilevskis's model implies that proper motion dispersion of cluster members is caused by observational and measurement errors assumed to be normally distributed. Thus the distribution of cluster stars is represented by a bivariate normal frequency function. The dispersion of field stars is due to not only the errors referred to above, but also to peculiar motion and differential galactic rotation. Therefore the field star distribution is not expected to be random, but rather to have a preferential direction and not normal distribution. However, in a first approximation a bivariate normal ellipsoidal distribution function was assumed for field stars, the major axis of the ellipse being parallel to the galactic plane.  Thus, Sanders's equations contain 8 unknown parameters to be determined: the number of cluster members, x and y components of two centroids, one circular and two elliptical dispersions. This system is solved by a maximum likelihood method (Sanders \\cite{Sanders}). The probability of stars being cluster members is calculated by frequency functions with determined parameters.  Slovak (\\cite{Slovak}) tested the Vasilevskis-Sanders method by modelling the proper motion distribution in the surroundings of a cluster. He proved the uniqueness and convergence of solutions of the system, provided that errors are represented by a Gaussian distribution and open clusters have no significant internal motion. But if there is noticeable motion within a cluster or it rotates then the above method fails.  Cabrera-Ca\\~{n}o and Alfaro (\\cite{Cabrera-Cano Alfaro1}) improved the numerical techniques for obtaining the above parameters. McNamara and Schneeberger (\\cite{McNamara Schneeberger}) showed that the final probabilities could be influenced by various weight groups. Zhao and He (\\cite{Zhao He}) provided a method for treating data with different accuracies.  However all these improvements did not treat the problem of star distribution model on which the membership probabilities are based. Even if the hypothesis of a normal distribution of field stars is realistic for some clusters, the centroids of field and cluster stars sometimes are too close to be well discriminated. The parametric model also fails when the cluster member-to-field star ratio is small. The Vasilevskis-Sanders method does not work in the case of significant internal motion in a cluster or its rotation.  To overcome some of the problems arising from the parametric Vasilevskis-Sanders method, especially the star distribution modelling, Cabrera-Ca\\~{n}o and Alfaro (\\cite{Cabrera-Cano Alfaro2}) developed a more general, non-parametric method of membership determination. Here no assumptions were made about the nature of cluster and field star distributions and allowed for the use of photometric data. Each cluster needed careful individual study.  So, the actual distribution of cluster and field stars, especially the latter, may not be fitted by a Vaselevskis-Sanders model; the centroids for the two groups may be too close to be distinguished; the reliability of the method depends on the cluster-to-field star ratio; in the case of significant internal motions or rotation of a cluster the traditional method fails.  In the next section we introduce a new method for the discrimination of cluster stars from surrounding field stars based on a cumulative effect using positions and proper motions. We do not try to overcome all the difficulties of traditional methods, though we offer another approach that enlarges the statistical distance between the two populations - cluster members and field stars - by revealing the group of stars with the least relative velocities. The advantage of this method is that no assumption is made about the distribution of field stars and determination of centroids is avoided. However, in order to determine probabilities we still have to assume a normal bivariate distribution for clusters. Another advantage of the method is that reference stars are not necessary: the discrimination of cluster members is most effective by rectangular coordinates. These features of the method allow us to increase the statistical distance between the two populations. The most noticeable advantage of the accumulation method is its ability to reveal dynamic structures within the clusters if there are any. ",
        "conclusions": "In the presented method of accumulation we assume that the cluster members moving in the Galaxy as a whole have similar velocities, while field stars show a wide range of velocities. The procedure - adding close velocities while compensating for disperced ones - should enhance the assembling of cluster stars, while distributing more sparsely the field ones, due to the cumulative effect. Thus this method effectively enlarges the statistical distance between physical members and field stars, so that a membership as well as a non-membership is more pronounced. This effect is better for larger cluster-to-field star ratios. For field stars no particular distribution is assumed, thus the centroids are not determined. However, when calculating probabilities, for physical members a normal bivariate distribution function is assumed. For rectangular coordinates no reference stars are needed in the procedure. The most interesting feature of the accumulation mathod is its ability to reveal more then one group of velocities, which has been shown with the example of Pleades.  The modified accumulation method was applied to 240 clusters from Dias's list. The probabilities calculated by the accumulation method showed satisfactory agreement with those obtained by the Vasilevskis-Sanders method for the majority of clusters. The poor agreements or disagreements can be ascribed to low cluster-to-field star ratios, or multiple dynamic structures.  The accumulation method can be further expanded using other variables like radial velocities and photometric quantities."
    },
    "0710/0710.0224_arXiv.txt": {
        "abstract": "{ We discuss the large scale properties of standard cold dark matter cosmological models characterizing the main features of the power-spectrum, of the two-point correlation function and of the mass variance.  Both the real-space statistics have a very well defined behavior on large enough scales, where their amplitudes become smaller than unity.  The correlation function, in the range $0<\\xi(r)<1$, is characterized by a typical length-scale $r_c$, at which $\\xi(r_c)=0$, which is fixed by the physics of the early universe: beyond this scale it becomes negative, going to zero with a tail proportional to $-(r^{-4})$.  These anti-correlations represent thus an important observational challenge to verify models in real space.  The same length scale $r_c$ characterizes the behavior of the mass variance which decays, for $r>r_c$, as $r^{-4}$, the fastest decay for any mass distribution. The length-scale $r_c$ defines the maximum extension of (positively correlated) structures in these models.  These are the features expected for the dark matter field: galaxies, which represent a biased field, however may have differences with respect to these behaviors, which we analyze.  We then discuss the detectability of these real space features by considering several estimators of the two-point correlation function. By making tests on numerical simulations we emphasize the important role of finite size effects which should always be controlled for careful measurements. ",
        "introduction": "In contemporary cosmological models the structures observed today at large scales in the distribution of galaxies in the universe are explained by the dynamical evolution of purely self-gravitating matter (dark matter) from an initial state with low amplitude density fluctuations, the latter strongly constrained by satellite observations of the fluctuations in the temperature of the cosmic microwave background radiation.  The other main observational elements for the understanding of the large scale structure of the universe is represented by the studies of galaxy correlations. Any theoretical model aiming to explain the formation of structures must be tested against the data provided by galaxy surveys which give the important bridge between the regimes characterized by large and small fluctuations.  Models of the early universe (see e.g.  Padmanabhan, 1993 and references therein) predict certain primordial fluctuations in the matter density field, defining the correlations of the initial conditions, i.e. at the time of decoupling between matter and radiation.  In the regime where density fluctuations are small enough, the correlation function of the present matter density field is simply related to one describing the initial conditions. In fact, according to the growth of gravitational instabilities in an expanding universe in the linear regime perturbations are simply amplified (see e.g., Peebles, 1980 and references therein).  Thus today at some large scales where the correlation function is still positive but with $\\xi(r)<1$ the imprint of primordial fluctuations should be preserved. In the region of strong non-linear fluctuations an analytical treatment to predict the behavior of the two-point correlation function has not been developed yet and, in general, one makes use of numerical simulations which provide a rich, but phenomenological, description of structure in the non-linear regime. It is in this regime, at small enough scales, where most observations have been performed until now.     We focus here on the type of correlations predicted in the linear regime by models of the early universe. While the characterization of correlations is usually done in terms of the power-spectrum of the density fluctuations a real space analysis turns out to be useful to point out some relevant features from an observational point of view (see, e.g., the discussion in Gabrielli et al., 2004).   Theoretical models of primordial matter density fields in the expanding universe are characterized by a single well-defined length scale, which is an imprint of the physics of the early universe at the time of the decoupling between matter and radiation (see e.g. Bond and Efstathiou 1984, and Padmanabhan 1993 for a general introduction to the problem). The redshift characterizing the decoupling is directly related to the scale at which the change of slope of the power-spectrum of matter density fluctuations $P(k)$ occurs, i.e. it defines the wavenumber $k_c$ at which there is the turnover of the power-spectrum between a regime, at large enough $k$, where it behaves as a negative power-law of the wave number $P(k) \\sim k^{m}$ with $-1<m\\le-3$, and a regime at small $k$ where $P(k)\\sim k$ as predicted by inflationary theories. Given the generality of this prediction, it is clearly extremely important to look for this scale in the data.    The exact location of this scale is related to several parameters, including the cosmological ones which describe the geometry of the universe at large scales (see e.g. Padmanabhan 1993, Tegmark et al. 2004 and Spergel et al. 2007 for a recent determination). We discuss in what follows that the scale $r_c$ corresponding to the wave-number $k_c$, in a particular variant of Cold Dark Matter (CDM) models --- the so-called $\\Lambda$CDM vanilla model --- is predicted to be $r_c \\approx 124$ Mpc/h\\footnote{For seek of clarity we have chosen the scale of distances normalized to the adimensional Hubble parameter $h$, which is defined from the Hubble's constant $H_0 = 100 h$ km/sec/Mpc.}. At this scale the real space correlation function crosses zero, becoming negative at larger scales. In particular the correlation function presents a positive power-law behavior at scales $r \\ll r_c$ and a negative power-law behavior at scales $r \\gg r_c$. Positive and negative correlations are exactly balanced in way such that the integral over the whole space of the correlation function is equal to zero. This is a global condition on the system fluctuations which corresponds to the fact that the distribution is super-homogeneous (or hyper-uniform), i.e. characterized by a sort of stochastic order and by fluctuations which are depressed with respect to, for example, a purely uncorrelated distribution of matter (Gabrielli, Joyce and Sylos Labini, 2002 --- see discussion below).  Note that the scale $r_c$ marks the maximum extension of positively correlated structures: beyond $r_c$ the distribution must be anti-correlated since the beginning, as the evolution time was not sufficient for the positive correlations to be developed. Thus this scale can be regarded as an upper limit to the maximum size of structures (with large of weak correlations) in the present universe. The possible discoveries of structures of larger size is still a challenging task for observational cosmology.   A relevant problem for the measurements of small amplitude values of the correlation function, i.e. when $\\xi(r) <1$, is represented by the characterization and the understanding of both the systematic biases which may affect the estimators of $\\xi(r)$ and the stochastic noise which perturbs any real determination. A study of this problems can be found, for example, in Kerscher (1999) and Kerscher et al. (2000) where it is shown that in general the biases in several estimators of the two-point correlation function are not negligible. In particular when there are structures of large spatial extension inside a given sample there can be non negligible biases affecting the determination of two-point properties.  We focus here on the systematic bias related to the effect of the so-called integral constraint, which distorts any estimator of the correlation function at large scales in any given sample. The integral-constraint represents an overall condition on any estimator of the correlation function which is due to the fact that the average density, estimated in any given sample, is in general different from its ensemble average value.  Here we treat explicitly the case for the simplest estimator of the two-point correlation function, {the so-called full-shell or minus estimator and and we illustrate the situation for the other estimatorsby studying artificial distributions. In particular we devote most attention to the estimator introduced by Davis and Peebles (1983), which is still very used in the literature, and to the estimator introduced by Landy and Szalay (1993), which is the most popular one. Kerscher et al. (2000) considered also other estimators, like the Hewett estimator (Hewett, 1982) and the Hamilton estimator (Hamilton, 1993) and have shown that the results obtained with he Landy and Szalay estimator are almost indistinguishable from the Hamilton estimator.  In this way we will be able to identify the problems related to the identification of correlations above the mentioned scale $r_c$: we will then propose several tests to be applied to the galaxy data, in order to define the strategy to study the correlation function at small amplitudes and larger distances in order to eventually detect the length scale $r_c$.   Up to now studies of the correlation function $\\xi(r)$ in galaxy samples have been limited to small scales, i.e. $0.1 < r \\ltapprox 30$ Mpc/h (i.e. Totsuji \\& Kihara, 1969, Davis and Peebles, 1983, Davis et al., 1988, Benoist et al., 1996 Park et al., 1994, Scranton et al., 2002}, Zehavi et al., 2002, Zehavi et al., 2004,  Ross et al., 2007) and only recently the volume covered by galaxy redshift samples is approaching a size which is large enough to make a robust estimation of the correlation function at scales of order 100 Mpc/h. When the Sloan Digital Sky Survey (SDSS) (York et al., 2000) will be completed by filling up the gap between the two main angular regions of observations, which are nowadays disjointed, the volume of the survey and the statistics of the number of objects in the samples would be large enough to test space correlations on scales of order $r_c$ or more. An exception to this situation is represented by the paper by Eisenstein, et al., (2005), who, by studying a sample of Luminous Red Galaxies (LRG) of the SDSS, have estimated the correlation function on scales of order 100 Mpc/h. These authors have however focused their attention to another real space feature of theoretical models: the so-called ``bump'' of the correlation function which corresponds in real space to the so-called Doppler peaks in the matter power-spectrum generated by the baryonic acoustic oscillations in the early universe. As we discuss below this bump, corresponding to a singular point of the correlation function (Gabrielli et al., 2004), is localized at scales of order of 100 Mpc/h and characterized by a small amplitude. This is a second important real-space scale of the theoretical correlation function which is localized at a scale slightly smaller than $r_c$. The detection of the baryonic bump is thus related to the detection of the scale $r_c$ as any finite-size effect perturbing the determination of the scale $r_c$ will, inevitably, also affect the determination of the baryonic bump.  In fact the baryonic bump can be seen as a small modification to the overall shape of the correlation function at scales of order $r_c$, to which we focus our attention here.     Note that, because of the very large scales, the acoustic signature and the zero point scale remain in the linear regime even today and they are weakly affected by non-linear effects (see Eisenstein et al., 2006). Thus real space and redshift space properties, at such large scales, should not differ substantially.   In Section 2 we introduce the basic definitions of the statistical quantities usually employed to characterize two-point properties in real and Fourier space. In Section 3 we discuss a simple functional behavior of the power-spectrum of matter density fluctuations which captures the main elements of a more realistic CDM power-spectrum. We discuss the real-space properties as represented by the two-point correlation function and we consider the problem of selection or biasing in the simplest theoretical scheme of biasing a correlated Gaussian field. In Section 4 we treat explicitly the case of a $\\Lambda$CDM matter density field characterizing in detail real space properties.  The main estimators of the two-point correlation function are discussed in Section 5 and in Section 6 we test these estimators in artificial distributions.  Finally in Section 7 we draw our main conclusions discussing the problems related to the estimations of two-point correlations in real galaxy samples. ",
        "conclusions": "We have considered the real-space properties of CDM density fields, focusing in detail in a particular variant known as $\\Lambda$CDM (vanilla) model. It is well known that the power-spectrum has typically a behavior $P(k) \\sim k^{m}$ with $-1<m\\le-3$ for large wavelengths $k > k_c$, and $P(k)\\sim k$ at smaller wavelengths $k <k_c$. We discussed that, correspondingly, the two-point correlation function shows approximatively a positive power-law behavior $\\xi(r) \\sim r^{-2}$ at small scales $r < r_c \\approx k_c^{-1}$ and a negative power-law behavior $\\xi(r) \\propto - r^{-4}$ at large scales $r>r_c$, where the zero-crossing occurs at about $r_c \\approx 124$ Mpc/h in the model considered. We discussed the fact that, globally, a system with this type of correlations belong to the category of super-homogeneous distributions, which are configurations of points more ordered than a purely uncorrelated (Poisson) distribution.  Correspondingly fluctuations are depressed with respect to the Poisson case, and the normalized mass variance, for instance, decay faster ($\\sigma^2(r) \\sim r^{-4}$) than for the Poisson case ($\\sigma^2(r) \\sim r^{-3}$). The condition of super-homogeneity is expressed by the condition that $P(k) \\rightarrow 0$ for $k \\rightarrow 0$, or alternatively that \\[ \\int_0^{\\infty}  \\xi(r) r^2 dr = 0 \\;. \\]  Following the work of Durrer et al. (2003) we have pointed out that the above condition is broken when one samples the distribution, as for example when the simplest biasing scheme of correlated Gaussian fields (introduced by Kaiser, 1984) is applied. This is particularly important for the behavior of the power-spectrum for $k<k_c$, which, under biasing, remains constant instead of going as $P(k) \\sim k$. The correlation function at large scales $r>r_c$ is instead expected to be linearly amplified with respect to the original one of the whole matter field. Thus the large scale negative tail $\\xi(r) \\sim -r ^{-4}$ is the main feature which one would like to detect in order to test theoretical models.  Given the fact that when $\\xi(r)$ becomes negative, it is characterized by a very small amplitude, the determination of the negative power-law tail represents a very challenging problem. We have discussed the fact that, at first approximation in a real measurement, one may treat the system as having positive correlations at small scales with an exponential cut-off at the scale $r_c$ and then it becomes uncorrelated (a situation which can be regarded as upper limit to the presence of anti-correlations). This implies that for $r_c > 124$ Mpc/h galaxy distribution should not present any positive correlation. Whether this behavior is compatible with the existences of structures of order 200 Mpc/h or more is an open problem which has to be addressed in the studies of forthcoming galaxy catalogs.  More in detail, one of the most basic results (see e.g., Peebles 1980) about self-gravitating systems, treated using perturbative approaches to the problem (i.e. the fluid limit), is that the amplitude of small fluctuations grows monotonically in time, in a way which is independent of the scale.  This linearized treatment breaks down at any given scale when the relative fluctuation at the same scale becomes of order unity, signaling the onset of the ``non-linear'' phase of gravitational collapse of the mass in regions of the corresponding size.  If the initial velocity dispersion of particles is small, non-linear structures start to develop at small scales first and then the evolution becomes ``hierarchical'', i.e., structures build up at successively larger scales. Given the finite time from the initial conditions to the present day, the development of non-linear structures is limited in space, i.e., they can not be more extended than the scale at which the linear approach predicts that the density contrast becomes of order unity at the present time. This scale is fixed by the initial amplitude of fluctuations, constrained by the cosmic microwave background anisotropies (Spergel et al., 2007), by the hypothesized nature of the dominating dark matter component and its correlation properties.  According to current models of CDM-type the scales at which non-linear clustering occurs at the present time (of order 10 Mpc) are much smaller than the scale $r_c \\approx 124$ Mpc/h (see e.g. Springel et al., 2005). Thus the region where the super-homogeneous features should still be in the linear regime, allowing a direct test of the initial conditions predicted by early universe models. The scale $r_c$ marks the maximum extension of positively correlated structures: beyond $r_c$ the distribution must be anti-correlated since the beginning, as there was no time to develop other correlations.  The possible presence of structures, which mark long-range correlations, whether or not of large amplitude, reported both by observations of galaxy distributions (like the Sloan Great Wall --- see Gott et al., 2005), by the detection of dark matter distributions (see e.g. Massey et al., 2007) and by the large void of radius $\\sim 140$ Mpc identified by Rudnick et al. (2007), is maybe indicating that positive correlations extend well beyond $r_c$.    We have discussed that an important finite size effect must be considered when estimating the correlation function, and which may mimic a break of the power-law behavior similar to the ones of CDM models at a scale of order $r_c$. This is related to the effect of the integral constraint in the estimators, namely the fact that the sample average, estimated in a finite sample, differs from the ensemble average, and can be finite-size dependent. This situation occurs when correlations (weak or strong) extend to scales larger than the sample size.    For these reasons, in order to study the two-point correlation function in real galaxy samples when its amplitude becomes smaller than unity, it is crucial to check whether the break of the power-law behavior has a finite size dependence or not, by choosing samples with different depth. In this perspective the assessment of the reality of the break of the two-point correlation function is the main observational point to be considered.  Once this will be clarified other features should be considered, as for the example the so-called baryonic bump, which is a very small perturbation to the overall shape of the correlation function at scales of order of the zero-point $r_c$. We will present a detailed analysis of the correlation properties of galaxy distribution in the SDSS catalog, considering specific tests for finite-size effects in the determination of the correlation function, in a forthcoming paper."
    },
    "0710/0710.2017_arXiv.txt": {
        "abstract": "{ H.E.S.S. observations of the old-age ($>$10$^4$~yr; $\\sim 0.5^\\circ$ diameter) composite supernova remnant (SNR) W~28 reveal very high energy (VHE) $\\gamma$-ray emission situated at its northeastern and southern boundaries. The northeastern VHE source (HESS~J1801$-$233) is in an area where W~28 is interacting with a dense molecular cloud, containing OH masers, local radio and X-ray peaks. The southern VHE sources (HESS~J1800$-$240 with components labelled A, B and C) are found in a region occupied by several HII regions, including the ultracompact HII region W~28A2. Our analysis of NANTEN CO data reveals a dense molecular cloud enveloping this southern region, and our reanalysis of EGRET data reveals MeV/GeV emission centred on HESS~J1801$-$233 and the northeastern interaction region. }  \\email{growell@physics.adelaide.edu.au}   \\begin{document} ",
        "introduction": "The study of shell-type SNRs at $\\gamma$-ray energies is motivated by the idea that they are the dominant sites of hadronic Galactic cosmic-ray (CR) acceleration to energies approaching the \\emph{knee} ($\\sim 10^{15}$~eV) and beyond, e.g. \\cite{Ginzburg:1}. CRs are then accelerated via the diffusive shock acceleration (DSA) process (eg. \\cite{Bell:1,Blandford:2}). Gamma-ray production from the interaction of these CRs with ambient matter and/or electromagnetic fields is a tracer of such particle acceleration, and establishing the hadronic or electronic nature of the parent CRs in any $\\gamma$-ray source is a key issue. Already, two shell-type SNRs, RX~J1713.7$-$3946 and RX~J0852.0$-$4622, exhibit shell-like morphology in VHE $\\gamma$-rays \\cite{HESS_RXJ1713_II,HESS_VelaJnr_II,HESS_RXJ1713_III} to 20~TeV and above. Although a hadronic origin of the VHE $\\gamma$-ray emission is highly likely in the above cases, an electronic origin is not ruled out.  W~28 (G6.4$-$0.1) is a composite or mixed-morphology SNR, with dimensions 50$^\\prime$x45$^\\prime$ and an estimated distance between 1.8 and 3.3~kpc (eg. \\cite{Goudis:1,Lozinskaya:1}). It is an old-age SNR (age 3.5$\\times 10^4$ to 15$\\times 10^4$~yr \\cite{Kaspi:1}), thought to have entered its radiative phase of evolution \\cite{Lozinskaya:1}. The shell-like radio emission \\cite{Long:1,Dubner:1} peaks at the northern and northeastern boundaries where interaction with a molecular cloud \\cite{Wootten:1} is established \\cite{Reach:1,Arikawa:1}. The X-ray emission, which overall is well-explained by a thermal model, peaks in the SNR centre but has local enhancements in the northeastern SNR/molecular cloud interaction region \\cite{Rho:2}. Additional SNRs in the vicinity of W~28 have also been identified: G6.67$-$0.42 and G7.06$-$0.12 \\cite{Yusef:1}. The pulsar PSR~J1801$-$23 with spin-down luminosity $\\dot{E} \\sim 6.2\\times 10^{34}$ erg~s$^{-1}$ and distance $d\\geq9.4$~kpc \\cite{Claussen:3}, is at the northern radio edge.  Given its interaction with a molecular cloud, W~28 is an ideal target for VHE observations. This interaction is traced by the high concentration of 1720~MHz OH masers \\cite{Frail:2,Claussen:1,Claussen:2}, and also the location of very high-density ($n>10^3$~cm$^{-3}$) shocked gas \\cite{Arikawa:1,Reach:1}. Previous observations of the W~28 region at VHE energies by the CANGAROO-I telescope revealed no evidence for such emission \\cite{Rowell:1} from this and nearby regions.  The High Energy Stereoscopic System (H.E.S.S.:  see \\cite{Hinton:1} for details and performance) has observed the W~28 region over the 2004, 2005 and 2006 seasons. After quality selection, a total of $\\sim$42~hr observations were available for analysis. Data were analysed using the moment-based Hillas analysis procedure employing {\\em hard cuts} (image size $>$200~p.e.), the same used in the analysis of the  inner Galactic Plane Scan datasets \\cite{HESS_GalScan,HESS_GalScan_II}. An energy threshold of $\\sim 320$~GeV results from this analysis. The VHE $\\gamma$-ray image in Fig.~\\ref{fig:tevskymap} shows that two source of VHE $\\gamma$-ray emission are located at the northeastern and southern boundaries of W~28. The VHE sources are labelled HESS~J1801$-$233 and HESS~J1801$-$240 where the latter can be further subdivided into three components A, B, and C. The excess significances of both sources exceed $\\sim$8$\\sigma$ after integrating events within their fitted, arcminute-scale sizes. Similar results were also obtained using an alternative analysis \\cite{Mathieu:1}. \\begin{figure*}[th] \\centering \\hbox{ \\begin{minipage}{0.55\\textwidth} \\includegraphics[width=\\textwidth]{icrc0129_fig1.eps} \\end{minipage} \\begin{minipage}{0.45\\textwidth} \\caption{H.E.S.S. VHE $\\gamma$-ray excess counts, corrected for exposure and Gaussian smoothed (with 4.2$^\\prime$ std. dev.). Solid green contours represent excess significance levels of 4, 5, and 6$\\sigma$, for an integrating radius $\\theta$=0.1$^\\circ$. The VHE sources HESS~J1801$-$233 and a complex of sources HESS~J1800-240 (A, B \\& C) are indicated. The thin-dashed circle depicts the approximate radio boundary of the SNR W~28 \\cite{Dubner:1,Brogan:1}. Additional objects indicated are: HII regions (black stars); W~28A2, {G6.1$-$0.6} % {6.225$-$0.569}; % The 68\\% and 95\\% location contours (thick-dashed yellow lines) of the $E>100$~MeV EGRET source {GRO~J1801$-$2320}; the pulsar {PSR~J1801$-$23} (white triangle). The inset depicts a pointlike source under identical analysis and smoothing as for the main image.} \\label{fig:tevskymap} \\end{minipage} } \\end{figure*} ",
        "conclusions": "H.E.S.S. and NANTEN observations reveal VHE emission in the W~28 region spatially coincident with molecular clouds. The VHE emission and molecular clouds are found at the northeastern boundary, and $\\sim 0.5^\\circ$ south of W~28 respectively. The SNR W~28 may be a source of power for the VHE sources, although there are additional potential particle accelerators in the region such as other SNR/SNR-candidates, HII regions and open clusters. Further details concerning these results and discussion are presented in \\cite{HESS_W28}."
    },
    "0710/0710.5727_arXiv.txt": {
        "abstract": "{}{We communicate the detection of soft (20--200\\,keV) $\\gamma$-rays from the pulsar and pulsar wind nebula of \\hbox{PSR~J1846$-$0258} and aim to identify the component of the system which is responsible for the $\\gamma$-ray emission.}{We combine spectral information from the \\integral $\\gamma$-ray mission with archival data from the \\chandra X-ray Observatory to pinpoint the source of soft $\\gamma$-ray emission.}{Our analysis shows that the soft $\\gamma$-rays detected from \\hbox{PSR~J1846$-$0258} include emission from both the pulsar and the pulsar wind nebula, but the measured spectral shape is dominated by the pulsar wind nebula.  We discuss \\hbox{PSR~J1846$-$0258} in the context of rotation powered pulsars with high magnetic field strengths and review the anomalously high fraction of spin-down luminosity converted into X- and $\\gamma$-ray emission in light of a possible overestimate of the distance to this pulsar.}{} ",
        "introduction": "The pulsar \\hbox{PSR~J1846$-$0258} (also known as \\hbox{AX~J1846.4$-$0258}) was discovered by \\citet{GotthelfVasishtBoylan-Kolchin2000} in the X-ray band and lies near the centre of the supernova remnant Kes 75 (SNR G29.7-0.3).  Recent X-ray imaging \\citep{HelfandCollinsGotthelf2003} shows that the pulsar is embedded in a pulsar wind nebula (PWN) which shows distinct physical structure.  No radio emission has been observed from \\hbox{PSR~J1846$-$0258}, but X-ray timing reveals the pulse period to be $P=324$\\,ms and the characteristic age is $P/(2\\dot{P})=728$\\,years -- the smallest characteristic age of any rotation-powered pulsar.  The inferred surface dipole magnetic field strength is $4.8\\times10^{13}$\\,G -- almost an order of magnitude greater than the magnetic fields for typical pulsars, placing it closer to the field strengths of magnetars.  The distance to the system is $\\sim19$\\,kpc as estimated from the neutral hydrogen density along the line of sight \\citep{BeckerHelfand1984}, which implies that the size of the supernova remnant (SNR) is extremely large for such a young pulsar.  It also implies that the efficiency of the conversion of the pulsar spin-down luminosity ($8.4\\times10^{36}$\\,erg\\,s$^{-1}$, using $P$ and $\\dot P$ from \\citealt{GotthelfVasishtBoylan-Kolchin2000}) to combined pulsar and PWN X-ray luminosity in the 0.5--10\\,keV energy range, is 20\\% (using the flux from this work) -- the largest of any rotation-powered pulsar.  This source is one of a growing class of rotation-powered pulsars with B-field strengths approaching those of magnetars.  In this paper we present the \\integral observations and results for \\hbox{PSR~J1846$-$0258}, link these results to previously published \\chandra results and follow up with a discussion of the properties of this source in comparison to other members of this class and examine the possibility of an overestimate in the distance to this pulsar.  \\begin{figure*} \\sidecaption \\includegraphics[width=12cm]{8432fig1.eps} \\caption{IBIS/ISGRI significance mosaic of the 20--100\\,keV energy band, showing \\hbox{PSR~J1846$-$0258} in the centre of the field.  False colour representation of the significance is displayed on a logarithmic scale.  The top right inset is a $50\\arcsec\\times50\\arcsec$ image from our reduction of \\chandra data in the 0.3--10\\,keV energy band.  One can clearly distinguish the bright pulsar and the surrounding synchrotron nebulosity.  The white ellipses on the \\chandra image indicate the extraction regions for the PWN and pulsar spectra as explained in Sect.~\\ref{chanresults}.} \\label{FigIntimg} \\end{figure*} ",
        "conclusions": "\\label{discuss}   PSR~J1846$-$0258 is one of a recently emerged and growing class of rotation-powered pulsars with inferred surface dipole magnetic fields approaching those of the magnetars.  A number of these pulsars have been discovered as a result of the Parkes multi-beam pulsar survey of the Galactic plane \\citep{ManchesterLyneCamilo2001}, for example \\hbox{PSR~J1119$-$6127} and \\hbox{PSR~J1814$-$1744} \\citep{CamiloKaspiLyne2000}, which have inferred magnetic fields of $4.1\\times10^{13}$\\,G and $5.5\\times10^{13}$\\,G respectively.  These are all young pulsars with characteristic ages of order a few thousand years and approximately half of them show X-ray emission (\\hbox{PSR~J1119$-$6127}, \\citealt{GonzalezSafi-Harb2003}; \\hbox{PSR~J1718$-$3718}, \\citealt{KaspiMcLaughlin2005}) with spectra harder than those of magnetars.  In fact, \\hbox{PSR~J1119$-$6127} shows a remarkable resemblance to \\hbox{PSR~J1846$-$0258} as far as spin characteristics and the properties inferred from them are concerned.  \\hbox{PSR~J1119$-$6127} has $P=408$\\,ms and $\\dot{P}=4.1\\times10^{-12}$\\,s\\,s$^{-1}$, giving a spin-down luminosity, characteristic age and inferred surface dipole magnetic field very close to those of \\hbox{PSR~J1846$-$0258} \\citep{CamiloKaspiLyne2000}.  This, however, is where the similarity ends.  Across the electromagnetic spectrum, these pulsars have very different attributes.  In the soft $\\gamma$-ray band covered by IBIS/ISGRI, \\hbox{PSR~J1846$-$0258} is clearly detected in the 20--100\\,keV band (see Table~\\ref{TabSpecPar}), while with similar exposure we can place a $2\\sigma$ upper limit of $3.3\\times10^{33}$\\,erg\\,s$^{-1}$ on any emission from the location of \\hbox{PSR~J1119$-$6127} (using a distance of  8.4\\,kpc, \\citealt{CaswellMcClure-GriffithsCheung2004}) in the same energy band.  Also in the X-ray band (0.5--10\\,keV), these two pulsars show contrasting spectral behaviour.  While \\hbox{PSR~J1846$-$0258} shows absorbed power law spectra from both the pulsar and PWN \\citep{HelfandCollinsGotthelf2003}, \\hbox{PSR~J1119$-$6127} shows a strong thermal spectrum superimposed on a hard power law \\citep{GonzalezKaspiCamilo2005}.  Although \\citet{GonzalezKaspiCamilo2005} attribute the hard power law component to the PWN, the thermal X-rays from the pulsar are difficult to interpret through conventional emission models, e.g. models for initial cooling of the entire neutron star result in a radius smaller than allowed by neutron star equations of state, while hot spot models result in unusually high temperature estimates.  It may be the case that, while these two pulsars show similar current spin characteristics, their evolutionary paths and physical properties may be very different, as reflected in the significantly different emission properties.  The most unusual feature of \\hbox{PSR~J1846$-$0258} is its efficiency in converting spin-down power, $\\dot E$, into X- and $\\gamma$-ray luminosity $L_{\\rm X}$ and $L_{\\rm \\gamma}$.  Using data from 41 pulsars, \\citet{PossentiCeruttiColpi2002} determined a general relation between the 2--10\\,keV X-ray luminosity of the combined pulsar and PWN, and the spin-down power.  According to this relation, and the value of $\\dot E$ determined from the spin characteristics, $L_{\\rm 2-10\\,keV}/\\dot{E}$ should be $\\sim 0.2$\\% for \\hbox{PSR~J1846$-$0258}.  The measured efficiency, at $L_{\\rm 2-10\\,keV}/\\dot{E}\\sim 12$\\%, when combined with the 20--100\\,keV measured in this work indicates $L_{\\rm 2-10,20-100\\,keV}/\\dot{E}\\sim27$\\%.  It is expected that the total efficiency in converting spin-down power into luminosity is actually much larger than this, as the luminosity contribution between 10--20\\,keV is not included and the soft $\\gamma$-ray flux extends beyond 100\\,keV (see Fig.~\\ref{FigSpec}).  The observation that the 2--10\\,keV spin-down conversion efficiency is so much larger than average, and the fact that this efficiency  becomes even larger upon inclusion of the 20--100\\,keV luminosity, may be a reason to call the distance estimate into question.  To this effect, we can compare the efficiency of \\hbox{PSR~J1846$-$0258} with two similar young pulsars in the Large Magellanic Cloud, for which the distance is well known at $\\sim50$\\,kpc.  Including luminosity contributions from both the pulsar and PWN \\hbox{PSR~J0537$-$6910} has $L_{\\rm 0.5-10\\,keV}/\\dot{E}=0.9$\\% \\citep{ChenWangGotthelf2006} and \\hbox{PSR~B0540$-$69} has $L_{\\rm 0.5-10\\,keV}/\\dot{E}=9.1$\\% \\citep{KaaretMarshallAldcroft2001}. Even though the estimate for \\hbox{PSR~B0540$-$69} is much larger than the X-ray efficiency of most pulsars, it is still only around half that of \\hbox{PSR~J1846$-$0258} ($L_{\\rm 0.5-10\\,keV}/\\dot{E}=20$\\%).  \\citet{BeckerHelfand1984} determined a distance of 19\\,kpc through neutral hydrogen measurements.  This distance, which agrees with the estimate of \\citet{Milne1979} based on the surface brightness -- diameter relation for supernovae and the neutral hydrogen measurements of \\citet{CaswellMurrayRoger1975}, does imply a very high mean expansion velocity for the supernova remnant \\citep{HelfandCollinsGotthelf2003}.  On the other hand, if the distance were overestimated, this would lead to inconsistency between the neutral hydrogen column as measured by \\citet{BeckerHelfand1984} and that deduced from fits to X-ray data of the PWN \\citep{HelfandCollinsGotthelf2003} -- such inconsistency is not observed.  \\citet{MortonSlaneBorkowski2007} also discuss the possibility of an incorrect distance estimate, but dismiss this option as it would double the already high inferred density of postshock gas in the SNR.  We can estimate a lower limit on the distance of the system by employing the half width at half maximum of the Si line as measured by \\citet{HelfandCollinsGotthelf2003}.  Assuming that the SNR has been expanding at this measured rate (1850\\,km\\,s$^{-1}$) for its entire lifetime (an upper limit of 884 years, \\citealt{LivingstonKaspiGotthelf2006}), we can estimate the distance from the angular size of the SNR on the sky to be $\\sim3$\\,kpc.  The real distance is certainly larger than this lower limit due to the fact that the SNR expansion velocity at earlier times was most likely larger than our current estimate.  Using this lower limit on the distance would provide $L_{\\rm 0.5-10\\,keV}/\\dot{E}\\sim 0.5$\\% -- more in line with the efficiencies of other rotation-powered pulsar systems.   This makes the prospect of a distance smaller than 19\\,kpc to \\hbox{PSR~J1846$-$0258} appealing, but difficult to substantiate in terms of the existing estimates \\citep{CaswellMurrayRoger1975,Milne1979,BeckerHelfand1984}.    Of the rotation-powered pulsars that exhibit X-ray emission, only very few have been detected above 10\\,keV.  Although a consistent treatment exploring the general properties of PWN in the 20--100\\,keV range will be presented as a future publication, we can see from the existing literature that the \\integral spectrum of \\hbox{PSR~J1617$-$5055}, a young rotation powered pulsar, shows a power law slope of $\\sim 2$ \\citep{LandiDeRosaDean2007}, while a combined IBIS/ISGRI and BeppoSAX spectrum of the PWN in \\hbox{PSR~B1509$-$58} \\citep{ForotHermsenRenaud2006} can be fit with a power law of photon index $\\sim2.1$ -- both close to the spectral shape of \\hbox{PSR~J1846$-$0258}.  In addition, the EGRET instrument aboard the Compton Gamma-Ray Observatory detected emission from $\\sim7$ rotation-powered pulsars and/or their PWN in the 30\\,MeV--20\\,GeV energy range \\citep{Roberts2005}. These pulsars have photon indices harder than 2.2, also consistent with the soft \\gm ray spectral slope of \\hbox{PSR~J1846$-$0258} in this work and with those \\integral observations quoted above.  However, smooth spectral coverage above 100\\,keV and into the GeV energy range is required to ascertain the presence or absence of spectral turnovers predicted by models of PWN emission mechanisms \\citep{ZhangHarding2000,ChengHoRuderman1986a}."
    },
    "0710/0710.5511_arXiv.txt": {
        "abstract": "% In hierarchical galaxy formation the stellar halos of galaxies are formed by the accretion of minor satellites and therefore contain valuable information about the (early) assembly process of galaxies. Our GHOSTS survey measures the stellar envelope properties of 14 nearby disk galaxies by imaging their resolved stellar populations with HST/ACS\\&WFPC2.  Most of the massive galaxies in the sample ($V_{\\rm rot}$$>$200 km/s) have very extended stellar envelopes with $\\mu(r)$$\\sim$$r^{-2.5}$ power law profiles in the outer regions. For these massive galaxies there is some evidence that the stellar surface density of the profiles correlates with Hubble type and bulge-to-disk ratio, begging the question whether these envelopes are more related to bulges than to a Milky Way-type stellar halo. Smaller galaxies ($V_{\\rm rot}$$\\sim$100 km/s) have much smaller stellar envelopes, but depending on geometry, they could still be more luminous than expected from satellite remnants in hierarchical galaxy formation models. Alternatively, they could be created by disk heating through the bombardment of small dark matter sub-halos. We find that galaxies show varying amounts of halo substructure. ",
        "introduction": "We select RGB stars from our CMDs and use those to trace the stellar surface density along the minor axis. RGB stars are ideal as they are abundant in our CMDs, are indicative of old stellar populations (as expected to be found in the outskirts of galaxies), and are representative of the underlying stellar mass. To map the surface brightness profiles in the central regions of the galaxies we use the integrated light from Spitzer/IRAC 4.5 micron observations. The Spitzer images provide near unobscured light profiles, even for our edge-on galaxies. We scale the RGB surface density star counts such that they match the IR luminosity profiles in the overlapping region. In this way we derive equivalent surface brightness profiles directly from the RGB star counts.  \\begin{figure} \\mbox{ \\epsfclipon \\epsfxsize=0.49\\textwidth \\epsfbox[83 207 500 550]{prof_min_linr00.eps} \\epsfclipon \\epsfxsize=0.49\\textwidth \\epsfbox[83 207 500 550]{prof_min_logr01.eps} } \\vspace*{-5mm} \\caption{ Minor axis surface density profiles of GHOSTS galaxies. The thin solid lines indicate the profiles derived from Spitzer/IRAC 4.5 micron images calibrated to Vega magnitudes (add about 3.5 mag to convert to Vega $V$-mag). The symbols connected with dotted lines represent RGB star count profiles, scaled to match the Spitzer data. To reduce confusion at small radii we only plot star counts beyond 12 kpc for the non-edge-on galaxies NGC\\,3031/M81 and NGC\\,5236/M83. On a linear radial scale (left diagram) exponential disks appear as straight lines, as indicated for IC5052 by the dot-dashed thick line. In the log-log plot on the right, where we have removed low mass galaxies for clarity, a straight line indicates a power law profile (e.g. thin dashed line = r$^{-2.5}$). Also show are an exponential disk (for NGC\\,0891, dot-dashed line) and a S\\'ersic profile with the typical parameters for a flattened stellar halo as modeled by \\citet{AbaNav06} (thick dashed line). \\vspace{-1mm} \\label{profs} } \\end{figure}  In Fig.\\,\\ref{profs} we show profiles of the edge-on galaxies analyzed so far, along with outer profiles for the more face-on galaxies NGC\\,3031/M81 and NGC\\,5236/M83% . The exponential thin disks only dominate the inner $\\sim$2--3 kpc (5-10 scale heights), while the extended components are evident at larger radii.  We find that eight of the nine galaxies analyzed thus far show components that are more extended than the exponential disks detected at small radii. Several of the most massive galaxies have very extended envelopes with stellar densities at 30 kpc that are 10--100 times higher than the contamination background, equivalent to $\\sim$29 $V$-mag arcsec$^{-2}$. NGC\\,5236/M83 is the only galaxy with a pure exponential disk to the last measured point at 20 kpc (more than 10 disk scale lengths).  \\subsubsection{The bulge-halo connection} In this section we explore the connection between bulges and the extended components. We model the minor axis profiles by combining a S\\'ersic profile and an exponential disk. Merger models show that the hot components resulting after a violent relaxation generally exhibit a S\\'ersic profile \\citep[e.g.,][and reference therein]{BarHer92,AbaNav06}.  If bulges and stellar envelopes are created by a collisionless merger processes, we thus expect their light to follow a S\\'ersic profile.  For a number of massive galaxies ($V_{\\rm rot}$$>$200 km\\,s$^{-1}$) we can fit the entire minor axis profile over a factor of 1000 in size ($\\sim$10$^{4.5}$ in surface density) by an exponential disk and a single S\\'ersic profile, representing both the inner bulge-like region and the outer envelope. This is, for instance, the case for the bulge dominated NGC\\,7814 or a galaxy like NGC\\,3031/M81, which has a power law outer envelope of rather steep slope. Other galaxies, like NGC\\,4565% , have too shallow an outer slope compared to their concentrated bulge region to be fitted by a S\\'ersic profile. NGC\\,891 can be fitted by an exponential disk and a S\\'ersic profile from central bulge to outer envelope if we ignore our outermost field and presume that the higher star density at 30 kpc is due to substructure. Finally NGC\\,5236/M83, which has only a small bulge, shows no sign of an outer envelope out to ten disk scale lengths.  The smallest galaxies in the sample ($V_{\\rm rot}$$\\simeq$100--120 km\\,s$^{-1}$) have small extended components, barely discernible above the background contamination (see Fig.\\,\\ref{profs}). The shape of the extended component is thus poorly constrained due to both the uncertainty in the background and low number statistics. The star counts can be fitted equally well by exponential, power law, and S\\'ersic law profiles. This feature could be the thick disk, as observations of the NGC\\,4244 major axis suggest the component is very flattened. However, the RGB main disk scale height is already twice that of the main sequence population and has been argued already represent the thick disk \\citep{Seth05II} with the feature observed here being an additional component. These additional components are most likely (depending on exact shape) more luminous than predicted in the hierarchical models of \\citet{PurBul07}, but could have been created by the bombardment of small dark matter sub-halos \\citep{KazBul07}.  Therefore, a number of the observed extended envelopes seem structurally directly related to the central bulge regions, like in NGC\\,7814. In other galaxies, where the bulge is too concentrated to be simply related to the outer envelope, we can suspect that secular evolution (e.g., bar driven central enhancement and thickening) can account for the extra (pseudo-)bulge light. NGC\\,4565 with its boxy bulge could nicely fall in this category. In small galaxies the extended envelope is unrelated to the central region, as these small galaxies have no bulge.  \\subsubsection{Envelope properties and halo models} Comparing the envelopes of the different galaxies we find that the two small galaxies have much smaller extended components than the larger galaxies, with surface densities that are lower relative to their disks. The more massive galaxies in our sample are very similar in terms of mass, luminosity, and scale size. Still, there is significant variation in outer envelope properties. At 20 kpc NGC\\,891, NGC\\,4565, and NGC\\,7814 have power law profiles with a slope of about -2.5. NGC\\,3031/M81 has a steeper profile, while at 20 kpc NGC\\,5236 is still dominated by the (face-on) disk. At first sight, the envelope luminosity at 20 kpc seems correlated with Hubble type and bulge-to-disk ratio, with the bulge dominated NGC\\,7814 being the brightest and the late-type spiral NGC\\,5236 showing no sign of an envelope at all. Although M31 also fits this trend, the Sab galaxy NGC\\,3031/M81 does not, as a steeper and fainter profile is evident at 20 kpc.  In Fig.\\,\\ref{profs} we also show a typical profile from the \\citet{AbaNav06} model of accreted stars. Our profiles are a bit shallower and mostly fainter than these models between 10 and 30 kpc. While the surface density normalization may be somewhat uncertain in the models, the shape is quite well constrained. It could be that the true halos only dominate at even larger radii and the slope becomes even shallower at larger radii. However, in hierarchical galaxy formation the halo and ``classical'' bulge are formed by the same merging process, so there is no reason to suspect a large structural change between bulge and halo. The S\\'ersic radii derived for our combined envelope and bulge fits are typically eight times smaller than those of \\citet{AbaNav06}. However, a number of simulation parameters affect the concentration of the accreted halos. Increasing star formation suppression in small sub-halos, such that only the most massive dark matter sub-halos contain stars, yields steeper and fainter envelopes \\citep{BekChi05}. Alternatively, the stars in the accreted satellites could sit deeper in the potential wells of their dark matter sub-halos than simulated in these models, thereby only being tidally stripped closer to the main galaxy, also resulting in more concentrated halos \\citep{BulJoh05}. ",
        "conclusions": ""
    },
    "0710/0710.0272_arXiv.txt": {
        "abstract": "{ANTARES is a large volume neutrino telescope currently under construction off La Seyne-sur-mer, France, at 2475m depth. Neutrino telescopes aim at detecting neutrinos as a new probe for a sky study at energies greater than 1 TeV. The detection principle relies on the observation, using photomultipliers, of the Cherenkov light emitted by charged leptons induced by neutrino interactions in the surrounding detector medium. Since late January 2007, the ANTARES detector consists of 5 lines, comprising 75 optical detectors each, connected to the shore via a 40 km long undersea cable. The data from these lines not only allow an extensive study of the detector properties but also the reconstruction of downward going cosmic ray muons and the search for the first upward going neutrino induced muons.The operation of these lines follows on from that of the ANTARES instrumentation line, which has provided data for more than a year on the detector stability and the environmental conditions. The full 12 line detector is planned to be fully operational early 2008.}     \\begin{document} ",
        "introduction": " ",
        "conclusions": "Great achievements have been made by the Antares collaboration in the last year. The detector is working in nominal mode with 5 lines and should be complete early 2008. Upward neutrino candidates have been found that validate the conceptual method and the chosen techniques. Very exciting times have started with a detector looking for neutrinos in a region of the celestial sky which has never been studied with such a level of sensitivity."
    },
    "0710/0710.0758_arXiv.txt": {
        "abstract": "The star HDE 226868 known as an optical counterpart of the black hole candidate Cyg X-1 has been observed in H$_\\alpha$ region using spectrograph at Ond\\v{r}ejov 2-m telescope. The orbital parameters are determined from He\\,I-line by means of the author's method of Fourier disentangling. Preliminary results are also presented of disentangling the H$_\\alpha$-line into a P-Cyg profile of the (optical) primary and an emission profile of the circumstellar matter (and a telluric component). ",
        "introduction": "The bright X-ray source Cyg X-1 has been identified with the star denoted as HDE 226868, V1357 Cyg or BD+34$^{\\circ}$3815 etc. An improvement of instrumentation of the Ond\\v{r}ejov 2-m telescope enabled to start with systematic observations of this target of magnitude V$\\simeq$8.9, B$\\simeq$9.6. With coordinates $\\alpha_{2000}=19^{h}58^{m}21.7^{s}$, $\\delta_{2000}=+35^{\\circ}12'6''$ it is well observable from Ond\\v{r}ejov mainly at summer.  It is known to be an interacting binary with period $P\\simeq 5.6 d$. The primary component is a supergiant of spectral type classified as B0 (or O9.7) Iab  with temperature $T_{\\rm eff}=30400\\pm500$ K and log $g=3.31\\pm0.07$. This primary, which nearly fills its Roche lobe, shows signs of variable strong stellar wind and an overabundace of He and heavier elements (cf. e.g. Karitskaya et al. 2007).  The secondary component invisible in optical radiation is a compact object, most probably a black hole. This companion, or its neighborhood emits a variable X-radiation, which is supposed to originate from an accretion disk fed by the stellar wind from the primary.  The X-radiation switches chaotically between two states. In the low/hard state the total X-ray flux is low and the spectrum is flat, so that the hard tail of X-radiation prevails. In the high/soft state the soft radiation is enhanced more, and consequently the spectrum has a steeper decrease toward the higher energies and hence the radiation is softer in the mean. Some intermediate states may also appear temporarily.  The X-ray flux is anticorrelated with the strength of emission in the H$_{\\alpha}$-line of the primary: in the X-low/hard state the H$_{\\alpha}$ emission is strong, while in the X-high/soft state the H$_{\\alpha}$ emission is weak.  The aim of the observational campaign at Ond\\v{r}ejov observatory was to improve orbital parameters of the system, to check a possible spectroscopic features connected with the circumstellar matter (either accretion disk around the black hole, gaseous streams or stellar wind) or with a possible third body, and to get line-profiles enabling a quantitative comparison with a model of the atmosphere and stellar wind of the primary. The first part of obtained spectra was provided for a study on Cyg X-1 organized in a wide international collaboration, the results of which should appear in Gies et al. (2007). In the present contribution, results obtained using the author's method of spectra disentangling from the same set of Ond\\v{r}ejov spectra are given. A more detailed study taking into account also recently obtained spectra is in progress. ",
        "conclusions": ""
    },
    "0710/0710.0875_arXiv.txt": {
        "abstract": "Theoretical arguments and indirect observational evidence suggest that the stellar initial mass function (IMF) may evolve with time, such that it is more weighted toward high mass stars at higher redshift. Here we test this idea by comparing the rate of luminosity evolution of massive early-type galaxies in clusters at $0.02\\leq z\\leq 0.83$ to the rate of their color evolution.  A combined fit to the rest-frame $U-V$ color evolution and the previously measured evolution of the $M/L_B$ ratio gives $x = -0.3^{+0.4}_{-0.7}$ for the logarithmic slope of the IMF in the region around 1\\,\\msun, significantly flatter than the present-day value in the Milky Way disk of $x= 1.3\\pm 0.3$. The best-fitting luminosity-weighted formation redshift of the stars in massive cluster galaxies is $3.7^{+2.3}_{-0.8}$, and a possible interpretation is that the characteristic mass $m_c$ had a value of $\\sim 2$\\,\\msun\\ at $z\\sim 4$ (compared to $m_c \\sim 0.1$\\,\\msun\\ today), in qualitative agreement with models in which the characteristic mass is a function of the Jeans mass in molecular clouds. Such a ``bottom-light'' IMF for massive cluster galaxies has significant implications for the interpretation of measurements of galaxy formation and evolution. Applying a simple form of IMF evolution to literature data, we find that the volume-averaged star formation rate at high redshift may have been overestimated (by a factor of $3-4$ at $z> 4$), and the cosmic star formation history may have a fairly well-defined peak at $z\\sim 1.5$. The $M/L_V$ ratios of galaxies are less affected than their star formation rates, and future data on the stellar mass density at $z>3$ will provide further constraints on IMF evolution. The formal errors likely underestimate the uncertainties, and confirmation of these results requires a larger sample of clusters and the inclusion of redder rest-frame colors in the analysis. ",
        "introduction": "The form of the stellar initial mass function (IMF) is of fundamental importance for many areas of astrophysics and a topic of considerable debate (see, e.g., {Schmidt} 1959; {Miller} \\& {Scalo} 1979; {Scalo} 1986; {Larson} 1998, 2003; {Kroupa} 2002; {Chabrier} 2003, for reviews). Measurements of the IMF are difficult and somewhat model-dependent as they require the conversion of the observed present-day luminosity function of a stellar population to its mass function at birth. Best estimates for the Galactic disk suggest that the IMF has a powerlaw slope at $m\\gtrsim 1$\\,\\msun, and  turns over at lower masses ({Kroupa} 2001; {Chabrier} 2003). This turnover can be modeled by a broken powerlaw ({Kroupa} 2001) or by a log-normal distribution with a characteristic mass $m_c$ ({Chabrier} 2003). The value of $m_c$ is $\\sim 0.1$\\,\\msun\\ in the disk of the Milky Way, with considerable uncertainty. The powerlaw slope at high masses is probably close to the {Salpeter} (1955) value of $x=1.35$, with an uncertainty of $\\sim 0.3$ ({Scalo} 1986; {Chabrier} 2003).  Although there is no  direct evidence for dramatic variations of the IMF within the present-day Milky Way disk (e.g., {Kroupa} 2001; {Chabrier} 2003), this does not preclude variations with time, metallicity, and/or environment.  In particular, {Larson} (1998, 2005) has argued that the characteristic turnover mass may be largely determined by the thermal Jeans mass, which strongly depends on temperature ($\\propto T^{3/2}$ at fixed density). In the context of this model one might expect that heating by ambient far-infrared radiation would disfavor the formation of low mass stars in extreme environments, such as in super star clusters and in the center of the Milky Way. Other models emphasize the role of turbulence as opposed to temperature in determining the distribution of protostellar clumps (e.g., {Padoan} \\& {Nordlund} 2002), and in these models the role of the environment may be less direct (see {McKee} \\& {Ostriker} [2007] for a recent review of various models to explain the characteristics of the IMF).  Observations may support the notion of a top-heavy (or ``bottom-light'') IMF in extreme environments. Some young super star clusters in M82 appear to have a top-heavy mass function (e.g., {Rieke} {et~al.} 1993; {McCrady}, {Gilbert}, \\& {Graham} 2003), as do clusters in the Galactic center region (e.g., {Figer} {et~al.} 1999; {Stolte} {et~al.} 2005; {Maness} {et~al.} 2007). The interpretation of observed mass functions is complicated by dynamical effects, which tend to make the mass function more top-heavy over time, in particular in the central regions of star clusters (see, e.g., {McCrady} {et~al.} 2003; {McCrady}, {Graham}, \\& {Vacca} 2005; {Kim} {et~al.} 2006). Recently Harayama, Eisenhauer, \\& Martins (2007) studied the IMF of NGC 3603, one of the most massive Galactic star-forming regions, out to large radii and conclude that its IMF is substantially flatter than Salpeter for masses $0.4 -20$\\,\\msun.  The IMF may also depend on redshift. At earlier times star formation presumably occurred more often in a burst mode than in a relatively gradual ``disk'' mode (e.g., {Steidel} {et~al.} 1996; {Blain} {et~al.} 1999b; {Lacey} {et~al.} 2007), which means that the IMF could generally be more skewed toward high mass stars at redshifts 1--3 and beyond. Furthermore, the average metallicity in star forming clouds was lower at higher redshift, which may have led to an extremely top-heavy IMF for the first generation of stars (e.g., {Abel}, {Bryan}, \\& {Norman} 2002; {Bromm}, {Coppi}, \\& {Larson} 2002). Finally, the cosmic microwave background (CMB) radiation sets a floor to the ambient temperature, and hence the Jeans mass, which scales with $(1+z)$. Beyond $z\\sim 2$ the CMB temperature exceeds the typical temperatures of dense prestellar cores in Galactic molecular clouds (e.g. {Evans} {et~al.} 2001; {Tafalla} {et~al.} 2004). Therefore, at sufficiently high redshift the characteristic mass may be expected to evolve roughly as $m_c \\propto (1+z)^{3/2}$, leading to IMFs which have a reduced fraction of low mass stars ({Larson} 1998). The effects of the CMB are even more pronounced when its influence on the pressure in star-forming clouds is taken into account, and {Larson} (2005) suggests that at $z=5$ the characteristic mass may be higher than today's value by as much as an order of magnitude. Such rapid evolution of the IMF would have important consequences for determinations of masses and star formation rates of distant galaxies, and for measurements of evolution in these properties.  It is very difficult to constrain the IMF at early times directly, as the light of high redshift galaxies is completely dominated by massive stars. The extremely blue rest-frame UV colors of galaxies at $z\\sim 6$ may imply a top-heavy IMF ({Stanway}, {McMahon}, \\&  {Bunker} 2005), although this is just one of several possible explanations. From observations of a lensed Lyman break galaxy there is some evidence that the {\\em slope} of the IMF at $z\\sim 3$ is similar to the Salpeter value at the high mass end ({Pettini} {et~al.} 2000), but there is essentially no information on stars with masses near or below 1\\,\\msun.  Fortunately, the form of the high redshift IMF has implications for the properties of galaxies at much lower redshift, as all stars with masses $\\lesssim 0.8$\\,\\msun\\ that formed in the history of the Universe are still with us today. Tumlinson (2007) finds that the properties of carbon-enhanced metal-poor stars in our Galaxy are best explained with a relatively high number of stars in the mass range 1--8\\,\\msun\\  at high redshift. Various other constraints obtained from galaxies at low redshift (including our own) are reviewed in {Chabrier} (2003). Of particular interest are the stellar populations of massive early-type galaxies, as they are very homogeneous and should reflect conditions in star forming regions at $z> 2$. Recently, {Cappellari} {et~al.} (2006) used the kinematics of elliptical galaxies to constrain the IMF, as the dynamical $M/L$ ratio provides an upper limit to the amount of mass that can be locked up in low mass stars. Current data appear to rule out a {Salpeter} (1955) (or steeper) IMF, but are consistent with {Kroupa} (2001) and {Chabrier} (2003) IMFs ({Cappellari} {et~al.} 2006).  In this paper, we provide new constraints on the IMF at high redshift by comparing the {\\em evolution} of the $M/L$ ratios of early-type galaxies to their color evolution. This method was first suggested by {Tinsley} (1980), but data of sufficient accuracy are only now becoming available. The method is sensitive to the IMF in the important mass range around 1\\,\\msun, where the effects of an evolving characteristic mass might be expected to manifest themselves.  A plan of the paper follows. In \\S\\,2 a relation between color evolution, luminosity evolution, and the logarithmic slope of the IMF $x$ is derived using stellar population synthesis models. In \\S\\,3 published data and archival HST images of galaxy clusters are used to construct the redshift evolution of the $U-V$ color-mass relation. In \\S\\,4 the color evolution from \\S\\,3 is combined with the previously measured evolution of the $M/L_B$ ratio. The relations from \\S\\,2 are then used to derive constraints on the IMF slope $x$ from the combined color and luminosity evolution. Section 5 is devoted to the (many) systematic uncertainties in the methodology and in the data, and \\S\\,6 asks whether our results are consistent with other constraints on the stellar populations of massive early-type galaxies. Although the constraints we derive in this paper are subject to many uncertainties, it is interesting to explore their consequences. In \\S\\,7 the fitting results of \\S\\,4 are interpreted in the context of an evolving characteristic mass $m_c$. The data on cluster galaxies are combined with previous constraints on $m_c$ for globular clusters and submm-galaxies, and a simple form of IMF evolution is proposed. This evolution is then applied to literature data on the evolution of the volume-averaged star formation rate and stellar mass density. The key results are summarized in \\S\\,8. We assume $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$, and $H_0=71$\\,\\kms\\,Mpc$^{-1}$ where needed. ",
        "conclusions": "\\label{conclusion.sec}  This paper compares the color evolution of massive cluster galaxies to their luminosity evolution, with the aim of constraining the form of the IMF at the time when the stars in these galaxies were formed. It is found that the evolution of the rest-frame $U-V$ color is not consistent with the previously determined evolution of the rest-frame $M/L_B$ ratio, unless the IMF slope is significantly flatter than the Salpeter value around 1\\,\\msun. For standard IMFs with a slope of 1.3 at $m\\geq 1$\\,\\msun\\ the luminosity evolution is too fast for the measured color evolution, and the implied stellar ages derived from $M/L$ evolution and color evolution are not consistent with each other. The only models that are able to fit the color evolution and the luminosity evolution simultaneously have IMF slopes of $\\sim 0$ around 1\\,\\msun\\ and mean luminosity-weighted stellar formation redshifts of $\\sim 4$ (for Solar metallicity).  This result is somewhat uncertain, as the currently available sample of cluster galaxies with accurate rest-frame $U-V$ colors and dynamical masses is somewhat limited and there are many systematic effects which may play a role. In particular, it is an open question whether stellar population synthesis models are able to predict color evolution with the required accuracy. The commonly used {Bruzual} \\& {Charlot} (2003) and {Maraston} (2005) models give broadly similar answers, but that may be because they share many of the same assumptions.  As discussed in \\S\\,\\ref{comp.sec} the higher stellar ages implied by a flat IMF are consistent with many other studies, which lends some credibility to the results presented here. Of particular importance is the agreement with the data on Balmer line strengths of {Kelson} {et~al.} (2001), as they do not suffer from the same systematic uncertainties as the color data. Formation redshifts substantially larger than two also fit more comfortably with the direct detection of old galaxies at high redshifts. A firm independent measurement of the star formation epoch of massive cluster galaxies, combined with a better understanding of selection effects at high redshift, would leave the IMF as the only free parameter and greatly simplify the problem.  The implications discussed in \\S\\,\\ref{imply.sec} are obviously somewhat speculative. Although the interpretation in terms of an evolving characteristic mass is physically plausible according to some models (e.g., {Larson} 2005), many other forms of the IMF are consistent with the data. The observations described in this paper are only sensitive to a narrow mass range near 1\\,\\msun, and the IMF proposed in Eq.\\ \\ref{mod.eq} represents a very substantial extrapolation. This is illustrated in Fig.\\ \\ref{summary.plot}: both the top-heavy IMF (red line) and the ``bottom-light'' IMF (green line) are consistent with the data presented in this paper. The main reason for preferring the bottom-light IMF over a top-heavy form is that the absolute $M/L$ ratios of galaxies are within a reasonable range. As shown in Fig.\\ \\ref{mlmc.plot} the $M/L_V$ ratios are similar to those implied by a standard {Chabrier} (2003) IMF, which means that they are consistent with dynamical measurements at $z=0$ ({Cappellari} {et~al.} 2006).  \\vbox{ \\begin{center} \\leavevmode \\hbox{% \\epsfxsize=8.5cm \\epsffile{f15.eps}} \\figcaption{\\small Illustration of the key results of this paper. In \\S\\,4 the slope of the IMF of massive cluster galaxies $x$ is estimated to be approximately $-0.3$ in a narrow region around 1\\,\\msun\\ (thick black line). The red dashed line and the green solid line show two possible interpretations: a global change of the slope of the IMF (with respect to a standard Chabrier or Salpeter IMF) at all masses (red dashed line), or a change in the characteristic mass (solid green line). The ``top-heavy'' interpretation is inconsistent with the dynamical $M/L$ ratios of nearby elliptical galaxies (see \\S\\,\\ref{absml.sec}), whereas the ``bottom-light'' interpretation is consistent with all data that we are aware of. The blue and yellow areas illustrate that stars with masses $\\gtrsim 10$\\,\\msun\\ drive star formation measurements, whereas stars with masses $1-5$\\,\\msun\\ drive $M/L$ measurements (see \\S\\,7.3-7.5). \\label{summary.plot}} \\end{center}}  An ``unintended'' effect of an evolving IMF of the form proposed here is that it reduces the discrepancy between the observed stellar mass density and the density implied by the cosmic star formation history. This result is in excellent (albeit qualitative) agreement with several other recent studies (e.g., Hopkins \\& Beacom 2006, Fardal et al.\\ 2007, P\\'erez-Gonz\\'alez et al.\\ 2007; Wilkins, Trentham, \\& Hopkins 2007; Dav\\'e 2007; see also, e.g., Fields 1999). The differences between a non-evolving IMF and an evolving IMF are fairly large at $z\\sim 4$ (as the effect on the star formation rate is strong and the effect on $M/L$ ratios is weak at that redshift), and it will be interesting to see where future measurements of the mass density at high redshift will fall in Fig.\\ \\ref{mdens.plot}. We note that the effects on the $M/L_V$ ratios are somewhat uncertain, as they rely on rather rudimentary stellar population synthesis modeling. The effects on star formation rates are more robust, and suggest that the cosmic star formation rate peaked at $z\\sim 1.5$.  This paper adds to previous theoretically and observationally motivated suggestions that the IMF may evolve with redshift (e.g., Worthey et al.\\ 1992; Larson 1998, 2005; Fields 1999; Blain et al.\\ 1999a; Baugh et al.\\ 2005; Stanway et al.\\ 2005; Hopkins \\& Beacom 2006; Tumlinson 2007; Lacey et al.\\ 2007; Fardal et al.\\ 2007; P\\'erez-Gonz\\'alez et al.\\ 2007). Although these studies vary greatly in their parameterization of IMF evolution and the range of stellar masses that are considered, they all suggest that the ratio of high-mass stars to low-mass stars was higher in the past. It should be pointed out that most of these papers invoke a change in the IMF as a ``last resort'' possibility, to explain data that are otherwise difficult to interpret. In the present study a different approach was followed, in that we set out with the specific purpose of constraining the slope of the IMF. An advantage of the applied method is that it is fairly direct, as the rate of luminosity evolution is determined by the number of stars as a function of mass. Disadvantages are that it is only sensitive to a very limited mass range (see Fig.\\ \\ref{summary.plot}); that it relies on stellar population synthesis models, which are not well calibrated in the relevant parameter range; that the progenitors of early-type galaxies may not be representative for the general population of high redshift galaxies; and that the currently available data are somewhat limited.  Accepting the possibility of an evolving characteristic mass, it is interesting to speculate what could be the cause, or causes. The proximate cause may well be a higher temperature in molecular clouds at high redshift, which would raise the Jeans mass and could inhibit the formation of low mass stars ({Larson} 1998, 2005). The ultimate cause could be the higher temperature of the cosmic microwave background, the fact that star formation tends to proceed in more extreme environments at higher redshift, or a combination. Available information on IR-bright galaxies suggests that dust temperatures in star burst galaxies are of order $30-40$\\,K ({Dunne} {et~al.} 2000; {Chapman} {et~al.} 2005), and hence exceed the CMB temperature for all relevant redshifts. However, it is as yet unclear what fraction of the total star formation has taken place in these extreme environments (see, e.g., {Reddy} {et~al.} 2007).   The analysis presented here can be improved in various ways. The number of clusters with accurate rest-frame $U-V$ is currently smaller than the number of clusters with accurate $M/L_B$ measurements, and this can be remedied by obtaining accurate (space-based) photometry in well-chosen filters of the remaining clusters in the vv07 sample. It is also important to measure the evolution in a redder rest-frame color, such as $V-I$. Redder color suffer less from possible contributions of hot stars, and their evolution is probably somewhat better calibrated in stellar population synthesis models. This requires very accurate photometry in the near-infrared, which should be possible with WFC3 on HST. On the modeling side, it would be helpful to implement more variations of the IMF in stellar population synthesis codes than the standard Salpeter, Kroupa, and Chabrier forms. Ultimately it may be fruitful (or prove necessary) to have the characteristic mass, or some other parameter describing the form of the IMF, as one of the ``standard'' parameters in these models, on a par with the age and metallicity."
    },
    "0710/0710.2041_arXiv.txt": {
        "abstract": "\\noindent The generalized Darmois--Israel  formalism for Einstein--Gauss--Bonnet theory is applied  to  construct thin-shell Lorentzian wormholes with spherical symmetry.  We calculate the energy localized on the shell, and we find that for certain values of the parameters wormholes could be supported by matter not violating the energy conditions. ",
        "introduction": "For spacetime dimension $D\\geq 5$  the  Einstein--Hilbert action of  gravity  admits quadratic corrections constructed  from coordinate-invariant tensors which scale as fourth derivatives of the metric. In particular, when $D=5$  the most general theory leading to  second order equations for the metric is the so-called Einstein--Gauss--Bonnet theory or Lovelock theory up to second order. This  class of model for higher dimensional gravity has been widely studied, in particular because it can be obtained in the low energy limit of  string  theory \\cite{1}. For spacetime dimensions $D<5$ the Gauss--Bonnet terms in the action represent a topological invariant.  The equations of gravitation admit  solutions, known as Lorentzian wormholes, which connect two regions of the same universe (or of two universes) by a throat, which is a minimal area surface \\cite{motho,book}. Such kind of geometries would present some features of particular interest, as for example the possibility of time travel (see Refs.  \\cite{morris-novikov}). But a central objection against the actual existence of wormholes is that  in Einstein gravity the  flare-out condition \\cite{hovis1} to be  satisfied at the throat  requires the presence of exotic matter, that is, matter violating the energy conditions \\cite{book}. In this sense,  thin-shell wormholes have the advantage that the exotic matter would be located only at the shell.  However, it has recently been shown  \\cite{gra-wi} that in pure Gauss--Bonnet gravity exotic matter is no needed  for wormholes to exist; in fact, they could exist even with no matter (see also Refs. \\cite{7,maeda}).  In this work we  thus study thin-shell wormholes in Einstein--Gauss--Bonnet gravity.  We focus in the amount of matter necessary for supporting the wormholes, without analyzing the microphysics explaining this matter. Differing from the approach in the related work Ref. \\cite{grg06}, where the Gauss--Bonnet terms were treated as an effective source for the Einstein's field equations,  here the Gauss--Bonnet contribution is treated  as an essentially  geometrical object. This requires a generalization \\cite{4} of the Darmois--Israel formalism \\cite{3} for thin shells, but provides a better physical understanding. In particular, we  show that for certain values of the parameters,  thin-shell wormholes could be supported by matter not violating the energy conditions. ",
        "conclusions": "Motivated by the results within pure Gauss--Bonnet gravity (i.e. without Einstein term) in Ref. \\cite{gra-wi}, here we evaluate the amount of exotic matter and the energy conditions, following the approach presented above in which the Gauss--Bonnet term is treated as a geometrical contribution in the field equations. Coming this contribution from the curvature tensor, this approach is clearly the most suitable to give a precise meaning to the characterization of matter supporting the wormhole. As we shall see, the results will considerably differ from those in Ref. \\cite{grg06}, where the Gauss--Bonnet term was treated as an effective source for the Einstein's field equations.  The {\\it weak energy condition} (WEC) states that for any timelike  vector $U^{\\mu}$ it must be $T_{\\mu\\nu}U^{\\mu}U^{\\nu}\\geq0$; the WEC also implies, by continuity, the {\\it null energy condition} (NEC), which means that for any null vector $k^{\\mu}$ it mus be $T_{\\mu\\nu}k^{\\mu}k^{\\nu}\\geq0$ \\cite{book}. In an orthonormal basis the WEC reads $\\rho\\geq0$, $\\rho+p_{l}\\geq0\\  \\forall\\, l$, while the NEC takes the form $\\rho+p_{l}\\geq0 \\ \\forall\\, l$. In the case of  thin-shell wormholes  the radial pressure $p_{r}$ is zero, while within Einstein gravity or even with the inclusion of a Gauss--Bonnet term in the way proposed in \\cite{grg06}, the surface energy density must fulfill  $\\sigma<0$, so that both energy conditions would be  violated. The sign of $\\sigma+p_{t}$  where $p_{t}$ is the transverse pressure  is not fixed, but it depends on the values of the parameters of the system.  In what follows we restrict to static configurations. The surface  energy density $\\sigma_{0}$ and the transverse pressure $p_{0}$ for a static  configuration ($b=b_0$, $\\dot{b}=0$, $\\ddot{b}=0$) are given by \\begin{equation} \\sigma_{0}=-\\frac{1}{8\\pi}\\left[ 6\\frac{\\sqrt{f(b_{0})}}{b_{0}}-2\\alpha\\sqrt{f(b_{0})}\\left(4\\frac{f(b_{0})}{b^{3}_{0}}-\\frac{12}{b^{3}_{0}}\\right)\\right],\\label{s0} \\end{equation} \\begin{equation} p_{0}=\\frac{1}{8\\pi}\\left[ 4\\frac{\\sqrt{f(b_{0})}}{b_{0}}+\\frac{f^{'}(b_{0})}{\\sqrt{f(b_{0})}}-2\\alpha\\left(2f^{'}(b_{0})\\frac{\\sqrt{f(b_{0})}}{b^{2}_{0}} -2\\frac{f^{'}(b_{0})}{b^{2}_{0}\\sqrt{f(b_{0})} }\\right)\\right].\\label{p0} \\end{equation} Note that the sign of the surface energy density is, in principle, not fixed. The most usual choice  for quantifying the  amount of exotic matter in a Lorentzian wormhole is the integral \\cite{nandi}: \\begin{equation} \\Omega= \\int (\\rho + p_{r})\\sqrt{|g|}\\,d^{4}x. \\end{equation} We can introduce a new radial coordinate ${\\cal R}=\\pm(r-b_{0})$ with $\\pm$ corresponding to each side of the shell. Then, because  in our construction the energy density is located on the surface, we can write $\\rho=\\delta({\\cal R})\\,\\sigma_{0}$, and because the shell does not exert radial pressure  the amount of exotic matter reads \\begin{equation} \\Omega=\\int^{2\\pi}_{0} \\int^{\\pi}_{0}\\int^{\\pi}_{0}\\int^{+\\infty}_{-\\infty}\\delta({\\cal R})\\,\\sigma_{0} \\sqrt{|g|}\\, d{\\cal R}\\,d\\theta\\,d\\chi\\ d\\varphi =2\\pi^{2}  b_{0}^3 \\sigma_{0}. \\end{equation} Replacing the explicit form of $\\sigma_{0}$ and $g$, we obtain the exotic matter amount as a function of the parameters  that characterize the configurations: \\begin{equation} {\\Omega}= -\\frac{3}{2}\\pi b^{2}_{0}\\sqrt{f(b_{0})} +2\\pi\\alpha\\sqrt{f(b_{0})}\\left[f(b_{0})-3)\\right], \\end{equation} where $f$ is given by (\\ref{f}). For  $\\Lambda=0$, in the limit $\\alpha\\rightarrow 0$ and Taylor expanding up to zeroth order we obtain the exotic matter for  the Reissner--N\\\"{o}rdstrom ($Q\\neq 0$) and Schwarzschild ($Q= 0$) geometries.     For $\\alpha\\neq 0$  we now find that there exist positive contributions to $\\Omega$; these come from the different signs in the expression (\\ref{s0}) for the surface energy density, because $\\Omega$ is proportional to $\\sigma_0$. We stress that this would not be possible if the standard Darmois--Israel formalism was applied, treating the Gauss--Bonnet contribution as an effective energy-momentum tensor, because this leads to $\\sigma_0\\sim - \\sqrt{f(b_0)}/b_0$ \\cite{grg06}. Now, once the explicit form of the function $f$ (with $\\Lambda=0$) is introduced, the condition $\\sigma_0>0$ leads to \\begin{equation} -8\\alpha-2b_0^2-b_0^2\\sqrt{1+\\frac{16M\\alpha}{\\pi b_0^4}-\\frac{8Q^2\\alpha}{3b_0^6}}>0,\\label{s+} \\end{equation} which can hold only for $\\alpha <0$. The subsequent  analysis is simplified  by considering the  case $Q=Q_c, \\,\\Lambda=0$; then there would be at most only one horizon in the original manifold, its radius being independent of the charge. But for this  charge  it can be shown that, for values of $\\alpha$ such that the horizon exists, it is not possible to fulfil Eq. (\\ref{s+}) for any wormhole radius larger than $r_{hor}$; the reason is that the horizon exists only for $\\alpha > -M/(3\\pi)$, which is not compatible with condition (\\ref{s+}) if $b_0$ is to be larger than the corresponding horizon radius. Instead, a simple numerical analysis shows that for $\\alpha$ slightly below  $-M/(3\\pi)$ both the singularity at $r\\neq 0$ and the horizon dissapear in the original manifold, so that the only condition to be fulfilled  is that given by Eq. (\\ref{s+}). And for $\\alpha <0$ it is always possible to choose $b_0$ such that this indeed happens, so  that   the existence of  thin-shell wormholes  is compatible with a positive surface energy density \\footnote{It is not difficult to see that for $Q$ slightly below $Q_c$ the same happens for larger  $|\\alpha|$.}.   In figures 1 and  2 we show the  amount $\\Omega$ as a function of the wormhole radius  for this  relatively large  value of $|\\alpha|$ (that is, $\\pi |\\alpha|$ of order $M$); though this would imply microscopic configurations or a scenario different from that suggested by present day observation, the  analysis shows that this is the most interesting situation. Besides the fact that $\\Omega$ results to be smaller when calculated by treating the Gauss--Bonnet contribution as a geometric object than in the case that it was treated as an effective energy-momentum tensor, this amount  is  smaller than which would be necessary in the five-dimensional  pure Reissner--N\\\"{o}rdstrom case (see Fig. 1). However, the central, remarkable, result is that  we have a region with $\\Omega >0$ (see Fig. 2), corresponding to  $\\sigma_0>0$; and that besides, from  Eqs. (\\ref{s0}) and (\\ref{p0}) we have $\\sigma_0+p_0=-(b_0/3)\\,d\\sigma_0/db_0$, which shows that for  wormhole radii such that $\\sigma_0 >0$ and ${d\\sigma_0}/{db_0}<0$ ($r_{wh}$ within the maximun and the zero of $\\sigma_0$ in Fig. 3) both the WEC and the NEC are satisfied. Thus, in the picture providing a clear meaning to matter in the shell, in Einstein--Gauss--Bonnet gravity the violation of the  energy conditions could be avoided, and  wormholes could be supported by ordinary matter.  \\begin{figure}[htp] \\centering \\includegraphics[height=8cm, width=12cm]{compara.eps} \\caption{The  amount $\\Omega$   is shown as a function of $r^2_{wh}/M$, for $Q=Q_c$ and  $\\alpha=-0.11 M$. The  dashed line  corresponds to the  five-dimensional Reissner--N\\\"{o}rdstrom case, the dotted line corresponds to considering  the Gauss--Bonnet term as a kind of effective source for the field equations, and the  solid  line shows the result obtained here with the generalized Darmois--Israel formalism for Einstein--Gauss--Bonnet theory.} \\end{figure} \\begin{figure}[htp] \\centering \\includegraphics[height=8.cm, width=8.5cm]{omega.eps} \\caption{The  amount $\\Omega$  is shown as a function of $r^2_{wh}/M$, for $Q=Q_c$ and  $\\alpha=-0.11 M$.  The plot shows the result obtained here with the generalized Darmois--Israel formalism for Einstein--Gauss--Bonnet theory.} \\end{figure} \\begin{figure}[htp] \\centering \\includegraphics[height=8cm, width=8cm]{wec.eps} \\caption{Energy conditions: the dashed line shows $\\tilde \\sigma_0=\\sqrt{M}\\sigma_0$, the dashed-dotted line shows $\\tilde p_0=\\sqrt{M}p_0$ and the solid line shows the sum $\\tilde \\sigma_0+\\tilde p_0$.} \\end{figure}"
    },
    "0710/0710.1873_arXiv.txt": {
        "abstract": "Precision measurement of the scalar perturbation spectral index, $n_s$, from the \\emph{Wilkinson Microwave Anisotropy Probe} temperature angular power spectrum requires the subtraction of unresolved point source power. Here we reconsider this issue, attempting to resolve inconsistencies found in the literature. First, we note a peculiarity in the \\emph{WMAP} temperature likelihood's response to the source correction: Cosmological parameters do not respond to increased source errors. An alternative and more direct method for treating this error term acts more sensibly, and also shifts $n_s$ by $\\sim0.3\\sigma$ closer to unity. Second, we re-examine the source fit used to correct the power spectrum.  This fit depends strongly on the galactic cut and the weighting of the map, indicating that either the source population or masking procedure is not isotropic. Jackknife tests appear inconsistent, causing us to assign large uncertainties to account for possible systematics. Third, we note that the \\emph{WMAP} team's spectrum was computed with two different weighting schemes: uniform weights transition to inverse noise variance weights at $l=500$.  The fit depends on such weighting schemes, so different corrections apply to each multipole range. For the Kp2 mask used in cosmological analysis, we prefer source corrections {$A=0.012\\pm0.005$ $\\mu$K$^2$} for uniform weighting and {$A=0.015\\pm0.005$ $\\mu$K$^2$} for $N_{\\rm obs}$ weighting. Correcting \\emph{WMAP}'s spectrum correspondingly, we compute cosmological parameters with our alternative likelihood, finding $n_s=0.970\\pm0.017$ and $\\sigma_8=0.778\\pm0.045$. This $n_s$ is only $1.8\\sigma$ from unity, compared to  the  $\\sim 2.6\\sigma$ \\emph{WMAP} 3-year result. Finally, an anomalous feature in the source spectrum at $l<200$ remains, most strongly associated with W-band. ",
        "introduction": "Measuring $n_s$, the spectral index of initial scalar fluctuations, which is scale invariant ($n_s = 1$) in the Harrison-Zeldovich model and slightly shallower in inflation models, is difficult, primarily because experimental systematics require control over a broad range of spatial scales.  In inflation, the deviation from unity closely relates to the inflationary potential and the number of $e$-folds of expansion, so a statistically robust measurement of $n_s \\neq 1$ places compelling constraints on the physics of the inflationary epoch.  Because all-sky measurements of the cosmic microwave background (CMB) access the largest observable scales in the universe, the angular power spectrum of the CMB, with a long lever arm, is crucial to such studies.  Indeed, the latest data release from the \\textit{Wilkinson Microwave Anisotropy Probe} (\\emph{WMAP}) claims $\\sim 2.6\\sigma$ deviation from the Harrison-Zeldovich spectrum \\citep{Spergel2007}. Unfortunately, the CMB is not a totally clean measurement.  For example, the well-known degeneracy with the optical depth since recombination ($\\tau$) makes precision measurement of $n_s$ impossible using CMB temperature anisotropies alone, and polarization is required to break it. Complicated noise properties and hints of unknown systematics in the \\emph{WMAP} measurement of large-scale polarization indicate that the systematic uncertainty in both $\\tau$ and $n_s$ should still be considered significant \\citep{Eriksen2007b}.  Another important, but under-appreciated, complication for the measurement of $n_s$ is additional power in the angular spectrum from unresolved, and unmasked, point sources.  At high $l$, this shot noise can significantly bias the power spectrum, and consequently $n_s$. The \\emph{WMAP} team devised a sensible prescription for dealing with this contaminant: 1) Use the spectral energy distribution measured from detected sources (and distinct from the CMB) to infer it for undetected ones; 2) measure the contamination using multifrequency data; 3) correct the spectrum; and 4) marginalize over the measurement error when computing the likelihood \\citep{Hinshaw2003,Hinshaw2007}.  \\citet{Huffenberger2004} found a level of source contamination consistent with the level in the first \\emph{WMAP} data release \\citep{Hinshaw2003}.  However, based on the three year temperature data \\citep{Hinshaw2007}, \\citet{Huffenberger2006} measured a point source spectrum with two irregularities.  First, at $l>200$ the spectrum is white, but with an amplitude below the value in the original preprint of \\citet{Hinshaw2007}. In the present work, we discovered a small error in the power spectra used for the \\citet{Huffenberger2006} estimate, which should have reported $A = 0.013 \\pm 0.001$ $\\mu$K$^2$ instead of $A = 0.011 \\pm 0.001$ $\\mu$K$^2$, still below the original WMAP value of $A = 0.017 \\pm 0.002$ $\\mu$K$^2$.  Prompted by our result, \\citet{Hinshaw2007} re-examined the issue, revising their value down somewhat and increasing the error bars, to $A = 0.014 \\pm 0.003$ $\\mu$K$^2$.  The \\citet{Spergel2007} bispectrum analysis indicates a non-Gaussianity consistent with these values, but lacks the statistical power of the multifrequency power spectrum comparison.  The second peculiarity is that the power at $100<l<200$ in \\citet{Huffenberger2006} was inconsistent, at strong statistical significance, with the rest of the white spectrum.  This paper again considers the power spectrum source correction procedure in detail.  We begin in Section \\ref{sec:ns_impact} with a study of the impact of the source correction on the scalar spectral index through the likelihood.  Following this, we probe the source amplitude in Section \\ref{sec:source_spectra}, examining the dependence of the fit on the sky weighting, mask, year of observation, and frequency dependence, and present our best estimates of the cosmological parameters.  These same tests probe the robustness of the $l<200$ feature.  Finally, we conclude in Section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have reanalyzed the source correction procedure for the three-year data release of \\emph{WMAP}.  First, we considered the impact of this procedure in the \\emph{WMAP} likelihood code. Surprisingly, we found that the \\emph{WMAP} likelihood does not react to changes in the point source correction error.  We have devised a modified likelihood which does respond as expected, although we note that more work is needed to completely validate this approach, as it couples to the important problem of how to approximate a non-Gaussian likelihood with fitting formulas over a wide multipole range.  To conclude $n_s < 1$, a precision measurement of the source contamination is required.  We note that the modes not contaminated favor $n_s$ consistent with unity.  Second, we found several indications that the unmasked source population in \\emph{WMAP} data is anisotropic.  This implies that the combined spectrum should be corrected differently in two multipole regions, based on the weighting of the map.  Anisotropy in the unmasked sources is unexpected, but can be turned to an advantage: By very carefully masking near the ecliptic poles and galaxy, or employing a wide galactic cut, the point source contamination can be cut substantially.  This gain must be weighed against the reduction of the sky area.  We note irregularities in jackknife tests of the source fit, grouped by time of observation or frequency band.  This prompted us to adopt large errors on the source fit, to account for systematics beyond the estimate of statistical error which accompanies our measurement.  This step is necessary to treat the source correction conservatively.  With the modified likelihood and reduced source amplitude from ignoring the excess at $l<200$, these enlarged error bars are responsible for raising our values of $n_s$ and $\\sigma_8$.  Finally, the previously noted anomalous $l < 200$ feature is still present, shows signs of being spatially associated with the galaxy, and is most strongly associated with the W band.  On the other hand, it does not appear to be associated with calibration or beam errors. It may represent an over-subtraction of the foreground template in W, although further investigation is warranted. However, the immediate conclusion is that this part of the spectrum should not be used to infer the point source amplitude at higher $l$'s."
    },
    "0710/0710.0514_arXiv.txt": {
        "abstract": "We study the phase-space behaviour of nearby trajectories in integrable potentials. We show that the separation of nearby orbits initially diverges very fast, mimicking a nearly exponential behaviour, while at late times it grows linearly. This initial exponential phase, known as Miller's instability, is commonly found in N-body simulations, and has been attributed to short-term (microscopic) N-body chaos.  However we show here analytically that the initial divergence is simply due to the shape of an orbit in phase-space. This result confirms previous suspicions that this transient phenomenon is not related to an instability in the sense of non-integrable behaviour in the dynamics of N-body systems. ",
        "introduction": "The problem of how exactly galaxies reach their final equilibrium configuration is still unsolved.  It is clear that, unlike for gases, two-body collisions between stars in galaxies are not the driving mechanism to reach a relaxed state, since the associated timescales are exceedingly large \\citep{bt}. In an attempt to explain the road to equilibrium from a statistical mechanics point of view, Lynden-Bell (1967) introduced the concept of ``violent relaxation''. In this context, the relaxation is reached through the effects of a ``violently changing'' gravitational field. However, the detailed physics of this process also remain to be understood \\citep{arad,valluri}.  Besides the statistical mechanics approach, it is also possible to study the problem of ``relaxation'' at the level of orbits. In this case, it is useful to introduce the concept of mixing, by which we mean how quickly nearby particle trajectories diverge in (phase-)space as a function of time. In the case of time-independent gravitational potentials it is customary to classify mixing into two types. If the particles move in an integrable potential, nearby orbits will diverge as a power-law in time, e.g. \\citet{hw}. This process is known as phase-mixing \\citep{bt}. However, when the potential admits a certain amount of chaos, there exist regions of phase-space where nearby orbits diverge exponentially, evidencing an extreme sensitivity to small changes in the initial conditions \\citep{Arnold}. This process is known as chaotic-mixing \\citep{Kandrup-Mahon,Kandrup}.  These mixing processes can also take place in a time-dependent gravitational potential, in which case the energies of the particles will not be constant. The degree of ``stickiness'', quantified by the time-evolution of the divergence of nearby orbits would then measure the degree of ergodicity of the mixing process.  In the case of chaotic mixing, this could lead to a system that does not have much memory of its evolutionary history. The timescales for evolution could be relatively short, and in principle, this process could be important in the path towards equilibrium for galaxies in the Universe \\citep{merritt,vm}.  Since the 1970s N-body simulations have become the standard tool for studies of the formation and dynamics of structures in the Universe. The question of whether they are a faithful representation of the Universe has always attracted significant attention. This is especially true in recent years \\citep{diemand,binney}, particularly with the finding that dark-matter halos have universal density profiles \\citep{nfw,moore,weinberg-a,weinberg-b}.  One of the first works to focus on how N-body systems evolve was \\citet{miller}. Using what must have been the very first computers in the world, he simulated a self-consistent system in virial equilibrium of 8 upto 32 particles distributed randomly in a cubic volume. Miller found that the trajectories of neighbouring particles initially diverged exponentially. This initial transient has been confirmed using numerical experiments with a significantly larger number of particles \\citep{lecar,Kandrup-Smith, vm, hm}, as well as with various degrees of numerical softening \\citep{ks01}. This implies that the initial exponential divergence cannot be purely attributed to the very grainy nature of the gravitational potential in Miller's experiments. Furthermore, even in high-resolution N-body realizations of well-behaved integrable systems such as the Plummer sphere, nearby orbits experiment a phase of exponential separation at very early times \\citep{kandrup-sideris}. This initial exponential divergence present in N-body simulations is now known as ``Miller's instability''. Understanding this puzzle is the focus of this paper.  That N-body systems would show a certain degree of chaoticity is not necessarily unexpected. However, it seems natural to expect that the larger the number of particles used to represent an otherwise integrable smooth gravitational potential, the more faithful the representation, and hence the lesser the degree of ``numerical'' chaos \\citep{quinlan-tremaine}. There is now significant evidence that when such a system is represented by a sufficiently large number of particles, it does tend to the behaviour expected from the collisionless Boltzmann equation \\citep{goodman,elzant,kandrup-sideris,sideris}.  Nevertheless, even in these high-resolution experiments the initial exponential growth phase is present \\citep{Kandrup-Smith,vm,ks01}. Furthermore, there is evidence \\citep{goodman,hm} that the rate of divergence associated to this phase increases in proportion to the number of particles used. Because Miller's instability only lasts for a very short timescale this does not imply that the system is (macroscopically) chaotic \\citep{vm}.  As stated by \\citet{elzant} it is likely that the ``mechanism leading to the short e-folding time in point particle systems is physically unimportant''.  So, while the existence of a continuum limit in N-body systems appears to be more or less established for long timescales, on short timescales Miller's instability remains a puzzle. The physical mechanism responsible for this was hitherto unknown. It seems quite unlikely that collisions between particles could be important on timescales as short as one-tenth of the crossing time of the system, as measured for example by \\citet{hm}. Microscopic chaos arising from ``white-noise'' or poor orbit integrations are also unlikely to be important on those timescales, particularly in integrable (well-behaved) potentials.  In this paper, we tackle this paradox by studying the {\\it initial} behaviour of nearby characteristics in an integrable smooth (and analytic) potential. Our aim is to understand how these nearby characteristics diverge on short timescales, and if they do so at nearly exponential rates. As we shall demonstrate below, this is indeed the case. The initial behaviour mimics an exponential divergence, but since the system is fully integrable this is not related to the presence of chaos. This near-exponential behaviour merely reflects the time evolution of an orbit in phase-space. This result shows that there is no need to introduce the concept of microscopic chaos, and confirms previous suspicions that this transient phenomenon is not related to an instability in the sense of non-integrable behaviour in the dynamics of N-body systems.  In this paper we describe the evolution of nearby orbits in phase-space, and in particular in configuration space, expanding upon a model developed by \\citet{hw} (hereafter HW). The details of this formalism are given in Sec.~\\ref{sec:aa}. In this Section we focus in detail on the behaviour of nearby orbits in a Plummer potential. In Sec.~\\ref{sec:comp} we summarize our results. ",
        "conclusions": "\\label{sec:comp}  Our analysis shows that the initial nearly-exponential divergence of nearby orbits in N-body systems is not due to chaos. It is present also in integrable smooth potentials, and it reflects a power-law divergence modulated by the shape of an orbit in phase-space.  It is interesting to note that the rates of divergence that we measure using our formalism are in very good agreement with those obtained by \\citet{hm} for N-body realizations of the same Plummer sphere. These authors find a characteristic e-folding time of $t_{cr}/20$ for systems with $N \\sim 10^5$ particles, which is very comparable to the values obtained in Section \\ref{sec:lyap}. They also find a weak dependence on $N$, which may also be readily understood within our framework. Such a dependence is induced by the very rapid decrease in the spatial density of the system.  If a ``relatively'' small number of particles is used in a N-body simulation, then the density cannot be mapped properly. For example, to measure a decline in the density of $10^{-5}$ on a timescale of $\\sim 3 t_{cr}$ as observed in Fig.~\\ref{fig:ex1}, N-body realizations with at least $10^{5}$ nearby particles are needed.  Previous works, including \\citet{miller} and \\citet{kandrup-sideris} have also noted an oscillatory behaviour in the divergence of nearby orbits.  Our analysis, as well as Figure \\ref{fig:ex1} show that this is due to the modulation produced by the periodicity of a regular orbit in phase-space. It is not, as suggested by \\citet{miller}, due to the formation of tight binaries in an N-body system. The fact that such behaviour was visible in the various N-body studies presented in the literature, in fact demonstrates that such N-body systems were faithful representations of the true (integrable) system, at least on short timescales. As stated by \\citet{Kandrup} and \\citet{vm}, the Lyapunov exponents need to be measured in the limit of infinite time intervals; short-time exponential-like divergences do not imply chaotic behaviour."
    },
    "0710/0710.3384_arXiv.txt": {
        "abstract": "We investigate the stellar populations of a sample of 162 Ly$\\alpha$ emitting galaxies (LAEs) at $z = 3.1$ in the Extended Chandra Deep Field South, using deep Spitzer IRAC data available from the GOODS and SIMPLE surveys to derive reliable stellar population estimates.  We divide the LAEs according to their rest-frame near-IR luminosities into IRAC-detected and IRAC-undetected samples.  About 70\\% of the LAEs are undetected in 3.6 \\micron\\ down to $m_{3.6} = 25.2$ AB. Stacking analysis reveals that the average stellar population of the IRAC-undetected sample has an age of $\\sim 200$ Myr and a mass of $\\sim \\pow{3}{8}$ \\Msun, consistent with the expectation that LAEs are mostly young and low-mass galaxies.  On the other hand, the IRAC-detected LAEs are on average significantly older and more massive, with an average age $\\ga 1$ Gyr and mass $\\sim 10^{10}$ \\Msun.  Comparing the IRAC colors and magnitudes of the LAEs to $z \\sim 3$ Lyman break galaxies (LBGs) shows that the IRAC-detected LAEs lie at the faint blue end of the LBG color-magnitude distribution, suggesting that IRAC-detected LAEs may be the low mass extension of the LBG population.  We also present tentative evidence for a small fraction ($\\sim 5\\%$) of obscured AGN within the LAE sample.  Our results suggest that LAEs posses a wide range of ages and masses. Additionally, the presence of evolved stellar populations inside LAEs suggests that the Ly$\\alpha$ luminous phase of galaxies may either be a long-lasting or recurring phenomenon. ",
        "introduction": "\\label{Intro} Lyman-alpha emitting galaxies (LAEs) are important tracers of galaxy formation.  The Ly$\\alpha$ emission is produced by on-going star formation in the galaxies, and the line emission enables discovery of objects that may be too faint to be seen in the continuum.  LAEs therefore offer an opportunity to probe the faint end of galaxy formation at high redshift, and may serve as building blocks of larger galaxies in a hierarchical universe.  In recent years, there have been great advancements in narrow-band surveys of LAEs.  Large samples of candidate LAEs now exist at redshifts $z = 3 - 7$ \\citep[e.g.][]{hu04, malho04, tanig05, shima06, dawso07, muray07, ouchi07}.  One outstanding question regarding the LAEs is their stellar population, which is important for understanding the physical nature of the LAEs and their connection with other high redshift galaxy populations such as Lyman break galaxies \\citep[LBGs;][]{shapl03, ando06, pente07, stanw07}.  A difficulty often encountered when studying the stellar population of high-redshift galaxies is the necessity to have observations at or redward of rest-frame optical in order to constrain the total stellar mass.  At $z \\ga 3$, this means having either deep near-IR or Spitzer IRAC observations.  Much progress has been made recently, however, and several studies have suggested that LAEs are small galaxies with masses $\\la 10^9 - 10^{10}$ \\Msun\\ \\citep{gawis06b, lai07, finke07, pirzk07, nilss07}.  In this study, we focus on a sample of 162 LAEs at $z = 3.1$ discovered in a narrow-band survey of the Extended Chandra Deep Field South (ECDF-S) by \\citet{gronw07}.  The narrow-band 4990 \\AA\\ imaging covers an area of 992 arcmin$^2$ and reaches a completeness limit of $\\sim \\pow{1.5}{-17}$ erg cm$^{-2}$ s$^{-1}$ (equivalent to a Ly$\\alpha$ luminosity $L_{{\\rm Ly}\\alpha} = \\pow{1.3}{42}$ erg s$^{-1}$).  Spectroscopic results confirm the robustness of the sample (details of the spectroscopy will be presented in \\citealt{gawis07} and Lira et al., in preparation).  Of the 52 candidates with sufficient S/N in the spectra to yield redshifts, 51 were confirmed to be $z=3.1$ LAEs (the remaining object is a $z=1.60$ AGN with C~III] 1909 \\AA\\ emission).  The most likely contaminants in our sample are $z=0.34$ [O~II] emitters.  However, there should be a negligible number of these objects in our sample because by requiring the observer's frame equivalent width to be $> 80$ \\AA, we eliminate all but the strongest (and rarest) [O~II] emitters \\citep{gronw07}.  The narrow-band imaging in ECDF-S is supplemented by optical and near-IR observations from MUSYC \\citep{gawis06a}, GOODS \\citep{dicki03}, and GEMS \\citep{rix04} surveys, and Spitzer observations from GOODS and SIMPLE\\footnotemark\\ (Damen et al., in preparation).  Using the multi-wavelength data available in ECDF-S, we study the rest-frame UV to near-IR properties of this large sample of LAEs.  We place special emphasis on the new Spitzer IRAC observations, which sample the rest-frame $0.9 - 2$ \\micron\\ emission from the $z = 3.1$ LAEs and provide valuable constraints on their stellar population and star formation history. \\footnotetext{http://www.astro.yale.edu/dokkum/SIMPLE/}  We assume a ${\\rm \\Lambda}$CDM cosmology with $\\Omega_{\\rm M} = 0.3$, $\\Omega_{\\rm \\Lambda} = 0.7$, and $h = 0.7$.  All magnitudes are in the AB system. ",
        "conclusions": "\\label{Disc} There are several previous studies on the stellar population of LAEs at $z \\sim 3$ and beyond \\citep{gawis06b, finke07, pirzk07, nilss07}. These previous studies found that LAEs have masses ranging from $10^6 - 10^9$ \\Msun, and ages ranging from $0.005 - 1$ Gyr.  Our best-fit average age of 160 Myr and mass of \\pow{3}{8} \\Msun\\ for the IRAC-undetected sample are broadly consistent with these previous studies.  In a further analysis of the IRAC-undetected sample presented in this paper, \\citet[submitted]{gawis07} fit two-component models and find that the IRAC-undetected LAEs have a total mass of \\pow{1}{9} \\Msun, with a young stellar component (age $\\sim 20$ Myr) accounting for $\\sim 20\\%$ of the total mass.  These results are consistent with ours which are derived using single-component models and should be interpreted as the average properties of the entire stellar population within the LAEs.  \\begin{figure} \\includegraphics[width=\\columnwidth]{f4.eps} \\caption{ A comparison of the $R-m_{3.6}$ colors and 3.6 \\micron\\ magnitudes between 8 \\micron\\ detected LBGs (green square), regular LBGs (black dots), and the IRAC-detected LAEs (red stars).  The solid line shows the $R<25.5$ selection limit for LBGs.  The data for the regular LBGs come from IRAC observations of the \\citet{steid03} LBG sample (Magdis et al., in preparation), and the 8 \\micron\\ LBG data come from combining the samples of Magdis et al.\\ and \\citet{rigop06}.  The mean colors and magnitudes of each sample are plotted as large filled symbols, with error bars showing the dispersion within the samples.  We additionally plot the stacked photometry of the IRAC-undetected sample, shown as the red filled red star with the faintest 3.6 \\micron\\ magnitude. \\label{compLBG} } \\end{figure}  The IRAC-detected LAEs are more luminous than the IRAC-undetected LAEs in both rest-frame UV and near-IR, and represent the massive end of the LAE mass spectrum.  Potentially, the IRAC-detected LAEs may provide a link between the LAEs and other galaxy populations.  For instance, they may be the $z \\sim 3$ analogues to the similarly massive IRAC-detected LAEs found at $z \\sim 5.7$ \\citep{lai07}.  In \\Fig{compLBG}, we compare the $R-m_{3.6}$ colors and 3.6 \\micron\\ magnitudes of the LAEs to 8 \\micron\\ detected LBGs and regular LBGs at $z \\sim 3$.  The $R-m_{3.6}$ color and 3.6 \\micron\\ magnitude correlate with stellar mass \\citep{rigop06}, with the 8 \\micron\\ LBGs being the most massive ($\\ga 10^{11}$ \\Msun), followed by LBGs ($\\sim 10^{10} - 10^{11}$ \\Msun), IRAC-detected LAEs ($\\sim 10^{10}$ \\Msun), and finally IRAC-undetected LAEs ($\\sim 10^8$ \\Msun).  The IRAC-detected LAEs occupy the faint and blue end of the LBG color-magnitude distribution, suggesting that they may be the lower mass extension of the LBG population.  This interpretation is supported by the fact that the inferred mass of the IRAC-detected LAEs ($\\sim 10^{10}$ \\Msun) is at the low end of the mass range found for LBGs ($\\sim 10^{10} - 10^{11}$ \\Msun; \\citealt{papov01, shapl01}). Furthermore, most LAEs in the present sample have (observer frame) optical colors that would allow them to be selected as LBGs, although the LAEs tend to be fainter in the rest-frame UV \\citep{gronw07, gawis06b}.  It has also been observed that $\\sim 20\\% - 25\\%$ of LBGs at $z \\sim 3$ exhibit Ly$\\alpha$ emission strong enough to be narrow band excess objects \\citep{steid00, shapl03}.  There is thus accumulating evidence suggesting that the IRAC-detected LAEs may be the bridge connecting the LAE and LBG populations.  Using the stellar mass estimates derived in the previous section and assuming a survey volume of \\pow{1.1}{5} Mpc$^3$ \\citep{gronw07}, we find stellar mass densities of $0.3 \\pm 0.3 \\times 10^6$ and $3 \\pm 1 \\times 10^6$ \\Msun\\ Mpc$^{-3}$ for the IRAC-undetected and IRAC-detected populations, respectively.  The errors in the stellar mass densities include uncertainties in the mass from the fits and uncertainties in the number density coming from Poisson fluctuations and a $\\sim 20\\%$ cosmic variance given the observed LAE bias of $\\sim 1.7$ \\citep[submitted]{gawis07}.  We should stress that the above stellar mass densities only account for LAEs with Ly$\\alpha$ luminosities above our survey completeness limit of \\pow{1.3}{42} erg s$^{-1}$.  The IRAC-undetected LAEs account for about $9 \\pm 10\\%$ of the total stellar mass in LAEs, even though they make up 2/3 of the total by number.  Compared to LBGs and DRGs at $z \\sim 3$, which have stellar mass densities of $\\sim \\pow{1}{7}$ and \\pow{7}{6} \\Msun\\ Mpc$^{-3}$ respectively \\citep{grazi07}, the stellar mass contained in LAEs is smaller, but not insignificant.  However, it is important to keep in mind that there may be substantial overlap between the LAEs and LBGs.  A better understanding of the overlap between these two populations is necessary before a direct comparison of the stellar mass densities can be made.  The present results show that the LAEs posses a wide range of masses and ages, from the massive and evolved IRAC-detected LAEs to the young and small IRAC-undetected LAEs.  The range of photometric properties shown in the LAE sample suggests that LAEs exhibit a continuum of properties between these two extremes.  The presence of both young and evolved stellar populations within the overall LAE population implies that the Ly$\\alpha$ luminous phase of galaxies may last $\\ga 1$ Gyr, or that the Ly$\\alpha$ luminous phase is recurring.  Interestingly, \\citet{shapl01} found that LBGs with best-fit stellar population ages $\\ga 1$ Gyr also show strong Ly$\\alpha$ emission in their spectra. The authors suggest that vigorous past star formation has destroyed and/or expelled the dust inside the galaxies, allowing the Ly$\\alpha$ photons to escape.  One characteristic of these evolved Ly$\\alpha$ emitting LBGs is that they tend to have more quiescent SFR than the rest of the population.  The IRAC-detected LAEs, with their evolved stellar population and low specific SFR, are consistent with this scenario.  If the LAE phase is recurring, then the young age of the IRAC-undetected LAEs implies that their number and stellar mass densities can be as much as a factor $\\sim 10$ higher (very roughly the ratio of the age of the universe to the stellar age).  Similarly young stellar populations ($\\la 100$ Myr) have also been found at higher redshifts \\citep{finke07, pirzk07, verma07}.  The stochastic nature of galaxies with short lifetimes suggests that there may be a related population of undetected pre or post-starburst galaxies that may contribute significantly to the stellar mass and star formation rate densities at $z \\ga 3$."
    },
    "0710/0710.1730_arXiv.txt": {
        "abstract": "The evolution of stellar collision products in cluster simulations has usually been modelled using simplified prescriptions. Such prescriptions either replace the collision product with an (evolved) main sequence star, or assume that the collision product was completely mixed during the collision.  It is known from hydrodynamical simulations of stellar collisions that collision products are not completely mixed, however. We have calculated the evolution of stellar collision products and find that they are brighter than normal main sequence stars of the same mass, but not as blue as models that assume that the collision product was fully mixed during the collision. ",
        "introduction": "The aim of the MODEST collaboration \\citep{article:modest1} is to model and understand dense stellar systems, which requires a good understanding of what happens when two single stars or binary systems undergo a close encounter. A possible outcome of such an encounter is a collision followed by the merging of two or more stars. This is a possible formation channel for blue straggler stars (\\emph{e.g.} \\citet{article:sills_on_axis}). Understanding the formation and evolution of blue stragglers is important for understanding the Hertzsprung-Russell diagram of clusters. ",
        "conclusions": "\\begin{figure} \\ifpdf \\includegraphics[width=\\textwidth]{glebbeek_fig_hrd} \\else \\includegraphics[angle=270,width=\\textwidth]{glebbeek_fig_hrd} \\fi \\caption{Colour-magnitude diagram of the open cluster M67 ($\\blacklozenge$). Overplotted are the locations of our models at $4 \\mathrm{Gyr}$, the age of M67. The black ({\\Large$\\bullet$, $\\blacktriangle$}) symbols are collisions from the M67 simulation. Two of these are double collisions, which are indicated by {\\Large$\\blacktriangle$}. The grey ({\\Large$\\bullet$}) symbols are from our larger grid.} \\label{fig:hrd_m67} \\end{figure}  Compared to normal stars, collision products are helium enhanced. Most of the helium enhancement is in the interior and does not affect the opacity of the envelope. The helium enhancement does increase the mean molecular weight and therefore the luminosity of the star. This decreases the remaining lifetime of collision products compared to normal stars of the same mass and can be important for the predicted number of blue stragglers from cluster simulations. The increased luminosity changes the distribution of blue stragglers in the colour-magnitude diagram, moving it above the extension of the main sequence.  The evolution track of a fully mixed model can be significantly bluer than a self-consistently calculated evolution track of a merger remnant. Fully mixed models are closer to the zero age main sequence.  Our grid of models covers most of the observed blue straggler region of M67 (Figure \\ref{fig:hrd_m67}). A better coverage of the blue part of the region can be achieved by increasing the upper mass limit in the grid. The brightest observed blue straggler falls outside the region of our grid because it requires at least a double collision to explain."
    },
    "0710/0710.1506_arXiv.txt": {
        "abstract": "We have modeled the emission of \\hdo\\ rotational lines from the extreme C-rich star IRC+10216. Our treatment of the excitation of \\hdo\\ emissions takes into account the excitation of \\hdo\\ both through collisions, and through the pumping of the $\\nu_2$ and $\\nu_3$ vibrational states by dust emission and subsequent decay to the ground state. Regardless of the spatial distribution of the water molecules, the \\hdo\\ $1_{10}-1_{01}$ line at 557 GHz observed by the {\\em Submillimeter Wave Astronomy Satellite} (SWAS) is found to be pumped primarily through the absorption of dust-emitted photons at 6 $\\mu$m in the $\\nu_2$ band. As noted by previous authors, the inclusion of radiative pumping lowers the ortho-\\hdo\\ abundance required to account for the 557 GHz emission, which is found to be $(0.5-1)\\times10^{-7}$ if the presence of \\hdo\\ is a consequence of vaporization of orbiting comets or Fischer-Tropsch catalysis. Predictions for other submillimeter \\hdo\\ lines that can be observed by the {\\em Herschel Space Observatory} (HSO) are reported. Multitransition HSO observations promise to reveal the spatial distribution of the circumstellar water vapor, discriminating among the several hypotheses that have been proposed for the origin of the \\hdo\\ vapor in the envelope of IRC+10216. We also show that, for observations with HSO, the \\hdo\\ $1_{10}-1_{01}$ 557 GHz line affords the greatest sensitivity in searching for \\hdo\\ in other C-rich AGB stars. ",
        "introduction": "\\label{sec:intro}  The discovery of water vapor emission in the extreme C-rich AGB star IRC+10216 with SWAS \\citep{mel01}, its confirmation by ODIN \\citep{has06}, and the subsequent detection of other O-bearing molecules like OH \\citep{for03}, H$_2$CO \\citep{for04}, and C$_3$O \\citep{ten06}, have challenged our current understanding of the chemistry in envelopes around C-rich AGB stars. According to standard models, essentially all oxygen nuclei are predicted to be locked into CO or SiO with no reservoir to form other O-bearing molecules (except for low abundances of species such as HCO$^+$ in the outer envelope where photochemistry is important). Therefore, the unexpectedly high abundances found for H$_2$O, OH, and H$_2$CO, indicate that several processes not included in standard models for C-rich environments are driving the oxygen chemistry, but the dominant water production mechanism is still a source of debate.   Four distinct mechanisms have been considered as possible sources of the observed water vapor in IRC+10216, each one making a specific prediction for the \\hdo\\ spatial distribution in the envelope: $(i)$ chemistry in the inner envelope, which would imply the presence of \\hdo\\ in the warmest and densest regions;  $(ii)$ vaporization of icy orbiting bodies that have survived from the main sequence into the C-rich AGB phase \\citep{mel01,for01}, predicting the release of \\hdo\\ at intermediate radii of $R_{int}=(1-5)\\times10^{15}$ cm; $(iii)$ grain surface reactions, such as the Fischer-Tropsch catalysis on the surfaces of small metallic grains \\citep{wil04}, which predicts \\hdo\\ to attain a nearly uniform abundance at radii larger than $R_{int}=(1.5-2)\\times10^{15}$ cm; $(iv)$ chemistry involving photodissociation products in the outermost regions of the envelope: a specific mechanism relying on the radiative association O+H$_2$$\\rightarrow$\\hdo+$\\gamma$ has been proposed by Ag\\'undez \\& Cernicharo (2006, hereafter AC06), and predicts \\hdo\\ to be present at radii higher than $R_{int}\\approx4\\times10^{16}$ cm. In all cases, \\hdo\\ is expected to have a uniform abundance from $R_{int}$ up to the external region where it is photodissociated producing OH.   The only water line detected so far in IRC+10216, the ortho-\\hdo\\ \\t110101\\ transition at 557 GHz, is the ground-state transition, with the upper level at only 27 K above the ground rotational level, and cannot discriminate -by itself- between the processes listed above. Nevertheless, the launch of the Herschel Space Observatory (HSO) will allow us to observe other \\hdo\\ lines in the submillimeter and far-infrared domains, permitting us to infer the region where the \\hdo\\ emission arises, and will potentially favor one of the proposed formation mechanisms. The Heterodyne Instrument for the Far Infrared (HIFI) onboard HSO will provide very high spectral resolution observations (0.14-1.0 MHz), thus permitting the lines to be velocity resolved. In this paper, we model the \\hdo\\ emission from IRC+10216 to show how the \\hdo\\ spatial distribution can be inferred from HSO multi-transition observations. Also, we explore which \\hdo\\ transition provides the most sensitive means of searching for water vapor around AGB stars other than IRC+10216; such a search would determine whether the occurence of \\hdo\\ in C-rich environments is widespread. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Model uncertainties}  Our models for IRC+10216 predict that, whatever the region where \\hdo\\ is formed or released to the outflow, \\hdo\\ is radiatively excited, which implies that model results are independent of the gas temperature and density profiles. This result relies on the extrapolation to higher $T_k$ of the collisional rates given by \\cite{phi96}, which will require further confirmation from new estimates of \\hdo-H$_2$ collisional rates at higher temperatures. Nevertheless, we find that collisional rates would have to be about one order of magnitude higher than estimated to compete efficiently with the radiative rates in the innermost regions of the envelope, which seems somewhat implausible.   A relatively uncertain parameter in our models is the assumed distance to IRC+10216, $D=170$ pc. It has been proposed that the distance may be substantially smaller, $D=100-150$ pc \\citep*{zuc86,gro92}. At 170 pc, the inferred stellar luminosity is $L_*\\approx2\\times10^4$ \\Lsun, which results in $L_*\\approx10^4$ \\Lsun\\ at $D=120$ pc, a value more similar to the inferred luminosities of other C-rich AGB stars with similar \\Mdot\\ \\citep*{sch06}. If $D=120$ pc, both \\Mdot\\ and the radiation density at 6 $\\mu$m are a factor of 2 lower than assumed, so that the radiative-to-collisional pumping rate ratio still remains unchanged. Since the closer proximity and weaker 6 $\\mu$m radiation density have opposite effects on the \\hdo\\ outflow rate required to match the SWAS 557 GHz flux, the latter will decrease by less than a factor 2, and thus the expected \\hdo\\ abundance will be a factor of $<2$ higher than in Table~\\ref{tab:h2omodels}. Therefore, we conservatively estimate an o-\\hdo\\ abundance in the range $(0.5-1)\\times10^{-7}$ for $R_{int}=2\\times10^{15}$ cm. Concerning the line ratios to be observed by HSO, we expect values similar to those obtained at $D=170$ pc if the line emission remains unresolved, i.e. for low values of $R_{int}$. For high values of $R_{int}$, beam effects will be more important at $D=120$ pc, diminishing the low-excitation line fluxes relative to the 557 GHz line.  A potentially more important source of uncertainty is the assumption of spherical symmetry in our models. Both the molecular and dust emission from IRC+10216 show evidence for departures from a smooth distribution in a spherically symmetric outflow; incomplete, discrete shells or arcs and clumpy structures are instead observed on a wide range of distances to the star \\citep*[e.g.][]{luc99,mau99,fon03,lea06}. Even more important, the OH line shapes observed in IRC+10216 strongly suggest an asymmetric distribution of OH in the outer regions of the envelope \\citep{for03}, thus also suggesting an asymmetric distribution of the parent \\hdo\\ molecule. These asymmetries may alter to some extent the \\hdo\\ line flux ratios calculated with the use of our spherically symmetric approach. While in spherical symmetry the emission from any line is isotropic, in filamentary structures or slabs of velocity-coherent gas the optically thick lines radiate preferentially in the direction perpendicular to the two faces of the sheet \\citep*{eli89}, whereas the emission from thinner lines will approach a more isotropic behavior. Future HSO observations will show the importance of the observed morphological complexity on the line fluxes by showing whether the different \\hdo\\ line flux ratios are consistent with a single value of $R_{int}$, or indicate a range of values.  Finally, the models assume a water shell with uniform \\hdo\\ abundance and sharp inner and outer radii; however, a finite abundance gradient obviously takes place at both the \\hdo\\ formation and dissociation regions, and variations of the \\hdo\\ abundance across the shell are also possible. These effects may also alter to some extent the expected line flux ratios. Uncertainties in the outer radius of the \\hdo\\ shell may also affect the expected fluxes from stars with low mass loss rates.    \\subsection{Water formation at inner or intermediate radii}  Both the cometary and Fischer-Tropsch catalysis hypothesis predict \\hdo\\ to be released or formed at intermediate radii of {\\it a few} $\\times10^{15}$ cm, and although no specific mechanism has been proposed for \\hdo\\ formation at the innermost regions ($R_{int}<10^{15}$ cm), this possibility cannot be rejected. The models shown in section~\\ref{sec:predictions} will permit us to discriminate, from HSO observations of IRC+10216, if there are significant amounts of \\hdo\\ in the innermost regions through the observation of mid-excitation \\hdo\\ transitions. More difficult will be a priori to discriminate between the cometary and Fischer-Tropsch catalysis propositions. Both predict similar values for $R_{int}$; nevertheless, if $R_{int}$ were found to be significantly higher than $2\\times10^{15}$ cm, the release of \\hdo\\ from comets could be favoured, unless some additional mechanism were found to shift $R_{int}$, within the Fischer-Tropsch catalysis framework \\citep{wil04}, outwards in the envelope. The search for water emission at 557 GHz in C-rich AGB stars other than IRC+10216 will also favour one of the two hypotheses: since IRC+10216 is at the high end of mass loss rates from C-rich stars, and the efficiency of Fischer-Tropsch catalysis to form \\hdo\\ molecules is expected to decrease with diminishing \\Mdot, one would expect in such a case a rate of detection significantly lower than that quantified in section~\\ref{sec:crich} for the cometary hypothesis, and one would expect a pronounced decline of the \\hdo\\ abundance with diminishing \\Mdot.  The ISO/LWS spectrum of IRC+10216 \\citep{cer96} shows an emission feature at 179.5 $\\mu$m, coincident with the wavelength of the \\t212101\\ o-\\hdo\\ transition, with a flux of $\\approx8\\times10^{-20}$ W cm$^{-2}$. This flux is very similar to that computed for the quoted line in model $B$ ($R_{int}=2.1\\times10^{15}$ cm). However, the far-infrared spectrum of IRC+10216 shows vibrationally excited rotational emission of HCN, and the combined emission of the $\\nu_1=1$ $J=19\\rightarrow18$ and $\\nu_3=1$ $J=19\\rightarrow18$ HCN lines, both emitting at 179.5 $\\mu$m, is expected to be also comparable to the measured line flux at 179.5 $\\mu$m \\citep{cer96}.  Given the uncertainties inherent to the HCN model in \\cite{cer96}, where the excited vibrational states are assumed to be thermalized, and given the high density of spectral lines in the IRC+10216 far-infrared spectrum, which makes it difficult to establish the contribution from the $\\nu_1=1$ and $\\nu_3=1$ rotational lines to the spectrum, the relative contribution from HCN and \\hdo\\ to the observed spectral feature is quite uncertain. On the other hand, the expected flux of the \\t303212\\ line at 174.6 $\\mu$m in model $B$ is $\\approx4\\times10^{-20}$ W cm$^{-2}$, below the 3-$\\sigma$ upper limit derived for that line from the ISO/LWS spectrum. These considerations suggest weakly that \\hdo\\ in IRC+10216 is formed or released at radial distances $R_{int}\\gtrsim2\\times10^{15}$ cm.  Our models show that the required \\hdo\\ abundance for $R_{int}$ in the range $(2-5)\\times10^{15}$ cm is $(0.5-1)\\times10^{-7}$ relative to H$_2$, which makes the requirements previously demanded for the cometary and Fischer-Tropsch catalysis hypothesis to work in IRC+10216 less restrictive \\citep{for01,wil04}. The water outflow rate is $(0.5-1)\\times10^{-5}$ M$_{\\oplus}$ yr$^{-1}$, which yields, within the framework of vaporization of icy bodies, a required total initial ice mass of $(0.5-10)$ M$_{\\oplus}$ for \\Mdot(\\hdo)/$M_0({\\rm ice})$ in the range $10^{-5}-10^{-6}$ yr$^{-1}$ \\citep{for01}. On the other hand, the Fischer-Tropsch catalysis on metallic grains will require a density of iron grains relative to total H nuclei of $n_g({\\rm Fe})/n_{{\\rm H}}=(0.5-1)\\times10^{-13}$ to explain the observed \\hdo\\ 557 GHz emission \\citep{wil04}.    \\subsection{Water formation in the outermost layers} \\label{sec:radassoc}    Multitransition HSO observations of \\hdo\\ in IRC+10216 will easily establish whether or not water is formed in the external layers of the envelope; the predicted fluxes in Fig.~\\ref{fig:flujosrint2} and the HSO-HIFI sensitivities in Table~\\ref{tab:h2otrans} indicate that only the \\t110101\\ o-\\hdo\\ line, and possibly the \\t111000\\ and \\t202111\\ p-\\hdo\\ lines, are detectable in one hour of observing time for $R_{int}\\gtrsim4\\times10^{16}$ cm.  For high $R_{int}$, the expected flux in the 557 GHz line can be obtained analytically. Since essentially all o-\\hdo\\ molecules are in the ground $1_{01}$ level (Fig.~\\ref{fig:pump557ghz}a), the o-\\hdo\\ \\t110101\\ line flux is derived from the radiative pumping rate given in eq.~(\\ref{gammar}) after multiplying $\\Gamma_r$ by 1.35 to account for the two radiative pumping routes that are relevant at high distances from the star (Fig.~\\ref{fig:pump557ghz}b): \\begin{eqnarray} F(1_{10}\\rightarrow1_{01})  =  1.4\\times10^{-21} \\times \\left( \\frac{4\\times10^{17} \\, {\\rm cm}}{R_{out}} \\right) \\nonumber \\\\ \\times \\left( \\frac{R_{out}}{R_{int}} -1 \\right)  \\times \\left( \\frac{X({\\rm o-H_2O})}{5\\times10^{-7}} \\right) \\,\\, {\\rm \\frac{W}{cm^{2}}}, \\label{fluxanal} \\end{eqnarray} where $X$ is the abundance relative to H$_2$. Equation~(\\ref{fluxanal}) is independent of the assumed distance $D$ to the star because, in order to match the observed mid-IR continuum, $\\Gamma_r\\propto D^{2}$ (eq.~\\ref{eqj}); it assumes a water shell with sharp edges at $R_{int}$ and $R_{out}$ and constant \\hdo\\ abundance; it ignores slight opacity effects in the ro-vibrational lines (section~\\ref{sec:110101}) as well as beam effects (both slightly raise the required abundance), and overestimates by less than 20\\% the 557 GHz line flux obtained in model $C$. In the limit $R_{out}\\rightarrow\\infty$, and using $F(1_{10}\\rightarrow1_{01})=10^{-20}$ W cm$^{-2}$ \\citep{mel01}, eq.~(\\ref{fluxanal}) gives \\begin{equation} \\chi({\\rm H_2O}) \\gtrapprox 2.4\\times10^{-7} \\times \\left( \\frac{R_{int}}{4\\times10^{16} \\, {\\rm cm}} \\right), \\label{fluxanallim} \\end{equation} where $\\chi({\\rm H_2O})$ is here the \\hdo\\ (ortho+para) abundance relative to H nuclei. Equation~\\ref{fluxanallim} gives the sharp-inner edge, lower limit for the \\hdo\\ abundance required to account for the observed 557 GHz line flux, as a function of $R_{int}$. AC06 have reported an abundance of $\\chi({\\rm H_2O})\\sim10^{-7}$ to account for the observed 557 GHz line flux in IRC+10216. However, the quoted abundance would imply an inner radius of $R_{int}\\lessapprox2\\times10^{16}$ cm, but the \\hdo\\ abundance shown by AC06 (their Fig. 7) decreases sharply at $r<4\\times10^{16}$ cm. Based on detailed modelling, we indicate that the $\\chi({\\rm H_2O})$ profile given by AC06 has to be shifted up by a factor of $\\approx2.5$ to account more accurately for the 557 GHz line flux measured by SWAS.   In the model proposed by AC06, atomic oxygen is produced in a shell by the photodissociation of CO (and particularly the $^{13}$CO isotopologue, which shields itself far less effectively than $^{12}$CO.)  Since the temperature is low ($\\sim 10$~K) within the shell where the atomic oxygen abundance is significant, the neutral-neutral reaction sequence $$\\rm O + H_2 \\rightarrow OH + H$$ $$\\rm  OH + H_2 \\rightarrow H_2O + H$$ is very slow --and therefore is negligible as a source of H$_2$O-- the first reaction being endothermic and the second --although exothermic-- possessing a substantial activation energy barrier.  AC06 therefore proposed the radiative association reaction $$\\rm O + H_2 \\rightarrow H_2O + \\gamma$$ as an alternative source of H$_2$O.  To match the SWAS- and ODIN-observed water line fluxes, AC06 had to posit that this reaction is relatively rapid at low temperature, with a rate coefficient $k_{ra} \\rm \\sim 10^{-15}\\,cm^3\\,s^{-1}$. According to our estimation above, $k_{ra}$ has to be a factor $\\approx2.5$ higher than this value to account for the measured \\t110101\\ \\hdo\\ flux. Nevertheless, and even with the use of the $k_{ra}$ estimation given by AC06, we find that this large reaction rate coefficient for the radiative association of O and H$_2$ is inconsistent with observations of H$_2$O and OH in at least one translucent molecular cloud, for which sufficient data exist to disipate any significant ambiguity: observations of the cloud along the sight-line to HD 154368 --carried out by \\cite{spa98} with the use of the Goddard High Resolution Spectrograph (GHRS) on the {\\it Hubble Space Telescope} (HST) -- yield a 3~$\\sigma$ upper limit on the water abundance that lies almost two orders of magnitude below the value that would obtain were $k_{ra}$ as large as $\\rm 10^{-15}\\,cm^3\\,s^{\\rm -1}$. The factor of discrepancy would rise above $150$ for $k_{ra} \\rm \\approx 2.5 \\times 10^{-15}\\,cm^3\\,s^{-1}$.  The best-fit model for the HD 154368 cloud obtained by \\cite{spa98} posits a plane parallel cloud of total visual extinction $A_V = 2.65$~mag in which the density of H nuclei is $n_{\\rm H} = 325\\, \\rm cm^{-3}$ \\citep[this is probably a lower limit at the cloud center, as the density inferred from the CO $J=1\\rightarrow0/3\\rightarrow2$ ratio is $\\sim10^3$ cm$^{-3}$; see][]{van91} and the external ultraviolet radiation field is 3 times mean interstellar value given by \\cite{dra78}: $I_{UV} = 3$.  Under these conditions, the destruction of H$_2$O is dominated by photodissociation at a rate $\\zeta = 5.9 \\times 10^{-10}{\\rm s}^{-1} \\times  I_{UV} \\times (\\exp\\{-1.7 A_{V1}\\} + \\exp\\{-1.7 A_{V2}\\})/2$ \\citep*{let00}\\footnote{We note that the factor of 1.7 in the exponential, widely used in the literature, yields \\hdo\\ photodissociation rates at the midplane of the plane-parallel cloud that are higher than the values reported by \\cite{rob91} by factors of 2.2 and 10 for $A_V^{tot}=1$ mag and $A_V^{tot}=10$ mag, respectively; therefore, our estimation for the $N({\\rm H_2O})/N({\\rm O})$ ratio predicted by the radiative association of O and H$_2$ is probably a conservative lower limit.}, where $A_{V1}$ is the visual extinction to one cloud surface and $A_{V2} = 2.65 - A_{V1}$ is the extinction to the other.  At the cloud center, the water photodissociation rate is $1.9 \\times 10^{-10}\\,\\,{\\rm s}^{-1}$.  For a radiative association rate of $k_{ra}$, the ratio of water vapor to atomic oxygen is therefore given by the expression $$n({\\rm H_2O})/n({\\rm O}) = k_{ra} n_{\\rm H} f_{\\rm H_2}/ \\zeta,$$ where $f_{\\rm H_2} \\equiv n({\\rm H_2}) / n_{\\rm H}$.  At the cloud center, the cloud is almost fully molecular, with $f_{\\rm H_2} \\sim 0.5$, and the above equation yields $n({\\rm H_2O})/n({\\rm O}) = 8.7 \\times 10^{-4} \\,(k_{ra}/\\rm 10^{-15}\\,cm^3\\,s^{\\rm -1})$.  In Fig.~\\ref{fig:nh2ono}, we show the predicted $n({\\rm H_2}) / n_{\\rm H}$ ratio for the best-fit \\cite{spa98} model, together with the $n({\\rm H_2O})/n({\\rm O})$ ratio that would result if $k_{ra}$ were $\\rm 10^{-15}\\,cm^3\\,s^{-1}$. The model predicts hydrogen to be predominantly in molecular form, in agreement with results by \\cite{sno96}. Averaging the $n({\\rm H_2O})/n({\\rm O})$ ratio over the entire sight-line, we obtain a column density ratio $N({\\rm H_2O})/N({\\rm O}) = 5.3 \\times 10^{-4} \\,(k_{ra}/\\rm 10^{-15}\\,cm^3\\,s^{\\rm -1})$.   The atomic oxygen column density along the HD 154368 sight-line is $N({\\rm O}) = 1.2 \\times 10^{18}\\,\\,{\\rm cm^{-2}}$ (Snow et al. 1996, from absorption line observations of the 1355 \\AA\\ OI] line). Based on a search for the $C^1B_1 - X^1A_1$ band of water vapor near 1240 \\AA, \\cite{spa98} obtained 3~$\\sigma$ upper limit on the water column density of $9 \\times 10^{12}\\,\\,{\\rm cm^{-2}}$, corresponding to $N({\\rm H_2O})/N({\\rm O}) = 7.5 \\times 10^{-6}$.  This 3~$\\sigma$ upper limit lies a factor 70 below the value that we would obtain were $k_{ra}$ equal to $10^{-15}\\rm \\,cm^3\\,s^{-1}$ as AC06 suggested, and places a 3~$\\sigma$ upper limit of $1.4 \\times 10^{-17}\\rm  \\,cm^3\\,s^{-1}$ on $k_{ra}$.  AC06 suggested that the freeze out of oxygen onto grain mantles could diminish the water vapor abundance in molecular clouds, as proposed by \\cite{ber00} to explain the low H$_2$O abundances measured by SWAS in dense clouds.  Ice absorption line observations of diffuse/translucent sight-lines, however, indicate that water ice is generally present only in clouds of $A_V \\ge 3$ \\citep{whi01}. Furthermore, in the specific case of HD 154368 under present consideration, the atomic oxygen is known from direct measurement to be $1.2 \\times 10^{18}\\,\\,{\\rm cm^{-2}}$.  The column density of H nuclei along this sight-line, $N_{\\rm H} = N({\\rm H}) + 2 N({\\rm H}_2)$, has been measured to be $4.2 \\times 10^{21}\\,\\,\\rm cm^{-2}$ \\citep{sno96}, so the mean line-of-sight oxygen abundance is $n({\\rm O})/n_{{\\rm H}} \\sim 3 \\times 10^{-4}$, a value that is entirely consistent with the abundances measured along diffuse sight-lines \\citep{mey98} and inconsistent with a significant depletion of oxygen onto ice mantles.   In summary, the upper limit on the water vapor abundance observed towards HD 154368 definitively rules out a rate coefficient for the radiative association reaction $\\rm O + H_2 \\rightarrow H_2O + \\gamma$ that is large enough to explain --in the context of the AC06 model-- the H$_2$O \\t110101\\ line strength measured by SWAS toward IRC+10216. If HSO observations would indicate that \\hdo\\ is formed in the external layers of IRC+10216, an explanation other than the radiative association proposed by AC06 would be required to avoid incompatibilities with observations toward HD 154368."
    },
    "0710/0710.4786_arXiv.txt": {
        "abstract": "The proper motion measurements for 143 previously known L and T dwarfs are presented. From this sample we identify and discuss 8 high velocity L dwarfs.  We also find 4 new wide common proper motion binaries/multiple systems. Using the moving cluster methods we have also identified a number of L dwarfs that may be members of the Ursa Major (age $\\approx$ 400 Myr), the Hyades (age $\\approx$ 625 Myr) and the Pleiades (age $\\approx$ 125 Myr) moving groups. ",
        "introduction": "Brown dwarfs may be thought of as failed stars. These low mass ($\\leq$70 M$_{\\rm Jup}$ \\citet{burrows01}), cool objects are the lowest mass objects that the star formation process can produce. The majority of the brown dwarfs that have been discovered to date are field objects discovered using surveys such as the Two Micron All Sky Survey (2MASS; \\citet{skrutskie06}, see \\citet{leggett02} for examples), the DEep Near-Infrared Sky survey (DENIS; \\citet{denis05}, see \\citet{delfosse99} for examples), the Sloan Digital Sky Survey (SDSS;\\citet{york00} see \\citet{hawley02} for examples) and the UKIRT Deep Infrared Sky Survey (UKIDSS; \\citet{lawrence06}, see \\citet{kendall07} for examples). However, to study brown dwarfs in depth, a knowledge of their age is essential, which means we must study brown dwarfs in open star clusters or moving groups.  Once a brown dwarf has been proved to belong to an open star cluster, or a moving group, then the age of the dwarf is known, allowing meaningful comparisons to evolutionary models to be made. The most recent example of this is the study done by \\citet{bannister07} who used existing proper motions and parallax measurements to show that a selection of field dwarfs in fact belong to the Ursa Major and Hyades moving groups. The importance of this study, is that these are the first brown dwarfs to be associated with an older cluster or group. Older clusters such as the Hyades are expected to contain very few or no brown dwarfs or low mass members, due to the dynamical evolution of the cluster over time \\citep{adams02}. However, these escaped low mass objects may remain members of the much larger moving group that surrounds the cluster.  To continue the  study started by \\citet{bannister07}, proper motions need to be measured for the majority of the field brown dwarfs currently known. This has been accomplished using the wide field camera (WFCAM, \\citet{casali07}) of the United Kingdom Infrared Telescope (UKIRT). Using these WFCAM images and existing data  we have measured proper motions for 143 L and T dwarfs listed in the online Dwarf Archive\\footnote{See http://spider.ipac.caltech.edu/staff/davy/ARCHIVE/, a webpage dedicated to L and T dwarfs maintained by C.\\ Gelino, D.\\ Kirkpatrick, and A.\\ Burgasser.}.  This proper motion data may be put to a number of uses. Taken with measured radial velocities and distances, it can yield all three components of velocity (U,V,W). Using reduced proper motion diagrams it can be used as an approximate measure of distance. However we have no radial velocities for these objects. These proper motion measurements can however also be used to help identify objects as members of a star cluster or members of a moving group via the moving cluster method.  Our proper motion data is discussed and listed in section 2 of this paper. From the proper motion measurements, we find 5 new wide common proper motion binaries/multiple systems. We also identify 8 high velocity L dwarfs, which are discussed in section 4. We suggest that these L dwarfs are probably old and belong to the thick disc or halo population of the galaxy. This in turn suggests that they are likely to be very faint stars rather than brown dwarfs. Finally in section 5 we identify a number of L and T dwarfs that may be members of the Hyades, Pleiades and Ursa Major moving groups . ",
        "conclusions": "We have measured the proper motions for 143 dwarfs from the Dwarf Archive. From these measurements, we find 4 new common proper motion wide binary or multiple systems. We also identify 8 high velocity dwarfs i.e. dwarfs with tangential velocities $\\geq$ 100 kms$^{-1}$. These dwarfs also have bluer than average \\textit{J}-\\textit{K} colours. We argue that these are probably thick disc objects with an age of order 10 Gyr. We estimate their luminosities which are $\\approx$10$^{-4}$L/L$_{\\odot}$. This suggests that they are probably very low luminosity stars rather than brown dwarfs. If so, they may be some of the dimmest stars found to date. Finally we have found 15 L dwarfs that are potential members of the Hyades moving group, 5 that are potential members of the Ursa Major moving group and 5 that are potential members of the Pleiades moving group. The next obvious step towards confirming membership of these groups is to measure parallaxes for these dwarfs. Parallaxes will allow accurate distances to be used to compare with the moving group distance. \\textit{Spitzer} 3.5, 4.49, 5.73 and 7.87 micron magnitudes will also be valuable for a fuller understanding of the high velocity metal poor dwarfs."
    },
    "0710/0710.3359_arXiv.txt": {
        "abstract": "The feasibility of a mean-field dynamo in nonhelical turbulence with superimposed linear shear is studied numerically in elongated shearing boxes. Exponential growth of magnetic field at scales much larger than the outer scale of the turbulence is found. The charateristic scale of the field is $\\lB\\propto S^{-1/2}$ and growth rate is $\\gamma\\propto S$, where $S$ is the shearing rate. This newly discovered shear dynamo effect potentially represents a very generic mechanism for generating large-scale magnetic fields in a broad class of astrophysical systems with spatially coherent mean flows. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.1440_arXiv.txt": {
        "abstract": "% The UV excess shown by elliptical galaxies in their spectra is believed to be caused by evolved low-mass stars, in particular sdB stars. The stellar system most similar to the ellipticals for age and metallicity, in which it is possible to resolve these stars, is the bulge of our Galaxy. sdB star candidates were observed in the color magnitude diagram of a bulge region by Zoccali et al.\\ (2003). The follow-up spectroscopic analysis of these stars confirmed that most of these stars are bulge sdBs, while some candidates turned out to be disk sdBs or cool stars.  Both spectroscopic and photometric data and a spectral library are used to construct the integrated spectrum of the observed bulge region from the UV to the optical: the stars in the color magnitude diagram are associated to the library spectra, on the basis of their evolutionary status and temperature. The total integrated spectrum is obtained as the sum of the spectra associated to the color magnitude diagram. The comparison of the obtained integrated spectrum with old single stellar population synthetic spectra calculated by Bruzual \\& Charlot~(2003) agrees with age and metallicity of the bulge found by previous work. The bulge integrated spectrum shows only a very weak UV excess, but a too strict selection of the sample of the sdB star candidates in the color magnitude diagram and the exclusion of post-Asymptotic Giant Branch stars could have influenced the result. ",
        "introduction": "\\label{sec:intro}  The UV excess that elliptical galaxies and bulge of spiral galaxies show in their spectra at $\\lambda$ shorter than 2300~\\AA~ was one of the most puzzling discoveries in the last 30 years, since these stellar systems are believed to be old and metal rich, without young and massive stars emitting most of their flux at short wavelength.  It is now widely accepted that this UV emission is caused by evolved low mass stars, in particular Extreme Horizontal Branch stars (EHB), called also sdB stars from their spectral classification. These stars are faint in the optical wavelength range and with the current instrumentation it is impossible to resolve them in the nearest galaxies. The stellar system most similar to the ellipticals for age and metallicity in which it is possible to resolve sdB stars is the bulge of our Galaxy. A sample of sdBs star candidates was observed in the Galactic bulge by Zoccali et al.~(2003)\\nocite{zoccali03}, by means of $V$ and $I$ photometry of the region MW05 from the ESO Imaging Survey (EIS\\footnote{\\tt http://www.eso.org/science/eis/}, the observations were taken with the Wide Field Imager, WFI@2.2m). These stars could be either highly reddened sdBs or cooler stars affected by lower reddening. A follow-up spectroscopic analysis of these stars has been necessary and observations at the Very Large Telescope (VLT) telescope were obtained. The data reduction and the comparison of the obtained spectra with models of hot evolved stars confirmed indeed that most of these stars are bulge sdBs, while some candidates turned out to be disk sdBs or cool stars (for more details, see Busso et al.~2005\\nocite{busso05}).  To be sure that the observed bulge region was not peculiar, other bulge fields were searched for sdB candidates: EIS photometric data of the bulge fields MW07 and MW08 were reduced and analyzed, finding that sdB star candidates are present also in these fields.  This work presents the procedure adopted (following the recipe as in Santos et al.~1995) to construct the integrated spectrum of the bulge region MW05. ",
        "conclusions": ""
    },
    "0710/0710.5780_arXiv.txt": {
        "abstract": "We report high-resolution spectroscopy of 125 field stars previously observed as part of the Sloan Digital Sky Survey and its program for Galactic studies, the Sloan Extension for Galactic Understanding and Exploration (SEGUE). These spectra are used to measure radial velocities and to derive atmospheric parameters, which we compare with those reported by the SEGUE Stellar Parameter Pipeline (SSPP). The SSPP obtains estimates of these quantities based on SDSS $ugriz$ photometry and low-resolution ($R \\sim 2000$) spectroscopy. For F- and G-type stars observed with high signal-to-noise ratios ($S/N$), we empirically determine the typical random uncertainties in the radial velocities, effective temperatures, surface gravities, and metallicities delivered by the SSPP to be 2.4 km s$^{-1}$, 130 K (2.2 \\%), 0.21 dex, and 0.11 dex, respectively, with systematic uncertainties of a similar magnitude in the effective temperatures and metallicities. We estimate random errors for lower $S/N$ spectra based on numerical simulations. ",
        "introduction": "Starting from the sixth public data release (DR-6; Adelman-McCarthy et al. 2007), the Sloan Digital Sky Survey (SDSS) provides estimates of the atmospheric parameters for a subset of the stars observed spectroscopically in the survey (those in the approximate range of temperature $4500 \\le T_{\\rm eff} \\le 7500$~K). Following completion of the main survey (SDSS-I), the SDSS instrumentation has been devoted to several programs, including SEGUE: Sloan Extension for Galactic Understanding and Exploration, a massive survey of the stellar content of the Milky Way. Collectively, the suite of computer programs employed to determine atmospheric parameters from SEGUE data is known as the SEGUE Stellar Parameter Pipeline (SSPP). Because each of the public data releases of the SDSS includes and supersedes previous releases, DR-6 also includes atmospheric parameters for archival stellar observations in SDSS-I. These stellar parameters are derived by a series of methods, some of which consider purely spectroscopic information (continuum-normalized spectra), solely photometry (available in the survey's $ugriz$ system for all targets), or a combination of photometry and spectroscopy. Paper I in this series describes the SSPP in detail (Lee et al. 2007a). Paper II compares the predictions of the SSPP radial velocities and atmospheric parameters with likely members of Galactic globular and open clusters (Lee et al. 2007b).  The SDSS uses a CCD camera (Gunn et al. 1998) on a dedicated 2.5m telescope (Gunn et al. 2006) at Apache Point Observatory, New Mexico, to obtain images in five broad optical bands ($ugriz$; Fukugita et al.~1996) over approximately 10,000~deg$^2$ of the high Galactic latitude sky. The survey data-processing software measures the properties of each detected object in the imaging data in all five bands, and determines and applies both astrometric and photometric calibrations (Lupton et al. 2001; Pier et al. 2003; Ivezi\\'c et al.~2004). Photometric calibration is provided by simultaneous observations with a 20-inch telescope at the same site (Hogg et al.~2001; Smith et al.~2002; Stoughton et al.~2002; Tucker et al.~2006). A technical summary is provided by York et al. (2000).  SDSS-I and the ongoing SEGUE survey have already built the largest-ever catalog of stars in the Milky Way. To date, this includes photometry in five bands for over 200 million stars and spectroscopy for nearly 300,000 stars (Adelman-McCarthy et al. 2007). The SDSS spectrographs deliver a resolving power $\\lambda/$FWHM $\\sim 2000$ over the wavelength range 380-900 nm. Data reduction is fully automated, and the SSPP employs the final products from the SDSS pipeline as input to produce atmospheric parameters (effective temperature, surface gravity, and metallicity) for stars with spectral types A, F, G, and K. The best results are obtained for F- and G-type stars spanning the effective temperature range $5000 < T_{\\rm eff} < 7000$~K.  The quality of the SSPP atmospheric parameters is evaluated using different approaches, as already described in Paper I: comparing with previously published spectral libraries, well-studied open and globular clusters, and with high-resolution observations of field stars. Existing spectral libraries are useful in order to evaluate and calibrate the SSPP methods that rely on spectroscopy alone. Allende Prieto et al. (2006) employed the low-resolution Indo-US library (Valdes et al. 2004), and high-resolution spectra from the Elodie library (Prugniel \\& Soubiran 2001) and the S$^4$N archive (Allende Prieto et al. 2004). Because the $ugriz$ system was introduced with the SDSS, the stars included in existing spectral libraries lack photometry in this system. In addition, these are relatively bright stars, typically with $V<14$ mag, brighter than the bright magnitude limit of the SDSS imaging. The bright magnitude limit for the SDSS is set by the saturation threshold of the detectors at the sidereal driftscan rate of the survey. Obtaining data for these brighter stars would require special-purpose observations with a very different instrument configuration, which would call into question their value as calibration observations for the otherwise homogeneous imaging survey.  Star clusters provide stringent tests of the SSPP, as the same metallicity should be derived for stars that explore wide ranges of masses and luminosities. Paper II in this series examines SSPP results for likely members of clusters included in DR-6. One cannot choose clusters with any given metallicity, but has to take what is provided by nature and accessible from Apache Point. Furthermore, the effective temperatures and surface gravities for the members of any given cluster are very strongly correlated, depending on age and chemical composition. This leads to a patchy coverage of the parameter space. Field stars, on the other hand, can be chosen to provide better coverage and, therefore, naturally complement the clusters. Among the stars spectroscopically observed with SDSS, those in the range $14 < V <16.5$ mag can be observed at high spectral resolution with large-aperture telescopes and modest integration times. Due to the vast size of the SDSS stellar sample, these stars can be selected to more uniformly cover the parameter space of stellar properties, and have the additional benefit that photometry is already available for them in the SDSS native system.  This paper, the third in the SSPP series, is devoted to the analysis of 125 SDSS stars newly observed at high-resolution with the Hobby-Eberly, Keck, and Subaru telescopes. Section 2 describes the sample selection and the observations. The determination of radial velocities and atmospheric parameters, based on these observations, are discussed in \\S 3 and \\S 4, respectively. Section 5 describes the results for several well-known standard stars observed with the Hobby-Eberly Telescope. Section 6 compares the parameters derived from high-resolution spectroscopy with those from the SSPP. Section 7 describes numerical experiments that explore how the parameters degrade at lower signal-to-noise ratios. Our conclusions are summarized in \\S 8. ",
        "conclusions": "We report on an analysis of high-resolution spectroscopic observations of a sample of stars previously observed with the SDSS instrumentation as part of SDSS-I or SEGUE. These new data are used to derive radial velocities and atmospheric parameters, and to scrutinize the performance of the SSPP Pipeline described in Paper I in this series. The sample we have examined includes 81 stars observed with the HET-HRS, 25 stars obtained with Keck-ESI, 11 stars observed with Keck-HIRES, and 9 stars from Subaru-HDS.  Through a comparison with external spectroscopic libraries, and by employing multiple methods of analysis for the HET sample, we estimate that our reference radial velocities are accurate to 1.6 km s$^{-1}$. Our values for the stellar atmospheric parameters, effective temperature, surface gravity, and metallicity, are accurate to 1.5 \\% ($\\sim 90$~K), 0.13~dex and 0.05~dex, respectively. These figures are derived from the comparison with the parameters for nearby stars in the S$^4$N catalog, but we find they are still valid for the moderately high $S/N$ of the HET spectra. Using the HET sample to benchmark the SSPP, subtracting in quadrature the uncertainties in the results for the former, we conclude that the SDSS/SEGUE radial velocities are typically accurate to 2.4 km s$^{-1}$ for high signal-to-noise SDSS spectra ($S/N > 50/1$). A similar comparison of the atmospheric parameters returned by the SSPP with those obtained from HET spectra leads to the conclusion that the SSPP effective temperatures, surface gravities, and metallicities for bright targets show random errors of 2.2\\% ($\\sim 130$ K), 0.21 dex, and 0.11 dex, respectively. Systematic offsets of a similar size are detected for the effective temperatures and metallicities. We evaluate the expected random uncertainties as a function of $S/N$ by repeating the analysis after introducing noise in the SDSS spectra. More extended tests are underway and will be reported elsewhere.  Our study also finds that the internal uncertainties delivered by the SSPP for both radial velocities and atmospheric parameters need to be systematically increased by a factor of $2-3$ in order to be consistent with our derived external errors. The uncertainties in the average SSPP atmospheric parameters are simply derived as the standard error of the mean for a Normal distribution from the multiple techniques applied to any particular target. The fact that many methods share the same spectroscopic indicators (e.g. Balmer lines or SDSS color indices to gauge $T_{\\rm eff}$), and models (e.g. Kurucz's model atmospheres) may cause unaccounted correlations that result in underestimated uncertainties.  The validation and calibration of the SSPP is an ongoing project. Several additional open and globular clusters have recently had data obtained with SDSS instrumentation, and will be considered in future papers. A sample of up to several hundred very low-metallicity stars from SDSS/SEGUE is presently being observed with the HET, which we will add to our calibration sample. Additional stars of intermediate metallicity, and with hotter and cooler temperatures than considered in the present work, will be added to our calibration sample based on observations with a number of large-aperture telescopes. Our goal is to produce an SSPP validation catalog for on the order of 500 stars, which will be used to refine and adjust the individual parameter estimation techniques employed by the SSPP, and thus establish a definitive atmospheric parameter estimation scale for application to the large (and growing) SDSS/SEGUE stellar samples, as well as to other future surveys."
    },
    "0710/0710.5922_arXiv.txt": {
        "abstract": "\\noindent We derive upper limits on the ratio $\\fgrbccsn (z)\\equiv \\rgrb(z) /\\rccsn(z)\\, \\equiv\\fgrbccsn(0)(1+z)^\\alpha$, the ratio of the rate, $\\rgrb$, of long-duration Gamma Ray Bursts (GRBs) to the rate, $\\rccsn$, of core-collapse supernovae (CCSNe) in the Universe ($z$ being the cosmological redshift and $\\alpha\\geq 0$), by using the upper limit on the diffuse TeV--PeV neutrino background given by the AMANDA-II experiment in the South Pole, under the assumption that GRBs are sources of TeV--PeV neutrinos produced from decay of charged pions produced in $p\\gamma$ interaction of protons accelerated to ultrahigh energies at internal shocks within GRB jets. For the assumed ``concordance model'' of cosmic star formation rate, $\\rsf$, with $\\rccsn (z) \\propto \\rsf (z)$, our conservative upper limits are $\\fgrbccsn(0)\\leq 5.0\\times10^{-3}$ for $\\alpha=0$, and $\\fgrbccsn(0)\\leq 1.1\\times10^{-3}$ for $\\alpha=2$, for example. These limits are already comparable to (and, for $\\alpha\\geq 1$, already more restrictive than) the current upper limit on this ratio inferred from other astronomical considerations, thus providing a useful independent probe of and constraint on the CCSN-GRB connection. Non-detection of a diffuse TeV--PeV neutrino background by the up-coming IceCube detector in the South pole after three years of operation, for example, will bring down the upper limit on $\\fgrbccsn (0)$ to below few $\\times10^{-5}$ level, while a detection will confirm the hypothesis of proton acceleration to ultrahigh energies in GRBs and will potentially also yield the true rate of occurrence of these events in the Universe. ",
        "introduction": "} Detection of supernova (SN) features in the afterglow spectra of several long duration (typically $>2\\s$) Gamma Ray Bursts (GRBs) in the past one decade has provided strong support to the hypothesis that a significant fraction, if not all, of the long duration GRBs arise from collapse of massive stars; see, e.g., Refs.~\\cite{Meszaros_araa02,Woosley-Bloom_araa06,DellaValle06} for recent reviews. The observed SN features in the GRB afterglow spectra are similar to those usually associated with core-collapse supernovae (CCSNe) of Type Ib/c (see, e.g., ~\\cite{Woosley-Heger06,DellaValle06}). The total energy (corrected for beaming) in keV--MeV gamma rays emitted by typical long-duration GRBs is of order $10^{51}\\erg$, which is roughly the same as the total explosion energy seen in typical CCSNe, although there exists considerable diversity in the energetics of both the SN and the GRB components in the SN-GRB associations observed so far. In particular, the estimated explosion energies of the SNe associated with the GRBs observed so far seem to be somewhat larger than those of normal SNe, leading to this ``special'' class of SNe being sometimes referred to as ``hypernovae''.  The broad class of observational results on SN-GRB associations can be understood within the context of the ``collapsar'' model~\\cite{collapsar_model_refs} in terms of a simple phenomenological picture (see, e.g., ~\\cite{Woosley-Zhang_astroph_0701320}) in which the core-collapse of a massive Wolf-Rayet star gives rise to two kinds of outflows emanating from the central regions inside the collapsed star: (a) a narrowly collimated and highly relativistic jet that is responsible for the GRB activity, the jet being driven, for example, by a rapidly rotating and accreting black hole formed at the center in the core-collapse process, and (b) a more wide-angled, quasi-spherical and non-relativistic (or at best sub-relativistic) outflow that goes to blow up the star and gives rise to the supernova. The energies channeled into these two components may in general vary independently, which may explain the diversity of energetics in the observed SN-GRB associations. Actually, depending on the energy contained in it the ``GRB jet'' may or may not be able to penetrate through the stellar material and emerge outside. Indeed, the fact that the SN-GRB associations observed so far involve CCSNe of Type Ib/c, but not of Type II, may be due to the inability of the GRB-causing jet to penetrate through the relatively larger amount of outer stellar material in the case of Type II SN as compared to that in SNe of Type Ib/c~\\cite{fn1}. Considering various factors that may govern the energy channeled into the GRB-causing jet, such as the mass and rotation rate of the black hole, accretion efficiency, efficiency of conversion of accretion energy into collimated relativistic outflow, and so on, Woosley and Zhang~\\cite{Woosley-Zhang_astroph_0701320} have obtained a rough lower limit of $\\sim10^{48}\\erg/\\s$ for the power required for the jet to be able to emerge from the star. This is consistent with the energetics of the GRB components of the SN-GRB associations observed so far.  While SN-GRB associations strongly support the stellar core-collapse origin of most long-duration GRBs, clearly, not all core-collapse events may result in a GRB --- the latter depends on whether or not the core-collapse event actually results in a ``central engine'' (a rotating black hole fed by an accretion disk in the above mentioned phenomenological picture, for example) that is capable of driving the required collimated relativistic outflow. In other words, while every long-duration GRB would be expected to be accompanied by a core-collapse supernova~\\cite{fn2}, the reverse is not true in general.  What fraction of all stellar core-collapse events in the universe produce GRBs? Methods based on astronomical observations generally indicate the ratio between the cosmic GRB rate and the cosmic Type Ib/c SN rate, $\\fgrbsnIbc$, to be in the range $\\sim 10^{-3}$ -- $10^{-2}$ for a wide variety of different assumptions on various relevant parameters such as those that characterize the cosmic star formation rate (SFR), initial mass function (IMF) of stars, masses of Type Ib/c SN progenitors, the luminosity function of GRBs, the beaming factor of GRBs (associated with the fact that individual GRB emissions are highly non-isotropic and confined to narrowly collimated jets covering only a small fraction of the sky), and so on; see, for example, \\cite{Guetta-DellaValle_06,Bissaldi_etal_07} and references therein. The dominant uncertainty in the estimate of $\\fgrbsnIbc$ comes from the uncertainties in the estimates of the local GRB rate and the average GRB beaming factor. However, irrespective of the exact value of the ratio $\\fgrbsnIbc$, it is clear that this ratio is significantly less than unity. This indicates that, apart from just being sufficiently massive stars, the GRB progenitors may need to satisfy additional special conditions. For example, it has been suggested~\\cite{Woosley-Bloom_araa06} that the degree of rotation of the central iron core of the collapsing star and the metalicity of the progenitor star may play crucial roles in producing a GRB.  In this paper, we discuss an alternative probe of the cosmic GRB rate that uses the predicted high energy (TeV--PeV) diffuse neutrino background produced by GRBs and the experimental upper limit on high energy diffuse neutrino background given by the AMANDA-II experiment in the South Pole~\\cite{amanda_II_limit}. Existence of a high (TeV--PeV) energy diffuse GRB neutrino background (DGRBNuB) due to $p\\gamma$ interactions of (ultra)high energy protons accelerated within GRB sources is a generic prediction~\\cite{Waxman-Bahcall_grbnu} in most currently popular models of GRBs. This DGRBNuB is subject to being probed by the currently operating and up-coming large volume (kilometer scale) neutrino detectors such as IceCube~\\cite{icecube}, ANITA~\\cite{anita}, ANTARES~\\cite{antares}, for example. Since neutrinos, unlike electromagnetic radiation, can travel un-hindered from the furthest cosmological distances, the DGRBNuB automatically includes the contributions from all GRBs in the Universe. Thus, an analysis of the DGRBNuB is likely to provide a good picture of the true rate of occurrence of these events in the Universe. Indeed, as we show in this paper, the upper limits on $\\fgrbccsn (0)$, the ratio of the local (i.e., redshift $z=0$) GRB to CCSN rates, derived here from the consideration of DGRBNuB, are, for a wide range of values of the relevant parameters, already more restrictive than the current upper limit on this ratio ($\\sim 2.5\\times10^{-3}$) inferred from other astronomical considerations~\\cite{Guetta-DellaValle_06,Bissaldi_etal_07}. Further, non-detection of a diffuse TeV--PeV neutrino background by the up-coming IceCube detector~\\cite{icecube} in the South Pole after three years of operation, for example, will imply upper limits on $\\fgrbccsn(0)$ at the level of few \\ $\\times10^{-5}$, while a detection of the DGRBNuB will provide strong support to the hypothesis of proton acceleration to ultrahigh energies within GRB jets.  Our use of the DGRBNuB in constraining the cosmic GRB rate is in the same spirit as efforts to constrain the cosmic star formation rate (and thereby the cosmic CCSN rate) by using the experimental upper limit (set by the Super-Kamiokande (SK) detector)~\\cite{SK_limit} on the predicted~\\cite{dsnub_original} low (few MeV) energy Diffuse Supernova Neutrino Background (DSNuB); see, for example, Refs.~\\cite{Ando-Sato_rev_04,Strigari_etal_05,Hopkins-Beacom_06}. Now that the cosmic SFR including its absolute normalization and thereby the cosmic CCSNe rate have got reasonably well determined by the recent high quality data from a variety of astronomical observations (see, e.g., ~\\cite{Hopkins-Beacom_06}) (which, by the way, predicts a DSNuB flux that is close to the SK upper limit, implying that the DSNuB is probably close to being detected in the near future), one can begin to think of using this SFR to constrain the ratio of the cosmic GRB rate to CCSNe rate by using the predicted DGRBNuB flux together with the recent upper limits on the diffuse high energy neutrino flux from neutrino telescopes.  We should emphasize here that the upper limits derived in this paper actually refer to the ratio of the rate of GRBs to that of {\\it all} CCSNe including those of Type Ib/c and Type II, although SN-GRB associations observed so far involve SNe of Type Ib/c only. It is known, however, that Type II SNe probably constitute as much as $\\sim$ 75\\% of all CCSNe; see, e.g., \\cite{Cappellaro_etal_07}. Thus, one can get the constraint on the GRB-to-SNIb/c ratio from the GRB-to-CCSNe ratio we obtain here by multiplying the latter by a factor of $\\sim 4$. Conversely, for later comparison, we shall take the ``observed'' value of the ratio $\\fgrbccsn(0)$ to be in the range $2.5\\times (10^{-4}$ -- $10^{-3})$~\\cite{Guetta-DellaValle_06,Bissaldi_etal_07}.  Below, we first briefly review the calculation of the DGRBNuB spectrum in section \\ref{sec:DGRBNuB_calc}. The resulting upper limits on $\\fgrbccsn$ obtained by comparing the DGRBNuB with the current upper limit from AMANDA-II experiment are discussed in section \\ref{sec:fgrbccsn_constraints} for various values of some of the relevant GRB parameters. Finally, in section \\ref{sec:summary} we summarize the main results and conclude. ",
        "conclusions": "} In this paper we have attempted to derive upper limits on the fraction $\\fgrbccsn$ of all stellar core-collapse events that give rise to GRBs, by using the current experimental upper limit on the high energy (TeV -- PeV) diffuse neutrino background given by the AMANDA-II experiment in the South Pole, under the assumption that GRBs are sources of such high energy neutrinos. High energy neutrinos are predicted to be produced within GRB jets through photopion production by protons and subsequent decay of the charged pions, provided protons are accelerated to ultrahigh energies at the internal shocks within GRB jets. In our calculation we have allowed for a possible evolution of the cosmic GRB rate relative to star formation rate. For a wide range of values of various parameters, the upper limits on $\\fgrbccsn(0)$ derived here from the AMANDA-II results are already more restrictive than the upper limit on this ratio inferred from other astronomical considerations, thus providing a useful independent probe of and constraint on the CCSN-GRB connection. The closeness of the upper limits on $\\fgrbccsn(0)$ derived here (in particular for the case of enhanced evolution of the GRB rate relative to the star formation rate at high redshifts) to the lower limit on this ratio inferred from various astronomical considerations seems to indicate that the predicted DGRBNuB flux should be detectable by the upcoming detectors such as IceCube which will have significantly improved sensitivity over that of AMANDA-II. On the other hand, non-detection of the DGRBNuB by the IceCube detector after three years of operation, for example, will give more stringent upper limits on $\\fgrbccsn$, but at the same time will also imply that either the values of $\\fgrbccsn$ inferred from direct astronomical observations have been significantly overestimated (which is possible, for example, due to incorrect estimates of the average GRB beaming factor) or that the assumption of proton acceleration to ultrahigh energies within GRB jets is invalid, or both of these. However, more precise determination of the distribution of some of the crucial GRB parameters such as the bulk flow Lorentz factor and variability timescale of the GRBs will be needed to reliably calculate the expected contribution of the GRBs to the high energy diffuse neutrino background, and thereby to determine the upper limits on $\\fgrbccsn$ more reliably. To conclude, then, the up-coming large volume neutrino telescopes hold immense promise of yielding significant information both on the nature of the fundamental physical process of particle acceleration in GRB sources as well as on the rate of occurrence of these events in the Universe.  One of us (PB) wishes to thank Nayantara Gupta for helpful clarifications."
    },
    "0710/0710.4029_arXiv.txt": {
        "abstract": "The cold dark matter (CDM) scenario generically predicts the existence of triaxial dark matter haloes which contain notable amounts of substructure. However, analytical halo models with smooth, spherically symmetric density profiles are routinely adopted in the modelling of light propagation effects through such objects. In this paper, we address the biases introduced by this procedure by comparing the surface mass densities of actual N-body haloes against the widely used analytical model suggested by Navarro, Frenk and White (1996) (NFW). We conduct our analysis in the redshift range of 0.0 - 1.5.  In cluster sized haloes, we find that triaxiality can cause scatter in the surface mass density of the haloes up to $\\sigma_+ = +60 \\%$ and $\\sigma_- = -70 \\%$, where the 1-$\\sigma$ limits are relative to the analytical NFW model given value. Subhaloes can increase this scatter to $\\sigma_+ = +70 \\%$ and $\\sigma_- = -80 \\%$. In galaxy sized haloes, the triaxial scatter can be as high as $\\sigma_+ = +80 \\%$ and $\\sigma_- = -70 \\%$, and with subhaloes the values can change to $\\sigma_+ = +40 \\%$ and $\\sigma_- = -80 \\%$.  We present an analytical model for the surface mass density scatter as a function of distance to the halo centre, halo redshift and halo mass. The analytical description enables one to investigate the reliability of results obtained with simplified halo models. Additionally, it provides the means to add simulated surface density scatter to analytical density profiles. As an example, we discuss the impact of our results on the calculation of microlensing optical depths for MACHOs in CDM haloes. ",
        "introduction": "The cold dark matter model, in which the non-baryonic part of the dark matter is assumed to consist of particles that were non-relativistic already at the time of decoupling, and that interact predominantly through gravity, has been very successful in explaining the formation of large-scale structures in the Universe \\citep[see e.g.][ for a review]{Primack}. In this scenario, both galaxies and galaxy clusters are hosted by CDM haloes, which formed hierarchically through mergers of smaller subunits.  Even though N-body simulations generically predict CDM haloes to be triaxial \\citep[e.g.][]{Jing & Suto} with substantial amounts of substructures left over from the merging process \\citep[e.g.][]{Moore et al.}, simplified halo models are often adopted in the modelling of light propagation through such objects. The most common approach is to treat dark matter haloes as spherical objects with smooth density profiles, usually either of the NFW \\citep{NFW} form, some generalization thereof \\citep{Zhao}, or that of a cored or singular isothermal sphere.  The light emitted from high-redshift objects such as quasars, supernovae, gamma-ray bursts, galaxies and galaxy clusters will typically have to pass through many dark matter haloes before reaching an observer on Earth. Several investigations have already indicated that smooth and/or spherical halo models may lead to incorrect results when treating the gravitational lensing effects associated with such foreground mass condensations \\citep[e.g][]{Bartelmann & Weiss,Dalal et al.,Oguri & Keeton,Hennawi et al.}  More realistic features like triaxiality and substructures can be included in gravitational lens calculations either by employing N-body simulations directly \\citep[e.g.][]{Bartelmann & Weiss,Seljak & Holz,Holopainen et al.} or by using analytical expressions for the halo shapes \\citep[e.g.][]{Kochanek,Golse & Kneib,Evans & Hunter,Chae} and subhalo properties \\citep[e.g.][]{Oguri b,Zackrisson & Riehm}. While N-body simulations often represent the safest choice, the approach is computationally demanding and does not always allow one to identify the features of the mass distribution responsible for a specific lensing effect. Methods which bring simple, analytical halo models into contact with the full phenomenology of the N-body simulations are therefore highly desirable.  In this paper, we focus on the projected mass density of CDM haloes as a function of distance from the halo centre. There are several situations in gravitational lensing when realistic estimates of the surface mass density (i.e. convergence) along a given line of sight through a dark halo may be important. Examples include the calculation of image separations in strong lensing by subhaloes located in the external potential of its host halo \\citep{Oguri b}, attempts to correct the luminosities of supernovae type Ia for the magnification by foreground haloes \\citep[e.g.][]{Gunnarsson} and estimates of the distribution of microlensing optical depths for high-redshift MACHOs \\citep[e.g.][]{Wyithe & Turner,Zackrisson & Riehm}. Other applications include the assessments of light propagation effects in models with non-zero coupling betweeen dark matter particles and photons \\citep[e.g.][]{Profumo & Sigurdson}.  Here, we use high-resolution, dissipationless N-body simulations of CDM haloes to investigate the errors in surface mass density introduced by treating these objects as spherical with smooth density profiles of the NFW type. Simple relations for the surface mass density error as a function of halo redshift and distance to the halo centre are presented, making it easy to investigate the reliablitiy of results obtained with simplified halo models.  On a related note, \\citet{KnebeWiessner} recently investigated the error introduced by spherically averaging an elliptical mass distribution. They found that for axis ratios typical for cosmological dark matter haloes, the variance in the local density can be as large as 50\\% in the outer parts. The current paper examines the problem of halo triaxiality from a slightly different point of view.  The N-body simulations used are described in Section 2. In Section 3, we describe the methods for extracting the halo sample. In Section 4, we compare the CDM surface mass densities obtained along random sightlines through the N-body haloes to the corresponding results obtained from smooth and spherical NFW models fitted to the same haloes. Section 5 presents a set of simple relations for the surface mass density errors introduced by this procedure as a function of distance to the halo centre and halo redshift. Section 6 discusses how these relations may be used in the context of optical depth estimates for MACHO microlensing. A number of caveats are discussed in Section 7. Section 8 summarizes our findings. ",
        "conclusions": "The results presented here do suffer from a number of shortcomings which should be pointed out. The simulations used are dissipationless. In reality, dark matter haloes contain baryons, and the dissipation and feedback associated with these will inevitably affect the overall potential of the system, and thereby the spatial distribution of the CDM. According to current models, baryonic cooling will increase the central density of the CDM \\citep[e.g.][]{Gnedin et al.} and also make the halo more spherical \\citep{Kazantzidis et al.}. The significance of these effects are, however, still difficult to predict reliably, as the gas dynamical simulations involved still suffer from so-called ``overmerging'' problems \\citep[e.g.][]{Balogh et al., Springel & Hernquist}.  When calculating the surface mass density profiles, we have moreover considered only the matter present within $r_\\mathrm{vir}$ of each halo, whereas simulations have shown that galaxy sized CDM haloes extend at least out to 2--3$r_\\mathrm{vir}$ \\citep{Prada et al.}. We restricted our analysis to $r_\\mathrm{vir}$ because it becomes increasingly demanding to separate halo particles from the background the further one wants to extend the analysis. Even in the presented case, we need to separate particles out to $\\sim 1.5 r_\\mathrm{vir}$ because the smoothed particles extend their influence inside the virial radius region even though they are positioned outside it. We tested our method out to 3$r_\\mathrm{vir}$, but the number counts of the halo particles at those distances are too low to produce reliable results. Our halo sample does not have the resolution needed for extending the analysis further than $r_\\mathrm{vir}$ safely.  The most significant limitation of this paper is the small number of haloes in our analysis. This is of course due to the limited resolution of the cosmological simulations we had access to. We would like to repeat our analysis with a more complete statistical sample of haloes, which would hopefully confirm our analytical description with smaller error bars.  We also note that our analytical description of the surface mass density is more reliable in the case in which subhaloes are {\\it excluded}. This is because subhaloes can introduce significant mass peaks to some radial bins. These peaks can lead to unwanted effects in the log-normal fitting procedure which is designed to handle relatively smooth and continuous mass distributions within a bin. Large subhalos can also disturb the NFW fits, at last in the low density regions. Thus, the use of the models which include subhaloes is discouraged."
    },
    "0710/0710.1110_arXiv.txt": {
        "abstract": "We examine the effects Lorentz violation on observations of cosmic microwave background radiation. In particular, we focus on changes in polarization caused by vacuum birefringence. We place stringent constraints on previously untested violations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.3862_arXiv.txt": {
        "abstract": "We study a large-scale instability in a sheared nonhelical turbulence that causes generation of large-scale vorticity. Three types of the background large-scale flows are considered, i.e., the Couette and Poiseuille flows in a small-scale homogeneous turbulence, and the \"log-linear\" velocity shear in an inhomogeneous turbulence. It is known that laminar plane Couette flow and antisymmetric mode of laminar plane Poiseuille flow are stable with respect to small perturbations for any Reynolds numbers. We demonstrate that in a small-scale turbulence under certain conditions the large-scale Couette and Poiseuille flows are unstable due to the large-scale instability. This instability causes formation of large-scale vortical structures stretched along the mean sheared velocity. The growth rate of the large-scale instability for the \"log-linear\" velocity shear is much larger than that for the Couette and Poiseuille background flows. We have found a turbulent analogue of the Tollmien-Schlichting waves in a small-scale sheared turbulence. A mechanism of excitation of turbulent Tollmien-Schlichting waves is associated with a combined effect of the turbulent Reynolds stress-induced generation of perturbations of the mean vorticity and the background sheared motions. These waves can be excited even in a plane Couette flow imposed on a small-scale turbulence when perturbations of mean velocity depend on three spatial coordinates. The energy of these waves is supplied by the small-scale sheared turbulence. ",
        "introduction": "Large-scale vortical structures are universal features observed in geophysical, astrophysical and laboratory flows (see, e.g., \\cite{L83,P87,C94,GLM97,T98,RAO98}). Formation of vortical structures is related to the Prandtl secondary flows (see, e.g., \\cite{P52,T56,P70,B87}). A lateral stretching (or ''skewing\") by an existing shear generates streamwise vorticity that results in formation of the first kind of the Prandtl secondary flows. In turbulent flow the large-scale vorticity is generated by the divergence of the Reynolds stresses. This mechanism determines the second kind of the Prandtl turbulent secondary flows \\cite{B87}.  The generation of large-scale vorticity in a homogeneous nonhelical turbulence with an imposed large-scale linear velocity shear has been recently studied in \\cite{EKR03}. Let us discuss a mechanism of this phenomenon. The equation for the mean vorticity ${\\bf W} = \\bec{\\nabla} {\\bf \\times} {\\bf U}$ read \\begin{eqnarray} {\\partial {\\bf W} \\over \\partial t} = \\bec{\\nabla} {\\bf \\times} ({\\bf U} {\\bf \\times} {\\bf W} + {\\bf F} - \\nu \\bec{\\nabla} {\\bf \\times} {\\bf W})  \\;, \\label{W10} \\end{eqnarray} where ${\\bf U}$ is the mean fluid velocity, ${\\bf F}_i = - \\nabla_j \\, \\langle u_i u_j \\rangle$ is the effective force caused by velocity fluctuations, ${\\bf u}$, and $ \\nu$ is the kinematic viscosity. The first term, ${\\bf U} {\\bf \\times} {\\bf W}$, in Eq.~(\\ref{W10}) determines laminar effects of the mean vorticity production caused by the sheared motions, while the effective force ${\\bf F}$ determines the turbulent effects on the mean fluid flow. Let us consider a simple large-scale linear velocity shear ${\\bf U}^{(s)} = (0, Sx, 0)$ imposed on the small-scale nonhelical turbulence. The equation for the perturbations of the mean vorticity, $\\tilde{\\bf W} = (\\tilde{W}_x(z), \\tilde{W}_y(z), 0)$, reads \\begin{eqnarray} {\\partial \\tilde{W}_x \\over \\partial t} &=& S \\, \\tilde{W}_y + \\nu_{_{T}} \\tilde{W}''_x \\;, \\label{E2}\\\\ {\\partial \\tilde{W}_y \\over \\partial t} &=& - \\beta_0 \\, S \\, l_0^2 \\, \\tilde{W}''_x + \\nu_{_{T}} \\tilde{W}''_y  \\;, \\label{E3} \\end{eqnarray} (see \\cite{EKR03}), where $\\tilde{W}'' = \\partial^2 \\tilde{W} /\\partial z^2$, $\\, \\nu_{_{T}}$ is the turbulent viscosity, $l_0$ is the maximum scale of turbulent motions and the parameter $\\beta_0$ is of the order of 1, and depends on the scaling exponent of the correlation time of the turbulent velocity field (see Sect. II). A solution of Eqs.~(\\ref{E2}) and~(\\ref{E3}) has the form $ \\propto \\exp(\\gamma t + i K_z z)$, where the growth rate of the large-scale instability is given by $\\gamma = \\sqrt{\\beta_0} \\, S \\, l_0 \\, K_z - \\nu_{_{T}} \\, K_z^2$ and $K_z$ is the wave number. The maximum growth rate of perturbations of the mean vorticity, $ \\gamma_{\\rm max} = \\beta_0 \\, (S \\, l_0)^2 / 4 \\nu_{_{T}}$, is attained at $ K_z = K_m = \\sqrt{\\beta_0} \\, S \\, l_0 /2 \\nu_{_{T}}$. This corresponds to the ratio $\\tilde{W}_y / \\tilde{W}_x = \\sqrt{\\beta_0} \\, l_0 \\, K_m \\approx S \\, \\tau_0$, where the time $\\tau_0 = l_{0} / u_0$ and $u_0$ is the characteristic turbulent velocity in the maximum scale $l_{0}$ of turbulent motions. Note that in a laminar flow this instability does not occur.  The mechanism of this instability is as follows (see \\cite{EKR03} for details). The first term, $S \\tilde{W}_y = ({\\bf W}^{(s)}\\cdot\\bec{\\nabla})~\\tilde{U}_x$, in Eq.~(\\ref{E2}) determines a ''skew-induced\" generation of perturbations of the mean vorticity $\\tilde{W}_x$ by stretching of the equilibrium mean vorticity ${\\bf W}^{(s)}= (0,0,S)$, where $\\tilde{\\bf U}$ are the perturbations of the mean velocity. In particular, the mean vorticity $\\tilde{W}_x {\\bf e}_x$ is generated from $\\tilde{W}_y {\\bf e}_y$ by equilibrium shear motions with the mean vorticity ${\\bf W}^{(s)}$, whereby $\\tilde{W}_x {\\bf e}_x \\propto ({\\bf W}^{(s)} \\cdot \\bec{\\nabla}) \\tilde{U}_x {\\bf e}_x \\propto \\tilde{W}_y {\\bf e}_y \\times {\\bf W}^{(s)} $. Here ${\\bf e}_x$, ${\\bf e}_y$ and ${\\bf e}_z$ are the unit vectors along $x$, $y$ and $z$ axes, respectively. On the other hand, the first term, $- \\beta_0 \\, S \\, l_0^2 \\, \\tilde{W}''_x$, in Eq.~(\\ref{E3}) determines a ''Reynolds stress-induced\" generation of perturbations of the mean vorticity $\\tilde{W}_y$ by the Reynolds stresses.  In particular, this term is determined by $ (\\bec{\\nabla} {\\bf \\times} {\\bf F})_y$. This implies that the component of the mean vorticity $\\tilde{W}_y {\\bf e}_y $ is generated by an effective anisotropic viscous term $ \\propto - l_0^2 \\, \\Delta \\, (\\tilde{W}_x {\\bf e}_x \\cdot \\bec{\\nabla}) \\, {U}^{(s)}(x) {\\bf e}_y \\propto - l_0^2 \\, S \\, \\tilde{W}''_x {\\bf e}_y .$ This instability is caused by a combined effect of the sheared motions (''skew-induced\" generation) and the ''Reynolds stress-induced\" generation of perturbations of the mean vorticity.  The mechanism for this large-scale instability in a sheared nonhelical homogeneous turbulence is different from that discussed in \\cite{MST83,KMT91,CMP94}, where the generation of large-scale vorticity in the helical turbulence occurs due to hydrodynamic alpha effect. The latter effect is associated with the hydrodynamic helicity of turbulent flow. In a nonhelical homogeneous turbulence this effect does not occur.  The large-scale instability in a nonhelical homogeneous turbulence has been studied in \\cite{EKR03} only for a simple case of unbounded turbulence with an imposed linear velocity shear and when the perturbations of the mean vorticity depend on one spatial variable $z$. In this study the theoretical approach proposed in \\cite{EKR03} is further developed and applied for comprehensive investigation of the large-scale instability for different situations with nonuniform shear, inhomogeneous turbulence and a more general form of the perturbations of the mean vorticity $\\tilde{\\bf W}({\\bf r})$ that depends on three spatial variables.  In the present study we consider three types of the background large-scale flows, i.e., the Couette flow (linear velocity shear) and Poiseuille flow (quadratic velocity shear) in a small-scale homogeneous turbulence, and the \"log-linear\" velocity shear in an inhomogeneous turbulence. We have derived new mean-field equations for perturbations of large-scale velocity which depend on three spatial coordinates in a small-scale sheared turbulence, for a nonuniform background large-scale velocity shear and for an arbitrary scaling of the correlation time $\\tau(k)$ of the turbulent velocity field.  The stability of the laminar Couette and Poiseuille flows in a problem of transition to turbulence has been studied in a number of publications (see, e.g., \\cite{DR81,SH01,CJJ03,BOH88,REM03,ESH07}, and references therein). It is known that laminar plane Couette flow and antisymmetric mode of laminar plane Poiseuille flow are stable with respect to small perturbations for any Reynolds numbers. A symmetric mode of laminar plane Poiseuille flow is stable when the Reynolds number is less than 5772 \\cite{CJJ03}. In laminar flows the Tollmien-Schlichting waves can be excited. The molecular viscosity plays a destabilizing role in laminar flows which promotes the excitation of the Tollmien-Schlichting waves (see, e.g., \\cite{SH01}). These waves are growing solutions of the Orr-Sommerfeld equation.  In the present study we have found a turbulent analogue of the Tollmien-Schlichting waves. These waves are excited by a small-scale sheared turbulence, i.e., by a combined effect of the turbulent Reynolds stress-induced generation of perturbations of the mean vorticity and the background sheared motions. The energy of these waves is supplied by the small-scale sheared turbulence. We demonstrate that the off-diagonal terms in the turbulent viscosity tensor play a crucial role in the excitation of the turbulent Tollmien-Schlichting waves. These waves can be excited even in a plane Couette flow imposed on a small-scale turbulence when perturbations of velocity depend on three spatial coordinates. When perturbations of large-scale velocity depend on one or two spatial coordinates the turbulent Tollmien-Schlichting waves can not be excited in a sheared turbulence. In the present study we show that the large-scale Couette and Poiseuille flows imposed on a small-scale turbulence can be unstable with respect to small perturbations. The critical effective Reynolds number (based on turbulent viscosity) required for the excitation of this large-scale instability, is of the order of 200.  This paper is organized as follows. In Sect. II the governing equations are formulated. In Sect. III we consider a homogeneous turbulence with a large-scale linear velocity shear (Couette flow), while in Sect. IV we study a homogeneous turbulence with a large-scale quadratic velocity shear (Poiseuille flow). In Sect. V we investigate formation of large-scale vortical structures in an inhomogeneous turbulence with an imposed nonuniform velocity shear. Finally, we draw conclusions in Sec.~VI. ",
        "conclusions": "In this study the theoretical approach proposed in \\cite{EKR03} is further developed and applied to investigate the large-scale instability in a nonhelical turbulence with a nonuniform shear and a more general form of the perturbations of the mean vorticity. In particular, we consider three types of the background large-scale sheared flows imposed on small-scale turbulence: Couette flow (linear velocity shear) and Poiseuille flow (quadratic velocity shear) in a small-scale homogeneous turbulence, and a more complicated nonuniform velocity shear with the logarithmic velocity profile near the boundaries matched with the linear shear velocity for the central part of the background flow. This nonuniform velocity shear is imposed on an inhomogeneous turbulence. The latter flow is typical for the atmospheric boundary layer.  We show that the large-scale Couette and Poiseuille flows imposed on a small-scale turbulence are unstable with respect to small perturbations due to the excitation of the large-scale instability. This instability causes generation of large-scale vorticity and formation of large-scale vortical structures. The size of the formed vortical structures in the direction of the background velocity shear is much larger than the sizes of the structures in the directions perpendicular to the velocity shear. Therefore, the large-scale structures formed during this instability are stretched along the mean sheared velocity. Increase of shear promotes the large-scale instability. The thresholds for the excitation of the large-scale instability in the value of shear and the aspect ratio of structures for Poiseuille background flow are larger than that for the Couette background flow. The growth rate of the large-scale instability for the inhomogeneous turbulence with the \"log-linear\" velocity shear is much larger than that for the Couette and Poiseuille background flows. The characteristic spatial and time scales for the instability are much larger than the characteristic turbulent scales. This justifies separation of scales which is required for the validity of the mean-field theory applied in the present study.  The large-scale instability results in excitation of the turbulent Tollmien-Schlichting waves. The  mechanism for the excitation of these waves is different from that for the Tollmien-Schlichting waves in laminar flows. In particular, the molecular viscosity plays a crucial role in the excitation of the Tollmien-Schlichting waves in laminar flows. Contrary, the turbulent Tollmien-Schlichting waves are excited by a combined effect of the turbulent Reynolds stress-induced generation of perturbations of the mean vorticity and the background sheared motions. The energy of these waves is supplied by the small-scale sheared turbulence, and the off-diagonal terms in the turbulent viscosity tensor play a crucial role in the excitation of the turbulent Tollmien-Schlichting waves.  Note that this study is principally different from the problems of transition to turbulence whereby the stability of the laminar Couette and Poiseuille flows are investigated (see, e.g., \\cite{DR81,SH01,CJJ03,BOH88,REM03,ESH07}, and references therein). Here we do not analyze a transition to turbulence. We study the large-scale instability caused by an effect of the small-scale anisotropic turbulence on the mean flow. This anisotropic turbulence is produced by an interaction of equilibrium large-scale Couette or Poiseuille flows with a small-scale isotropic background turbulence produced by, e.g., a steering force. The anisotropic velocity fluctuations are generated by tangling of the mean-velocity gradients with the velocity fluctuations of the background turbulence \\cite{EKR03,EKRZ02}.  The \"tangling\" mechanism is an universal phenomenon that was introduced in \\cite{W57,BH59} for a passive scalar and in \\cite{G60,M61} for a passive vector (magnetic field). The Reynolds stresses in a turbulent flow with a mean velocity shear is another example of tangling anisotropic fluctuations \\cite{L67}. For instance, these velocity fluctuations are anisotropic in the presence of shear and have a steeper spectrum $\\propto k^{-7/3}$ than, e.g., a Kolmogorov background turbulence (see, e.g., \\cite{L67,WC72,SV94,IY02,EKRZ02}). The anisotropic velocity fluctuations determine the effective force and the Reynolds stresses in Eq.~(\\ref{B15}). This is the reason for the new terms $\\propto \\beta_n \\, l_0^2$ appearing in Eqs.~(\\ref{AA1})-(\\ref{A2}).  The obtained results in this study may be of relevance in different turbulent astrophysical, geophysical and industrial flows. Turbulence with a large-scale velocity shear is a universal feature in astrophysics and geophysics. In particular, the analyzed effects may be important, e.g., in accretion disks, extragalactic clusters, merged protostellar and protogalactic clouds. Sheared motions between interacting clouds can cause an excitation of the large-scale instability which results in generation of the mean vorticity and formation of large-scale vortical structures (see, e.g., \\cite{P80,ZN83,C93}). Dust particles can be trapped by the vortical structures to enhance agglomeration of material and formation of particle clusters \\cite{BS95,BR98,EKR98,CH00,JAB04}.  The suggested mechanism can be used in the analysis of the flows associated with Prandtl's turbulent secondary flows (see, e.g., \\cite{P52,B87}). However, in this study we have investigated only simple physical mechanisms to describe an initial (linear) stage of the formation of vortical structures. The simple models considered in this study can only mimic the flows associated with turbulent secondary flows. Clearly, the comprehensive numerical simulations of the nonlinear problem are required for quantitative description of the turbulent secondary flows."
    },
    "0710/0710.2864_arXiv.txt": {
        "abstract": "We present results from two high--contrast imaging surveys that exploit a novel technique, L--band angular differential imaging. Our first survey targeted 21 young stars in the $\\beta$~Pic and Tuc--Hor moving groups with VLT/NACO reaching typical sensitivities  of $<$1~M$_{\\mathrm Jup}$ at $r>20$~AU. The statistical analysis of the null result demonstrates that the giant planet population is truncated at 30~AU or less (90\\% confidence level). Our second, on--going MMT/Clio survey utilizes the unique sensitivity achieved in the L--band for old planets to probe all M--dwarf stars within 6~pc. The proximity of these targets enables direct imaging of planets in orbits like Jupiter for the first time --- a key step for directly imaging giant planets. ",
        "introduction": "The study of exoplanetary systems is arguably the most rapidly developing field in modern astrophysics. Surprisingly, much progress has been made without directly imaging a single planet: radial velocity/microlensing and primary and secondary planet eclipses provide limited, but valuable insights. Direct imaging of planetary systems will have a fundamental impact on the field --- a single, low--resolution 3--5 $\\mu$m spectrum of a planet may carry more information than all existing Spitzer transit  photometry combined. As was the case in the search for the first brown dwarf, or for radial velocity and planet transit techniques, achieving the first firm detection is a very difficult and often frustrating challenge. But these investments paid off rapidly by opening whole new classes of objects for study. ",
        "conclusions": "We present results from a novel high--contrast imaging technique. Our NACO survey of 21 nearby young stars demonstrates that the giant planet population does not extend beyond 30 AU and suggests a cut-off at radius $<$ 15 AU. Most previous imaging surveys have not detected planets because they targeted young stars ($>$ 20 pc) forcing them to probe orbital radii $>$20~AU. Our ongoing MMT 6pc volume--limited survey is probing --- for the first time --- the massive giant planet population around the closest stars with orbital radii $>$3."
    },
    "0710/0710.5169_arXiv.txt": {
        "abstract": "We construct new models of black hole-neutron star binaries in quasiequilibrium circular orbits by solving Einstein's constraint equations in the conformal thin-sandwich decomposition together with the relativistic equations of hydrostationary equilibrium. We adopt maximal slicing, assume spatial conformal flatness, and impose equilibrium boundary conditions on an excision surface (i.e., the apparent horizon) to model the black hole. In our previous treatment we adopted a ``leading-order\" approximation for a parameter related to the black-hole spin in these boundary conditions to construct approximately nonspinning black holes. Here we improve on the models by computing the black hole's quasilocal spin angular momentum and setting it to zero. As before, we adopt a polytropic equation of state with adiabatic index $\\Gamma=2$ and assume the neutron star to be irrotational. In addition to recomputing several sequences for comparison with our earlier results, we study a wider range of neutron star masses and binary mass ratios. To locate the innermost stable circular orbit we search for turning points along both the binding energy and total angular momentum curves for these sequences. Unlike for our previous approximate boundary condition, these two minima now coincide. We also identify the formation of cusps on the neutron star surface, indicating the onset of tidal disruption. Comparing these two critical binary separations for different mass ratios and neutron star compactions we distinguish those regions that will lead to a tidal disruption of the neutron star from those that will result in the plunge into the black hole of a neutron star more or less intact, albeit distorted by tidal forces. ",
        "introduction": "Coalescing black hole-neutron star (BHNS) binaries, as well as other compact binaries composed of neutron stars and/or black holes, are among the most promising sources of gravitational waves for both ground-based \\cite{LIGO,GEO,TAMA,VIRGO} and space-based laser interferometers \\cite{LISA,DECIGO}. BHNS binary mergers are also candidate central engines of short-hard gamma-ray bursts (SGRBs) (see, e.g. \\cite{LeeR07} and references cited therein). The remnants of both BHNS binary mergers \\cite{FaberBST06,ShibaU0607} and binary neutron star mergers \\cite{ShibaT06,PriceR06,OechsJ06,HMNS} are feasible progenitors for SGRBs because both may result in black holes surrounded by hot, massive accretion disks with very little, if any, baryon contamination along the polar symmetry axis.  Motivated by these factors, considerable effort has gone into the study of BHNS binaries. Most approaches to date assume Newtonian gravity in either some or all aspects of the calculation (see, e.g.~\\cite{Chand69,Fishb73,LaiRS93,LaiW96,TanigN96,Shiba96,UryuE99,WiggiL00,IshiiSM05,Mille05} for quasiequilibrium calculations and \\cite{Mashh75,CarteL83,Marck83,LeeK99,Lee00,RosswSW04,KobayLPM04,RantsKLR07} for dynamical simulations). More recently, several groups have also studied BHNS binaries in a fully relativistic framework, both for quasiequilibrium models \\cite{Mille01,BaumgSS04,TanigBFS05,TanigBFS06,Grand06,TanigBFS07} and dynamical simulations \\cite{FaberBSTR06,FaberBST06,SopueSL06,LofflRA06,ShibaU0607}.  Our group has pursued a systematic approach to developing increasingly realistic models of BHNS binaries in quasiequilibrium circular orbits. Our first studies \\cite{BaumgSS04,TanigBFS05,FaberBSTR06,FaberBST06} assumed extreme mass ratios, i.e., black hole masses that are much greater than the neutron star mass. While this is a very natural first step from a computational point of view, binaries with comparable masses are much more interesting from the perspective of ground-based gravitational wave observations and for the launching of SGRBs. More recently we have therefore relaxed this assumption and have extended our results to the case of comparable-mass BHNS binaries \\cite{TanigBFS06,TanigBFS07}.  Specifically, in \\cite{TanigBFS07} (hereafter Paper I) we constructed quasiequilibrium models by solving Einstein's constraint equations in the conformal thin-sandwich formalism, assuming conformal flatness and maximal slicing, together with the relativistic equations of hydrostationary equilibrium. We accounted for the black hole by excising a coordinate sphere and imposing the equilibrium black-hole boundary conditions of Cook and Pfeiffer \\cite{CookP04}. This original version implemented a ``leading-order\" approximation to nonspinning black holes, which equates an otherwise undetermined spin parameter $\\Omega_r$ that appears in the boundary condition for the shift vector with the orbital angular velocity seen by an inertial observer at infinity, $\\Omega$. As for the original irrotational binary black hole models of \\cite{CookP04}, this condition does not lead to simultaneous turning points of the binding energy and the total angular momentum in constant-mass sequences in Paper I. Such simultaneous turning points are expected for those sequences if they are truly in quasiequilibrium \\cite{dMOmegadJ}. An improvement over this condition, namely to iterate over $\\Omega_r$ until the quasilocal spin angular momentum of the black hole vanishes, was suggested and implemented for binary black holes by \\cite{CaudiCGP06}.  In this paper we reconstruct quasiequilibrium models of BHNS binaries using the same techniques as in Paper I, but with the improved black hole spin angular velocity condition as suggested by \\cite{CaudiCGP06}. We then compute sequences of BHNS binaries in quasicircular orbits for a wider range of neutron star masses and binary mass ratios than in Paper I, focusing our attention on irrotational neutron stars orbiting nonspinning black holes. Here we focus only on the irrotational state for the neutron star because it is astrophysically considered to be more realistic in a BHNS binary \\cite{Kocha92,BildsC92,FaberBSTR06}. On the other hand, we will compute the case of spinning black holes in future work. As was the case for the irrotational black hole binaries constructed in \\cite{CaudiCGP06}, we find that this improved condition for the spin parameter of the black hole $\\Omega_r$ does lead to simultaneous turning points in the binding energy and the total angular momentum along constant-mass sequences.  The paper is organized as follows. We briefly review the basic equations in Section II. We present numerical results in Section III, and outline some qualitative considerations concerning the fate of BHNS binaries in Section IV. In Section V we summarize our findings. Throughout this paper we adopt geometrized units with $G=c=1$, where $G$ denotes the gravitational constant and $c$ the speed of light. Latin and Greek indices denote purely spatial and spacetime components, respectively. ",
        "conclusions": "We have constructed new quasiequilibrium configurations of black hole-neutron star binaries in general relativity. We have solved the Einstein constraint equations in the conformal thin-sandwich formalism coupled with the equations of relativistic hydrostationary equilibrium. In Paper I, we set the spin angular velocity parameter of the black hole equal to that of the orbital angular velocity in order to produce a nonspinning black hole in the ``leading-order\" approximation \\cite{CookP04}, while in this paper we compute this parameter by requiring the quasilocal spin angular momentum of the black hole to be zero \\cite{CaudiCGP06}. We have also improved the formulation of the gravitational field equations and obtained more accurate results than in Paper I.  As an indication of the improvements in these calculations, a post-Newtonian analysis predicts smaller binary eccentricities for these new BHNS models than for those computed in Paper I (\\cite{Will07}, compare \\cite{BertiIW07}). In \\cite{BertiIW07}, Berti {\\it et al.} fit numerical results for the binding energy and angular momentum of binaries in circular orbit to post-Newtonian expressions for binaries that are not necessarily in circular orbit. Deviations between the the two approaches then lead to non-zero eccentricities in the post-Newtonian expressions. These eccentricities are smaller for our new results than for those of Paper I. We also remark on another finding of \\cite{BertiIW07}, namely that for a given neutron star mass $\\bar{M}_{\\rm ADM,0}^{\\rm NS}$ and a given value of $\\Omega M_0$, the eccentricities in BHNS models, though small, are found to be larger than in binary neutron star models \\cite{TanigG0203}. This suggests a larger deviation from quasiequilibrium for BHNS binaries than binary neutron stars. But, for BHNS binaries with a mass ratio of $\\hat{q} = 5$, these parameters correspond to a larger binary separation than for binary neutron stars with a mass ratio of $\\hat{q} = 1$ (compare Eq.~(\\ref{eq:Kepler})). For similar numerical resources, this larger binary separation leads to a larger numerical error, which may explain the larger eccentricity found by \\cite{BertiIW07}, at least in part.  In addition to recomputing several sequences we presented in Paper I, we have constructed sequences for a wider range of neutron star masses and binary mass ratios, employing a $\\Gamma=2$ polytropic neutron-star equation of state throughout. We computed several constant-mass sequences, for various mass ratios and neutron star compactions, and searched for the appearance of a cusp at the neutron star surface -- indicating the onset of tidal disruption -- and turning points on the binding energy and angular momentum curves -- identifying the ISCO. We also included some qualitative fits that allow for a simple prediction of those binary parameters separating these two different outcomes of binary coalescence. Unlike in our earlier findings, we found simultaneous turning points along the binding energy and angular momentum quasiequilibrium curves."
    },
    "0710/0710.0785_arXiv.txt": {
        "abstract": "{\\it NeXT} (New X-ray Telescope) is the next Japanese X-ray astronomical satellite mission after the {\\it Suzaku} satellite. {\\it NeXT} aims to perform wide band imaging spectroscopy. Due to the successful development of a multilayer coated mirror, called a supermirror, {\\it NeXT} can focus X-rays in the energy range from 0.1~keV up to 80~keV. To cover this wide energy range, we are in the process of developing a hybrid X-ray camera, Wideband X-ray Imager (WXI) as a focal plane detector of the supermirror. The WXI consists of X-ray CCDs (SXI) and CdTe pixelized detectors (HXI), which cover the lower and higher X-ray energy bands of 0.1--80~keV, respectively. The X-ray CCDs of the SXI are stacked above the CdTe pixelized detectors of the HXI. The X-ray CCDs of the SXI detect soft X-rays below $\\sim 10$~keV and allow hard X-rays pass into the CdTe detectors of the HXI without loss. Thus, we have been developing a ``back-supportless CCD'' with a thick depletion layer, a thinned silicon wafer, and a back-supportless structure. In this paper, we report the development and performances of an evaluation model of CCD for the SXI, ``CCD-NeXT1''. We successfully fabricated two types of CCD-NeXT1, unthinned CCDs with 625-$\\mu$m thick wafer and 150-$\\mu \\rm m$ thick thinned CCDs. By omitting the polishing process when making the thinned CCDs, we confirmed that the polishing process does not impact the X-ray performance. In addition, we did not find significant differences in the X-ray performance between the two types of CCDs. The energy resolution and readout noise are $\\sim 140$~eV (FWHM) at 5.9~keV and $\\sim 5$ electrons (RMS), respectively. The estimated thickness of the depletion layer is  $\\sim 80 ~\\mu \\rm m$. The performances almost satisfy the requirements of the baseline plan of the SXI. ",
        "introduction": "{\\it NeXT} (New X-ray Telescope) is the sixth Japanese X-ray astronomical satellite mission, which is proposed to be launched around 2012. A Hard X-ray Telescope (HXT), supermirror onboard {\\it NeXT} (see Tawara {\\it et~al.} (2003) \\cite{tawara03} and references therein), has a large collecting area for X-rays in the energy from 0.1 to 80~keV. In particular, the HXT has a high reflectivity even in the hard X-ray band above 10~keV. Previous satellite missions have not had X-ray focusing optics capable of observations in this band. {\\it NeXT} is designed to be the first to perform imaging and spectroscopic observations in the energy band above 10~keV. In order to meet the energy range covered by the HXT, we have been developing a Wideband X-ray Imager (WXI).  The first successful space flight use of X-ray CCDs as photon counting and spectroscopic imagers was the SIS aboard $ASCA$ \\cite{byrke91}. Since then, X-ray CCDs have become standard focal plane detectors for X-ray telescopes in the X-ray energy band of 0.1--10~keV, and have been adopted as the principal detectors of recent X-ray observatories such as the ACIS of {\\it Chandra} \\cite{garmire03}, the EPIC of {\\it XMM-Newton} \\cite{struder01,turner01}, and the XIS of {\\it Suzaku} \\cite{koyama07} because X-ray CCDs have well balanced spectroscopic, imaging, and time resolution performances. However, to achieve a quantum efficiency of 10\\% for X-rays with an energy of 40~keV, a $\\sim 1000~{\\rm \\mu m}$ depletion layer is required, which is nearly impossible. A detector with a high-Z material is essential for observations in the hard X-ray band above 10--20~keV. On the other hand, the performances of imaging and spectroscopy below 10~keV of high-Z solid detectors such as CdTe detectors are poorer than those of X-ray CCDs, suggesting that a single detector cannot cover the entire 0.1--80~keV band with the best X-ray performance. Thus, we have been developing a hybrid camera, the Wide band X-ray Imager (WXI), which combines an X-ray CCD and a CdTe pixelized detector \\cite{Takahashi1999,tsuru01,tsuru04,takahashi04}. Holland (2003) \\cite{adholland03} has also reported the first laboratory demonstration of such a hybrid detector with a thinned X-ray CCD, which is operated in front of CZT detector.  The WXI consists of two sub-instruments; the Soft X-ray Imager (SXI) and the Hard X-ray Imager (HXI); overviews can be found in Tsuru {\\it et~al.} (2005) \\cite{tsuru05} and Takahashi {\\it et~al.} (2004) \\cite{takahashi04s}, respectively. The SXI and HXI are the upper and lower parts of the WXI, respectively. The SXI consists of X-ray CCDs with a thick depletion layer for the lower energy band below 10--20~keV. The HXI is based on CdTe pixelized detectors, which cover the hard X-rays above 10--20~keV.  We have developed a new type of CCD for the SXI, a ``Back-Supportless CCD'' (BS-CCD), in which the back supporting package under the imaging area of the CCD is removed. Most X-rays absorbed in the field-free region of the CCD are undetected and lost. Hence, we also removed the field-free region as much as possible. The BS-CCD of the SXI is placed over the CdTe pixelized detectors of the HXI. Soft X-rays are detected in the BS-CCD, while hard X-rays penetrate through the BS-CCD and are detected by the CdTe pixelized detectors. Thus, both soft and hard X-rays are detected without loss.    As previously reported by Tsuru {\\it et~al.} (2005) \\cite{tsuru05} in detail, we have been developing a BS-CCD for the SXI following two plans of a rather conservative ``baseline plan'' and an innovative ``goal plan'' in parallel. Table~\\ref{tab:SXIgoala} shows the specifications of the BS-CCD of the two plans along with those of the XIS \\cite{koyama07}, which is one of the most excellent X-ray CCDs currently in orbit.  In the goal plan, we realize a fully-depleted back-illuminated type of BS-CCD with a very thick depletion layer of $\\sim 200~{\\rm \\mu m}$ by adopting a p-channel device. We have already successfully fabricated test devices with a  full depletion layer, which was $\\sim$ 200-$\\mu{\\rm m}$ thick. The details and status of the developments are reported elsewhere \\cite{kamata04,takagi05,kamata06,matuura06,takagi06,tsuru06}.   In the baseline plan, we developed BS-CCDs based on a natural extension of our successful developments of CCD-CREST/CREST2 and MAXI-CCD in order to minimize the risk involved with their development \\cite{bamba01,tsunemi05,tomida00,miyata02}. We adopted a front-illuminated type of CCD mainly due to the manufacturing process \\cite{takagi05}. The thickness of the depletion layer is designed to be 70--80$~{\\rm \\mu m}$ or more. We have already successfully developed a small test model of BS-CCD and confirmed that the thinning processes of the wafer and the back-supportless structure do not degrade the performance \\cite{takagi05}. After the successfully developing the small test model, we constructed an evaluation model, ``CCD-NeXT1''. Following CCD-NeXT1, we will develop a flight model, ``CCD-NeXT2'', which matches the specifications of the SXI shown in Table~1. In this paper, we report the development and the performance of CCD-NeXT1\\footnote{Note that  part of the results reported herein have already been reported as a contributing paper of a SPIE conference \\cite{OzawaSPIE06_CCD-NeXT1}.}. ",
        "conclusions": "\\begin{itemize}  \\item Following the successful development of the test models of BS-CCD for SXI onboard the {\\it NeXT} satellite, we developed an evaluation model, CCD-NeXT1. We fabricated two types of CCD-NeXT1. One was a type of un-thinned CCD with 625-$\\mu$m thick wafers. The other type was a thinned CCD with 150-$\\mu \\rm m$ thickness.  \\item We processed thinned CCD-NeXT1 devices by omitting the polishing process in the thinning process. The evaluation of the devices confirmed that omitting the polishing process did not impact on the X-ray performance.  \\item We did not observe a significant difference in the X-ray performance of the unthinned and the thinned devices. The energy resolution and the readout noise were $\\sim$140~eV (FWHM) at 5.9~keV and $\\sim$5 electrons (RMS), respectively. The detection efficiency of the $\\rm ^{109}Cd$ photons indicates that the depletion layer is $\\sim$ 80-${\\rm \\mu m}$ thick. This performance meets the requirements for the baseline plan of the SXI.  \\end{itemize}"
    },
    "0710/0710.2780_arXiv.txt": {
        "abstract": "}[2]{{\\footnotesize\\begin{center}ABSTRACT\\end{center} \\vspace{1mm}\\par#1\\par \\noindent {~}{\\it #2}}}  \\newcommand{\\TabCap}[2]{\\begin{center}\\parbox[t]{#1}{\\begin{center} \\small {\\spaceskip 2pt plus 1pt minus 1pt T a b l e} \\refstepcounter{table}\\thetable \\\\[2mm] \\footnotesize #2 \\end{center}}\\end{center}}  \\newcommand{\\TableSep}[2]{\\begin{table}[p]\\vspace{#1} \\TabCap{#2}\\end{table}}  \\newcommand{\\FigCap}[1]{\\footnotesize\\par\\noindent Fig.\\  % \\refstepcounter{figure}\\thefigure. #1\\par}  \\newcommand{\\TableFont}{\\footnotesize} \\newcommand{\\TableFontIt}{\\ttit} \\newcommand{\\SetTableFont}[1]{\\renewcommand{\\TableFont}{#1}}  \\newcommand{\\MakeTable}[4]{\\begin{table}[p]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newcommand{\\MakeTableTop}[4]{\\begin{table}[t]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newcommand{\\MakeTableSep}[4]{\\begin{table}[p]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}} \\newcommand{\\TabCapp}[2]{\\begin{center}\\parbox[t]{#1}{\\centerline{ \\small {\\spaceskip 2pt plus 1pt minus 1pt T a b l e} \\refstepcounter{table}\\thetable} \\vskip2mm \\centerline{\\footnotesize #2}} \\vskip3mm \\end{center}}  \\newcommand{\\MakeTableSepp}[4]{\\begin{table}[p]\\TabCapp{#2}{#3}\\vspace*{-.7cm} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newfont{\\bb}{ptmbi8t at 12pt} \\newfont{\\bbb}{cmbxti10} \\newfont{\\bbbb}{cmbxti10 at 9pt} \\newcommand{\\uprule}{\\rule{0pt}{2.5ex}} \\newcommand{\\douprule}{\\rule[-2ex]{0pt}{4.5ex}} \\newcommand{\\dorule}{\\rule[-2ex]{0pt}{2ex}} \\def\\thefootnote{\\fnsymbol{footnote}}  \\newenvironment{references}% { \\footnotesize \\frenchspacing \\renewcommand{\\thesection}{} \\renewcommand{\\in}{{\\rm in }} \\renewcommand{\\AA}{Astron.\\ Astrophys.} \\newcommand{\\AAS}{Astron.~Astrophys.~Suppl.~Ser.} \\newcommand{\\ApJ}{Astrophys.\\ J.} \\newcommand{\\ApJS}{Astrophys.\\ J.~Suppl.~Ser.} \\newcommand{\\ApJL}{Astrophys.\\ J.~Letters} \\newcommand{\\AJ}{Astron.\\ J.} \\newcommand{\\IBVS}{IBVS} \\newcommand{\\PASP}{P.A.S.P.} \\newcommand{\\Acta}{Acta Astron.} \\newcommand{\\MNRAS}{MNRAS} \\renewcommand{\\and}{{\\rm and }} {Period--luminosity (PL) relations of variable red giants in the Large (LMC) and Small Magellanic Clouds (SMC) are presented. The PL diagrams are plotted in three planes: $\\log P$--$K_S$, $\\log P$--$W_{JK}$, and $\\log P$--$W_I$, where $W_{JK}$ and $W_I$ are reddening free Wesenheit indices. Fourteen {\\it PL} sequences are distinguishable, and some of them consist of three closely spaced ridges. Each of the sequences is fitted with a linear or quadratic function. The similarities and differences between the {\\it PL} relations in both galaxies are discussed for four types of red giant variability: OGLE Small Amplitude Red Giants (OSARGs), Miras and Semiregular Variables (SRVs), Long Secondary Periods (LSPs) and ellipsoidal variables.  We propose a new method of separating OSARGs from non-variable stars and SRVs. The method employs the position in the reddening-free {\\it PL} diagrams and the characteristic period ratios of these multiperiodic variables. The {\\it PL} relations for the LMC OSARG are compared with the calculated relations for RGB models along isochrones of relevant ages and metallicities. We also compare measured periods and amplitudes of the OSARGs with predictions based on the relations valid for less luminous solar-like pulsators.  Miras and SRVs seem to follow {\\it PL} relation of the same slopes in the LMC and SMC, while for LSP and ellipsoidal variables slopes in both galaxies are different. The {\\it PL} sequences defined by LSP variables and binary systems overlap in the whole range of analyzed wavebands. We put forward new arguments for the binary star scenario as an explanation of the LSP variability and elaborate on it further. The measured pulsation to orbital period ratio implies nearly constant ratio of the star radius to orbital distance, $R/A\\approx0.4$, as we find. Combined effect of tidal friction and mass loss enhanced by the low-mass companion may explain why such a value is preferred.}{Stars: AGB and post-AGB -- Stars: late-type -- Stars: oscillations -- Magellanic Clouds} ",
        "introduction": "The first attempt to determine period--luminosity (PL) relation for long period variables (LPVs) was made by Gerasimovi\\v{c} (1928), who noticed that Mira stars with longer periods are on average fainter at visual wavelengths. This result has been confirmed by subsequent studies (\\eg Wilson and Merrill 1942, Osvalds and Risley 1961, Clayton and Feast 1969), however the scatter of this period--luminosity dependence turned out to be very large.  First tight {\\it PL} relation for LPVs was discovered for Mira stars at near infrared (NIR) wavebands (Glass and Lloyd Evans 1981). This {\\it PL} law, based on only 11 Miras in the LMC, was refined by extensive studies of Feast \\etal (1989) and Hughes and Wood (1990). The second, parallel {\\it PL} sequence, occupied by semiregular variables (SRVs), was identified by Wood and Sebo (1996). This sequence was shifted relative to the Miras' ridge toward shorter periods by a factor of two.  However, the subject of {\\it PL} distribution of LPVs has progressed rapidly over recent years when large microlensing surveys (MACHO, OGLE, EROS, MOA) published long-term photometry of huge number of stars. Complex structure of the {\\it PL} distribution was demonstrated for the first time by Cook \\etal (1997), who published {\\it PL} diagram for variable stars detected during the MACHO survey in the LMC. A series of three or four {\\it PL} sequences defined by LPVs can be distinguished in that diagram.  Sharper picture was presented by Wood \\etal (1999), who distinguished and described five {\\it PL} sequences (denoted as A--E) in the period--Wesenheit index plane. Wood (2000) showed similar distribution in the $\\log P$--$K$ diagram. These results were then confirmed by many studies based on observations originated in various sources (Cioni \\etal 2001, 2003, Noda \\etal 2002, Lebzelter \\etal 2002, Ita \\etal 2004, Groenewegen 2004, Fraser \\etal 2005).  Kiss and Bedding (2003, 2004) used OGLE data to reveal new features in the {\\it PL} distribution. They noticed that sequence B consists of two closely spaced parallel ridges (Ita \\etal 2004 denoted the additional sequence as C$'$). Below the tip of the red giant branch (TRGB) Kiss and Bedding (2003) found three sequences shifted in $\\log P$ relative to stars brighter than TRGB. It was the definitive proof that stars in the first ascent Red Giant Branch (RGB) pulsate similarly to objects being in the Asymptotic Giant Branch (AGB) phase.  The Optical Gravitational Lensing Experiment (OGLE) collected unprecedented amount of photometric data of stars in the Large and Small Magellanic Clouds. Both galaxies have been constantly monitored since 1997 and at present time this is the best and longest available photometric dataset for analyzing huge number of variable red giants. Our studies on LPVs resulted in many discoveries, regarding also the {\\it PL} relations.  Soszy{\\'n}ski \\etal (2004a) showed that OGLE Small Amplitude Red Giants (OS\\-ARGs) constitute separate class of variable stars, with different structure in the {\\it PL} plane than ``classical'' SRVs and Miras. We indicated two previously overlooked {\\it PL} relations -- the longest (${\\rm a}_1$) and the shortest (${\\rm a}_4$) period sequences followed by AGB OSARGs. We also suggested a method of empirical division between RGB and AGB OSARGs fainter than TRGB.  Red giants revealing ellipsoidal modulation caused by binarity were analyzed by Soszy{\\'n}ski \\etal (2004b). It was shown that, if true orbital periods are considered, the {\\it PL} relation of ellipsoidal variables (sequence~E) is a direct continuation of sequence~D occupied by mysterious Long Secondary Period (LSP) variables. This is a hint that the LSP phenomenon may be related to binarity, but taking into account available radial velocity measurements, the secondary component usually must be a low mass object, possibly former planet. This idea was supported by Soszy{\\'n}ski (2007) who discovered in some LSP variables ellipsoidal-like and eclipsing-like modulations with periods equal to LSPs.  In Soszy{\\'n}ski \\etal (2005) we again increased complexity of the PL distribution of LPVs. Each of the NIR {\\it PL} sequences C$'$, C and D in the LMC (occupied by SRV, Miras and LSP variables) split into two separate ridges in the period -- optical Wesenheit index plane, what corresponds to the spectral division into oxygen-rich (O-rich) and carbon-rich (C-rich) AGB stars. Thus, we found a new photometric method of distinguishing between these two populations.  In this paper we describe in details the {\\it PL} relations of variable red giants in both Magellanic Clouds. We show new details in the {\\it PL} plane and compare {\\it PL} distribution in the LMC and SMC. The paper is organized as follows. Section~2 gives details of the observations and data reduction. In Section~3 the {\\it PL} relations are presented with a description of their derivation. A discussion about four types of red giant variability -- OSARGs, Miras/SRVs, LSPs and ellipsoidal modulation -- is given in Sections~4--7. Section~8 summarizes and concludes the paper. ",
        "conclusions": "In this paper we showed the most complex structure of the {\\it PL} distribution presented so far. The Wood's five ridges turn out to be an overlap of fourteen sequences (if consider closely spaced {\\it PL} relations of OSARGs, the number of sequences exceeds twenty). In order to help recognizing {\\it PL} relations with published {\\it PL} laws we provide Table~3 with appropriate identifications.  Diagrams employing period ratios (similar to the Petersen diagram) were used as a tool for discriminating OGLE Small Amplitude Red Giants (OSARGs). We compared three sequences of the {\\it PL} relations for the LMC OSARGs in the RGB phase with the calculated relation for first three radial modes using isochrone calculations. We found an essential agreement with our knowledge about metallicities and ages of red giants population in the LMC. However, there are also discrepancies which may suggest the necessity of refinement of the knowledge and/or stellar models. This is a potential application of the OSARGs.  \\MakeTableTop{l@{\\hspace{4pt}}c@{\\hspace{4pt}} l@{\\hspace{8pt}} c@{\\hspace{8pt}} c@{\\hspace{8pt}} c@{\\hspace{8pt}}}{12.5cm} {Labels of the {\\it PL} relations in this and previous papers} {\\hline \\noalign{\\vskip3pt} \\multicolumn{3}{c}{this paper} & Wood \\etal 1999 & Kiss and Bedding 2003 & Ita \\etal 2004 \\\\ \\noalign{\\vskip3pt} \\hline \\noalign{\\vskip3pt} &     & ${\\rm b}_1$ &   & $R_1$ &                  \\\\ & RGB & ${\\rm b}_2$ & B & $R_2$ & $B^-$ \\\\ &     & ${\\rm b}_3$ & A & $R_3$ & $A^-$ \\\\ OSARGs &     & ${\\rm a}_1$ &   &                           &                  \\\\ &     & ${\\rm a}_2$ & B & 2O                        & $B^+$ \\\\[-1ex] & \\raisebox{1.5ex}{AGB} & ${\\rm a}_3$ & A & 3O      & $A^+$ \\\\ &     & ${\\rm a}_4$ &   &                           &                  \\\\ \\noalign{\\vskip3pt} \\hline \\noalign{\\vskip3pt} &  & C$_{\\rm O}$ & C & F  & C \\\\[-1ex] Miras    & \\raisebox{1.5ex}{O-rich} & C$'_{\\rm O}$ & B & 1O & C$'$ \\\\ and SRVs &  & C$_{\\rm C}$ & C & F  & C \\\\[-1ex] & \\raisebox{1.5ex}{C-rich} & C$'_{\\rm C}$ & B & 1O & C$'$ \\\\ \\noalign{\\vskip3pt} \\hline \\noalign{\\vskip3pt} & O-rich & D$_{\\rm O}$ & D & $L_2$ & D \\\\[-1ex] \\raisebox{1.5ex}{LSPs} & C-rich & D$_{\\rm C}$ & D & $L_2$ & D \\\\ \\noalign{\\vskip3pt} \\hline \\noalign{\\vskip3pt} \\multicolumn{2}{l}{Ell. and Ecl.} & E & E$^*$ & $L_1^*$ & E$^*$ \\\\ \\noalign{\\vskip3pt} \\hline \\multicolumn{6}{l}{$^*$ -- the sequence is shifted due to halving the orbital periods.} }  Most likely mechanism responsible for pulsation in the OSARGs is the stochastic excitation. Therefore we looked at these objects as solar-like pulsators. The range of periods agrees with predictions based on extrapolation of the relations found for much less luminous stars. The amplitudes are lower than predicted by the Kjeldsen and Bedding (1995) formula, where the the amplitude rises linearly with the luminosity-to-mass ratio. However, there are calculations predicting a slower amplitude rise. Testing theory of stochastic excitation is another possible application of OSARGs.  We would like to bring the reader's attention to NIR Wesenheit index ($W_{JK}$). The sequences in the period--$W_{JK}$ plane are generally better defined than those at $K_S$ magnitudes. Wesenheit index is a reddening independent quantity, so even heavy reddened Miras fall on the {\\it PL} sequence. Therefore, the period--$W_{JK}$ relations can be used as a distance indicators without correcting them for reddening. Moreover, O- and C-rich Miras and SRV obey very similar relations in the period--$W_{JK}$ plane, while at $K_S$ the {\\it PL} relations are significantly different for both populations. We conclude that $W_{JK}$ index can be a useful tool for studying LPVs.  Comparison of the {\\it PL} relations in NIR supports the binary star scenario as the explanation of the LSP variability. We further develop this scenario taking into account data on the ratios of the short period (pulsational) and LSP variability. For OSARGs these ratio is nearly constant which translates into nearly constant ratio of the stellar to orbital radius the value of about 0.4. We proposed that the small mass companion position is determined by the balance between mass loss and tidal effects. This scenario requires that the proximity of the companion enhances mass loss.  Moreover, it requires that a substantial fraction (majority) of red giants have such companions and that their mass cannot be too small.  \\vspace*{9pt}  \\Acknow{The paper was supported by the Foundation for Polish Science through the Homing Program and by MNiSW grants: 1P03D01130 and N20303032/4275.  This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation.}  \\vspace*{9pt}"
    },
    "0710/0710.5505_arXiv.txt": {
        "abstract": "We analyze a toy swiss-cheese cosmological model to study the averaging problem.  In our swiss-cheese model, the cheese is a spatially flat, matter only, Friedmann-Robertson-Walker solution (\\textit{i.e.,} the Einstein--de Sitter model), and the holes are constructed from a Lema\\^{\\i}tre-Tolman-Bondi solution of Einstein's equations. We study the propagation of photons in the swiss-cheese model, and find a phenomenological homogeneous model to describe observables. Following a fitting procedure based on light-cone averages, we find that the the expansion scalar is unaffected by the inhomogeneities (\\textit{i.e.,} the phenomenological homogeneous model is the cheese model). This is because of the spherical symmetry of the model; it is unclear whether the expansion scalar will be affected by non-spherical voids.  However, the light-cone average of the density as a function of redshift is affected by inhomogeneities.  The effect arises because, as the universe evolves,  a photon spends more and more time in the (large) voids than in the (thin) high-density structures.  The phenomenological homogeneous model describing the light-cone average of the density is similar to the $\\Lambda$CDM concordance model.  It is interesting that although the sole source in the swiss-cheese model is matter, the phenomenological homogeneous model behaves as if it has a dark-energy component.  Finally, we study how the equation of state of the phenomenological homogeneous model depends on the size of the inhomogeneities, and find that the equation-of-state parameters $w_{0}$ and $w_{a}$ follow a power-law dependence with a scaling exponent equal to unity. That is, the equation of state depends linearly on the distance the photon travels through voids. We conclude that within our toy model, the holes must have a present size of about $250$ Mpc to be able to mimic the concordance model. ",
        "introduction": "Most, if not all, observations are consistent with the cosmic concordance model, according to which  one-fourth of the present mass-energy of the universe is clustered and dominated by cold dark matter (CDM).  The remaining three-quarters is uniform and dominated by a fluid with negative pressure (dark energy, or $\\Lambda$).  While the standard $\\Lambda$CDM model seems capable of accounting for the observations, 95\\% of the mass-energy of the present universe is unknown. This is either a feature, and we are  presented with the opportunity of discovering the nature of dark matter and dark energy, or it is a bug, and nature might be different than described by the $\\Lambda$CDM model. Regardless, until such time as dark matter and dark energy are completely understood, it is useful to look for alternative cosmological models that fit the data.  One non-standard possibility is that there are large effects on the {\\it observed} expansion rate (and hence on other observables) due to the back-reaction of inhomogeneities in the universe (see, \\textit{e.g.,} Ref.\\ \\cite{kmr,notari,rasa, buchert} and references therein).  The basic idea is that all evidence for dark energy comes from the observational determinations of the expansion history of the universe.  Anything that affects the observed expansion history of the universe alters the determination of the parameters of dark energy; in the extreme it may remove the need for dark energy.  The ``safe'' consequence of the success of the concordance model is that the isotropic and homogeneous $\\Lambda$CDM model is a good  {\\it phenomenological} fit to the real inhomogeneous universe. And this is, in some sense, a verification of the cosmological principle:  the inhomogeneous universe can be described by means of an isotropic and  homogeneous solution.  However, this does not imply that a primary source of dark energy  exists, but only that it exists as far as the  phenomenological fit is concerned. For example, it is not straightforward that the universe is accelerating. If dark energy does not exist at a fundamental level, its presence in the concordance model would tell us that the pure-matter inhomogeneous model has been renormalized, from the phenomenological point of view  (luminosity-distance and redshift of photons), into a homogeneous $\\Lambda$CDM model.  The issue is the observational significance of the back-reaction of inhomogeneities. Our point of view is tied to our past light cone: we focus on the effects  of large-scale nonlinear inhomogeneities on observables such as the  luminosity-distance--redshift relation. We will not discuss averaged domain dynamics, even though if we think it is a crucial step in understanding how General Relativity effectively works in a  lumpy universe \\cite{ellis, buchert_new}.  Following this approach, we built in Ref.\\ \\cite{marra-sc} a particular swiss-cheese model, where the cheese consists of a spatially flat, matter only Friedmann-Robertson-Walker (FRW) solution and the holes are constructed out of a Lema\\^{i}tre-Tolman-Bondi  (LTB) solution of Einstein's equations. We attempted to find a model that was solvable and ``realistic'' (even if still toy), rather than finding a model with interesting volume-averaged dynamics. The model, however, will turn out to be useful to investigate light-cone averages.  It has been indeed shown that the LTB solution can be used to fit the observed $d_{L}(z)$  without the need of dark energy (for example, see Ref.\\ \\cite{alnes0602}).  To achieve this result, however, it is necessary to place the observer at the center of a rather large-scale underdensity. To overcome this  fine-tuning problem we built a swiss-cheese model with the observer in the cheese looking through a series of holes.  In Ref.\\ \\cite{marra-sc} we studied this model in detail and discussed the effects of large-scale nonlinear inhomogeneities on observables such as the luminosity-distance--redshift relation.  We found that inhomogeneities are able (at least partly) to mimic the effects of dark energy.  In this paper we will analyze the same swiss-cheese model through the fitting scheme developed by Ellis and Stoeger \\cite{ellis-f} in order to better understand how inhomogeneities renormalize the (matter only) swiss-cheese model allowing us to avoid a physical dark-energy component.  We think that this model fits well in that context and therefore we  might be able to shed some light on the important topics discussed there. We will propose a fitting procedure based on light-cone averages.  The paper is organized as follows: In Sec.\\ \\ref{model} we will specify the parameters of our swiss-cheese model and summarize the main results  obtained in Ref.\\ \\cite{marra-sc}. In Sec.\\ \\ref{fitti}, we develop our fitting procedure, and in  Sec.\\ \\ref{disco} we discuss our results. Then, in Sec.\\ \\ref{dressing} we study the dependence of the best-fit  parameters on the size of the holes. Conclusions are given in Sec.\\ \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  The aim of this investigation was to understand the role of large-scale non-linear  cosmic inhomogeneities in the interpretation of observational data. We focused on an exact (if toy) solution, based on the Lema\\^{\\i}tre-Tolman-Bondi (LTB) model. This solution has been studied extensively in the literature  \\cite{alnes0607, notari-mansouri, alnes0602, celerier, mansouri, flanagan, rasanen, tomita, chung, nambu}. It has been shown  that it can be used to fit the observed $d_{L}(z)$  without the need of dark energy (for example in Ref.\\ \\cite{alnes0602}).  To achieve this result, however, it is necessary to place the observer at the center of a rather large-scale underdensity. To overcome this  fine-tuning problem we built a swiss-cheese model, placing the observer  in the cheese and having the observer look through the holes in the swiss-cheese as pictured in Fig.\\ \\ref{schizzo}.  In Sec.\\ \\ref{model} we defined the model and  described its dynamics:  it is a swiss-cheese model where the cheese is made of the usual FRW solution and the holes are made of a LTB solution. The voids inside the holes  are expanding faster than the cheese. We reported also the results for $d_{L}(z)$ obtained in Ref.\\ \\cite{marra-sc}, to which we refer the reader for a more thorough analysis. We found that redshift effects are suppressed because of a compensation effect due to spherical symmetry. However, we found interesting effects in the calculation of the angular distance: the evolution of the inhomogeneities bends the photon path compared to the FRW case. Therefore, inhomogeneities will be able (at least partly) to mimic the effects of dark energy.  After having analyzed the model from the observational point of view,  we set up in Section \\ref{fitti} the fitting problem in order to better  understand how inhomogeneities renormalize the matter swiss-cheese model allowing us to eschew a primary dark energy. We followed the scheme developed in Ref.\\ \\cite{ellis-f}, but modified in the way  to fit the phenomenological model to the swiss-cheese one. We chose a method that is intermediate between the fitting approach  and the averaging one: we fitted with respect to light-cone averages.  In particular, we focused on the expansion and the density. While the expansion behaved as in the FRW case because of the compensation  effect mentioned above, we found that the density behaved differently thanks to  its intensiveness to that compensation effect: a photon is spending more and more time in the (large) voids than in the (thin) high density  structures. This effect is not directly linked to the one giving us an  interesting $d_{A}$. The best fit we found for holes of $r_{h}=250$ Mpc is $w_{0}=-1.03$ and $w_{a}=2.19$; qualitatively similar to the concordance model.  The flow chart of Fig.\\ \\ref{schema} summarizes the results obtained. The insensitivity to the compensation effect made us think that a swiss cheese made of spherical symmetric holes and a swiss cheese without an exact spherical symmetry  would share the same light-cone averaged density.  Knowing the behavior of the density we are therefore able to know the one of the Hubble parameter that will be the one of the FRW solution with an phenomenological source characterized by the fit equation of state. In this way we can think to go beyond the main limitation of this model, that is, the assumption of spherical symmetry. From this point of view, the light-cone averaged density can be seen as a  tool in performing this step.  Summarizing: \\begin{itemize}  \\item We started with a swiss-cheese model based on spherically symmetric holes. A photon, during its journey through the swiss cheese, undergoes  a redshift which is not affected by inhomogeneities. However the photon is spending more and more time in the voids than in the structures. The lack of an effect is due to the the assumption of spherical symmetry. We focus on this because a photon spending most of its time in voids should have a different redshift history than a photon propagating in a homogeneous background.  \\item Assuming that the density is a quantity that does not heavily depend  on the assumption of spherical symmetry, we  tried to resolve the issue by focusing on  the density alone and getting from it the expansion (and therefore the  redshift history).  \\item This resulted in a swiss-cheese model with holes that effectively are not perfectly spherical. In this model the redshift history of a photon depends on the time passed inside the voids.  \\item In practice this means that we will use the phenomenological best-fit model  found, that is, we will use a model that behaves similarly to the concordance  model.  \\end{itemize}  Then, in Section \\ref{dressing} we studied how the equation of state of  a phenomenological model with only one effective source depends on the size of the inhomogeneity. We found that $w^{R}_{0}$ and $w^{R}_{a}$ follow a power-law dependence with the same scaling exponent which is equal to unity. That is, the equation of state depends linearly on the distance the photon travels through voids.  We finally asked which size of the holes will give us a phenomenological model able to mimic the concordance model.  We found that for $n=1$, that is for a holes of radius $r_{h}=250$ Mpc, we have $w_{0}=-1.03$ and $w_{a}=2.19$."
    },
    "0710/0710.5219_arXiv.txt": {
        "abstract": "{} {The analysis of near infrared spectropolarimetric data at the internetwork at different regions on the solar surface could offer constraints to reject current modeling of these quiet areas.} {We present spectro-polarimetric observations of very quiet regions for different values of the heliocentric angle for the Fe\\,{\\sc i} lines at $1.56$ $\\mu$m, from disc centre to positions close to the limb. The spatial resolution of the data is $0.7-1''$. We analyze direct observable properties of the Stokes profiles as the amplitude of circular and linear polarization as well as the total degree of polarization. Also the area and amplitude asymmetries are studied.} {We do not find any significant variation of the properties of the polarimetric signals with the heliocentric angle. This means that the magnetism of the solar internetwork remains the same regardless of  the position on the solar disc. This observational fact discards the possibility of modeling the internetwork as a Network-like scenario. The magnetic elements of internetwork areas seem to be isotropically distributed when observed at our spatial resolution.} {} ",
        "introduction": "The presence of magnetic fields in the solar internetwork was discovered more than 30 years ago \\citep{livingston_75,smithson_75}. Since that work, improvements on the instrumentation have allowed the observation of the full Stokes vector in those regions with a signal-to-noise ratio good enough to retrieve the magnetic field vector from the observational data. Nevertheless, the internetwork magnetic topology remains yet very  indistinct and shows a growing  complexity as we improve the quality of the data. \\cite{martin_87} observed that in her videomagnetographs at a spatial resolution of $3''$ the longitudinal component of internetwork magnetic fields was present everywhere in the solar disc. She immediately concluded that these magnetic structures should have very small scales and a very tangled geometry, giving as an example a scenario in which the magnetic fields in the internetwork would consist of a maze of small loops.  As spectro-polarimeters have made possible the precise detection of these signals, the efforts have concentrated on the study of the distribution of magnetic fields at disc centre. Different methodologies (mainly based upon the use of the Zeeman and Hanle effects) shed light on  different aspects of the internetwork magnetism. Concerning the Zeeman effect, the Stokes $Q$, $U$ and $V$ signals detected on the most widely used spectral lines (Fe\\,{\\sc i} at $1.5$ $\\mu$m and 630 nm) are very weak at the best spatial resolutions of 0.5-1$''$. Moreover, the Stokes $I$ profile seems to come from a field free atmosphere, as expected from weakly polarized media \\citep{jorge_99}. The filling factor of the magnetic elements retrieved from the analysis of these signals is always around 2 \\% \\citep{khomenko_03, jorge_ita_03, marian_spw4}. The rest of the resolution element would be filled with very weak magnetic fields. Another possibility is that the real element is filled with mixed polarity magnetic fields which would partially cancel out due to the lack of spatial resolution. The magnetic field strength distributions recovered for this 2 \\% of the resolution element appear to show a preference for magnetic fields around the equipartition field at photospheric heights and even weaker \\citep{khomenko_03, marian_spw4, julio07}. Additionally, \\cite{andres_07} have shown the first direct observational evidence of flux cancellation in the internetwork (note that their internetwork region is surrounded by a very enhanced network structure), showing that more than 95 \\% of the magnetic flux is cancelled in the resolution element ($\\sim 1''$). This means that the above-mentioned distributions could not give a complete vision of the internetwork magnetism. \\cite{andres_07} give an amount of $250$ G for the mean magnetic field in the resolution element (note that the values of the magnetic flux density obtained by means of the Zeeman effect are below 10 Mx/cm$^2$). This would mean that the magnetism of the internetwork could play an important role on the solar global magnetism. This has also been pointed out by works using the Hanle effect \\citep{javier_04}.  The study of regions in different positions of the solar disc could represent a strong constraint to reject models for the solar internetwork. Using the 630 nm lines with a spatial resolution of $\\sim 1''$, \\cite{lites_02} built a histogram of both circular and linear polarization signals in two quiet regions, one at disc centre and another at $\\mu=0.82$ (being $\\mu$ the cosine of the heliocentric angle). He did not observe significant linear polarization signals in neither of the two regions.  However, Figure 9 in \\cite{lites_02} shows that the histograms of circular polarization do not present any variation for those signals whose integrated signal is below 0.005. \\cite{meunier_98}, studying integrated polarization, and  \\cite{harvey_07} using magnetograms, show the presence of a horizontal component of the magnetic field everywhere in the solar disc. In this paper we present the first study of the internetwork at several positions on the solar disc using high quality $0.8''$ spectro-polarimetric data. ",
        "conclusions": "We have analyzed high quality spectro-polarimetric  data of the Fe\\,{\\sc i} lines at $1.56$ $\\mu$m in order to infer information about  the internetwork magnetism at different positions on the Sun's surface. The whole field of view presents significant signal, meaning that the magnetic fields pervade the observed areas.  We have found that the circular and linear polarization amplitudes do not have any clear dependence on the heliocentric angle. This fact goes against a Network-like scenario for the internetwork: quasi--vertical flux tubes cannot explain this observational result, nor in fact any field topology with a preferred orientation within the field-of-view. An isotropical distribution of magnetic fields, oriented in all directions in the whole field of view, is on the other hand expected to show this behaviour. \\cite{marian_07} found that at least 10-20 \\% of the magnetic flux in the internetwork is connected by low-lying loops. Consequently, the scenario proposed by \\cite{martin_87} of an internetwork characterized by a myriad of small loops is a very reasonably idea that is compatible with all the observational constraints presented in this work. Of course one can think of other scenarios that are compatible with the observations. \\citep{stenflo_87, rafa_04, javier_04} adopt turbulent internetwork magnetism, \\cite{jorge_00} proposes a micro-structuration of the atmosphere (MISMA) to explain all the magnetic phenomena on the solar surface. All of them are compatible with the presented results and our efforts should be headed towards finding more constraints to reject some of them and strengthen others.   The size of the magnetic structures can also be constrained by the information presented in this study. First, improving the spatial resolution, we do not see a global increase in the signals. \\cite{lites_04} found no increment of the magnetic flux density from $1\"$ to $0.6\"$. Recently \\cite{lites_07} computed a magnetic flux density of about $11$ Mx/cm$^2$ at 0.3$''$ using HINODE's data, which is compatible with the value of $10$ Mx/cm$^2$ found by \\cite{martin_87} at $3''$. This means that, either the magnetic field structures are already resolved at $\\sim 0.5''$ or we are very far from resolving them. The fact that the polarimetric signals do not vary along the solar surface would point towards very small structures as the responsibles for the internetwork magnetism. The size of these magnetic structures is something not yet constrained by the observations.  We have presented a study of high quality spectro-polarimetric data in different positions of the solar surface, from the disc centre towards $\\mu=0.28$. This is the first step of the study of the variation on the magnetism of the internetwork with the heliocentric angle. Much more work has to be done by retrieving physical parameters as the magnetic field strength vector, magnetic flux, etc. to really constrain the modeling of the internetwork and reject models that are not compatible with the results."
    },
    "0710/0710.4134_arXiv.txt": {
        "abstract": "Since 1999, a radial velocity survey of 179 red giant stars is ongoing at Lick Observatory with a one month cadence.  At present $\\sim$20$-$100 measurements have been collected per star with an accuracy of 5 to 8 m\\,s$^{-1}$. Of the stars monitored, 145 (80\\%) show radial velocity (RV) variations at a level $>$20 m\\,s$^{-1}$, of which 43 exhibit significant periodicities. Here, we investigate the mechanism causing the observed radial velocity variations. Firstly, we search for a correlation between the radial velocity amplitude and an intrinsic parameter of the star, in this case surface gravity ($\\log g$). Secondly, we investigate line profile variations and compare these with theoretical predictions. ",
        "introduction": "Since 1999, a radial velocity survey of 179 red giant stars is ongoing at Lick Observatory, using the 60 cm Coud\\'e Auxiliary Telescope (CAT) in conjunction with the Hamilton echelle spectrograph (R $\\approx$ 60\\,000).  These stars have been selected from the Hipparcos catalogue \\cite{esa1997}, based on the criteria described by \\cite{frink2001}. The selected stars are all brighter than 6~mag, are presumably single and have photometric variations $< 0.06$~mag in V.  The system with an iodine cell in the light path has been developed as described by \\cite{marcy1992} and \\cite{valenti1995}. With integration times of up to thirty minutes for the faintest stars ($m_{v}$ = 6 mag) we reach a signal to noise ratio of about $80-100$ at $\\lambda = 5500$ \\AA , yielding a radial velocity precision of $5-8$ m\\,s$^{-1}$.  The initial aim of the survey was to check whether red giants would be stable enough to serve as reference stars for astrometric observations with SIM/PlanetQuest \\cite{frink2001}. In \\cite{hekker2006} it is shown that a large fraction of the red giants in a specific part of the absolute magnitude vs. B-V colour diagram are stable to a level of 20 m\\,s$^{-1}$ and could be effectively searched for long period companions, as is required for astrometric reference stars. For other stars in the sample the radial velocity variations are larger, and for 43 stars these show significant periodicities. So far, sub-stellar companions have been announced for two stars from the present sample ($\\iota$ Dra \\cite{frink2002} and $\\beta$ Gem \\cite{reffert2006}).  Here, we investigate which physical mechanism causes the observed radial velocity variations. In cases for which we do not find a significant periodicity in the observed radial velocity variations, an intrinsic mechanism such as spots or pulsations, possibly multi-periodic, seems most likely. On the other hand, the periodic radial velocity variations can be caused by sub-stellar companion, an intrinsic mechanism, or by both these mechanisms simultaneously. In Section 2 we search for a relation between the amplitude of the radial velocity variations and an intrinsic parameter, i.e.~$\\log g$. In Section 3 we investigate line shape variations and compare these with theoretical predictions. Our conclusions are presented in Section 4. A more extended paper on this subject is submitted \\cite{hekker2007b}. ",
        "conclusions": "There exists a clear correlation between the half peak-to-peak values of the radial velocity variations and the surface gravity of the red giant stars. This is a strong indication that the observed radial velocity variations are caused by a mechanism intrinsic to the star.  Companions and an intrinsic mechanism might be present in stars with periodic radial velocity variations and $\\log g > 1.6$ dex, while for stars with $\\log g<1.6$ dex solely an intrinsic mechanism seems most likely.  We have investigated whether there is a relation between the radial velocity amplitude and the amplitude of the line depth residuals. From theory we find that the rotational velocity, intrinsic and equivalent line width largely influence the line depth residual and no clear correlation could be identified.  In the theoretical models temperature variations due to the pulsations is ignored, but these temperature variations influence the intrinsic and equivalent line width of the spectral line, which influences the line depth residuals. This might be an explanation why the theoretical models do not overlap with the observations."
    },
    "0710/0710.4072_arXiv.txt": {
        "abstract": "{We investigate correlations between the optical linear polarization position angle and the orientation of the host galaxy/extended emission of Type~1 and Type~2 Radio-Loud (RL) and Radio-Quiet (RQ) quasars. We have used high resolution Hubble Space Telescope ({\\it HST}) data and deconvolution process to obtain a good determination of the host galaxy orientation. With these new measurements and a compilation of data from the literature, we find a significant correlation between the polarization position angle and the position angle of the major axis of the host galaxy/extended emission. The correlation appears different for Type~1 and Type~2 objects and depends on the redshift of the source. Interpretations in the framework of the unification model are discussed.}    \\addkeyword{Quasars : general} \\addkeyword{Polarization}   \\begin{document} ",
        "introduction": " ",
        "conclusions": "We can summarize our results as follow : \\begin{asparaitem} \\item{} While Type~2 RL and RQ quasars are known to exhibit an anti-alignment between the major axis of hteir host/extended emission and their optical polarization, we find that an alignment is mostly observed for Type~1 quasars. \\item{} The redshift dependence of the alignment effect, and lack of correlation with the near-IR $PA_{host}$, suggest that it might be related to the rest-frame extended UV/blue emission of quasars. \\item{}We show that these observations can be interpreted in the framework of the unification model + a two component scattering model. \\end{asparaitem}"
    },
    "0710/0710.3378_arXiv.txt": {
        "abstract": "In order to prepare the analysis of COROT data, it has been decided to build a simple tool to simulate the expected light curves. This simulation tools takes into account both instrumental constraints and astrophysical inputs for the COROT targets. For example, granulation and magnetic activity signatures are simulated, as well as p modes, with the expected photon noise. However, the simulations rely sometimes on simple approach of these phenomenons, as the main goal of this tool is to prepare the analysis in the case of COROT data and not to perform the most realistic simulations of the different phenomenons. ",
        "introduction": "Simulating the data that a space instrument like COROT will provide might look presomptuous. Indeed, it is certainly, when comparing to previous comparable instruments like IPHIR or GOLF. These two examples show that the nominal behaviour of the instrument is not always reached, but this does not prevent this instrument to provide very interesting data. However, despite some technical problems, IPHIR and GOLF yielded a wealth of scientific results. Thus, what is the interest of simulating COROT data? How close to reality these simualtions will get? This might not be the most important fact as the preparation of these simulations will help us to prepare the analysis of real data and to be ready in case of unexpected technical behaviour of the instrument perturbating the data, or unexpected physical behaviour of the targets of the instrument. A consequence of that is that the simulation tool must include technical and physical aspects, making the task even more difficult. These aspects cover: photon noise, p modes excitation, granulation signal, stellar activity signal, orbital perturbations, stellar rotation...  The software presented here is freely available at:\\\\ {\\tt www.lesia.obspm.fr/$\\sim$corotswg/simulightcurve.html} ",
        "conclusions": "This simulator software will continue to evolve with time. As indicated above, intensity modulation due to starspots will be included, as well as other stellar or instrumental signals, as for example instrumental perturbations due to orbital vraiations. Moreover, this effort of simulation will not end with the delivery of first data but will be continued after that. The comparison with real data will allow to check for the validity of physical hypothesis used to simulate the different signals of astrophysics origin in the data. This shoud bring a great amount of information on our knowmedge of these often not well known phenomena, which stellar simulation is often derived from the solar case. In parallel, the simulation of instrumental components of the signal will be improved to help the interpretation of real data. All these reasons justify in our opinion the need for the simulation tool presented here."
    },
    "0710/0710.4244_arXiv.txt": {
        "abstract": "We summarise the mathematical foundation of the holographic method of measuring the reflector profile of an antenna or radio telescope.  In particular, we treat the case, where the signal source is located at a finite distance from the antenna under test, necessitating the inclusion of the so-called Fresnel field terms in the radiation integrals. We assume a ``full phase'' system with reference receiver to provide the reference phase.  We describe in some detail the hardware and software implementation of the system used for the holographic measurement of the 12m ALMA prototype submillimeter antennas. We include a description of the practicalities of a measurement and surface setting.  The results for both the VertexRSI and AEC (Alcatel-EIE-Consortium) prototype ALMA antennas are presented. ",
        "introduction": "\\label{intro}  \\PARstart{L}{arge} reflector antennas, as those used in radio astronomy and deep-space communication, generally are composed of a set of surface panels, supported on three or more points by a support structure, often called the backup structure. After assembly of the reflector it is necessary to accurately locate the panels onto the prescribed paraboloidal surface in order to obtain the maximum antenna gain. The fact that some antennas have a \"shaped\"contour is irrelevant for the purpose of our discussion. We are concerned with describing a method which allows us to derive the position of the individual panels in space and compute the necessary adjustments of their support points to obtain a continuous surface of a certain prescribed shape.  The analysis by Ruze \\cite{Ruze1966} of the influence of random errors in the reflector contour on the antenna gain indicates that the RMS error should be less than about one-sixteenth of the wavelength for acceptable performance. Under the assumption that the errors are small compared to a wavelength, randomly distributed with RMS value \\(\\epsilon \\), have a correlation length {\\bfseries c} which is much larger than the wavelength \\(\\lambda \\), and much smaller than the reflector diameter D, the relative decrease in aperture efficiency (or gain) can be expressed by the simple formula  \\begin{equation} \\frac{\\eta_A}{\\eta_{A0}} = \\exp\\left\\{-{\\left(\\frac{4\\pi\\epsilon}{\\lambda}\\right)}^2\\right\\}, \\label{eq:etaa} \\end{equation}  \\noindent{where} \\({{\\eta }_{A0}}\\) is the aperture efficiency of the perfect reflector. An error $\\epsilon$ of \\(\\lambda \\)/40 is required to limit the gain loss to 10 percent; with an error of \\(\\lambda \\)/16 the gain is decreased to about half of the maximum achievable.  The setting of the reflector panels at accuracies better than 100 \\(\\mu \\)m has required the development of measuring methods of hitherto unsurpassed accuracy. It should be noted that these measurements need to be done ``in the field'', which in the case of millimeter radio telescopes generally means the hostile environment of a high mountain site.  One versatile, and by now widely used method is normally called ``radio holography''. The method makes use of a well-known relationship in antenna theory: the far-field radiation pattern of a reflector antenna is the Fourier Transformation of the field distribution in the aperture plane of the antenna. Note that this relationship applies to the amplitude/phase distributions, not to the power pattern. Thus, if we can measure the radiation pattern, in amplitude and phase, we can derive by Fourier Transformation the amplitude and phase distribution in the antenna aperture plane with an acceptable spatial resolution. Bennett \\etal \\cite{Bennett1976} presented a sufficiently detailed analysis of this method to draw the attention of radio astronomers. Scott \\& Ryle \\cite{Scott1977} used the new Cambridge 5 km array to measure the shape of four of the eight antennas, using a celestial radio point source and the remaining antennas to provide the reference signal. Simulation algorithms were developed by Rahmat-Samii \\cite{RahmatSamii1985} and others, adding to the practicability of the method.  Using the giant water vapour maser at 22 GHz in Orion as a source Morris \\etal \\cite{Morris1988} achieved a measurement accuracy of 30 $\\mu$m and were able to set the surface of the IRAM 30-m millimeter telescope to an accuracy of better than 100 $\\mu$m.  \\pubidadjcol  Artificial satellites, radiating a beacon signal at a fixed frequency have also been used as farfield (${R_f} = \\frac{2 {D^2}}{\\lambda}$) signal sources. Extensive use has been made of synchronous communication satellites in the 11 GHz band \\cite{Godwin1986}, \\cite{Rochblatt1992}. These transmitters of course do not provide the range of elevation angles accessible with cosmic sources. Some satellites, notably the LES (Lincoln Experimental Satellite) 8 and 9, have been used for radio holography of millimeter telescopes \\cite{Baars1999}. They provided a signal at the high frequency of 37 GHz and with their geo-synchronous orbit moved over some 60 degrees in elevation angle. Unfortunately, both satellites are no longer available. Lacking a sufficiently strong source in the farfield, we have to take recourse to using an earth-bound transmitter. In practice these will be located at a distance of several hundreds of meters to a few kilometers and be at an elevation angle of less than 10 degrees. Clearly, these are in the nearfield of the antenna, requiring significant corrections to the received signals.  In particular, the phase front of the incoming waves will not be plane and it contains higher order terms in the radial coordinate of the antenna aperture. These must be corrected before the Fourier transformation can be applied. We treat these corrections in detail in this paper.  Successful measurements on short ranges have been reported for the University of Texas millimeter telescope \\cite{Mayer1983}, the IRAM 30-m telecope \\cite{Morris1988b}, the JCMT \\cite{Hills2002} and the ASTE antenna of NAOJ \\cite{Ezawa2000}.  ALMA (Atacama Large Millimeter Array) is a new large aperture synthesis array for submillimeter astronomy consisting of 50 high accuracy antennas of 12 m diameter. The instrument is under construction at 5000 m altitude in the Atacama desert of northern Chile. ALMA is a collaboration of North America and Europe with participation of Japan.  Two prototype antennas were procured and erected at the site of the Very Large Array of NRAO in New Mexico. The results of an extensive evaluation program of these antennas has been presented by Mangum \\etal \\cite{Mangum2006}. The reflector surface accuracy was specified at 20-25 $\\mu$m, requiring a measurement method with an accuracy of 10 $\\mu$m or better. This was achieved with a near-field holographic system using a transmitter at a wavelength of 3 mm and at a distance of only 315 m from the antennas at an elevation angle of 9 degrees. Here we describe these measurements in some detail. ",
        "conclusions": "\\label{conclusions}  We have successfully performed a holographic measurement and consecutive panel setting of the reflectors of the two ALMA prototype antennas to an accuracy of better than 20 $\\mu$m.  Our estimated measurement accuracy is approximately 5 $\\mu $m. The data collection and analysis software packages are easy to use and provide quick results of the measurements, directly usable for a panel adjustment setting. We consider this system suitable for the routine setting of the ALMA production antennas to the goal of 20 $\\mu$m accuracy in an acceptable time span. Modern survey equipment enables contractors to deliver reflectors with an accuracy of 50-60 $\\mu$m without undue cost. Although the holography system can easily start with a much larger error, in the former case it is feasible to reach the specification with only one panel setting based on holography. We note that these measurements, being performed at one elevation angle only, do not provide information on the gravitationally induced deformation as function of elevation angle.  In summary:  \\begin{enumerate} \\item The holography system has functioned according to specification and has enabled us to measure the surface of the antenna reflector with a repeatability of better than $10\\mu$m. \\item As shown in Figs.~\\ref{fig:holosVertex} and \\ref{fig:holosAEC}, we have set both antenna surfaces to and accuracy of 16-17 $\\mu$m RMS. This will provide an aperture efficiency of about 65 percent of that of a perfect reflector at the highest observing frequency of 950 GHz. \\item The small differences in the surface maps obtained over several days of measurement are consistent with the measurement repeatability and at best marginally significant. If taken at face value, they indicate that the deformations of the reflector under varying wind and temperature influence are fully consistent with, and probably well within, the specification.  This excellent behaviour over time is more important than the actual achieved surface setting. We stopped iteration of the settings after having achieved the goal of less than 20 $\\mu$m. \\item Further information on the performance of the ALMA Prototype Antennas can be found in \\cite{Mangum2006}. \\end{enumerate}  \\appendices"
    },
    "0710/0710.2654_arXiv.txt": {
        "abstract": "In this letter we briefly describe the first results of our numerical study on the possibility of magnetic origin of relativistic jets of long duration gamma ray bursters within the collapsar scenario. We track the collapse of massive rotating stars onto a rotating central black hole using axisymmetric general relativistic magnetohydrodynamic code that utilizes a realistic equation of state of stellar matter, takes into account the cooling associated with emission of neutrinos, and the energy losses due to dissociation of nuclei. The neutrino heating is not included. We describe the solution for one particular model where the progenitor star has magnetic field $B=3\\times10^{10}$G. The solution exhibits strong explosion driven by the Poynting-dominated jets whose power exceeds $2\\times10^{51}\\,\\mbox{erg/s}$. The jets originate mainly from the black hole and they are powered via the Blandford-Znajek mechanism. ",
        "introduction": "\\label{introduction} The phenomenon of Gamma Ray Burst (GRB) has been puzzling astrophysicists for many years since its discovery in 1970s~\\cite{KSO73,MGI74}. The recent identification of long duration GRBs with supernovae (see Della Valle 2006, and Woosley \\& Bloom 2006 for full review) means that we are dealing with enormous amount of energy, $10^{51}-10^{52}\\mbox{erg}$, released within a very short time, 2-100 seconds, in the form of highly relativistic collimated outflow \\cite{P05}.  Most of the current GRB studies are focused on the physics associated with production of gamma rays in such flows and their interaction with the interstellar medium or the stellar wind of the supernova progenitor. However, the central question in the problem of GRBs is undoubtedly the nature of their central engines. These powerful jets have to be produced as a result of stellar collapse, most likely by the relativistic object, neutron star or black hole (BH), formed in the center, and make their way through the massive star unscathed, remaining well collimated and highly relativistic.  The most popular model of central engine is based on the ``failed supernova'' scenario of stellar collapse, or ``collapsar'', where the iron core of progenitor star forms a BH \\cite{W93}.  If the progenitor is non-rotating then its collapse is likely to continue in a ``silent'' manner until the whole star is swallowed by the BH.  If, however, the specific angular momentum in the equatorial part of stellar envelope exceeds that of the last stable orbit of the BH then the collapse becomes highly anisotropic.  While in the polar region it may proceed more or less uninhibited, for a while, the equatorial layers form dense and massive accretion disk. The gravitational energy released in the disk can be very large, more then sufficient to stop the collapse of outer layers and drive GRB outflows, presumably in the polar direction where density is much lower \\cite{MW99}. In addition, there is plenty of rotational energy in the BH itself  \\begin{equation} E\\sub{rot} = \\frac{M\\sub{bh}c^2}{2} \\left\\{ 2-\\left[ \\left(1+\\sqrt{1-a^2}\\right)^2+a^2 \\right]^{1/2} \\right\\}, \\end{equation} where $M\\sub{bh}$ is the BH mass and $a\\in(-1,1)$ is its dimensionless rotation parameter. For $M\\sub{bh}=3M_{\\sun}$ and $a=0.9$ this gives the enormous value of $E\\sub{rot} \\simeq8\\times10^{53}$erg.  The three currently actively discussed mechanisms of powering GRB jets in the collapsar scenario are the heating via annihilation of neutrinos produced in the disk \\cite{MW99}, the magnetic braking of the disk \\cite{BP82,UM06}, and the magnetic braking of the BH \\cite{BZ77}.  The potential role of neutrino mechanism is rather difficult to assess as this requires accurate treatment of neutrino transport in a complex dynamic environment of collapsar. The long and complicated history of numerical studies of neutrino-driven supernova explosions teaches us to be cautious.  Numerical simulations by MacFadyen \\& Woosley\\shortcite{MW99} and Aloy et al.\\shortcite{AIMGM00} have demonstrated that sufficiently large energy deposition in the polar region above the disk may indeed result in fast collimated jets. However, the neutrino transport has not been implemented in these simulations and the energy deposition was based simply on expectations.  When Nagataki et al.\\shortcite{NTMT07} utilized a simple prescription for neutrino transport in their code they found that neutrino heating was insufficient to drive polar jets.  A number of groups have studied the collapsar scenario using Newtonian MHD codes and implementing the Paczynski-Witta potential in order to approximate the gravitational field of central BH \\cite{PMAB03,FKYHS06,NTMT07}. In this approach it is impossible to capture the Blandford-Znajek effect and only the magnetic braking of the accretion disk can be investigated. The general conclusion of these studies is that the accretion disk can launch magnetically-driven jets provided the magnetic field in the progenitor core is sufficiently strong. Unfortunately, the jet power has not been given in most of these papers and is difficult to evaluate from the published numbers.  In the simulations of Proga et al.\\shortcite{PMAB03} the jet power at $t\\simeq 0.25$s is $\\simeq 10^{50}\\mbox{erg}/$s.  The initial magnetic field in these simulations is monopole with $B\\simeq 2\\times 10^{14}$G at $r=3r_g$, where $r_g=GM\\sub{bh}/c^2$ (private communication).  \\begin{figure*} \\includegraphics[width=57mm]{figures/exp2.png} \\includegraphics[width=57mm]{figures/exp4.png} \\includegraphics[width=57mm]{figures/exp3.png} \\caption{Solution immediately before the explosion (t=0.24s).  Left panel: the baryonic rest mass density, $log_{10}\\rho$, in $g/cm^3$ and the magnetic field lines; Middle panel: the ratio of gas and magnetic pressures, $log_{10}P/P_m$, and velocity direction vectors; Right panel: the ratio of azimuthal and poloidal magnetic field strengths, $log_{10}B^\\phi/B_p$, and the magnetic field lines.  } \\label{f0} \\end{figure*}   The study of collapsars in full GRMHD is still in its infancy. Sekiguchi \\& Shibata\\shortcite{SS07} studied the collapse of rotating stellar cores and formation of BH in the collapsar scenario. Their results show powerful explosions soon after the accretion disk is formed around the BH and the free falling plasma of polar regions collides with this disk. These explosions are driven by the heat generated as a result of such collision. However, the authors have not accounted for the neutrino cooling and the energy losses due to photo-dissociation of atomic nuclei. and the explosions could be similar in nature to the ``successful'' prompt explosions of early supernova simulations \\cite{bethe}. Mizuno et al.\\shortcite{MYKS04a,MYKS04b} carried out GRMHD simulations in the time-independent space-time of a central BH.  The computational domain did not include the BH ergosphere and thus they could not study the role of the Blandford-Znajek effect~\\cite{K04a}. The energy losses have not been included and the equation of state (EOS) was a simple polytrope. These simulations were run for a rather short time, $\\simeq 280 r_g/c$ where $r_g=GM/c^2$, and jets were formed almost immediately due to unrealistically strong initial magnetic field.  In this letter we describe the first results of axisymmetric GRMHD simulations of collapsars where we use realistic EOS~\\cite{TS00}, include the energy losses due to neutrino emission (assuming optically thin regime) and photo-dissociation of nuclei (see the details of micro-physics in Komissarov \\& Barkov 2007 ), use the computational domain that includes the BH horizon and its ergosphere, and run simulations for a relatively long physical time, up to 0.5s. The neutrino heating is not included.   \\begin{figure*} \\includegraphics[width=57mm]{figures/rho1.png} \\includegraphics[width=57mm]{figures/rho2.png} \\includegraphics[width=57mm]{figures/rho3.png} \\caption{Solution on different scales at $t=0.45$s. The colour images show the baryonic rest mass density, $log_{10}\\rho$ in g/cm$^3$, the contours show the magnetic field lines, and the arrows show the velocity field.} \\label{f1} \\end{figure*}  \\begin{figure*} \\includegraphics[width=57mm]{figures/b1.png} \\includegraphics[width=57mm]{figures/b2.png} \\includegraphics[width=57mm]{figures/b3.png} \\caption{The inner region at $t=0.45$s.  Left panel: the magnetization parameter, $log_{10}P/P_m$, and the magnetic field lines; Middle panel: the ratio of azimuthal and poloidal magnetic field strengths, $log_{10}B^\\phi/B_p$, and the magnetic field lines; Right panel: the magnetic field strength, $log_{10}(B)$, and the magnetic field lines. } \\label{f2} \\end{figure*} ",
        "conclusions": "Our results provide strong support to the idea that magnetic fields can play a crucial role in driving powerful GRB jets and associated stellar explosions not only in the magnetar model but also in the collapsar model. The main energy source for the jets and explosions in our simulations is the rotational energy of black hole and it is released via the Blandford-Znajek mechanism. The measured rate of energy release, $\\dot{E} \\ge 2\\times10^{51}\\mbox{erg}\\,\\mbox{s}^{-1}$, can explain the energetics of even the shortest of long duration GRBs. The fact that the rotational energy of black hole, $E\\sub{bh}\\simeq\\mbox{few}\\times10^{53}\\mbox{erg}$, exceeds the typical explosion values derived from observations, $E\\simeq 10^{52}\\mbox{erg}$, suggests a self-regulating process in which the black hole activity ceases when the blast wave terminates further mass supply to the accretion disk.  The full details of the simulations together with the results of parameter study will be presented elsewhere."
    },
    "0710/0710.5418_arXiv.txt": {
        "abstract": "By applying recent results for the slab correlation time scale onto cosmic ray scattering theory, we compute cosmic ray parallel mean free paths within the quasilinear limit. By employing these results onto charged particle transport in the solar system, we demonstrate that much larger parallel mean free paths can be obtained in comparison to previous results. A comparison with solar wind observations is also presented to show that the new theoretical results are much closer to the observations than the previous results. ",
        "introduction": "Cosmic rays (CRs) interacting with turbulent magnetic fields get scattered and accelerated (see Melrose 1968, Schlickeiser 2002). The theoretical description of these scattering and acceleration processes are essential for understanding the penetration and modulation of low-energy cosmic rays in the heliosphere, the confinement and escape of galactic cosmic rays from the Galaxy, and the efficiency of diffusive shock acceleration mechanisms.  A key factor in CR scattering are the properties of the magnetic fields. A standard approach is the assumption of a superposition of  a mean magnetic field $\\vec{B}_0 = B_0 \\vec{e}_z$ and a turbulent component $\\delta \\vec{B} (\\vec{x})$. Whereas the mean field can easily be meassured in the solar system (here we find approximatelly $B_0 \\approx 4-5 nT$), the turbulent component has to be emulated by turbulence models. In the literature there is no consensus available about the true turbulence properties (see Cho \\& Lazarian 2005 for a review). In the solar system, however, some turbulence properties such as the wave spectrum can be obtained from meassurements (see e.g. Denskat \\& Neubauer 1983, Bruno \\& Carbone 2005).  More unclear are the orientation of the turbulence wave vectors (also refered to as turbulence geometry) and the dynamical decorrelation of the magnetic fields. In a recent CR diffusion study (Shalchi et al. 2006) a slab/2D composite model was combined with a nonlinear anisotropic dynamical turbulence (NADT) model. This model can be used to reproduce meassured CR mean free paths parallel and perpendicular to the mean field $\\vec{B}_0$. The authors of this article assumed that the slab correlation time scale is independent of the wave vector $\\vec{k}$.  In a recent study (Lazarian \\& Beresnyak 2006), however, it was shown that the slab time scale is indeed $\\vec{k}-$dependent.\\footnote{That study put to the test the idea of the damping of slab perturbations by the ambient turbulence in Yan \\& Lazarian (2002), Farmer \\& Goldreich (2004).} More precisely it was found that $t_c^{-1} = \\gamma_c = v_A \\sqrt{k_{\\parallel} / L}$; here we used the correlation time $t_c$, the correlation rate $\\gamma_c$, the Alfv\\'en speed $v_A$, and the outer scale of the turbulence $L$. It is the purpose of this article to apply this new result of the slab correlation time scale onto cosmic ray parallel diffusion. A comparison with solar wind observations of the parallel mean free path is also presented. It is demonstrated that we can find a much larger parallel mean free path if we employ the correlation time scale of Lazarian \\& Beresnyak (2006).  In Section 2 we explain the turbulence model that is used in this article. In Section 3 a quasilinear description of cosmic ray scattering is combined with this turbulence model to derive analytic forms of the pitch-angle diffusion coefficient and the parallel mean free path. In Section 4 we evaluate these formulas numerically to compute diffusion coefficients and we also provide a comparison with previous results and solar wind observations. In the closing Section 5 our results are summerized. ",
        "conclusions": "The theoretical explanation of measured parallel mean free paths in the solar system is a fundamental problem of space science. In a recent article (Shalchi et al. 2006) it has been demonstrated the these observations can indeed be reproduced theoretically. By using recent results of turbulence theory (Lazarian \\& Beresnyak 2006) we further improved the dynamical correlation function which is a key input in transport theory considerations. It is demonstrated in this article that the improved slab correlation time scale (see Eq. (\\ref{Alexmodel})) leads to a much larger parallel mean free path (see Fig. \\ref{nadtf4}). This effect is important since it was argued in several previous articles that the theoretical parallel mean free path is too small (Palmer 1982, Bieber et al. 1994) in comparison with solar wind observations.  Another problem of cosmic ray scattering theory is the importance of nonlinear effects. Whereas we have applied QLT in the current article it was argued in other papers (e.g. Shalchi et al. 2004) that nonlinear effects are important for parallel diffusion. However, these nonlinear effects are directly related to the interaction between charged particles and 2D modes. These 2D modes were neglected since we assumed pure slab fluctuations. Therefore, QLT can be applied and the results presented in this article should be valid. For non-slab models, where 2D modes are present, however, the applicability of QLT is questionable. It has to be subject of future work to explore the validity of QLT for realistic turbulence models such as dynamical turbulence models in non-slab geometry."
    },
    "0710/0710.4522_arXiv.txt": {
        "abstract": "We directly measure the evolution of the intergalactic Lyman-$\\alpha$ effective optical depth, $\\tau_{\\rm eff}$, over the redshift range $2 \\leq z \\leq 4.2$ from a sample of 86 high-resolution, high-signal-to-noise quasar spectra obtained with Keck/ESI, Keck/HIRES, and Magellan/MIKE. We find that our estimates of the quasar continuum levels in the \\Lya~forest obtained by spline fitting are systematically biased low, but that this bias can be accounted for using mock spectra. The mean fractional error $\\langle \\Delta C/C_{\\rm true} \\rangle$ is $<1\\%$ at $z=2$, 4\\% at $z=3$, and 12\\% at $z=4$. We provide estimates of the level of absorption arising from metals in the \\Lya~forest based on both direct and statistical metal removal results in the literature, finding that this contribution is $\\approx6-9\\%$ at $z=3$ and decreases monotonically with redshift. The high precision of our measurement indicates significant departures from the best-fit power-law redshift evolution, particularly near $z=3.2$. ",
        "introduction": "The evolution of the intergalactic medium (IGM) as traced by the \\Lya~forest provides a powerful record of the thermal and radiative history of the Universe. This power owes to our ability to measure the \\Lya~opacity of the IGM as a function of redshift, as well as to the relatively simple physics of the \\Lya~forest. In fact, cosmological simulations in which the forest arises from absorption by smooth density fluctuations imposed on the warm photoionized IGM as a natural consequence of hierarchical structure formation within cold dark matter models \\citep[e.g.,][]{1996ApJ...457L..51H}, have been remarkably successful at reproducing the properties of the absorption observed in actual quasar spectra. This synergy between theory and observations make the \\Lya~forest a particularly compelling probe of the diffuse Universe.\\\\ \\\\ In this work, we present a direct precision measurement of the effective \\Lya~optical depth and its evolution over the redshift range $2\\leq z \\leq4.2$ from a sample of 86 high-resolution, high signal-to-noise quasar spectra obtained obtained with Keck/ESI (16), Keck/HIRES (44), and with Magellan/MIKE (26). The full details have been reported in \\citep{2007arXiv0709.2382F}. We assume a cosmology with $(\\Omega_{m},~\\Omega_{b},~\\Omega_{\\Lambda},~h,~\\sigma_{8})=(0.27,~0.046,~0.73,~0.7,~0.8)$ \\citep[][]{2007ApJS..170..377S}. ",
        "conclusions": ""
    },
    "0710/0710.3591_arXiv.txt": {
        "abstract": "The near-infrared spectrum of (50000) Quaoar obtained at the Keck Observatory shows distinct absorption features of crystalline water ice, solid methane and ethane, and possibly other higher order hydrocarbons. Quaoar is only the fifth Kuiper belt object on which volatile ices have been detected. The small amount of methane on an otherwise water ice dominated surface suggests that Quaoar is a transition object between the dominant volatile-poor small Kuiper belt objects (KBOs) and the few volatile-rich large KBOs such as Pluto and Eris. ",
        "introduction": "While once Pluto and Triton were the only objects in the outer solar system known to contain volatile ices on their surfaces, the recent discoveries of frozen methane on the large Kuiper belt objects (KBOs) Eris, Sedna, and 2005 FY9 have shown that these objects are part of a new class of surface volatile rich bodies in the outer solar system \\citep{2005ApJ...635L..97B, 2006A&A...445L..35L,2005A&A...439L...1B, ebspec}. In contrast to these bodies with detectable volatiles, spectral observations of small KBOs over the past decade have found that most of these objects either contain varying amounts of involatile water ice on their surfaces or have flat spectra with no identifiable features \\citep{Kris}. To understand the dichotomy between volatile rich and volatile free surfaces in the outer solar system, \\citet{2007ApJ...659L..61S} constructed a simple model of atmospheric escape of volatile ices over the age of the solar system. They found that while most KBOs are too small and hot to retain their initial volatile ices to the present day, a small number are large and cold enough to retain these ices on their surfaces.  As anticipated, the model suggests that the largest KBOs, Eris, Pluto, and Sedna are all expected to retain surface volatiles, while the vast majority of the other known objects in the Kuiper belt are expected to have lost all surface volatiles. Two known intermediate-sized KBOs are predicted to be in the transition region where they may have differentially lost some volatile ices (N$_2$) but retained others (CH$_4$). One of these transition objects, 2005 FY9, with a diameter of $\\sim$1450 km \\citep{2007astro.ph..2538S} does indeed appear to contain abundant CH$_4$ but be depleted in N$_2$.  The other object that appears to be in the volatile non-volatile transition region is Quaoar, with a diameter of 1260$\\pm 190$ km\\citep{2004AJ....127.2413B}. The infrared spectrum of Quaoar does not resemble that of 2005 FY9, however. Quaoar's spectrum is dominated by absorptions due to involatile water ice, which is not detected at all on 2005 FY9. In addition, \\citet{2004Natur.432..731J} reported the detection of an absorption feature near 2.2 $\\mu m$ that they attributed to ammonia hydrate. They also detected the presence of crystalline water ice which, at the $\\sim$ 40 K radiative equilibrium temperature of Quaoar, is thought to be converted to amorphous water ice on a relatively short ($\\sim 10$ Myr) timescale by cosmic ray bombardment. The crystallinity of the water ice and the detection of the 2.2 $\\mu m$ feature that they attributed to ammonia hydrate led \\citet{2004Natur.432..731J} to suggest that Quaoar may have experienced relatively recent cryovolcanic activity.  In this paper, we present a new infrared spectrum of Quaoar with a signal-to-noise in the K-band six times greater than that of \\citet{2004Natur.432..731J} and model the ices present on the surface. ",
        "conclusions": "With significantly higher signal-to-noise in the 2.0 - 2.4 $\\mu m$ region, the 2.2 $\\mu m$ absorption feature on Quaoar previously identified as ammonia hydrate \\citep{2004Natur.432..731J} is clearly seen to be due to methane ice. No compelling evidence is seen for the presence of ammonia. The presence of crystalline water ice on the surface of Quaoar still remains unexplained because it is expected that ice should currently exist in the amorphous form on the $\\sim$ 40 K surface of Quaoar. However, the presence of the 1.65 $\\mu m$ absorption feature due to crystalline water ice in the spectrum of every well observed water ice rich KBO (even down to diameters of only a few hundred kilometers) \\citep{Kris} suggests that exotic processes such as cryovolcanism are unlikely to be required. The presence of crystalline water ice on so many small outer solar system bodies may indicate that our current understanding of the physics of the crystalline/amorphous phase transition may not be complete. The spectrum of Quaoar is consistent with that of a cold geologically dead object slowly losing the last of its volatile ices by escape in a tenuous, perhaps patchy, atmosphere.   Ethane is an expected by-product of irradiation of methane ice \\citep{2003NIMPB.209..283B}. The presence of ethane on Quaoar and on 2005 FY9 supports the suggestion of \\citet{ebspec} that these irradiation products are preferentially seen on bodies with large abundances of pure methane rather than on the bodies where the methane is diluted in nitrogen. Quaoar also appears to be rich in more complex irradiation products. Quaoar is the only water ice rich KBO which has a red color in the visible. Other water ice rich KBOs like Orcus, Charon, and 2003 EL61 and its family of collisional fragments are all blue in the visible \\citep{Kris}. Quaoar's red surface is likely due to the continued irradiation of methane, ethane, and their products on the surface \\citep{2006ApJ...644..646B}.  While methane on Quaoar is sufficiently volatile that it is likely to seasonally migrate if Quaoar has a moderate obliquity, ethane and the other irradiation products are essentially involatile at Quaoar's temperature. Quaoar is therefore likely to have an irregular covering of irradiation products, perhaps leading to rotational variability in its visible color and in the abundance of ethane. Continued observations of this object will provide insight into the nature of the volatile non-volatile transition and atmospheric escape in the outer solar system.   {\\it Acknowledgments:} We thank an anonymous referee for a helpful review. E.L.S. is supported by a NASA Graduate Student Research Fellowship.  The 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 the National Aeronautics and Space Administration.  The observatory was made possible by the generous financial support of the W.M. Keck Foundation."
    },
    "0710/0710.1846_arXiv.txt": {
        "abstract": "{The theory of stellar evolution can be more closely tested if we have the opportunity to measure new quantities. Nowadays, observations of galactic RR Lyr stars are available on a time baseline exceeding 100 years. Therefore, we can exploit the possibility of investigating period changes, continuing the pioneering work started by V.~P.~Tsesevich in 1969.} {We collected the available times of maximum brightness of the galactic RR Lyr stars in the GEOS RR Lyr database. Moreover, we also started new observational projects, including surveys with automated telescopes, to characterise the O--C diagrams better.} {The database we built has proved to be a very powerful tool for tracing the period variations through the ages. We analyzed 123 stars showing a clear O--C pattern (constant, parabolic or erratic) by means of  different least--squares methods.} {Clear evidence of period increases or decreases at constant rates has been found, suggesting evolutionary effects. The median values are $\\beta$=+0.14~d~Myr$^{-1}$ for the 27 stars showing a period increase and $\\beta$=--0.20~d~Myr$^{-1}$ for the 21 stars showing a period decrease. The large number of RR Lyr stars showing a period decrease (i.e., blueward evolution) is a new and  intriguing result. There is an excess of RR Lyr stars showing large, positive $\\beta$ values. Moreover, the observed $\\beta$ values are slightly larger than those predicted by theoretical models.} {} ",
        "introduction": "RR Lyr variables are low--mass stars in a core helium burning phase; they fill the part of the HR diagram where the horizontal branch intersects the classical instability strip. The crossing of the instability strip can take place in both directions; as a consequence, the periods will be either increasing, if the stars evolve from blue to red, or decreasing, if they evolve from red to blue. Despite its importance as a test for the stellar evolution theory, the {\\it observed} rate of the period changes is still an unknown quantity. It has been measured in globular clusters (e.g., Smith \\& Sandage \\cite{mfifteen}, Lee \\cite{lee}, Jurcsik et al. \\cite{omega}), comparing periods determined in different epochs. However, this task has not been undertaken yet  on the wide population of galactic RR Lyr stars, though most of them have been studied for tens of years, several of them since the end of the XIX$^{\\rm th}$ century, and therefore the available data are almost continuous. One of the critical points is to determine the importance of the negative rates, i.e., the period decreases. Such rates have been observed in some cluster stars (Jurcsik et  al. \\cite{omega}), but it is unclear in which evolutionary phase they occur (Sweigart \\& Renzini \\cite{random2}).  In the present paper, we derive the period variation rate of the best observed RR Lyr stars belonging to the  RRab sub-class. For this purpose we use the {\\it GEOS RR Lyr database} which is described in Sect.~\\ref{sect_grrdb}. The observations of the TAROT telescopes coordinated in the {\\it GEOS RR Lyr Survey} give a strong impulse to this analysis, providing a large number of times of maximum brightness in the recent years. They complete the effort of the amateur observers, in particular in the European associations BAV (Bundesdeutsche Arbeitsgemeinschaft f\\\"ur Ver\\\"anderliche Sterne) and GEOS (Groupe Europ\\'een d'Observations Stellaires), which have been  surveying these stars for years. These new observations are described in Sect.~\\ref{sect_grrs}. Next, Sect.~\\ref{sect_datan} shows how the data were analysed. In Sect.~\\ref{sect_cstrate}, we analyse the stars with constant period variation. Section~\\ref{sect_BLZH} is devoted to the study of some particularities encountered during the present work: Blazhko and light--time effects. Finally, a discussion of the results is developed in Sect.~\\ref{sect_discuss} in the light of stellar evolution. ",
        "conclusions": "The analysis of O--C variations over a timescale of 100 years or more has proved to be a powerful tool for providing quantitative tests of the stellar evolution of the horizontal branch stars. We can stress some well--established observational facts: \\begin{enumerate} \\item RRab stars experiencing blueward evolution (i.e., period decreases) are quite common, only slightly less so than RRab stars experiencing redward evolution (i.e., period increases). The period ranges covered by the two groups are very similar and the mean and median period values are nearly coincident; \\item the absolute values of period changes are larger than expected. Also in the extreme case of the rapid--evolving star \\object{SV Eri} the rate is much larger than expected; \\item the O--C behavior can be very complicated in some cases, with abrupt or irregular period variations, rather than monotonic. The regular variations caused by the light--time effect (and hence duplicity), often invoked to explain large O--C excursions, are not convincingly observed  in our extensive sample.  The physical explanations should be searched in the stellar structure; \\item Blazhko effect is often superimposed on secular changes, but the monotonic trend due to evolutionary variations still remains visible. \\end{enumerate} As a general conclusion about the comparison between our observational results and the theoretical predictions, we claim that there is a very powerful feedback between the two approaches. In particular, theoretical investigations should take into account that we have observational evidence of many RR Lyr stars showing blueward evolution. The theoretical models should also match the observed $\\beta$ values in a more satisfactory way, as these seem to be higher than expected, both for redward and blueward evolutions. On the other hand, the observational effort to monitor RR Lyr stars should be continued, possibly extended to stars at fainter magnitudes, and the accuracy of the maximum time determinations should be improved (by using automated telescopes), thus obtaining the same information on a shorter time baseline."
    },
    "0710/0710.1077_arXiv.txt": {
        "abstract": "In this work we give special attention to the bimetric theory of gravitation with massive gravitons proposed by Visser in 1998. In his theory, a prior background metric is necessary to take in account the massive term. Although in the great part of the astrophysical studies the Minkowski metric is the best choice to the background metric, it is not possible to consider this metric in cosmology. In order to keep the Minkowski metric as background in this case, we suggest an interpretation of the energy-momentum conservation in Visser's theory, which is in accordance with the equivalence principle and recovers naturally the special relativity in the absence of gravitational sources. Although we do not present a general proof of our hypothesis we show its validity in the simple case of a plane and dust-dominated universe, in which the `massive term' appears like an extra contribution for the energy density. ",
        "introduction": "\\label{intro} Could the graviton have a non-zero rest mass? The observations have shown that this is a possibility. One of the most accurate bounding on the mass of the graviton comes from the observations of the planetary motion in the solar system. Variations on the third Kepler law comparing the orbits of Earth and Mars can lead us to $m_g < 7.8 \\times 10^{-55}g$ \\cite{Talmadge88}. Another bound comes from the analysis of galaxy clusters that lead to $m_g <  2 \\times 10^{-62}g$ \\cite{Goldhaber74} which is considerably more restrictive but less robust due to the uncertainties in the content of the universe in large scales. Studying rotation curves of galactic disks, \\cite{JC} has found that we should have a massive graviton of $m_g \\ll 10^{-59}g$ in order to obtain a galactic disk with a scale length of $b\\backsim10$ kpc.  The above tests are obtained from static fields based on deviations of the newtonian gravity. In the weak field limit has been proposed \\cite{Finn2002} to constraint $m_g$ using data on the orbital decay of binary pulsars. From the  binary pulsar  PSR B1913+16 (Hulse-Taylor pulsar) and PSR B1534+12 it is found the limit $m_g < 1.4 \\times 10^{-52}g$, which is weaker than the bounds in static field.  It is worth recalling that the mass term introduced via a Pauli-Fierz (PF) term in the linearized approximation produces a theory whose predictions do not reduce to those of general relativity for $m_g \\rightarrow 0$. This is the so called van Dam Veltmann Zakharov discontinuity \\cite{Veltman1970}. Moreover the Minkowski space as background metric is unstable for the PF theory \\cite{Gruzinov2005}. However, there is no reason to prefer the PF term over any other non-PF quadratic terms.  It is important to emphasize that these mass terms do not have clear extrapolation to strong fields. A way to do that was proposed by Visser \\cite{visser98}. To generalize the theory to strong fields, Visser makes use of two metrics, the dynamical metric ($g_{\\mu\\nu}$) and a non-dynamical background metric ($\\left( g_0\\right) _{\\mu\\nu}$) that are connected by the mass term. Although adding a prior geometry is not in accordance with the usual foundations underlying Einstein gravity, it keeps intact the principles of equivalence (at least in its weak form) and general covariance in the Visser's work. Some interesting physical features emerge from the theory such as extra states of polarizations of the gravitational waves \\cite{wayne2004}.  In the present article, we explore some aspects which are not treated by Visser in his original paper. In the great part of the astrophysical studies the Minkowski metric is the most appropriate choice to the background metric. However, in the study of cosmology, it is not possible to consider this kind of metric, and we need some prior considerations regarding a background metric. Once this problem emerges from the coupling of the two metrics and the energy's conservation condition, we analyze an alternative interpretation of this condition. We also show that this interpretation is in accordance with the equivalence principle and recovers naturally the special relativity in the absence of gravitational sources. Arguments in favor of a Minkowskian background metric in Visser's theory are also considered.  This paper is organized as follows: in section \\ref{sec:1} we show how to introduce a mass for the graviton through a non-PF term. We present the strong field extrapolation as given by Visser in section \\ref{sec:2}. In section \\ref{sec:3} we show that the theory is not in accordance with a Minkowski background metric in the study of cosmology. In section \\ref{sec:4} we re-interpret the stress-energy conservation in order to keep Minkowski as background in any case. In particular, we show that our re-interpretation is in accordance with the equivalence principle. In section \\ref{sec:5} we show why Minkowski is the most natural choice to the background metric. We briefly study some cosmological consequences of our interpretation of the energy-momentum conservation in section \\ref{sec:6}. And finally, we present our conclusions in the last section. ",
        "conclusions": "\\label{sec:7}  Our interpretation of the energy-momentum conservation in the Visser's massive gravity is in accordance to the equivalence principle and recover naturally the results of special relativity in the absence of gravitational sources.  The point of view considered in this paper allow us to consider Minkowski as background metric in Visser's theory in all astrophysical cases including cosmology.  This new interpretation may lead to interesting cosmological results once we can construct a cosmological model in a theory with massive gravitons with a Minkowski background. Additional contributions to the cosmological fluids will appear due to the modifications in the interaction potential, which, maybe, would be a way of treat the dark-energy problem. The analyses of the theory in the absence of gravitational sources lead us to exclude the de Sitter space-time as a vacuum solution of the massive gravity, once a constant $\\Lambda$ term is rigorously zero in a flat background.  Another interesting feature is that our interpretation of the energy conservation in strong fields is independent of the form of the tensor which interact with the perfect-fluid tensor, so this can be used to other models with additional energy-momentum contribution."
    },
    "0710/0710.0247_arXiv.txt": {
        "abstract": "We  present high spatial resolution spectroscopic measurements of dynamic fibrils (DFs) in the Ca~{\\small{II}}~8662~{\\AA} line. These data show clear Doppler shifts in the identified DFs, which demonstrates that at least a subset of DFs are actual mass motions in the chromosphere. A statistical analysis of 26 DFs reveals a strong and statistically significant correlation between the maximal velocity and the deceleration. The range of the velocities and the decelerations are substantially lower, about a factor two, in our spectroscopic observations compared to the earlier results based on proper motion in narrow band images. There are fundamental differences in the different observational methods; when DFs are observed spectroscopically the measured Doppler shifts are a result of the atmospheric velocity, weighted with the response function to velocity over an extended height. When the proper motion of DFs is observed in narrow band images, the movement of the top of the DF is observed. This point is sharply defined because of the high contrast between the DF and the surroundings. The observational differences between the two methods are examined by several numerical experiments using both numerical simulations and a time series of narrow band H$\\alpha$ images. With basis in the simulations we conclude that the lower maximal velocity is explained by the low formation height of the Ca~IR~line. We conclude that the present observations support the earlier result that DFs are driven by magneto-acoustic shocks exited by convective flows and p-modes. ",
        "introduction": "The dynamical nature of the chromosphere is obvious when the Sun is imaged in the line center of strong chromospheric spectral lines, most commonly in the H$\\alpha$ line \\citep[e.g.,][]{2006Noort}. One of the dominating features is a vast number of thin (0.2-1 Mm) omnipresent jet like structures \\citep{1968Beckers}. On the quiet solar limb they are commonly known as spicules, while on the quiet disk they are often called mottles, and finally in active regions they are known as active region fibrils or dynamic fibrils (DFs). The nomenclature can be confusing, but there are strong indications that these structures are physically closely related \\citep{1994Tsiropoula,1995Suematsu,2001Christo,2007Rouppe}.  Recently, a combination of high--resolution observations and advanced numerical modeling have shown that DFs are most likely driven by shocks that form when photospheric oscillations leak into the chromosphere along inclined flux tubes \\citep{1990Suematsu,2004dePontieu,2006Hansteen,2007dePontieu}. The inclination of the magnetic field lowers the acoustic cutoff frequency sufficiently to allow p--modes with the dominant low frequencies to propagate along flux tubes \\citep{1973Mich,1977Bel}.  These insights into the formation of DFs have become possible because of recent developments in observational techniques, such as bigger telescopes combined with real time wavefront corrections by adaptive optics (AO) systems \\citep[e.g.][]{2000Rimmele,SSTAO}, and post-processing methods \\citep[e.g.][]{speckle,momfbd} which have made observations of these jet structures much more reliable.  These developments have spurred several authors to focus on the detailed understanding of DFs \\citep[e.g.,][]{2003dePontieu,2004dePontieu, 2005dePontieu,2007dePontieu,2004Kostas,2006Hansteen,2006deWijn,2007Julius, 2007Lars}. One of the important results of this work is that the DFs are driven by and can channel photospheric oscillations into the chromosphere and the corona.  \\begin{figure*}[!ht] \\includegraphics[width=\\textwidth]{f1.eps} \\caption{ Spectrograms of the H$\\alpha$ line (panel a) and the Ca~{\\small{II}}~8662~{\\AA} line (panel b). Notice the highly dynamical line center of the Ca line. The corresponding MFBD processed slitjaw image (panel c) and the narrow band H$\\alpha$ DOT image (panel d). The position of the spectrograph's slit is marked with a line in the DOT image. } \\label{plotone} \\end{figure*}  In two papers \\citet{2006Hansteen} and \\citet{2007dePontieu} used high spatial and high cadence observations of the H$\\alpha$ line center together with realistic simulations to investigate the nature of DFs. One of their conclusions was that the DFs follow parabolic paths along their axis with decelerations lower than the solar gravitational deceleration. Previous observations did not allow an accurate determination of the nature of the trajectory because of lower quality data and line--of--sight (LOS) effects that were difficult to estimate \\citep{1988Nishikawa,1995Suematsu}.  \\citet{2006Hansteen} and \\citet{2007dePontieu} further report regional differences between DFs observed in two different plage areas. The regional differences are explained by different inclination angles of the magnetic fields in the two regions. In this way the magnetic topology of the solar atmosphere works as a filter, where only waves with certain periods can leak through. The simulations, spanning from the upper convection zone to the corona, reproduce the observed correlations between the maximum velocities and decelerations in DFs leading to the conclusion that DFs are formed by chromospheric shocks driven by global p-modes and convective flows.  In this paper we add to the understanding of the DFs, by analyzing high spatial and high cadence spectrograms of the Ca {\\small{II}} 8662 {\\AA} line, put into context by simultaneous H$\\alpha$ spectrograms and narrow-band images. Furthermore, the observational results are compared with the numerical simulations of \\citet{2006Hansteen}.  In \\S~\\ref{Obs} we describe the observing program and instrumentation. The data reduction method is described in \\S~\\ref{Data}. In \\S~\\ref{Obsres} we show the results of the observations. The main observational errors are discussed in \\S~\\ref{error}. To get a better understanding of the observations we present several numerical experiments in \\S~\\ref{Sim}. Finally, we summarize the results in \\S~\\ref{Con}. ",
        "conclusions": "\\label{Con}  In this work we have presented co-spatial and co-temporal narrow band H$\\alpha$ images and spectroscopic measurements of the Ca~{\\small{II}}~8662~{\\AA} line. These observations have been used to identify 26 DFs, and measure their Doppler shifts. A reduced--$\\chi^2$ analysis shows that the time evolution of the Doppler shifts are well approximated by a linear fit, if the measurement errors are about $2$~km\\,s${}^{-1}$. Using this approximation we derive values for the decelerations and maximal velocities for each DF. Scatter plots of the deceleration and maximal velocity show a strong positive correlation between the two. We also observe weak correlations between the deceleration and lifetime and the maximal velocity and lifetime. These results are supporting the shock-wave theory as explanation model for the DFs \\citep{2006Hansteen, 2007dePontieu,2007Lars}.  Furthermore, the Doppler shifts show that at least a subset of DFs are caused by mass moving up and down in the atmosphere. The values of the maximum velocity and decelerations are all somewhat lower than earlier reported values. Using numerical experiments we have explained the differences in the two observational sets with the intrinsic differences in observational methods. Earlier observations have used the high contrast seen between the top of the DF and the background for measuring the proper motion of the DF. This high contrast is caused by the intensity increase due to contributions from the transition region. In the present observations the DF motion is measured using Doppler shifts, which are affected by the atmospheric conditions over the formation height. The formation height of the Ca~{\\small{II}}~8662~{\\AA}~line is much lower than the transition region. Since the shock amplitude is increasing from the formation height for the Ca IR line to the transition region we necessarily measure lower velocities using spectroscopy. The difference in maximal velocities derived from the two methods in our simulations is about a factor two, which is about the same as observed.  The one dimensional nature of the slit spectrograph somewhat affects our results, but experiments with narrow band images show that this is not altering the results significantly."
    },
    "0710/0710.5712_arXiv.txt": {
        "abstract": "The cosmological concordance model contains two separate constituents which interact only gravitationally with themselves and everything else, the dark matter and the dark energy. In the standard dark energy models, the dark matter makes up some 20\\% of the total energy budget today, while the dark energy is responsible for about 75\\%. Here we show that these numbers are only robust for specific dark energy models and that in general we cannot measure the abundance of the dark constituents separately without making strong assumptions. ",
        "introduction": "We cosmologists are very proud that our field has finally reached the status of ``precision science'' over the last decade. The quality of the current observations of the cosmic microwave background (CMB), the galaxy distribution and the luminosity distance to type Ia supernovae (SN-Ia) is indeed impressive, and has allowed the construction of a concordance model in which the universe contains the known, baryonic, matter (5\\% of the energy density today), radiation (negligible energy density today), dark matter (20\\%) and dark energy (75\\%).  The need for dark matter became apparent long ago in order to explain the motion of galaxies in clusters \\cite{darkmat} and the observed galaxy rotation curves. In cosmology, it is often modelled as a pressureless fluid with negligible interactions.  Much more recently, less than ten years ago, new SN-Ia data \\cite{sn1a} convinced the majority of cosmologists that dark energy was needed as well. Until today the nature of the dark energy is a deep mystery.  Although many models have been proposed, there are none that can explain its current abundance in a natural way. The alternative to the model building is a more phenomenological approach, where one measures the physical properties of the dark energy. To this end, one introduces a completely general fluid and tries to determine its characteristics from observations. ",
        "conclusions": "We have seen that cosmology cannot measure separately the properties of the dark matter and of a general dark energy component. In order to do that, we either need to impose additional assumptions, for example that the dark energy is a scalar field, or else we need a non-gravitational measurement of the dark matter properties, specifically of its contribution to the total energy density of the universe. One possibility is a detection of supersymmetry at LHC, which may in turn determine the abundance and mass of the lightest stable SUSY particle, one of the best candidates for the dark matter.  As a corollary, if the abundance determined in this way is not the one expected within the $\\Lambda$CDM cosmological concordance model, one possible explanation is an evolving dark energy.  We can also consider the degeneracy as a test of the generality of the different approaches to measure the dark energy equation of state. Since no analysis so far seems to have found it, we can only wonder what else has been overlooked.  Finally, although we show here that one never can prove experimentally from cosmological data alone that the dark energy is a cosmological constant, it is remarkable that a model containing just cold dark matter and $\\Lambda$ fits the data so well. From a model selection point of view $\\Lambda$CDM is still the preferred model because of its simplicity.  \\ack It is a pleasure to thank Luca Amendola, Ruth Durrer, Domenico Sapone and Anze Slosar for interesting discussions. MK acknowledges funding by the Swiss NSF."
    },
    "0710/0710.0137_arXiv.txt": {
        "abstract": "We present multiple epochs of H$\\alpha$ spectroscopy for 47 members of the open cluster NGC 3766 to investigate the long term variability of its Be stars.  Sixteen of the stars in this sample are Be stars, including one new discovery.  Of these, we observe an unprecedented 11 Be stars that undergo disk appearances and/or near disappearances in our H$\\alpha$ spectra, making this the most variable population of Be stars known to date.  NGC 3766 is therefore an excellent location to study the formation mechanism of Be star disks.  From blue optical spectra of 38 cluster members and existing Str\\\"omgren photometry of the cluster, we also measure rotational velocities, effective temperatures, and polar surface gravities to investigate the physical and evolutionary factors that may contribute to the Be phenomenon.  Our analysis also provides improvements to the reddening and distance of NGC 3766, and we find $E(B-V) = 0.22 \\pm 0.03$ and $(V-M_{\\rm V})_0 = 11.6 \\pm 0.2$, respectively.  The Be stars are not associated with a particular stage of main-sequence evolution, but they are a population of rapidly rotating stars with a velocity distribution generally consistent with rotation at $70-80$\\% of the critical velocity, although systematic effects probably underestimate the true rotational velocities so that the rotation is much closer to critical.  Our measurements of the changing disk sizes are consistent with the idea that transitory, nonradial pulsations contribute to the formation of these highly variable disks. ",
        "introduction": "\\setcounter{footnote}{3} NGC 3766 is a rich, young open cluster in the Carina spiral arm that is well known for its high content of Be stars \\citep{slettebak1985}, and many previous studies of this cluster have focused on the characteristics of these stars to identify their evolutionary status. The cluster has been the target of numerous photometric studies \\citep{ahmed1962, yilmaz1976, shobbrook1985, shobbrook1987, moitinho1997, piatti1998, tadross2001, mcswain2005b}.  But despite these intensive investigations, the cluster's age and distance remain somewhat uncertain; measurements of its age range from 14.5 to 25 Myr (WEBDA\\footnote{The WEBDA database is maintained by E.\\ Paunzen and is available online at http://www.univie.ac.at/webda/navigation.html.}; \\citealt{lynga1987, moitinho1997, tadross2001}), and its distance is between 1.5 and 2.2 kpc.  The reddening $E$($B$-$V$) is between 0.16 and 0.22 (see the discussion of \\citealt{moitinho1997}).  Spectroscopic investigations of NGC 3766 have targeted a limited sample of cluster members, focusing primarily on the Be star and supergiant populations (\\citealt{harris1976}; \\citealt{mermilliod1982} and references therein; \\citealt{slettebak1985, levesque2005}).  Even the eclipsing double-lined spectroscopic binary BF Centauri (= HD 100915), a member of NGC 3766, has been largely neglected by modern spectroscopic observations (\\citealt{clausen2007} and references therein).  For most cluster members, no detailed information about their physical characteristics such as temperature, gravity, rotation, and metallicity are known.  In this work, we present red and blue optical spectra for both normal B-type and Be stars in the cluster.  Like many prior studies of NGC 3766, our primary goal is to investigate the Be star population; but unlike other works, we achieve a more complete understanding of this subset of B stars by comparing these emission-line objects to their non-emission counterparts.  Therefore we present measurements of the effective temperature, $T_{\\rm eff}$, surface gravity, $\\log g$, and in most cases the projected rotational velocity, $V \\sin i$, for 26 normal B stars and 16 Be stars in NGC 3766.  We use these results to improve the known reddening and distance to the cluster.  From multiple epochs of H$\\alpha$ spectroscopy, we also investigate the variability of the circumstellar disks and estimate the disk mass loss/gain rates for 11 Be stars.  Finally, we use the observed disk masses and angular momenta to show that nonradial pulsations are a possible origin for the disks, and they probably fill during short-lived bursts of mass flow from the stellar surface. ",
        "conclusions": "Our spectroscopic analysis of NGC 3766 has revealed that Be stars may be much more common than we originally thought.  In our photometric study of NGC 3766 \\citep{mcswain2005b}, we found up to 13 Be stars (5 definite, 8 uncertain) out of an expected 191 B-type stars, not counting the one Be star that saturated our photometry.  The new total of 16 Be stars is 23\\% greater.  Among these 16 Be stars, 2--5 of them appear to have almost no disk at any given time, and an additional 2--4 have extremely subtle emission in their H$\\alpha$ line profile that could easily be mistaken for other phenomena (such as NRP manifesting themselves as bumps moving across the line or SB2 line blending). Therefore 25--50\\% of the Be stars may go undetected in a single spectroscopic observation, and photometric snapshots are even less likely to discern such weak emitters.  We note four stars (Nos.\\ 27, 45, 49, and 77) that were found to be possible or likely Be stars in the photometric study by \\citet{shobbrook1985, shobbrook1987}, but they never showed emission during our observations and thus remain unconfirmed.  The existence of transitory, weak disks (especially Nos.\\ 130 and 196) could mean that many more Be stars are waiting to be discovered.  For our total sample of 48 Southern open clusters in our photometric survey, we found a low Be fraction of $2-7$\\% \\citep{mcswain2005b}. Considering the very weak disks that are observed in NGC 3766 and the exceptionally high variability among the cluster's Be population, the total fraction of Be stars could be much greater.  We are currently performing a similar spectroscopic study of several other clusters from our survey, and we will address those results in a future paper.  While the Be stars of NGC 3766 are not distinguishable from normal B-type stars by their evolutionary states, they do form a population of rapidly rotating stars.  With two exceptions, their measured velocities are consistent with a uniform population of rapid rotators having $V = 0.7-0.8 \\; V_{\\rm crit}$.  Gravitational darkening and weak emission in the \\ion{He}{1} lines may mean that these velocities are underestimated by as much as 33\\% \\citep{townsend2004}, so the true $V_{\\rm rot}$ is probably at least $0.84 \\; V_{\\rm crit}$.  From the measured changes in the disks' masses and angular momenta, NRP are a capable source for the mass flow into the equatorial plane.  The pulsations may be a transitory phenomenon, however, and the variable nature of the Be stars probably reflects dramatic changes in the surface activity."
    },
    "0710/0710.2903_arXiv.txt": {
        "abstract": "We use the kinetic theory of nucleation to explore the properties of dust nucleation in sub-saturated vapors. Due to radiation losses, the sub-critical clusters have a smaller temperature compared to their vapor. This alters the dynamical balance between attachment and detachment of monomers, allowing for stable nucleation of grains in vapors that are sub-saturated for their temperature. We find this effect particularly important at low densities and in the absence of a strong background radiation field. We find new conditions for stable nucleation in the $n-T$ phase diagram. The nucleation in the non-LTE regions is likely to be at much slower rate than in the super-saturated vapors. We evaluate the nucleation rate, warning the reader that it does depend on poorly substantiated properties of the macro-molecules assumed in the computation. On the other hand, the conditions for nucleation depend only on the properties of the large stable grains and are more robust. We finally point out that this mechanism may be relevant in the early universe as an initial dust pollution mechanism, since once the interstellar medium is polluted with dust, mantle growth is likely to be dominant over non-LTE nucleation in the diffuse medium. ",
        "introduction": "Dust particles are one of the fundamental components of the interstellar medium (ISM) and an ever-present worry for observers due to their opacity at optical and UV wavelengths (Cardelli, Clayton \\& Mathis 1989). The ISM of the Milky Way is polluted by a mixture of grains made of a variety of materials, likely dominated by carbonaceous grains, silicates, and small PAHs particles (Mathis, Rumpl \\& Nordsieck 1977; Weingartner \\& Draine 2001). The dust properties are supposed to be the result of dust formation in the outflows of evolved stars (e.g. Salpeter1977; Stein \\& Soifer 1983; Mathis 1990; Whittet 1992; Draine 2003, and references therein) and subsequent evolution, and eventual dissolution, in the ISM, mainly as the effect of shock waves that destroy the grains through sputtering (Draine 1989; McKee 1989; Edmunds 2001). Alternative dust production sites are supernova explosions (Kozasa, Hasegawa \\& Nomoto 1989, 1991; Todini \\& Ferrara 2001; Nozawa et al. 2003; Schneider, Ferrara \\& Salvaterra 2004; Bianchi \\& Schneider 2007) and quasar outflows (Elvis, Marengo \\& Karovska, 2002).  The theory of dust nucleation in astrophysics is heavily influenced by the theory of the nucleation of phase transitions in super-saturated vapors (Becker \\& Doring 1935; Feder et al. 1966; Abraham 1974). The theory had mild success in reproducing nucleation rates, but is still controversial in many aspects, especially because it extrapolates the properties of macroscopic bodies to clusters of few molecules and because it extends the thermodynamic approach to systems with a handful of particles. To add to these problems, astrophysical dust nucleation requires chemical reactions, since grains of materials that do not have a vapor state do nucleate (think, for example, to the nucleation of olivines from silicon oxides and metals; Draine 1979; Gail \\& Sedlmayr 1986). An alternative approach is the so-called kinetic theory, which describes nucleation as the result of attachment and detachment of monomers from a seed cluster of $n$ particles (atoms, molecules or radicals; Nowakowski \\& Ruckenstein 1991ab).  Both the thermodynamic and the kinetic nucleation theory have been developed in conditions of {\\it true equilibrium}, i.e., when the two phases have the same temperature. In the astrophysical scenario, however, the temperature of the dust grains can be sensibly lower than the temperature of the gas in which they are embedded due to efficient radiation cooling (e.g., Draine 1981). This would seem to be irrelevant to nucleation theory, since a vapor needs to be already nucleated in order to have grains that can be colder than the gas phase. Even a sub-saturated vapor, however, has a large number of unstable clusters that form by random association of monomers (and rapidly evaporate). In this paper we study the effect of cooling of these proto-clusters in a sub-saturated vapor, and the effect this has on the balance between attachment and detachment of monomers. Using the kinetic theory of nucleation, we find that even largely sub-saturated vapors can nucleate, provided they are not immersed in a strong radiation field.  We compute, albeit under some controversial assumptions, the non-LTE nucleation rate. We show that, even though it is not as large as in super-saturated vapors, it can produce dust grains at a rate that can reproduce the average dust grain density in the Milky Way over a timescale of several million years. In addition we show that non-LTE effects can increase the nucleation rate in super-saturated vapors. Non-LTE nucleation could therefore provide a slow channel for dust formation, in which dust is built over a relative long time in a slowly evolving region. Such an example could be the outflow from AGN nuclei (Elvis et al. 2002). Such evolution is different from the one envisaged in the classical dust factories -- AGB star atmospheres and supernov\\ae\\ -- where dust nucleation is rapid but short lived since the favorable conditions are rapidly lost.  This paper is organized as follows: in \\S~2 we briefly review the classical kinetic theory of nucleation; in \\S~3 we compute the dust grain temperature and in \\S~4 we compute the new conditions for nucleation. In \\S~5 we consider the nucleation rate and discuss our results in \\S~6. ",
        "conclusions": "We have considered the effect of grain (cluster, droplet) cooling in the nucleation of liquid and solid phases in vapors. We find that the effect can be dramatic on the nucleation rate and on the nucleation phase diagram, allowing for nucleation in large regions of the parameter space that are classically considered to be non-nucleating.  As is in general true for nucleation, there are several limits and approximations that we should bear in mind when considering the theory from the quantitative point of view. As exemplified by the comparison of the data with the theory in Fig.~\\ref{fig:j}, several orders of magnitude can separate the nucleation rate prediction from the observations. The controversial points are:  \\begin{itemize} \\item {\\it Sticking coefficients} --- Most of the figures and computations in this work assume $k_s=1$. This is not always true (Batista et al. 2005). In addition to its dependence on temperature, the sticking coefficient may depend on the size of the cluster. A big cluster could more easily absorb the extra kinetic energy of the incoming monomer, compared to a small cluster (K00), and therefore $k_s$ may be significantly smaller than unity for very small clusters. \\item {\\it Capillary approximation} --- The capillary approximation, i.e., the assumption that the surface tension does not depend on the cluster size, is very controversial, and a change in the surface energy for very small clusters could result in big changes on the nucleation rate. For example, the discrepancy in Fig.~\\ref{fig:j} could be solved by assuming a larger surface tension for the very small water droplet. In addition, macromolecules do not even have a properly defined surface, and the whole concept does not apply. Finally, the exponential factor in Eq.~\\ref{eq:sigma} depends on the assumption that the clusters are spherical. A different ratio of the surface to the volume would modify this term. This is likely for small graphite clusters, since graphite tends to aggregate in a planar form. \\item {\\it Detachment rate for small clusters} --- In this paper, and in most nucleation theory, the detachment rate is computed by propagating to very small clusters the detachment rate of macroscopic bodies. It is very likely that, in the very small limit, the detachment is governed by completely different processes. Let us analyze the two limits. For a macroscopic body, the number of monomers is so large that monomers with a statistically higher energy can detach since their energy is larger than the binding energy. In the opposite limit of a dimer, the detachment has to be due to an external action: either a collision with a fast monomer or with a photon or with another cluster. In this limit destructive collisions have to be considered. \\item {\\it Cooling of macromolecules} --- This is a new problem that arises when the cooling of the grains is considered. We have assumed that the grains cool as modified black bodies down to the smallest sizes. When the number of monomers in the cluster becomes small, the cooling will not be through a continuum spectrum, but through lines and bands. A more refined treatment of cooling is necessary to compute accurate nucleation rates. \\item It should finally be kept in mind that we allowed for the complete cooling of the grains, neglecting the effects of a background radiation field in setting a lower limit to the temperature. In addition, we neglected the fact that the temperature of small grains is largely stochastic and that at high temperature some collisions between grains and gas particles can result in sputtering rather than accretion. \\end{itemize}  Despite all these caveats, the main result of this paper holds: the region where a vapor spontaneously nucleate is not limited to the classic region, when nucleation takes place in thermal equilibrium and $T_g=T_s$. Allowing for the clusters to cool inhibits the detachment of monomers from the clusters and allow for nucleation at higher temperatures and lower densities than in the classical scenario. The nucleation rates tend to be orders of magnitude smaller than those derived in thermal equilibrium, but provide a non-negligible correction to the LTE rate for mildly super-saturated vapors with $S\\gsim1$, especially at low temperature.  In the astrophysical scenario small nucleation rates are not a big worry. The ISM of our Galaxy contains approximately 1 per cent of its mass in dust grains (Mathis et al. 1977). This corresponds to approximately one dust particle every cubic meter (assuming a power-law grain size distribution as in Mathis et al. 1977). However, in the present day Universe the ISM is already polluted with dust and in the presence of seed grains it is likely that the process of mantle growth dominates over non-LTE nucleation in the diffuse medium. The process of non-LTE nucleation may therefore be important at high redshift, when the ISM is first polluted with metals by supernova explosions. It is unclear if supernov\\ae\\ do generate dust by themselves, and even more whether the generated dust can survive the sputtering in the forward-reverse shock systems (Bianchi \\& Schneider 2007; Nath, Laskar \\& Shull 2007). In the case that supernov\\ae\\ do mainly pollute the ISM with metals but with no or a negligible quantity of dust, non-LTE nucleation could become the dominant process of dust nucleation in the early universe. Detailed estimates of nucleation in the various scenarios require a more detailed understanding of the properties of the very small nuclei and are beyond the scope of this paper."
    },
    "0710/0710.2418_arXiv.txt": {
        "abstract": "{ We report here on a new VHE source, HESS~J1908+063, disovered during the extended H.E.S.S. survey of the Galactic plane and which coincides with the recently reported MILAGRO unidentified source MGRO~J1908+06. The position, extension and spectrum measurements of the HESS source are presented and compared to those of MGRO~J1908+06. Possible counterparts at other wavelenghts are discussed. For the first time one of the low-lattitude MILAGRO sources is confirmed.  }  \\begin{document} ",
        "introduction": "H.E.S.S. observations of the inner Galactic plane in the [$270^{\\circ}$, $30^{\\circ}$] longitude range have revealed more than two dozens of  new VHE sources, consisting of shell-type SNRs, pulsar wind nebulae, X-ray binary systems, a putative young star cluster, etc, and yet unidentified objects (see e.g. \\cite{HESSScanII} and  \\cite{HESSSurveyICRC07} in these proceedings for a summary).  The extended  H.E.S.S. survey in the [$30^{\\circ}$-$60^{\\circ}$] longitude range performed between 2005 and 2007 overlaps with regions covered by the MILAGRO sky survey at longitudes greater than $30^{\\circ}$. The latter experiment has recently reported \\cite{MILAGRO} three low-latitude sources including,  MGRO~J1908+06, detected after seven years of operation (2358 days of data) at 8.3$\\sigma$ (pre-trials) confidence level. MGRO~J1908+06, of which the extension remains unknown but bounded to a maximum diameter of 2.6$^{\\circ}$, is located near the galactic longitude $\\sim 40^{\\circ}$ and hence is covered by the H.E.S.S. galactic plane survey.  A new H.E.S.S. source, HESS~J1908+063, which coincides with MGRO~J1908+06, is presented here. Its position, size and spectrum are measured and compared to the MILAGRO source. Possible counterparts at other wavelengths are discussed in the light of the H.E.S.S. measurements. ",
        "conclusions": ""
    },
    "0710/0710.2126_arXiv.txt": {
        "abstract": "$\\mu$ Orionis was identified by spectroscopic studies as a quadruple star system.  Seventeen high precision differential astrometry measurements of $\\mu$ Ori have been collected by the Palomar High-precision Astrometric Search for Exoplanet Systems (PHASES).  These show both the motion of the long period binary orbit and short period perturbations superimposed on that caused by each of the components in the long period system being themselves binaries.  The new measurements enable the orientations of the long period binary and short period subsystems to be determined.  Recent theoretical work predicts the distribution of relative inclinations between inner and outer orbits of hierarchical systems to peak near 40 and 140 degrees. The degree of coplanarity of this complex system is determined, and the angle between the planes of the A-B and Aa-Ab orbits is found to be $136.7 \\pm 8.3$ degrees, near the predicted distribution peak at 140 degrees; this result is discussed in the context of the handful of systems with established mutual inclinations.  The system distance and masses for each component are obtained from a combined fit of the PHASES astrometry and archival radial velocity observations. The component masses have relative precisions of $5\\%$ (component Aa), $15\\%$ (Ab), and $1.4\\%$ (each of Ba and Bb).  The median size of the minor axes of the uncertainty ellipses for the new measurements is 20 micro-arcseconds ($\\microas$).  Updated orbits for $\\delta$ Equulei, $\\kappa$ Pegasi, and V819 Herculis are also presented. ",
        "introduction": "$\\mu$ Orionis (61 Ori, HR 2124, HIP 28614, HD 40932) is a quadruple star system that has been extensively studied by radial velocity (RV) and differential astrometry.  It is located just North of Betelgeuse, Orion's right shoulder (left on the sky); $\\mu$ Ori is a bright star that is visible to the unaided eye even in moderately light-polluted skies.  \\cite{Frost1906} discovered it to be a short period (four and a half day) single-lined spectroscopic binary; this was component Aa, whose short-period, low mass companion Ab has never been detected directly.  \\cite{Aitken1914} discovered it also had a more distant component (B) forming a sub-arcsecond visual binary.  Much later, \\cite{Fekel1980} found B was itself a short-period (4.78 days) double-lined spectroscopic binary, making the system quadruple; these stars are designated Ba and Bb.  Most recently, \\cite{Fekel2002} (hereafter F2002) reported the astrometric orbit of the A-B motion, double-lined RV orbits for A-B and the Ba-Bb subsystem, and a single-lined RV orbit for the Aa-Ab subsystem.  F2002 estimate the spectral types as A5V (Aa, an Am star), F5V (Ba), and F5V (Bb), though they note these are classifications are less certain than usual due to the complexity of the system.  For a more complete discussion of the history of $\\mu$ Ori, see F2002.  Until now, astrometric observations have only been able to characterize the long period A-B motion, lacking the precision necessary to measure the astrometric perturbations to this orbit caused by the Aa-Ab and Ba-Bb subsystems.  The method described by \\cite{LaneMute2004a} for ground-based differential astrometry at the $\\sim 20$ $\\microas$ level for sub-arcsecond (``speckle'') binaries has been used to study $\\mu$ Ori during the 2004-2007 observing seasons.  These measurements represent an improvement in precision of over two orders of magnitude over previous work on this system.  The goal of the current investigation is to report the center-of-light (photocenter) astrometric orbits of the Aa-Ab and Ba-Bb subsystems.  This enables measurement of the coplanarities of the A-B, Aa-Ab, and Ba-Bb orbits.  The masses and luminosity ratio of Aa and Ab are measured for the first time.Also presented are updated orbits for the PHASES targets $\\delta$ Equ, $\\kappa$ Peg, and V819 Her.  Astrometric measurements were made at the Palomar Testbed Interferometer \\citep[PTI;][]{col99} as part of the Palomar High-precision Astrometric Search for Exoplanet Systems (PHASES) program \\citep{Mute06Limits}.  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. ",
        "conclusions": "The center-of-light astrometric motions of the Aa-Ab and Ba-Bb subsystems in $\\mu$ Ori have been constrained by PHASES observations.  While four degenerate orbital solutions exist, two of these can be excluded with high reliability based on mass-luminosity arguments, and the fact that Ab is not observed in the spectra.  Ba and Bb are stars of a class (mid-F dwarfs) whose properties have been well established by studying other binaries.  Their association with Aa and Ab, which are members of more poorly studied classes (Am and late K dwarfs) allows a better understanding of those objects in a system which can be assumed to be coevolved.  The orbital solution finds masses and luminosities for all four components, the basic properties necessary in studying their natures.  Complex dynamics must occur in $\\mu$ Ori.  The Ba-Bb orbital plane is nearly perpendicular to that of the A-B motion, and certainly undergoes Kozai-type inclination-eccentricity oscillations.  It is possible that the mutual inclination of the A-B pair and Aa-Ab subsystem is a result of KCTF effects over the system's evolution.  Finally, it is noted that the orbits in the $\\mu$ Ori system are quite non-coplanar.  This is in striking contrast with the planets of the solar system, but follows the trend seen in triple star systems.  With the solar system being the only one whose coplanarity has been evaluated, it is difficult to draw conclusions about the configurations of planetary systems in general.  It is important that future investigations evaluate the coplanarities of extrasolar planetary systems to establish a distribution.  Whether that distribution will be the same or different than that of their stellar counterparts may point to similarities or differences in star and planet formation, and provide a key constraint on modeling multiple star and planet formation."
    },
    "0710/0710.0854_arXiv.txt": {
        "abstract": "{Many thermally emitting, isolated neutron stars have magnetic fields that are larger than $10^{13}$~G. A realistic cooling model that includes the presence of high magnetic fields should be reconsidered.} {We investigate the effects of an anisotropic temperature distribution and Joule heating on the cooling of magnetized neutron stars.} {The 2D heat transfer equation with anisotropic thermal conductivity tensor and including all relevant neutrino emission processes is solved for realistic models of the neutron star interior and crust.} {The presence of the magnetic field affects significantly the thermal surface distribution and the cooling history during both, the early neutrino cooling era and the late photon cooling era.} { There is a large effect of Joule heating on the thermal evolution of strongly magnetized neutron stars. Both magnetic fields and Joule heating play an important role in keeping magnetars warm for a long time. Moreover, this effect is important for intermediate field neutron stars and should be considered in radio--quiet isolated neutron stars or high magnetic field radio--pulsars.} ",
        "introduction": "Observation of thermal emission from neutron stars (NSs) can provide not only information on the physical properties such as the magnetic field, temperature, and chemical composition of the regions where this radiation is produced but also information on the properties of matter at higher densities deeper inside the star \\citep{Yakovlev2004,Page2006}. To derive this information, we need to calculate the structure and evolution of the star, and compare the theoretical model with the observational data. Most previous studies assumed a spherically symmetric temperature distribution. However, there is increasing evidence that this is not the case for most nearby neutron stars whose thermal emission is visible in the X-ray band of the electromagnetic spectrum  \\citep{Zavlin2007,Haberl2007}. The anisotropic temperature distribution may be produced not only in the low density regions where the spectrum is formed and preliminary investigations had focused their attention, but also in intermediate density regions, such as the solid crust, where a complicated magnetic field geometry could cause a coupled magneto-thermal evolution. In some extreme cases, this anisotropy may even be present in the poorly known interior, where neutrino processes are responsible  for the energy removal.  The observational fact that most thermally emitting isolated NSs have magnetic fields larger than $10^{13}$ G \\citep{Haberl2007}, which is sometimes confirmed by spin down measurements, leads to the conclusion that a realistic cooling model must not avoid the inclusion of the effects produced by the presence of high magnetic fields. The transport processes that occur in the interior are affected by these strong magnetic fields and their effects are expected to have observable consequences, in particular for highly magnetized NSs or magnetars. Moreover, the large surface magnetic field strengths inferred from the observations probably indicate that the interior field could reach even larger values, as theoretically predicted by some models \\citep{TD1993}.  The presence of a magnetic field affects the transport properties of all plasma components, especially the electrons. In general, the motion of free electrons perpendicular to the magnetic field is quantized in Landau levels, and the thermal and electrical conductivities exhibit quantum oscillations. In the limit of a strongly quantizing field, in which almost all electrons populate the lowest level, such as in the envelope of a NS, a quantum description is necessary to calculate the thermal and electrical conductivities. Earlier calculations by \\cite{Canuto1970} and \\cite{Itoh1975} concluded that the electron thermal conductivity is strongly suppressed in the direction perpendicular to the magnetic field and increased along the magnetic field lines, which reduces the thermal insulation of the envelope ({\\it heat blanketing}). Thus, there is an anisotropic heat transport in the NS's envelope governed by the magnetic field geometry,  that produces a non-uniform surface temperature.  The anisotropy in the surface temperature of a NS appears to be confirmed by the analysis of observational data from isolated NSs (see \\cite{Zavlin2007} and \\cite{Haberl2007} for reviews on the current status of theory and observations). The mismatch between the extrapolation to low energy of fits to the X-ray spectra, and the observed Rayleigh-Jeans tail in the optical band ({\\it optical excess flux}), cannot be addressed using a unique temperature. Several simultaneous fits to multiwavelength spectra of \\rxdieciocho \\citep{Pons2002,Truemper2004}, \\rbdoce \\citep{Schwope2005,Schwope2007},  and \\rxcerosiete \\citep{Perez2006} are explained by a small hot emitting area $\\simeq$ 10--20 km$^2$, and an extended cooler component. Another piece of evidence that strongly supports the nonuniform temperature distribution are pulsations in the X-ray signal  of some objects of amplitudes $\\simeq$ 5--30 $\\%$, some of which have irregular light curves that point towards a non-dipolar temperature distribution. All of these facts reveal that the idealized picture of a NS with a dipolar magnetic field  and uniform surface temperature is oversimplified.  In a pioneering work, \\cite{Greenstein1983} obtained the temperature at the surface of a NS as a function of the magnetic field inclination angle in a simplified plane-parallel approximation. This model was applied to different magnetic field configurations and the observational consequences of a non-uniform temperature distribution were analyzed in the pulsars Vela and Geminga among others \\citep{Page1995}. \\cite{Potekhin2001} improved the former calculations including realistic thermal conductivities. Nevertheless, the temperature anisotropy as generated in the envelope may be insufficiently to be consistent with the observed thermal distribution and, in this case, should originate deeper inside the NS \\citep{Geppert2004,Azorin2006}.  Crustal confined magnetic fields could be responsible for the surface thermal anisotropy. In the crust,  even if a strong magnetic field is present, the electrons occupy a large number of Landau levels and the classical approximation remains valid during a long time in the thermal evolution. The magnetic field limits the movement of electrons in the direction perpendicular to the field and, since they are the main carriers of the heat transport, the thermal conductivity in this direction is highly suppressed, while remaining almost unaffected along the field lines. Temperature distributions in the crust were obtained as stationary solutions of the diffusion equation with axial symmetry \\citep{Geppert2004}. The approach assumes an isothermal core and a magnetized envelope as an inner and outer boundary condition, respectively. The results show important deviations from the crust isothermal case for crustal confined magnetic fields with strengths larger than $10^{13}$ G and temperatures below $10^{8}$ K. Similar conclusions were obtained considering not only poloidal but also toroidal components for the magnetic field \\citep{Azorin2006, Geppert2006}. This models succeeded in explaining simultaneously the observed X-ray spectrum, the optical excess, the pulsed fraction, and other spectral features for some isolated NS such as \\rxcerosiete \\citep{Perez2006} and \\rxdieciocho  \\citep{Geppert2006}.  Non-uniform surface temperature in NSs was studied by different authors using simplified models \\citep{ShibYak1996,Potekhin2001}. Although these models can provide useful information, a detailed investigation of heat transport in 2D must be completed to obtain more solid conclusions. However, this is not the only effect that must be revisited to study the cooling of NSs. For isolated NSs, different relevant magnetic field dissipation processes were identified \\citep{Goldreich1992}. The {\\it Ohmic} dissipation rate is determined by the finite conductivity of the constituent matter. In the crust, the electrical resistivity is mainly due to electron-phonon and electron-impurity scattering processes \\citep{Flowers1976}, resulting in more efficient Ohmic dissipation than in the fluid interior. The strong temperature dependence of the resistivity leads to rapid dissipation of the magnetic energy in the outermost low-density regions during the early evolution of a hot NS, which becomes less relevant as the star cools down. Joule heating in the crustal layers due to Ohmic decay was thought to affect only the late photon cooling era in old NS ($\\geq 10^7$ yr), and to be an efficient mechanism to maintain the surface temperature as high as $\\simeq 10^{4-5}$ K for a long time \\citep{Miralles98}. \\cite{Page2000} studied the 1-D thermal evolution of NSs combined with an evolving Stokes function that defines a purely poloidal, dipolar magnetic field. The Joule heating rate was evaluated averaging the currents over the azimuthal angle. However, for strongly magnetized NS, Joule heating can be important much earlier in the evolution. In a recent work, \\cite{Kaminker2006} placed a heat source inside the outer crust of a young, warm, magnetar of field strength $5\\times10^{14}$ G. To explain observations, they concluded that the heat source should be located at a density $\\lesssim 5\\times 10^{11}$ g~cm$^{-3}$, and the heating rate should be $\\gtrsim 10^{20}$ erg~cm$^{-3}$~s$^{-1}$ for at least $5\\times10^4$ yr. Anisotropic heat transport is neglected in these simulations, which were performed in spherical symmetry, assuming that it will  not affect the results in the early evolution. Nevertheless we will show that, in 2D simulations, the effect of anisotropic heat transport is important.  In addition to purely Ohmic dissipation, strongly magnetized NSs can also experience a {\\it Hall drift} with a drift velocity proportional to the magnetic field strength. Although the Hall drift conserves the magnetic energy and it is not a dissipative mechanism by itself, it can enhance the Ohmic decay by  compressing the field into small scales, or by displacing currents to regions with higher resistivity, where Ohmic dissipation is more efficient. Recently, the first 2D-long term simulations of the magnetic field evolution in the crust studied the interplay of Ohmic dissipation and the Hall drift effect \\citep{PonsGeppert2007}. It was shown that, for magnetar field strength, the characteristic timescale during which Hall drift influences Ohmic dissipation is of about $10^{4}$ yr. All of these studies imply that both field decay and Joule heating play a role in the cooling of neutron stars born with field strengths $\\geq 10^{13}$ G.  We will show that, during the neutrino cooling era and the early stages of the photon cooling era, the thermal evolution is coupled to the magnetic field decay, since both cooling and magnetic field diffusion proceed on a similar timescale ($\\approx 10^{6} $ yr). The energy released by magnetic field decay in the crust could be an important heat source that modifies or even controls the thermal evolution of a NS. Observational evidence of this fact  is shown in \\cite{PonsLink2007}. They found a strong correlation between the inferred magnetic field and the surface temperature for a wide range of magnetic fields: from magnetars ($\\geq 10^{14}$ G), through radio-quiet isolated neutron stars ($\\simeq 10^{13}$ G) down to some ordinary pulsars ($\\leq 10^{13}$ G). The main conclusion is that, rather independently from the stellar structure and the matter composition, the correlation can be explained by the decay of currents on a timescale of $\\simeq 10^{6}$ yr.  The aim of the present work is to study in a more consistent way the cooling of a realistic NS under the effects of large magnetic fields, including the effects of an anisotropic temperature distribution and Joule heating in 2D simulations. As a first step towards a fully coupled magneto-thermal evolution, a phenomenological law for the magnetic field decay is considered.  This article is structured as follows. In Sect.~2 we discuss the equations governing the magnetic field structure and evolution, while Sect.~3 is devoted to the thermal evolution equations. Sect.~4 presents the microphysics inputs. Sect.~5 and 6 contain our results for weakly and strongly magnetized NSs, respectively. In Sect. 7, we focus on the effects of field decay and Joule heating on the cooling history of a NS. Finally, in Sect. 8 we present the main conclusions and perspectives of the present work. ",
        "conclusions": "We have presented a thorough study of the thermal evolution of neutron stars including some of the most intriguing effects of magnetic fields. Our results were based on two-dimensional cooling simulations of realistic models that account for the anisotropy in the thermal conductivity tensor. In the first part of the paper,  we revisited the classical scenario with low magnetic fields and presented the input microphysics, working assumptions, and the baseline models. As an interesting byproduct, we reconsidered the growth of the crust and of the superconducting region in the NS core, and found that there are situations in which both growth rates are comparable. The main body of the work was aimed at the discussion of the two principal effects of magnetic fields: the anisotropic surface temperature distribution and the additional heating by magnetic field decay. We found that, even for purely dipolar fields, an inverted temperature distribution is plausible at intermediate ages. Thus the  surface temperature distribution of neutron stars with high magnetic fields, even in the axisymmetric case, may be quite different  from the model with two hot polar caps and a cooler equatorial region. The irregular light curves of some isolated neutron stars, for instance RBS1223, \\citep{Schwope2005,Kaplan2005} are an indication of such complex structures.  The main result of this work is that, in NSs born as magnetars, Joule heating has an enormous effect on the thermal evolution.  Moreover, this effect is important for intermediate field stars. If the magnetic field is supported by crustal currents, this effect can reach a maximum because two combined factors enhance the efficiency of the heating process: i) more heat is released into the crust, in the regions of higher resistivity close to the surface, and ii) large non radial components of the field channel the heat towards the surface, instead of being lost by neutrinos in the core. As expected, it becomes clear that magnetic fields and Joule heating are playing a key role keeping magnetars warm for a long time, but it is likely that the same effect, although quantitatively smaller, must be considered in radio--quiet isolated NSs or high magnetic field radio--pulsars.  Another aspect that should be considered when we try to explain observations using theoretical cooling curves is  that for many objects the age is estimated assuming that the loss of angular momentum is entirely due to dipolar radiation from a magnetic dipole (spin-down age). In the case of a decaying magnetic field, the {\\em spin down age}, seriously overestimates the {\\em true age} \\citep{GO1970}. Therefore, the cooling evolution time should be corrected,  according to the prescription for magnetic field decay, to compare our model accurately with observations. A detailed comparison of the cooling curves  obtained in this work with observational sources can be found in \\cite{Aguilera2008}.   Our last striking remark is that the occurrence of direct URCA or, in general, fast neutrino cooling in NS may be hidden by a combination of effects due to strong magnetic fields. Our conclusion is that the most appropiate candidates to monitor as rapid coolers are NSs with fields lower than $10^{13}$ G. Otherwise, we may be misled in our interpretation of the temperature-age diagrams.  The main drawback of our work is that we are not yet able to return a fully consistent simulation of the magneto-thermal coupled evolution of temperature and magnetic field. In the near future, we plan to extend this study by coupling our thermal diffusion code to the consistent evolution of the magnetic field in the crust given by the Hall induction equation. That approach will permit the accurate evaluation of the heating rates, including the non-linear effects associated with the Hall--drift in the NS crust. We believe, however, that the phenomenological parameterization employed in this paper, reproduces qualitatively the results expected in a real case.  We have provided another step towards understanding the cooling of neutron stars, by pointing out a number of important features that must be more carefully considered in future work."
    },
    "0710/0710.0406.txt": {
        "abstract": "We report results from a search for massive and evolved galaxies at $z \\ga 5$ in the Great Observatories Origins Deep Survey (GOODS) southern field. Combining HST ACS, VLT ISAAC and Spitzer IRAC broad--band photometric data, we develop a color selection technique to identify candidates for being evolved galaxies at high redshifts. The color selection is primarily based on locating the Balmer--break using the K- and 3.6$\\mu$m bands. Stellar population synthesis models are fitted to the SEDs of these galaxies to identify the final sample. We find 11 candidates with photometric redshifts in the range $4.9 \\leq z < 6.5$, dominated by an old stellar population, with ages 0.2$-$1.0 Gyr, and stellar masses in the range $(0.5 - 5) \\times 10^{11}$ M$_{\\odot}$. The majority of the stars in these galaxies were formed at $z > 9$ and the current star formation activity is in all cases, except two, a few percent of the inferred early star formation rate. One candidate has a spectroscopically confirmed redshift, in good agreement with our photometric redshift. The galaxies are very compact, with half--light radii in the observed $K-$band smaller than $\\sim 2$ kpc. Seven of the 11 candidates are also detected at 24$\\mu$m with the MIPS instrument on Spitzer. %%% By itself, the 24$\\mu$m emission could potentially be interpreted as PAH emission from a dusty starburst at $z \\sim 2-3$, however, it is also consistent with the presence of an obscured AGN at $z \\ga 5$. Indeed, for the $z \\ga 5$ solutions, all the MIPS detected galaxies, except two, %BBG\\#3348 and JD2, have relatively high internal extinction. While we favor the obscured AGN interpretation, based on the model SED fits to the optical/UV, we define a 'no--MIPS' sample of candidates in addition to the full sample. Results will be quoted for both samples. % We estimate the completeness of the Balmer break galaxy sample to be $\\sim$40\\% (an upper limit). The comoving number density of galaxies with a stellar mass $\\ga 10^{11}$ M$_{\\odot}$, at an average redshift $\\bar{z} = 5.2$, is $3.9 \\times 10^{-5}$ Mpc$^{-3}$ (no--MIPS sample: $1.4 \\times 10^{-5}$ Mpc$^{-3}$). The corresponding stellar mass density is $8 \\times 10^{6}$ M$_{\\odot}$ Mpc$^{-3}$ (no--MIPS sample: $6.2 \\times 10^{6}$ M$_{\\odot}$ Mpc$^{-3}$). The estimated stellar mass density at $\\bar{z} = 5.2$ is $2-3$\\% of the present day total stellar mass density and $20-25$\\% of the stellar mass density at $z \\sim 2$. If the stellar mass estimates are correct, the presence of these massive and evolved galaxies when the universe was $\\sim$1 Gyr old could suggest that conversion of baryons into stars proceeded more efficiently in the early universe than it does today. ",
        "introduction": "An important goal of observational cosmology is to understand how stars are assembled into galaxies and how this is related to the evolution of dark matter halos. In prevailing hierarchical models, star formation starts out in low mass systems, which build more massive galaxies through sequential merging (e.g. White \\& Rees 1978; Somerville 2004). In this picture, the most massive galaxies are found at relatively low redshifts. Recently, a significant population of galaxies with stellar mass $\\sim 10^{11}$ M$_{\\odot}$ has been found at $z \\sim 2 - 3$ (cf. Franx et al. 2003; Glazebrook et al. 2004; Fontana et al. 2004; Yan et al. 2004; Daddi et al. 2005a; Rudnick et al. 2006; van Dokkum et al. 2006). Stellar population synthesis models combined with broad-band photometric data show that many of these galaxies contain an old stellar population, with ages indicating a star formation phase within 1$-$3 Gyr after the Big Bang. Moreover, a number of submillimeter detected galaxies at $z \\sim 2 - 3$, which are known to be massive systems, based on their inferred molecular gas and dynamical mass estimates (cf. Greve et al. 2005), also appear to contain an old stellar population with mass $\\sim 10^{11}$ M$_{\\odot}$ (Borys et al. 2005). Therefore, a consensus seems to be emerging, that the most massive galaxies seen today, formed the bulk of their stars within the first $\\sim$3 Gyr of cosmic history (cf. Cimatti et al. 2004; Daddi et al. 2005b; Juneau et al. 2005). However, it is not known how these stars were assembled into their present host galaxies, whether this was done during multiple merger events, as proposed in hierarchical models, or if the stars and their host galaxy are co-eval. In view of the early formation epoch implied for many of these massive galaxies, the question whether the formation is hierarchical or monolithic becomes a matter of semantics as the merger time scale becomes comparable to the dynamical time scale.  Recent ultra--deep surveys, done at wavelengths stretching from the UV to mid--infrared, have resulted in detection of galaxies and AGNs at even higher redshifts, reaching into the era of re-ionization. One example is HUDF--JD2, in the Hubble Ultra Deep Field, which Mobasher et al. (2005) identify as a candidate for a massive, evolved galaxy at $z=6.5$. The age of this galaxy is estimated to be $\\ga$600 Myrs, with a stellar mass of $\\sim 6 \\times 10^{11}$ M$_{\\odot}$, much larger than the stellar mass of the Milky Way. The implied age of this galaxy means that  the bulk of the stars were formed on a short time scale just a few hundred million years after the recombination era. Other recent studies have used data from the Spitzer Space Telescope to analyze the stellar masses and ages for galaxies at $z > 5$ (cf. Yan et al. 2005, 2006; Eyles et al. 2005, 2006; Stark et al. 2006; Verma et al. 2007). The inferred stellar masses range from $1 - 10 \\times 10^{10}$ M$_{\\odot}$ and ages of several $10^8$ years. In several cases, galaxies have spectroscopically determined redshifts. Another spectroscopically confirmed galaxy is the gravitationally lensed object HCM06 at $z = 6.56$ (Hu et al. 2002), with a stellar mass of a few $10^{10}$ M$_{\\odot}$ and an age of $\\sim$300 Myr (cf. Chary, Stern \\& Eisenhardt 2005; Schaerer \\& Pell\\'{o} 2005).  The presence of these massive and old galaxies at $z \\ga 5$ holds important clues for understanding how the first galaxies formed and how the galaxy population in general has evolved with cosmic time. In order to determine whether a significant population of massive and old galaxies exists at $z > 5$, and to derive the parameters characterizing this population, we need a selection method that specifically targets and selects evolved stellar systems at very high redshifts, using broad-band photometric data available from deep multiwavelength surveys. The presence of old galaxies at high redshift can not efficiently be inferred using the normal Lyman drop-out technique. The drop-out technique has proven to be efficient in detecting galaxies that are actively forming stars and contain relatively small amounts of dust, but it is ineffective for detecting galaxies without strong UV continuum.  In this paper we develop a method for selecting galaxies dominated by a stellar population older than $\\sim$100 Myr and situated at $z \\ga 5$, and discuss the results and its implications. The technique is primarily based on detecting the presence of a well--developed Balmer break, redshifted to $\\sim$3$\\mu$m,  that can be probed by the K$_\\mathrm{s} - 3.6\\mu$m color index. A second color index is used to further isolate the old high-z galaxies from foreground `contaminants'. The color signature of the Balmer break has previously been used to select galaxies at redshifts $z \\sim 1 - 3$ (Franx et al. 2003; Daddi et al. 2005a; Adelberger et al. 2004). By choosing a suitable filter combination, the Balmer break can be used to select galaxies at any redshift, in a manner similar to the Lyman--break technique. In this paper we will refer to the galaxies selected through this technique at $z > 5$ as Balmer--break galaxies (BBGs).  The paper is structured as follows: % In Sect~\\ref{data} we present our sample and photometric data. % Sect.~\\ref{selecting} discusses the Balmer break feature, the stellar population synthesis models used and examines the confidence of the model fitting procedure using Monte Carlo simulations. In this section we also discuss degeneracies and define the final color selection criteria used in this paper. % In Sect.~\\ref{results} we present the Balmer--break candidates selected using our color criteria and model fits of synthetic stellar spectra. We derive associated physical parameters from the models and discuss the Spitzer/MIPS 24$\\mu$m detections. In this section we also discuss the individual sources and assign a confidence classification to each source based on its likelihood to have the correct redshift. % In Sect.~\\ref{testing} we apply our model fitting to galaxies with known spectroscopic redshift and assess the reliability of the estimated parameters. % In Sect.~\\ref{errors} we discuss different sources of errors and derive the completeness of our sample. % We discuss our results in Sect.~\\ref{discussion} and compare the number density of Balmer--break galaxies with the expected number density of dark matter halos. % Sect.~\\ref{summary} gives a summary of our results. % We adopt H$_0 = 72$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m = 0.3$, and $\\Omega_{\\Lambda} = 0.7$ throughout this paper. All magnitudes are in the AB system (Oke 1974). ",
        "conclusions": "\\label{discussion}  The existence of massive and evolved galaxies at redshifts $z \\ga 5$, when the universe was $\\la 1$ Gyr old, seems surprising at first sight. In the hierarchical scenario for galaxy formation, the majority of massive galaxies are assembled at relatively low redshifts. However, the presence of massive galaxies at high redshifts poses a fundamental problem for hierarchical models only if their number density exceeds that of correspondingly massive dark matter halos (e.g. Somerville 2004). In Sect.~\\ref{completeness} we derived the number and mass density of the Balmer--break galaxies, using our sample of 11 galaxies, as well as for a more restricted sample only containing those candidates which are not detected with MIPS at 24$\\mu$m, the 'no--MIPS sample'.  %Dark matter halos By equating the co--moving number density of the Balmer--break galaxies with the expected density of dark matter halos at the same redshift, we can estimate the the maximum halo mass associated with the BBG's. Using the Sheth--Tormen modified Press--Schechter formalism (Sheth \\& Tormen 1999), with the dark matter halo concentration predicted for the revised value of the power spectrum normalization $\\sigma_{8}=0.74$ (Spergel et al. 2007), and the  estimated lower limit to the number density of BBGs ($3.9 \\times 10^{-5}\\,\\eta^{-1}$ Mpc$^{-3}$), we predict a halo mass of $M_{\\mathrm{h}} = 1.0 \\times 10^{12}$ M$_{\\odot}$ (assuming the standard $\\Lambda$CDM model)\\footnote{The estimated halo mass is a non-linear function of the incompleteness coefficienct $\\eta$. For instance, with $\\eta = 0.5$, the corresponding halo mass becomes $8 \\times 10^{11}$ M$_{\\odot}$.}. Using the average stellar mass for the BBGs, we get $M_{*}/M_{\\mathrm{h}} \\approx 0.20$. Considering the no--MIPS sample, with a number density $1.4 \\times 10^{-5}\\,\\eta^{-1}$ Mpc$^{-3}$, the corresponding halo mass is $1.3 \\times 10^{12}$ M$_{\\odot}$, giving a ratio $M_{*}/M_{\\mathrm{h}} \\approx 0.08$. This estimate of the halo mass assumes that all available $\\sim 10^{12}$ M$_{\\odot}$ halos at $z \\sim 5.2$ are associated with Balmer--break galaxies. If a fraction of these halos would host lower mass stellar systems, such as Lyman--break galaxies, the $M_{*}/M_{\\mathrm{h}}$ ratio for the Balmer--break galaxies would have to increase accordingly.  The ratio of the baryonic--to--total mass can be expressed in terms of a star formation efficiency parameter, $\\beta$ ($M_{*} = \\beta\\,M_{\\mathrm{baryon}}$: the fraction of baryons turned into stars over the life time of the galaxy), and the stellar mass, M$_{*}$, as, $$ M_{\\mathrm{baryon}}/M_{\\mathrm{total}} = \\beta^{-1}M_{*}/M_{\\mathrm{total}} = \\kappa $$ where $M_{\\mathrm{total}} = \\beta^{-1}\\,M_{*} + M_{\\mathrm{h}}$. We can then write $$ M_{*}/M_{\\mathrm{h}} = \\beta\\,\\kappa/(1 - \\kappa). $$ Adopting $\\kappa = 0.17$ from the WMAP3 results (Spergel et al. 2007), we get $M_{*}/M_{\\mathrm{h}} = 0.20\\,\\beta$. Klypin, Zhao \\& Somerville (2002) estimate the total (virial) and baryonic mass of the Milky Way and M31 galaxies and find $M_{*}/M_{\\mathrm{h}} = 0.06 - 0.08$, implying that in this case $\\beta = 0.3 - 0.4$. For the Balmer--break galaxies, we find $\\beta \\approx 0.4 - 1.0$, where the lower value corresponds to the no--MIPS sample. If we only consider the no--MIPS sample, the baryonic fraction is comparable to local galaxies. However, for the full sample, the results suggests that the BBGs at $z \\approx 5.2$ contain a higher fractions of baryons than galaxies at $z \\approx 0$. Another possible explanation for the high baryonic fraction is that the number density of dark matter halos at high redshift is underestimated by the Sheth--Tormen analysis, or that we have systematically overestimated the stellar masses of the BBGs by a factor $\\ga 2$. %%% Using a Chabrier or Kroupa initial mass function with a less steep low--mass end, could lower the estimated stellar masses by a factor 1.5--1.8 (see Sect.~\\ref{systematics}). %%%  It would be more instructive to compare the $M_{*}/M_{\\mathrm{h}}$ ratio to that of massive elliptical galaxies at $z \\approx 0$. However, the evidence for dark matter in elliptical galaxies is still circumstancial and limited to the central regions. Using planetary nebulae and globular clusters as kinematic probes, it has been possible to push the analysis to $\\sim$5 $R_{\\mathrm{eff}}$ (e.g. Romanowsky 2003; Richtler 2004). While the number of ellipticals studied in detail is still small, the general conclusion is that most have surprisingly weak dark matter halos, i.e. large $M_{*}/M_{\\mathrm{h}}$ ratios. It remains to be determined whether the inferred $M_{*}/M_{\\mathrm{h}}$ ratio for Balmer--break galaxies is consistent with local giant elliptical galaxies.   %Stellar mass density The stellar mass density of the universe from redshifts 0 to 6 has been estimated by several groups, using different samples and methods (e.g. Bell et al. 2003; Dickinson et al. 2003; Rudnick et al. 2003, 2006; Fontana et al. 2006; Yan et al. 2006). Some of these results are listed in Table~\\ref{stellarmass}, as a comparison with the results obtained for the BBGs. Most of the values listed in Table~\\ref{stellarmass} are based only on the objects observed and are lower limits. In a few cases, the mass function has been integrated to obtain the total stellar mass (e.g. Dickinson et al. 2003). In the local universe, the global stellar mass density is $(3-4) \\times 10^{8}$\\,M$_{\\odot}$ Mpc$^{-3}$, while it decreases to $\\sim 0.3 \\times 10^{8}$ M$_{\\odot}$ Mpc$^{-3}$ at $z = 2.5-3$. In Sect.~\\ref{completeness} we found that the stellar mass density of the 11 BBG candidates is $8 \\times 10^{6}\\,\\eta^{-1}$ M$_{\\odot}$ Mpc$^{-3}$. This is $\\sim 2-3$\\% of the present day total stellar mass density. Restricting the comparison to large early type galaxies in the local universe, that is, galaxies at least as massive as our BBG sample, the percentage increases to $\\sim 4-6$\\%. Comparing with the stellar mass density at $z \\sim 2$, the BBG sample already comprise $20-25$\\% of the total stellar mass found at this redshift. For the no--MIPS sample, the stellar mass density is $\\sim$5 times smaller, and in this case the comparison with stellar mass densities at lower redshifts has to be corrected accordingly.  The galaxies found in this study are remarkable in that they contain a large stellar mass, have small physical sizes and that their main epoch of star formation occured at $z \\ga 10$. Galaxies with similar properties have, however, also been found by others. In a recent paper, McClure et al. (2006) searched for Lyman--break galaxies in the UKIDSS ultra deep survey, and found 9 candidates with $z > 5$. Their stellar masses are $>5 \\times 10^{10}$ M$_{\\odot}$ and their ages range from 50--500 Myr. Overall, these galaxies have properties similar to our Balmer--break galaxies. The number density for the $z > 5$ galaxies found by McClure et al. is $\\sim$4x smaller than what we find in this paper. However, the different selection process, the fact that the UKIDSS sample does not include IRAC data and the large completeness corrections needed, makes a comparison difficult. A number of massive galaxies at $z > 4$ were also found by Fontana et al. (2006) using the GOODS-MUSIC sample. Their broad--band photometric data set consists of 14 bands, including the 4 IRAC bands. The objects were identified by fitting template SED based on Bruzual \\& Charlot (2003) models to all galaxies in the sample. The best-fitting SEDs for the high redshift objects suggest that they are passively evolving galaxies, characterized by a very short time scale for star formation or by a constant star formation and a large amount of dust extinction. The stellar masses found are in excess of $10^{11}$ M$_{\\odot}$. Hence, massive and passively evolving galaxies at $z \\sim 5$ are found in several studies. A direct comparison of the results is presently not practical as different selection criteria are used, and the completeness corrections are presently poorly defined.  Another surprising property of the Balmer--break galaxies is their compact sizes. As derived in Sect.~\\ref{candidateparameters}, the typical half--light radius is $\\la$2 kpc. Although this is larger than what is expected from the size vs. redshift relation derived for UV bright galaxies at similarly high redshift (Ferguson et al. 2004; Bouwens et al. 2004; Dahlen et al. in prep.), the stellar masses of the Balmer--break galaxies are at least 10 times higher. No massive compact galaxies of this type has been found in the local universe. However, compact galaxies with a large stellar mass have been found at $z \\sim 1.4$ (Trujillo et al. 2006) and at $z\\sim 2.5$ (Daddi et al. 2005b; Zirm et al. 2007; Toft et al. 2007). These galaxies are massive ($M_{*} > 10^{11}$ M$_{\\odot}$), with no sign of AGN activity and contain a passively evolving stellar population, similar to the Balmer--break galaxies. The effective radius of these galaxies, measured at rest--frame optical wavelengths, are typically $\\la 1$ kpc (Zirm et al. 2007; Toft et al. 2007), or 3--6 times smaller than local counterparts of similar stellar mass. It is hypothesized that the on--set of rapid star formation in these systems quench the star formation process, leading to very compact objects. These galaxies, as well as the Balmer--break galaxies, cannot represent fully assembled systems and must undergo subsequent evolution in their structural parameters in order to resemble local galaxies with the same stellar mass.  %Implications for the reionization of the IGM The presence of a population of massive galaxies that underwent a period of intense star formation at $z \\sim 10 - 25$ is likely to have ramifications to the reionization of the intergalactic medium (IGM). Panagia et al. (2005) calculated that the star formation associated with the formation of the massive $z = 6.5$ galaxy HUDF--JD2 (Mobasher et al. 2005), could significantly contribute to the reionization of the IGM. The main uncertainties were the escape fraction of the Lyman continuum photons and the volume density of objects similar to JD2. With the new sample of post--starburst galaxies with formation redshifts in the same range as JD2, it is possible to address this question. The integrated output of Lyman continuum photons from the Balmer--break candidates depends only on the total stellar mass and the assumed IMF (Panagia et al. 2005). Because the average stellar mass for the BBG candidates is about a factor 2 smaller than for JD2, assuming the same IMF, the average number of Lyman continuum photons is likewise a factor of 2 lower. Panagia et al. (2005) concluded that if each field of 2\\ffam5 $\\times$ 2\\ffam5 contained a source like JD2, then these sources account for at least $\\sim$20\\% of the reionization of the IGM. A higher percentage is possible if the escape fraction is higher and/or the IGM is clumped. In the present case, we have a total area that is 25 times larger and  a total ionizing photon output $\\sim$10 times larger than in the case of JD2. Hence, the implication is that the BBG sources can account for $\\sim$10\\% or more of the photons needed for reionization, depending on the poorly constrained parameters describing the Lyman continuum escape fraction and the clumpiness of the IGM itself. The implications for reionization are discussed in more detail in Panagia et al. (in prep.)."
    },
    "0710/0710.1255_arXiv.txt": {
        "abstract": "I~have re-visited the spatial distribution of stars and high-mass brown dwarfs in the $\\sigma$~Orionis cluster ($\\sim$3\\,Ma, $\\sim$360\\,pc). The input was a catalogue of 340 cluster members and candidates at separations less than 30\\,arcmin to $\\sigma$~Ori~AB. Of them, 70\\,\\% have features of extreme youth. I~fitted the normalised cumulative number of objects counting from the cluster centre to several power-law, exponential and King radial distributions. The cluster seems to have two components: a dense core that extends from the centre to $r \\approx$ 20\\,arcmin and a rarified halo at larger separations. The radial distribution in the core follows a power-law proportional to $r^1$, which corresponds to a volume density proportional to $r^{-2}$. This is consistent with the collapse of an isothermal spherical molecular cloud. The stars more massive than 3.7\\,$M_\\odot$ concentrate, however, towards the cluster centre, where there is also an apparent deficit of very low-mass objects ($M <$ 0.16\\,$M_\\odot$). Last, I~demonstrated through Monte Carlo simulations that the cluster is azimuthally asymmetric, with a filamentary overdensity of objects that runs from the cluster centre to the Horsehead~Nebula. ",
        "introduction": "The $\\sigma$~Orionis region in the {Ori~OB~1~b} association is finally becoming recognised as one of the most important young open clusters, with an age of only about 3\\,Ma. In the discovery paper, Garrison (1967) used the term ``clustering'' to refer to an agglomeration of fifteen B-type stars surrounding and including the multiple star {$\\sigma$~Ori}. Afterwards, Lyng\\aa~(1981) tabulated $\\sigma$~Orionis in his catalogue of open clusters. Since the rediscovery of the cluster by Wolk (1996) and its subsequent study in depth, which has revealed the most numerous and best known substellar population (B\\'ejar et~al. 1999; Zapatero Osorio et~al. 2000, 2002; Caballero et~al. 2007), only a few authors have investigated the $\\sigma$~Orionis spatial distribution. In particular, B\\'ejar et~al. (2004) and Sherry et~al. (2004) analysed the radial distribution of $\\sigma$~Orionis cluster members and candidates in annuli of width $\\Delta r$ as a function of the separation $r$ to $\\sigma$~Ori~AB. To maximise the number of objects per annulus and minimise the Poissonian errors, $\\Delta r$ must be wide. This leads to have few annuli (no more than 12 in the $r$ = 0--30\\,arcmin interval) to fit to a suitable radial profile (exponential decay -- B\\'ejar et~al. 2004; King -- Sherry et~al.~2004). Both studies agree that the cluster may extend only up to $\\sim$25--30\\,arcmin. The low surface density of cluster members at larger separations, the sharp increase of extinction due to the nearby {Horsehead Nebula}-{Flame Nebula}-{IC~434} complex and the closeness to (or even overlapping with) other stellar populations in the Orion Belt surrounding {Alnitak} ($\\zeta$~Ori) and {Alnilam} ($\\epsilon$~Ori) prevent from suitably broaden the radial distribution analysis (Caballero 2007a). At the canonical heliocentric distance to $\\sigma$~Orionis of 360\\,pc (e.g. Brown, de~Geus \\& de~Zeeuw 1994), the cluster would have an approximate radius of~3\\,pc.  In spite of the agreement on the size of $\\sigma$~Orionis, the fits and the profiles in B\\'ejar et~al. (2004) and Sherry et~al. (2004) seem to be rather incomplete and inappropiate, respectively. On the one hand, the King models were designed for tidally truncated globular clusters (King 1962, 1966; Meylan 1987), and have also been satisfactorily used for describing galaxies (e.g. Kormendy 1977; Binggeli, Sandage \\& Tarenghi 1984). These systems have had enough time to be isothermal, on the contrary to very young open clusters like $\\sigma$~Orionis, where only gravitational relaxation by initial mixing may have occurred (King 1962). On the other hand, B\\'ejar et~al. (2004) exclusively focused on the cluster substellar population. Besides, the exponential fit in B\\'ejar et~al. (2004) only accounted for the five innermost annuli, which leaded to a high uncertainty in the derived parameters. Last, in the works by Sherry et~al. (2004) and B\\'ejar et~al. (2004), the input list of cluster members and candidates came from $VRI/IZ$ optical surveys. Many sources in both analysis had no near-infrared or spectroscopic follow-up.  For a correct study of the spatial distribution in $\\sigma$~Orionis, it is therefore necessary to use new fitting radial profiles and an input catalogue as comprehensive as possible. It must cover a wide mass interval. Maximum completeness and minimum contamination of the catalogue are also desired. These requirements are verified by the {\\em Mayrit} catalogue, which tabulates 339 $\\sigma$~Orionis members and candidates in a 30\\,arcmin-radius circular area centred on $\\sigma$~Ori~AB (Caballero 2007c). Of them, 241 display features of extreme youth (e.g. OB spectral types, Li~{\\sc i} in absorption, H$\\alpha$ in strong emission, spectral signatures of low gravity, near- and mid-infrared excesses due to discs). The catalogue covers three orders of magnitude in mass, from the $\\sim$18+12\\,$M_\\odot$ of the O9.5V+B0.5V binary $\\sigma$~Ori~AB to the $\\sim$0.033\\,$M_\\odot$ of the brown dwarf {B05~2.03--617} (Caballero \\& Chabrier, in~prep.). Accounting for $\\sigma$~Ori~A and B as different objects separated by $\\sim$0.25\\,arcsec, then the equatorial coordinates of 340 young stars, brown dwarfs and cluster member candidates are available. I will use this input catalogue to investigate the radial and azimuthal distribution of objects in the $\\sigma$~Orionis~region. ",
        "conclusions": "The $\\sim$3\\,Ma-old $\\sigma$~Orionis cluster is a perfect laboratory of star formation. I~have investigated the radial distribution of 340 cluster members and candidates in a 30\\,arcmin-radius area centred on $\\sigma$~Ori~AB, taken from Caballero (2007c). The analysis has covered a mass interval from the 18+12\\,$M_\\odot$ of $\\sigma$~Ori~AB to the $\\sim$0.03\\,$M_\\odot$ of the faintest brown dwarf detectable by DENIS. The cluster shows a clear radial density gradient, quantified by the ${\\mathcal Q}$-parameter, that accounts for the mean separation between members and the Euclidean minimum spanning tree of the cluster. I~have calculated the functional relations between normalised cumulative numbers of objects counting from the cluster centre, $f(r)$, and surface densities, $\\sigma(r)$. Cumulative distribution functions as these avoid many problems associated with binning. Among the studied radial (power-law, exponential and King) profiles, the best fit is for a composite power-law distribution of cluster members with a core and a rarified halo. The core extends up to $\\sim$20\\,arcmin from the cluster centre and is nicely modelled by a surface density $\\sigma(r) \\propto r^{-1}$, that corresponds to a volume density $\\rho(r) \\propto r^{-2}$. This volume density matches, in its turn, the radial profile in a cluster formed from the collapse of a self-gravitating, isothermal sphere. The most massive $\\sigma$~Orionis stars deviate, however, from the general trend and are much more concentrated towards the cluster centre. There is also an apparent deficit of very low-mass stars and high-mass brown dwarfs (0.16\\,$M_\\odot \\ga M \\ga$ 0.035\\,$M_\\odot$) in the innermost 4\\,arcmin and an excess in the annulus at 6--7\\,arcmin to the central Trapezium-like system. Last, there is a significant azimuthal asymmetry due to a filament-shape overdensity of objects that connects the cluster centre with a part of the Horsehead Nebula. This discovery supports the formation scenarios that predict burst of star formation in filamentary~gas."
    },
    "0710/0710.3911_arXiv.txt": {
        "abstract": "% We follow the evolution of helium stars of initial mass $(2.2 - 2.5)\\,M_\\odot$, and show that they undergo off-center carbon burning, which leaves behind ${\\mathbf \\sim 0.01\\,M_\\odot}$ of unburnt carbon in the inner part of the core. When the carbon-oxygen core grows to Chandrasekhar mass, the amount of left-over carbon is sufficient to ignite thermonuclear runaway. At the moment of explosion, the star will possess an envelope of several $0.1\\,M_{\\odot}$, consisting of He, C, and possibly some H, perhaps producing a kind of peculiar SN. Based on the results of \\citet{Waldman2007} for accreting white dwarfs, we expect to get thermonuclear runaway at a broad  range of $\\rho_c \\approx (1 - 6) \\times 10^9 \\mathrm{ g\\,cm^{-3}}$, depending on the amount of residual carbon. We verified the feasibility of this scenario by showing that in a close binary system with initial masses $(8.5 + 7.7)\\,M_{\\odot}$\\ and initial period of 150 day the primary produces a helium remnant of $2.3\\,M_{\\odot}$\\ that evolves further like the model we considered. ",
        "introduction": "Type Ia supernovae (SN~Ia) have a relatively small dispersion of luminosity (the standard deviation in peak blue luminosity is $\\sigma_B \\approx 0.4 - 0.5$ mag., \\citet{Branch1993ApJL}) and are being used as distance indicators (``standard candles''), having especial significance in the effort of determining the cosmological parameters of our universe.  The long-standing explanation of the SN~Ia phenomenon has been the explosive burning of degenerate carbon in the core of a carbon-oxygen white dwarf, which becomes unstable as it grows to Chandrasekhar mass ($M_{Ch}$) either by accretion from a binary companion or by a merger of two white dwarfs, following the angular momentum loss from the system by gravitational wave radiation. However, theoretical models are still far from self-consistently producing an evolutionary path towards the progenitor and reproducing crucial features of the observational data, such as the composition of the ejecta. For a detailed review of the above see, e.g., \\citet{Leibundgut2000A&ARv, Hillebrandt&Niemeyer2000ARA&A, Filippenko2005ASSL}.  As well, SN~Ia can not be regarded as perfectly homogeneous class, since their Hubble diagram  exhibits scatter larger than the photometric errors, while spectroscopic and photometric peculiarities have been noted with increasing frequency in well-observed SN~Ia \\citep[e.g., ][]{Filippenko2005ASSL}.  Therefore, there is an obvious need for progenitor scenarios that could explain the diversity among SN~Ia. Several explanations have been suggested, such as variations in the metallicity of the progenitor, in the carbon to oxygen ratio at its center, or in the central density at the time of ignition \\citep[e.g., ][]{Timmes2003ApJ, Ropke_etal2006A&A, Lesaffre2006MNRAS}. The variation of the latter two parameters is expected to result from the variation in the initial white dwarf mass and in the accretion history.  In this work we follow the evolution of helium stars with initial mass $\\approx (2.2 - 2.5)\\,M_\\odot$ and show that they might reach thermonuclear explosion and perhaps account for some of the peculiar SNe. ",
        "conclusions": ""
    },
    "0710/0710.5126_arXiv.txt": {
        "abstract": "Our campaign of deep monitoring observations with {\\it Chandra} of the nearby elliptical galaxy NGC 3379 has lead to the detection of nine globular cluster (GC) and 53 field low mass X-ray binaries (LMXBs) in the joint {\\it Hubble}/{\\it Chandra} field of view of this galaxy. Comparing these populations, we find a highly significant lack of GC LMXBs at the low (0.3-8~keV) X-ray luminosities (in the $\\sim 10^{36}$ to $\\sim 4\\times10^{37}$ erg s$^{-1}$ range) probed with our observations. This result conflicts with the proposition that all LMXBs are formed in GCs. This lack of low-luminosity sources in GCs is consistent with continuous LMXB formation due to stellar interactions and with the transition from persistent to transient LMXBs. The observed cut-off X-ray luminosity favors a predominance of LMXBs with main-sequence donors instead of ultra-compact binaries with white-dwarf donors; ultra-compacts could contribute significantly only if their disks are not affected by X-ray irradiation. Our results suggest that current theories of magnetic stellar wind braking may work rather better for the unevolved companions of GC LMXBs than for field LMXBs and cataclysmic variables in the Galaxy, where these companions may be somewhat evolved. ",
        "introduction": "Since their discovery in the Milky Way (see Giacconi 1974), the origin and evolution of Low-mass X-ray binaries (LMXBs) has been the subject of much discussion. LMXBs are found in both the stellar field and globular clusters (GCs). Their incidence per unit stellar mass is much higher in GCs than in the field, requiring a special formation mechanism, presumably dynamical (Clark 1975; Katz 1975). The high efficiency of these dynamical processes  led to the suggestion that all LMXBs may form in GCs and then disperse in the field (see e.g., Grindlay 1984; Grindlay \\& Hertz 1985); however, some primordial field binaries are also expected to evolve into LMXBs, suggesting that there may be two populations of these sources, with distinct formation and evolutionary histories (see review by Verbunt \\& van den Heuvel 1995).  The discovery of X-ray source populations in early-type galaxies with {\\it Chandra} has provided a wider and more diverse observational basis for the study of LMXBs, their association with GCs and the role that GCs may play in LMXB formation (see review Fabbiano 2006). However, until recently, only the most luminous extra-Galactic LMXBs with ($\\sim$0.3-8~keV) $L_X \\geq$ a few $10^{37}$erg s$^{-1}$ have been observed, and therefore the study of the GC-LMXB association has been limited to systems with X-ray luminosities in the upper range of Galactic LMXB luminosities.  With our deep 337 ks {\\it Chandra} ACIS-S3 observations of the  unperturbed elliptical galaxy NGC 3379 in the nearby poor group Leo (D$\\sim$11 Mpc), we can now pursue this study in a  luminosity range more typical of the well-studied Galactic LMXBs. In NGC 3379, The GC-LMXB association has been previously studied by Kundu, Maccarone \\& Zepf (2007, hereafter KMZ), using the first shorter {\\it Chandra} observation of this galaxy, which has a typical source detection threshold of $\\sim 1-2 \\times 10^{37}$ erg s$^{-1}$  (KMZ). Our data set is $\\sim$10 times deeper (337 ks), allowing the detection of LMXBs at luminosities of $\\sim 10^{36}$ erg s$^{-1}$. ",
        "conclusions": "Our campaign of deep monitoring observations of the nearby elliptical galaxy NGC 3379 with {\\it Chandra} ACIS-S3 has led to the detection of  nine  GC LMXBs in the field studied optically with {\\it Hubble} WFPC2,  seven of which were previously detected in a much shorter {\\it Chandra} exposure (1/10 of the exposure time; KMZ). The comparison of GC and field LMXB statistics in the joint {\\it Hubble}-{\\it Chandra} field of view demonstrates a relative lack of GC LMXB at luminosities below $\\sim 4\\times10^{37}$ erg s$^{-1}$; field and GC LMXBs instead are know to have closely matching XLFs above this luminosity (Kim E. et al 2006; KMZ). The dearth of low-luminosity GC LMXBs in NGC 3379 is consistent with a similar suggested behavior in NGC 3115 (KMZ), and with the XLFs of field and GC LMXBs in the Milky Way and M31 (Voss \\& Gilfanov 2007). These differences between low-luminosity field and GC XLFs falsify suggested theories that {\\it all} LMXBs may have been originated in GCs.  The luminosity-dependent differences of field and GC XLFs cannot be explained by high luminosity GCs containing multiple LMXBs, because we find clear evidence of source variability for seven out of the nine sources invalidating this hypothesis. Persistent behavior of high luminosity GC sources, compared with transient field sources of similar high luminosity may explain the discrepancy as excess of luminous GC LMXBs. However, the detection of three candidate transient source (with peak luminosity greater than $10^{37}$ erg s$^{-1}$) in the GC LMXB population of NGC 3379 may not support this explanation.  The lack of low-luminosity sources in GCs is consistent with the prediction of Bildsten and Deloye (BD4) based on the transition of sources from persistent to transients due to the thermal disk instability. However, the value of the observed XLF cut-off is not  consistent with their suggestion that GC LMXBs are dominated by ultra-compact binaries, but instead favors LMXBs with H-rich MS donors. The $\\sim 4\\times10^{37}$ erg s$^{-1}$ luminosity cut-off is also consistent with current theories of magnetic stellar wind braking, suggesting that this effect may work rather better for the unevolved companions of GC LMXBs than for field LMXBs and cataclysmic variables in the Galaxy, where these companions may be somewhat evolved.  While our results firmly establish a dearth of GC sources in NGC3379 at low luminosity, the accurate luminosity (and uncertainty) of the GC XLF  cut-off will need the formal analysis of the NGC3379 LMXB luminosity function (Kim et al in preparation). The forthcoming analysis of the very deep {\\it Chandra} observations of NGC4278, a GC-rich galaxy (see Kim, D.-W. et al 2006), will provide in the near future stronger constraints on the potential 'universality' and value of the GC LMXB XLF cut-off luminosity."
    },
    "0710/0710.0968_arXiv.txt": {
        "abstract": "{Gamma-ray binaries have been established as a new class of sources of very high energy (VHE, $>$100~GeV) photons. These binaries are composed of a massive star and a compact object. The gamma-rays are probably produced by inverse Compton scattering of the stellar light by VHE electrons accelerated in the vicinity of the compact object. The VHE emission from LS~5039 displays an orbital modulation. } {The inverse Compton spectrum depends on the angle between the incoming and outgoing photon in the rest frame of the electron. Since the angle at which an observer sees the star and electrons changes with the orbit, a phase dependence of the spectrum is expected.} {A procedure to compute anisotropic inverse Compton emission is explained and applied to the case of \\ls. The spectrum is calculated assuming the continuous injection of electrons close to the compact object: the shape of the steady-state distribution depends on the injected power-law and on the magnetic field intensity.} {Compared to the isotropic approximation, anisotropic scattering produces harder and fainter emission at inferior conjunction, crucially at a time when attenuation due to pair production of the VHE gamma-rays on star light is minimum. The computed lightcurve and spectra are very good fits to the HESS and EGRET observations, except at phases of maximum attenuation where pair cascade emission may be significant for HESS. Detailed predictions are made for a modulation in the GLAST energy range. The magnetic field intensity at periastron is 0.8$\\pm0.2$~G. } {The anisotropy in inverse Compton scattering plays a major role in \\ls.  A simple model reproduces the observations, constraining the magnetic field intensity and injection spectrum.  The comparison with observations, the derived magnetic field intensity, injection energy and slope suggest emission from a rotation-powered pulsar wind nebula. These results confirm gamma-ray binaries as promising sources to study the environment of pulsars on small scales.} ",
        "introduction": "Gamma-ray binaries have been established in the past couple of years as a new class of sources of very high energy (VHE, $>$100~GeV) photons. They are characterized by a large gamma-ray luminosity above an MeV, at the level of or exceeding their X-ray luminosity. At present, all three such systems known (\\object{\\ls}, \\object{\\psrb}\\ and \\object{\\lsi}, recently possibly joined by \\object{Cyg X-1}) comprise a massive star \\citep{2005Sci...309..746A,2005A&A...442....1A,2006Sci...312.1771A,2007arXiv0706.1505M}. The compact object in \\psrb\\ is a 48-ms, young radio pulsar. The VHE emission arises from the interaction of the relativistic wind from this pulsar, extracting rotational energy from the neutron star, with the stellar wind from its companion \\citep{1994ApJ...433L..37T}. Particles gain energy at the shock between the winds, resulting in a small-scale pulsar wind nebula \\citep{1981MNRAS.194P...1M}. The particles radiate away their energy as they are entrained in the shocked flow, forming a comet-like trail of emission behind the pulsar \\citep{2006A&A...456..801D}.  The nature of the compact object and origin of the VHE emission remains controversial in \\ls\\ and \\lsi, although recent observations indicate the radio emission of \\lsi\\ behaves like the comet tail expected in the pulsar scenario \\citep{2006smqw.confE..52D}. Alternatively, the VHE emission could originate from particles accelerated in a relativistic jet, the energy source being accretion onto a black hole or neutron star \\citep{2006ApJ...643.1081D,2006A&A...451..259P}. The rationale being that there is evidence for particle acceleration in the jets of microquasars and active galactic nuclei.  However, hard evidence for accretion occuring in either \\ls\\ or \\lsi\\ has been hard to come by \\citep[e.g.][]{2005A&A...430..245M} and the similarities between the three systems (and differences with the usual microquasars) do not argue in favour of the accretion/ejection scenario \\citep{2006A&A...456..801D}.  Regardless of the actual powering mechanism, some particles must be accelerated to high energies to generate the VHE gamma-rays. If these particles are leptons, the only viable gamma-ray radiation mechanism is inverse Compton scattering on the stellar photons.  The massive stars in gamma-ray binaries have effective temperatures of several tens of thousand K and radii of about 10~$R_\\odot$, yielding luminosities of the order of $10^{39}$~erg~s$^{-1}$. This provides a huge density of stellar photons in the UV band that VHE leptons may up-scatter, much greater than any other possible source of target photons (e.g. synchrotron or bremsstrahlung emission).  The emitted VHE photons also have enough energy to produce $e^+e^-$ pairs with the UV stellar photons. Most of the VHE flux may therefore be lost to the observer if the source is behind the star and VHE photons have to travel through the stellar light. Gamma-ray attenuation has been shown to lead to a modulation of the VHE flux with minimum absorption (maximum) at inferior (superior) conjunction \\citep{2005ApJ...634L..81B,2006A&A...451....9D}.  HESS observations have indeed shown a stable modulation of the VHE flux from \\ls\\ on the orbital period with a maximum around inferior conjunction. This suggests attenuation plays a role and that the source of VHE gamma-rays cannot be more than about an AU from the binary (or attenuation would be too weak to modulate the flux). However, a non-zero flux is detected at superior conjunction where a large attenuation is expected, possibly because of pair cascading. Moreover, the spectral changes that are reported do not fit with an interpretation based on pure attenuation of a constant VHE source spectrum \\citep{2006A&A...460..743A}.  Inverse Compton scattering also has a well-known dependence on the angle $\\Theta$ between incoming and outgoing photon.  The photon flux from the star being anisotropic, the resulting inverse Compton emission will depend on the angle at which it is observed. Hence, a phase-dependent VHE spectrum will be observed even if the distribution of particles is isotropic and remains constant throughout. This effect has previously been investigated in \\psrb\\ by \\citet{2000APh....12..335B} who calculated the radiative drag on the unshocked pulsar wind from scattering of stellar light, using results from \\citet{1989ApJ...343..277H}.  The drag produces a Compton gamma-ray line with a strong dependence on viewing angle.  This work purports to explain the HESS observations of LS~5039 using a combination of anisotropic inverse Compton scattering and attenuation in the simplest way possible. The aim is to constrain the underlying particle distribution and/or powering mechanism. \\S2 derives the main equations governing anisotropic Compton scattering in the context of gamma-ray binaries and discusses the principal characteristics to expect. \\S3 presents the application to the case of LS~5039. The lightcurve and spectra observed by the HESS collaboration are reproduced by a model taking into account the photon field anisotropy and the attenuation due to pair creation. \\S4 concludes on the origin of the VHE emission from this system. ",
        "conclusions": "The anisotropic behaviour of inverse Compton scattering has a major influence on the emission from gamma-ray binaries. In these sources, the massive star provides a large source of seed photons with energies around an electron-volt. If high energy electrons are accelerated in the vicinity of the compact object, then the angle between the star, compact object and observer changes with orbital phase. The variation in viewing angle leads to a strong modulation in both the intensity and spectral shape of the scattered radiation.  Scattering stellar photons to the TeV range requires very energetic electrons with Lorentz factors $\\gamma_{\\rm e}\\approx 10^6-10^7$. The scattering therefore occurs in the Klein-Nishina regime. In this case, the anisotropy results, at inferior conjunction, in a harder and fainter spectrum than predicted using an isotropic approximation for the incoming photons.  Crucially, inferior conjunction also corresponds to the phase at which the produced VHE gamma-rays are less attenuated by pair production on stellar photons.  At other phases the emitted spectrum is close to the one obtained using the isotropic photon field approximation and can be severely attenuated by pair production. The result is a complex interplay that reduces the amplitude of the variations expected from a pure attenuation model and a hardening at inferior conjunction.  The \\ls\\ lightcurve and spectra were modelled using a simple-minded leptonic model. The electrons are assumed to be accelerated efficiently in a small zone in the vicinity of the compact object with a standard $\\gamma_{\\rm e}^{-p}$ power-law. Radiative losses due to inverse Compton emission and synchrotron emission generate a distinctive steady-state electron distribution in this environment dominated by stellar photons. The distribution has a prominent hardening between the energy at which inverse Compton losses enter the Klein-Nishina regime ($\\gamma_{\\rm KN}\\approx 6~10^4$ in \\ls) and the energy at which synchrotron losses take over ($\\gamma_{\\rm S}\\approx 10^7$ for a 1~G field). This is for instance the distribution found in the vicinity of the pulsar wind shock but it applies equally well to any leptonic model where particles are accelerated close to the compact object. The magnetic field was allowed to vary as the inverse of the orbital separation, as expected from a pulsar wind nebula. The model has only three parameters: the intensity of the magnetic field, the normalization of the electron distribution and the slope $p$ of the injected power-law $\\gamma_{\\rm e}^{-p}$.  The cutoff in the very high energy gamma-ray spectrum is very sensitive to the magnetic field intensity, via the location of $\\gamma_{\\rm S}$ in the electron distribution. Fitting the high-state spectrum seen by HESS gives a rather constrained magnetic field intensity at periastron of 0.8$\\pm0.2$~G. This value compares well with the values found using simple pulsar wind models which give 5 $(\\dot{E}_{36} \\sigma_3)^{1/2} R_{11}^{-1}$~G, where $\\dot{E}_{36}$ is the pulsar spindown power in units of 10$^{36}$~erg~s$^{1}$, $\\sigma_3$ is the ratio of magnetic to kinetic energy in the pulsar wind in units of 10$^{-3}$ and $R_{11}$ is the distance of the shock to the pulsar in units of 10$^{11}$~cm. Fitting the HESS high-state spectrum also sets the injection slope to $p=2\\pm0.3$, close to the canonical value for shock acceleration. The normalization of the electron distribution implies an injection rate of 10$^{36}$~erg~s$^{-1}$ for a radiative zone of 3~$10^{11}$~cm. These results are remarkably consistent with the expectations for a pulsar wind model.  The spectrum is also found to fit extremely well the EGRET observations, adding credence to the reliability of this simple approach. The model predicts a strong variation in the GLAST band with a softening from high to low flux below a GeV (where synchrotron emission dominates the spectrum) but a hardening above a GeV (where inverse Compton emission dominates the spectrum). The HESS low-state spectrum is not explained to satisfaction. The model fits nicely the EGRET measurements but produces too many gamma-rays at 5-10~TeV. A possible solution is a more complex orbital phase-dependence of the electron distribution at selected phases.  Another solution is that the low-state spectrum corresponds to phases of strong attenuation and that emission from the created pairs contribute significantly to the spectrum. Additional HESS observations near minimum flux would be welcomed.  The orbital modulation of the HESS emission is easily reproduced. A well-defined peak is predicted between phases 0.7-0.9 for which evidence may already be seen in the data. The lightcurve at GLAST energies is anti-correlated with the HESS lightcurve and has a peak at periastron, where the stellar photon density is maximum, and a minimum at inferior conjunction because of the anisotropic effects in inverse Compton scattering. The GLAST spectrum below 1~GeV should be influenced by the tail of the synchrotron emission from the highest energy electrons. The peak synchrotron emission is at about $100$~MeV for maximally accelerated electrons, regardless of magnetic field. Hence, if this component is detected, it will provide evidence that electrons are indeed accelerated with extreme efficiency in this source.  Similar results for the magnetic field intensity and particle energy are found when a lower inclination is used, i.e. implying a black hole compact object rather than a neutron star. In this case, the emission is thought to arise from a relativistic jet powered by accretion onto the black hole. Within the assumptions of this work on the particle distribution, it is difficult to argue that a significant part of the emission occurs far along a jet since this does not naturally reproduce neither the spectrum nor the lightcurve measured by HESS. Most of the emission should still occur close to the compact object. However, unlike in the case of a pulsar wind nebula, there is no independent theoretical expectations in support of the magnetic field intensity (certainly smaller than its equipartition value in the accretion flow) and particle energy that are derived. Therefore, the pulsar wind nebula model appears favoured independently of other possible considerations.  Despite the complexity of the phenomena involved in pulsar wind nebula emission, it is found that the peculiar environment of a gamma-ray binary, most prominently the enormous luminosity of the massive companion, severely constrains the number of degrees-of-freedom in the model. A simple model suffices to reproduce most of the observations. The value of the magnetic field at the shock is found to be tightly constrained by the HESS observations to 0.8$\\pm$0.2~G and the injection spectrum slope to $p=2\\pm0.3$. These results confirm that gamma-ray binaries are promising sources to study the environment of pulsars on very small scales."
    },
    "0710/0710.3193_arXiv.txt": {
        "abstract": "In this paper a neutron star with an inner core which undergoes a phase transition, which is characterized by conformal degrees of freedom on the phase boundary, is considered. Typical cases of such a phase transition are e.g. quantum Hall effect, superconductivity and superfluidity. Assuming the mechanical stability of this system the effects induced by the conformal degrees of freedom on the phase boundary will be analyzed. We will see that the inclusion of conformal degrees of freedom is not always consistent with the staticity of the phase boundary. Indeed also in the case of mechanical equilibrium there may be the tendency of one phase to swallow the other. Such a shift of the phase boundary would not imply any compression or decompression of the core. By solving the Israel junction conditions for the conformal matter, we have found the range of physical parameters which can guarantee a stable equilibrium of the phase boundary of the neutron star. The relevant parameters turn out to be not only the density difference but also the difference of the slope of the density profiles of the two phases. The values of the parameters which guarantee the stability turn out to be in a phenomenologically reasonable range. For the parameter values where the the phase boundary tends to move, a possible astrophysical consequence related to sudden small changes of the moment of inertia of the star is briefly discussed. ",
        "introduction": "One of the most intriguing features of neutron stars is the possibility of having an inner core which undergoes a (colour) superconductivity and/or superfluidity phase transition due to the extremely high pressure and density. This fact was first pointed out by Migdal \\cite{Mi59} (there is a huge amount of literature on this subject with a little hope of providing one with a complete list of references: for two detailed reviews, see\\ \\cite% {Pet92} \\cite{Ra99} and references therein; for general relativistic formalisms suitable for dealing with neutron stars within these scenarios see \\cite{CL98} \\cite{AC07} and references therein). Such phase transitions inside neutron stars could lead to interesting observational effects relevant in cosmology \\cite{Sigl:2006ur} and through the quasi normal modes of the stars as well as the related gravitational radiation (see, for instance, \\cite{ST02} \\cite{MP03}; detailed reviews on the subjects are \\cite% {KS99} \\cite{Ste03} \\cite{ST03}). Recently, there has been also pointed out the intriguing possibility to have a quantum hall phase of gluons and quarks in the inner core of a neutron star (see, in particular, \\cite{IM05} \\cite{Iw05}% ). The discovery of \\textit{magnetars} \\cite{DT92}, highly magnetized neutron stars whose magnetic fields can reach $10^{18}G$ in the core, makes the arising of Quantum Hall features in the inner core of a neutron star not unlikely. All the here described types of phase transitions have in common to be characterized by the presence of conformal degrees of freedom on the phase boundary predicted by QFT\\footnote{% The fact that the boundary degrees of freedom are conformal can be seen \"heuristically\" in the case of a superconducting phase. Due to the Meissner effect, the current must circulate on the phase boundary of the superconductive phase. Phenomenologically, there is no dissipation, this means that there is no physical scale over which the boundary excitations can dissipate energy. The lack of a characteristic scale for the boundary theory is related to conformal symmetry. In the case of Quantum Hall Effect, this can be easily seen from the fact that it is well described by a Chern-Simons action in the bulk. Thus, the theory induced on the boundary is the Wess-Zumino-Witten action which also has conformal symmetry.} (see, for instance, \\cite{Wi90} \\cite{We96}).         The main goal of this paper is to study under which conditions the phase boundary inside a neutron star, which undergoes such a phase transition with conformal boundary degrees of freedom, can be static in general relativity, assuming mechanical equilibrium of the system\\footnote{ The analysis of the mechanical stability of neutron stars has been performed in a huge number of papers.  To provide one with a complete list of references is a completely hopeless task; representative papers are, for instance, \\cite{EKO91} \\cite% {ABR00} \\cite{CL98} and references therein)}. Indeed also if the system is mechanically stable i.e. there is no pressure inducing the collapse or expansion of the core, there can be another sort of instability: there may be the tendency of one phase to prevail over the other near the boundary. This would imply that the boundary gets shifted. This situation is analogous to a bubble chamber where a small perturbation induces a phase transition. To study this phenomenon  one must take into account the contribution of the nontrivial traceless boundary stress tensor to the Einstein field equations. This problem is equivalent to solving Israel's junction conditions. It is necessary to assume a dynamical phase boundary and to check under which conditions such a configuration is a solution of the Einstein field equations (by solving the junction conditions). It turns out that the equation of motion for the dynamical phase boundary is equivalent to the dynamics of a classical point particle in an effective potential. There exist a static stable equilibrium configuration if the effective potential has a local minimum for a suitable negative value of the potential. It is shown that the existence of a local minimum is regulated both by the difference in the mass density between the two phases and by difference of the slope of the two density profiles. A third parameter which determines the form of the potential is the energy of the shell. The non trivial effects of conformal boundary degrees of freedom have been pointed out, in the context of black hole physics, in \\cite% {Cao7} in which they are related to the arising (after the event horizon is formed) of Quantum Hall features \\cite{Lau99} \\cite{Wi90}\\ related to the strong attractive nature of the gravitational field acting on Fermions inside a collapsed neutron star.  The structure of the paper is the following: in Section 2, the assumptions of the present paper are explained. In Section 3 the junction conditions derived from the Einstein equations for a neutron star with a phase boundary are solved. In section 4 the range of the physical parameters, characterizing our configuration, for which the phase boundary is in stable equilibrium is determined. In section 5 an astrophysical implication is discussed. Section 6 the conclusions and perspectives are drawn. ",
        "conclusions": "We have studied a model of a neutron star with an inner core which undergoes a phase transition. It has been shown that, even assuming mechanical equilibrium, it is not always consistent to assume a static phase boundary once conformal boundary degrees of freedom are taken into account.   In order to study the staticity of such a phase boundary one must take into account the general relativistic effects of these boundary degrees of freedom by including a nontrivial stress tensor on the junction. To the best of the authors knowledge, such a consistency analysis with conformal boundary degrees of freedom has not been performed previously. The astrophysical consequences of the non-static regime, related to sudden changes of the moment of inertia of the star, are worth further investigating."
    },
    "0710/0710.1196_arXiv.txt": {
        "abstract": "HI features near young star clusters in M81 are identified as the photodissociated surfaces of Giant Molecular Clouds (GMCs) from which the young stars have recently formed. The HI column densities of these features show a weak trend, from undetectable values inside $R = 3.7$\\ kpc and increasing rapidly to values around $3 \\times 10^{21}\\ \\mbox{cm}^{-2}$ near $R \\approx 7.5$\\ kpc. This trend is similar to that of the radially-averaged HI distribution in this galaxy, and implies a constant area covering factor of $\\approx 0.21$ for GMCs throughout M81. The incident UV fluxes \\Gnaught\\ of our sample of candidate PDRs decrease radially.  A simple equilibrium model of the photodissociation-reformation process connects the observed values of the incident UV flux, the HI column density, and the relative dust content, permitting an independent estimate to be made of the total gas density in the GMC. Within the GMC this gas will be predominantly molecular hydrogen. Volume densities of $1 < n < 200\\ \\mbox{cm}^{-3}$ are derived, with a geometric mean of $17\\ \\mbox{cm}^{-3}$. These values are similar to the densities of GMCs in the Galaxy, but somewhat lower than those found earlier for M101 with similar methods. Low values of molecular density in the GMCs of M81 will result in low levels of collisional excitation of the CO(1-0) transition, and are consistent with the very low surface brightness of CO(1-0) emission observed in the disk of M81. ",
        "introduction": "\\label{sec:intro}  Giant molecular clouds (GMCs) in the interstellar medium (ISM) are generally accepted as the birthplaces of new stars. The most massive of these new stars produce copious amounts of far ultraviolet (FUV) radiation which will, in turn, photodissociate the molecular gas of the parent GMCs, producing ``blankets'' of HI (and other atomic species) on their surfaces.  \\citet{all1986} were the first to present evidence that major features in the HI distribution of the nearby spiral galaxy M83 (NGC 5236), namely, the inner HI spiral arms observed with the Very Large Array (VLA), were the result of photodissociation of \\Htwo on galactic scales. \\citet{all1997} confirmed that HI features existed in M81 (NGC 3031) on scales of $\\approx 150$ pc which were qualitatively consistent with the expected morphology of large, low density photodissociation regions (PDRs), and explicitly related those HI features to nearby bright sources of Far-UV (FUV) radiation found on images of the galaxy from the Ultraviolet Imaging Telescope (UIT).  \\citet{smi2000} used a simple, but quantitative, model for the equilibrium physics of photodissociation regions and applied it to VLA-HI and UIT-FUV observations of M101 (NGC 5457). Their work showed that this approach provides an entirely new method for determining the volume density of molecular gas in star-forming GMCs of galaxies, a method which is no longer dependent on the (often poorly known) excitation conditions for line emission by specific molecular tracers.  In this paper we return to another spiral galaxy, M81, and carry out a quantitative analysis of the GMCs in this galaxy using the methods described by \\citet{smi2000}. This new analysis has been made possible by new data on M81 which was not available to \\citet{all1997}; new, higher-resolution, high-sensitivity VLA-HI observations, and sensitive new FUV imagery from the GALEX satellite.  We have also sought an independent verification that the very basis of our approach is valid, namely, that the HI in the immediate vicinity of FUV concentrations is indeed produced in PDRs. To this end we have examined the Spitzer/IRAC data on M81 for evidence of mid-IR emission by polycyclic aromatic hydrocarbons (PAHs) which are also thought to be tracers of PDRs.  M81 has been probed extensively for the molecular tracer CO. No global CO map of M81 has been published to this date, and the detected CO emission is found to be very weak and spotty (see e.g. \\citet{kna2006} and references therein). Since PDRs also produce CO emission \\citep[e.g.][]{all2004}, comparing the CO results to the results in this paper is therefore of considerable interest.  The outline of this paper is as follows: Section \\ref{sec:data} contains a description of the data we used, followed by a brief theoretical description of PDRs in Section \\ref{sec:theory} and the application of the method in Section \\ref{sec:method}. The results are presented in Section \\ref{sec:results}. The results are discussed and the conclusions are briefly summarized in Section \\ref{sec:discussionconclusions}. ",
        "conclusions": "\\label{sec:discussionconclusions}  The candidate PDRs in M81, which were selected on the basis of their FUV emission, seem to fit the photodissociation model well. Our results show no systematically different properties of the parent GMCs in different parts of M81. The total hydrogen volume density is roughly constant, even as the underlying HI, FUV and dust-to-gas ratio vary.  The cloud densities we find are lower than the range of values (30 - 1000 cm$^{-3}$) found in M101 \\citep{smi2000}. The observed values of $N_{HI}$ of individual PDRs are similar to those seen in M101. The observed HI columns in both galaxies abruptly increase at the same normalized radius (0.3) and appear to decline somewhat beyond 0.7, assuming an $R_{25}$ of 30.3 kpc for M101 \\citep{vil1992}.  The range of observed HI columns is also the same. The downward trend in \\Gnaught\\ also is consistent with \\citet{smi2000}, when no internal extinction correction is applied to their data. Figure \\ref{fig:Rnorm_n} shows that the densities of the GMCs in M81 do not appear to change with galactocentric radius, consistent with the M101 results.  \\citet{kau1999} modeled the expected CO intensity from a PDR for a range of incident FUV fluxes and cloud densities. The range of our results would be consistent with modeled CO intensities below $5 \\times 10^{-8}~ \\mbox{ergs cm}^{-2}~ \\mbox{s}^{-1}~ \\mbox{sr}^{-1}$, or $1 \\times 10^{-8}~ \\mbox{ergs cm}^{-2}~ \\mbox{s}^{-1}~ \\mbox{sr}^{-1}$ for the vast majority of sources (6.4 K km s$^{-1}$). Figure 1 in \\citet{all2004} shows that for the range of values in Figure \\ref{fig:selection}, the modeled CO intensities are independent of $n$. The low volume densities we find are consistent with a lack of CO emission in M81 as discussed by \\citet{kna2006}, and which those authors attribute to insufficient excitation. We note that the volume number densities of colliding \\Htwo\\ molecules in the GMCs of M81 is 1/2 of the values for $n$ calculated here. The mean value we have found for $n$ in the GMCs of M81 translates into a mean number density for $n_{H_2} \\approx 10 \\mbox{cm}^{-3}$, well below the values required for collisional excitation of CO molecules. In this case, the CO emission in M81 is in general \\textit{subthermal}. Their reported upper limit for $I_{CO}$ is 1.03 K km s$^{-1}$ for the regions they investigated, near our source no. 18. The \\Gnaught\\ and $n$ that we find there are consistent to that value of $I_{CO}$ (and somewhat higher) per Figure 1 in \\citet{all2004}, when beam dilution is taken into account. The low CO emission in M81 is also explored in \\citet{cas2007}, who again point to a lack of excitation of the molecular gas. They find no molecular gas in the nucleus of M81. The absence of FUV sources to excite the gas is consistent with that finding.  After accounting for observational and projection effects (see the previous Section), we note that in the nearby galaxies in general on which our method is applicable, it is most sensitive to a combination of low \\Gnaught\\ and low $n$.  Summarizing, our conclusions are:  \\begin{itemize} \\item We selected a number of discrete FUV sources in M81, which we consider to be potential PDRs on the surfaces of the parent GMCs. \\item The total hydrogen volume densities of GMCs close to clusters of young stars in M81 are in the range of 1 $< n <$ 200 cm$^{-3}$ with a geometric mean of 17 cm$^{-3}$. This is approximately ten times lower than GMCs in M101 studied with the same method. \\item The low GMC volume densities are consistent with a lack of CO emission in M81. \\item M81's GMCs have a filling factor of $\\approx$ 0.21 within $R_{25}$. \\item No candidate PDRs are found in M81 within $R < 0.3 R_{25}$. \\item We have provided a thorough analysis of the observational selection effects on our results and conclude that, while such effects are (necessarily) present in our results, our main conclusions as to the range and values of the total volume densities of GMCs in M81 are not affected. \\item The presence of PAH emission in the neighborhood of our candidate PDRs lends support to our view that the HI patches near FUV sources are indeed produced by photodissociation. PAH emission occurs near almost all UV sources. PAH and HI emission coincide in more than half of our sources. \\end{itemize}"
    },
    "0710/0710.1475_arXiv.txt": {
        "abstract": "Gamma ray bursts have been divided into two classes, long-soft gamma ray burst and short-hard gamma ray burst according to the bimodal distribution in duration time. Due to the harder spectrum and the lack of afterglows of short-hard bursts in optical and radio observations, different progenitors for short-hard bursts and long-soft bursts have been suggested. Based on the X-ray afterglow observation and the cumulative red-shift distribution of short-hard bursts, \\citet{Nak06} found that the progenitors of short-hard bursts are consistent with old populations, such as mergers of binary neutron stars. Recently, the existence of two subclasses in long-soft bursts has been suggested after considering multiple characteristics of gamma-ray bursts, including fluences and the duration time.  In this work, we extended the analysis of cumulative red-shift distribution to two possible subclasses in L-GRBs. We found that two possible subclass GRBs show different red-shift distributions, especially for red-shifts $z > 1$. Our results indicate that the accumulative red-shift distribution can be used as a tool to constrain the progenitor characteristics of possible subclasses in L-GRBs. ",
        "introduction": "Traditionally, gamma-ray bursts (denoted as GRBs) have been divided into two classes which are separated by the duration time of gamma-ray bursts $T_{90}=2$ sec \\citep{Kou93}. Since the short-duration gamma-ray bursts had relatively hard observed spectra than long-duration gamma-ray bursts, two classes are usually called as long-soft gamma ray burst (denoted as L-GRB) and short-hard gamma-ray burst (denoted as SHB). From the afterglow observations of L-GRBs, the association between the L-GRBs and supernovae/hypernovae has been suggested. These observations suggest that the progenitors of L-GRBs are massive giant stars which are collapsing. One of the best candidate is the Collapsar model \\citep{Woo93,Mac99} in which rapidly rotating black hole with Kerr parameter about 0.8 is formed after GRB. Recently, spinning black holes with high Kerr parameters, $a_\\star \\approx 0.8$, have been observed in soft X-ray black hole binaries \\citep{McC06,Sha06}. These observations support the idea that the rapidly rotating black holes in soft X-ray black hole binaries are the remnants of L-GRBs \\citep{Bro00,Lee02}, providing the natural sources for Collapsars.  Due to the differences in the observational characteristics of SHBs compared to those of L-GRBs, such as lack of afterglow observations in optical and radio band and harder spectrum, etc., different progenitors for SHBs have been suggested. Mergers of binary neutron stars \\citep{Eic89,Nar92} are among the best candidates. After recent X-ray afterglow observations by Swift and HETE-II, the associations between SHBs and the host galaxies were reported (see Table~\\ref{tab1} and references there in). Based on the X-ray afterglow observation of SHBs and the cumulative red-shift distribution of SHBs, \\citet{Nak06} found that the progenitors of SHBs are consistent with old populations. This finding supports mergers of binary neutron stars as progenitors of SHBs.  Recently, the existence of two subclasses in L-GRBs has been suggested after considering both fluences and the duration time \\citep{Cha07}. They divided L-GRBs into Clusters II and III which were separated by the fluences combined with the duration time. They classified L-GRBs with higher fluences as Cluster III and suggested that their progenitors are the massive collapsing stars, such as Collapsar \\citep{Woo93,Mac99}. As a possible origin of Cluster II GRBs, they suggested neutron-star, white-dwarf binaries. If there exist two subclasses in L-GRBs as \\citet{Cha07} suggested and only Cluster III GRBs are associated with massive collapsing giants which are the progenitors of supernovae or hypernovae, the observed red-shift distribution of two subclasses might show different behaviour.   In this work, we extended the work of \\citet{Nak06} to two possible subclasses in L-GRBs. We found that Clustters II and III show different red-shift distribution at high red-shift $z>1$. Our results indicate that the accumulative red-shift distribution can be used as a tool to constrain the lifetimes of possible subclasses in GRBs. In Sec.~\\ref{sec-rate}, the numerical methods which we used in our analysis are summarized \\citep{Nak06}. The estimated lifetime of SHBs are summarized in Sec.~\\ref{sec-SHB}. We confirmed that the progenitors of SHBs are consistent with old population with lifetime $\\tau_\\ast=6.5$ Gyr, indicating that SHBs are from the mergers of binary neutron stars. In Sec.~\\ref{sec-LGRB}, the cumulative red-shift distributions of two possible subclasses in L-GRB have been summarized. Our results show that there are clear differences in the red-shift distributions of two possible subclasses in L-GRBs for the red-shift  $z>1$. This suggests that the cumulative red-shift distribution can be used as a tool to distinguish the subclasses in L-GRBs. Our final conclusion follows in Sec.~\\ref{sec-con}. ",
        "conclusions": "\\label{sec-con}  In this work, by extending the work of Nakar et al. \\citep{Nak06,Nak07}, we investigated the constraints on the progenitor lifetime of SHBs. We confirmed that the SHBs are consistent with old population with mean time $\\tau_\\ast=6.5$ Gyr. We also confirmed that this conclusion is independent of the details of the cosmology models. This results support the conclusion that the progenitors of SHBs are consistent with old population, such as merging compact star binaries (neutron star--neutron star binaries or neutron star--black hole binaries) \\citep{Nak06,Nak07,Bet07,Lee07}.  We also investigated two subclasses Cluster II and III L-GRBs \\citep{Cha07} using the same method of \\citet{Nak06}. We found that the power-law distribution is more consistent with Cluster III GRBs than with Cluster II GRBs. We also found that the cumulative red-shift distribution of two clusters show clear difference in the red-shift region $z>1$. This result supports the existence of two subclasses in L-GRBs. However, in order to draw firm conclusions, many effects which are not included in our analysis, e.g., the connection between GRBs and the metallicity of the host galaxy, different characteristics of cosmological GRBs and subluminous GRBs, observability premium for different type of GRB progenitors, etc., have to be considered."
    },
    "0710/0710.1828_arXiv.txt": {
        "abstract": "We present high signal-to-noise spectrophotometric observations of seven luminous \\HII\\ galaxies. The observations have been made with the use of a double-arm spectrograph which provides spectra with a wide wavelength coverage, from 3400 to 10400\\,\\AA\\ free of second order effects, of exactly the same region of a given galaxy. These observations are analysed applying a methodology designed to obtain accurate elemental abundances of oxygen, sulphur, nitrogen, neon, argon and iron in the ionized gas. Four electron temperatures and one electron density are derived from the observed forbidden line ratios using the five-level atom approximation. For our best objects errors of  1\\% in t$_e$([O{\\sc iii}]), 3\\% in t$_e$([O{\\sc ii}]) and 5\\% in t$_e$([S{\\sc iii}]) are achieved with a resulting accuracy of  7\\% in total oxygen abundances, O/H.  The ionisation structure of the nebulae can be mapped by the theoretical oxygen and sulphur ionic ratios, on the one side,  and the corresponding observed emission line ratios, on the other --  the $\\eta$ and $\\eta$' plots --. The combination of both is shown to provide a means to test photo-ionisation model sequences currently applied to derive elemental abundances in \\HII\\ galaxies. ",
        "introduction": "When studying evolution two types of ages should be distinguished: the chronological and the evolutionary ages. In the case of galaxies, estimates of the chronological age can be obtained analyzing, for example, the age distribution of their stellar population while the evolutionary age can be estimated from, for example, the metal content of their interstellar medium.  \\HII\\ galaxies, the subclass of Blue Compact Dwarf galaxies (BCDs) which show spectra with strong emission lines similar to those of giant extragalactic \\HII\\ regions \\citep[GEHRs;][]{1970ApJ...162L.155S,1980ApJ...240...41F}, have the lowest metal content of any starforming galaxy suggesting that they are among the youngest or less evolved galaxies known \\citep{2007ApJ...654..226R, 1972ApJ...173...25S}. After the findings that a considerable number of the objects observed at intermediate and high redshifts seem to have properties similar to the \\HII\\ galaxies we know in the Local  Universe, it has been suggested that these objects might have been very common in the past and some of them may have evolved to other kind of objects \\citep{1995ApJ...440L..49K}. In order to detect these evolutionary effects we need to compare the properties of \\HII\\ galaxies both in the Local Universe and at higher redshifts. We therefore need to know the true distribution functions of their properties among which the chemical abundances are of the greatest relevance.  Spectrophotometry of bright \\HII\\ galaxies in the Local Universe allows the determination of abundances from methods that rely on the measurement of emission line intensities and atomic physics. This is referred to as the \"direct\" method. In the case of more distant or intrinsically fainter galaxies, the low signal-to-noise obtained with current telescopes precludes the application of this method and empirical ones based on the strongest emission lines are required. The fundamental basis of these empirical methods is reasonably well understood \\citep[see e.g.][]{2005MNRAS.361.1063P}. The accuracy of the results however depends on the goodness of their calibration which in turn depends on a well sampled set of precisely derived abundances by the \"direct\" method so that interpolation procedures are reliable. Enlarging the calibration range is also important since, at any rate, empirically obtained relations should never be used outside their calibration validity range.  The precise derivation of elemental abundances however is not a straightforward matter. Firstly, accurate measurements of the emission lines are needed. Secondly, a certain knowledge of the ionisation structure of the region is required in order to derive ionic abundances of the different elements and in some cases photoionisation models are needed to correct for unseen ionisation states. An accurate diagnostic requires the measurement of faint auroral lines covering a wide spectral range and their accurate (better than 5\\%) ratios to Balmer recombination lines. These faint lines are usually about 1\\% of the  H$\\beta$ intensity. The spectral range must include from the UV [O{\\sc ii}]\\,$\\lambda$\\,3727\\,\\AA\\ doublet, to the near IR [S{\\sc iii}] $\\lambda\\lambda$\\,9069,9532\\,\\AA\\ lines. This allows the derivation of the different line temperatures: T$_e$([O{\\sc ii}]), T$_e$([S{\\sc ii}]), T$_e$([O{\\sc iii}]), T$_e$([S{\\sc iii}]), T$_e$([N{\\sc ii}]), needed in order to study the temperature and ionisation structure of each \\HII\\ galaxy considered as a multizone ionised region.      Unfortunately most of the available \\HII\\ galaxy spectra have only a restricted wavelength  range (usually from about 3600 to 7000\\,\\AA), consequence of observations with single arm spectrographs, 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; Stoughton et al.\\ 2002) \\nocite{2002AJ....123..485S} spectra do not cover simultaneously the 3727 [O{\\sc ii}] and  the 9069 [S{\\sc iii}] lines, they only represent an average inside a 3\\,arcsec fibre and reach the required S/N only for the brightest objects.  It is important to realise that the combination of accurate spectrophotometry and wide spectral coverage cannot be achieved  using single arm spectrographs where, in order to  reach the necessary spectral resolution, the wavelength range must be split into several independent observations. In those cases, the quality of the spectrophotometry is at best doubtful mainly because the different spectral ranges are not observed simultaneously. This problem applies to both objects and calibrators. Furthermore one can never be sure of observing exactly the same region of the nebula in each spectral range. To avoid all these problems the use of double arm spectrographs is required.         In this work we present simultaneous blue and red observations obtained with the double arm TWIN spectrograph at the 3.5m telescope of the Spanish-German Observatory of Calar Alto. These data are of a sufficient quality as to allow the detection and measurement of several temperature sensitive lines and add to the still scarce base of precisely derived abundances. In the next section we describe some details regarding the selection of the sample as well as the observations and data reduction. The results are presented in section 3. Sections 4 and 5 are devoted to the analysis of these results which are compared with previous data in section 6. Section 7 is devoted to the discussion of our results and finally, our conclusions are summarized in section 8.     \\begin{table*} \\centering \\caption[]{Journal of observations.} \\label{jour} \\begin{tabular} {@{}c c c c c c} \\hline \\multicolumn{1}{c}{Object  ID}  &  spSpec SDSS   & hereafter ID &  Date &  Exposure (s)  & Seeing (\\arcsec)\\\\ \\hline \\uno    & spSpec-52790-1351-474  & \\unoc    & 2006 June 25  &  5 $\\times$ 1800  & 0.9-1.2 \\\\ \\dos    & spSpec-52721-1050-274  & \\dosc    & 2006 June 23  &  4 $\\times$ 1800  & 0.8-1.1 \\\\ \\tres   & spSpec-52765-1293-580  & \\tresc   & 2006 June 22  &  4 $\\times$ 1800  & 0.8-1.2 \\\\ \\cuatro & spSpec-52072-0617-464  & \\cuatroc & 2006 June 24  &  6 $\\times$ 1800  & 1.0-1.4 \\\\ \\cinco  & spSpec-52377-0624-361  & \\cincoc  & 2006 June 23  &  5 $\\times$ 1800  & 0.8-1.1 \\\\ \\seis   & spSpec-52791-1176-591  & \\seisc   & 2006 June 25  &  5 $\\times$ 1800  & 0.9-1.2 \\\\ \\siete  & spSpec-51818-0358-472  & \\sietec  & 2006 June 22  &  5 $\\times$ 1800  & 0.8-1.1 \\\\ \\hline \\end{tabular} \\end{table*}   \\begin{table*} \\centering \\caption[]{Right ascension, declination, redshift and SDSS photometric magnitudes obtained using the SDSS explore tools$^a$.} \\label{obj} \\begin{tabular} {l c c c c r r c c} \\hline \\multicolumn{1}{c}{Object  ID}   & \\multicolumn{1}{c}{RA} & \\multicolumn{1}{c}{Dec} &  redshift  & \\multicolumn{1}{c}{u} & \\multicolumn{1}{c}{g} & \\multicolumn{1}{c}{r} & \\multicolumn{1}{c}{i}   &    \\multicolumn{1}{c}{z} \\\\ \\hline \\unoc    &  14$^h$\\,55$^m$\\,06\\fs06 & 38$^{\\circ}$\\,08\\arcmin\\,16\\farcs67 &  0.028 & 18.25 & 17.57 & 17.98 & 18.23 & 18.18 \\\\ \\dosc    &  15$^h$\\,09$^m$\\,09\\fs03 & 45$^{\\circ}$\\,43\\arcmin\\,08\\farcs88 &  0.048 & 18.57 & 17.72 & 18.19 & 17.87 & 17.94 \\\\ \\tresc   &  15$^h$\\,28$^m$\\,17\\fs18 & 39$^{\\circ}$\\,56\\arcmin\\,50\\farcs43 &  0.064 & 18.54 & 17.88 & 18.17 & 17.52 & 17.99 \\\\ \\cuatroc &  15$^h$\\,40$^m$\\,54\\fs31 & 56$^{\\circ}$\\,51\\arcmin\\,38\\farcs98 &  0.011 & 19.11 & 18.91 & 18.97 & 19.53 & 19.46 \\\\ \\cincoc  &  16$^h$\\,16$^m$\\,23\\fs53 & 47$^{\\circ}$\\,02\\arcmin\\,02\\farcs36 &  0.002 & 16.84 & 16.45 & 16.77 & 17.35 & 17.43 \\\\ \\seisc   &  16$^h$\\,57$^m$\\,12\\fs75 & 32$^{\\circ}$\\,11\\arcmin\\,41\\farcs42 &  0.038 & 17.63 & 17.03 & 17.27 & 17.15 & 17.15 \\\\ \\sietec  &  17$^h$\\,29$^m$\\,06\\fs56 & 56$^{\\circ}$\\,53\\arcmin\\,19\\farcs40 &  0.016 & 18.05 & 17.26 & 17.21 & 17.38 & 17.24 \\\\ \\hline \\multicolumn{9}{l}{$^a$http://cas.sdss.org/astro/en/tools/explore/obj.asp} \\end{tabular} \\end{table*} ",
        "conclusions": "We have performed a detailed analysis of newly obtained spectra of seven \\HII\\ galaxies selected from the Sloan Digital Sky Survey Data Release 3. The spectra cover from 3400 to 10400 \\AA\\ in wavelength at a FWHM resolution of about 2000 in the blue and 1500 in the red spectral regions.  The high signal-to-noise ratio of the obtained spectra allows the measurement of four line electron temperatures: T$_e$([O{\\sc iii}]), T$_e$([S{\\sc iii}]), T$_e$([O{\\sc ii}]) and T$_e$([S{\\sc ii}]), for all the objects of the sample with the addition of T$_e$([N{\\sc ii}]) for one of the objects. These measurements and a careful and realistic treatment of the observational errors yield total oxygen abundances with accuracies between 7 and 12\\%. The fractional error is as low as 1\\% 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$_e$([O{\\sc iii}]), but increases to up to 30\\% for the O$^{+}$/H$^{+}$ ratio. The accuracies are lower in the case of the abundances of sulphur (of the order of 25\\% for S$^+$ and 15\\% for S$^{2+}$) due to the presence of larger observational errors both in the measured line fluxes and the derived electron temperatures. The error for the total abundance of sulphur  is also larger than in the case of oxygen (between 15\\% and 20\\%) due to the uncertainties in the ionisation correction factors.  This is in contrast with the unrealistically small errors quoted for line temperatures other than T$_e$([O{\\sc iii}]) in the literature, in part due to the commonly assumed methodology of deriving them from the measured T$_e$([O{\\sc iii}]) through a theoretical relation. These relations are found from photoionization model sequences and no uncertainty is attached to them although large scatter is found when observed values are plotted; usually the line temperatures obtained in this way  carry only the observational error found for the T$_e$([O{\\sc iii}]) measurement and does not include the observed scatter, thus heavily understimating the errors in the derived temperature.  In fact, no clear relation is found between T$_e$([O{\\sc iii}]) and T$_e$([O{\\sc ii}]) for the existing sample of objects confirming our previous results. A comparison between measured and model derived T$_e$([O{\\sc ii}]) shows than, in general, model predictions overestimate this temperature and hence underestimate the O$^+$/H$^+$ ratio. This, though not very important for high excitation objects, could be of some concern for lower excitation ones for which total O/H abundances could be underestimated by up to 0.2 dex. It is worth noting that the objects observed with double-arm spectrographs, therefore implying simultaneous and spatially coincident observations over the whole spectral range,  show less scatter in the T$_e$([O{\\sc iii}])\\,-\\,T$_e$([O{\\sc ii}] plane clustering around the N$_e$\\,=\\,100 cm$^{-3}$ photo-ionisation model sequence. On the other hand, this small scatter could partially be due to the small range of temperatures shown by these objects due to possible selection effects. This small temperature range does not allow either to investigate the metallicity effects found in the ralations between the various line temperatures in recent photo-ionisation models by Izotov et al.\\ (2006).  Also, the observed objects, as well as those in Paper I, though showing  Ne/O and Ar/O relative abundances typical of those found for a large \\HII\\ galaxy sample (P\\'erez-Montero et al.\\ 2007), show higher than typical N/O abundance ratios that would be even higher if the [O{\\sc ii}] temperatures would be found from photo-ionisation models. We therefore conclude that approach of deriving the O$^+$ temperature from the O$^{2+}$ one should be discouraged if an accurate abundance derivation is sought.  These issues could be addressed by re-observing the objects in Table \\ref{temp} , which cover an ample range in temperatures and metal content, with double arm spectrographs. This sample should be further extended  to obtain a self consistent sample of about 50 objects with high S/N and excellent spectrophotometry covering simultaneously from 3600 to 9900\\,\\AA\\ This simple and easily feasible project would provide important scientific return in the form of critical tests of photoionisation models.          The O$^{+} $/O$^{2+} $ and S$^{+} $/S$ ^{2+} $ ratios for all the observed galaxies, except one,  cluster around a value of the ``softness parameter\" $\\eta$ of 1 implying high values of the stellar ionising temperature. For the discrepant object, showing a much lower value of $\\eta$,  the intensity of the [O{\\sc ii}]\\,$\\lambda\\lambda$\\,7319,25\\,\\AA\\ lines are affected by atmospheric absorption lines. When the observational counterpart of the ionic ratios is used, this object shows a ionisation structure similar to the rest of the observed ones. This simple exercise shows the potential of checking for consistency in both the $\\eta$ and $\\eta$' plots in order to test if a given assumed ionisation structure is adequate. In fact, these consistency checks show that the stellar ionising temperatures found for the observed \\HII\\ galaxies using the ionisation structure predicted by state of the art ionisation models result too low when compared to those implied by the corresponding observed emission line ratios.  Therefore, metallicity calibrations based on abundances derived according to this conventional method are probably bound to provide metallicities which are systematically too high and should be revised."
    },
    "0710/0710.4956_arXiv.txt": {
        "abstract": "We present Submillimeter Array observations of the z=3.91 gravitationally lensed broad absorption line quasar \\apm\\, which spatially resolve the 1.0~mm (200~$\\mu$m rest-frame) dust continuum emission. At 0\\fasec4 resolution, the emission is separated into two components, a stronger, extended one to the northeast ($46\\pm5$~mJy) and a weaker, compact one to the southwest ($15\\pm2$~mJy). We have carried out simulations of the gravitational lensing effect responsible for the two submm components in order to constrain the intrinsic size of the submm continuum emission. Using an elliptical lens potential, the best fit lensing model yields an intrinsic (projected) diameter of $\\sim$80~pc, which is not as compact as the optical/near-infrared (NIR) emission and agrees with previous size estimates of the gas and dust emission in \\apm. Based on our estimate, we favor a scenario in which the 200~$\\mu$m (rest-frame) emission originates from a warm dust component (T$_{\\rm d}$=150-220~K) that is mainly heated by the AGN rather than by a starburst (SB). The flux is boosted by a factor of $\\sim$90 in our model, consistent with recent estimates for \\apm. ",
        "introduction": "\\apm\\, (=APM08) is a strongly lensed broad absorption line (BAL) quasar \\citep{irwin98,lewis98} with a very powerful active galactic nucleus \\citep[AGN;][]{soifer87}.  This combination makes it an extremely bright object despite its redshift of $z$=3.91 \\citep{downes99}.  Its bolometric luminosity is thought to be $\\sim$5$\\times$10$^{13}$~L$_{\\odot}$, amplified by up to a factor of 100 by a foreground galaxy which has yet to be identified.  Due to the large amplification, broad absorption lines have been detected even in the X-ray \\citep{chartas02}. Ground-based, {\\it Chandra} and HST observations have suggested that \\sapm\\, consists of at least two components separated by $\\sim$$0\\farcs4$. \\citet[=I99]{ibata99} and \\citet[=E00]{egami00} later detected a third image (their image C), and argued that it is likely a third lensed image of the background QSO instead of the lensing galaxy.  \\cite{lewis02b} obtained optical spectra of the three images using HST/STIS, and showed that the three spectra are quite similar.  This indicates that the image C is indeed a third lensed image of \\sapm. Due to the magnification by the lens, several molecular lines have been detected in \\sapm\\, \\citep[e.g.,][] {weiss07,guelin07,riecher06,wagg06,garcia06,wagg05,lewis02a,downes99}, suggesting a molecular gas mass of M$_{\\rm gas}$(H$_2$)=8$\\times$10$^{10}$~$m^{-1}$~M$_\\odot$ \\citep[][$m$$\\equiv$magnification factor]{riecher06}.  The dust mass is estimated to be M$_{\\rm dust}$=5$\\times$10$^8$~$m^{-1}$~M$_\\odot$ \\citep[][W07]{downes99,weiss07}. However, despite the increasing amount of mm-data, no tight constraints have been set on the size of the dust and/or gas emission which may help to characterize the heating source of the dust and gas, i.e., whether it is in the form of an AGN and/or starburst (SB). Only recently have W07 presented some indirect evidence through brightness temperature arguments that the molecular line and dust emission originate from a compact region with a radius of 100-200~pc. In this {\\it Letter}, we present observations from the Submillimeter Array (SMA) with sufficient resolution (0\\fasec4) to resolve the 1.0~mm dust continuum emission in \\sapm, and lensing models that constrain its intrinsic source size. ",
        "conclusions": "\\label{con} We have detected the 200~$\\mu$m rest-frame continuum emission in \\sapm\\, using the SMA with an angular resolution of $\\sim$0\\fasec4. The two (main) lensed images are clearly separated with a combined flux of $\\sim$60$\\pm$12~mJy.  Simulations of the gravitational lensing effect in this system yield a diameter of the intrinsic (submm) continuum emitting region of $\\sim$80~pc and magnification factor of 90. Our data and simulations seem to be only consistent with a lensing scenario including a third lensed image of \\sapm\\, close to NE if the position of the compact optical/NIR emission and the position of the extended submm emission are offset from each other before lensing by $\\sim$0\\fasec003 ($\\equiv$21 pc). Further (sub)mm observations may be beneficial to determine whether such a positional offset is indeed necessary or more complicated lens potentials have to be considered. Given our size estimate, we favor a scenario in which the 200~$\\mu$m emission originates from a warm dust component that is supposed to be mainly heated by the AGN rather than by a SB."
    },
    "0710/0710.1097_arXiv.txt": {
        "abstract": "We report extensive photometry of the dwarf nova V419 Lyr throughout its 2006 July superoutburst till quiescence. The superoutburst with amplitude of $\\sim 3.5$ magnitude lasted at least 15 days and was characterized by the presence of clear superhumps with a mean period of $P_{\\rm sh}=0.089985(58)$ days ($129.58\\pm0.08$ min). According to the Stolz-Schoembs relation, this indicates that the orbital period of the binary should be around 0.086 days i.e. within the period gap.  During the superoutburst the superhump period was decreasing with the rate of $\\dot{P}/P_{sh}=-24.8(2.2)\\times10^{-5}$, which is one of the highest values ever observed in SU UMa systems. At the end of the plateau phase, the superhump period stabilized at a value of 0.08983(8) days.  The superhump amplitude decreased from 0.3 mag at the beginning of the superoutburst to 0.1 mag at its end. In the case of V419 Lyr we have not observed clear secondary humps, which seems to be typical for long period systems.  \\noindent {\\bf Key words:} Stars: individual: V419 Lyr -- binaries: close -- novae, cataclysmic variables ",
        "introduction": "Dwarf novae -- a subclass of Cataclysmic Variable stars -- are quite well studied interacting binary systems composed of late-type red dwarf secondary and white dwarf primary stars (Warner 1995, Hellier 2001). Matter transferred from the red dwarf forms an accretion disc around the white dwarf. Although in the last decade significant progress has been made in explaining the behaviour of dwarf novae light curves, some physical processes ongoing in these systems are still not fully understood (see for example Smak 2000, Schreiber and Lasota 2007). In particular, the thermal-tidal instability model of Osaki (1996, 2005) describing the phenomenon of superoutbursts and superhumps may be tested by examination of SU UMa-type dwarf novae light curves. Additionally, objects near and inside the so called period gap are very important from an evolutionary point of view. Those systems give us an unprecedented opportunity to study the evolution of dwarf novae.  V419 Lyr is a poorly studied cataclysmic variable discovered by Kurochkin (1990) and originally classified as a Z Cam-type dwarf nova. Later, Nogami et al. (1998) caught this object in outburst and found superhumps in its light curve. Detection of superhumps together with characteristic properties of the outburst allowed them to classify V419 Lyr as a SU UMa-type dwarf nova, but short coverage of the eruption did not allow accurate determination of the superhump period. Nevertheless, there was a strong suggestion that V419 Lyr has one of the longest orbital periods known among SU UMa variables.  This object has been monitored at various photometric bands by the Variable Star Network (VSNET) (see for example Kato et al. 2004a). The observations from that program enabled a tentative determination of the supercycle period to be about $\\sim 340$ days (Katysheva and Pavlenko 2003). Moreover, Morales-Rueda and Marsh (2002) obtained a spectrum of V419 Lyr during outburst showing a relatively broad absorption feature around 430-440 nm.  In this work we present an analysis of photometric data collected during the 2006 July superoutburst of V419 Lyr. The data are much richer than previous studies and provide us with an opportunity to determine parameters describing this system more precisely. ",
        "conclusions": "The orbital period of V419 Lyr is unknown. However it is possible to estimate its value using the relation in Stolz and Schoembs (1984) connecting the period excess $\\epsilon$ defined as $P_{\\rm sh}/P_{\\rm orb}-1$ with the orbital period of the binary. This empirical relation is as follows:  \\begin{equation} \\epsilon = 0.858(11) \\cdot P_{\\rm orb} - 0.0282(2) \\end{equation}  Using the definition of $\\epsilon$ and knowing $P_{\\rm sh}$ for V419 Lyr, we were able to estimate the orbital period as $P_{\\rm orb}\\approx 0.086$ days. This is slightly longer than two hours which indicates that V419 Lyr is a dwarf nova in the period gap.  Many characteristics of V419 Lyr are typical of SU UMa stars. It goes into superoutburst every year or so, the eruption lasts about two weeks and has an amplitude of $\\sim 3.5$ mag. Superhumps appear shortly after the beginning of the superoutburst and have a maximum amplitude of 0.3 mag, which decreases to 0.1 mag at the end of the outburst. In addition to its long orbital period, V419 Lyr is unusual in two other properties. Its superhump period derivative has one of the largest negative values known and it shows only a weak trace of secondary humps in the final stages of the superoutburst.  \\bigskip \\noindent {\\bf Acknowledgments.} ~We acknowledge generous allocation of  the Warsaw Observatory 0.6-m telescope time. This work used the online service of the VSNET and AAVSO. We would like to thank Prof. J\\'ozef Smak for fruitful discussions."
    },
    "0710/0710.1574_arXiv.txt": {
        "abstract": "{We present the results of a {\\it Chandra} soft X-ray observation of the spectacular ionization cone in the nearby Seyfert~2 galaxy NGC~5252. As almost invariably observed in obscured AGN, the soft X-ray emission exhibits a remarkable morphological concidence with the cone ionized gas as traced by HST O[{\\sc iii}] images. Energy-resolved images and high-resolution spectroscopy suggest that the X-ray emitting gas is photoionized by the AGN, at least on scales as large as the innermost gas and stellar ring ($\\le$3~kpc). Assuming that the whole cone is photoionized by the AGN, we reconstruct the history of the active nucles in the last $\\sim$10$^5$~years.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.0602_arXiv.txt": {
        "abstract": "We describe the discovery of HAT-P-4b, a low-density extrasolar planet transiting BD+36~2593, a $V=11.2$\\,mag slightly evolved metal-rich late F star. The planet's orbital period is $3.056536\\pm0.000057$\\,d with a mid-transit epoch of $2,454,245.8154\\pm0.0003$ (HJD). Based on high-precision photometric and spectroscopic data, and by using transit light curve modeling, spectrum analysis and evolutionary models, we derive the following planet parameters: \\mpl$=0.68\\pm0.04$\\,\\mjup, \\rpl$=1.27\\pm0.05$\\,\\rjup, \\rhopl$=0.41\\pm0.06$\\,\\gcmc\\ and $a=0.0446\\pm0.0012$\\,AU\\@. Because of its relatively large radius, together with its assumed high metallicity of that of its parent star, this planet adds to the theoretical challenges to explain inflated extrasolar planets. ",
        "introduction": "\\label{sec:intro} In the course of our ongoing wide field planetary transit search program HATNet \\citep{bakos04}, we have discovered a large radius and low density planet orbiting an $11$th magnitude star BD+36~2593. This planet is the fifth member of a group of low-density transiting exoplanets. The combination of its low mass and the relatively high metallicity and age of the parent star makes theoretical interpretation of its large radius difficult. In this Letter we describe the observational properties of the system and derive the physical parameters both for the host star and for the planet. We also briefly comment on the theoretical status of inflated extrasolar planets. ",
        "conclusions": "\\label{sec:disc}  We presented the discovery data and derived the physical parameters of HAT-P-4b, an inflated planet orbiting BD+36~2593. Among the 20 transiting planets discovered so far, there are five with \\rhopl~$\\lesssim0.4$\\,\\gcmc. All others have at least $50$\\% higher densities. For ease of comparison, \\tabr{puffy} lists the relevant properties of the five inflated planets. It is remarkable how similar these planets are (except for TrES-4 that has distinctively low density). With its Safronov number of $0.036$, HAT-P-4b belongs to the Class II planets according to the recent classification of \\cite{hansen07} and (together with TrES-4) further strengthens the mysterious dichotomy of the known transiting planets in this parameter. The parent star of HAT-P-4b is among the largest radii, largest mass, lowest gravity and highest metallicity transiting planet host stars.  \\notetoeditor{This is the intended place of \\tabr{puffy}}  \\begin{deluxetable}{lcccccc} \\tabletypesize{\\scriptsize} \\tablecaption{ Comparison of the properties of inflated planets.\\tablenotemark{a} \\label{tab:puffy}} \\tablewidth{0pt} \\tablehead{ \\colhead{Name}          & \\colhead{P}             & \\colhead{a}             & \\colhead{M}             & \\colhead{R}             & \\colhead{$\\rho$}        & \\colhead{$\\log g$}     \\\\ \\colhead{}              & \\colhead{(d)}           & \\colhead{(AU)}          & \\colhead{(M$_J$)}       & \\colhead{(R$_J$)}       & \\colhead{(\\sc{cgs})}    & \\colhead{(\\sc{cgs})} } \\startdata WASP-1b    & 2.52 & 0.038 & 0.87 & 1.40 & 0.39 & 3.04 \\\\ HAT-P-4b   & 3.06 & 0.045 & 0.68 & 1.27 & 0.41 & 3.02 \\\\ HD~209458b & 3.53 & 0.045 & 0.64 & 1.32 & 0.35 & 2.96 \\\\ TrES-4     & 3.55 & 0.049 & 0.84 & 1.67 & 0.22 & 2.87 \\\\ HAT-P-1b   & 4.47 & 0.055 & 0.53 & 1.20 & 0.38 & 2.96 \\\\ \\enddata \\tablenotetext{a} {Data from \\cite{shporer07}, this paper, \\cite{burro07}, \\cite{mandu07} and \\cite{winn07}. From top to bottom, metallicities for the parent stars are: $0.23$ \\citep{stempels07}, $0.24$, $0.02$ $0.0$ (adopted) and $0.13$.} \\end{deluxetable}  Current models of irradiated giant planets are able to match the observed radii of most of the planets without invoking any additional heating mechanism. Higher metallicity cases, such as the present one, however, may pose problems (assuming that the planet and star have similar metallicities). More metals imply two opposite effects on the radius: (i) inflating it due to higher opacities in the envelope; (ii) shrinking it due to the higher molecular weight of the interior and the possible development of a large high density core. These effects have been discussed recently by \\cite{burro07}. Since WASP-1b is similar in several aspects (i.e., irradiance, metallicity) to HAT-P-4b, we consider the coreless models of WASP-1b as shown in Fig.~7 of \\cite{burro07}. It seems that HAT-P-4b can be fitted by near solar metallicity coreless models, assuming that its age is not too much greater that $4$\\,Gyr. We also refer to the layered convective mechanism of \\cite{chabrier07} that gives an alternative explanation for planets with inflated radii.  We conclude that more definite statements on the relation of the observations and planet structure theories can be made only by reaching higher accuracy in the observed star/planet parameters. Nevertheless, HAT-P-4b (together with WASP-1b) does not seem to support the existence of a simple relation between host star metallicity and planet's core mass \\citep[see][]{guillot06,burro07}."
    },
    "0710/0710.2741_arXiv.txt": {
        "abstract": "We study the cosmological evolution of an induced gravity model with a self-interacting scalar field $\\sigma$ and in the presence of matter and radiation. Such model leads to Einstein Gravity plus a cosmological constant as a stable attractor among homogeneous cosmologies and is therefore a viable dark-energy (DE) model for a wide range of scalar field initial conditions and values for its positive $\\gamma$ coupling to the Ricci curvature $\\gamma \\sigma^{2}R$. ",
        "introduction": "Several years ago a model for a varying gravitational coupling was introduced by Brans and Dicke \\cite{BD}. The model consisted of a massless scalar field whose inverse was associated with the gravitational coupling. Such a field evolved dynamically in the presence of matter and led to cosmological predictions differing from Einstein Gravity (EG) in that one generally obtained a power-law dependence on time for the gravitational coupling. In order to reduce such a strong time dependence in a cosmological setting, while retaining the Brans-Dicke results in the weak field limit, several years ago a simple model for induced gravity \\cite{CV,TV} involving a scalar field $\\sigma$ and a quartic $\\lambda\\sigma^{4}/4$ potential was introduced. This model was globally scale invariant (that is did not include any dimensional parameter) and the spontaneous breaking of scale invariance in such a context led to both the gravitational constant and inflation, through a non-zero cosmological constant \\cite{Zee}. The cosmological consequences of introducing matter as a perturbation were studied leading to a time dependence for the scalar field and consistent results. Since the present observational status is compatible with an accelerating universe dominated by dark energy (DE) we feel that the model should be studied in more detail. \\\\ In an EG framework quadratic or quartic potentials for canonical scalar fields are consistent with DE only with an extremely fine tuning in the initial conditions leading to slow-roll evolution until the present time. Indeed, a massive or a self-interacting scalar field in EG, respectively behave as dust or radiation during the coherent oscillation regime. In contrast with this, induced gravity with a self-interacting $\\lambda\\sigma^4/4$ potential has as attractor EG plus a cosmological constant on breaking scale invariance. The simple model we consider illustrates how non-trivial and non perturbative dynamics can be obtained in the context of induced gravity DE, and more generally within scalar-tensor DE. ",
        "conclusions": "We have shown how a simple model of self-interacting induced gravity \\cite{CV,TV} is a viable DE model, for tiny values of the self-coupling ($\\lambda \\sim {\\cal O}(10^{-128})$). The model has a stable attractor towards EG plus $\\Lambda$ and can be very similar to the $\\Lambda$CDM model for the homogeneous mode, on taking into account the Solar system constraints quoted in \\cite{Bertotti:2003rm,Eubanks,gdotref}. At the cosmological level, in the presence of radiation and dust, it is interesting that for such a simple potential (quartic for induced gravity or interacting for the equivalent Brans-Dicke model) the model has an attractor towards EG plus $\\Lambda$, very differently from the case of a massless scalar - i.e. $\\lambda = 0$ -, for which there is no mechanism of attraction towards EG. The attractor mechanism towards GR and the onset to acceleration are both inevitably triggered by the same mechanism, i.e. scale symmetry breaking in this induced gravity model. \\begin{figure}[t!!] \\centering \\epsfig{file=ws_phi4.eps, width=8.5 cm} \\caption{Evolution of $w_{\\rm DE}$ as defined in (\\ref{rhopfake}) for different choices of $\\gamma$.\\label{ws}}. \\end{figure} In contrast with EG, the choice of a runaway potential \\cite{PR} for quintessence is not mandatory (see however \\cite{others} for a study of these potential in the context of scalar-tensor theories). If the final attractor is an accelerated universe, the only constraints on parameters and initial conditions come from observations. We have shown that late time cosmology - after recombination, for instance - is mostly insensitive with respect to the initial time derivative of $\\sigma$. For this reason, the full set of parameters and initial conditions of the model are fully specified on taking the observed values for $G, H_0, \\Omega_\\Lambda$.\\\\ We have discussed in detail such model in the context of Einstein gravity, i.e. keeping the Newton constant (approximately) fixed at its actual value. The model predicts an equation of state of the equivalent Einstein model with $w_{DE}$ slightly less than $-1$ at present and homogeneous cosmology by itself seems able to constrain $\\gamma \\lesssim {\\cal O}(10^{-3})$, although a full statistical analysis is clearly beyond the scope of this work.\\\\ It is interesting to also study also what other potentials, besides the simple potential $\\lambda \\sigma^4/4$ employed in this paper, will also be compatible with the observed evolution of the universe.\\\\ \\\\ {\\bf Acknowledgment.} We wish to thank the Referee for helpful and constructive criticism."
    },
    "0710/0710.0058_arXiv.txt": {
        "abstract": "The Wilkinson Microwave Anisotropy Probe (WMAP) three year results are used to constraint non-minimal inflation models. Two different non-minimally coupled scalar field potentials are considered to calculate corresponding slow-roll parameters of non-minimal inflation. The results of numerical analysis of parameter space are compared with WMAP3 data to find appropriate new constraints on the values of the non-minimal coupling. A detailed comparison of our results with previous studies reveals the present status of the non-minimal inflation model after WMAP3.\\\\ {\\bf PACS}:\\, 04.50.+h,\\, 98.80.-k\\\\ {\\bf Key Words}: Scalar-Tensor Gravity,\\, Inflation,\\, WMAP3 Data ",
        "introduction": "Non-minimal coupling (NMC) of scalar field with gravity is necessary in many situations of physical and cosmological interest. There are several compelling reasons to include an explicit non-minimal coupling in the action ( see for example [1,2,3,4] and references therein ). NMC arises at the quantum level when quantum corrections to the scalar field theory are considered. It is necessary also for the renormalizability of the scalar field theory in curved space. In most theories used to describe inflationary scenarios, it turns out that a non-vanishing value of the coupling constant is unavoidable [2]. In general relativity, and in all other metric theories of gravity ( in which the scalar field is not part of the gravitational sector ), the coupling constant necessarily assumes a non-vanishing value[5]. The study of the asymptotically free theories in an external gravitational field with a Gauss-Bonnet term shows a scale dependent coupling parameter. For instance, asymptotically free grand unified theories have a non-minimal coupling depending on a renormalization group parameter that converges to the value of $\\frac{1}{6}$ or to any other initial conditions depending on the gauge group and on the matter content of the theory[6]. In view of the above results and several other motivations( see for example [7,8] and references therein), it is then natural to incorporate an explicit NMC between scalar field and Ricci scalar in the inflationary paradigm. Generally, with non-minimally coupled scalar field it is harder to achieve accelerated expansion of the universe[2,7]. Part of this difficulty is related to the more sophisticated machinery of fine tuning. Nevertheless, over the last decade several non-minimal inflation scenario have been proposed to find reliable framework for issues such as graceful exit of inflationary phase and the observational constraints on the values that non-minimal coupling can attain in order to have successful inflationary scenario are discussed [8-16].  The recent astronomical observations, specially high precision data of WMAP3 [17] have opened new doors in the field of observational cosmology. As a result, these data have significant impact on the inflation paradigm. In this regard, these high precision data will give more accurate bounds on the values of non-minimal coupling in a typical non-minimal inflation model. The purpose of this letter is to study impact of WMAP3 and non-minimal inflation. Considering some well-known inflationary potentials, we explore new observational constraints on the values of non-minimal coupling to have successful non-minimal inflation. By definition of an effective scalar field potential, we show that there is a region in parameter space that inflation is driven by the non-minimal coupling term. A detailed study of spectral index and its running shows the spontaneous exit of inflationary phase ( without any mechanism ) in a suitable region of the parameter space. We also compare our results with the results of previous studies. This comparison reveals the present status of non-minimal inflation paradigm after WMAP3. ",
        "conclusions": ""
    },
    "0710/0710.3167_arXiv.txt": {
        "abstract": "In this brief proceedings article I summarize the review talk I gave at the IAU 246 meeting in Capri, Italy, glossing over the well-known results from the literature, but paying particular attention to new, previously unpublished material.  This new material includes a careful comparison of the apparently contradictory results of two independent methods used to simulate the evolution of binary populations in dense stellar systems (the direct $N$-body method of \\cite{2007ApJ...665..707H} and the approximate Monte Carlo method of \\cite{2005MNRAS.358..572I}), that shows that the two methods may not actually yield contradictory results, and suggests future work to more directly compare the two methods. ",
        "introduction": "Globular clusters are observed to contain significant numbers of binary star systems---so many, in fact, that they must have born with binaries (\\cite{1992PASP..104..981H}). Their presence in clusters is important for two complementary reasons.  Through super-elastic dynamical scattering interactions, they act as an energy source which may postpone core collapse, and may be the dominant factor in setting the core radii of observed Galactic globulars.  Similarly, the dense stellar environment and increased dynamical interaction rate in cluster cores is responsible for the high specific frequency of stellar ``exotica'' found in clusters, including low-mass X-ray binaries (LMXBs), cataclysmic variables (CVs), blue straggler stars (BSSs), and recycled millisecond pulsars (MSPs). ",
        "conclusions": "\\label{sec:summary}  In this proceedings article I have very briefly discussed the connection between binary stars and globular cluster dynamics, moving from basic physics to current research in the span of a few paragraphs.  A thorough, easily readable, and fairly recent discussion of the material can be found in \\cite{2003gmbp.book.....H}.  The primary new material presented here is a comparison in phase space of the seemingly contradictory binary population evolution simulations of \\cite{2005MNRAS.358..572I} and \\cite{2007ApJ...665..707H}, showing that they may in fact both represent the same underlying physics. In other words, new simulations must be performed to better compare the two very different methods."
    },
    "0710/0710.3351_arXiv.txt": {
        "abstract": "We present the results of \\spi\\ mid-infrared spectroscopic observations of two highly-obscured massive X-ray binaries: \\igrj\\ and \\gx.  Our observations reveal for the first time the extremely rich mid-infrared environments of this type of source, including multiple continuum emission components (a hot component with $T$ $>$ 700~K and a warm component with $T$ $\\sim$ 180~K) with apparent silicate absorption features, numerous \\ion{H}{1} recombination lines, many forbidden ionic lines of low ionization potentials, and pure rotational \\hh\\ lines.  This indicates that both sources have hot and warm circumstellar dust, ionized stellar winds, extended low-density ionized regions, and photo-dissociated regions. It appears difficult to attribute the total optical extinction of both sources to the hot and warm dust components, which suggests that there could be an otherwise observable colder dust component responsible for the most of the optical extinction and silicate absorption features. The observed mid-infrared spectra are similar to those from Luminous Blue Variables, indicating that the highly-obscured massive X-ray binaries may represent a previously unknown evolutionary phase of X-ray binaries with early-type optical companions.  Our results highlight the importance and utility of mid-infrared spectroscopy to investigate highly-obscured X-ray binaries. ",
        "introduction": "\\label{sec_intro}  Recently, a large number of highly-obscured (e.g., \\nh\\ $\\ge$ 10$^{23}$ cm$^{-2}$) massive X-ray binaries have been discovered by the \\integ\\ hard X-ray ($\\ge$ 15 keV) satellite \\citep{wet03}.  The prototypical case is \\igrj\\ which shows variable high obscuration in the X-ray, sometimes reaching \\nh\\ $\\simeq$ 2 $\\times$ 10$^{24}$ \\citep{coet03, walteret03}.  Its bright optical and near-infrared (IR) counterpart is an early B-type supergiant star with numerous emission lines \\citep{fc04}.  Interestingly the obscuration toward \\igrj\\ obtained in the optical and near-IR wavebands (\\av\\ $\\sim$ 18) is almost two orders of magnitude smaller than that inferred from the X-rays, which suggests that the extreme obscuration seen in the X-ray is intrinsic only to the X-ray source.  However, \\av\\ $\\sim$ 18 is still  greater than the interstellar obscuration, indicating the existence of substantial circumstellar material around the supergiant companion.  In order to fully understand the implications of the \\integ\\ discoveries --- specifically if they imply existence of a separate class of highly-obscured X-ray binaries --- we must investigate any similarities between the new \\integ\\ sources and previously known sources.  As already suggested by \\citet{ret03}, the X-ray pulsar \\gx\\ appears similar to the new \\integ\\ highly-obscured massive X-ray binaries: it has variable X-ray obscuration of \\nh\\ $\\simeq$ 10$^{23}$--10$^{24}$ cm$^{-2}$, the compact source is a neutron star, and the optical companion is an early B-type supergiant \\citep[or hypergiant;][]{ket95}, like \\igrj.  Recently, \\citet[][; hereafter Paper~I]{kmr06}, using optical, near-, and mid-IR ($\\le$~20~$\\mu$m) spectral energy distributions (SEDs), have found that both sources have strong mid-IR excesses that they identify as continuum emission from hot dust.  This is suggestive that they have very similar circumstellar material which may be related to their strong X-ray obscuration.  In this {\\em Letter}, we present the results of \\spi\\ mid-IR spectroscopic observations of \\igrj\\ and \\gx, showing that both sources indeed have very similar rich mid-IR properties previously unknown for X-ray binaries. ",
        "conclusions": "\\label{sec_dis_sum}  Our \\spi\\ spectroscopic observations of \\igrj\\ and \\gx\\ have revealed for the first time the rich mid-IR environment of highly-obscured X-ray binaries. This includes two dust components with prominent silicate absorption, numerous \\ion{H}{1} recombination lines, many forbidden ionic  lines, and pure rotational \\hh\\ lines. Based on the observed spectra, we infer the following components for \\igrj\\ and \\gx: (1) hot ($T$ $>$ 700~K) and warm ($T$ $\\sim$ 180~K) circumstellar dust; (2) ionized stellar winds responsible for the \\ion{H}{1} lines; (3) extended low-density ionized regions for the forbidden lines; and (4) photo-dissociated regions associated with the PAH, \\hh\\ and possibly the \\silii\\ line emission.  For the forbidden lines, all the detected lines have relatively low ionization potentials like in starburst galaxies where the radiation is relatively soft (compared with active galactic nuclei, for instance).  This may indicate that the illumination of hard X-rays from the central compact X-ray source is not primarily responsible for the forbidden line emission.  However, the inferred temperature for the radiation exciting the \\neii\\ and \\neiii\\ lines is hotter than the stellar photospheres, so there can be some contribution from the compact object. Considering that \\niii\\ and \\feii\\ were detected only in \\igrj\\ while \\feiii\\ was only in \\gx, the radiation field of \\igrj\\ may be softer than \\gx, as demonstrated by the small temperature difference between the two sources (see \\S~\\ref{sec_highres}).  Perhaps the most natural explanation for the origin of the hot dust component relies on dust formation in the dense outflows from the early-type companions of \\igrj\\ and \\gx, as is the case in most sgB[e] stars (i.e., B-type supergiants with forbidden emission lines). However, the origin of the warm circumstellar dust component is very uncertain, and B[e] stars seldom show evidence for this type of warm dust. If both the hot and warm dust components have spherical shell geometres around the central star, then the associated optical extinctions are: \\av\\ $\\simeq$ 4 $\\times$ 10$^{-4}$ $Q_{\\rm abs}$ ($M_{\\rm d}$/10$^{-6}$ \\msun) ($T_{\\rm d}$/100 K)$^4$ ($L_{\\rm UV}$/10$^{39}$ ergs s$^{-1}$), where $M_{\\rm d}$ and $T_{\\rm d}$ are the mass and temperature of the dust components, $Q_{\\rm abs}$ $<$ 1 is the dust absorption coefficient, and $L_{\\rm UV}$ is the ultra-violet luminosity of the central star. This gives the optical extinctions \\av\\ $\\ll$ 1 for both the hot and warm dust components of \\igrj\\ and \\gx. (Here we use 1 $\\times$ 10$^{39}$ ergs s$^{-1}$ for $L_{\\rm UV}$ for both sources.) Therefore, under the assumption of spherical shell geometry, both the hot and warm dust components of \\igrj\\ and \\gx\\ contribute very little to the total optical extinction, suggesting that the hot and warm dust components are {\\em NOT} strongly associated with the silicate absorption features. This also applies to the dust associated with the warm extended \\hh\\ gas since the optical extinction from this component is tiny (i.e., Av $\\ll$ 1; see \\S~\\ref{sec_highres}). What's then the origin of the silicate absorption features? If it is due to the foreground ISM, we would expect to see the $\\rm CO_2$ ice feature around 15 $\\mu$m, especially for \\igrj, based on the intensity ratio between the silicate absorption feature and ice feature found in the ISM \\citep{ket05}. The absence of the ice feature in our spectra supports the interpretation that the silicate absorption features are probably not associated with the foreground ISM. One possibility may be the existence of an undisclosed colder (e.g., $\\ll$ 100~K) circumstellar dust component which is responsible for the silicate absorption features and most of the optical extinction. We need further longer wavelength observatons to confirm this possibility. Considering that the 9.7~$\\mu$m silicate absorption feature represents the oxygen-rich material, the existence of the 9.7 micron silicate absorption feature may indicate that the origin of the potential colder dust component is related to the nucleosynthesis of the progenitors of \\igrj\\ and \\gx\\ (or their companions).  The optical/near-IR companion of \\igrj\\ is a sgB[e] star.  Such stars are known to have hot ($T$ $\\sim$ 1000~K) circumstellar dust, and probably are evolving into Luminous Blue Variables (LBVs) or Wolf-Rayet stars.  Our mid-IR spectra of \\igrj\\ and \\gx\\ are very similar to that of the LBV P~Cygni which shows many \\ion{H}{1} lines and forbidden ionic lines \\citep{let96b}.  The difference is that while the mid-IR continuum of P~Cygni is due to the free-free emission in stellar winds, the main mid-IR emission of \\igrj\\ and \\gx\\ is from multiple dust continua.  (The SEDs of both sources observed here are inconsistent with the free-free emission, as mentioned in Paper~I.) However, we note that some LBV stars have also been observed to have thermal mid-IR dust emission \\citep{let96a}. Based on the fact that B[e] stars seldom show mid-IR forbidden line emission, the B-type supergiant (or hypergiant) companions of \\igrj\\ and \\gx\\ may be in the evolutionary track to LBVs, implying that the extremely high obscuration seen in some massive X-ray binaries may be a phenomenon associated with the evolutionary phase.  This scenario is also consistent with the fact that the highly-obscured massive X-ray binary Cygnus~X-3 has a Wolf-Rayet star companion, together with the circumstellar dust emission of $T$ $\\sim$ 250~K \\citep{ket02}."
    },
    "0710/0710.4601_arXiv.txt": {
        "abstract": "We constructed some main-sequence mergers from case A binary evolution and studied their characteristics via Eggleton's stellar evolution code. Both total mass and orbital angular momentum are conservative in our binary evolutions. Assuming that the matter from the secondary homogeneously mixes with the envelope of the primary and that no mass are lost from the system during the merger process, we found that some mergers might be on the left of the zero-age main sequence as defined by normal surface composition (i.e helium content $Y=0.28$ with metallicity $Z=0.02$ for Pop I) on a colour-magnitude diagram(CMD) because of enhanced surface helium content. The study also shows that central hydrogen content of the mergers is independent of mass. Our simple models provide a possible way to explain a few blue stragglers (BSs) observed on the left of zero-age main sequence in some clusters, but the concentration toward the blue side of the main sequence with decreasing mass predicted by Sandquist et al. will not appear in our models. The products with little central hydrogen in our models are probably subgiants when they are formed, since the primaries in the progenitors also have little central hydrogen and will likely leave the main sequence during merger process. As a consequence, we fit the formula of magnitude $M_{\\rm v}$ and $B-V$ of the mergers when they return back to thermal equilibrium with maximum error 0.29 and 0.037, respectively.  Employing the consequences above, we performed Monte Carlo simulations to examine our models in an old open cluster NGC 2682 and an intermediate-age cluster NGC 2660. Angular momentum loss (AML) of low mass binaries is very important in NGC 2682 and its effect was estimated in a simple way. In NGC 2682, binary mergers from our models cover the region with high luminosity and those from AML are located in the region with low luminosity, existing a certain width. The BSs from AML are much more than those from our models, indicating that AML of low mass binaries makes a major contribution to BSs in this old cluster. Our models are corresponding for several BSs in NGC 2660. At the region with the most opportunity on the CMD, however, no BSs have been observed at present. {\\bf Our results are well-matched to the observations if there is $\\sim 0.5M_\\odot$ of mass loss in the merger process, but a physical mechanism for this much mass loss is a problem.} ",
        "introduction": "Much evidence shows that primordial binaries make an important contribution to blue stragglers (BSs) \\cite{fer03,dpa04,map04}. At present, a few BSs, i.e. F190, $\\theta$ Car, have already been confirmed to be in binaries by observations, and their formation may be interpreted by mass transfer between the components of a binary. Whereas in intermediate-age and old open and globular clusters, the number of observed close binaries among well-studied BSs is consistent with the hypothesis of binary coalescence. For example, Mateo et al. \\shortcite{mat90} made a comparison of the number of close binaries with the total number of BSs in NGC 5466 and found that it is an acceptable claim that all non-eclipsing BSs are formed as the result of mergers of the components in close binaries, though the possibility of other mechanisms to produce BSs cannot be ruled out due to the large uncertainties in their analysis. Monte-Carlo simulations of binary stellar evolution \\cite{pol94} also show that binary coalescence may be an important channel to form BSs in some clusters (e.g. with an age greater than 40 Myr). Meanwhile, the arguments in theory show that W UMa binaries (low-mass contact binaries) must eventually merge into a single star \\cite{web76,web85,ty87,mat90}. Observationally, the lack of radial velocity variations for most BSs further indicates that binary coalescence may be more important than mass transfer for BS formation \\cite{str93,pol94}. FK Comae stars are generally considered to be direct evidence for binary coalescence \\cite{str93}. The smallest mass ratio of components among observed W UMa systems to date is about 0.06. All of the above show that it is important to study the remnants of close binaries. However the merge process is complicated and the physics during the process is still uncertain. Recently, Andronov, Pinsonneault \\& Terndrup \\shortcite{apt06} studied the mergers of close primordial binaries by employing the angular momentum loss rate inferred from the spindown of open cluster stars. Their study shows that main sequence mergers can account for the observed number of single BSs in M67 and that such mergers are responsible for at least  one third of the BSs in open clusters older than 1 Gyr. The physics of mergers are limiting case treatments in the study of Andronov, Pinsonneault \\& Terndrup \\shortcite{apt06}. Based on previous studies of contact binaries and some assumptions, we construct a series of merger models in this paper, to study the structure and evolution of the models and show some comparisons with observations.  Case A binary evolution has been well studied by Nelson \\& Eggleton \\shortcite{nel01}. They defined six major subtypes for the evolution (AD, AR, AS, AE, AL and AN) and two rare cases (AG and AB). Three of the subtypes (AD, AR, AS) lead the binary contact as both components are main--sequence stars and two cases (AE and AG) reach contact with one or both components having left the main sequence. As there is no description for weird objects except for two merged main--sequence stars, merger products (except for two main-sequences stars) are generally assumed to have terminated their evolution \\cite{pol94}, i.e. they have left the main sequence and cannot be recognized as BSs. Here we are interested in the cases of two main-sequence stars, i.e. cases AS, AR and AD. If $t_{\\rm dyn}$, $t_{\\rm KH}$, $t_{\\rm MS}$ represent the dynamic timescale, thermal timescale and main sequence timescale of the primary (the initial massive star,*1), respectively, the following shows a simple definition of the three evolutionary cases: AD--dynamic Roche lobe overflow (RLOF), $\\dot{M}>M/t_{\\rm dyn}$; AR--rapid evolution to contact,  $\\dot{M}>M/t_{\\rm KH}, t_{\\rm contact} -t_{\\rm RLOF}(*1)<0.1t_{\\rm MS}(*1)$; AS--slow evolution to contact, $t_{\\rm contact} -t_{\\rm RLOF}(*1)>0.1t_{\\rm MS}(*1)$, where $t_{\\rm RLOF}$ and $t_{\\rm contact}$ are the ages at which RLOF begins and the binary comes into contact, respectively. In case AD, the core of the secondary spirals in quickly and stays in the center of the merger. The merger then has a chemical composition similar to that of the primary, resembling the result of smoothed particle hydrodynamic calculations \\cite{lrs96,sill97,sill01}. We therefore studied just the systems in cases AR and AS for this work.  \\section {Assumptions} Using the stellar evolution code devised by Eggleton \\shortcite{egg71,egg72,egg73}, which has been updated with the latest physics over the last three decades \\cite{han94,pol95,pol98}, we re-calculate the models of cases AS and AR with primary masses between 0.89 and $2M_\\odot$ until the systems become contact binaries. The structures of the primaries and the compositions of the secondaries are stored to construct the merger remnants.  Before the system comes into contact, the accreting matter is assumed to be deposited onto the surface of the secondary with zero falling velocity and distributed homogeneously all over the outer layers. The change of chemical composition on the secondary's surface caused by the accreting matter is \\begin{equation} {\\partial X_i / \\partial t }={(\\partial M /\\partial t)/[(\\partial M /\\partial t){\\rm d}t+M_{\\rm s}}] \\cdot (X_{i{\\rm a}}-X_{i{\\rm s}}), \\end {equation} where $\\partial M /\\partial t$ is the mass accretion rate, $X_{i{\\rm a}}$  and $X_{i{\\rm s}}$ are element abundances of the accreting matter and of the secondary's surface for species $i$, respectively, and $M_{\\rm s}$ is the mass of the outermost layer of the secondary. The value of $M_{\\rm s}$ will change with the moving of the non-Lagrangian mesh as well as the chosen model resolution, but it is so small ($\\sim 10^{-9}-10^{-12} M_{\\odot}$) in comparison with $(\\partial M /\\partial t){\\rm d}t$ ($\\sim 10^{-3}-10^{-5} M_{\\odot}$) during RLOF that we may ignore the effect of various $M_{\\rm s}$ on element abundances. Before and after RLOF, we get $\\partial X_{\\rm i}/\\partial t =0$ from the equation, which is reasonable in the absence of mixing \\cite{ch04}.  The merger models are constructed based on the following assumptions: (i) contact binaries with two main-sequence components coalesce finally and the changes of structures of individual components during coalescence are ignored; (ii) the matter of the secondary is homogeneously mixed with that of the primary beyond the core-envelope transition point, which separates the core and the envelope of the mass donor; (iii) the system mass is conserved.  Firstly, we present a brief discussion on these assumptions. Webbink \\shortcite{web76} studied the evolutionary fate of low-mass contact binaries, and found that a system cannot sustain its binary character beyond the limits set by marginal contact evolution ($\\mu =M_1/(M_1+M_2)=1.0$). He stated that a contact binary will very likely coalesce as the primary is still on the main sequence in a real system. Up to now, it is widely believed that case AD probably leads to common envelope, spiral-in, and coalescence on quite a short timescale. The final consequences of AS and AR are not very clear, but Eggleton \\shortcite{egg00} pointed out that systems undergoing AR or AS evolution may maintain a shallow contact (perhaps intermittently) as the mass ratio becomes more extreme, and finally coalesce. Recent study on W UMa (Li, Zhang \\& Han, 2005) also shows that these systems will be eventually coalescence. The merged timescale, i.e. the time from a binary contact to coalescence, is important here. If it is too long, the structures of both components will change remarkably and the system may have not completed coalescence within the cluster age. There are many {\\bf conflicting estimates} for the timescale, however, from observations and theoretical models of the merger process. {\\bf Early observational estimates range from} $10^7$--$10^8$ yr in various environments \\cite{van79,eggen89}. The following study explored the average age about $5 \\times 10^8$ yr \\cite{van94,dry02}. Bilir et al. \\shortcite{bil05} pointed out that the age difference between field contact binaries and chromospherically active binaries, 1.61 Gyr, is likely an upper limit for the contact stage by assuming an equilibrium in the Galaxy, whereas the study of W UMa by Li, Han \\& Zhang \\shortcite{lhz04} suggested a much longer timescale, about 7 Gyr. We adopt the empirically estimated values in this paper {\\bf (i.e from $5 \\times 10^7$ to $1 \\times 10^9$ yrs)} and ignore the changes of structure of individual components during merger process. For low-mass contact binaries, the common envelope is convective \\cite{web77}, and the matter in it is thus homogeneous. If a system mimics shallow contact during coalescence, it is reasonable to assume that the matter of the secondary mixes with the envelope homogeneously. Van't Veer \\shortcite{van97} found that the mass loss from the system during coalescence is at a rate of about $2 \\times 10^{-10}M_\\odot{\\rm yr}^{-1}$ by observations. If we consider that the coalescence time is $5 \\times 10^8$ yr in a binary, only $0.1M_\\odot$ is lost from the system as the binary finally becomes a single star. We then roughly assume that the mass is conservative during coalescence. However mass loss might be an important way to carry orbital angular momentum away from the binary in this process.  Secondly, we discuss the choice of the core-envelope transition point which separates the core and the envelope in the primary. Many characteristics of the merger are relevant to the choice, e.g. the chemical composition in the envelope, evolutionary track on Hertzsprung-Russel diagram, and some observational characteristics. Unfortunately, one cannot find the core-envelope transition point in a main-sequence star as easily as in evolved stars because the density profile, as well as many other thermodynamic quantities (entropy, pressure, temperature etc.), is smooth and does not have a deep gradient for main-sequence stars. Chen \\& Han \\shortcite{ch05} studied the influences of core-envelope transition point on the mergers of contact binaries with two main-sequence components. They found that one may ignore the effects which result from different choices of the transition point on colours and magnitudes of the merger if it is outside the nuclear reaction region of the primary, which is commonly considered as the nearest boundary of the secondary reaches in cases AS and AR. In this paper, the core-envelope transition is determined as the point within which the core produces 99 per cent of total luminosity. This choice is generally outside the nuclear reaction regions and has little effect on the final results.  Finally the merger remnant is constructed as follows: it has the total mass of the system and a chemical composition within $M_{\\rm 1c}$ similar to the core of the primary. The chemical composition in the envelope of the merger is given by \\begin{equation} X_i=(M_{i2}+M_{i1\\rm b})/(M_2+m_{\\rm b}), \\end{equation} where $M_{i2}$ and $M_{i1\\rm b}$ denote the total masses of species $i$ of the secondary and of the primary's envelope, respectively. $m_{\\rm b}$ is the envelope mass of the primary. There might be a region in which the helium abundance is less than that of the outer region. The matter in this region then has a lower mean molecular weight than that in the outer region. This results in secular instability and thermohaline mixing \\cite{kip80,ulr72}. We include it as a diffusion process in our code \\cite{ch04}.  In the models of Nelson \\& Eggleton \\shortcite{nel01}, both total mass and angular momentum are conservative. It was mentioned by the authors, however, that these assumptions were only reasonable for a restricted range of intermediate masses, i.e spectra from about G0 to B1 and luminosity class III-V. Observationally, some low mass binaries with late-type components show clear signs of magnetic activity, which indicates that the systems evolve by way of a scenario implying angular momentum loss (AML) by magnetic braking \\cite{mes84}. Magnetized stellar winds probably do not carry off much mass, but they are rich in angular momentum because of magnetic linkage to the binaries. For close binaries, rotation is expected to synchronize with orbital period, so AML is at the expense of the orbital angular momentum, resulting in orbital decaying and the components approaching each other. A detached binaries, then, may become contact and finally coalesce at or before the cluster age \\cite{ste95}. There are a number of subjects including the treatment of AML \\cite{lhz04,ste06,mk06,dek06}. For simplicity, the conservative assumption is also adopted in our binary evolutions. In old clusters, however, AML of low mass binaries is very important and {\\bf we estimate its importance in another way (see section 4.2).} ",
        "conclusions": "In Sect.4, we notice that the timescale from contact to complete coalescence, $t_{\\rm cc}$, strongly affects the initial parameter space of primordial binaries which eventually produce single BSs in a cluster. On the other hand, there are {\\bf some conflicting estimates} for $t_{\\rm cc}$ from observations and theoretical models. {\\bf In this paper we adopt empirical values, i.e. $t_{\\rm cc}$ is short in comparison to the evolution timescale of both components in a binary, and ignore the changes during the merger process. In this section, we will first discuss the consequences of a long $t_{\\rm cc}$. Because of evolution of both components during the merger process, the primaries have lower central hydrogen content and the matter from the secondaries have larger He content. The former results in a redder colour for the mergers while the latter makes the mergers bluer. So the final positions of the mergers are possibly similar to those shown in this paper, except that the primaries have left the main sequence at final coalescence. This case will appear in the mergers with little central hydrogen. For example, a star with $2M_\\odot$ may evolve from ZAMS to exhausted of central hydrogen in $10^9$ yr, and then none of the mergers from binaries with primary' masses larger than $2M_\\odot$ will be on the main sequence if $t_{\\rm cc}= 1 \\times 10^9$ yr. For the primaries with very little hydrogen in the center at contact, the mergers may never be on the main sequence even in the cases of short $t_{\\rm cc}$. The long $t_{\\rm cc}$ also delays the appearance of the mergers and shortens their timescales on the main sequence. The latter has not been exactly expressed in our models, and therefore we just see that the mergers from a long $t_{\\rm cc}$ have larger luminosities as shown in section 4.}  {\\bf In our binary evolutions, we have not included AML, which exists in low-mass binaries and may be the main course making the binaries change from detached to contact and finally coalesce, resulting in a large contribution to BSs in old clusters, e.g NGC 2682. In young and intermediate-age clusters, however, AML has little contribution to the birthrate of BSs, since (a) the time is not long enough for binaries to go from detached to contact and (b) the mass of the mergers is probably less than the turnoff of the cluster even though their parents may coalesce in the cluster age. So, we simply estimated the effect of AML in NGC 2682 while negecting it in NGC 2660.}  The mass loss during the merger process can also affect our result, mainly the location on the CMD of the products. As shown in NGC 2660, no BSs have been observed in the region with the most opportunity from our models. Because of mass loss, the mergers will be fainter than those given in the paper. However the faintness will be slight since the mass loss is not vast from both observations and smooth particle hydrodynamic simulations \\cite{lrs96,sill97,sill01}. The lost mass may carry some angular momentum out from the parent binary. By analyzing the BS spectra from Hubble Space Telescope (there is an apparent continuum deficit on the short-wavelength side of Balmer discontinuity ), De Marco et al. \\shortcite{dm04} argued that some BSs might be surrounded by a circumstellar disk. However, Porter \\& Townsend \\shortcite{pt05} showed that the flux deficits may be attributed wholly to rapid rotation. The rotation rates needed are of the order of those found in the study of De Macro et al. \\shortcite{dm05}. Whether the flux deficits shortward of the Balmer jump are induced by a circumstellar disk or rapid rotation, it provides a possible explanation for the orbital angular momentum of the system after coalescence. Such a large mass loss as shown in NGC 2660 (about $0.5 M_\\odot$), however, is a problem and should be explained reasonably in physics.  Based on some assumptions, we studied the mergers of close binaries from AS and AR evolution by detailed evolutionary calculations. The products from our models may stay on the left of the ZAMS and have no central concentration with decreasing mass. Because of the {\\bf development} of the convective core, the mergers with little central hydrogen (less than 0.01) in our models have unusually long timescales on the main sequence ($\\sim 10^8$ yrs). These objects are probably subgiants as they are formed, since the primaries in the progenitors also have little central hydrogen and may have left the main sequence during merger process.  The mergers from our models stay on the main sequence for a timescale in order of $10^8$ yrs. Some low-mass mergers may stay on the MS for about $10^9$ yrs. The timescale is similar to that of W UMa stars from observations, and therefore we may roughly estimate the contribution to BSs from AS and AR via the number of W UMa systems in a cluster. The estimation, however, is not absolutely since both of the two timescales have wide ranges and large uncertainties, and we cannot rule out other methods for creating W UMa systems except for AS and AR. Comparison to observations indicates that our models (binary coalescence from AS and AR) are not important for the produce of BSs in old open clusters, while likely play a critical role in some younger open clusters.  We performed Monte Carlo simulations to examine our models in an old open cluster NGC 2682 and in an intermediate-age cluster NGC 2660. The effect of AML was estimated in NGC 2682 in a simple way, where the mergers are replaced with ZAMS models. In NGC 2682, binary mergers from our models cover the region with high luminosity and those from AML are located in the region with low luminosity, existing a certain width. The BSs from AML are much more than those from our models, indicating that AML of low mass binaries makes a major contribution to BSs in this cluster. Our models are corresponding for several BSs in NGC 2660. In the region with the most opportunity on CMD, however, no BSs have been observed. {\\bf Our results are well-matched to the observations if there are $\\sim 0.5M_\\odot$ of mass loss in the merger process, but a physical mechanism for this much mass loss is a problem.}"
    },
    "0710/0710.5607_arXiv.txt": {
        "abstract": "Extreme gravitational lensing refers to the bending of photon trajectories that pass very close to supermassive black holes and that cannot be described in the conventional weak deflection limit. A complete analytical description of the whole expected phenomenology has been achieved in the recent years using the strong deflection limit. These progresses and possible directions for new investigations are reviewed in this paper at a basic level. We also discuss the requirements for future facilities aimed at detecting higher order gravitational lensing images generated by the supermassive black hole in the Galactic center. ",
        "introduction": "All known astrophysical cases of bending of photon trajectories by gravitational fields can be interpreted using the conventional weak deflection paradigm, originally formulated by Einstein \\cite{Ein}. Yet it is clear that photons passing very close to black holes suffer very large deflections, which must be addressed in a full general relativistic context. However, integrating null geodesics in full general relativity usually leads to very complicated analytical formulae or heavy numerical codes, which typically obscure the basic physical interpretation. The Strong Deflection Limit (SDL) is an analytical tool that allows to derive simple analytical formulae describing the higher order images appearing in extreme gravitational lensing. In this work we review the main progresses achieved in the recent years in this technique and its application to the most interesting physical case: the black hole in the Galactic center (Sgr A*). All relevant formulae derived in previous works are recalled and restated here in the simplest possible form.  This work is structured as follows: \\S~2 explains the basic phenomenology of extreme lensing; \\S~3 introduces the SDL for spherically symmetric black holes; \\S~4 generalizes the method to spinning black holes; \\S~5 applies the formulae to Sgr A*, examining possible sources for extreme gravitational lensing; \\S~6 generalizes the method to sources very close to the black hole; \\S~7 contains the conclusions. ",
        "conclusions": "Extreme gravitational lensing is a very spectacular though elusive phenomenon, which demands a great effort in order to be observed. The most promising extreme lens is represented by the supermassive black hole in the Galactic center, which anyway requires resolutions of order microarcseconds for the direct observation of higher order images. Suitable known sources for gravitational lensing are giant stars in the IR band and low mass X-ray binaries in the X-ray band.  In this paper we have reviewed recent results on gravitational lensing in the Strong Deflection Limit, which is an approximation devoted to the analytical determination of the properties of higher order images. Within this framework, it has been proved that the logarithmic divergence in the deflection angle is a universal feature of all black hole metrics. There exists a general method to determine the SDL coefficients of the deflection angle for any given spherically symmetric black hole metric. This method has been applied to numerous metrics. Higher order images are formed by photons performing one or more loops around the black hole before reaching the observer. The position and the flux ratios of higher order images depend on the specific metric through the SDL coefficients and are thus able to track any possible deviations from general relativity in the strong field regime.  The SDL can be extended to spinning black holes, where several new features emerge, such as extended and shifted caustics and additional images. The position and the shape of the caustics can be expressed by simple analytical formulae and the lens equation for sources near to caustics, which represent the most physically relevant situation, can be solved.  The extension of the SDL to the case of sources very close to black holes opens the way to even more interesting investigations of extreme gravitational lensing. In fact, whereas direct observation of higher order images is very difficult and probably very far to come in the future, the contribution of higher order images to currently observed phenomena involving sources very close to black holes is already measurable at present time. Light curves of flares born in the accretion disk and spectral measurements of Iron K-lines are heavily influenced by the presence of higher order images. With the SDL setup in its updated version, it is possible to study such phenomena, bearing in mind that the large model dependence remains the main uncertainty in all theoretical attempts to interpret the complicated physics of the supermassive black holes environment."
    },
    "0710/0710.5431_arXiv.txt": {
        "abstract": "Warm dark matter is consistent with the observations of the large-scale structure, and it can also explain the cored density profiles on smaller scales. However, it has been argued that warm dark matter could  delay the star formation.  This does not happen if warm dark matter is made up of keV sterile neutrinos, which can decay into X-ray photons and active neutrinos. The X-ray photons have a catalytic effect on the formation of molecular hydrogen, the essential cooling ingredient in the primordial gas. In all the cases we have examined, the overall effect of sterile dark matter is to facilitate the cooling of the gas and to reduce the minimal mass of the halo prone to collapse.  We find that the X-rays from the decay of keV sterile neutrinos facilitate the collapse of the gas clouds and the subsequent star formation at high redshift. ",
        "introduction": "Both cold and warm dark matter models agree with the observed structure on the large scales.  However, there are several inconsistencies between the predictions of the cold dark matter (CDM) model and the observations~\\cite{cdm_problems}.   The low cutoff in dark matter contents of dwarf spheroids, the smoothness of our dark matter halo, and the old globular clusters (observed in Fornax) resisting the infall into the center by dynamical friction~\\cite{cdm_problems}, all can be explained by warm dark matter (WDM) because it suppresses the structure on scales that are smaller than the free-streaming length.  While the suppression of the small-scale structure is desirable, it has been argued that ``generic'' WDM (for example, gravitino) can slow down structure formation and delay reionization of the universe, which can lead, in turn, to an inconsistency with the reionization redshift obtained by the WMAP \\cite{Yoshida}. This problem can be alleviated in the case of the WDM  in the form of sterile neutrinos with mass of several keV and a small mixing angle with the ordinary neutrino \\cite{bier,stas,stas07} because such sterile neutrinos can decay and produce photons that catalyze the formation molecular hydrogen and speed up the star formation.  In the absence of metals, gas cooling is mainly due to the collisional excitation of H$_{2}$, its subsequent spontaneous de-excitation, and photon emission. In the primordial gas clouds, hydrogen molecules can be formed only in reactions involving $e^{-}$ or H$^+$ as a catalyst. Thus, an X-ray radiation can increase the production of the H$_{2}$ by enhancing the ionization fraction, which subsequently leads to speed up of the gas cooling and star formation. Although sterile neutrinos are stable on cosmological time scales, they nevertheless decay. The decay channel important for us is that of decay into one active neutrino and one photon, i.e., $\\nu_s \\rightarrow \\nu_a \\gamma$, where the photon energy is half of the sterile neutrino mass, $E_{0} \\approx m_{s}c^2/2$. These decays produce an X-ray background radiation that increases the production of molecular hydrogen and can induce a rapid and prompt star formation at high redshift.  Sterile dark matter has a firm motivation from particle physics~\\cite{dw,Kusenko:2006rh}.  The discovery of the neutrino masses implies the existence of right-handed gauge-singlet fields, all or some of which can be lighter than the electroweak scale.  These sterile neutrinos can be produced in the early universe by different mechanisms, for example, from neutrino oscillations~\\cite{dw} or from the Higgs decays~\\cite{Kusenko:2006rh}, or from the couplings to a low-scale inflaton~\\cite{Shaposhnikov:2006xi}.   The same particles, produced in a supernova, could account for the supernova asymmetries and the pulsar kicks~\\cite{kus}, and can play a role in the formation of super-massive black holes in the early universe~\\cite{puzzle}.  We will examine the thermal evolution of the gas clouds, taking into account both effects of the sterile neutrino decays, namely, the ionization and heating of the gas. We follow the evolution of the baryonic top-hat overdensity, the gas temperature and the H$_{2}$ and $e^{-}$ fraction. In order to perform the calculation we have incorporated to our previous code \\cite{stas}, the effects of sterile neutrino decays within collapsing halos and absorption of the X-ray background from sterile neutrinos by He atoms in the intergalactic medium. Our goal is to juxtapose the evolution of the gas temperature in the primordial clouds in the CDM model and the WDM model with keV sterile neutrinos and estimate of the minimal mass able to collapse at a given redshift. ",
        "conclusions": ""
    },
    "0710/0710.2027_arXiv.txt": {
        "abstract": "{ The exciting results from H.E.S.S. point to a new population of $\\gamma$-ray sources at energies E$>$10~TeV, paving the way for future studies and new discoveries in the multi-TeV energy range. Connected with these energies is the search for sources of PeV cosmic-rays (CRs) and the study of multi-TeV $\\gamma$-ray production in a growing number of astrophysical environments. {\\em TenTen} is a proposed stereoscopic array (with a suggested site in Australia) of modest-sized (10 to 30m$^2$) Cherenkov imaging telescopes with a wide field of view (8$^\\circ$ to 10$^\\circ$ diameter) optimised for the E$\\sim$10 to 100~TeV range. {\\em TenTen} will achieve an effective area of $\\sim$10~km$^2$ at energies above 10~TeV. We outline here the motivation for {\\em TenTen} and summarise key performance parameters. }  \\email{growell@physics.adelaide.edu.au}   \\begin{document} ",
        "introduction": " ",
        "conclusions": "We have outlined the motivation for a new array of IACTs achieving 10~km$^2$ at $E>10$~TeV and described some important performance parameters. This array, known as {\\em TenTen}, could also be considered complementary to future MeV to TeV $\\gamma$-ray instruments such as GLAST, HESS-II and MAGIC-II. Studies are currently underway to further optimise individual telescopes (optics, electronics, camera design), overall layout parameters, and site potential in Australia.   \\scriptsize"
    },
    "0710/0710.5788_arXiv.txt": {
        "abstract": "The observed increase in star formation efficiency with average cloud density, from several percent in whole giant molecular clouds to $\\sim30$\\% or more in cluster-forming cores, can be understood as the result of hierarchical cloud structure if there is a characteristic density as which individual stars become well defined. Also in this case, the efficiency of star formation increases with the dispersion of the density probability distribution function (pdf). Models with log-normal pdf's illustrate these effects. The difference between star formation in bound clusters and star formation in loose groupings is attributed to a difference in cloud pressure, with higher pressures forming more tightly bound clusters. This correlation accounts for the observed increase in clustering fraction with star formation rate and with the observation of Scaled OB Associations in low pressure environments. ``Faint fuzzie'' star clusters, which are bound but have low densities, can form in regions with high Mach numbers and low background tidal forces.  The proposal by Burkert, Brodie \\& Larsen (2005) that faint fuzzies form at large radii in galactic collisional rings, satisfies these constraints. ",
        "introduction": "\\label{sect:intro}  Stars form in concentrations with a range of densities, from star complexes, OB associations, and T Tauri associations at the low end to compact clusters and super-star clusters (SSC) at the high end. The overall structure is usually hierarchical (Scalo 1985; Feitzinger \\& Galinski 1987; Ivanov et al. 1992; Gomez et al. 1993; Battinelli, Efremov \\& Magnier 1996; Bastian, et al. 2005, 2007; Elmegreen et al. 2006; see reviews in Efremov 1995; Elmegreen et al. 2000, Elmegreen 2005), and this hierarchy continues even inside the youngest clusters (Testi et al. 2000; Heydari-Malayeri et al. 2001; Nanda Kumar, Kamath, \\& Davis 2004; Smith et al. 2005a; Gutermuth et al. 2005; Dahm \\& Simon 2005; Stanke, et al. 2006; see review in Allen et al. 2006).  Most likely, the hierarchy in stars comes from a hierarchy in the gas (St\\\"utzki et al. 1998; Dickey et al. 2001), which is the result of turbulence compression and gravitational contraction that is self-similar over a wide range of scales (see review in Mac Low \\& Klessen 2004). Clusters form in the densest parts of this gas and lose their initial sub-structure as the stellar orbits mix (for simulations of star formation in clusters, see Klessen \\& Burkert 2000; Bonnell, Bate, \\& Vine 2003; Li, et al. 2004; Tilley \\& Pudritz 2004; Klessen et al. 2005; V\\'azquez-Semadeni, Kim, \\& Ballesteros-Paredes 2005; Bonnell \\& Bate 2006; Li \\& Nakamura 2006).  Hunter (1999) and Ma\\'iz-Apell\\'aniz (2001) noted that some massive star-forming regions (which they called ``scaled OB associations'', or SOBA's) do not form dense clusters while others with the same total mass do (e.g., the SSC's). We would like to understand this difference. Obviously the density of the gas is involved, as dense clusters require dense gas, but the distinction between SSCs and SOBAs should also be related to the efficiency of star formation, because clusters forming at low efficiency disperse quickly when the gas leaves (Lada, Margulis \\& Dearborn 1984). At very low efficiency, stars form individually without passing through an embedded cluster phase. Recent observations of giant molecular clouds (GMCs) show star formation at both high and low densities, with some embedded stars in dense clusters and others more dispersed (Megeath et al. 2004; J{\\o}rgensen et al. 2006, 2007).  Larsen \\& Brodie (2000) discovered ``faint fuzzies'' at intermediate radii in the disks of the S0 galaxies NGC 1023 and NGC 3384, and suggested they are old clusters with unusually large radii (7-15 pc) and low densities. They are gravitationally bound because of their large ages (8-13 Gyr; Brodie \\& Larsen 2002), and they are as massive as SSC's and halo globular clusters. Such clusters appear to represent an intermediate stage between dispersed and bound star formation. They appear to be too low in average density to have had time for core collapse and envelope expansion as in models of globular clusters by Baumgardt et al. (2002).  Burkert, Brodie \\& Larsen (2005) suggested they formed in a collisional ring interaction between two galaxies.  What determines the relative proportion of dispersed and clustered star formation? Larsen \\& Richtler (2000) showed that clustering on a galactic scale, measured as the fraction of uv light in the form of massive young clusters, increases as the star formation rate increases. This could be the result of a selection effect if starbursts are active for less than a cluster destruction time. On the other hand, the clustering fraction could depend on pressure. Higher pressure makes the cool phase of gas denser, which promotes more clustering, and locally high pressures trigger star formation on the periphery of GMCs, making clusters in the dense gas (e.g., Comer\\'on, Schneider, \\& Russeil  2005; Zavagno et al. 2006). The Larsen \\& Richtler correlation could then follow from the mutual correlation between pressure and star formation rate with gas column density. Faint fuzzies are a counter example, however: they formed bound but presumably at low pressure to have such low central column densities now. Can all of these clustering types be understood as a continuum of properties in a universal physical model?  There have been several attempts to explain the difference between clustered and distributed star formation based on numerical simulations. Klessen, Heitsch, \\& Mac Low (2000) suggested that clusters form in non-magnetic gas when the turbulence driving scale is large. Heitsch, Mac Low \\& Klessen (2001) noted that non-magnetic turbulence driven on large scales produces a clustered collapse, while magnetic turbulence in supercritical clouds produces a more distributed collapse. Mac Low (2002) suggested that stars form in clusters when there is no turbulent support and they form disbursed when there is. V\\'azquez-Semadeni, Ballesteros-Paredes, \\& Klessen (2003) suggest this transition from no global turbulent support to support corresponds to an increase in the Mach number and a decrease in the sonic scale, which is the length where the size-linewidth relation gives a Mach number of unity. Large sonic scale (low Mach number) corresponds to a sonic mass larger than the thermal Jeans mass, which means a lack of global support and the formation of a cluster. Low sonic scale corresponds to the dispersed formation of stars, one for each tiny compressed region where the mass exceeds the local thermal Jeans mass. Li, Klessen \\& Mac Low (2003) suggested that the equation of state determines the stellar clustering properties: soft equations produce dense clusters while hard equations produce isolated stars. These suggestions all apply to initially uniform media. External compression of a cloud into a massive dense core can also make a cluster; most young clusters are in high-pressure regions like OB associations.  Simulations of turbulent media produce stars in compressed regions that act as seeds for the small-scale gravitational collapses that follow (V\\'azquez-Semadeni, Ballesteros-Paredes, \\& Klessen 2003; Clark \\& Bonnell 2005). These simulations also have probability distribution functions (pdfs) for density that are either log-normal or log-normal with a power-law tail at high density, especially when self-gravity is important (e.g., Li, Klessen \\& Mac Low 2003).  The efficiency of star formation is then proportional to the fraction of the gas in a dense form. Here we examine variations in this fraction as functions of average density and velocity dispersion, and as a function of the local density inside a cloud. The efficiency is taken to be the ratio of the stellar mass to the gas mass during a complete star-forming event. It generally increases with cloud density from a few percent in GMCs (Williams \\& McKee 1997) to several tens of percent in cluster-forming cores (e.g., Lada \\& Lada 2003). It may reach $\\sim50$\\% or more inside the densest star-forming cores. We explain this increase as a result of hierarchical structure, regardless of the dynamics and mechanisms of star formation, and we show that for log-normal or similar density pdf's, as expected in turbulent media, the mass fraction of regions with high efficiency increases with the Mach number and, independently, with the average density. This result may explain the Larsen \\& Richtler (2000) correlation as well as the observed variations in clustering properties with pressure. We also show that at high Mach number, bound clusters can form with relatively low average densities, thereby explaining faint fuzzies. These are all consequences of star formation in the dense cores of clouds that are structured by turbulence. They result primarily from the geometry of the gas, which is somewhat universal, and should be nearly independent of the gas dynamics or the strength of the magnetic field.  We make an important assumption that gravitational contraction and star formation can occur in regions that are either larger or smaller than the sonic scale. This means we assume that contraction to one or more stars can occur in a supersonically turbulent region. V\\'azquez-Semadeni, Ballesteros-Paredes, \\& Klessen (2003) suggest that if a cloud is supported by turbulence, then only regions smaller than a sonic scale and more massive than the thermal Jeans mass are unstable to form stars. Padoan (1995) was the first to consider this condition. However, clouds are probably not supported for any significant time by turbulence, and even if they were, it is only necessary that a clump mass exceed the turbulent Jeans mass for self-gravitational forces to exceed inertial forces. GMCs for example, have comparable self-gravitational and turbulent energy densities and yet are much larger than the sonic length. Our assumption is contrary to that in Krumholz \\& McKee (2005), who assume the same as V\\'azquez-Semadeni et al..  We are consistent with McKee \\& Tan (2003), however, as they consider the collapse of a highly turbulent core to make a massive star. Saito, et al. (2006), for example, observe star formation in massive turbulent cores. Thus the sonic length should not provide a threshold for star formation.  Our primary condition for star formation is that the gas density exceed some fixed value, taken here to be $10^5$ cm$^{-3}$. This is the density at which HCN gives a nearly constant star formation rate per unit gas mass (Gao \\& Solomon 2004a,b; Wu et al. 2005) and at which a variety of microscopic processes conspire to shorten the magnetic diffusion time (Elmegreen 2007). Regions with this density should have a wide range of Mach numbers but a nearly universal efficiency, according to the Gao \\& Solomon and Wu et al. observations. The density of $\\sim10^5$ cm$^{-3}$, when converted to 5900 M$_\\odot$ pc$^{-3}$, is also typical for star clusters, as most of those surveyed by Tan (2007), which span a factor of $10^5$ in mass, have about this average density. The cluster density equals the gas density times the efficiency, and the efficiency has to exceed $\\sim10-30$\\% for a bound cluster to form. Thus, the gas density for both cluster formation and high efficiency appears to be around $10^5$ cm$^{-3}$ or, possibly, $10^6$ cm$^{-3}$. We note that the long timescale derived for HCN gas as the ratio of the total HCN mass divided by the total star formation rate (Gao \\& Solomon 2004b; Wu et al. 2005) is not the duration of star formation in any one place, but is the HCN consumption time. As long as a new HCN region is formed somewhere each time an old HCN region disperses, the HCN consumption time can be long even when each region of star formation lasts for a short time (see Elmegreen 2007 for a discussion of star formation timescales).  As a result, the average efficiency of star formation in any one HCN region can be moderately large even if the average HCN consumption time is long. An average efficiency of $\\sim5-10$\\% would be reasonable considering that most star-forming regions leave unbound clusters after the gas leaves (Lada \\& Lada 2003), most regions are observed at half their total ages, and star formation typically accelerates over time (Palla \\& Stahler 2000). With constant acceleration, only $\\sim1/4$ of the total stars form in the first 1/2 of the total time. The efficiency that determines whether a bound cluster will remain is taken here to be 14.4\\% (see below), and the peak efficiency in a single-star core is taken to be 50\\% at $10^5$ cm$^{-3}$ density. These values are uncertain and are used here only to illustrate how cluster formation might scale with the velocity dispersion and density of the cool component of the interstellar medium. ",
        "conclusions": "Stars form bound clusters where the total efficiency is high.  The efficiency should be proportional to the mass fraction of a cloud in the form of dense cores where the individual stars form. This mass fraction increases with the cloud density in hierarchical clouds. It also increases with the Mach number of the turbulence because stronger shocks at higher Mach numbers compress the gas to a wider range of densities. As a result, a log-normal density pdf becomes flatter below the threshold density of star formation, and this means the mass fraction is higher at each density below this value. When the Mach number is high (really, when the dispersion of the density pdf is high), the efficiency is high even at a fairly low average cloud density, and so a high fraction of star formation ends up bound. An extreme example of this trend is the type of cluster called a faint fuzzie. We propose that faint fuzzies form in moderately low density clouds with moderately high Mach numbers. The galactic tidal ring environment proposed by Burkert et al. (2005) is an example of a region that would have such clouds.  The mass fraction of clumps at a particular high density of star formation also increases with the average density of the ISM because then the whole density pdf shifts toward higher values. The combination of high densities and high Mach numbers, characteristic of starburst regions, makes for a high fraction of star formation in bound clusters.  On the other hand, relatively low Mach numbers and/or low densities should produce stars in a more dispersed way, as in the low-density regions of molecular clouds or in low surface brightness galaxies and regions of galaxies. This combination of parameters corresponds to a low ISM pressure, so we infer that low pressure regions, which means those with a low gas column densities and low star formation rates per unit area, should produce relatively fewer bound clusters and relatively more unbound associations.  There is a physical explanation for the trends discussed here. Consider a moderately low density cloud with a low turbulent Mach number. The compressions inside that cloud will be modest and most of them will not reach a level of density or enhanced magnetic diffusion rate that allows gravitational collapse before the compression ends.  Then few stars will form and the efficiency will be low (unless the cloud continues to makes these weak compressions for a very large number of crossing times, which seems unlikely). Now consider this same cloud with the same size and average density but with a higher Mach number (this will require a higher ISM pressure). The stronger compressions will more easily produce high-density cores in which gravity is important and magnetic diffusion is fast. More stars will form and the efficiency will be high. The only difference between these two examples is the Mach number, which affects the range of densities in the compressed regions. In terms of the density pdf, higher Mach numbers produce a broader pdf at the same average density.  In terms of star formation in bound clusters, high pressure regions having clouds with high Mach number produce a higher fraction of their stars in bound clusters.  The model predicts that young bound clusters in starburst regions, or in regions of high ISM pressures or Mach numbers, will have a wider range of densities than bound clusters in more quiescent, low-pressure regions. This is partly because the cloud densities should be higher in starburst regions, so the resulting clusters can be denser overall, but it is also because the lower-density parts of starburst clouds can form stars with a high efficiency, leaving bound clusters when the gas leaves rather than dispersed OB associations. As there is generally more mass at low density than at high density, the net distribution of cluster density could shift toward lower values when the ISM density pdf is broad. According to the model, this density shift is relative to the mean ISM density, and it should be measured only before significant cluster expansion. Low surface brightness galaxies should make a preponderance of unbound OB associations rather than bound clusters, and the young clusters they form should have a relatively narrow range of densities compared to normal."
    },
    "0710/0710.0955_arXiv.txt": {
        "abstract": "{} {We present the optical classification and redshift of 348 X-ray selected sources from the XMM-Newton Bright Serendipitous Survey (XBS) which contains a total of 400 objects (identification level = 87\\%). About 240 are new identifications. In particular, we  discuss in detail the classification criteria adopted for the Active Galactic Nuclei population.} {By means of systematic spectroscopic campaigns and through the literature search we have collected an optical spectrum for the large majority of the sources in the XBS survey and applied a well-defined classification ``flow-chart''. } {We find that the AGN represent the most numerous population  at the flux limit of the XBS survey ($\\sim$10$^{-13}$ erg cm$^{-2}$ s$^{-1}$) constituting 80\\% of the  XBS sources selected in the 0.5-4.5 keV energy band and 95\\% of the ``hard'' (4.5-7.5 keV) selected objects. Galactic sources populate significantly the 0.5-4.5 keV sample (17\\%) and only marginally (3\\%) the 4.5-7.5 keV sample. The remaining sources in both samples are clusters/groups of galaxies and normal galaxies (i.e. probably not powered by an AGN). Furthermore, the percentage of type~2 AGN (i.e. optically absorbed AGNs with A$_V>2$mag) dramatically increases going from the 0.5-4.5 keV sample (f=N$_{AGN 2}$/N$_{AGN}$=7\\%) to the 4.5-7.5 keV sample (f=32\\%). We finally propose two simple diagnostic plots that can be easily used to obtain the spectral classification  for relatively low redshift AGNs even if the quality of the spectrum is not good. } {} ",
        "introduction": "In the last few years {\\it XMM-Newton} and {\\it Chandra} telescopes have represented an excellent tool to survey the hard X-ray sky at all fluxes, from relatively bright (10$^{-13}$ erg cm$^{-2}$ s$^{-1}$, e.g. Della Ceca et al. 2004 and references therein), to medium (10$^{-13}$ erg cm$^{-2}$ s$^{-1}$-10$^{-14}$ erg cm$^{-2}$ s$^{-1}$, e.g. Barcons et al. 2007 and references therein) and deep (10$^{-14}$-10$^{-16}$  erg cm$^{-2}$ s$^{-1}$, Brandt \\& Hasinger 2005; Worsley et al. 2005 and references therein) fluxes. At the energies ($\\sim$0.5-10 keV) covered by the instruments on board  these two telescopes, Active Galactic Nuclei (AGN) can be efficiently selected and studied even when affected by large levels of absorption (up to N$_H\\sim$10$^{24}$ cm$^{-2}$, corresponding to an optical absorption of A$_V\\sim$500 mag). This important characteristic, combined with the good/excellent spatial and energy resolutions of the detectors, makes the ongoing surveys a fundamental tool for AGN studies. At the same time, these new surveys represent an observational challenge at wavelengths different from the X-ray ones: multiwavelength follow-ups of X-ray sources, particularly in the optical domain, are decisive to derive the distance and to understand the properties of the selected objects but they also require large fractions of dedicated observing time at different telescopes. Probably one of the most challenging and time-consuming efforts is the optical spectroscopic follow-up of the selected sources.  One of the  primary goals of all these hard X-ray surveys is to explore the population of absorbed AGN and, to this end, an optical classification that can reliably separate between optically absorbed and non-absorbed objects is always required. Two important limits, however, affect the spectroscopic follow-ups of deep and, in part, medium surveys: first, the optical counterparts are often too faint to be spectroscopically observed even at the largest optical telescopes currently available; second, even when a spectrum can be obtained, its quality is not always good enough to provide the critical pieces of information that are required to assess a reliable optical classification. These two problems often  limit the final scientific results that are based on the optical classification of medium/deep surveys.  On the contrary, bright surveys offer the important possibility of obtaining a reliable optical classification for virtually all (with some exceptions, as discussed in the next sections) the selected sources. The disadvantage of dealing with shallow/wide-angle samples is that the techniques to observe efficiently many sources at once, like Multi-objects or fibers-based methods,  cannot be applied for the optical follow-up, given the low space-density of sources at bright X-ray fluxes. The only suitable method, the ``standard'' long-slit technique, requires many independent observing nights to achieve the completion of the optical follow-up.  In this paper we present and discuss in detail the optical classification process of the {\\it XMM-Newton} Bright Serendipitous Survey (XBS, Della Ceca et al. 2004), which currently represents the widest (in terms of sky coverage) among the existing {\\it XMM-Newton} or {\\it Chandra} surveys for which  a spectroscopic follow-up has been almost completed. The aim of the paper, in particular, is to provide not only a generic classification of the sources and their redshift but also a quantification, in the limits of the available data, of the corresponding threshold in terms of level of optical absorption.  The paper is organized as follows: in Section~2 we describe the XBS survey, in Section~3 we describe the process of identification of the optical counterpart, in Section~4 and Section~5 we respectively  summarize our own spectroscopic campaigns carried out to collect the data as well as the data obtained from the literature. In Section~6 we shortly discuss the data reduction and analysis of the optical spectra and in Section~7 we give the details on the classification criteria adopted for the sources in the XBS survey. In Section~8 we propose two diagnostic plots that can be used to easily classify the sources into type~1 and type~2 AGN. The resulting catalogue is presented in Section~9 while in Section~10 we briefly discuss the optical breakdown and the redshift distribution of the sources. The conclusions are finally summarized in Section~11 .  Throughout this paper H$_0$=65 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda}$=0.7 and $\\Omega_M$ = 0.3 are assumed. ",
        "conclusions": "We have presented the details of the identification work of the sources in the XBS survey, which is composed by two complete flux limited samples, the BSS and the HBSS sample, selected in the 0.5-4.5 keV and 4.5-7.5 keV band respectively. We have secured a redshift and a spectroscopic classification for 348 (including data from the literature) out of 400 sources,  corresponding to 87\\% of the total list of sources and to 87\\% and 97\\% considering the BSS and HBSS samples separately.  The  results of the identification work can be summarized as follows:  \\begin{itemize}  \\item We have quantified the criteria used to distinguish optically absorbed AGN (i.e. type~2) from optically non-absorbed (or moderately absorbed) AGN (type~1) and we have shown that the adopted dividing line between the two classes of sources corresponds to an optical extinction of A$_V\\sim$2 mag, which translates into an expected column density of N$_H\\sim$4$\\times$10$^{21}$ cm$^{-2}$, assuming a Galactic A$_V$/N$_H$ relationship.  \\item About 10\\% of the extragalactic sources (35 objects in total) show an optical spectrum which is highly contaminated by the star-light from the host galaxy. These sources have been studied in detail in a companion paper (Caccianiga et al. 2007). Using the X-ray data we have found an elusive AGN in 33 of these objects and we have classified them into type~1 and type~2 AGN according to the value of N$_H$ measured from the X-ray spectrum. To this end, we have used a N$_H$=4$\\times$10$^{21}$ cm$^{-2}$ dividing value  which matches (assuming the standard Galactic A$_V$/N$_H$ relation) the value of A$_V$ (=2 mag) adopted with the optical classification.  \\item We have then proposed two simple diagnostic diagrams. The first one, based on the 4000\\AA\\ break and the [OIII]$\\lambda$5007\\AA\\ equivalent width, can reliably distinguish between type~1 and type~2 AGN if the host galaxy does not dominate the optical spectrum. The second uses the H$\\alpha$ and [OIII]$\\lambda$5007\\AA\\ line equivalent widths to classify into type~1 and type~2 the elusive AGN sources in which a possibly broad H$\\alpha$ emission line is detected.  \\item We find that the AGN represent the most numerous population  at the flux limit of the XBS survey ($\\sim$10$^{-13}$ erg cm$^{-2}$ s$^{-1}$) constituting 80\\% of the  XBS sources selected in the 0.5-4.5 keV energy band and 95\\% of the ``hard'' (4.5-7.5 keV) selected objects. Galactic sources populate significantly the 0.5-4.5 keV sample (17\\%) and only marginally (3\\%) the 4.5-7.5 keV sample. The remaining sources in both samples are clusters/groups of galaxies and normal galaxies (i.e. probably not powered by an AGN).  \\item As expected, the percentage of type~2 AGN dramatically increases going from the 0.5-4.5 keV sample (f=N$_{AGN 2}$/N$_{AGN}$=7\\%) to the 4.5-7.5 keV sample (f=32\\%). A detailed analysis on the intrinsic (i.e. taking into account the selection effects) relative fraction of type~1 and type~2 AGN will be be presented in a forthcoming paper (Della Ceca et al. 2007, in prep.).  \\end{itemize}"
    },
    "0710/0710.2175_arXiv.txt": {
        "abstract": "Voids are a dominant feature of the low-redshift galaxy distribution.  Several recent surveys have found evidence for the existence of large-scale structure at high redshifts as well. We present analytic estimates of galaxy void sizes at redshifts $z \\sim  5 - 10$ using the excursion set formalism.  We find that recent narrow-band surveys at $z \\sim 5 - 6.5$ should find voids with characteristic scales of roughly $20$ comoving $\\Mpc$ and maximum diameters approaching $40 \\Mpc$.  This is consistent with existing surveys, but a precise comparison is difficult because of the relatively small volumes probed so far.  At $z \\sim 7 -10$, we expect characteristic void scales of $\\sim 14 - 20$ comoving $\\Mpc$ assuming that all galaxies within dark matter haloes more massive than $10^{10} \\Msol$ are observable.  We find that these characteristic scales are similar to the sizes of empty regions resulting from purely random fluctuations in the galaxy counts.  As a result, true large-scale structure will be difficult to observe at $z \\sim 7 -10$, unless galaxies in haloes with masses $\\la 10^9 \\Msol$ are visible. Galaxy surveys must be deep and only the largest voids will provide meaningful information.  Our model provides a convenient picture for estimating the ``worst-case'' effects of cosmic variance on high-redshift galaxy surveys with limited volumes. ",
        "introduction": "The complex network of filaments and voids observed in the present-day Universe is believed to have formed from an initially homogeneous distribution of matter.  In hierarchical models of structure formation, tiny perturbations seeded by the inflationary epoch grew through gravitational instability, collapsing first on smaller scales to form haloes. The subsequent merging and clustering of smaller haloes resulted in the formation of highly structured large-scale systems.  Perhaps the most striking characteristic of the Universe today is the prevalence of large and nearly spherical voids in the galaxy distribution.  The scales of these voids can be enormous. Indeed, \\citet{HandV2004} report characteristic radii of $R \\sim 15h^{-1}~\\Mpc$ with maximum scales approaching $R \\sim 25 h^{-1} \\Mpc$ in the 2dF Galaxy Redshift Survey.  The characteristics of voids and the galaxies that populate them have been the subject of numerous theoretical and observational studies throughout the years \\citep{gregory78, kirshner81, delapparent86, vogeley94, HandV2004, Conroy2005}.  To date, these studies have mostly focused on low redshifts.  However, it is clear that voids should appear at higher redshifts also.  The DEEP2 survey indicates that voids exist at redshifts of $z \\sim 1$ \\citep{Conroy2005}.  Surprisingly, a handful of recent Lyman $\\alpha$ emitter (LAE) surveys have found hints of large-scale structure at redshifts around $z \\sim 5$ \\citep{Shimasaku2003, Shimasaku2006, Hu2004, Ouchi2005}.  The modelling of voids poses an interesting theoretical problem. There have been numerous studies utilising $N$-body simulations \\citep{MandW2002,Benson2003,Gottlober2003,Colberg2005}.  While these simulations are invaluable tools for understanding the details of void dynamics, they are computationally expensive due to the large volumes and high dynamic range required to include a representative sample of voids while also resolving the much smaller galaxies that define them.  Analytic methods provide a useful alternative.  Perhaps the most promising analytic model of void abundances is the excursion set approach taken by \\citet{SandV2004}.  They argue that voids actually provide deeper insight into large-scale structure than halo formation itself. Their assertion is based on the fact that underdense regions generally tend to evolve toward a spherical geometry, making the idealisation of spherical expansion more reasonable.  In contrast, gravitationally bound objects typically have geometries that are far from spherical.  Approximating gravitational collapse with the spherical model may be highly inaccurate, which partially explains the discrepancies between the \\citet{PandS1974} halo mass function and simulations \\citep{SandT1999,Sheth2001,Jenkins2001}.  A key disadvantage to the approach of \\citet{SandV2004} is the difficulty in relating their definition of voids to observational studies. As we will discuss in \\S \\ref{VoidDef}, \\citet{SandV2004} use the dark matter underdensity to define voids.  \\citet{FandP2006} extend their model to define voids in terms of the local \\emph{galaxy} underdensity. They predict characteristic void sizes of $R \\sim 10 h^{-1} \\Mpc$ at the present day -- nearly as large as observed voids.  In what follows, we present analytic estimates of void size distributions at redshifts between $z \\sim 5 - 10$.  Our aim is to provide a convenient basis of comparison for current and future high-redshift observations -- presumably, though not limited to, LAE surveys.  As such, we consider the effects of statistical fluctuations in the galaxy counts and the abundance of Ly $\\alpha$ emitting galaxies on void observations.  LAE surveys have become an invaluable tool in cosmological studies.  In addition to building larger samples at $z \\sim 5$, observers have pushed the threshold to redshifts as high as $z \\sim 7 - 10$ \\citep{WandC2005,Cuby2007,Stark2007,Ota2007}.  Surveys at these redshifts could potentially reveal important information on the epoch of reionization. Indeed, the observed clustering properties of LAEs could someday be a powerful probe of the epoch \\citep{furl04-lya, Furlanetto2006,McQuinn2006,McQuinn2007, mesinger07}. Regions of ionized hydrogen grow quickly around clustered galaxies as reionization progresses.  When these regions are large enough, Ly $\\alpha$ photons are sufficiently redshifted before they reach neutral hydrogen gas, allowing them to avoid absorption in the intergalactic medium (IGM).  Sources within overdense regions are therefore more likely to be observed relative to void galaxies, resulting in a large-scale modulation of the number density.  One method to quantify such clustering is with void statistics, as first attempted by \\citet{McQuinn2007}.  This provides comparable power to correlation function measurements of the galaxies.  However, taking full advantage of this technique requires a deeper understanding of voids in the underlying galaxy distribution; our calculations aim to provide such a baseline model.  The remainder of this paper is organized in the following manner.  In \\S \\ref{ExcursionSet}, we briefly  review the basic principles of the excursion set formalism.  In section \\ref{VoidDef}, we present the \\citet{FandP2006} definition of voids in terms of the local galaxy underdensity.  Section \\ref{VoidDistributions} contains the main results of this paper: void size distributions at $z = 4.86 - 10$.  In \\S \\ref{StochasticVoids}, we estimate the typical sizes of voids that result from random fluctuations in the galaxy distribution and develop an alternative definition of voids.    In \\S \\ref{VisibleFrac}, we explore the assumption that only a certain fraction of galaxies are actually visible in LAE surveys.  Section \\ref{Observations} contains a rough comparison of our calculations to high redshift Ly $\\alpha$ surveys.  Finally, we offer concluding remarks in \\S \\ref{Discussion}.  In what follows, we assume a cosmology with parameters $\\Omega_m = 0.24,~\\Omega_{\\Lambda} = 0.76,~\\Omega_b = 0.042,~H = 100h~\\mathrm{km~s}^{-1}~\\Mpc^{-1}$ (with $h = 0.73$), $n = 0.96$, and $\\sigma_8 = 0.8$, consistent with the latest measurements \\citep{Spergel2007}. All distances are reported in comoving units. ",
        "conclusions": "\\label{Discussion}  We have calculated void size distributions at $z = 4.86$--$10$ using the excursion set model developed by \\citet{SandV2004} and \\citet{FandP2006}.  The latter found characteristic void radii of $R \\approx  7-14 \\Mpc$ at $z = 0$.  For the observational sensitivities assumed in this paper, we obtained characteristic void radii that are very similar:  $R \\approx 7-10 \\Mpc$ for redshifts between $z = 4.86$ and $z = 10$.  These results are virtually independent of the void-crushing barrier (for any reasonable choice).  We have shown that characteristic void scales actually increase with redshift for a fixed halo mass threshold due to a decreased number density and increased bias with respect to the underlying matter density.  Following recent studies on the abundances of low-redshift LAEs, we explored the possibility that only a fraction $f_{vis}$ of galaxies are sampled in LAE surveys.  This has only a small effect on the void size distribution but increases the contamination of void samples by stochastic fluctuations.  In section \\ref{StochasticVoids}, we have explored stochastic fluctuations in the galaxy distribution.  These fluctuations, although inherently different from the \"real\" voids we model in this paper, will result in large empty regions in the sky. Stochastic voids can therefore contaminate real void samples and lead to erroneous conclusions on the formation of large-scale structure.  We have estimated the typical scale of these regions to be slightly smaller than the characteristic scale of true voids at $z \\sim 5$. At $z \\sim 10$, the situation depends on the particular choice of $m_{min}$.  For $m_{min} \\sim 10^{10} \\Msol$, stochastic voids are typically the same scale as real voids.  The increased importance of stochastic fluctuations will make the identification of large-scale structure at this redshift difficult.  Attempts to do so must observe halos near the minimum mass to form stars, $\\sim 10^8$--$10^9 \\Msol$, in order for true voids to dominate the observed distribution.  We found that a large fraction of real voids in our fiducial model contain no visible galaxies, adding to the difficulties in differentiating them from stochastic fluctuations.  We have presented a modified definition of voids that incorporates both stochastic and real voids and so is easier to compare to the limited observational samples thus far available.  In our new approach, we defined voids in terms of the probability for a region to be empty.  We found that the modified void distributions are more sharply peaked and have characteristic scales that are comparable to the fiducial model.  We have also attempted to visually compare our results to the most recent narrow-band filter surveys at $z = 5.7$.  While we found no inconsistencies, it is difficult to draw any decisive conclusions because of small-number statistics, projection effects, and lower-redshift contaminants.  Obviously, a more systematic approach is required.  Future surveys promise to provide better statistics and increased sample volumes for studies on high-redshift voids.  In the context of next generation surveys for high-redshift galaxies, our model is useful for gauging the impact of cosmic variance.  Consider a fictitious survey at $z = 10$ with a detection threshold of $m_{min} = 10^{10} \\Msol$.  Figure \\ref{PoissonPlot} illustrates that stochastic voids with $R \\sim 10 - 20 \\Mpc$ will dominate the sky distribution.  Therefore, one must either search for voids with $R > 20 \\Mpc$ or search deeper for significantly smaller sources.  The latter may be possible if the sources observed by \\citet{Stark2007} are indeed at $z \\sim 9$, in which case they imply that halos near $\\la 10^9 \\Msol$ are visible \\citep{mesinger07}.  However, high-redshift galaxies are so highly biased that even with deep observations, a substantial fraction of the Universe is filled with empty or nearly-empty regions.  For example, at $z=10$ and $m_{min} = 10^{10} \\Msol$, $\\sim 37\\%$ of space is filled by regions that are at least 80\\% underdense in galaxies and at least 20 Mpc across -- or fully 7 arcmin.  With the small fields of view available to near-infrared detectors, this suggests that either many independent fields must be observed or a large contiguous volume surveyed to be guaranteed of detecting a reasonable number of sources.  Finally, we have neglected reionization and its effect on the appearance of large-scale structure.  Regions of neutral hydrogen are expected to modulate the LAE density on large scales and accentuate the appearance of structure \\citep{furl04-lya, Furlanetto2006,McQuinn2006,McQuinn2007, mesinger07}. Although the precise time frame is currently unknown, quasar observations and cosmic microwave background measurements have provided some evidence that reionization occurred between $z \\sim 6 - 10$ (e.g, \\citealt{Fan2006,Page2007,MandH2004,MandH2007}).  Interestingly, \\citet{Kashikawa2006} found a significant high-luminosity suppression in the LAE luminosity function between $z = 5.7$ and $z = 6.5$. Whether or not reioinization is responsible for this effect is currently unclear \\citep[no such suppression was observed by][]{dawson07}.  Because IGM absorption modulates the LAE density on large scales, we would expect reionization to have a substantial effect on the observed void sizes in such narrow-band surveys (it should not affect galaxies identified through broadband effects).  Of course, the plots in Figure \\ref{Z7to10} provide analytic estimates only of the \\emph{intrinsic} void size distributions.  They provide a basis for comparison with high-redshift surveys in order to determine whether the observed features are easily attributable to the large-scale clustering alone.  It therefore helps illuminate efforts to use voids to constrain the IGM properties during reionization, as first attempted by \\citet{McQuinn2007}."
    },
    "0710/0710.3345_arXiv.txt": {
        "abstract": "Determining the final spin of a black-hole (BH) binary is a question of key importance in astrophysics% . Modelling this quantity in general is made difficult by the fact that it depends on the 7-dimensional space of parameters characterizing the two initial black holes. However, in special cases, when symmetries can be exploited, the description can become % simpler. For black-hole binaries with unequal masses but with equal spins which are aligned with the orbital angular momentum, we show that the use of recent simulations and basic but exact constraints derived from the extreme mass-ratio limit allow to model this quantity with a simple analytic expression. Despite the simple dependence, the expression models very accurately all of the available estimates, with errors of a couple of percent at most. We also discuss how to use the fit % to predict when a Schwarzschild BH is produced by the merger of two spinning BHs, when the total angular momentum of the spacetime ``flips'' sign, or under what conditions the final BH is ``spun-up'' by the merger. % Finally, suggest an extension of the fit % to include unequal-spin binaries, thus potentially providing a \\textit{complete} description of the final spin from the coalescence of generic black-hole binaries with spins aligned to the orbital angular momentum. ",
        "introduction": "\\label{intro}  The determination of the final spin of a BH binary is a question of key importance in astrophysics. Modelling this in general is made difficult by the fact that it depends on the 7-dimensional space of parameters characterizing the two initial BHs. However, in special cases, when symmetries can be exploited, the description can be much simpler.  Several recent studies have shed light on the remnant of the merger process. Using conservation principles, Hughes and Blandford~\\citep{Hughes_Blandford:2003} argued that mergers rarely lead to rapidly rotating objects. \\citet{Gonzalez:2006md} numerically evolved a sequence of non-spinning unequal-mass BHs, arriving at detailed estimates of the radiated energy and angular momentum.  In a series of papers~\\citep{Koppitz:2007ev, Pollney:2007ss, Rezzolla_etal:2007a} we have studied the parameter space of mergers of equal-mass BH binaries whose spins are aligned with the orbital angular momentum but otherwise arbitrary.  The findings agree well with independent numerical evolutions~\\citep{Campanelli_etal:2007,Herrmann:2007ac}, as well as more recent studies of models with initial spins up to $J/M^2=0.8$~\\citep{Marronetti_etal:2007}. An important result of these studies has been the determination of simple (quadratic) fitting formulas for the recoil velocity and spin of the merger remnant as a function of the initial BH parameters~\\citep{Rezzolla_etal:2007a}.  A number of analytical approaches have been developed over the years to determine the final spin of a binary coalescence~\\citep{ Damour:2001tu, Buonanno_Damour:2000, Buonanno_etal:2006, Damour_Nagar:2007a, Boyle:2007sz}. Very recently, an interesting method, inspired by the dynamics of a test particle around a Kerr BH, has been proposed for generic binaries~(\\citet{Buonanno_etal:2007b}, BKL hereafter). The approach assumes that the angular momentum of the final BH is the sum of the individual spins and of the orbital angular momentum of a test particle on the last-stable orbit of a Kerr BH with the same spin parameter as that of the final BH.  Here, we combine the data obtained in recent simulations to provide a phenomenological but analytic estimate for the final spin in a binary BH system with arbitrary mass ratio and spin ratio, but in which the spins are constrained to be parallel to the orbital angular momentum. Our numerical simulations have been carried out using the CCATIE code~\\citep{Pollney:2007ss}. In addition to the data presented in \\citet{Rezzolla_etal:2007a}, we add three simulations of equal-mass, high-spin binaries and three simulations of unequal-mass, spinning binaries (see Table~\\ref{tableone}).  Other data is taken from % unequal-mass, nonspinning binaries~\\citep{Gonzalez:2006md,Berti_etal:2007,Buonanno_etal:2007a}, and of equal-mass, spinning binaries~\\citep{Rezzolla_etal:2007a,Marronetti_etal:2007}; all of the AEI data is summarized in Table~\\ref{tableone}. To avoid the possible contamination from the errors associated with high-spin binaries reported by~\\citet{Marronetti_etal:2007}, we have not considered binaries with initial spin $|J/M^2| \\geq 0.75$ reported in the literature~\\citep{Campanelli_etal:2007,Marronetti_etal:2007}. We have, however, considered estimates of high-spin binaries (\\textit{cf.}, Table~\\ref{tableone}), for which we know the spins remain essentially constant prior to merger, with changes less than $0.5\\%$~\\citep{Pollney:2007ss}, and that are very well captured by the fit. ",
        "conclusions": "Modelling the final spin in a generic binary BH merger is not trivial given the large space of parameters on which this quantity depends. We have shown that the results of recent simulations combined with basic but exact considerations derived from the EMRL allow us to model this quantity with a simple analytic expression in the case of BH binaries having unequal masses and unequal spins which are aligned with the orbital angular momentum. When compared with all other estimates coming either from numerical calculations or from approximation techniques, the estimates of the 2D fit show differences which are of few percent at most.     \\bigskip \\noindent We % thank A. Buonanno, T. Damour, S. A. Hughes, L. Lehner, A. Nagar, and B. S. Sathyaprakash for discussions. We are grateful to D. Merritt for pointing out an error in the interpretation of our results. The computations were performed on the supercomputers at AEI, LITE, LSU, LONI and NCSA. Support comes also through the DFG grant SFB/TR% ~7."
    },
    "0710/0710.1340_arXiv.txt": {
        "abstract": "Optically thin two-temperature accretion flows may be thermally and viscously stable, but acoustically unstable. Here we propose that the O-mode instability of a cooling-dominated optically thin two-temperature inner disk may explain the 23-day quasi-periodic oscillation (QPO) period observed in the TeV and X-ray light curves of Mkn~501 during its 1997 high state. In our model the relativistic jet electrons Compton upscatter the disk soft X-ray photons to TeV energies, so that the instability-driven X-ray periodicity will lead to a corresponding quasi-periodicity in the TeV light curve and produce correlated variability. We analyse the dependence of the instability-driven quasi-periodicity on the mass (M) of the central black hole, the accretion rate ($\\rm{\\dot{M}}$) and the viscous parameter ($\\alpha$) of the inner disk. We show that in the case of Mkn~501 the first two parameters are constrained by various observational results, so that for the instability occurring within a two-temperature disk where $\\alpha=0.05-1.0$, the quasi-period is expected to lie within the range of 8 to 100 days, as indeed the case. In particular, for the observed 23-day QPO period our model implies a viscosity coefficient $\\alpha \\leq 0.28$, a sub-Eddington accretion rate $\\dot{M} \\simeq 0.02\\,\\dot{M}_{\\rm Edd}$ and a transition radius to the outer standard disk of $r_0 \\sim 60 \\,r_g$, and predicts a period variation $\\delta P/P \\sim 0.23$ due to the motion of the instability region. ",
        "introduction": "Strong variability is one of the common observational properties of TeV emitting Active Galactic Nuclei (AGNs) \\cite{cat99}. In many cases, the highly variable high-energy { gamma-rays and the X-rays appear to be correlated with no time delays evident on day-scales}, suggesting that the $\\gamma$-rays { may} result from inverse Compton upscattering of lower energy photons. { The first TeV blazar, for example, that was observed simultaneously in multiple bands from radio to TeV gamma-rays is Mrk~421. The first campaign, conducted in 1994 \\cite{mac95} on Mrk~421, revealed correlated variability between the keV X-ray and TeV gamma-ray emission. The gamma-ray flux varied by an order of magnitude on a timescale of 2 days and the X-ray flux was observed to double in 12 hr. On the other hand, the high-energy gamma-ray flux observed by EGRET, as well as the radio and UV fluxes, showed less variability than the keV or TeV bands. Another multiwavelength campaign on Mrk~421 performed in 1995 revealed another coincident keV/TeV flare \\cite{buk96,taka96}. The UV and optical bands also showed correlation during the flares. With the detection of TeV emission from Mrk~501 \\cite{qui96}, several multiband campaigns were organized on Mrk~501 to verify whether such a behavior is a general property of TeV-emitting blazars or whether it is unique to Mrk 421.}   The gamma-ray blazar Mkn~501, detected as a strong TeV emitter in 1995 \\cite{qui96}, is one of the closest ($z = 0.0337$) and brightest BL Lacertae objects. Historically (i.e., prior to 1997), its spectral energy distribution (SED) $\\nu\\,F_{\\rm \\nu}$ resembles that of X-ray selected BL Lac objects, having a peak in the extreme UV--soft X-ray energy band \\cite{kat99,sam96}. { Earlier} optical observations of Mkn~501 have shown variations of up to $1.^{m}3$ and polarized emission up to $P_{\\rm opt} \\simeq 7 \\%$ \\cite{fan99}.  During its active state in 1997, where Mkn~501 was monitored in the 2-10 keV X-ray band by the all sky monitor(ASM) on board RXTE and in the TeV energy band by several Cherenkov telescope groups \\cite{aha99a,cat97,dja99,kraw00,pro97,rxte99}, both X-rays and TeV gamma-rays increased by more than 1 order of magnitude from quiescent level \\cite{cat97,pia98}. Analysis of the X-ray and TeV data showed, that the variations were strongly correlated between both bands, { yet only marginally correlated with the optical UV band} \\cite{cat97,dja99,kraw00,pet00}. { While the synchrotron emission peaked below 0.1 keV in the quiescent state, in 1997 it peaked at ~100 keV. This is the largest shift ever observed for a blazar \\cite{pia98}. Earlier investigations of the 1997 April flare in Mkn~501 \\cite{pia98}, showed that its origin may be related to a variation of $\\gamma_{max}$ together with an increased luminosity and a flattening of the injected electron distribution.}   During the 1997 high state, the X-ray and TeV light curves displayed a quasi-periodic signature \\cite{hay98,kra99,pro97}. A detailed periodicity analysis, based on the TeV measurements obtained with all Cherenkov Telescopes of the HEGRA-Collaboration and the X-ray data with RXTE, was performed by Kranich et al. \\cite{kra99,kra01} and more recently also by Osone \\cite{oso06} including Utah TA data. The results indeed strongly suggest the existence of a 23 day periodicity, with a combined probability of $p \\simeq 3 \\times 10^{-4}$ in both the TeV and the X-ray light curves covering the same observational period \\cite{kra99,kra01}. \\footnote{{ Note that no QPOs have been seen by MAGIC during 1998-2000, when the source was not very active \\cite{alb07}, suggesting that the QPOs only occur during a very active stage.}} { Rieger and Mannheim ~(2000) have shown that the origin of these QPOs may be related to the presence of a binary black hole in the center of Mkn~501. While this may well be possible, we explore here an alternative scenario, where the observed QPOs are related to accretion disk instabilities.}  In a seminal paper, Shapiro et al.~(1976) (SLE) have pointed out that a hot (Compton cooling-dominated) optically thin two-temperature accretion disk may be present in the inner region of the standard optically thick disk \\cite{sha73}, whenever the SLE inequality $\\alpha^{1/4} \\dot{M}_{*}^{2}M_{*}^{-7/4} \\geq 0.6$ is satisfied, where $\\dot{M}_{*} = \\dot{M}/(10^{17} \\rm{g} \\rm{s}^{-1})$, $M_{*}=M/(3 M_{\\odot})$, $\\dot{M}$ is the accretion rate of the disk, $M$ the mass of the central black hole, and $\\alpha$ the viscosity parameter, constrained by the model to lie within the range $ 0.05 < \\alpha < 1$. Later work \\cite{pir78,pri76} has shown that the SLE configuration might be thermally unstable (although less than the standard disk) if the simple standard viscosity prescription is employed. However, relatively small changes in the viscosity law can already ensure a stable configuration \\cite{pir78}, in particular, if stabilizing effects of a disk wind are fully taken into account. A kind of SLE two-temperature disk structure may thus well exist in the inner region of BL Lac type objects \\cite{wan91,wanu91,cao03} and provide a possible explanation for the X-ray and TeV variability phenomenon in Mkn~501. For, firstly, Compton processes in a inner two-temperature disk with electron temperature $T_e$ of about $10^{9}$ K (and ion temperature one or two orders of magnitude higher) can produce (steep) X-ray power-law spectra, in contrast to the standard disk model that can only produce emission up to the optical-UV band \\cite{sha76,wan91}. Secondly, analysis of the linear stability of an optically thin two-temperature disk around a compact object shows that the disk is subject to a radial pulsational instability (inertial-acoustic mode instability) \\cite{wu97}. This possible mode of pulsational overstability, in which radial disk oscillations with local Keplerian frequencies become unstable against axisymmetric perturbations, occurs if the viscosity coefficient increases sufficiently upon compression \\cite{kat78}. In this case, thermal energy generation due to viscous dissipation increases during compression, leading to amplification of the oscillations analogous to the role played by nuclear energy generation in stellar pulsations. As we demonstrate below, the occurrence of such a type of disk oscillation may well account for some quasi-periodic time variability in AGNs in a way similar to Galactic black hole candidates \\cite{blu84,hon92,man96,yan97}. ",
        "conclusions": "The X-ray to TeV spectra of TeV blazars have been often interpreted within a one-zone SSC model, e.g. \\cite{blo96,der97,mas97,tav98,tav01}. However, if a two-temperature disk structure is present in some of these sources, the real situation may be more complex as the X-ray radiation from a two-temperature disk, for example, may represent an additional, non-negligible source of seed photons for the inverse Compton scattering to TeV energies, in particular during active source stages, likely to be associated with changes in accretion history. Hence our model assumes, that a pulsational instability occurring within a two-temperature disk leads to observable, periodic variations of its X-ray radiation field. Part of this periodically modulated X-ray emission will enter the jet and (in addition to direct synchrotron photons) serves as seed photons for Compton-upscattering to TeV energies { similar as in \\cite{der92,der93}}. Our model thus takes it that both, a direct synchrotron self Compton and an external Compton contribution are relevant for modelling the SED of Mkn~501, the seed photons for external Compton-upscattering consisting of both, the infrared-optical seed photons from the (quasi-steady) SS disk component and the variable X-ray photons from the two-temperature disk component. A related, but more simpler scenario, assuming the seed photons for inverse Compton scattering to be provided by a (direct) synchrotron plus an quasi-steady flux component comparable to the observed infrared-optical flux, has been proposed by Pian et al.~(1998) in order to account for the different degrees of SED variations of Mkn~501 at X-ray and sub-X-ray energies during the April 1997 outburst (see also \\cite{ghi98,kat99}). Detailed analysis indeed suggests that one-component SSC models cannot fit both the April 1997 SEDs and the lightcurves from X-ray to TeV, and that (at least) an additional, moderately variable low energy component contributing in the energy range between 3-25 keV is required \\cite{kraw02,mas04}. A similar conclusion seems to hold for the strong X-ray outburst observed in July 1997 \\cite{lam98}. Interestingly, observations in 1998 also provide evidence for an additional component in the optical regime, possible associated with the SS disk component in our hybrid disk model { (see also \\cite{kat01} for an alternative interpretation)}: As shown by Massaro et al.~\\cite{mas04} optical to X-ray data taken in June 1998 indicate that the optical spectrum is steep and does not match the low energy extrapolation of the X-ray spectrum, hence suggesting the presence of different emission components in the optical and in the X-ray regime as naturally expected in our model.  Monitors in both the X-rays and the TeV emission show evidence for a 23-day periodicity during the 1997 high state. As demonstrated above the 23-day period in the X-ray light curve may be caused by a pulsational instability in two-temperature accretion disk, and via the inverse Compton process result in the same periodicity in the TeV light-curve. A pulsational instability occurring in a two-temperature disk with transition radius $r_0 \\sim (48-65)\\, r_{g}$  will result in a recurrence timescale of 8 to 100 days.  Based on the observed TeV variability we have employed a characteristic black hole mass of $9 \\times 10^7 M_{\\odot}$ in our calculations. We note that quite different central mass estimates for Mkn~501 have been claimed in the literature, ranging from several times $10^7$ (mainly based on high energy emission properties) up to $10^9\\,M_{\\odot}$ (based on host galaxy observations), see e.g. \\cite{cao02,dej99,fal02,fan05,rie03}. However, as shown by Rieger \\& Mannheim~\\cite{rie03} uncertainties associated with host galaxy observations may easily lead to an overestimate of the central black hole mass in Mkn~501 by a factor of three and thus reduce the implied central mass to $\\simeq (2-3)\\times 10^8\\,M_{\\odot}$, { a value in fact recently confirmed by an independent analysis of central mass constraints derived from host galaxy observations \\cite{woo05}.} Moreover, as argued by the same authors some of the apparent disagreement in central mass estimates may possibly be resolved if a binary black hole system exists in the center of Mkn~501, see also \\cite{rie00,rie03,vil99}, similar as in the case of OJ~287 \\cite{sil88}. For example, if Mkn~501 harbours a binary system with a more massive primary black hole of $\\lppr 10^{9} M_{\\odot}$ and a less massive (jet-emitting) secondary black hole of $\\sim 10^{8}\\,M_{\\odot}$, the mass ratio $\\rho=m/M$ would be of order 0.1, which may compare well with the result $\\rho<0.25$ estimated for OJ~287 \\cite{liu02}. While our characteristic black hole mass employed falls well within the above noted range, we note that the SLE pulsational instability model may still work successfully, if a higher black hole mass is used. For example, if one adopts $M = 3 \\times 10^{8} M_{\\odot}$ and $P=23$ days, one obtains $\\alpha \\leq 0.07$, $r_{0*}=46$ and $\\dot{M}\\simeq 0.008 \\dot{M}_{\\rm Edd}$.  Our analysis is based on a specific disk model (SLE) which is open to questions, in particular with respect to its possible stability properties. We note however, that a relatively small change in the usually employed viscosity description may already lead to a thermally stable configuration \\cite{pir78}. On the other hand, it may as well be possible that the SLE configuration represents a quasi-transient phenomenon associated with those changes in accretion history that probably initiate the high states. An alternative (inner) disk configuration of interest may be represented by an optically thin two-temperature ADAF solution \\cite{nar98,yi99}. Such a configuration can exist for accretion rates $\\dot{m}= \\dot{M}/\\dot{M}_{\\rm Edd}$ below a critical rate $\\dot{m}_{\\rm crit} \\simeq 0.3\\,\\alpha^2 \\simeq 0.019$, where the canonical ADAF value of $\\alpha =0.25$ has been employed \\cite{nar98}. ADAFs are generally less luminous than a standard disk, with the typical ADAF luminosity given by $L_A \\simeq 0.02\\,(\\dot{m}/\\alpha^2)\\,\\dot{M}\\,c^2$ \\cite{yi99}. Using the constraints above, the possible ADAF luminosity for Mkn~501 becomes $L_A \\leq 1.3 \\times 10^{43}$ erg/s, which is already about an order of magnitude smaller than required by the X-ray analysis. This suggests that -- at least during its high state -- an optically thin ADAF is not a viable option for Mkn~501.  Based on Eq. (7) we can estimate the variation rate of the period due to the motion of the instability \\begin{eqnarray} \\delta P/P = {\\frac{3}{2}} P{\\frac{1}{r}} {\\frac{\\partial{r}}{\\partial{t}}}\\,. \\end{eqnarray} Using $v_{r} = {\\frac{\\partial{r}}{\\partial{t}}}$, $\\dot{M} = 4 \\pi \\rho h r v_{r}$ and Eqs. (4) and (5), one finds \\begin{eqnarray} \\delta P/P  =0.44 ~\\alpha^{-7/6}~\\dot{M_{24}}^{5/6}~M_{7}^{-5/6}~r_{*}^{-1/4} \\zeta^{-1/6} \\end{eqnarray} For $\\alpha$ = 0.28 the relevant parameters result in $\\delta P/P= 0.23$. From the period analysis performed by Kranich et al.~\\cite{kra99} a $3\\sigma$ deviation in period corresponds to 6.67 days. If we take this deviation as the intrinsic variation on the periodicity (P), then a $\\delta P/P$ of ${\\frac{6.67}{23}}=0.29$ can be estimated from the results by Kranich et al.~\\cite{kra99}. As this estimate assumes that the deviation of the period is only affected by the motion of the instability, while it may in fact be caused by more than one effect, our theoretical $\\delta P/P$ should not be greater than the observational results. We conclude that the observed 23-day QPOs in Mkn~501 might be caused by the instability of a two-temperature accretion disk. The model presented here may thus offers an alternative explanation to the binary-driven helical jet model of Rieger \\& Mannheim~(2000). Comprehensive computational modelling of the pulsational instability in a two-temperature, cooling dominated disk will be essential to verify this in more detail.  Our model predicts that a period correlation in the X-ray and $\\gamma$-ray should always be present during an active source stage, while the period of the QPOs may vary as the instability region could change from one high state to the other."
    },
    "0710/0710.4086_arXiv.txt": {
        "abstract": "{ Based on the results of applying the extended ADC emission model to three Z-track sources: GX\\th 340+0, GX\\th 5-1 and Cyg\\th X-2, we propose an explanation of the Z-track sources in which the Normal and Horizontal Branches are dominated by the increasing radiation pressure of the neutron star. The emitted flux becomes several times super-Eddington at the Hard Apex and Horizontal Branch and we suggest that the inner accretion disk is disrupted by this and that part of the accretion flow is diverted vertically. This position on the Z-track is exactly the position where radio emission is detected showing the presence of jets. We thus propose that high radiation pressure is a necessary condition for the launching of jets. We also show that flaring must consist of unstable nuclear burning and that the mass accretion rate per unit emitting area of the neutron star $\\dot m$ at the onset of flaring agrees well with the critical theoretical value at which burning becomes unstable. ",
        "introduction": "% \\label{sect:intro}  The Z-track sources are the brightest group of Galactic low mass X-ray binaries (LMXB) containing a neutron star persistently emitting at the Eddington luminosity or several times this. The sources trace out a Z-shape in hardness-intensity (Hasinger et al. 1989) clearly showing that strong physical changes take place, probably at the inner disk and neutron star, but a convincing explanation of the Z-track phenomenon does not exist. The majority of LMXB are of the Atoll class which show somewhat different shapes in hardness-intensity which are also not understood, and neither is the relation between the two classes making our understanding of LMXB very incomplete.  Moreover, it is well-known that the Z-track sources are detected as radio emitters, but in one branch only, the horizontal branch. Not only is radio detected, but striking results from the VLA show the release of a massive radio condensation from the source Sco\\th X-1 (Fomalont et al. 2001). Because radio is detected essentially in one branch only, the sources offer the possibility of determining the conditions found in this branch distinguishing it from the other two branches, and so finding the conditions necessary for jet formation.    Possible ways of understanding the Z-track sources are by theoretical approaches, timing studies or spectral studies. A theoretical model for the Z-track sources was produced by Psaltis et al. (1995) based on a magnetosphere of the neutron star and the changing properties and geometry of this as the mass accretion rate changed. However, the model assumed that the Comptonized emission observed in the spectra (of all LMXB) originated in a small central region close to the neutron star, and this is inconsistent with our more recent measurements of Comptonizing region size (Church \\& Ba\\l uci\\'nska-Church 2004, below).  Extensive timing studies have been made to investigate QPO variations around the Z-track (e.g. van der Klis et al. 1987), but this has not revealed the nature of the Z-track. Previous spectral fitting has applied the Eastern model (Done et al. 2002, Agrawal \\& Sreekumar 2003; di Salvo et al. 2002) which assumes the X-ray emission consists of disk blackbody emission plus non-thermal emission from a small central Comptonizing region. However, our work over a period of 10 years with the dipping class of LMXB provides strong evidence that the source of Comptonized emssion, the ADC, is very extended, typically having a radial extent that is 15\\% of the accretion disk size, but increasing with source luminosity, and this is inconsistent with the Eastern model. As a result we have proposed the ``extended ADC'' emission model consisting of blackbody emission from the neutron star plus Comptonized emission from an extended ADC (Church Ba\\l uci\\'nska-Church 1995). Moreover, the pattern of parameter changes obtained by fitting the Eastern model to the Z-track sources is not very easy to interpret and does not immediately suggest a convincing physical explanation. Thus in the present work, we take the approach of applying the extended ADC model for the first time to the Z-track sources and we present the results of applying this model to the sources GX\\th 340+0, GX\\th 5-1 and Cygnus\\th X-2.  \\begin{figure*}[!ht]                                                           % \\begin{center} \\includegraphics[width=60mm,height=140mm,angle=270]{church_2007_01_fig1a}   % \\includegraphics[width=60mm,height=80mm,angle=270]{church_2007_01_fig1b} % \\caption{Top: Background-subtracted and deadtime-corrected PCA light curve of the 1997 September observation of GX\\th 340+0 with 64 s binning. Bottom: the corresponding variation of hardness ratio \\hbox{(7.3 -- 18.1 keV)/(4.1 -- 7.3 keV)} with intensity.} \\label{} \\end{center} \\end{figure*} ",
        "conclusions": "We show that the radiation pressure of the enitting part of the neutron star is very strong at the hard apex and horizontal branch of the Z-track in three sources, exactly correlating with the parts of the Z-track where radio emission is observed showing the presence of jets, and we suggest that strong radiation pressure is a necessary condition for jet formation."
    },
    "0710/0710.1947_arXiv.txt": {
        "abstract": "% {} {We study the variability of the Fe 6.4 KeV emission line from the Class I young stellar object Elias 29 in the $\\rho$ Oph cloud.} {We analysed the data from Elias 29 collected by \\xmm\\ during a nine-day, nearly continuous observation of the $\\rho$ Oph star-forming region (the Deep Rho-Oph X-ray Observation, named {{\\sc Droxo}}). The data were subdivided into six homogeneous time intervals, and the six resulting spectra were individually analysed} {We detect significant variability in the equivalent width of the Fe 6.4 keV emission line from Elias 29. The 6.4~keV line is absent during the first time interval of observation and appears at its maximum strength during the second time interval (90 ks after Elias 29 undergoes a strong flare). The X-ray thermal emission is unchanged between the two observation segments, while line variability is present at a 99.9\\% confidence level. Given the significant line variability in the absence of variations in the X-ray ionising continuum and the weakness of the photoionising continuum from the star's thermal X-ray emission, we suggest that the fluorescence may be induced by collisional ionisation from an (unseen) population of non-thermal electrons. We speculate on the possibility that the electrons are accelerated in a reconnection event of a magnetically confined accretion loop, connecting the young star to its circumstellar disk.} {} ",
        "introduction": "\\label{sec:intro}  The X-ray emission from young stellar objects (YSOs) at CCD-resolution is usually modelled as thermal emission from a hot plasma in coronal equilibrium, with higher characteristic temperatures than observed in older and less active stars.  An interesting deviation from a pure thermal X-ray spectrum is the presence of fluorescent emission from neutral (or weakly ionised) Fe as shown by the presence of the 6.4 keV line. This was first detected by \\citet{ikt01} in the X-ray emission of the YSO YLW16A in $\\rho$-Oph, during a large flare: in addition to the Fe\\,{\\sc xxv} complex at 6.7 keV, a 6.4 keV emission line was clearly visibe. Such fluorescence line is produced when energetic X-rays photoionise cold material close to the X-ray source, and it is therefore a useful diagnostic tool of the geometry of the X-ray emitting source and its surroundings.  Since 2001, detections of the Fe K fluorescent emission line at 6.4-keV in the spectra of YSOs have been reported by a number of authors. \\citet{tfg+05} has identified seven sources with an excess emission at 6.4 keV among 127 observations of YSOs within the COUP observation of Orion; \\citet{fms+05} report 6.4 keV fluorescent emission in Elias 29 in $\\rho$-Oph both during quiescent and flaring emission, unlike all other reported detection of Fe fluorescent emission in YSOs that were made during intense flaring; \\citet{gfm+07} have detected Fe 6.4 keV emission from a low-mass young star in Serpens, during an intense, long-duration flare. Recently, \\cite{sc2007} have reported intense Fe fluorescent emission in the spectrum of V 1486 Ori during a strong flare, when the plasma reached a temperature in excess of 10 keV.  The 6.4 keV fluorescent line has been detected in different classes of X-ray emitters: X-ray binaries, active galactic nuclei (AGNs), massive stars, supernova remnants, and the Sun itself during flares. In the case of the Sun, the fluorescing material is the solar photosphere, in the YSOs, however, indications are that the material in the circumstellar disk and its related accretion structures could be responsible for the fluorescence. The typical equivalent width of the 6.4 keV emission line in the studies mentioned above is of the order of 150~eV and is too large to be explained with fluorescent emission in the stellar photosphere or in diffuse circumstellar material (e.g. \\citealp{tfg+05}; \\citealp{fms+05}).  This scenario implies that the disk is ``bathed'' in high-energy X-rays emitted by the star, with significant astrophysical implications; for instance, X-rays, in addition to cosmic rays, would play an important role in photoionising the circumstellar material around young star and thus in coupling the gas to the ambient magnetic field (as suggested by e.g. \\citealp{gfm00}). \\citet{cbt+02} suggest that the ``hot'' component they observe in the disk, in the infrared, is heated by the stellar high-energy emission.  \\begin{figure*} \\begin{center} \\leavevmode \\epsfig{file=7899fg01.ps, width=17.0cm, bbllx=0,bblly=190,bburx=842,bbury=1000, angle=270, clip=}  \\caption{The light curve of Elias 29 over the nine days of the \\droxo\\ observation from PN ({\\em top}), MOS1 ({\\em bottom-left}), and MOS2 ({\\em bottom-right)}. The line with error bars gives the background-subtracted light curve of the source, while the thinner line without error bars gives the total counts (source plus background). The 6 time intervals that we selected for the spectral analysis, on the basis of the PN data, are indicated.}  \\label{fig:lc} \\end{center} \\end{figure*}  Observations of the time variability of the Fe 6.4 keV emission in YSOs would provide useful constraints on the geometry and sizes of the star-accretion disk system and of other circumstellar structures such as funnel flows, jets, or wind columns. We present here the results of a time-resolved spectral study of the X-ray emission of Elias 29, during $\\sim 9$ days of nearly continuous observation by \\xmm\\ in the context of the ultra-deep observation of $\\rho$-Oph, named \\droxo\\ (from Deep Rho-Oph X-ray observation -- \\citealp{psf+07}).  We investigated the presence of variations in its strong Fe 6.4 keV emission over the 9-day time scale covered by the observations.  This paper is structured as follows. After a summary, below, of the properties of Elias 29, the observations and data analysis are briefly presented in Sect.\\,\\ref{sec:obs}. Results are summarised in Sect.\\,\\ref{sec:res}; the simulations carried out to assess the reliability of the line detections and the significance of its variability are described in Sect.\\,\\ref{sec:sim}. The results are discussed in Sect.\\,\\ref{sec:disc}.  \\subsection{Elias 29 (GY214)}  Elias 29 (16:27:09.4, $-$24:37:18.9) with a bolometric luminosity $L = 26-27.5~L_{\\sun}$ (\\citealp{bak+01}; \\citealp{nts06}) is the most luminous Class I YSO in the $\\rho$-Oph cloud. \\cite{mhc98} used the luminosity in the Br$\\gamma$ line to determine the object's accretion luminosity at $L_{\\rm acc}$ = $15-18 L_{\\sun}$, which makes it the source with the highest accretion luminosity in their sample. More recently, \\citet{nts06} used the luminosity of the hydrogen recombination lines to derive an accretion luminosity of 28.8~$L_{\\sun}$.  Using millimeter interferometric observations, \\citet{bhc+02} resolved the emission from the disk and the envelope surrounding Elias 29, showing that the disk is in a relatively face-on orientation ($i < 60^{\\circ}$), which explains many of the remarkable observational features of this source, such as its flat spectral energy distribution, its brightness in the near-infrared, the extended components found in speckle interferometry observations, and its high-velocity molecular outflow. Their best-fitting disk model has an inner radius of 0.01 AU, outer radius of 500 AU, and a mass $M = 0.012 M_{\\sun}$.  Elias 29 was previously observed in X-rays with ASCA, \\chandra\\, and \\xmm.  In the \\chandra\\ observation (\\citealp{ikt01}), the source quiescent phase is characterised by a temperature of 4.3~keV and luminosity of $2.0 \\times 10^{30}$~\\es, fully consistent with the values derived from the subsequent \\xmm\\ observations by \\citet{ogm05}: $kT = (3.6-5.1)$ keV, $N({\\rm H}) = (4.4-5.3) \\times 10^{22}$ cm$^{-2}$, $Z=(0.8-1.3)~Z_{\\sun}$, and $L_{\\rm X} = 2.8 \\times 10^{30}$ \\es.  The source was seen flaring during one of the ASCA observations (\\citealp{kkt+97}) and during the \\chandra\\ observation. The two flares had similar intensity and duration with an $e$-folding time of $\\sim 10$~ks (\\citealp{tik+00}; \\citealp{ikt01}). ",
        "conclusions": "So far, the Fe 6.4 keV emission in YSOs has been explained in terms of fluorescent emission from the photoionised (colder) material in the circumstellar disk. This scenario, however, cannot easily explain the observed variability of the Fe 6.4 keV emission in Elias 29, which occurs in the absence of signficant variations of the observed X-ray continuum. The equivalent width of the line at its maximum strength of $\\sim 250$ eV is also not easily reconciled with a photoionisation scenario. An alternative line-formation mechanism is collisional excitation by a population of non-thermal electrons. We suggest that these electrons could be accelerated by magnetic reconnection events in the accretion tubes that connect the star to its circumstellar disk. The electrons are decelerated in situ by the accreting material, ionising it and causing the observed Fe 6.4 keV emission."
    },
    "0710/0710.0566_arXiv.txt": {
        "abstract": "We present new observations of the fundamental ro-vibrational CO spectrum of V1647 Ori, the young star whose recent outburst illuminated McNeil's Nebula. Previous spectra, acquired during outburst in 2004 February and July, had shown the CO emission lines to be broad and centrally peaked---similar to the CO spectrum of a typical classical T Tauri star.   In this paper, we present CO spectra acquired shortly after the luminosity of the source returned to its pre-outburst level (2006 February) and roughly one year later (2006 December and 2007 February). The spectrum taken in 2006 February revealed blue-shifted CO absorption lines superimposed on the previously observed CO emission lines. The projected velocity, column density, and temperature of this outflowing gas was 30 km s$^{-1}$, $3^{+2}_{-1}\\times10^{18}$ cm$^{-2}$, and 700$^{+300}_{-100}$ K, respectively. The absorption lines were not observed in the 2006 December and 2007 February data, and so their strengths must have decreased in the interim by a factor of 9 or more.  We discuss three mechanisms that could give rise to this unusual outflow. ",
        "introduction": "McNeil's Nebula was recently illuminated by the outburst of V1647 Ori (McNeil 2004), a young star that is embedded in the Lynds 1630 dark cloud and coincides with the 850$\\micron$ continuum source OriBsmm55 (Mitchell et al. 2001). V1647 Ori has a flat SED in the mid-infrared, and is thus a Class I young stellar object (YSO; Andrews et al. 2004). V1647 Ori underwent a similar outburst as recently as 1966 (Aspin et al. 2006), indicating that V1647 Ori is also an EXor. Such pre-main sequence stars undergo eruptive events that dramatically increase their luminosity for periods of months to years (Hartmann 1998). The outbursts are thought to be triggered by a rapid increase in the stellar accretion rate (Hartmann \\& Kenyon 1996).  The eruption in November 2003 of V1647 Ori lasted two years.  During the outburst, the star brightened by a factor of 50 in  X-rays (Kastner et al. 2006), a factor of 250 in the red (6 mag in the $R_C$ band; Fedele et al. 2007a; Brice\\~{n}o et al. 2004), a factor of 15 in the near-IR($\\sim$3 mag in the $J$, $H$, and $K$ bands; Reipurth \\& Aspin 2004), and a factor of $\\sim$15 at wavelengths from $3.6\\micron$ to $70\\micron$ (Muzerolle et al. 2005). From the overall brightening of the source, Muzerolle et al. (2005) concluded that the bolometric luminosity increased by a factor of 15 to 44$L_\\sun$ (see also Andrews et al. 2004) and that the stellar accretion rate increased from $\\sim10^{-7} M_{\\sun}$ yr$^{-1}$ to $\\sim10^{-5} M_{\\sun}$ yr$^{-1}$.  Similarly, Gibb et al. (2006) inferred a stellar accretion rate of $3-6\\times10^{-6} M_{\\sun}$ yr$^{-1}$ from the luminosity of the Br$\\gamma$ emission one year later. This is somewhat larger than the typical accretion rate of a young low mass star ($10^{-8}-10^{-7} M_{\\sun}$ yr$^{-1}$; Bouvier et al. 2007), yet lower than is expected for a star of the FUor type ($\\sim10^{-4} M_{\\sun}$ yr$^{-1}$; Hartman \\& Kenyon 1996).  During the onset of the outburst of V1647 Ori, observations of atomic lines with P Cygni profiles provided evidence for a hot ($T\\sim$10,000 K), high velocity ($v = -400$ km s$^{-1}$) wind (Brice\\~{n}o et al. 2004; Reipurth et al. 2004; Vacca et al. 2004; Walter et al. 2004; Ojha et al. 2006; Fedele et al. 2007a) with a mass-loss rate of $\\dot{M}_{\\rm wind}$=4$\\times$10$^{-8}$ \\Mdot (Vacca et al. 2004). This mass loss rate is much lower than that of the typical FUor (Hartmann \\& Kenyon 1996) and comparable to that of a classical T Tauri star (cTTS; Hartigan et al. 1995). The absorption component of the P Cygni profile of several lines (e.g. Pa$\\beta$) disappeared within a few months following the peak of the outburst in early 2004 (Gibb et al. 2006). However, P Cygni structure in the H$\\alpha$ profile indicated that a weaker wind continued throughout the outburst phase (Ojha et al. 2006; Fedele et al. 2007a).  In contrast to the hydrogen and helium lines,  the fundamental near-infrared ro-vibrational emission lines of CO, observed 2004 February 27, were broad, centrally peaked, and compatible in their intensity with an excitation temperature of 2500 K (Rettig et al. 2005).  The width of the lines was shown to be consistent with Keplerian orbital motion of the gas within the inner disk surrounding the central star, similar to the broad emission line profiles that are observed around cTTSs and Herbig Ae/Be stars (HAeBes; Najita et al. 2003; Blake \\& Boogert 2004).  A later observation, obtained on 2004 July 30, showed that the CO lines remained broad but the temperature of the gas decreased to 1700 K (Gibb et al. 2006).   Neither observation showed any indication of CO in an outflow, as a blue shifted absorption component was not detected.  We report followup observations of CO from V1647 Ori, which were acquired 2006 February, 2006 December, and 2007 February. By the time of our initial 2006 observation, V1647 Ori had returned to quiescence and the absorption component in the H$\\alpha$ line profile had disappeared (Fedele et al. 2007a and references therein). Presumably the accretion rate had fallen by two orders of magnitude to its pre-outburst accretion rate as the star faded to its pre-outburst brightness (see Muzerolle et al. 2005).  As suggested by the work of Najita et al. (2003),  only a continued decrease in CO line intensity from the warm gas in the disk was therefore expected.  However, the first post-outburst observation revealed the striking metamorphosis of these lines from centrally peaked emission features to emission lines with blue-shifted absorption. Subsequently, by late 2006 and early 2007, the CO emission lines returned to their original centrally peaked structure, indicating that the production of the outflow diminished within one year of the end of the outburst.  In this paper we discuss  three scenarios that can give rise to such a phenomenon. ",
        "conclusions": "During quiescence, V1647 Ori is a class I YSO (Andrews et al. 2004). However, its mass-loss rate during outburst was similar to a strongly accreting cTTS (Vacca et al. 2004). While outflows from cTTSs and HAeBes are common, fundamental ro-vibrational CO emission lines with a blue-shifted absorption component have not been observed around any of the more than 300 such sources observed to date (Najita et al. 2000, 2003; Blake \\& Boogert 2004; Rettig et al. 2006; Brittain et al. 2007; J. Brown private communication). Further, Class I YSOs such as GSS 30 IRS 1, HL Tau and RNO 91 do not show ro-vibrational CO emission lines with blue-shifted absorption components (Pontoppidan et al. 2002, Brittain et al. 2005, and Rettig et al. 2006, respectively). In contrast to these systems, the FUor V1057 Cyg has shown blue-shifted overtone ro-vibrational CO absorption lines (Hartmann et al. 2004).  While the similarity of the unusual CO outflows is intriguing, there are important differences between FUors such as V1057 Cyg and EXors such as V1647 Ori. First, FUors have greater accretion rates ($\\sim$10$^{-4}$ \\Mdot) and more extreme winds than EXors (Hartmann \\& Kenyon 1996). Secondly, CO is always and {\\it only} detected in absorption toward FUors, and is thought to originate in the accretion disk (Calvet et al. 1993; Calvet et al. 1991). Despite the significant differences between V1057 Cyg and V1647 Ori, both stars have undergone major eruptions that have generated outflows in CO, suggesting that a different mass-loss mechanism may be at work in these systems than in other young stars.  Given the unusual nature of the outflowing CO from V1647 Ori, it is of interest to consider the mechanisms that could give rise to this phenomenon. One possibility is that the absorbing CO condensed out of the hot outflowing wind that was observed during the outburst.  There were two outflow components noted in the H$\\alpha$ absorption feature: a variable component at 400 km s$^{-1}$ and a steady component at 150 km s$^{-1}$ (Fedele et al. 2007a).  If the lower-velocity gas decelerated at a constant rate to $30$ km s$^{-1}$ (the velocity of the outflowing  CO), by the time of our second CO observation in 2004 July the wind would have expanded to 10 AU. There was no evidence of blue-shifted CO absorption on this date (Fig. 1). By 2006 February, when the absorption was observed, the moderate velocity wind would have expanded to nearly 40 AU. However, it seems unlikely that the wind could still retain a kinetic temperature of 700K at a distance of 40AU from the star. Furthermore, it is not clear why CO would condense out of this outflow to reveal warm, blue shifted absorption but not out of any of the other outflows with similar or even greater mass-loss rates. Thus we conclude that this scenario is unlikely.  A second possibility is that the CO absorption formed in a shell of material that was swept up by the atomic wind.  In this case the absorption did not appear until the column density of material was sufficient to produce measurable absorption, and the heating was the result of the interaction of the wind with the shell. When the mass loss decreased at the end of the outburst phase, this heating was eventually shut down.  If the CO/H$_2$ ratio in such a shell was similar to that of a dense molecular cloud, 1.5 $\\times$10$^{-4}$, then the column density of gas was 2$\\times$10$^{22}$ cm$^{-2}$. Adopting a normal interstellar extinction--to--gas ratio, i.e., $A_V=5.6\\times 10^{22}  N_{\\rm H}$ mag cm$^2$ atom$^{-1}$ (Bohlin et al. 1978), we find that the observed column density of CO corresponded to $\\sim$20 mags of visible extinction. This is a lower limit, as it is possible that selective dissociation of gas in the nebula could drive down the relative abundance of CO.  However, the extinction measured on the line of sight toward V1647 Ori appears to have remained relatively unchanged over the entire course of the outburst at $\\sim$11 mags (Vacca et al. 2004; Gibb et al. 2006), of which 6.5 mags is due to the nebula (Fedele et al. 2007b).  There is no evidence for an additional 10--20 mags of extinction toward V1647 Ori, and so we conclude that this scenario is also unlikely.  A final scenario we consider is that the CO outflow was launched in response to the reorganization of the stellar magnetic field following the sharp drop in the accretion rate.  Pre-main sequence stars tend to have kilogauss magnetic fields which mediate stellar accretion and outflows (e.g. Johns-Krull et al. 2000).  This stellar magnetic field truncates the accretion disk where the ram pressure from accretion balances the magnetic pressure from the magnetosphere (Camenzind 1990).  Consequently, the truncation radius of the disk, $R_T$, is inversely and nonlinearly proportional to the accretion rate, $\\dot{M}$, and given by $ R_T \\propto \\dot{M}^{-2/7}$ (e.g. Bouvier et al. 2007).  Thus the two order of magnitude change in the stellar accretion rate experienced by V1647 Ori would have resulted in the truncation radius being shifted by nearly a factor of four.  Two years into the outburst, V1647 Ori rapidly faded by a factor of 40 in the R$_C$-band in just 180 days to return to its pre-outburst level (Fedele et al. 2007a).  The mid-infrared flux also returned to its pre-outburst level, indicating that the drop in the optical/near-infrared lightcurve was due to the intrinsic fading of the source (Fedele et al. 2007a). This sharp drop in the luminosity of the source indicates that the accretion rate fell rapidly as the star returned to quiescence. In response to the drop in ram-pressure from the accretion flow, the truncation radius was pushed back by the magnetosphere. Evidence from simulations suggests that disk-magnetosphere systems tend to form outflows when the system is undergoing the greatest amount of dynamical rearrangement (e.g., Fig. 7 in Balsara 2004). It is possible that the realignment of the magnetic field as it pushed out against the circumstellar disk resulted in a warm, shortlived outflow, an outflow that is not observed toward any other cTTSs or HAeBes.  While virtually all accreting low-mass stars drive outflows, the CO outflow from V1647 Ori is highly unusual. Indeed, the transformation of centrally peaked ro-vibrational CO emission lines to CO emission lines with blue-shifted absorption is unique.  We suggest that the mechanism responsible for producing this outflow is distinct from the one that drives the outflow from typical cTTSs. The coincidence between the rapid fading of V1647 Ori and the subsequent observation of the CO outflow hints at a connection. Better sampling of the fundamental ro-vibrational CO spectrum of EXors as they brighten and fade is crucial for determining whether this coincidence is significant. The rapid and dramatic change in the accretion rate that characterizes the EXor phenomenon provides an important opportunity to study the interplay between stellar accretion, the inner disk, and outflows. This insight is key to reaching a satisfactory theoretical understanding of these events."
    },
    "0710/0710.2563_arXiv.txt": {
        "abstract": "In this study we compile for the first time comprehensive data sets of solar and stellar flare parameters, including flare peak temperatures $T_p$, flare peak volume emission measures $EM_p$, and flare durations $\\tau_f$ from both solar and stellar data, as well as flare length scales $L$ from solar data. Key results are that both the solar and stellar data are consistent with a common scaling law of $EM_p \\propto T_p^{4.7}$, but the stellar flares exhibit $\\approx 250$ times higher emission measures (at the same flare peak temperature). For solar flares we observe also systematic trends for the flare length scale $L(T_p) \\propto T_p^{0.9}$ and the flare duration $\\tau_F(T_p) \\propto T_p^{0.9}$ as a function of the flare peak temperature. Using the theoretical RTV scaling law and the fractal volume scaling observed for solar flares, i.e., $V(L) \\propto L^{2.4}$, we predict a scaling law of $EM_p \\propto T_p^{4.3}$, which is consistent with observations, and a scaling law for electron densities in flare loops, $n_p \\propto T_p^2/L \\propto T_p^{1.1}$. The predicted ranges of electron densities are $n_p \\approx 10^{9-10}$ cm$^{-3}$ for solar nanoflares at $T_p=1$ MK, $n_p \\approx 10^{10-11}$ cm$^{-3}$ for typical solar flares at $T_p=10$ MK, and $n_p \\approx 10^{11-12}$ cm$^{-3}$ for large stellar flares at $T_p=100$ MK. The RTV-predicted electron densities were also found to be consistent with densities inferred from total emission measures, $n_p=\\sqrt{EM_p/q_V V}$, using volume filling factors of $q_V=0.03-0.08$ constrained by fractal dimensions measured in solar flares. Solar and stellar flares are expected to have similar electron densities for equal flare peak temperatures $T_p$, but the higher emission measures of detected stellar flares most likely represents a selection bias of larger flare volumes and higher volume filling factors, due to low detector sensitivity at higher temperatures. Our results affect also the determination of radiative and conductive cooling times, thermal energies, and frequency distributions of solar and stellar flare energies. ",
        "introduction": "Scaling laws provide important diagnostics and predictions for specific physical models  of nonlinear processes such as self-organized criticality, turbulence, diffusion, plasma heating, and particle accleration. These models have been widely applied in plasma physics, astrophysics, geophysics, and the biological sciences. Here we investigate scaling laws of physical parameters in solar and stellar flares, which should allow us to decide whether solar and stellar flare data are consistent with the same physical flare process.  The scaling of solar and stellar flare data has been pioneered by Stern (1992), Feldman et al.~(1995b), and Shibata \\& Yokoyama (1999; 2002), who showed evidence for a nonlinear scaling between the flare volume emission measure $EM_p$ and the flare peak temperature $T_p$. These parameters have been measured in solar flares with instruments like {\\sl Skylab}, {\\sl GOES}, {\\sl Yohkoh/SXT} (Soft X-ray Telescope), and {\\sl RHESSI (Ramaty High Energy Solar Spectroscopic Imager)}, and in stellar flares with {\\sl ASCA}, {\\sl BeppoSAX}, {\\sl Einstein}, {\\sl EUVE}, {\\sl EXOSAT}, {\\sl Ginga}, {\\sl HEAO}, {\\sl ROSAT}, {\\sl Chandra}, and {\\sl XMM-Newton}. Compilations of solar flare parameters have been presented in Aschwanden (1999), while stellar flare parameters were compiled in a recent review by G\\\"udel (2004). In this paper we present for the first time this host of mostly new measurements ``on the same page'' and investigate commonalities and differences between the scaling of solar and stellar flares.  In Section 2 we present the statistical correlations found in stellar flare data, while the corresponding counterparts of solar flare data are shown in Section 3. In Section 4 we present theoretical modeling of the data, using the well-known RTV law, the generalization with gravitational stratification and spatially non-uniform heating, the fractal flare volume scaling, and volume filling factor. In Section 5 we discuss the differences between solar and stellar scaling laws, the consistency between two different electron density measurement methods, and a previously derived ``universal scaling law'' for solar and stellar flares. Section 6 summarizes our conclusions. ",
        "conclusions": "We compiled directly observed parameters from solar and stellar flares, such as the volume peak emission measure $EM_p$, flare peak temperature $T_p$, flare duration $\\tau_f$, and flare length scale $L$ (the latter only for solar flares). A prominent statistical correlation is found between the volume emission measure $EM_p$ and flare peak temperature $T_p$, which scales as $EM_p \\approx T_p^{4.7}$ for both solar and stellar flares. Another recent study demonstrated that the flare volume has a fractal scaling, $V(L)\\propto L^{2.4}$, rather than the generally used Euclidian scaling of $V(L)\\propto L^3$. Applying the RTV scaling law, combined with the fractal volume scaling and the statistical $L-T_p$ correlation $L(T_p) \\propto T_p^{0.9}$, leads directly to a theoretically predicted scaling law of $EM_p \\propto T_p^{4.3}$, which explains the observed correlations in both solar and stellar flares.  A second result we find is an unexplained offset by a factor of about 250 between solar and stellar flares at the same temperature, which is likely due to a selection bias for stellar flare events with larger volume filling factors and larger spatial scales. Interestingly, however, this selection bias does not affect the overall $EM-T$ relationship and the lower threshold has a similar functional dependence of $EM_{min} \\propto T_p^{4.7}$, probably because the detector sensitivities are dropping off with a similar function with higher temperatures.  A third result is that our model of fractal flare volume scaling provides realistic estimates of volume filling factors, and thus of flare densities. We find that the electron densities in solar flare loops can be predicted based on our fractal scaling in close agreement to the predictions of the RTV law. The agreement of the predicted electron densities with both methods agrees always better than an order of magnitude (Fig.~7), although the absolute magnitude varies by three orders of magnitude between the smallest nanoflares and the largest solar flares, i.e., $n_e \\approx 10^9-10^{12}$ cm$^{-3}$. Since the RTV scaling law (combined with the observed $L \\propto T_p$ correlation) predicts about a linear relationship between the electron densities and flare peak temperatures, i.e., $n_p \\propto T_p$, we expect up to an order of magnitude higher electron densities in the largest stellar flares due to the higher temperature than in solar flares.  The determination of correct scaling laws allows us also to infer realistic estimates of the (conductive and radiative) flare cooling times, which can be tested from the e-folding decay time of individual peaks in (solar and stellar) flare light curves.  The scaling laws allow us also to eliminate temperature biases in the statistics of the total thermal energy of flares, i.e., $E_T \\propto n_p T_p V \\propto EM_p T_p/n_p$. Since we find that the electron density (corrected for a fractal filling factor) scales approximately as $n_p \\propto T_p^{1.1}$, the thermal energy scales approximately as $E_T \\propto EM_p$, and thus the observed total emission measure $EM_p$ can be used as a good proxy for the thermal flare energy $E_T$. Such unbiased frequency distributions of flare energies $N(E_T)$ permit us then to determine whether there is more energy in large or small flares, an important test for nanoflare heating theories."
    },
    "0710/0710.0799_arXiv.txt": {
        "abstract": "{ Non thermal emission from galaxy clusters demonstrates the existence of relativistic particles and magnetic fields in the Intra Cluster Medium (ICM). Present instruments do not allow to firmly establish the energy associated to these components. In a few years gamma ray observations will put important constraints on the energy content of non thermal hadrons in clusters, while the combination of radio and hard X-ray data will be crucial to measure the energy content in the form of relativistic electrons and magnetic field. SIMBOL-X is expected to drive an important breakthrough in the field also because it is expected to operate in combination with the forthcoming low frequency radio telescopes (LOFAR, LWA). In this contribution we report first estimates of {\\it statistical properties} of the hard X--ray emission in the framework of the {\\it re-acceleration model}. This model allows to reproduce present radio data for Radio Halos and to derive expectations for future low frequency radio observations, and thus our calculations provide hints for observational strategies for future radio and hard--X-ray combined observations. ",
        "introduction": "Clusters of galaxies represent the largest virialized structures in the present Universe. Rich clusters have typical total masses of $10^{15} M_{\\odot}$, mostly in the form of dark matter, while $\\sim 5\\%$ of the mass is in the form of a hot ($T \\sim 10^8 K$), tenuous ($n_{gas} \\sim 10^{-3}-10^{-4} cm^{-3}$), X-ray emitting gas. In terms of energy density, the gas is typically heated to roughly the virial temperature, but there is also room to accomodate a non-negligible amount of non-thermal energy.  Clusters are ideal astrophysical environments for particle acceleration and cosmic rays (CR) accelerated within the cluster volume are expected to be confined for cosmological times (e.g., Blasi, Gabici, Brunetti 2007, BGB07, for a review). The bulk of the energy of these CRs is expected in protons since they have radiative and collisional life--times much longer than those of the electrons. While present gamma ray observations can only provide upper limits to the average energy density of CR protons in the ICM (e.g. Reimer et al. 2004), evidence of a non-thermal component is in fact obtained from radio observations of a fraction of galaxy clusters showing synchrotron emission on Mpc scales : Radio Halos, fairly symmetric sources at the cluster center, and Radio Relics, elongated sources at the cluster periphery (e.g., Feretti 2005).  Although the bulk of present data comes from radio observations, theoretically a substantial fraction of the non thermal radiation is expected from inverse Compton (IC) scattering of the photons of the cosmic microwave background (e.g., Sarazin 1999). Measuring IC emission from clusters in the hard X--rays is extremely important to derive the energy density of emitting electrons and the strength of the magnetic field when these measures are combined with radio data. Despite the poor sensitivity of present and past hard X--ray telescopes, several groups have claimed detection of hard X--ray emission (HXR) in a few massive clusters (e.g., Fusco-Femiano et al.~2004; Petrosian et al.~2006; Rephaeli et al.~2006; see also Rossetti \\& Molendi 2004 and Fusco-Femiano et al.~2007 for a discussion on the strength of the HXR detection in the Coma cluster).  Thanks to its sensitivity and capability to perform hard X--ray imaging SIMBOL--X will open a new era in the study of non thermal radiation from galaxy clusters. In this contribution we report first expectations on the Luminosity Functions (LFs) and number counts of HXR from clusters. We calculate only the contribution to the IC spectrum from electrons re-accelerated by turbulence in the ICM which are the responsible for the origin of Radio Halos in the context of the {\\it re-acceleration scenario}. ",
        "conclusions": "In this contribution we report first expectations for HXRs from galaxy clusters, a more detailed study will be reported in a forthcoming paper. Calculations are performed in the framework of the {\\it re-acceleration model} assuming physical parameters which allow the {\\it re-acceleration model} to reproduce present data of the statistical behaviour of giant Radio Halos.  The strength of the magnetic field in the ICM is a crucial parameter in our calculations and we have shown that SIMBOL-X will provide unique constraints. By assuming a value of the magnetic field averaged in Mpc$^3$ volume of $\\approx$0.2$\\mu$G we find that SIMBOL-X will discover HXRs in $\\approx$30--100 clusters at z$\\leq$0.2 ."
    },
    "0710/0710.0367_arXiv.txt": {
        "abstract": "{ Combined X-ray synchrotron and inverse-Compton $\\gamma$-ray observations of pulsar wind nebulae (PWN) may help to elucidate the processes of acceleration and energy loss in these systems. In particular, such observations provide constraints on the particle injection history and the magnetic field strength in these objects. The newly discovered TeV $\\gamma$-ray source HESS\\,J1718$-$385 has been proposed as the likely PWN of the high spin-down luminosity pulsar PSR\\,J1718$-$3825. The absence of previous sensitive X-ray measurements of this pulsar, and the unusual energy spectrum of the TeV source, motivated observations of this region with \\emph{XMM-Newton}. The data obtained reveal a hard spectrum X-ray source at the position of PSR\\,1718$-$3825 and evidence for diffuse emission in the vicinity of the pulsar. We derive limits on the keV emission from the centroid of HESS\\,J1718$-$385 and discuss the implications of these findings for the PWN nature of this object.  } ",
        "introduction": "Young pulsars drive relativistic winds into their environments, confinement of which leads to the production of extremely broadband emission via the synchrotron and inverse-Compton (IC) processes \\citep[see][for a recent review]{PWN:review}. The most prominent PWN, the Crab Nebula, is detected in all wavebands from the radio to TeV $\\gamma$-rays~\\cite{Whipple:crab}, with the transition from synchrotron to IC emission at $\\sim$1~GeV. The recent increase in sensitivity of ground-based TeV $\\gamma$-ray instruments has led to a rapid increase in the number of putative PWN in this waveband. These objects are characterised by diffuse, typically offset, nebulae around high spin-down luminosity pulsars. The archetype of this new object class is the PWN G\\,18.0---0.7/HESS\\,J1825$-$137. G\\,18.0---0.7 is a $\\sim$5$'$ long asymmetric X-ray synchrotron nebula associated with the middle-aged (characteristic spin-down age $\\tau$=21~kyr) pulsar PSR\\,B1823$-$13~\\cite{XMM:1825}.  The IC nebula HESS\\,J1825$-$137 is much larger ($\\sim$100$'$ at 1~TeV) but exhibits energy-dependent morphology, shrinking towards the pulsar at high energies~\\cite{HESS:1825p2,HESS:1825icrc}, suggestive of cooling of the highest-energy (X-ray synchrotron emitting) electrons away from the pulsar.  The TeV $\\gamma$-ray source HESS\\,J1718$-$385 was discovered in deep observations of the supernova remnant RX\\,J1713.7$-$3946 using H.E.S.S. in 2004-2005~\\cite{HESS:twopwn}.  The absence of other potential counterparts and the relatively compact nature of the source ($9'\\times4'$ rms) make an association with PSR\\,J1718$-$3825 ($8'$ from the centroid of the TeV source) plausible.  The TeV source is unusual in its sharply peaked spectral energy distribution (SED), which is similar to that of the $\\gamma$-ray nebula of the Vela pulsar~\\cite{HESS:velax}.  The $\\gamma$-ray emission from these objects is commonly attributed to IC scattering of relativistic electrons \\citep[see][for an alternative view]{Horns:VelaX}. In this scenario the spectral break seen at $\\sim$10~TeV in these objects can be interpreted as a signature of electron cooling. However, PSR\\,J1718$-$3825 (estimated distance  4.2~pc) has a characteristic spin-down age (90~kyr) almost an order of magnitude greater than that of the Vela pulsar, making such a high energy break very surprising.  The search for a possible X-ray counterpart to HESS\\,J1718$-$385 is important for two reasons: firstly, to verify the identification of the TeV source as the PWN of PSR\\,J1718$-$3825 and secondly, to explore the physical conditions and electron energy distribution in the putative nebula. As no sensitive X-ray observations of the PSR\\,J1718$-$3825/HESS\\,J1718$-$385 region existed, \\emph{XMM-Newton} was used to observe this region in September 2006. ",
        "conclusions": "The discovery of hard spectrum X-ray emission from the vicinity of PSR\\,J1718$-$3825, and the evidence for a diffuse halo around the pulsar strongly suggest the existence of a synchrotron nebula around this pulsar. This discovery strengthens the association of the $\\gamma$-ray source HESS\\,J1718$-$385 to PSR\\,J1718$-$3825, but the relationship of the X-ray emission to the $\\gamma$-ray source is not straightforward. The overall asymmetry of the nebula with respect to the pulsar is consistent with the idea of SNR expansion into a non-uniform molecular environment \\citep[see for example][]{blondin01:PWN}. The very different morphologies in the two wavebands suggest that either electrons of rather different energies are responsible for the two sources and/or that the magnetic field strength within the nebula is highly non-uniform. As the target for IC emission is the CMBR and other large-scale radiation fields, the IC flux $F_{\\rm IC}$ is simply proportional to the number of radiating electrons, $n_{e}$, whereas the synchrotron flux goes as: $F_{\\rm synch}\\propto B^{2}n_{e}$. In either case the situation may be rather similar to that of HESS\\,J1825$-$137 or indeed HESS\\,J1813$-$178 \\cite{Funk:1813}, with the lifetime of TeV $\\gamma$-ray emitting electrons being longer than the age of the pulsar, and having time to propagate over distances of several parsecs. The SED of the source is presented in Figure~\\ref{fig:sed}. Three features are of note:  \\begin{figure}[t] \\centering \\resizebox{0.9\\hsize}{!}{\\includegraphics{Figure3}} \\vspace{-1mm} \\caption{Spectral energy distribution for the pulsar wind nebula of PSR\\,J1718$-$3825. The de-absorbed spectrum of emission from within $1'$ of the pulsar is shown together with limits for diffuse emission from the region covered by HESS\\,J1718$-$385. Three sets of illustrative synchrotron and inverse Compton model curves are shown, based on assumptions of: {\\bf A)} Mono-energetic 70~TeV electrons injected over a $10^{4}$ year period, $B = 5\\mu$G, {\\bf B)} 70 TeV electrons, $B=20\\,\\mu$G, $t=8$ years and {\\bf C)} An electron energy distribution following a power-law ($\\alpha=1.8$) with an abrupt cut-off at 100~TeV. These curves are calculated as described in \\citet{GC:Hinton07}. } \\label{fig:sed} \\end{figure}  1) a hard IC spectrum at TeV energies, with a peak at $\\sim$10~TeV. This suggests that the electrons responsible for this emission are uncooled. This is rather surprising given the spin-down age of the pulsar (90 kyr): cooling on the CMBR alone would result in a spectral break at 2~TeV after 90~kyr, inconsistent with the $\\gamma$-ray data. Indeed ages $>40$ kyrs appear to be excluded by the data. A true age of $\\sim$10~kyr could be explained by a birth period for the pulsar very close to its current period of 75~ms, or breaking deviating significantly from the pure magnetic dipole case (as it seems may commonly be the case, see \\citet{Kramer:Spindown}) \\footnote{We further note that the projected length of the $\\gamma$-ray nebula ($\\sim$10~pc) is roughly half that of HESS\\,J1825$-$137, despite the fact that PSR\\,J1718$-$3825 is apparently a factor 5 older. Another possibility is that HESS\\,J1825$-$137 represents only the youngest part of a larger, softer spectrum, $\\gamma$-ray nebula.}. The shape of the TeV spectrum appears to be consistent with a constant injection of (mono-energetic) $\\sim$70~TeV electrons (curve {\\bf A}), or with a hard power-law (index $\\approx$1.8) with a sharp cut-off around 100 TeV ({\\bf C}).  2) hard spectrum X-ray  emission from the pulsar vicinity with a much lower energy flux. The spectrum from within $1'$ of PSR\\,J1718$-$3825 shown in Figure~\\ref{fig:sed} represents the combined flux of pulsar itself and the inner PWN. The pulsar contribution is certainly less than half of the total emission (as is clear from the flux in the annulus surrounding the pulsar) and for typical systems of this type represents only $\\sim$20\\% of the PWN emission in the $>2$keV range \\citep{Kargaltsev:ChandraPWN}. Assuming the PWN is dominant, the hard spectrum suggests either higher electron energies or larger magnetic fields in this region ({\\bf B}) in comparison to those found in the $\\gamma$-ray nebula. The lower flux can be explained if only recently injected electrons are confined in the region around the pulsar. Whilst the two-zone scenario illustrated by curves {\\bf A} and {\\bf B} is clearly grossly oversimplified, is does appear that the data are consistent with the idea that electrons with a relatively narrow energy distribution rapidly escape from a high $B$-field region close to the pulsar into the extended nebula seen in $\\gamma$-rays. Another plausible scenario is that the injection spectrum of electrons has changed significantly over the lifetime of the pulsar.  3) the energy flux level of diffuse X-ray emission from the HESS\\,J1718$-$385 region exceeds the TeV flux by not more than a factor $\\sim$2. The energy distribution of electrons is essentially fixed by the TeV data. The diffuse X-ray limits can therefore be used constrain the $B$-field in the extended  nebula to be not much greater than $5\\,\\mu$G (curves {\\bf C}+{\\bf A}), close to the mean value of the ISM. We note that this constraint comes principally from the diffuse limit at 2--4.5 keV which is relatively independent of the assumed absorbing column.   In conclusion, the diffuse X-ray emission around XMMU\\,171813.8$-$382517 and HESS\\,J1718$-$385 appear to represent different zones in the PWN of the middle-aged pulsar PSR\\,J1718$-$3825. Future studies of this complex system are certainly well motivated. For example, with the superior angular resolution of Chandra, the contribution of the pulsar itself to the non-thermal emission could be separated from that of the PWN."
    },
    "0710/0710.5774_arXiv.txt": {
        "abstract": "We review the properties of carbon-sequence ([WC]) Wolf-Rayet central stars of planetary nebulae (CSPNe). Differences  between the subtype distribution of [WC] stars and their massive WC cousins are discussed. We conclude that [WO]-type differ from early-type [WC] stars as a result of weaker stellar winds due to high surface gravities, and that late- and early-type [WC] and [WO] stars generally span a  similar range in abundances, X(He)$\\sim$X(C)$\\gg$X(O), consistent with a late thermal pulse, and likely progenitors to PG1159 stars. ",
        "introduction": "%  This review discusses the properties of the small fraction of central stars of Planetary Nebulae (CSPNe) which share a spectroscopic appearance with massive, carbon-sequence (WC-type) Wolf-Rayet stars. Massive WC stars are the chemically evolved descendents of  initially very massive O stars  ($M_{\\rm init} \\geq 25 M_{\\odot}$) exhibiting the C and O products of core helium burning, plus a unique emission line spectral appearance due to fast, dense stellar wind outflows. Such stars are young, with ages of only a few Myr of which several hundred cases are known within the Milky Way, supplemented by thousands more known in external star-forming galaxies (Crowther 2007; Hamann these proc.).  CSPNe possessing a similar spectral morphology to WC stars are denoted [WC] and are at a post-Asymptotic Giant Branch (AGB) phase in the late stages of evolution of low or intermediate mass stars ($M_{\\rm init} \\sim 1-5 M_{\\odot}$?) with only a few dozen examples known in the Milky Way plus a handful in the Magellanic Clouds. With respect to normal H-rich CSPNe, the unusual surface chemical composition of [WC] stars apparently results from a late thermal pulse (LTP), causing a H-deficient surface, and likely connection with other H-deficient stars, most notably PG1159 stars (Werner et al. these proc.) ",
        "conclusions": "We discuss physical and wind properties of Wolf-Rayet CSPNe drawn from the recent literature plus new analyses for a range of subtypes. In general the higher ionization lines seen at earlier spectral type indicates an increased stellar temperature, although very early [WO] subtypes  are favoured for hot CSPNe with weak winds, with [WC4--5] subtypes resulting for hot CSPNe with stronger winds. [WC8--9] subtypes correspond to lower temperature CSPNe with strong winds, with much lower temperatures indicated in [WC10] stars. We explain the differences in subtype distribution between Galactic disk CSPNe and massive WC stars as a result of decreased wind densities in the former, owing to increased surface gravities. Massive WC stars  in the LMC differ from  Galactic WC stars through reduced wind densities caused by lower metallicities, rather than increased surface gravities.     There does {\\it not} appear to be a systematic difference between carbon-to-helium mass fractions in late to early-type CSPNe, at least for subtypes for which we are able to employ common diagnostics ([WC9] to [WO1]). Consequently, evolution from late-type [WC] through early-type [WC] and [WO] to PG1159 stars appears to be consistent with most abundance patterns, and expectations for a late thermal pulse. Indeed, we note that similar analysis tools have been applied to the post He-flash system V605 Aql, which has now evolved through to a early-type [WC] spectral type over the past 80 years and shares a similar abundance pattern with X(He):X(C):X(O) = 54:40:5\\% (Clayton et al. 2006)."
    },
    "0710/0710.5542_arXiv.txt": {
        "abstract": "There are periodic solutions to the equal-mass three-body (and $N$-body) problem in Newtonian gravity. The figure-eight solution is one of them. In this paper, we discuss its solution in the first and second post-Newtonian approximations to General Relativity. To do so we derive the canonical equations of motion in the ADM gauge from the three-body Hamiltonian. We then integrate those equations numerically, showing that quantities such as the energy, linear and angular momenta are conserved down to numerical error. We also study the scaling of the initial parameters with the physical size of the triple system. In this way we can assess when general relativistic results are important and we determine that this occur for distances of the order of $100M$, with $M$ the total mass of the system. For distances much closer than those, presumably the system would completely collapse due to gravitational radiation. This sets up a natural cut-off to Newtonian $N$-body simulations. The method can also be used to dynamically  provide initial parameters for subsequent full nonlinear numerical simulations. ",
        "introduction": "The closest star to the solar system, Alpha Centauri, is a triple system, so is Polaris and HD 188753. Triple stars and black holes are common in globular clusters~\\cite{Gultekin:2003xd, Miller:2002pg}, and galactic disks. Triple black hole mergers can be formed in galaxy merger~\\cite{Valtonen96} and a triple quasar, representing a triple supermassive black hole system has been recently discovered~\\cite{Djorgovski:2007ka}.   Full numerical simulations of black holes made possible only in the last couple of years have already produced numerous astrophysically interesting results, among them, the orbital hangup and respect of the cosmic censorship hypothesis for spinning black holes~\\cite{Campanelli:2006uy,Campanelli:2006fg,Campanelli:2006fy}, precession and spin-flips~\\cite{Campanelli:2006fy}, and the discovery~\\cite{Campanelli:2007ew} of large recoil velocities in highly-spinning black hole mergers up to $4,000$ km/s~\\cite{Campanelli:2007cga}.   The 2005 breakthroughs in Numerical Relativity~\\cite{Pretorius:2005gq,Campanelli:2005dd,Baker:2005vv}, not only provided a solution to the long standing two-body problem in General Relativity, but it also proved applicable to the black hole - neutron star binaries~\\cite{Faber:2007dv} and recently to the three (and $N$) - black holes systems~\\cite{Campanelli:2007ea}.   In general, the solution of three-body problem in Newtonian gravity can be chaotic. There are however, periodic orbits in the problem of three equal masses on a plane. One of the most surprising solution is a figure-eight % orbit. The three bodies chase each other forever around a fixed eight-shaped curve. This was found first by Moore~\\cite{Moore:1993} and discussed with the proof of the existence in Ref.~\\cite{Chenciner:2000}. Heggie~\\cite{Heggie:2000} also estimates the probability for such systems to occur in a galaxy.   Because of effects such as the perihelion shift, it was unclear if the figure-eight orbits would exist in a low post-Newtonian expansion, even if it consist of only conservative terms. Imai, Chiba and Asada succeeded in obtaining the figure-eight solution in a first post-Newtonian order approximation by finding the general relativistic corrections to the Newtonian initial conditions. In Ref.~\\cite{Chiba:2006ad} they also estimated the periodic gravitational waves from this system.   In Ref.~\\cite{Imai:2007gn} was used the Euler-Lagrange equations of motion in an approximation to first post-Newtonian order. In our paper we instead assume the Hamiltonian formulation to derive the equations of motion. We start from the Hamiltonian given in Ref.~\\cite{Schaefer87} (with typos corrected in our Appendix). We derive the equations of motion in this formalism, which are different from those used in Ref.~\\cite{Imai:2007gn} and have the virtue of explicitly satisfying the constants of motion of the problem, and thus being more amenable to numerical integration.   The paper is organized as follows. In Section~\\ref{sec:EOM}, we summarize the equations of motion to be solved numerically in order to obtain the figure-eight orbits. The starting point is the three-body Hamiltonian in the first post-Newtonian approximation. In Section~\\ref{sec:1PN}, we discuss the initial conditions for the figure-eight solutions. We study the scaling relation between the orbital radius and the linear momenta. From this analysis, we can estimate when general relativistic effects are important. In Section~\\ref{sec:2PN}, we extend our calculation to the second post-Newtonian order and in Section~\\ref{sec:DIS}, we summarize the results of this paper and discuss some remaining problems. The 2PN three-body Hamiltonian is explicitly given in the Appendix. Throughout this paper, we use units in which $c=G=1$. ",
        "conclusions": "\\label{sec:DIS}  In this paper we have used the figure-eight orbits as a theoretical lab to test the properties of the low post-Newtonian expansions of General Relativity. We have found that those closed orbits exists for three (and presumably $N$) bodies. We have provided an improved first-post-Newtonian order formalism for deriving the equations of motion that satisfy the Hamiltonian (the linear and angular momenta) constraint to round-off error. The subsequent numerical evolution is well behaved during for more than $t\\sim10,000m$. We have also extended this analysis to the $2PN$ corrections, still giving a conservative system of equations. In the process of finding the figure-eight solutions by trial of different initial momenta we also showed (numerically) the stability of the orbit against small perturbations.  This method is particularly useful to determine, dynamically (as an alternative to determine them through families of initial data~\\cite{Campanelli:2005kr}), initial orbital parameters for subsequent full numerical evolution~\\cite{Campanelli:2007ea}, when the holes are close enough that general relativistic effects can no longer be ignored. Note that our method fully takes into account the three-body post-Newtonian interactions unlike other simulations that approximate the problem in successive two-body problems~\\cite{Aarseth:2007wv}.   It is interesting to note here that the scaling fits~(\\ref{eq:2PNfit}) give a practical way to determine when relativistic or Newtonian approaches are appropriate. For $\\lambda=1$ we have that the ratio of the first coefficient, $0.01617654493$ (Newtonian) to the second coefficient $0.002017242451$ first-post-Newtonian is nearly $0.12/\\lambda$ and the second coefficient to the third one $0.0002463605227$ (dominated by second-post-Newtonian) is also approximately $0.12/\\lambda$. This indicates that post-Newtonian corrections are important. For $\\lambda=1$ the distance between the initial bodies is $200m$, what indicates that for nearly $67M$ with $M\\approx3m$ the total mass of the system has strong post-Newtonian effects. For $\\lambda \\gg 1$ Newtonian gravity should describe the system accurately, while for $\\lambda<1$ general relativistic effects should be very important, eventually leading to the total collapse of the system. It is interesting to remark here that most of the $N$-body codes use some sort of regularization of the Newtonian gravity for very close encounters \\cite{aarseth-03}, instead the natural way to regularize these close encounters \\cite{Campanelli:2007ea} is given by the General Theory of Relativity, and as we show here, the post-Newtonian corrections are already non-negligible at separations of the order of $100M$. In any case, for most of the astrophysical encounters this is way too short distance, but it can obviously be reached in systems involving black holes and neutron stars."
    },
    "0710/0710.5632_arXiv.txt": {
        "abstract": "{Homogeneous anisotropic turbulence simulations are used to determine off-diagonal components of the Reynolds stress tensor and its parameterization in terms of turbulent viscosity and $\\Lambda$-effect. The turbulence is forced in an anisotropic fashion by enhancing the strength of the forcing in the vertical direction. The Coriolis force is included with a rotation axis inclined relative to the vertical direction. The system studied here is significantly simpler than that of turbulent stratified convection which has often been used to study Reynolds stresses. Certain puzzling features of the results for convection, such as sign changes or highly concentrated latitude distributions, are not present in the simpler system considered here.} ",
        "introduction": "The Reynolds stress, described by the correlation of fluctuating velocity components, $Q_{ij} = \\overline{u_i u_j}$, is one of the most important generators of differential rotation in stars (R\\\"udiger \\cite{Ruediger1989}). These stresses have been studied with the help of 3D convection simulations (e.g.\\ Pulkkinen et al.\\ \\cite{Pulkkinenea1993}; Chan \\cite{Chan2001}; K\\\"apyl\\\"a et al.\\ \\cite{Kaepylaeea2004}; R\\\"udiger et al.\\ \\cite{Ruedigerea2005}). These results have revealed some surprising features such as the peaking of the horizontal stress $Q_{xy}$ very close to the equator, and a positive (outward) flux for rapid rotation. Both of these results are at odds with theoretical considerations (Kitchatinov \\& R\\\"udiger \\cite{KitcRued1993}). Furthermore, disentangling of the diffusive (turbulent viscosity) and non-diffusive ($\\Lambda$-effect) parts of the stress is difficult from convection simulations.  Here, we present preliminary results from anisotropic homogeneous, isothermal, non-stratified turbulence simulations in which diffusive and non-diffusive effects can be studied separately. Imposing a linear shear flow on top of isotropically driven turbulence allows the study of turbulent viscosity without $\\Lambda$-effect. On the other hand, using a special form of forcing, anisotropic homogeneous turbulence can be generated. Rotation is added to study the $\\Lambda$-effect. A simple analytical closure model, based on the minimal tau-approximation (hereafter MTA, see e.g. Blackman \\& Field \\cite{BlackField2002}; Brandenburg et al.\\ \\cite{Brandea2004}), is used to compare with simulations in the cases with rotation. ",
        "conclusions": "\\label{sec:conclusions} Shear flow turbulence simulations show that the ratio of turbulent to molecular viscosity increases linearly up to ${\\rm Re} \\approx 30$ with $\\nu_{\\rm t}/\\nu\\approx1.5\\,\\mbox{Re}$. For the largest Reynolds number the scaling seems somewhat shallower but the present data is not yet sufficient to substantiate this. The $\\Lambda$-effect from homogeneous, anisotropic turbulence does not exhibit the puzzling features found in convection simulations. Further study is required in order to understand which of the neglected physics is responsible for the lack of these features.  The MTA-closure is able to reproduce many of the qualitative aspects of the simulation results including a maximum of the horizontal stress at about $30^\\circ$ latitude, with its largest value for $\\mbox{Co}\\approx0.5$. In the model, the vertical stress can have a maximum away from the equator for $\\mbox{Co}\\ga0.2$, which is not seen in the simulations. Nevertheless, both simulations and model have the largest vertical stress for $\\mbox{Co}\\approx0.3$. However, the model generally overestimates the magnitudes of the stresses. More detailed analysis of the simulation and closure results will be presented in a future publication."
    },
    "0710/0710.2538_arXiv.txt": {
        "abstract": "We analyze the possibility of delensing Cosmic Microwave Background (CMB) polarization maps using foreground weak lensing (WL) information. We build an estimator of the CMB lensing potential out of optimally combined projected potential estimators to different source redshift bins. Our estimator is most sensitive to the redshift depth of the WL survey, less so to the shape noise level. Estimators built using galaxy surveys like LSST and SNAP recover up to 80-90\\% of the potential fluctuations power at $l\\leq 100$ but only \\ax 10-20\\% of the small-angular-scale power ($l\\leq 1000$). This translates into a 30-50\\% reduction in the lensing $B$-mode power. \\par We illustrate the potential advantages of a 21-cm survey by considering a fiducial WL survey for which we take the redshift depth \\zmax and the effective angular concentration of sources \\nbar as free parameters. For a noise level of 1 $\\mu$K arcmin in the polarization map itself, as projected for a CMBPol experiment, and a beam with $\\theta_{\\rm\\scriptscriptstyle{FWHM}}$=10 arcmin, we find that going to \\zmax=20 at \\nbar=100 \\gal yields a delensing performance similar to that of a quadratic lensing potential estimator applied to small-scale CMB maps: the lensing $B$-mode contamination is reduced by almost an order of magnitude. In this case, there is also a reduction by a factor of \\ax4 in the detectability threshold of the tensor $B$-mode power. At this CMB noise level, the $B$-mode detection threshold is only $3\\times$ lower even for perfect delensing, so there is little gain from sources with $z_{max}>20$. The delensing gains are lost if the CMB beam exceeds $\\sim 20$ arcmin.  The delensing gains and useful $z_{max}$ depend acutely on the CMB map noise level, but beam sizes below 10 arcmin do not help. Delensing via foreground sources does not require arcminute-resolution CMB observations, a substantial practical advantage over the use of CMB observables for delensing. ",
        "introduction": "\\label{I} The anisotropies in the CMB have been long recognized as a major probe for cosmology. The WMAP satellite has measured the temperature fluctuations up to $l_{max} <=$1000 and has confirmed the cosmological standard model of a power-law, flat, $\\Lambda$CDM universe.  Much hope lies with CMB polarization measurements because they have the potential to unveil some of the unknowns of inflation. While some predictions of inflation (such as nearly flat space curvature, nearly scale-invariant power spectrum, and Gaussianity of the primordial fluctuations) have been confirmed by CMB and large scale structure experiments, there is another prediction which is yet to be verified. This is the existence of an almost scale-invariant spectrum of gravitational waves, whose amplitude is directly related to the energy scale of inflation. The inflationary gravitational waves are tensor perturbations to the metric and we expect them to leave a curl-like signature in the CMB polarization field, i.e. a $B$-mode pattern. Density fluctuations, arising from scalar perturbations to the metric, create a gradient-like component in the polarization field, the $E$ mode. In the linear regime, density fluctuations do not create a $B$ mode, so a detection of the latter would confirm the existence of primordial gravitational waves; it would also provide the energy scale of inflation, which we could use to distinguish between different inflationary scenarios.\\par CMB polarization measurements are difficult to carry out, primarily for two reasons. Firstly, the amplitude of the signal is very small: for example, the scalar $E$-mode power is 2-3 orders of magnitude smaller than the scalar temperature power, tensor $B$-mode much lower. Secondly, the polarization foregrounds are very poorly understood and they dominate the CMB signal at almost all relevant frequencies. Therefore, foreground removal and detector sensitivity are two of the most stifling limitations of a polarization experiment.\\par $B$-mode measurements are additionally obstructed by WL contamination, especially of the recombination signal (l$\\leq$100). Thus CMB polarization measurements are able to provide inflationary insights to the extent to which we can: 1. remove the foreground contribution to the overall signal and 2. delens the $B$ mode power and extract the tensor contribution to it. The delensing process consists of the reconstruction of the lensing potential from chosen observables. The estimated lensing potential is then used to evaluate the WL-created $B$-mode signal and to subtract it from the measured $B$-mode map. \\par In this paper we probe the ability of weak lensing (WL) galaxy surveys to delens the CMB in the absence of foregrounds. To be specific, we try to answer two questions: what attributes should a galaxy or 21~cm WL survey and also a CMB polarization mission have in order to detect the tensor $B$ mode? What is the minimum amplitude of the $B$ mode, expressed in terms of the tensor-to-scalar ratio, $r$, that can be detected using galaxy or recombination observations for delensing?  We take three examples of surveys to illustrate how our estimator works: the ground-based Large Synoptic Survey Telescope (LSST \\footnote{\\tt {http://www.lsst.org/}}), the space-based Supernova Acceleration Probe (SNAP \\footnote{\\tt {http://snap.lbl.gov/}}) and a toy model mimicking recombination-era 21-cm observations that we mention in more detail in section \\S\\ref{IV}. The outline of the paper is as follows. In \\S\\ref{II} we describe the WL contamination of the tensor $B$ mode. In \\S\\ref{III} we present our minimum-variance lensing estimator. In \\S\\ref{IV} we determine the minimum detectable $r$ when we use this estimator to delens.  In \\S\\ref{V} we discuss our results and draw conclusions. Let us now briefly mention similar work existing in the literature.  There has been a vast and impressive amount of work on the topic of CMB delensing and $r$-detection. The great majority of this work uses the CMB observables $\\Theta, E, B$ to reconstruct the projected potential. Averaging over various quadratic combinations of the temperature field (e.g. see the work of \\citet{1999PhRvD..59l3507Z}, \\citet{1998A&A...338..767B}, \\citet{2001PhRvD..64h3005H}, \\citet{2001ApJ...557L..79H}) and the polarization field (e.g. \\citet{2000PhRvD..62d3517G}, \\citet{2002ApJ...574..566H}, \\citet{2003PhRvD..67l3507K}) has been thoroughly considered for the projected potential reconstruction. To give a quick summary: one can build minimum-variance, unbiased estimators, using certain field statistics, as shown by \\citet{2001ApJ...557L..79H}. For a post-Planck experiment (sensitivity of 0.3 $\\mu$K arcmin and beam size of 3 arcmin), \\citet{2002ApJ...574..566H} found that the most efficient of these estimators can map the potential up to l$\\leq$1000. \\citet{2002PhRvL..89a1304K} and \\citet{2002PhRvL..89a1303K} used this last estimator to predict the minimum detectable $r$ as a function of CMB experimental characteristics. There is another, more promising method for lensing potential reconstruction, based on likelihood techniques. \\citet{2003PhRvD..67d3001H} have built a maximum-likelihood estimator for the convergence field using temperature maps and have found its performance similar to that of the quadratic estimator introduced by \\citet{2001ApJ...557L..79H}. \\citet{2003PhRvD..68h3002H} found that the same maximum likelihood estimator built from polarization maps is even more effective: there is an order of magnitude reduction in the mean squared error in the lensing reconstruction compared to the quadratic estimator method, if the survey characteristics are adequate: sensitivity of 0.25 $\\mu$K arcmin and a beam size of 2 arcmin. \\citet{2005PhRvD..72l3006A} analyze the detectability of tensor $B$ modes in the presence of polarized dust emission, as a function of sky coverage; \\citet{2006JCAP...01..019V} do a study of optimal surveys for $B$-mode detection, considering both dust and synchrotron emissions.  One disadvantage of the reconstruction methods presented so far is that they require high-resolution CMB maps: they use arc-minute structures of the CMB fields, to reconstruct degree-scale maps of the deflection field, as explained by \\citet{2001ApJ...557L..79H}. \\citet{2005PhRvL..95u1303S} point out that one can use non-CMB observables to delens the CMB; in this case the requirement for high angular resolution of the CMB mission can be relaxed significantly. These authors determine lower limits for the detectable $r$ using the 21 cm radiation emitted by neutral hydrogen atoms to delens the CMB. We follow a similar approach here, but employ foreground galaxies instead of 21-cm emission as the source plane for delensing. ",
        "conclusions": "\\label{V} In this paper we have examined the possibility of delensing $B$-mode polarization maps using galaxy WL surveys. We have proposed a weighted combination of projected potential estimators to different source redshift bins which optimally reconstructs the projected potential seen by the CMB. We have used three fiducial surveys to exemplify our estimator: LSST, SNAP and a generic survey relevant mostly to future 21-cm studies. These examples have different source redshift distribution, source density \\nbar, and sky coverage $f_{\\rm sky}$, and have enabled us to test the effect of each of these factors on the lensing potential estimator and also on its ability to reduce the WL contamination of $B$ mode maps. Using a $\\Delta \\chi^{2}$ test for the delensed $B$ mode field, we have determined the minimum value of the tensor-to-scalar ratio statistically distinguishable from 0 in optimally delensed data. Throughout this paper we have ignored the polarized foreground contamination and we have also made the assumption of Gaussianity for the delensed $B$ mode map. While this assumption may be inaccurate, especially for the cases when the efficiency of delensing is low, we expect the qualitative results of our work to hold.  \\par The lensing potential estimator is sensitive mostly to the redshift depth of the WL survey and reconstructs best the large angular scale multipoles. The performance of the estimator improves continually as higher redshift source galaxies are used. In the limit \\zmax$\\rightarrow z_{\\scriptscriptstyle\\rm CMB}$ and \\nbar$\\rightarrow\\infty$, the reconstruction is perfect. An experiment like SNAP recovers \\ax 90\\% of the CMB projected potential power for $l\\leq 100$ and \\ax 20\\% for $l\\leq 1000$. The performance of LSST is a little worse, because it detects sources at an average redshift of $\\langle z \\rangle$=1, compared to $\\langle z \\rangle$=1.5 of SNAP. A box distribution (constant redshift distribution of source galaxies), provides better CMB potential estimator: if we go as far as \\zmax=20, there is an order of magnitude improvement over SNAP, for the same angular concentration of galaxies.  However, in this last case, the reduction in $r_{min}$, compared to the case where no delensing is done, is only of a factor of \\ax 7 for $w^{-1/2}=0.1 \\mu$K arcmin and $\\theta_{\\rm\\scriptscriptstyle{FWHM}}$=10 arcmin. This happens because low-$l$ lensing $B$-modes are generated by beating of gravitational potential modes and $E$-modes on scales coresponding to $l$\\ax1000, where the potential estimator is less faithful.  In the case of LSST and SNAP, the reduction in $r_{min}$ relative to the no delensing case is only by \\ax15\\% and \\ax50\\% for the same CMB noise level and beam.  We stress that delensing with foreground weak lensing offers significant gains in $r_{min}$ (factors of 5 or more) only when three conditions are met: \\begin{enumerate} \\item the noise in the CMB map is sufficiently low, $\\sim 1\\,\\mu$K arcmin; \\item the beam size of the CMB map is $<20$ arcmin; \\item the lensing source distribution extends to $z\\sim 20$ or greater. \\end{enumerate} Delensing is not relevant to tensor mode detection if $w^{-1/2}>10 \\mu$K arcmin or if $\\theta_{\\rm\\scriptscriptstyle{FWHM}}$ \\gt$2^{\\,\\circ}$. Also, there is no advantage in lowering the beam size beyond $\\theta_{\\rm\\scriptscriptstyle{FWHM}}$=10 arcmin: for our delensing method, a 1 arcmin beam will do just as well as a 10 arcmin beam. This is perhaps the most important feature of our delensing technique, as both the quadratic estimator of \\citet{2001ApJ...557L..79H} and the maximum likelihood estimator proposed by \\citet{2003PhRvD..68h3002H} require beam sizes of 2-3 arcmin to yield their best performance. For the CMBPol mission, (which might reach a noise of 1 $\\mu$K arcmin) delensing with a box of \\zmax=20 results in $r_{min}\\approx 2\\times 10^{-5}$. If no delensing is applied, then $r_{min}\\approx 8\\times10^{-5}$, a factor of 4 worse. CMBPol detector noise keeps $r_{min}>7\\times 10^{-6}$ even with perfect delensing, so there is less incentive to acquire delensing source planes above \\zmax=20. Also for the CMBPol noise level, we have investigated the impact of $\\bar n$ and have concluded that as long as we go to high redshift ($z_{max}>10$), even with a low density of source galaxies, we can still improve $r_{min}$ by a factor of a few."
    },
    "0710/0710.1011_arXiv.txt": {
        "abstract": "We present a 40 minute time series of filtergrams from the red and the blue wing of the \\halpha\\ line in an active region near the solar disk center. From these filtergrams we construct both Dopplergrams and summed ``line center'' images. Several dynamic fibrils (DFs) are identified in the summed images. The data is used to simultaneously measure the proper motion and the Doppler signals in DFs. For calibration of the Doppler signals we use spatially resolved spectrograms of a similar active region. Significant variations in the calibration constant for different solar features are observed, and only regions containing DFs have been used in order to reduce calibration errors. We find a coherent behavior of the Doppler velocity and the proper motion which clearly demonstrates that the evolution of DFs involve plasma motion. The Doppler velocities are found to be a factor 2--3 smaller than velocities derived form proper motions in the image plane. The difference can be explained by the radiative processes involved, the Doppler velocity is a result of the local atmospheric velocity weighted with the response function. As a result the Doppler velocity originates from a wide range in heights in the atmosphere. This is contrasted by the proper motion velocity which is measured from the sharply defined bright tops of the DFs and is therefore a very local velocity measure. The Doppler signal originates from well below the top of the DF. Finally we discuss how this difference together with the lacking spatial resolution of older observations have contributed to some of the confusion about the identity of DFs, spicules and mottles. ",
        "introduction": "\\label{sec:intro}  The solar chromosphere owes its name to the reddish rim that appears above the lunar limb during solar eclipses. This reddish color mostly stems from the Balmer {\\halpha} spectral line which makes this line one of the most important chromospheric diagnostics. Due to the highly dynamic state of the chromosphere and strong NLTE effects, the line formation processes are still not yet fully understood \\citep[e.g.][]{radyn,2006Leenaarts}. This is an important shortcoming in our interpretation tools which makes \\halpha\\ observations traditionally difficult to interpret.  Due to the highly fibrilar structure of the chromosphere \\citep{1908Hale}, a strong influence from magnetic fields on the chromosphere has been suspected for about a century.  The most common of these fibrilar magnetic fine structures are the jet-like structures known as spicules, mottles, and dynamic fibrils (DFs). In short, spicules are traditionally observed at the limb, mottles on disk in the quiet Sun, and DFs in active regions.  Whether or not these structures are manifestations of the same phenomenon viewed at different angles have been the subject of a long standing discussion \\citep[e.g.][]{1968Beckers,1992Grossman,1994Tsiropoula,1995Suematsu, 2001Christo,2007Rouppe}. One important argument against these structures being caused by the same mechanism has been the difference in the measured absolute velocities \\citep{1992Grossman}. Other authors have done direct measurements of mottles crossing the limb \\citep{2001Christo}. They also state that since both proper motions and Doppler motions are used in the comparisons, systematic errors are probably introduced. Such errors might also be amplified by the rather limited spatial resolution of some of the data sets used.  The detailed analysis of DFs has accelerated in recent years \\citep[e.g.][]{2004dePontieu,2006deWijn,2006Hansteen,2007dePontieu, 2007Lars} due to major advances in both observational techniques and simulation efforts.  One of the main conclusions from these studies is that DFs are driven by magneto-acoustic shocks caused by p-mode oscillations and convective flows leaking into the chromosphere.  In a recent study, \\citet[][from now on paper 1]{2007bLangangen} presented spectroscopic analysis of DFs as seen in one of the Ca~{\\small{II}}~IR lines. Numerical analysis of the line formation process showed a much lower DF velocity derived from Doppler measurements as compared to the proper motion velocity. This was found to be due to both the low formation height and the extensive width of the contribution function of the Ca~{\\small{II}}~IR line. Furthermore, the DFs analyzed in paper 1 showed mass motion, thus ruling out any ionization/temperature wave as explanation model for DFs \\citep[e.g.][and references therein]{2000Sterling}.  With the advantage of well sampled spectral line profiles, the number of analysed DFs in paper~1 was rather modest due to the limited spatial coverage of the spectrograph slit. With the current data we exploit the wide spatial coverage of a tunable filter instrument, at the expense of limited spectral resolution.  \\begin{figure*}[!ht] \\includegraphics[width=\\textwidth]{f1.eps} \\caption{ Field of view (FOV) for one wideband image (left) and the corresponding sum of the two narrow band filtergrams (right). In the narrow band image several fibrilar structures are seen, and some DF axes are marked (solid white lines) for illustrative purposes.  } \\label{plotone} \\end{figure*}  In this paper we add to the understanding of these jet like structures by analysis of Dopplergrams obtained in an active region close to the disk center, hence the observed jet structures are commonly known as DFs. In \\S\\ref{sec:obs} we describe the observational program and the instrumentation. The data reduction and the calibration method is explained in \\S\\ref{sec:datareduction}. In \\S\\ref{sec:Obsres} we present the results of our measurements. We discuss our results in \\S\\ref{Disc} and finally we summarize the results in \\S\\ref{Sum}. ",
        "conclusions": "\\label{Disc}  The correlation between the maximum velocities and decelerations found from the proper motion measurements is similar to the correlation found by \\citet{2006Hansteen} and \\citet{2007dePontieu}. This is illustrated in Fig.~\\ref{plotseven} by the grey-scaled cloud shown in the background. This correlation between the deceleration and the maximum velocity is known to be the signature of shock waves being the driving mechanism of DFs \\citep{2006Hansteen,2007dePontieu,2007Lars}.  Further support for this driving mechanism comes from the fact that we find a coherent behavior between the evolution of the Doppler signal and the proper motion for a large fraction of the DFs. This is a strong indication that there is actual plasma motion occurring during the life time of DFs. This supports the findings of paper~1, but based on a much larger sample.  The Doppler measurements show a similar correlation as for the proper motion, but with much lower absolute values for both the decelerations and maximum velocities. One possible explanation for these lower values could be high inclination angles of the DF trajectories with the LOS. One could expect to be able to derive the full trajectory vector by combining the two measured deceleration components. This naive method would give very high inclination angles, typically $75^\\circ$. We know, however, that this can not be the true inclination angle, since the Doppler velocity is a result of the local atmospheric velocity weighted with the response function to velocity over an extended height. In contrast, the measured proper motion is very local due to the high contrast boundary between the top of the fibril and the surroundings. Combining these two measurements leads to highly overestimated inclination angles. The difference in absolute values must be considered in the context of the results from Paper~1. The lower Doppler velocities found from the Ca~{\\small{II}}~IR line was explained by a combination of lower formation height and extended formation range. This is probably also the case for the \\halpha\\ line, but the formation height usually extends over a larger height range as compared to the Ca~{\\small{II}}~IR line.  The lacking Doppler shifts in ${\\sim}15$\\% of the DFs can either be caused by very high inclination angles, or their driving mechanism is fundamentally different and the evolution of these DFs is not a result of mass motion. We believe that high inclination angles combined with the uncertainties in the measurements is a more plausible explanation for the lacking Doppler signals. The identification method of DFs introduces a bias toward the more inclined DFs. There are a number of suggestive cases where DFs are visible in the Dopplergrams but no clear signature can be seen in the corresponding intensity images. We refrain from measuring these DFs since this will complicate a comparison with other data sets. Furthermore, the identification of these DFs is not objective and we expect the measurement errors to be unacceptably high since low inclination angles would lead to potentially high Doppler velocities. Due to the lacking spectral sampling this could lead to strong saturation effects in the measured Doppler velocities.  \\subsection{Spicules, mottles and fibrils}  The identification of the disk counterpart of spicules was already an important question forty years ago \\citep[e.g][]{1968Beckers}. One of the main problems was to reconcile the velocities measured in spicules with those measured in mottles. This problem was also the main concern of \\citet{1992Grossman}. They conclude that since the velocities are much larger in spicules than in mottles, the two could not be the same structure seen at different viewing angles. They, however, admit that the seeing might impair their results if the structures observed were smaller than $1\\arcsec{}$. Later studies of mottles and spicule properties lead to the conclusion that spicules and mottles are in fact the same feature seen at different angles \\citep{1993Tsiropoula,1994Tsiropoula}. \\citet{1994Tsiropoula} showed, using cloud modelling, that the proper motion and the cloud velocities were consistent. In a more recent work, \\citet{2001Christo} use a limb darkening correction method to directly observe mottles crossing the limb. They argue that the main reason for earlier confusion is caused by the lacking spatial resolution of the observations.  In our study it is clear that the Doppler signal originates from spatially resolved structures. The excellent quality of the observations largely removes the errors due to lacking spatial resolution. Assuming reasonable inclination angles, i.e. the mean angles being not very large nor very small, we can conclude that the Doppler velocities are typically a factor of $\\sim 2$--$3$ smaller than the corresponding proper motion. A similar, but larger difference was reported by \\citet{1994Tsiropoula}, we believe that this difference can be attributed to worse spatial resolution.  As discussed above, radiative transfer processes are the fundamental reason for the Doppler velocities being lower than the proper motion. We argue that the fundamental differences between Doppler and proper motion velocities that we find for DFs, also are valid for similar measurements on spicules and mottles. Besides the arguments of lacking spatial resolution, we believe that this difference was an important contributor to the earlier confusion about the unification of mottles and spicules. Similar work on spatially resolved limb spicules is needed to finally settle this discussion."
    },
    "0710/0710.4919_arXiv.txt": {
        "abstract": "Quantum gravity may remove classical space-time singularities and thus reveal what a universe at and before the big bang could be like. In loop quantum cosmology, an exactly solvable model is available which allows one to address precise dynamical coherent states and their evolution in such a setting. It is shown here that quantum fluctuations before the big bang are generically unrelated to those after the big bang. A reliable determination of pre-big bang quantum fluctuations would require exceedingly precise observations. ",
        "introduction": "Cosmology, as the physics of the universe as a whole, places special limitations on scientific statements to be reasonably inferred from observations. On a general basis, this was discussed in \\cite{Uncertain} who noted that some general properties seem to arise automatically during cosmological evolution. Their verification for our universe then does not reveal anything about its possibly more complicated past. For instance, isotropization \\cite{Isotropize} demonstrates that the observation of a nearly isotropic present universe does not present much useful information on an initial state at much earlier times. Similarly, one may add decoherence to the list which can be seen as doing the same for the observation of a nearly classical present universe: Most initial states would decohere and arrive at a semiclassical state; thus its observation does not rule out stronger quantum behavior at earlier stages. Given this, it is surprising to see recent discussions about a universe before the big bang in several approaches to quantum gravity. Not just statements about the possibility of the universe having existed before the big bang are being made, which could in principle be inferred from a general analysis of equations of motion, but even assumptions on its classicality or claims about the form of its state such as its fluctuations at those times are put forward. Even if this may be possible theoretically, it raises the question how much one can really learn about a pre-big bang universe from observations.  A solvable model of quantum cosmology is used here to draw cosmological conclusions within this model about the possible nature of the big bang and a universe preceding it. The viewpoint followed has been spelled out in \\cite{BeforeBB} where also some of the results have already been reported. This model does show a bounce at small volume instead of the classical singularity present in solutions of general relativity. We therefore use it to shed light on the general questions posed above. The promotion of this specific model is not intended as a statement that its bounce would be generic for quantum gravity even within the same framework. In fact, the conclusion of the bounce in this and related models available so far is based on several specific properties which prevent a generic statement about this form of singularity removal. We rather take the following viewpoint: Assume there is a theoretical description of a bouncing universe; what implications can be derived for its pre-bounce state?  The specific model used is distinguished by a key fact which makes it suitable for addressing such a general question: As a solvable model, the system displays properties of its dynamical coherent states which, to some degree, are comparable to those of the harmonic oscillator. There is no back-reaction of quantum fluctuations or higher moments of a state on the time evolution of its expectation values. Their trajectory is thus independent of the spreading of a state if it occurs, a property which is well known for the harmonic oscillator (for instance, as a simple consequence of the Ehrenfest theorem).  This behavior is certainly special compared to other quantum systems, but the precise derivation of properties of dynamical coherent states it allows are nevertheless important. For instance, a large part of theoretical quantum optics is devoted to a computation of fluctuations in squeezed coherent states of the harmonic oscillator. Similarly, the solvable model studied here allows detailed calculations of its dynamical coherent states which are of interest for quantum cosmology. In particular, the solvable model we will be using eliminates the classical big bang singularity by quantum effects. Dynamical coherent states highlight the behavior of fluctuations of the state of the universe before and after the big bang.  Care is, however, needed for the physical interpretation of the results. While quantum optics allows one to prepare a desired state and perform measurements on it, quantum cosmology has to make use of the one universe state that is given to us. Unlike quantum optics, where states can be prepared to be close to harmonic oscillator states, we cannot realize states of the solvable quantum cosmological model. The real universe is certainly very different from anything solvable, and so the availability of a certain feature or numerical result in a solvable model is unlikely to be related to a realistic property of the universe. Thus, as spelled out in \\cite{BeforeBB}, we focus on pessimism in our analysis of the solvable model: the {\\em in}ability of making certain predictions even in a fully controlled model is likely to be a reliable statement much more generally; adding complications of a real universe would only make those predictions even more difficult.  For solvable systems of this kind it is easier by far to solve equations of motion for expectation values and fluctuations directly, rather than taking the detour of a specifically represented wave function. Properties of coherent states are then determined by selecting solutions of fluctuations which saturate uncertainty relations. We will first illustrate this procedure for the harmonic oscillator with an emphasis on squeezed states. We present this brief review in section \\ref{Squeezed} intended as an introduction of the methods then used in quantum cosmology.  The main part of this paper is an application of those methods to the solvable system of quantum cosmology, as well as a physical discussion. Instead of different cycles of a harmonic oscillator, in quantum cosmology we will be dealing with the pre- and post-big bang phase of a universe. Just as fluctuations in a squeezed harmonic oscillator state can oscillate during the cycles, fluctuations in quantum cosmology may change from one phase to the next. Our main concern will be the reliability of predictions about the precise state of the universe before the big bang, based on knowledge we can achieve after the big bang. ",
        "conclusions": ""
    },
    "0710/0710.3364_arXiv.txt": {
        "abstract": "The focusing of the radiation generated by a polarization current with a superluminally rotating distribution pattern is of a higher order in the plane of rotation than in other directions. Consequently, our previously published asymptotic approximation to the value of this field outside the equatorial plane breaks down as the line of sight approaches a direction normal to the rotation axis, \\ie is nonuniform with respect to the polar angle. Here we employ an alternative asymptotic expansion to show that, though having a rate of decay with frequency ($\\mu$) that is by a factor of order $\\mu^{2/3}$ slower, the equatorial radiation field has the same dependence on distance as the nonspherically decaying component of the generated field in other directions: it, too, diminishes as the inverse square root of the distance from its source. We also briefly discuss the relevance of these results to the giant pulses received from pulsars: the focused, nonspherically decaying pulses that arise from a superluminal polarization current in a highly magnetized plasma have a power-law spectrum (\\ie a flux density $S\\propto\\mu^\\alpha$) whose index ($\\alpha$) is given by one of the values $-2/3$, $-2$, $-8/3$, or $-4$. ",
        "introduction": "Radiation by polarization currents whose distribution patterns move faster than light {\\it in vacuo} has been the subject of several theoretical and experimental studies in recent years \\cite{BessarabAV:FasEsi,ArdavanA:Exponr,BessarabAV:Expser,BolotovskiiBM:Radssv,% BolotovskiiBM:Radbcm,ArdavanH:Genfnd,ArdavanH:Speapc,ArdavanH:Morph}. When the motion of its source is accelerated, this radiation exhibits features that are not shared by any other known emission. In particular, the radiation from a rotating superluminal source consists, in certain directions, of a collection of subbeams whose azimuthal and polar widths narrow (as $R\\subP^{-3}$ and $R\\subP^{-1}$, respectively) with distance $R\\subP$ from the source \\cite{ArdavanH:Morph}. Being composed of tightly focused wave packets that are constantly dispersed and reconstructed out of other waves, these subbeams neither diffract nor decay in the same way as conventional radiation beams. The field strength within each subbeam diminishes as $R\\subP^{-1/2}$, instead of $R\\subP^{-1}$, with increasing $R\\subP$ \\cite{ArdavanH:Genfnd,ArdavanH:Speapc,ArdavanH:Morph}.  In earlier treatments \\cite{ArdavanH:Speapc,ArdavanH:Morph}, we evaluated the field of a superluminally rotating extended source by superposing the fields of its constituent volume elements, \\ie by convolving its density with the familiar Li\\'enard-Wiechert field of a rotating point source. This Li\\'enard-Wiechert field is described by an expression essentially identical to that which is encountered in the analysis of synchrotron radiation, except that its value at any given observation time receives contributions from more than one retarded time. The multivalued nature of the retarded time gives rise to the formation of caustics. The wave fronts emitted by each constituent volume element of a superluminally moving accelerated source possess a cusped envelope on which the field is infinitely strong (see Figs.~1 and 4 of Ref.~\\cite{ArdavanH:Morph}). Correspondingly, the Green's function for the problem is nonintegrably singular for those source elements that approach the observer along the radiation direction with the speed of light and zero acceleration at the retarded time (see Fig.~3 of Ref.~\\cite{ArdavanH:Morph}): the cusp of the envelope of wave fronts emanating from each such element is a spiraling curve extending into the far zone that passes through the position of the observer. When the source oscillates at the same time as rotating, the Hadamard finite part of the divergent integral that results from convolving the Green's function with the source density has a rapidly oscillating kernel for a far-field observation point. The stationary points of the phase of this kernel turn out to have different orders depending on whether the observer is located in or out of the equatorial plane.  To reduce the complications posed by the higher-order stationary points of this phase, we restricted the asymptotic evaluation of the radiation integral thus obtained in Refs.~\\cite{ArdavanH:Speapc,ArdavanH:Morph} to observation points outside the plane of rotation, \\ie to spherical polar angles $\\theta\\subP$ that do not equal $\\pi/2$. The purpose of this paper is to evaluate the field of a superluminally rotating extended source also for the smaller class of observers at polar coordinate $\\theta\\subP=\\pi/2$ and to obtain, thereby, a more global description of the nonspherically decaying radiation that is generated by such a source. The asymptotic expansion presented in Refs.~\\cite{ArdavanH:Speapc,ArdavanH:Morph} breaks down in the case of a subbeam that is perpendicular to the rotation axis because there is a higher-order focusing associated with the waves emitted by those source elements whose actual speeds (rather than the line-of-sight components of their speeds) equal the speed of light as they approach the observer with zero acceleration.  Here, we present a brief account of the background material on the radiation field of a rotating superluminal source in Section 2, and the asymptotic evaluation of this field for an equatorial observer in Section 3. In Section 4, we give a description of the spectral properties of the nonspherically decaying component of this radiation in the light of the present results and those obtained in Refs.~\\cite{ArdavanH:Speapc,ArdavanH:Morph}, and discuss the relevance of these properties to pulsar observations. ",
        "conclusions": ""
    },
    "0710/0710.2542_arXiv.txt": {
        "abstract": "{INTEGRAL, the European Space Agency's $\\gamma$-ray observatory, tripled the number of super-giant high-mass X-ray binaries (sgHMXB) known in the Galaxy by revealing absorbed and fast transient  (SFXT) systems.} {In these sources, quantitative constraints on the wind clumping of the massive stars could be obtained from the study of the hard X-ray variability of the compact accreting object.} {Hard X-ray flares and quiescent emission of SFXT systems have been characterized and used to derive wind clump parameters.} {A large fraction of the hard X-ray emission is emitted in the form of flares with a typical duration of 3 ks, frequency of 7 days and luminosity of $10^{36}$ erg/s. Such flares are most probably emitted by the interaction of a compact object orbiting at $\\sim10~R_*$ with wind clumps ($10^{22-23}$ g) representing a large fraction of the stellar mass-loss rate. The density ratio between the clumps and the inter-clump medium is  $10^{2-4}$ in SFXT systems. } {The parameters of the clumps and of the inter-clump medium, derived from the SFXT flaring behavior, are in good agreement with macro-clumping scenario and line driven instability simulations. SFXT have probably a larger orbital radius than classical sgHMXB.} ",
        "introduction": "\\label{sec:intro}  Stellar winds have profound implications for the evolution of massive stars, on the chemical evolution of the Universe and as source of energy and momentum in the interstellar medium.  Photons emitted by  massive stars directly transfer momentum to the stellar wind through absorption in numerous Doppler shifted optically thick spectral lines \\citep{cak1975} and drive the wind to highly supersonic speeds. Such line driven stellar winds are very unstable \\citep{ocr1988,feldmeier1995}. The collisions between high speed and low density material with slower gas trigger strong shocks, form high density contrasts and wind clumps and lead to thermal X-ray emission.  There are multiple observational lines of evidence for wind clumping in massive stars from wind line profiles \\citep{Bouret2005,Fullerton2006}, discrete absorption components \\citep{PrinjaHowarth1986}, variable line profiles \\citep{Lepine1999,Markova2005}, polarimetry \\citep{Lupie1987, Davies2007}, X-rays continuum \\citep{CassinelliOlson1979} and line emission \\citep{OskinovaFeldmeierHamann2006}.  Wind clumping has an important effect on mass-loss rate diagnostics that depends on the square of the wind density. It leads to a reduction of the stellar mass-loss rate estimates by a factor $f_V^{-0.5}$ where $f_V$ is the clump volume filling factor \\citep{Abbott1981,Fullerton2006}. A reduction by a factor $>3$ becomes problematic for massive star evolution \\citep{Hirschi2007}. Optically thick clumps may however help to reconcile wind clumping and usual mass-loss rates by reducing spectral features and free-free emission \\citep{OskinovaHamannFeldmeier2007}.  Indirect measures of the structure of massive star winds are possible in binary systems through the analysis of the interaction between the stellar wind and the companion or its stellar wind. In colliding wind binaries, \\cite{Schild2004} and \\cite{Pollock2005} have pointed out that the observed X-ray column densities are much lower than expected from smooth winds and that clumped winds are the probable solution \\citep[see also][]{pittard2007}. In wind-fed accreting High-Mass X-Ray Binaries (HMXB), clumping has been invoked to explain the orbital variability of X-ray line profiles emitted by the photo-ionized wind \\citep{Sako2003,Vandermeer2005}.  In this  paper we aim at studying the clumping of stellar wind in HMXB using the hard X-ray variability observed by the IBIS/ISGRI instrument \\citep{ubertini03AA,lebrun03AA} on board the INTErnational Gamma-Ray Astrophysics Laboratory \\citep[INTEGRAL,][]{winkler03AA} which could be used to probe clump parameters and density contrasts in the stellar wind. This study follows a first paper  \\citep{Leyder2007} interpreting the flaring behavior of \\object{IGR J08408$-$4503} in term of wind clumping.  Classical wind-fed super-giant HMXB (sgHMXB) are made of a compact object orbiting within few stellar radii from a super-giant companion. Recently INTEGRAL almost tripled the number of sgHMXB systems known in the Galaxy and revealed a much more complex picture with two additional families of sources: the highly absorbed persistent systems \\citep{walter04esasp,walter2006} and the super-giant fast X-ray transient (SFXT) systems \\citep{Negueruela2006}.  The highly absorbed systems have orbital and spin periods similar to those of classical Roche-lobe underflow sgHMXB, however the absorbing column densities are much higher than observed on average in classical systems \\citep{walter2006}. The fast transient systems are characterized by fast outbursts, by a low quiescent luminosity and by super-giant OB companions \\citep{Sguera2006,Negueruela2007}.  The sample of sources and the INTEGRAL data analysis are described in Sect. 2. Their hard X-ray variability is discussed in the context of clumpy stellar winds in Sect. 3. Finally Sect. 4 summarizes our principal conclusions. ",
        "conclusions": "INTEGRAL tripled the number of super-giant HMXB systems known in the Galaxy and revealed two new populations: the absorbed and the fast transient (SFXT) systems. The typical hard X-ray variability factor is $\\lesssim 20$ in classical and absorbed systems and $\\gtrsim 100$ in SFXT. We have also identified some ``intermediate'' systems with smaller variability factors that could be either SFXT or classical systems.  The SFXT behavior is best explained by the interaction between the accreting  compact object and a clumpy stellar wind \\citep{IntZand2005,Leyder2007}. Using the hard X-ray variability observed by INTEGRAL in a sample of SFXT we have derived typical wind clump parameters. The compact object orbital radius are probably relatively large ($10~R_*$) and the clumps which generate most of the hard X-ray emission have a size of a few tenth of $R_*$. The clump mass is of the order or $10^{22-23}~\\rm{g}$ (for a column density of $10^{22-23}~\\rm{cm}^{-2}$) and the corresponding mass-loss rate is $10^{-(5-6)}~\\rm{M_{\\odot}/y}$. At the orbital radius, the clump separation is of the order of $R_*$ and their volume filling factor is $0.02$. Depending how the clump density varies with radius, the average volume filling factor could be as large as 0.1.  These parameters are in good agreement with the macro clumping scenario proposed by \\cite{OskinovaHamannFeldmeier2007}.  The observed ratio between the flare and quiescent count rates indicate density ratios between the clumps and the inter-clump medium which vary between 15 to 50 in ``Intermediate'' systems and $10^{2-4}$ in SFXT. Such ratios and the observed clump densities are in reasonable agreement with the predictions of line driven instabilities at large radii \\citep{Runacres2005}.  The main difference between classical sgHMXB and SFXT could be the orbital radius of the compact object. At small orbital radius ($R_{orb}\\sim2~R_*$) the systems are persistent and luminous. At larger radius and if wind clumping takes place the fast transient SFXT behavior is observed."
    },
    "0710/0710.0547_arXiv.txt": {
        "abstract": "We present new $H$-band echelle spectra, obtained with the NIRSPEC spectrograph at Keck II, for the massive star cluster ``B'' in the nearby dwarf irregular galaxy NGC~1569. From spectral synthesis and equivalent width measurements we obtain abundances and abundance patterns. We derive an Fe abundance of [Fe/H]=$-0.63\\pm0.08$, a super-solar [$\\alpha$/Fe] abundance ratio of $+0.31\\pm0.09$, and an O abundance of [O/H]=$-0.29\\pm0.07$. We also measure a low  $\\rm ^{12}C/^{13}C\\approx 5\\pm1$ isotopic ratio. Using archival imaging from the Advanced Camera for Surveys on board HST, we construct a colour-magnitude diagram (CMD) for the cluster in which we identify about 60 red supergiant (RSG) stars, consistent with the strong RSG features seen in the $H$-band spectrum. The mean effective temperature of these RSGs, derived from their observed colours and weighted by their estimated $H$-band luminosities, is 3790 K, in excellent agreement with our spectroscopic estimate of $T_{\\rm eff} = 3800\\pm200$ K. From the CMD we derive an age of 15--25 Myr, slightly older than previous estimates based on integrated broad-band colours. We derive a radial velocity of $\\langle v_r \\rangle=-78\\pm 3$ km/s and a velocity dispersion of $9.6\\pm0.3$ km/s. In combination with an estimate of the half-light radius of $0\\farcs20\\pm0\\farcs05$ from the HST data, this leads to a dynamical mass of $(4.4\\pm1.1)\\times10^5$ M$_\\odot$.  The dynamical mass agrees very well with the mass predicted by simple stellar population models for a cluster of this age and luminosity, assuming a normal stellar IMF. The cluster core radius appears smaller at longer wavelengths, as has previously been found in other extragalactic young star clusters. ",
        "introduction": "Massive star clusters are potentially useful test particles for constraining the evolutionary histories of their host galaxies. They can remain observable for the entire lifetime of a galaxy (as illustrated by the old \\emph{globular clusters} which surround every major galaxy), and they are bright enough to be studied in detail well beyond the Local Group. The internal velocity broadening of their spectra is typically only a few km/s, making detailed abundance analysis at high spectral resolution feasible. In a previous paper we have taken a first step towards exploiting this potential by analysing $H$ and $K$-band spectra of a young massive star cluster (YMC) in the nearby spiral galaxy NGC~6946 \\citep{lar06}. Here we apply a similar analysis to one of the young star clusters in the nearby (post-) starburst galaxy NGC~1569.  NGC~1569 was one of the first galaxies in which the presence of exceptionally bright, young star clusters was suspected. A spectrum of one of the two bright ``stellar condensations'' in NGC~1569 was obtained already by \\citet{mayall35}, although the true nature of these objects was probably first discussed in detail by \\citet{as85}.  They found that the spectra are of composite nature, and also noted that observations with the Hubble Space Telescope (HST) would definitively settle the issue whether or not these objects are really ``super star clusters''.  Indeed, pre-refurbishment mission HST observations by \\citet{oco94} settled the issue by showing that both objects are extended, with half-light radii of about two pc. Based on high-dispersion spectroscopy from the HIRES spectrograph on the Keck~I telescope, a dynamical mass of $\\approx 10^6$ M$_\\odot$ was soon after estimated for the brighter of the two objects, NGC~1569-A \\citep{hf96,stern98}.  The clusters in NGC~1569 are prime candidates for abundance analysis in the near-infrared where they are very bright. The IR holds a particular advantage over optical studies for NGC~1569 due to the large amount of foreground extinction. \\citet{gg02} found the integrated $H$-band spectra of both clusters A and B to be well approximated by red supergiant templates.  Although NGC~1569-A is the brighter of the two, it is actually a binary cluster itself with some evidence for a (small) age difference between the two components \\citep{guido97,maoz01,origlia2001}. Whether or not the two components are physically connected or the result of a chance projection is unknown.  In this paper we concentrate on NGC~1569-B.  In addition to providing insight into the histories of their host galaxies, young star clusters with masses in excess of $10^5$ M$_\\odot$ also offer an excellent opportunity to study large samples of coeval massive stars. Such clusters are rare in the Milky Way and even in the Local Group, so it is necessary to extend the search to a larger volume.  The most massive known young star clusters in the Milky Way disk are Westerlund 1 \\citep{clark05} and an object discovered in the 2MASS survey \\citep{figer06}, both of which may have masses approaching $10^5$ M$_\\odot$. Young clusters of similar masses are found in the Large Magellanic Cloud \\citep{vdb99}, but these pale in comparison with objects like NGC~1569-A and NGC~1569-B.  We have obtained new NIRSPEC $H$-band spectra of NGC~1569-B, optimised for abundance analysis, and we additionally present photometry for \\emph{individual} stars in the cluster derived from archival observations with the high resolution channel (HRC) of the Advanced Camera for Surveys (ACS) on HST.  The HST data provide an independent verification of the stellar parameters (notably $T_{\\rm eff}$ and $\\log g$) derived from the spectral analysis, and also allow us to compare the observed colour-magnitude diagram (CMD) with standard isochrones.  Finally, we use new measurements of the structural parameters and velocity dispersion to derive the cluster mass and mass-to-light ratio and compare with predictions by simple stellar population (SSP) models.  We begin by briefly describing the data in \\S\\ref{sec:data}. The analysis and our main results then follow in \\S\\ref{sec:results}, where we first discuss the near-infrared spectroscopy and the abundance analysis (\\S\\ref{sec:abundance}). We then proceed to construct a CMD from the HST imaging (\\S\\ref{sec:cmd}) from which we derive stellar parameters for the red supergiants (RSGs) that are compared against the spectroscopic results (\\S\\ref{sec:speccmp}) and theoretical isochrones (\\S\\ref{sec:interp}). In \\S\\ref{sec:virmass} we measure structural parameters for NGC~1569-B and combine these with velocity dispersion measurements to derive a dynamical mass estimate. Finally, some additional discussion and a summary are given in \\S\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  We have presented a detailed investigation of the stellar content and other properties of the massive star cluster 'B' in NGC~1569. Using new high S/N $H$-band echelle spectroscopy from the NIRSPEC spectrograph on Keck~II, we have carried out abundance analysis of red supergiant stars in the cluster. We find an iron abundance of [Fe/H] = $-0.63\\pm0.08$, close to that of SMC field stars. Our estimate of the oxygen abundance, [O/H] = $-0.29\\pm 0.07$, is about a factor of two higher than that derived for H{\\sc ii} regions \\citep{sdk94,dev97,ks97}, and the resulting super-solar [O/Fe] abundance is unlike that observed in the Magellanic Clouds and young stellar populations in the Milky Way. However, according to our measurements NGC~1569-B also differs from these galaxies by having a higher [$\\alpha$/Fe] ratio.  The difference between our [O/H] measurement for NGC~1569-B and the O abundance derived for H{\\sc ii} regions is greater than the formal uncertainties on either value, and our oxygen abundance is higher than the nebular abundances in most dwarf galaxies \\citep[e.g.][]{vad07}. This raises the question whether one (or both) methods suffers from possible systematic errors, or if there might be a genuine difference between cluster and H{\\sc ii} region abundances in NGC~1569 -- perhaps due to an ab initio oxygen enhancement of the gas out of which the cluster formed. Since oxygen is the only element in common between the cluster and H{\\sc ii} region data, a detailed comparison of the abundance patterns is unfortunately not possible. However, it is of interest to note that the very strong Wolf-Rayet features in the spectrum of cluster A also suggest a high metallicity \\citep{maoz01}.  NGC~1569 is a popular target for testing models of chemical evolution, and it is generally recognised that the strong galactic wind observed in the galaxy must play an important role. \\citet{martin02} observed a super-solar $[\\alpha/{\\rm Fe}]$ ratio for the wind, as we do for NGC~1569-B. However, it is unclear to what extent the two are related as material ejected in the outflow may not participate in star formation, and in any case not before it has had time to cool down and fall back into the galaxy. The chemical evolution of an NGC~1569-type galaxy has been modelled in detail by \\citet{rec06}, assuming a variety of bursty star formation histories. Their chemo-dynamical simulations show that elements produced in the last burst of star formation generally do not get mixed with the ISM in the galaxy, but instead get injected into the hot gas phase. None of their models produce an O abundance as high as the one observed by us, and the model that comes closest ($12 + \\log$ (O/H) $\\sim 8.4$) severely underpredicts the N/O abundance of the H{\\sc ii} regions observed by \\citet{ks97}. Generally, the best fits to the H{\\sc ii} region abundances are obtained for a ``gasping'' star formation history, but these models all predict a maximum $12 + \\log$ (O/H) $\\approx$ 8.1. Similarly, the O abundance predicted by the models of \\citet{romano06} for NGC~1569 never exceeds $12 + \\log$ (O/H) = 8.41. We speculate that the galactic wind which is responsible for removing metals may be less efficient in doing so near the bottom of the potential well where the massive clusters formed. Unfortunately, none of the models include predictions for the wide range of other elements we are observing here.  The Small Magellanic Cloud presents an interesting comparison case. Early spectroscopic studies and Str{\\\"o}mgren photometry indicated that the young cluster NGC~330 is about 0.5 dex more metal-poor than the surrounding field \\citep[and references therein]{gr92}. High-dispersion spectroscopy by \\citet{hill99} showed a smaller and only marginally significant difference between the field stars and NGC~330, although still in the sense that the cluster is more metal-poor than the field ([Fe/H] = $-0.82\\pm0.10$ vs.\\ [Fe/H] = $-0.69\\pm0.11$). The [O/Fe] abundance ratios in these stars were all found to be sub-solar ([O/Fe] $\\approx$ $-$0.15 to $-$0.3 dex), as in the LMC and in young Galactic supergiants \\citep{hbs97}.  \\citet{hill99} also found the $\\alpha$-elements abundances (Mg, Ca, Ti) relative to Fe to be around Solar for NGC~330.  \\citet{gw99} found slightly lower iron abundances of [Fe/H] = $-0.94\\pm0.02$ for K supergiants in NGC~330 than \\citet{hill99} and an [O/Fe] ratio closer to Solar. They derived somewhat enhanced  [Ca/Fe], [Si/Fe] and [Mg/Fe] ratios ($+0.18$, $+0.32$ and $+0.11$), in contrast to sub-solar ratios derived for B stars.  While there is some evidence for a difference in metallicity between NGC~330 and the SMC field, the situation there seems to contrast with the case of NGC~1569-B where we find a \\emph{higher} oxygen abundance of the cluster compared to the H{\\sc ii} regions and an \\emph{enhanced} [O/Fe] ratio. Our mean [$\\alpha$/Fe] = $+0.31\\pm0.09$ is also higher than that derived for the stars in NGC~330. One significant difference may be the mass of NGC~1569-B, which is probably an order of magnitude greater than that of NGC~330 \\citep{fb80}.   The spectral analysis returns a mean $T_{\\rm eff} = 3800\\pm200$ K and $\\log g \\approx 0.0$ for the RSGs in NGC~1569-B. We have checked these values using resolved photometry of the RSGs in the cluster, derived from archival HST/ACS images. About 60 RSGs are easily identifiable in the colour-magnitude diagram, and from their $m_{\\rm F555W} - m_{\\rm F814W}$ colours we derive a mean effective temperature of $\\langle T_{\\rm eff}\\rangle = 3850$ K and $\\log g = 0.1$. An even closer match to the spectroscopic $T_{\\rm eff}$ is obtained if we weigh the $T_{\\rm eff}$ estimates for the individual stars by their $H$-band luminosities, in which case we get $\\langle T_{\\rm eff} \\rangle = 3790$ K.  We have compared the CMD with $Z=0.004$ and $Z=0.008$ isochrones from the Padua and Geneva groups.  Both sets of models provide the best match to the observed colours and magnitudes of the RSGs for $Z=0.008$ and ages of 15--25 Myrs, but no isochrone can reproduce the observed CMD in detail. For these sub-solar metallicities, the models predict higher (by up to a few 100 K) mean effective temperatures $\\langle T_{\\rm eff} \\rangle$ for the RSGs than observed.  Since the RSGs generally become cooler with increasing metallicity, models of higher metallicity are favoured, but models of Solar metallicity would be required to match the observed colours. For such models, however, the \\emph{blue} supergiants become much too cool to match the observations. The observed ratio of blue to red supergiants (BSG/RSG = $0.39\\pm0.10$) is also significantly lower than most of the model predictions.  Although we are detecting an impressive number of RSGs, it should be noted that since we are only resolving stars located well outside the half-light radius, the actual number of RSGs in NGC~1569-B is likely to be more than a factor of two greater than the number we have detected.  We derive a velocity dispersion of $9.6\\pm0.3$ km/s from the integrated spectrum of NGC~1569-B, somewhat higher than the value of 7.5 km/s of \\citet{gg02}. Combining this with an estimate of the cluster half-light radius of $0\\farcs20\\pm0\\farcs05$ measured on the ACS images, we obtain a dynamical mass of $(4.4\\pm1.1)\\times10^5 \\left(\\frac{D}{2.2 {\\rm Mpc}}\\right) M_\\odot$. This is in excellent agreement with the mass predicted by simple stellar population models for a cluster of this age and luminosity and a standard IMF.  A correction for mass segregation would bring the two estimates into even closer agreement.  Our analysis of structural parameters reveals a decrease in core radius with wavelength for NGC~1569-B. This trend is consistent with the colour gradients observed by \\citet{hunter00} and by us in the HRC data when variations in the PSF are taken into account. A similar colour gradient was observed in an YMC in NGC~6946 \\citep{laretal01}.  The correlation between core radius and wavelength is also in the same sense as that found for M82-F by \\citet{mgv05} and may be an indication that mass segregation (primordial or dynamical) is present in NGC~1569-B. However, this needs to be confirmed by a more detailed analysis, including a modelling of how mass segregation would translate into observables such as core- and half-light radius.  As a final note, we find it fascinating to contemplate that only a couple of decades ago, it was still uncertain whether NGC~1569-A and NGC~1569-B were in fact star clusters in NGC~1569 or mere foreground stars. Today, the Advanced Camera for Surveys on HST has made it possible to not only settle this question definitively, but to study the individual stars that make up these clusters."
    },
    "0710/0710.0637_arXiv.txt": {
        "abstract": "We combine high-resolution images in four optical/infra-red bands, obtained with the laser guide star adaptive optics system on the Keck Telescope and with the Hubble Space Telescope, to study the gravitational lens system \\lensname (lens redshift~\\zdvalue, source redshift~\\zsvalue). We show that (under favorable observing conditions) ground-based images are comparable to those obtained with HST in terms of precision in the determination of the parameters of both the lens mass distribution and the background source. We also quantify the systematic errors associated with both the incomplete knowledge of the PSF, and the uncertain process of lens galaxy light removal, and find that similar accuracy can be achieved with Keck LGSAO as with HST. We then exploit this well-calibrated combination of optical and gravitational telescopes to perform a multi-wavelength study of the source galaxy at $0\\farcs01$ effective resolution.  We find the \\sersic index to be indicative of a disk-like object, but the measured half-light radius (\\re=\\revalue) and stellar mass (\\stellarmass=\\stellarmassvalue) place it more than three sigma away from the local disk size-mass relation. The \\lensname source has the characteristics of the most compact faint blue galaxies studied, and has comparable size and mass to dwarf early-type galaxies in the local universe. With the aid of gravitational telescopes to measure individual objects' brightness profiles to 10\\% accuracy, the study of the high-redshift size-mass relation may be extended by an order of magnitude or more beyond existing surveys at the low-mass end, thus providing a new observational test of galaxy formation models. ",
        "introduction": "\\label{sect:intro}  Galaxies do not appear in arbitrary combinations of luminosity, mass and shape, but instead obey empirical scaling relations (such as the Fundamental Plane for early-type galaxies). Explaining the origin, and cosmic evolution, of the scaling relations is a fundamental goal of galaxy formation theories.  As far as disk galaxies are concerned, the hierarchical structure formation scenario predicts a correlation between size and stellar mass, with width depending on the distribution of the initial spin of the dark halos \\citep{F+E80}. At any given mass, the expected distribution of sizes is well-approximated by a log-normal distribution.  Qualitatively, this prediction is quite robust, although the exact forms of the correlation and the distribution depend on the details of baryonic processes such as energy feedback from star formation and bulge instability \\citep{MMW98,She++03,Ton++06,Dut++07,S+B07}. Therefore, measuring the shape and width of the correlation provides not only a test of the standard paradigm, but also valuable information on the poorly-understood baryonic processes happening at sub-galactic scales.  From an empirical point of view, the relation between size, luminosity (or equivalently surface brightness) and stellar mass is well established for disk galaxies in the local Universe \\citep[\\eg][]{She++03,Dri++05}. Analysis of suitable objects in the Sloan Digital Sky Survey shows that at any given mass (luminosity) the distribution of galaxies is indeed well-approximated as log-normal, although the scaling with mass of the characteristic size and the width of the distribution are non-trivial. Defining disk galaxies as those being well-fit by a single \\sersic component with index $n<2.5$, \\citet{She++03} find that above a characteristic stellar mass ($\\log M_{*,0}/\\msun\\sim 10.6$ corresponding to approximately M$_{r,0}=-20.5$), size scales rapidly with stellar mass ($R\\sim M_*^{0.39}$) and the scatter is relatively small ($\\sigma_{\\ln R}\\sim 0.34$). Below the characteristic stellar mass the correlation flattens ($R\\sim M_*^{0.14}$) and the scatter increases significantly ($\\sigma_{\\ln R}\\sim 0.47$).   At intermediate redshift ($0.1\\lsim z \\lsim 1$) the nature and interpretation of the size-luminosity or size-mass relation is more uncertain. Several authors \\citep[\\eg][]{Fer++04,Bar++05,Tru++06,Mel++06} have used Hubble Space Telescope images to determine the sizes of intermediate and high ($z \\gsim 1$) redshift galaxies, down to the resolution and completeness limits of HST (roughly equivalent to 1~kpc and 10$^{10} \\msun$). Recent studies conclude, taking selection effects into account, that there is significant evolution in the size-luminosity relation \\citep{Bar++05,Tru++06,Mel++06}. However, it is hard to disentangle luminosity evolution from size evolution, to ensure that samples at different redshifts are directly comparable, and to compare results from different studies, as the selection criteria are often similar but not identical (\\eg color vs. morphology; morphology determined via \\sersic index vs.\\ bulge to disk decomposition vs.\\ concentration parameter vs. visual classification). Overall, it appears that disk galaxy evolution cannot be explained by pure luminosity or pure size evolution, but requires a combination of both. In contrast, the relation between size and stellar mass appears to have changed very little since $z\\sim1$ \\citep{Bar++05}, much less than would be expected in the naive model where stellar mass and size are proportional to the virial mass and radius (and hence size is expected to scale as $H(z)^{-\\frac{2}{3}}$, where $H(z)$ is the Hubble parameter). Rather, galaxies appear to be growing ``inside-out'' in scale radius as their stellar mass increases such that the size-mass relation is preserved over cosmic time \\citep{Bar++05}. It has been suggested that galaxy evolution models that take into account the ever-increasing concentration of dark matter halos, and the further effect of baryons via adiabatic contraction could provide the physics required to reproduce the observed trend \\citep{Som++06}, although this may make it more difficult to reproduce simultaneously other scaling laws, for example the Tully-Fisher \\citep{T+F77} relation \\citep{Dut++07}.  Lower mass ($M_* \\lsim 10^{10} \\msun$) galaxies are even less well understood. While the local size-mass relations of \\citet{She++03} for low ($n<2.5$) and high ($n>2.5$) \\sersic index objects diverge, the interpretation of \\sersic index as a morphological galaxy classifier becomes more uncertain at lower masses \\citep[e.g.\\ ][]{CCD92,TBB04}. At the same time, the measurement of the structural parameters themselves becomes harder as the galaxy size decreases. Nevertheless, such small galaxies are important objects to understand: the luminous compact blue galaxies first noted by \\citet{K+K88} appear in large numbers at intermediate redshifts in deep HST images \\citep[\\eg][]{Noe++06,Raw++07}, but evolve very rapidly to vanishing abundance in the local universe. What becomes of these objects, which represent sites of small-scale but vigorous star formation, is a topic of some debate, with dwarf spheroids \\citep[\\eg][]{Koo++94,Noe++06} and the bulges of disk galaxies \\citep[\\eg][]{Ham++01,Raw++07} the principle candidates.  Gravitational lensing is a powerful tool with which to extend the investigation of scaling laws over cosmic time \\citep[\\eg][]{Tre07}. On the one hand, the lensing geometry provides a precise and almost model-independent measure of total mass of the lens galaxy. Since the lens galaxies are mostly early-type galaxies (or the bulges of spirals), this gives a new handle on the mass profile of these systems \\citep{T+K04,Koo++06} and hence, for example, on the relationship between stellar and total mass \\citep{Bol++07}. On the other hand, the background source is typically magnified by a factor of $\\sim$10, mostly in the form of a stretch along the azimuthal direction.  While lensing preserves surface brightness, the increase in apparent size of the lensed source means that the number of pixels at any one surface brightness also increases, such that the isophotes are observed at higher signal-to-noise. Thus, gravitational lenses act as natural telescopes, allowing one to gain a factor of $\\sim10$ in sensitivity and spatial resolution, and thus improve markedly our ability to study the size and dynamical mass (through rotation curves) of intermediate and high redshift galaxies. For example, studies of the internal structure of faint blue galaxies~\\citep{Ell97}, and in particular the most compact of these \\citep{Koo++94}, are currently limited by the resolution of HST \\citep{Phi++97}. When magnified by a gravitational lens, such objects become well-resolved. Thanks to the dedicated efforts of several groups, the number of known gravitational lenses is increasing dramatically: it is now possible to envision statistical studies of relatively large sample of lensing or lensed galaxies in the near future.  In this paper we present multi-color high-resolution images of the gravitational lens system \\lensname \\citep{Bol++06}, obtained with both the Hubble Space Telescope and with the Laser Guide Star Adaptive Optics (LGSAO) System on Keck II. The scientific goal of the analysis of this case study is two-fold. First, we perform a detailed comparison of the results of the lens modeling across bands, showing that -- when a bright nearby star is available for tip-tilt correction and conditions are favorable -- the most important parameters can be measured with comparable accuracy with HST and Keck-LGSAO. Second, we exploit this particular cosmic telescope to achieve super-resolution of the source galaxy. \\citep[See][ for Keck LGSAO observations of a lens with a point-like source.]{McK++07} With a lens magnification of $\\mu \\gsim 10$, the resolution of the HST and Keck images ($\\sim 0\\farcs1$ FWHM) corresponds to a physical scale of ($0.66{\\rm kpc} / \\mu \\approx 0.05{\\rm kpc}$) at the redshift of the source {\\zs~=~\\zsvalue}, comparable to the resolution attainable from the ground when studying galaxies in the Virgo Cluster in $1$~arcsec seeing. We derive the \\sersic index, size, and stellar mass of the source, and show that using gravitational telescopes the size-mass relation may be extended by an order of magnitude in size with respect to current studies, thus allowing one to probe, for example, whether the change in slope and intrinsic scatter below the characteristic mass persists to higher redshifts.  This paper is organized as follows. After describing the observations in section~\\ref{sect:obs}, we outline in sections~\\ref{sect:psf} and~\\ref{sect:subtraction} two sources of systematic error and our strategies for dealing with them, before explaining our modeling methodology in section~\\ref{sect:modelmethod}. In sections~\\ref{sect:modelresults} and ~\\ref{sect:source} we present our results, which are then discussed (section~\\ref{sect:discuss}) before we draw conclusions in section~\\ref{sect:conclude}. Throughout this paper magnitudes are given in the AB system. We assume a concordance cosmology with matter and dark energy density $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$, and Hubble constant H$_0$=70 kms$^{-1}$Mpc$^{-1}$. ",
        "conclusions": "\\label{sect:conclude}  We find that high quality images from \\nirc are capable of providing very similar precision on simple lens and source model parameters to typical datasets from \\acs and \\nicmos. The data themselves contain information about the most appropriate PSF model to use, to the extent that a set of nearby unsaturated stars can  be fruitfully compared using suitable statistics that are  sensitive to the goodness-of-fit. We estimate that even for the LGSAO imaging  this way of modeling the PSF allows a photometric precision of 0.05 mag. However, the calibration of isothermal However, the calibration of \\emph{isothermal} galaxy-scale gravitational lenses as cosmic telescopes is very likely limited by the subtraction of the lens galaxy light. We estimate that this procedure introduces up to 0.1 magnitudes of systematic error into the source galaxy photometry. However, this is still smaller than the error introduced by the assumption of an isothermal density profile for the lens itself.  With this in mind we draw the following conclusions about the source behind \\lensname:  \\begin{itemize}  \\item Our photometry is robust enough to permit a reconstruction of the SED, and we find a stellar mass of (\\stellarmassvalue). This is a factor of 5 smaller than the completeness limit of the GEMS disk galaxy analysis of \\citet{Bar++05}, and also smaller than the least massive spheroid at this redshift studied by~\\citep{McI++05}.  \\item The \\sersic profile parameters of the source can be measured to high accuracy. We find an effective radius of (\\revalue) ($\\approx 0.09$ arcsec with $\\sim10$\\% accuracy), and a \\sersic index of (\\bluesindexvalue) in the \\Iband ($\\sim$ rest-frame~B), and that these values change little over the rest-frame optical range.  \\item This very small galaxy lies approximately 3-sigma below the local size-mass relation for disks. However, it shares the properties of the smallest of the compact narrow emission line galaxies of \\citet{Koo++94}, and, despite its low \\sersic index, is more typical of the dwarf early-type galaxies observed in the Virgo cluster~\\citep{Fer++06} and the ``elliptical'' galaxies studied by~\\citet{McI++05} at high redshift.  \\end{itemize}  While the planned  statistical analysis of a large sample of lensed galaxies will rely on the detailed understanding of the selection function, it is clear that the magnifying effect of gravitational lenses allows us to extend current size-mass studies to smaller sizes and lower masses than would otherwise be available, posing fresh challenges to models of galaxy formation and evolution."
    },
    "0710/0710.2318_arXiv.txt": {
        "abstract": "% Proposed mechanisms for the formation of km-sized solid planetesimals face long-standing difficulties.  Robust sticking mechanisms that would produce planetesimals by coagulation alone remain elusive.  The gravitational collapse of smaller solids into planetesimals is opposed by stirring from turbulent gas.   This proceeding describes recent works showing that ``particle feedback,\" the back-reaction of drag forces on the gas in protoplanetary disks, promotes particle clumping as seeds for gravitational collapse.  The idealized streaming instability demonstrates the basic ability of feedback to generate particle overdensities.  More detailed numerical simulations show that the particle overdensities produced in turbulent flows trigger gravitational collapse to planetesimals.  We discuss surprising aspects of this work, including the large (super-Ceres) mass of the collapsing bound cluster, and the finding that MHD turbulence aids gravitational collapse. ",
        "introduction": "Coagulation is the dominant mechanism for the growth of dust grains via van der Waals forces  \\citep{dt97} and for the growth of solid protoplanets by gravitational binding \\citep{gls04}.  But sticking is difficult in the ~mm--km size range, as confirmed by an extensive body of experimental work \\citep{wb06}.  Moreover, incremental growth of planetesimals leads to the rapid ($\\sim10^2$ yr)  inspiral of particles near a meter in size.  An alternative hypothesis \\citep[hereafter GW]{saf69, gw73} proposes that km-sized planetesimals formed from the gravitational collapse of smaller solids.  This mechanism overcomes the sticking and radial drift obstacles in one fell swoop.  However stirring by turbulent gas opposes gravitational collapse. GW noted that their dense particle midplane would trigger a turbulent boundary layer via Kelvin-Helmholtz instabilities.    \\citet{sw80} argued that this particle-driven turbulence (an example of drag force feedback) would stir up the particle midplane enough to prevent gravitational collapse.  This appeared to be a fatal flaw of the GI hypothesis, since particle settling was a self-limiting process.  \\citet{sek98} and \\citet{ys02} showed that particles could actually help their cause of becoming planetesimals.  When the surface density  of solids relative to gas is larger (by factors of few) than solar abundances, then vertical shear is no longer strong enough to overcome the anti-buoyancy of the dense midplane layer.  The ability of particles to stir themselves is limited.  However these works did not model the back reaction of drag forces on the gas in detail.  Indeed most analyses of midplane Kelvin-Helmholtz instabilities \\citep[with the exception of][]{jhk06} reduce the particle and gas dynamics to a simplified set of equations which omit relative motion and prevent a detailed examination of the effects of drag forces.  Section \\ref{sec:SI} describes the streaming instability \\citep[hereafter YG]{yg05} which takes the opposite approach of neglecting stratification and vertical shear to consider two-way drag forces in a self-consistent, if incomplete, model.  This idealized system shows that particle feedback triggers spontaneous particle clumping. The detailed 3D simulations of  \\citet[hereafter JOMKHY]{nature07} include vertical stratification, self-gravity, MHD turbulence and multiple particle sizes, see \\S\\ref{sec:KS}  These models show that clumps of 15 -- 60 cm boulders, augmented by feedback effects, collapse gravitationally into bound clusters, which should continue to contract into planetesimals. ",
        "conclusions": "The long-standing mystery of planetesimal formation suddenly appears much less daunting, if no less fascinating.  For decades, the obstacles seemed insurmountable: low sticking efficiencies, rapid radial migration and the inevitability of turbulent stirring.  Now it is clear that turbulent stirring does not necessarily prevent gravitational collapse of solid particles into planetesimals.  As discussed in \\S\\ref{sec:surprise}, more turbulence can promote gravitational collapse in some cases!  The breakthoughs in particle-gas dynamics are related to the role of particle feedback.  Once thought only to hinder particle settling by triggering Kelvin-Helmholz instabilities, we now realize that feedback triggers and augments particle clumping.  The link between feedback and clumping has  many pillars of support: the idealized streaming instability (YG, YJ, JY), Kelvin-Helmholz instabilities with uniform rotation \\citep{jhk06}, and now 3D stratified models of Keplerian disk midplanes (JOMKHY and its supplement).  Much work and many uncertainties remain.  The large initial sizes assumed by JOMKHY might be unrealistic, either because coagulation stalls at smaller sizes, or if less violent gravitational collapse occurs first (see work on dissipative gravitational collapse by \\citealp{war76,y05a}).  The fact that most primitive undifferentiated meteorites betray no structures larger than mm-sized chondrules suggest this concern may be valid.  However, particles which are small and tightly coupled to the gas disk are less effected by the dynamical clumping mechanisms discussed here, at least at the resolutions studied to date.  Perhaps more surprises are needed."
    },
    "0710/0710.1631_arXiv.txt": {
        "abstract": "We present a comprehensive study of accretion activity in the most underdense environments in the universe, the voids, based on the SDSS DR2 data.  Based on investigations of multiple void regions, we show that Active Galactic Nuclei (AGN) are definitely common in voids, but that their occurrence rate and properties differ from those in walls. AGN are more common in voids than in walls, but only among moderately luminous and massive galaxies ($M_r < -20$, log $M_*/M_{\\sun} < 10.5$), and this enhancement is more pronounced for the relatively weak accreting systems (i.e., $L_{\\rm [O III]} < 10^{39}$ erg s$^{-1}$).  Void AGN hosted by moderately massive and luminous galaxies are accreting at equal or lower rates than their wall counterparts, show lower levels of obscuration than in walls, and similarly aged stellar populations.  The very few void AGN in massive bright hosts accrete more strongly, are more obscured, and are associated with younger stellar emission than wall AGN.  These trends suggest that the accretion strength is connected to the availability of fuel supply, and that accretion and star-formation co-evolve and rely on the same source of fuel.  Nearest neighbor statistics indicate that the weak accretion activity (LINER-like) usually detected in massive systems is not influenced by the local environment.  However, H {\\sc ii}s, Seyferts, and Transition objects are preferentially found among more grouped small scale structures, indicating that their activity is influenced by the rate at which galaxies interact with each other. These trends support a potential H{\\sc~ii}$\\rightarrow$Seyfert/Transition Object$\\rightarrow$LINER evolutionary sequence that we show is apparent in many properties of actively line-emitting galaxies, in both voids and walls. The subtle differences between void and wall AGN might be explained by a longer, less disturbed duty cycle of these systems in voids. ",
        "introduction": "The regions that are apparently devoid of galaxies \\citep{kir81} and clusters \\citep{ein80}, the voids, are arguably the best probes of the effect of the environment and cosmology on galaxy formation and evolution.  If, as suggested by the standard cosmological paradigm, structure in the present-day universe formed through hierarchical clustering, with small structures merging to form progressively larger ones, galaxies in the currently most underdense regions must be the least ``evolved'' ones, as they must have formed at later times than those in the dense regions.  Therefore, void and cluster galaxies must follow different evolutionary paths.  Disturbing processes like stripping and harassment, that operate preferentially in crowded environments, should occur rarely in voids.  Studies of the properties of the void galaxies, in contrast to those in relatively crowded regions, or walls, should provide some of the strongest constraints for distinguishing the intrinsic properties, which characterize a galaxy when it is first assembled, from properties that have been externally induced, over the whole history the universe: the ``nature versus nurture'' problem.    Statistically significant conclusions regarding the distinctness of the void galaxies relative to those in denser regions, the ``walls'' hereafter, emerged only recently, with the advent of large surveys such as SDSS and 2dF.  Such data, and in particular SDSS, offered for the first time the possibility to find and analyse both photometrically and spectroscopically, large samples of extremely low density regions (i.e., $\\delta\\rho/\\rho < -0.6$ measured on a scale of $7 h^{-1}$ Mpc, Rojas et al. 2004, 2005), and allowed for accurate estimates of the void galaxy luminosity and mass functions \\citep{hoy05, gol05}.  These studies show that void galaxies are fainter, bluer, have surface brightness profiles more similar to those of late-type systems, and that their specific star formation rates are higher than those in denser regions.  The mass and luminosity functions are found to be clearly shifted towards lower characteristic mass and fainter magnitudes ($M^*$).  Moreover, the faint end slopes of the wall and void luminosity functions are very similar which suggests that voids are not dominated by an excess population of low-luminosity galaxies.  Consistently, no significant excess in the amount of dark matter is apparent.  This means that, although largely devoid of light, the most underdense regions conform to a galaxy formation picture which is clearly not strongly biased.    All these peculiarities demonstrate that the cosmological evolution of void systems is different from that of those living in environments of average cosmic densities.  Given the tight correlations between the mass of black holes ($M_{\\rm BH}$) and the dispersions and the masses of the galactic bulges within which they reside \\citep{mag98, fer00, geb00, mar03}, one would then expect that the growth of massive BHs in galaxy centers (and therefore the accretion process within active galactic nuclei, AGN), also differs among distinct environs. Extension of environmental studies of AGN properties to extreme regions like cosmic voids is thus crucial to understand the co-evolution of galaxies and their central BHs.  Moreover, while there is general agreement that the growth of black holes must be closely related to galaxy assembly \\citep{silk98, kau00, beg05}, there is no consensus as to how exactly accretion and star formation are coupled. The void galaxies could be, arguably, the best test-bed for understanding whether these processes are synchronized, or precede one another, and whether feedback from the actively growing BHs facilitates star formation (e.g., by dynamically compressing gas clouds through radio jets), or suppresses it (e.g., by blowing away the gas).    To date, studies of the spectral properties of the void AGN remain limited to individual voids, e.g., the Bootes void \\citep{kir81}, permitting the identification of only a few AGN among only a few dozen void galaxies \\citep{cru02}.  Quite surprisingly, such investigations find that the AGN fraction and their emission-line properties are similar in voids and in their field counterparts.   Moreover, their associated stellar populations appear to share similar characteristics in the two extreme environs.  The conclusions of these studies are based on small number statistics and do not however exclude the hypothesis that the void emission-line activity, whether originating in star-formation or accretion, could be connected with, e.g., filaments within voids; such structures would provide local environs similar to those in the field.  The present SDSS samples of voids and void galaxies offer us for the first time the possibility to test and observationally constrain such ideas.  It is important to note that previous investigations of the environmental dependence of nuclear activity in the relatively nearby universe do not reach the extreme spatial densities representative of voids.  For example, in \\citet{kau04}, the lowest density regions include over 25\\% of the galaxies, which is more than 3 times more galaxies than the true void regions encompass.  Their conclusions are interesting, and an extension of such an investigation at truly low densities is clearly desirable.  In particular, it is important to quantify the degree to which the finding that, at fixed stellar mass, twice as many galaxies host strong-lined AGN in low-density regions than in high, extends to cosmic voids.  Our work provides such an analysis.  We employ in this work the most accurately classified samples of voids identified within SDSS to date, that yield $\\sim 10^3$ void galaxies. Motivated by our recent study on the AGN clustering phenomena \\citep{con06a}, which shows that there are differences in the large scale structure of active galaxies, and that their clustering amplitude correlates with their strength or rate of accretion, and possibly with the availability of fuel, we compare void and wall active galaxies of different types as classified based on their emission-line properties.  Through such a comparison we aim to understand: 1) how the large and small-scale structures influence accretion onto their central black holes, and 2) to what degree AGN activity is triggered by interactions or mergers between galaxies.  To answer these questions, we investigate the occurrence rate of different types of spectrally defined AGN, and how their accretion activity relates to their associated black hole mass, the mass and the age of their associated stellar populations, host morphology, brightness, and nearest-neighbor distance.   We organize the paper as follows.  In Section ~\\ref{data} we present the void and wall sample selection, and the spectral classification we use in defining various types of actively emitting galaxies.  We compare and discuss the AGN occurrence rate in voids and walls, both globally and at fixed host properties in Section ~\\ref{fractions}.  We examine in Section ~\\ref{voidaccretion} the accretion rates, the fuel supply, and the properties of the associated star-formation in void and wall actively line emitting systems, while in Section ~\\ref{nnprop} we discuss potential differences in their small scale environments.  We summarize our findings in ~\\ref{discussion}, and discuss the possible implications on the nature of the power sources in the low luminosity AGN, the AGN--host connection, and current models of galaxy formation.  In particular, we show empirical evidence for a possible evolutionary sequence that links different types of strong line-emitting galaxies defined based on their spectral characteristics.  Throughout this work, unless otherwise noted, we assume $\\Omega_m = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and $H_0 = 100h$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "\\label{discussion}   \\subsection{\\sc Summary of Results} \\label{results}  We use the largest sample of voids and void galaxies yet defined to investigate the galactic nuclear accretion phenomenon in the most underdense regions, in relation to their more populous counterparts. By employing spectroscopic and photometric data based on the SDSS DR2 catalog, and in particular measurements available in the Garching catalog, we conduct a comparative analysis of void and wall systems of different radiative signature.  We find that all types of Low Luminosity AGN exist in void regions. However, their occurrence rate and intrinsic properties show variation from their wall counterparts.  The differences between the wall and void AGN seem to be driven by the properties of their hosts, which are correlated with (or governed by) their small scale environment. Following is a summary of our main results:  (i) Among moderately bright or fainter galaxies ($10 < log (M_*/M_{\\sun}) < 10.5$, $M_r \\ga -20$), the rate of occurrence of AGN is higher in voids than in walls.  The most common accretion activity in voids is of medium power, with $L_{\\rm [O III]} \\sim 10^{38}$ erg s$^{-1}$.  For the more luminous massive hosts, due to small number statistics, the relative prevalence of accretion activity in voids versus walls remains poorly constrained.  (ii) The majority of void AGN, which are hosted by $M_r \\ga -20$ galaxies, have the tendency to accrete at lower rates than in those in walls. This behavior seems related to the fact that the void systems show less obscuration and, perhaps, less dense emitting gas.  That the stellar populations associated with void and wall AGN are similarly aged suggests that fuel might be equally available for accretion in void and wall galaxies of similar properties (i.e., $M_r$), but that fuel is less efficiently driven towards the nucleus in void galaxies.  (iii) The few void AGN hosted by bright, massive galaxies ($M_r \\la -20$, log $M_*/M_{\\sun} > 10.5$) are LINERs that show peculiarly higher accretion rates, larger amounts of obscuring matter and more recent star-formation than in their wall counterparts. These particular systems reinforce the general trends other objects show:  higher accretion rates are invariably associated with younger stellar populations and higher obscuration.    These trends suggest that the amount of obscuration could be a measure of the available fuel for both star formation and accretion.  (iv) The radio activity of line-emitting galaxies appears both less frequent and weaker in voids than in walls.  Were we able to support these differences with statistically significant measurements, they would imply that central radio activity in wall systems, including H {\\sc ii}s, is more pronounced because it builds on contributions from accretion that remains optically obscured and therefore undetected.  (v) Nearest neighbor statistics show that the type of emission-line activity is correlated with the small-scale local environment.  The star-forming regions (the H{\\sc~ii}s), populate the most crowded sub-regions of voids while populating relatively sparse regions in walls; both are environments where low-mass galaxies recently formed. The weakly active galaxies (LINERs) live within the clusters in walls but the most rarefied regions in voids.  This finding is puzzling and suggests that these systems, which are generally old, were probably not aware of their environments when they formed.  Actively accreting systems (Seyferts and possibly the Ts) inhabit intermediate regions, which are relatively dense galaxy neighborhoods in voids but are of average density in walls.  (vi) These correlations among the type and strength of galactic nuclear activity, incidence rates of different types, and their small and large scale environments, suggest an H{\\sc~ii}$\\rightarrow$S/T$\\rightarrow$LINER evolutionary scenario in which interaction is responsible for propelling gas towards the galaxy centers, triggering star formation and feeding the active galactic nucleus.   \\subsection{\\sc The H{\\sc~ii}$\\rightarrow$S/T$\\rightarrow$LINER Evolutionary Sequence} \\label{sequence}  Figure ~\\ref{h2stl} illustrates how various intrinsic and host properties of actively line-emitting galaxies follow this H {\\sc ii}$\\rightarrow$S/T$\\rightarrow$LINER sequence. The early stages of such objects manifest themselves as H{\\sc~ii}, as the accreting source remains heavily embedded in dust.  As the star-burst fades in time, the dominance of the Seyfert-like excitation in systems of generally small but actively accreting black holes becomes more evident.  Successive evolution reveals aging stellar populations associated with objects spectrally classified as Transition objects, that are still showing signs of accretion, followed by LINERs, whose stars are predominantly old and whose accretion onto already grown-up BHs is close to minimal.  Note that this H{\\sc~ii}$\\rightarrow$S/T$\\rightarrow$LINER progression is very similar in walls and voids. The lower accretion rates and the higher frequency of actively accreting systems in void versus wall galaxies of similar properties indicate a potential delay in the AGN dominance phase within voids.  Thus, void AGN progress through the H {\\sc~ii}$\\rightarrow$S/T$\\rightarrow$LINER sequence more slowly, while the sequence is similar for both void and wall galaxies.  This picture fits well the observed properties of each type of galaxy nucleus:  \\begin{itemize} \\item That H{\\sc~ii}-type of activity is significantly more frequent in voids than in walls suggests that their void-like environments, in which they have closer (both 1st and 3rd) neighbors than other types of objects have, are essential in triggering their activity. In other words, close encounters that produce either harassment, or major and/or minor mergers may be an important cause for igniting both accretion and star formation.  \\item Seyferts' environments in both voids and walls are intermediate between those of H{\\sc~ii}s and LINERs, regardless of their brightness.  If the Seyfert-like activity is triggered by interactions, probably the same ones that turn on the H{\\sc~ii}s, there must be a time lag between the onset of the star-burst and when accretion becomes dominant, or simply observable.  Such a time interval corresponds to a period of aging of the stellar population (as seen in the differences in $D4000$ and H$\\delta_{\\rm A}$ between H{\\sc~ii}s and Seyferts), when the post starburst fuel becomes increasingly available for accretion.  Moreover, this progression develops relatively uninterrupted in voids, and therefore, possibly, at a slower pace than it would in walls where close encounters or other types of interactions can either accelerate or terminate it.  \\item Void and wall Ts are barely distinguishable in their physical characteristics.  In both voids and walls, their nn-statistics and intrinsic and host properties are intermediate between those of LINERs and H{\\sc~ii}s, for any given range of $M_r$.  Their BHs are apparently growing (their accretion activity is stronger than that of Ls but weaker than that of Seyferts), their fuel supply seems plentiful (they are found in some of the most obscured systems), and they are associated with (quite massive) stellar populations that are generally younger than those of LINERs, but older than the majority of star-forming systems.  It is among Ts however that the most massive void accreting BHs are observed; this might suggest that, in the proposed H {\\sc~ii}$\\rightarrow$S/T$\\rightarrow$LINER sequence, massive void galaxies reach the low accretion rate (i.e., LINER) phase later than in walls.  \\item Whether Seyferts or Transition objects are first in this sequence remains ambiguous.  Both their intrinsic properties and nn-statistics are very similar, and remain intermediate between those of H {\\sc~ii}s and LINERs.  The H$\\alpha$/H$\\beta$ Balmer decrements and the $d_{\\rm 1nn}$ are the only parameters that show a ``jump'' in the otherwise smooth H{\\sc~ii}$\\rightarrow$S$\\rightarrow$T$\\rightarrow$LINER sequence manifested by other properties of these systems.  Further investigations of these objects should address the differences between them in terms of a possible evolutionary progression.  \\item Although we do not provide any quantitative estimates of the time spent in or during the various phases we propose here, this this analysis shows that both void and wall galaxies follow the same cycle. The AGN evolution does not affect the gravitational environment.  To the contrary, it is the environment that sets the time scale for evolution along such a sequence; it seems to take longer to march through the different phases in voids than in walls, but the physics is the same.  The large scale clustering is consistent with this picture: LINERs are {\\it now} more clustered because objects in dense regions underwent the H{\\sc~ii}$\\rightarrow$S/T$\\rightarrow$LINER evolution more quickly; the higher rate of galaxy-galaxy interactions speeds up the way AGN proceed through the sequence.  Hosts whose central regions are now in the H {\\sc~ii} phase will always be less clustered than current LINERs.  \\end{itemize}  Although far from being complete, this proposed evolutionary sequence is engaging and offers a comprehensive picture for the co-evolution of AGN and their host galaxies.  The broad idea that mergers trigger star formation and that the AGN appears afterwards, in fact shutting off the star formation because of feedback, has been discussed previously in the literature.  For example, N-body simulations by \\citet{byr86} and \\citet{her95} showed that interactions drive gas toward the nucleus and can produce intense star formation followed by an AGN. More recent state-of-the-art hydrodynamical models \\citep{dimat05, spr05, hop05, hop06} show that during mergers, the BH accretion peaks considerably {\\it after} the merger started, and {\\it after} the star-formation rate has peaked.  However, whether early bright quasars and later, dimmer AGN obey similar physics needs still to be addressed.  The H{\\sc~ii}$\\rightarrow$S/T$\\rightarrow$L sequence that this study reveals, based on the smooth alignment of several of their spectral properties, may be the first empirical evidence for an analogous duty cycle in high redshift bright systems and in nearby galaxies hosting weak quasar-like activity.  This scenario can also accommodate the rather inconclusive findings regarding the role of mergers in activating AGN: their hosts do not show evidence for bars \\citep{mul97, lai02} or disturbances caused by galaxy-galaxy interactions \\citep{mal98}, exhibit morphologies very similar to those of field galaxies \\citep{derob98a, derob98b}, and pair counting in both optical and IR remain inconclusive as possible excess of companions are sometimes found \\citep{dah84, raf95}, but not always \\citep{fue88, lau95}.  Moreover, \\citet{sch01, sch04} show that claims of evidence of interactions in the literature could be attributed to selection effects.  The H{\\sc~ii}$\\rightarrow$S/T$\\rightarrow$LINER cycle suggests that the majority of AGN might be detected only a certain period after the interaction, allowing time for the starburst to fade and for the BH accretion to gain strength.  One might argue that that other forms of evolution could also exist for these emission-line galaxies, as opposed to this simple progression.  We note that this proposed sequence does not imply that every HII galaxy at the present epoch must necessarily become a LINER; it is certainly possible that some systems go through H{\\sc~ii}-only phases, or L-only phases, which might not be part of the larger progression.  We would however like to emphasize that the timescales necessary to transform from one galaxy type to the next are quite reasonable.  At a first glance, this is somewhat surprising given the relatively large range in BH masses ($2 \\times 10^7 M_{\\sun}$ for H{\\sc~ii}s, to $2 \\times 10^8 M_{\\sun}$ for LINERs, inferred from $\\sigma_*$'s assuming, e.g., Tremaine et al. 2002), and their significantly low accretion rates ($L/L_{\\rm Edd} \\la 0.005$; see Figure ~\\ref{h2stl}, and note that log $L[{\\rm O III}]/\\sigma_*^4 = -1$ corresponds to approximately $L/L_{\\rm Edd} = 0.05$, according to Kewley et al. 2006); apparently, for a canonical value for the accretion efficiency, i.e., 10\\%, it takes approximately a Hubble time to e-fold in BH mass.  The key is in the fact that the low luminosity AGN accrete inefficiently, as very little energy generated by accretion is radiated away (the optically thin cooling time of the gas is longer than the inflow time).  Studies show that, in these cases, the radiative efficiency can be as low as $10^{-6}$ \\citep{ree82, nar94, qua03}.  With a (not a necessarily extreme) efficiency value of $10^{-5}$ then, the e-folding time in the BH mass is $\\approx 4.5 \\times 10^2/l$ years, and thus, as little as few Myrs for, e.g., $l = L/L_{\\rm Edd} = 0.0005$.  This (not necessarily extreme) example shows then that such an evolutionary scenario is actually physically possible.  The peculiarity of LINER environments is a new and intriguing result and clearly shows that this type of activity is not controlled by its surrounding environment.  Galaxies hosting LINERs could be associated with high initial density peaks in the dark matter distribution, which evolved subject to their large scale environment.  Being generally massive systems to start with, LINERs' hosts would be prone to accreting material around them.  In voids, this ``cleaning'' enterprise would contribute to emptying the already rarefied neighboring space, leaving little or insufficient material for future formation of (massive, bright) galaxies; they would thus end up in the most underdense void neighborhoods.  In walls, and in particular within clusters, accretion of surrounding material would make a small difference as the matter density is higher.  Simulations of dark matter halos and correlations with the properties of their inhabiting galaxies should be able to address these ideas.    Further tests of this H{\\sc~ii}$\\rightarrow$S/T$\\rightarrow$LINER evolutionary scenario are clearly needed.  These tests require larger samples that allow separation into different morphological subsamples, and observables that parametrize the galaxy morphology better than the concentration index.  Moreover, we need better constraints on the BH masses and consequently the Eddington rates for the H{\\sc~ii}s in particular, or the late type galaxies in general.  When available, analysis of such parameters would shed light on the assumed coevality of star-formation and Seyfert-like BH accretion in centers of galaxies, removing ambiguities regarding the initial stages of the H{\\sc~ii}$\\rightarrow$S/T$\\rightarrow$LINER progression.   \\vspace{1cm}  SPECIAL NOTE: During the review process of this paper, we became aware of another piece of work which introduces the same idea of an evolutionary sequence, ``from star formation via nuclear activity to quiescence`` \\citep{sch07}.  Interestingly, although their general approach, analysis, and samples used, are quite different from ours, the measurements on which the evidence for a ``star-forming, transition object, Seyfert, LINER'' sequence is built shares a common set of parameters with those used in our analysis: stellar velocity dispersions, ages of starbursts, and reddening.  The global frame in which this evolutionary sequence in presented is however definitely different.  While Schawinski et al. take a step forward toward understanding this possible time sequence among early-type galaxies and provide an in-depth analysis of the possible time scales involved in this picture, our paper offers a broader perspective on how such an evolutionary sequence constrains the galaxy evolution models as we provide important links to the environment.      \\vspace{1cm}"
    },
    "0710/0710.3634_arXiv.txt": {
        "abstract": "It is believed that quark matter can exist in neutron star interior if the baryon density is high enough. When there is a large isospin density, quark matter could be in a pion condensed phase. We compute neutrino emission from direct Urca processes in such a phase, particularly in the inhomogeneous Larkin-Ovchinnikov-Fulde-Ferrell (LOFF) states. The neutrino emissivity and specific heat are obtained, from which the cooling rate is estimated. ",
        "introduction": "\\label{introduction}  The phase structure of quantum chromodynamics (QCD) is one of the most challenging problem in particle and nuclear physics. We schematically illustrate in Fig.\\ \\ref{phase} the phase diagrams in $T-\\m_B$ and $T-\\m_I$ plots, where $T$, $\\m_B$ and $\\m_I$ are the temperature, baryon and isospin (or equivalently electron) chemical potentials respectively. In the $T-\\m_B$ diagram, the left panel of Fig.\\ \\ref{phase}, the hadronic phase locates at low $T$ and low $\\m_B$ region and undergoes a phase transition or a crossover to the deconfined quark phase at certain critical temperature $T_c$ or baryon chemical potential $\\m_{Bc}$ of the orders of $T_c\\sim 200$ MeV or $\\m_{Bc}\\sim 1$ GeV. At very high temperature, the quark-gluon-plasma (QGP), made of free quarks and gluons, forms. At asymptotically high $\\m_B$ but low $T$, the ground state of QCD is the color-flavor-locked (CFL) superconductor \\cite{alford1999} where the condensation of quark pairs spontaneously breaks color and chiral symmetries. At intermediate $T$ and $\\m_B$, although quarks and gluons are deconfined they are still strongly coupled. In this regime, many QCD phases are proposed in recent years, such as, at low $T$, two-flavor color superconductivity (2SC) \\cite{Ruester:2004eg}, gapless 2SC (g2SC) \\cite{shovkovy2003}, gapless CFL (gCFL)\\cite{alford2004}, spin-1 color superconductor \\cite{iwasaki1995,schafer2000,schmitt2002}, kaon condensation in the CFL phase \\cite{kaon-condensation}, et al.. For reviews of color superconductivity, see, e.g. Ref. \\cite{csc-reviews}. There may exist resonance states at intermediate $T$, such as strongly coupled QGP (sQGP) \\cite{shuryak2006} at low $\\m_B$ or the pseudo-gap phase at moderate $\\m_B$ \\cite{kitazawa2005}. In the $T-\\m_I$ diagram, the right panel of Fig.\\ \\ref{phase}, the hadron phase is in the region with low $T$ and low $\\m_I$ while the QGP phase locates at very high $T$. At low $T$, when $\\m_I$ grows above the value of the pion mass $m_\\p$, the ground state turns out to be a Bose-Einstein condensation (BEC) of pions, as $\\m_I$ increases further and is larger than about 230 MeV, the pion BEC crossover smoothly into the BCS superfluid of quark-anti-quark pairs with the condensate $\\langle \\ub i\\g_5 d\\rangle$ or $\\langle \\db i\\g_5u\\rangle$ \\cite{son2001}. At intermediate $T$, resonance states such as sQGP may occur at low $\\m_I$ while the states with strong fluctuations of thermally excited mesons or Cooper pairs are possible at moderate $\\m_I$.     \\begin{figure}[!htb] \\begin{center} \\includegraphics[width=6cm]{phase1.eps} \\includegraphics[width=6cm]{phase2.eps} \\caption{The schematic phase diagrams of QCD on $T-\\m_B$ and $T-\\m_I$ planes.} \\label{phase} \\end{center} \\end{figure}  In this paper we consider the quark matter cores in neutron stars. The neutrino emission from direct Urca processes $d\\ra u+e^-+\\nb, u+e^-\\ra d+\\n$ is the most efficient way of cooling in quark matter. Sch\\\"afer and Schwenzer summarized the neutrino emissivities and specific heats for a variety of color superconducting phases of quark matter \\cite{schafer2004}. Due to beta equilibrium the isospin chemical potential $\\m _I$ is nonzero, quark matter could be a pion BEC ($\\m_I<230$ MeV) or a BCS superfluid ($\\m_I > 230$ MeV) when $\\m _I > m_\\p$ \\cite{son2001}. On the other hand, the baryon chemical potential is still large (in the order of 1 GeV) and makes a big mismatch between the fermi surfaces of the pairing quarks $u(\\bar u)$ and $\\bar d(d)$, thus the BEC or BCS state is gapless. Such a gapless phase is stable in the BEC region but unstable in the BCS one with respect to the formation of nonzero LOFF momentum. In a previous work \\cite{huang2007}, we have studied the neutrino emissivity and cooling rate due to Urca processes for the gapless pion condensed quark matter in the BEC region. In this paper we study the neutrino emission in the LOFF phase. We work in two-flavor case in the moderate baryon density, where the role of strange quarks is not important.  Our units are $\\hbar=k_B=c=1$ except particular specifications. As a convention, we denote a 4-momentum as $K^\\mu =(k_0,\\mathbf{k})$, and its 3-momentum magnitude as $k=|\\mathbf{k}|$.   \\section {Quark Propagator\\label{quark}}  Our starting point is the two flavor Nambu-Jona-Lasinio Lagrangian of QCD \\begin{eqnarray} \\label{NJL} \\cl=\\jb(i\\g^\\m\\pt_\\m+\\m\\g_0-m_0)\\j+g[(\\jb\\j)^2+(\\jb i\\vec{\\t}\\g_5\\j)^2], \\end{eqnarray} where $\\j=(u,d)^{\\rm T}$ is the quark fields, $g$ is the coupling constant and $\\vec{\\t}$ is the Pauli matrices. We have introduced the chemical potential matrix in flavor space, $\\m=\\mathrm{diag}(\\m_u, \\m_d)=(\\m+\\d\\m, \\m-\\d\\m)=(\\m_B/3+\\m_I/2, \\m_B/3-\\m_I/2)$ with $\\m_B, \\m_I$ the baryon and isospin chemical potential respectively. We assume that the $\\b$-equilibrium is reached $\\m_d=\\m_u+\\m_e$, which gives $\\m_I=-\\m_e$. In the chiral limit and without the chemical potentials, the Lagrangian (\\ref{NJL}) respects the symmetry $U_B(1)\\otimes SU_V(2)\\otimes SU_A(2)$ corresponding to the baryon number, isospin vector and pseudovector conservation respectively. However, the presence of the chemical potentials explicitly break the isospin symmetry down to $U_V(1)$ with the conserved quantum number of $\\t_3$, and chiral symmetry down to $U_A(1)$ with the conserved quantum number of $i\\g_5\\t_3$. By introducing the chiral condensate $\\s=-2g\\lan\\jb\\j\\ran$ and pion condensates $\\p^-=\\D e^{-2i\\bl\\cdot\\bx}=4g\\langle\\ub i\\g_5 d\\ran$, $\\pi^+=\\D e^{2i\\bl\\cdot\\bx}=4g\\langle\\db i\\g_5 u\\ran$, we arrive at the mean field Lagrangian \\begin{eqnarray} \\cl_\\mf=\\jb\\lb\\begin{array}{cc}i\\g^\\m \\pt_\\m+\\m_u\\g_0-m & \\p^+i\\g_5 \\\\ \\p^-i\\g_5 & i\\g^\\m \\pt_\\m+\\m_d\\g_0-m \\end{array}\\rb\\j -\\frac{\\s^2+\\D^2}{4g}, \\end{eqnarray} with the effective quark mass $m=m_0+\\s$. There are mismatches in Fermi surfaces of anti-u and d quarks or anti-d and u quarks by the baryon chemical potential, hence pion condensates with nonzero total momentum or the LOFF states may be favored. The formation of the condensate $\\s\\sim\\lan\\jb\\j\\ran$ breaks the $U_A(1)$ chiral symmetry spontaneously with the Goldstone boson $\\p_0\\sim\\jb i\\g_5\\t_3\\j$, and that of pion condensates $\\D\\sim \\langle\\db i\\g_5u\\ran\\sim \\langle\\ub i\\g_5d\\ran$ break the $U_V(1)$ isospin symmetry spontaneously. The translational and rotational symmetries are spontaneously broken by nonzero LOFF momentum $\\bl$. The partition function of the system can be written as a functional integral \\begin{equation} Z=\\int [d\\jb ] [d\\j ] \\exp{\\left( \\int_0^\\b d\\t\\int d^3\\bx\\cl_\\mf\\right) } . \\end{equation} Rewriting the quark fields via a gauge transformation which leaves the partition function unchanged, $\\c_u(x)=u(x)e^{-i\\bl\\cdot\\bx}$, $\\c_d(x)= d(x)e^{i\\bl\\cdot\\bx}$ (we still call $\\c _{u,d}$ quark fields), the inverse quark propagator in flavor and momentum space reads \\begin{eqnarray} \\label{s-inverse} S^{-1}(K)=\\lb\\begin{array}{cc}\\g^\\m K_\\m-\\bl\\cdot\\g+\\m_u\\g_0-m & \\D i\\g_5 \\\\ \\D i\\g_5 & \\g^\\m K_\\m+\\bl\\cdot\\g+\\m_d\\g_0-m \\end{array}\\rb. \\end{eqnarray} Note that $K^\\m$ represents the 4-momentum of the $\\c$ fields instead of the $\\j$ fields. The propagator is written as \\begin{eqnarray} \\label{propagator} S(K)=\\lb\\begin{array}{cc} S_{uu}(K) & S_{ud}(K) \\\\ S_{du}(K) & S_{dd}(K) \\end{array}\\rb. \\end{eqnarray} A straightforward calculation from Eq. (\\ref{s-inverse}) gives the four elements \\begin{eqnarray} S_{uu}(K)&=&\\frac{(\\g^\\m K_{+\\m}+m)(K_-^2-m^2-\\D^2)+2\\D^2\\g^\\m l_\\m}{(K_+^2-m^2+\\D^2)(K_-^2-m^2+\\D^2)-\\D^2[(K_++K_-)^2-4m^2]},\\non S_{dd}(K)&=&\\frac{(\\g^\\m K_{-\\m}+m)(K_+^2-m^2-\\D^2)-2\\D^2\\g^\\m l_\\m}{(K_+^2-m^2+\\D^2)(K_-^2-m^2+\\D^2)-\\D^2[(K_++K_-)^2-4m^2]},\\non S_{ud}(K)&=&\\frac{(\\g^\\m K_{+\\m}+m)(K_-^2-m^2-\\D^2)+2\\D^2\\g^\\m l_\\m}{(K_+^2-m^2+\\D^2)(K_-^2-m^2+\\D^2)-\\D^2[(K_++K_-)^2-4m^2]}\\non &&\\times \\frac{\\g^\\m K_{-\\m}-m}{K_-^2-m^2}i\\g^5\\D,\\non S_{du}(K)&=&\\frac{(\\g^\\m K_{-\\m}+m)(K_+^2-m^2-\\D^2)-2\\D^2\\g^\\m l_\\m}{(K_+^2-m^2+\\D^2)(K_-^2-m^2+\\D^2)-\\D^2[(K_++K_-)^2-4m^2]}\\non &&\\times \\frac{\\g^\\m K_{+\\m}-m}{K_+^2-m^2}i\\g^5\\D,\\non \\end{eqnarray} where $K_\\pm^\\m=(k_0+\\m\\pm\\d\\m,\\bk\\pm\\bl)$ and $l^\\m=(\\d\\m, \\bl)$. The excitation spectra of the quasi-particles can be obtained by solving the equation $\\det S^{-1}(k_0,\\bk)=0$ or equivalently the roots of the denominator of the propagator for $k_0$, \\begin{eqnarray} 0&=&(K_+^2-m^2+\\D^2)(K_-^2-m^2+\\D^2)-\\D^2[(K_++K_-)^2-4m^2]\\non &\\approx&[(k_0+\\m+\\ve^-_{\\bk,\\bl})^2-(\\ve_{\\bk,\\bl}^+ +\\d\\m)^2 -\\D^2]\\non &&\\times [(k_0+\\m-\\ve_{\\bk,\\bl}^-)^2-(\\ve_{\\bk,\\bl}^+-\\d\\m)^2-\\D^2], \\end{eqnarray} where $\\ve_{\\bk,\\bl}^\\pm=(E_{\\bk+\\bl}\\pm E_{\\bk-\\bl})/2$ with $E_\\bk\\equiv\\sqrt{\\bk^2+m^2}$. To arrive at the last line, we have taken the assumption that both $\\D$ and $\\bl$ are small comparing to the quark Fermi momenta in the LOFF phase. We have four excitation branches, $E_r^a(\\bk,\\bl)=-r\\sqrt{(\\ve_{\\bk,\\bl}^+ - a\\d\\m)^2+\\D^2}-(\\m - a\\ve_{\\bk,\\bl}^-)$, with $a,r=\\pm$. Taking the same approximation to the numerators in the elements of the quark propagator, we neglect the terms proportional to $\\D^2\\bl$. Thus we rewrite the elements of the quark propagator as, \\begin{eqnarray} \\label{suu-sdd} S_{uu}(K)&\\simeq&\\sum_{a,r=\\pm}\\frac{B_{r}^{a}(\\bk ,\\bl )\\L^{a}_{\\bk+\\bl }\\g_0}{k_0-E_r^a(\\bk,\\bl)},\\non S_{dd}(K)&\\simeq&\\sum_{a,r=\\pm}\\frac{B_{-r}^{a}(\\bk ,\\bl )\\L^{-a}_{\\bk -\\bl }\\g_0}{k_0-E_r^a(\\bk,\\bl)}, \\end{eqnarray} where we have introduced the energy projectors, \\begin{equation} \\L^a_{\\bk }=\\frac{1}{2}\\left[1+a\\frac{\\g_0(\\g\\cdot\\bk+m)}{E_\\bk}\\right], \\end{equation} and the Bogoliubov coefficients, \\begin{equation} \\label{bog-coeff} B^{a}_r(\\bk,\\bl)=\\frac 12 \\left[1-ar\\frac{\\varepsilon ^+_{\\bk,\\bl}-a\\delta\\mu} {\\sqrt{(\\varepsilon ^+_{\\bk,\\bl}-a\\delta\\mu)^2+\\Delta^2 }}\\right], \\end{equation} ",
        "conclusions": ""
    },
    "0710/0710.3896_arXiv.txt": {
        "abstract": "We analyze the mean rest-frame ultraviolet (UV) spectrum of Type Ia Supernovae (SNe) and its dispersion using high signal-to-noise Keck-I/LRIS-B spectroscopy for a sample of 36 events at intermediate redshift ($\\overline{z}$=0.5) discovered by the Canada-France-Hawaii Telescope Supernova Legacy Survey (SNLS).  We introduce a new method for removing host galaxy contamination in our spectra, exploiting the comprehensive photometric coverage of the SNLS SNe and their host galaxies, thereby providing the first quantitative view of the UV spectral properties of a large sample of distant SNe~Ia.  Although the mean SN~Ia spectrum has not evolved significantly over the past 40\\% of cosmic history, precise evolutionary constraints are limited by the absence of a comparable sample of high quality local spectra.  The mean UV spectrum of our $z\\simeq$0.5 SNe~Ia and its dispersion is tabulated for use in future applications. Within the high-redshift sample, we discover significant UV spectral variations and exclude dust extinction as the primary cause by examining trends with the optical SN color. Although progenitor metallicity may drive some of these trends, the variations we see are much larger than predicted in recent models and do not follow expected patterns.  An interesting new result is a variation seen in the wavelength of selected UV features with phase. We also demonstrate systematic differences in the SN~Ia spectral features with SN lightcurve width in both the UV and the optical.  We show that these intrinsic variations could represent a statistical limitation in the future use of high-redshift SNe~Ia for precision cosmology.  We conclude that further detailed studies are needed, both locally and at moderate redshift where the rest-frame UV can be studied precisely, in order that future missions can confidently be planned to fully exploit SNe~Ia as cosmological probes. ",
        "introduction": "\\label{sec:introduction}  Supernovae of Type Ia (SNe~Ia) are now well-established as cosmological distance indicators. In addition to the original surveys by the Supernova Cosmology Project \\citep[SCP;][]{1997ApJ...483..565P,1999ApJ...517..565P} and the High-Z Supernova Search Team \\citep{1998ApJ...507...46S,1998AJ....116.1009R}, a new generation of SN~Ia surveys is underway both locally \\citep{2002SPIE.4836...61A,2005coex.conf..525L,2006PASP..118....2H} and at higher redshifts \\citep{2006A&A...447...31A,2007ApJ...659...98R,2007ApJ...666..694W}. Despite the availability of independent probes of the presence and properties of dark energy from studies of the cosmic microwave background \\citep{2007ApJS..170..377S} and galaxy redshift surveys \\citep{2002MNRAS.330L..29E,2005MNRAS.362..505C,2005ApJ...633..560E}, the luminosity distance--redshift relation for SNe~Ia provides the only {\\it direct} evidence for a cosmic acceleration.  The detection and characterization of dark energy, via measurements of the average cosmic equation of state parameter $<$$w$$>$, requires the precision measurement of SNe~Ia to redshifts $z\\simeq$0.5--1, sampling the epoch of cosmic acceleration \\citep{2006A&A...447...31A}. However, more precise constraints on the nature of dark energy, for example evidence for any variation in $w$ with redshift, requires extending these studies to redshift $z$$>$1 \\citep{2004ApJ...607..665R,2007ApJ...659...98R} where the early effects of deceleration may be detectable. As projects are developed which plan to probe SNe~Ia beyond $z$=1 for this purpose \\citep[e.g][]{2005NewAR..49..346A,2006SPIE.6265E..67B}, it becomes important to understand the possible limitations of using SNe~Ia as distance probes.  Key issues relating to the diversity of SNe~Ia and their possible evolution with redshift as a population, together with the limiting effects of dust and/or color corrections to their photometric properties, are particularly crucial to understand.   Several local studies \\citep{1995AJ....109....1H,1999AJ....117..707R, 2000AJ....120.1479H,2001ApJ...554L.193H, 2005A&A...433..807M,2005ApJ...634..210G} have already indicated correlations between SN~Ia properties and host galaxy morphologies. More recently, \\citet{2006ApJ...648..868S} have shown that the properties of distant SNe~Ia appear to be a direct function of their local stellar population, with the distribution of light curve widths and hence peak luminosities correlating with the host galaxy specific star-formation rate.  This work also determined that the rate of SNe~Ia per unit stellar mass of their host galaxies is larger in actively star-forming galaxies, suggesting many must be produced quite rapidly in recently formed stellar populations, perhaps suggestive of more than one progenitor mechanism. The authors conclude that SNe~Ia may well be a bimodal or a more complex population of events \\citep[see also][]{2005ApJ...629L..85S,2006MNRAS.370..773M}.  Such diversity in the properties of SNe~Ia could have far-reaching implications, particularly if the {\\it mix} of mechanisms or delay times within the broad population gradually changes with look-back time \\citep[e.g.][]{2006MNRAS.370..773M,2006ApJ...648..868S,2007astro.ph..1912H}.  These recent developments, which illustrate how improved precision reveals new physical correlations in the SN~Ia population, raise the broader question of whether future SN~Ia experiments might be limited in precision by variations of a systematic nature within the population, for example with redshift, which cannot be removed via empirical correlations.  Detailed local surveys such as the LOSS/KAIT \\citep{2005coex.conf..525L} and CfA surveys \\citep{1999AJ....117..707R,2006AJ....131..527J} have presented valuable data on the homogeneity and trends in the SNe~Ia population. Further promising work is being undertaken via the Supernova Factory \\citep{2002SPIE.4836...61A} and the Carnegie Supernova Project \\citep{2006PASP..118....2H}. Important though these continued programs will be, they are insufficient to address all possible concerns about the use of SNe~Ia as precision tools in cosmology.  Comparable studies at intermediate redshift\\footnote{Defined here to represent the range 0.2$<$$z$$<$0.7.} will be particularly important in order to address questions relating to possible evolutionary effects and environmental dependencies. In addition, it is not always practical at low redshift to cover the full wavelength range necessary to test for systematic trends.  In this paper we analyze high signal-to-noise ratio rest-frame ultraviolet (UV) spectra of a large sample of intermediate redshift SNe~Ia drawn from the Canada-France-Hawaii Telescope Supernova Legacy Survey \\citep[SNLS;][]{2006A&A...447...31A}, a rolling search which is particularly effective for locating and studying events prior to their maximum light. Our aim is obtain a substantially higher signal-to-noise in the spectra than that typically obtained during spectroscopic programs to type SNe and measure redshifts. We target the UV wavelength region because in this wavelength region the SN spectrum is thought to provide the most sensitive probe of {\\it progenitor metallicity} \\citep[e.g.][]{1998ApJ...495..617H,2000ApJ...530..966L}, a variable which may shed light on the possibility of progenitor evolution. The time-dependent UV spectrum is also needed in estimating ``cross-band'' $k$-corrections, particularly at redshifts $z>1$ where optical bandpasses probe the rest-frame near-UV \\citep{2004ApJ...607..665R,2007ApJ...659...98R}. Little is known about the properties and homogeneity of the UV spectra of SNe~Ia, largely because of the absence of suitable instruments for studying this wavelength range in local events. Although some local SN~Ia UV spectra are available from International Ultraviolet Explorer (IUE) or Hubble Space Telescope (HST) satellite data \\citep[e.g.][]{1991ApJ...371L..23L,1993ApJ...415..589K,1995ESASP1189.....C}, the bulk of the progress now possible in this area can be provided from optical studies of intermediate-redshift events with large ground-based telescopes.  The goals of this paper are thus to address the question of both the diversity and possible physical evolution in the intermediate redshift SN~Ia family. We compare the rest-frame UV behavior of local SNe~Ia with that derived for intermediate-redshift ($z\\simeq$0.5) events where the rest-frame UV enters the region of high efficiency of the Keck LRIS-B spectrograph.  We also study the degree to which the UV spectra at intermediate redshift represent a homogeneous population, independent of other variables such as the physical environment and light curve stretch.  A plan of the paper follows. In $\\S$~\\ref{sec:selection-cfhtls-sne} we introduce the salient features of the SNLS and our method for selecting SNe~Ia for detailed study. In $\\S$~\\ref{sec:keck-observations} we discuss the Keck spectroscopic observations and their reduction, including the treatment of host galaxy subtraction and flux calibration. In $\\S$~\\ref{sec:analyses} we consider our sample with respect to the broader set of SNe found by SNLS, ensuring it is a representative subset in terms of various observables, and discuss existing local UV spectra. In $\\S$~\\ref{sec:results}, we undertake the detailed analysis. First we compare the UV spectra found in our sample with those found locally. We then examine the diversity of intermediate-redshift SNe~Ia in various ways and correlate the UV variations with the light curves of the SNe and the properties of the host galaxies. We discuss these trends in terms of progenitor mechanisms in $\\S$~\\ref{sec:discussion} and examine the implications in terms of possible long term limitations of SNe~Ia as probes of dark energy.  We also present the mean phase-dependent SN Ia spectrum and its uncertainties for use in future work. ",
        "conclusions": "\\label{sec: conclusions}  We summarize our findings as follows:  \\begin{enumerate}  \\item{} We have secured high signal-to-noise ratio Keck spectra for a sample of 36 intermediate redshift SNe~Ia, observed at various phases, spanning the redshift range 0.15$<z<$0.7, and drawn from the Supernova Legacy Survey (SNLS). We demonstrate via inspection of the SN properties that our Keck sample is a reasonably fair subset of the larger sample of distant SNe~Ia being studied by the SNLS.  \\item{} We develop a new method for removing host galaxy contamination from our spectra based on measures of the galaxy and SN photometry. These refinements to traditional spectral reduction techniques allow us to achieve host-galaxy subtracted and flux-calibrated rest-frame spectra of high quality, extending down to rest-frame wavelengths of 2900\\AA.  \\item{} Although no strong evidence is found for spectral evolution in the mean early-phase and maximum light spectra, when compared to local data, such evolutionary tests are hampered by the paucity of quality data at low redshift and a significant scatter in the spectra shortward of 4000\\AA. We find no evidence for evolution internal to our data.  We argue that the well-used local UV spectral template \\citep{2002PASP..114..803N} is likely to be less representative than the mean spectrum compiled from the Keck data which we tabulate with the measured dispersion for use in future cosmological applications.  \\item{} Our principal finding is a large scatter from one SN to the next in the rest-frame UV spectrum even after standard dust corrections are made. By constructing various photometric bandpasses that avoid uncertainties arising from differential $k$-corrections associated with the range of redshifts in our sample, we show that while we can reproduce the stretch-dependent trends seen locally at 3500-4000\\AA, the scatter at 3000-3400\\AA\\ is 3-5 times larger.  \\item{} Although progenitor metallicity may drive some of the trends seen in the Keck data, the UV variations are much larger than in contemporary models which span the expected metallicity range. Moreover, the UV spectrum also changes with phase in a manner which is not consistent with models. We conclude there are significant variations in the UV properties of SNe~Ia which are not accounted for by either the presently-employed empirical trends or the available SN~Ia models.  \\item{} As an illustration of the importance of understanding these new results, we calculate the error arising from the use of a single UV spectral template for calculating the cross-color $k$ correction, a correction essential for constructing the SN~Ia Hubble diagram as a probe of the expansion history. The dispersion arising from our UV spectra, if not randomly distributed along the Hubble diagram, presents an uncertainty 2-3 times larger than than would be necessary for recovering the equation of state parameter $w$ to 5\\% using SNe~Ia at $z\\simeq$1. We conclude that further detailed studies are essential if SNe~Ia are to be useful for precision measures of dark energy.   \\end{enumerate}"
    },
    "0710/0710.1541_arXiv.txt": {
        "abstract": "By considering simple, but representative, models of brane inflation from a single brane-antibrane pair in the slow roll regime, we provide constraints on the parameters of the theory imposed by measurements of the CMB anisotropies by WMAP including a cosmic string component. We find that inclusion of the string component is critical in constraining parameters. In the most general model studied, which includes an inflaton mass term, as well as the brane-antibrane attraction, values $n_{\\rm s}<1.02$ are compatible with the data at $95\\%$ confidence level. We are also able to constrain the volume of the warped throat region (modulo factors dependent on the warp factor) and the value of the inflaton field to be $<0.66M_{\\rm P}$ at horizon exit. We also investigate models with a mass term. These observational considerations suggest that such models have $r< 2\\times 10^{-5}$, which can only be circumvented in the fast roll regime, or by increasing the number of antibranes. Such a value of $r$ would not be detectable in any CMB polarization experiment likely in the near future, but the B-mode signal from the cosmic strings could be detectable. We present forecasts of what a similar analysis using PLANCK data would yield and find that it should be possible to rule out $G\\mu > 6.5\\times 10^{-8}$ using just the TT, TE and EE power spectra. ",
        "introduction": "The inflationary paradigm is strongly supported by observations of the cosmic microwave background (CMB) made by the COBE and WMAP satellites~\\cite{Smoot:1992td,Spergel:2006hy}. However, it is still a paradigm in search of a specific model based on fundamental physics. Brane inflation~\\cite{Dvali:1998pa} has emerged as one of the most popular ways of embedding inflation within string theory. It uses a natural candidate for the inflaton: the field which describes the brane-antibrane separation. This field has a non-trivial potential due to the attractive brane-antibrane interaction which is flattened by the effect of the compactification of the extra dimensions \\cite{Burgess:2001fx}, and the geometry of the branes \\cite{GarciaBellido:2001ky}.  Cosmic strings are also a natural occurrence within this model and are result of inhomogeneities in the tachyon field of the brane-antibrane pair \\cite{C_Strings1,C_Strings2}. These are represented by a complex field with a non-trivial potential which supports the formation of codimension 2 defects which have been shown to exhibit the correct properties of lower dimension branes \\cite{Tachyon_Strings}. Since reheating of the universe in these model proceeds via the annihilation of the brane-antibrane pair, the formation of cosmic strings is expected at the end of the inflation and the strings will be naturally located at the bottom of the throat as their tension is also warp-factor dependent.  The most complete model of brane inflation incorporating moduli stabilization has been proposed in ref.~\\cite{Kachru:2003sx}. In this model inflation happens naturally as a mobile brane falls down the warped throat, being attracted by an antibrane stuck at the bottom of the throat. Antibranes have a warp-factor dependent potential, and therefore minimize their energy by moving to the bottom of the throat,  the region of strongest warping. The mobile brane is not affected by the warping. One can understand this by considering the warped geometry as being generated by a large stack of branes. An antibrane is attracted to the stack and will move towards it, that is towards largest warping. The brane feels no force from the stack of branes (it is BPS with respect to them), only from the antibrane located at the bottom of the throat. This situation can be described by a simple scalar field, since in this model all moduli are stabilized and only the brane-antibrane separation is evolving.  Constraints on the cosmic string tension, $G\\mu$, come from a variety of observations. Of most interest here are those which are most robust. In particular, we will concentrate on the constraints which result from their inclusion as a sub-dominant component in the angular power spectrum of anisotropies of the cosmic microwave background (CMB)~\\cite{CHMb,WB,Wyman:2005tu,Seljak:2006bg,Bevis:2006mj,Battye:2006pk,Bevis:2007gh}. As pointed out recently~\\cite{Battye:2006pk,Bevis:2007gh} in the context of the third year WMAP data, larger values of the spectral index of density fluctuations, $n_{\\rm s}$, are compatible with observations if a sub-dominant string component with around $5-10\\%$ of the large-scale amplitude is included. This is a generic feature of any inflation model which produce strings. In ref.~\\cite{Battye:2006pk} accurate constraints on the coupling constant, $\\kappa$ and mass scale, $M$, relevant to the simplest models of supersymmetric F- and D-term hybrid inflation were computed, taking into account the fact that the observed power spectrum is described by 3 parameters, $G\\mu$, $n_{\\rm s}$ and the power spectrum amplitude, $P_{\\cal R}$, each of which can be derived from $\\kappa$ and $M$. In this model dependent approach more powerful constraints are possible. We will adapt the same approach, where relevant to the case of brane inflation in this paper. ",
        "conclusions": "\\label{sec:discussion}  In this paper we have presented limits and constraints on the parameters (and derived parameters) of brane inflation models in the slow-roll regime when a cosmic string component is included in the fitting. Important aspects of the results are an increased range of acceptable values of $n_{\\rm s}$, limits on $\\gamma$, which is related to the volume of the internal space, $\\beta$ the inflaton mass parameter and a constraint on the value of $\\phi_{\\rm e}<M_{\\rm P}$ irrespective of any theoretical considerations with the consequent implications for $r$.  We note that one way out of the bounds which we have discussed here is to consider the large $\\beta$ regime, where there are again viable inflationary models~\\cite{Bean:2007hc}. These models fall in the category of inflationary models with general speed of sound studied in~\\cite{Peiris:2007gz}.  In this case fast-roll inflation is possible due to the kinetic term actually being of the Dirac-Born-Infeld type (DBI)~\\cite{Silverstein:2003hf,Alishahiha:2004eh,Kecskemeti:2006cg}. This modifies some of the preceding discussion. The Lagrangian for DBI inflation is\\be {\\mathcal L} = \\frac{1}{f\\left(\\phi\\right)}\\left[1-\\sqrt{1- f\\left(\\phi\\right)g^{\\mu\\nu}\\partial_{\\mu}\\phi\\partial_{\\nu}\\phi}\\right] - V\\left(\\phi\\right)\\,, \\ee where the function $f\\left(\\phi\\right)$ depends on the throat geometry and the potential $V(\\phi)$ would be given by (\\ref{masspot}).  In this case, the expression for the scalar-to-tensor ratio is modified to $r=16c_{\\rm s}\\epsilon$ where the parameter $\\epsilon$ is defined as a generalized slow-roll parameter $\\epsilon = 2c_{\\rm s}M_{\\rm P}^2(H^{\\prime}/H)^2$ and the speed of sound is \\be c_{\\rm s}^{2} = 1- f\\left(\\phi\\right)g^{\\mu\\nu}\\partial_{\\mu}\\phi\\partial_{\\nu}\\phi\\,. \\ee Since $c_{\\rm s}^2\\le 1$ one might think that $r$ would be small, however $\\epsilon$ can be much larger in fast roll models. This was emphasized in ref.~\\cite{Bean:2007hc} who also pointed out in the case where $r$ is small the non-Gaussianity of the density fluctuations, quantified by $f_{\\rm NL}$, would be particularly high due to consistency relation~\\cite{Lidsey:2006ia} \\be 1-n_{\\rm s}=0.4r\\sqrt{f_{\\rm NL}}\\,, \\ee which was derived in the equilateral triangle limit in momentum space.  As $\\beta$ increases in the slow-roll $\\text{regime}^{\\footnotemark[1]}$ \\footnotetext[1]{Simultaneously with our work, Ref.\\cite{Bean:2007eh} appeared where an alternative model, the so-called infrared DBI brane inflation, is compared with observations. For this class of models an even stronger bound on the cosmic string tension is found, $G\\mu < 10^{-14}$.} the value of $n_{\\rm s}$ increases beyond that which is compatible with the data. However, at some value of $\\beta$ the effects of the DBI kinetic term kick in and the potential (\\ref{masspot}) becomes dominated by the term $\\propto\\beta$. Hence $P_{\\cal R}$ and $n_{\\rm s}$ are those of simple Klein-Gordon field, albeit modified by $c_{\\rm s}$. This breaks the link between $G\\mu$ and $r$, and therefore the constraints discussed in the previous sections only apply in the slow roll regime.  We note that all the brane inflation models constructed so far require some amount of  fine-tuning to work. In the present case, the tuning corresponds to the value of $\\beta$ being low. The effects of moduli stabilization have led to a large value of the $\\eta$ parameter, making fine-tuning necessary in models of F-term inflation~\\cite{Kachru:2003sx,McAllister:2005mq}. D-term inflation models do not suffer form the same $\\eta$-problem but threshold corrections to the superpotential \\cite{Berg:2004ek,Berg:2005ja} generate a large mass for the inflaton, making fine-tuning necessary in these models as well. A possible way to alleviate the fine-tuning problem is made possible by a remarkable property of a class of multi-brane models~\\cite{Cline:2005ty}. Usually one balances the effect of the volume stabilization mechanism against the Coulombic brane-antibrane interaction to obtain a potential with a flat enough region to support inflation~\\cite{Burgess:2004kv}. The model requires a fine balancing of the two effects. For generic values of the parameters the potential will either be too steep, or it will feature a local de-Sitter minimum where the mobile branes being separated from the antibranes by a potential barrier. However, if one or more branes tunnel out of the local minimum and annihilates with the antibranes, the height of the barrier decreases, and for a critical number of branes it disappears, resulting in a monotonic but almost flat potential. Inflation then proceeds via slow-roll as the remaining branes roll towards the antibranes stuck at the bottom of the warped throat.  We will investigate this type of model using the techniques applied here in a future study.  A similar, but complementary study, to ours for the case $\\beta=0$ but not including the string component has been performed in ref.~\\cite{Lorenz:2007ze}. In order to get meaningful constraints they applied the bound on the exit scale as suggested in ref.\\cite{Baumann:2006cd} which is discussed in section~\\ref{sec:strat}. This study is relevant to the domain where the string tension is weakened by a large number of antibranes. We estimate that if $M>36$ these constraints will apply.  Finally we comment that there may be limits on $G\\mu$ which come from pulsar timing if the cosmic string network achieves scaling by the creation of loops and the subsequent emission of radiation (see ref.~\\cite{Caldwell:1996en} and references therein). We caution that these should  be considered to be less robust since they are more strongly effected by the small-scale dynamics, such as loop formation, of the cosmic string network. This is not completely understood in the case of standard cosmic strings in 3+1 dimensions; the situation with respect to the higher dimension cosmic strings is even less clear. Nonetheless, limits on the energy density in gravitational waves, $\\Omega_{\\rm g}h^2$, have substantially improved in recent times. Probably the most reliable limit is $\\Omega_{\\rm g}h^2<2\\times 10^{-8}$ at frequencies $f=2\\times 10^{-9}{\\rm Hz}$~\\cite{Jenet:2006sv}. Limits on $G\\mu$ from such a bound were considered in ref.~\\cite{Battye:2006pk} and they are dependent on the loop production size relative to the horizon $\\alpha$. If $\\alpha<10^{-4}$ then $G\\mu>10^{-6}$ is excluded and hence the limit from CMB anisotropy is strongest, whereas for $\\alpha>10^{-4}$ one finds that $G\\mu> 10^{-10}/\\alpha$ is excluded, which would be tighter than the CMB limit for $\\alpha>10^{-3}$. It is clear that an improved understanding of the loop production mechanism coupled with expected improvements in the bound on $\\Omega_{\\rm g}h^2$ could lead to a more powerful constraint than is expected from PLANCK."
    },
    "0710/0710.3772_arXiv.txt": {
        "abstract": "Procyon A is a bright F5IV star in a binary system. Although the distance, mass and angular diameter of this star are all known with high precision, the exact evolutionary state is still unclear. Evolutionary tracks with different ages and different mass fractions of hydrogen in the core pass, within the errors, through the observed position of Procyon A in the Hertzsprung-Russell diagram. For more than 15 years several different groups have studied the solar-like oscillations in Procyon A to determine its evolutionary state. Although several studies independently detected power excess in the periodogram, there is no agreement on the actual oscillation frequencies yet. This is probably due to either insufficient high-quality data (i.e., aliasing) or due to intrinsic properties of the star (i.e., short mode lifetimes). Now a spectroscopic multi-site campaign using 10 telescopes world-wide (minimizing aliasing effects) with a total time span of nearly 4 weeks (increase the frequency resolution) is performed to identify frequencies in this star and finally determine its properties and evolutionary state. ",
        "introduction": "The bright F5 subgiant Procyon A is the primary of  an astrometric binary system with a white dwarf in a 40 year orbit. Procyon A is the brightest northern-hemisphere asteroseismology candidate with well-determined characteristics, such as distance, mass and angular diameter. Brown {\\etal} \\cite{brown1991} were among the first to observe an excess power between 0.5 and 1.5 mHz in radial velocity observations confirmed by several other radial velocity studies, e.g. \\cite{martic1999}, \\cite{bouchy2002}, \\cite{kambe2003}, \\cite{martic2004}, \\cite{eggenberger2004}, \\cite{claudi2005}, \\cite{leccia2007}. So far these studies have independently revealed detections of power excess, but there is no agreement yet on the actual oscillation frequencies. This may be due to aliases present in the spectral window, short mode lifetimes, shifts from the asymptotic relation due to avoided crossings, or any combination of these factors. Although the frequencies are not yet known in detail, most studies obtain a large frequency spacing of $55 \\pm 1$ $\\mu$Hz.  \\begin{table}[h!] \\caption{\\label{stelpar} Stellar parameters of Procyon A from \\cite{allende2002}.} \\begin{center} \\begin{tabular}{lr@{$\\pm$}l} \\br Mass [M$_{\\odot}$]& 1.42 & 0.06\\\\ T$_{\\rm eff}$ [K] & 6512 & 49 \\\\ Radius [R$_{\\odot}$] & 2.071 & 0.020\\\\ $\\rm[Fe/H]$  [dex] & $-$0.09 & 0.03 \\\\ $v_{\\rm rot}\\sin i$ [km\\,s$^{-1}$] & 3.16 & 0.50\\\\ \\br \\end{tabular} \\end{center} \\end{table}  Another point of discussion is the fact that \\cite{matthews2004} did not detect any power excess in their MOST photometry and published Procyon A as a flat liner. Other photometric studies such as WIRE \\cite{bruntt2005} and a reanalysis of the MOST 2004 data \\cite{regulo2005} claim to detect power excess in the same region as the radial velocity studies. A recent reanalysis of the MOST 2004 data and the analysis of new MOST data taken in 2005 reinforce the null detection of p-modes \\cite{guenther2007}. This issue is discussed more extensively by \\cite{bedding2005}, who claim that the non-detection of oscillations in Procyon by the MOST satellite is fully consistent with the ground based radial-velocity studies due to a combination of several noise sources and the low photometric amplitude of the oscillations.  Procyon A is in a very interesting evolutionary state near the end of its main sequence life. The stellar parameters (see Table~\\ref{stelpar}) are all known to high precision, but several evolutionary tracks with different ages and different hydrogen core mass fractions overlap, within the errors, with its position in the HR-diagram. The exact evolutionary state of a star can be revealed by means of asteroseismology and therefore the oscillation frequencies are needed.  To determine the actual frequencies of Procyon A a ground-based multi-site campaign using 10 telescopes with a total time span of nearly 4 weeks was performed from December 28, 2006 till January 22, 2007. Here we present first results of this campaign. In Section 2 the campaign is described, while Section 3 discusses the way the data of the different telescopes are combined, and a final power spectrum is obtained. Some concluding remarks and future prospects are provided in Section 4. ",
        "conclusions": "The Procyon campaign presented here is the largest spectroscopic campaign, so far, aimed at the detection and identification of solar-like oscillations. Using 10 telescopes over a time span of nearly 4 weeks provided us with a unique data set with high time coverage and frequency resolution. The details of the data processing methods will be fully described by \\cite{arentoft2008}, while the oscillation frequencies extracted from the full data set acquired during the spectroscopic Procyon campaign will be presented by \\cite{bedding2008}.    \\begin{figure} \\begin{center} \\includegraphics[width=\\linewidth]{powersmall.eps} \\caption{\\label{power}Final power spectrum of Procyon A with window-optimised weights. The data are high-pass filtered to remove low-frequency drifts.} \\end{center} \\end{figure}"
    },
    "0710/0710.1777_arXiv.txt": {
        "abstract": "Studies of stellar magnetism at the pre-main sequence phase can provide important new insights into the detailed physics of the late stages of star formation, and into the observed properties of main sequence stars. This is especially true at intermediate stellar masses, where magnetic fields are strong and globally organised, and therefore most amenable to direct study. This talk reviews recent high-precision ESPaDOnS observations of pre-main sequence Herbig Ae-Be stars, which are yielding qualitatively new information about intermediate-mass stars: the origin and evolution of their magnetic fields, the role of magnetic fields in generating their spectroscopic activity and in mediating accretion in their late formative stages, and the factors influencing their rotational angular momentum. ",
        "introduction": "\\subsection{Magnetism and rotation in the main sequence A and B stars}  Between about 1.5 and 10 M$_{\\odot}$, at spectral types A and B, about 5 \\% of main sequence (MS) stars have magnetic fields with characteristic strengths of about 1kG. Such stars also show important chemical peculiarities and are thus usually called the magnetic chemically peculiar Ap/Bp stars. The strength of the magnetic fields of these stars cannot be explained by an envelope dynamo as in the sun. Until now, the most reliable hypothesis has been to assume a fossil origin for these magnetic fields. This hypothesis implies that the stellar magnetic fields are relics from the field present in the parental interstellar cloud. Its also implies that magnetic fields can (at least partially) survive the violent phenomena accompanying the birth of stars, and can also remain throughout their evolution and until at least the end of the MS, without regeneration.  According to the fossil field model, we should observe magnetic fields in some pre-main sequence (PMS) stars of intermediate mass, the so-called Herbig Ae/Be stars. However no magnetic field was observed up to recently in these stars \\cite[except HD~104237,][]{donati97}. Can we obtain some observational evidence of the presence of magnetic fields during the PMS phase of evolution, as predicted by the fossil field hypothesis? If some Herbig Ae/Be stars are discovered to have magnetic fields, is the fraction of magnetic to non-magnetic Herbig stars the same as the fraction for main sequence stars? Is the magnetic field in Herbig stars strong enough to explain the strength of that of Ap/Bp stars?  Chemical peculiarities and magnetism are not the only characteristic properties observed in the Ap/Bp stars. Most magnetic MS stars have rotation periods (typically of a few days) that are several times longer than the rotation periods of non-magnetic MS stars (a few hours to one day). It is usually believed that magnetic braking, in particular during PMS evolution, when the star can exchange angular momentum with its massive accretion disk, is responsible for this low rotation \\cite[][]{stepien00,stepien02}. An alternative involves a rapid dissipation of the magnetic field during the early stages of PMS evolution for the fastest rotators, due to strong turbulence induced by rotational shear developed under the surface of the stars, as the convection do in the solar-type stars \\cite[see e.g.][]{lignieres96}. In this scenario, only slow rotators could retain their initial magnetic fields, and evolve as magnetic stars to the main sequence. So the question to be addressed is the following: does the magnetic field control the rotation of the star, or else does the rotation of the star control the magnetic field? We propose that this question can be answered by studying rotation and magnetic fields in Herbig Ae/Be stars.   \\subsection{The Herbig Ae/Be stars}  The Herbig Ae/Be stars are intermediate-mass pre-main sequence stars, and therefore the evolutionary progenitors of the MS A and B stars. They are distinguished from the classical Be stars by their IR emission and the association with nebulae, characteristics which are due to their young age.  They display many observational phenomena often associated with magnetic activity. First, high ionised lines are observed in the spectra of some stars \\cite[e.g.][]{bouret97,roberge01}, and X-ray emission have been detected, coming from some Herbig stars \\cite[e.g.][]{hamaguchi05}. In active cool stars, many of these phenomena are produced in hot chromospheres or coronae. Some authors mentioned rotational modulation of resonance lines which they speculate may be due to rotation modulation of winds structured by magnetic field \\cite[][]{praderie86,catala89,catala91,catala99}.  In the literature we find many clues of the presence of circumstellar disks around these stars, from spectroscopic data showing  strong emission, and also from photometric data \\cite[e.g.][]{mannings97,mannings00}. Recently, using coronagraphic data and interferometric data, some authors have also found direct evidence of circumstellar disks around these stars \\cite[][]{grady99,grady00,eisner03}. A careful study of these disks shows that they have similar properties to the disk of their low mass counterpart \\cite[][]{natta01}, the T Tauri stars, whose the emission lines are explained by magnetospheric accretion models \\cite[][]{konigl91,muzerolle98,muzerolle01}. Finally \\citet{muzerolle04} have sucessfully applied their magnetospheric accretion model to Herbig stars to explain the emission lines in their spectra.  For all these reasons we suspect that the Herbig stars may host large-scale magnetic fields that should be detectable with current instrumentation. However, many authors tried to detect such fields without much success \\cite[][]{catala93,catala99,donati97,hubrig04,wade07}. But in 2005, a new high-resolution spectropolarimeter, ESPaDOnS, has been installed at the canada-france-hawaii telescope (CFHT). We therefore decided to proceed to survey many Herbig stars in order to investigate rotation and magnetism in the pre-main sequence stars of intermediate mass. ",
        "conclusions": "We used the new spectropolarimeter ESPaDOnS installed at the CFHT to proceed in a survey of the Herbig stars, in order to investigate their rotation and magnetic field. We discovered four magnetic stars whose field topology is similar to the MS magnetic A-B stars. We also show that the magnetic intensities of these fields and the proportion of magnetic Herbig stars can explain the magnetic intensity and the proportion of magnetic fields among the MS stars, in the context of fossil field model. We therefore bring fundamental arguments in favour of this hypothesis.  The four magnetic Herbig stars are slow rotators ($v\\sin i < 41$ km.s$^{-1}$) which supports that magnetic Herbig Ae/Be stars are the progenitors of the magnetic Ap/Bp stars. Among these magnetic stars two have very low $v\\sin i$ ($< 10$ km.s$^{-1}$) and are very young (age$ < 2.8$ Myr). Assuming these stars are true slow rotators, this implies that there exists a braking mechanism which acts very early during the PMS evolution of the intermediate mass stars. We could think that this braking mechanism has a magnetic origin, although among the undetected stars we also observe same stars with small $v\\sin i$ ($\\sim 15$ km.s$^{-1}$). The nature of the braking mechanism requires addition study.  Finally it has been proposed that all Herbig stars should undergo magnetospheric accretion, as the spectra of all stars show similar emission. However, we calculated the minimum polar magnetic intensity of the magnetic Herbig stars and of two well-constrained undetected stars to get magnetospheric accretion, using three different models which have been developed for the T Tauri stars \\cite[][]{konigl91,cameron93,shu94}. Considering the intensity of the magnetic stars, as well as the maximum magnetic intensity of the two well-constrained undetected stars, if they host a magnetic field, our first conclusion is that magnetospheric accretion cannot occur in all the Herbig stars (a statistic study taking into account all the undetected stars is in progress in order to confirm this result). Therefore, either the models that we used are not well adapted to the Herbig stars, or the emission lines are not only produced by magnetospheric accretion. A thorough observational and theoretical study of the emission in the spectra, as well as the surroundings of the Herbig Ae/Be stars is necessary to better understand the interaction of these stars with their surroundings. \\vspace{1cm}"
    },
    "0710/0710.4370_arXiv.txt": {
        "abstract": "Based on a simulation of galaxy formation in the standard cosmological model, we suggest that a consistent picture for Gamma-Ray Bursts and star formation may be found that is in broad agreement with observations: {\\it GRBs preferentially form in low metallicity environments and in galaxies substantially less luminous that $L_*$.} We find that the computed formation rate of stars with metallicity less than $0.1\\zsun$ agrees remarkably well with the rate evolution of Gamma-Ray Bursts observed by Swift from $z=0$ to $z=4$, whereas the evolution of total star formation rate is weaker by a factor of about $4$. Given this finding, we caution that any inference of star formation rate based on observed GRB rate may require a more involved exercise than a simple proportionality. ",
        "introduction": "The intriguing observational linkage between long duration Gamma-Ray Bursts (GRBs) and core-collapse supernovae (e.g., Stanek \\etal  2003; Hjorth \\etal 2003) suggests that the progenitors of GRBs may be very massive stars. This possible connection was predated by a proposed unified picture (Cen 1998). As such, it has been hoped that GRBs may be a good tracer of cosmic star formation (Wijers \\etal 1998; Totani 1999; Lamb \\& Reichart 2000; Blain \\& Natarajan 2000; Porciani \\& Madau 2001; Daigne \\etal 2006; Coward 2006; Le \\& Dermer 2007; Li 2007). However, recent observations indicate that typical GRBs may prefer relatively low metallicity environments ($\\sim 0.1\\zsun$) (Fynbo \\etal 2003; Le Floch \\etal 2003; Christensen \\etal 2004; Fruchter \\etal 2006; Stanek \\etal 2006) and host galaxies significantly less luminous than $L_*$ (Fruchter \\etal 1999,2006), although there is evidence that the actual spread in metallicity may be wide (Berger \\etal 2006; Prochaska 2006; Wolf \\& Podsiadlowski 2007).   The aim of this {\\it Letter} is to first address the issue of consistency of GRB environment with respect to metallicity and galaxy luminosity, i.e., the galaxy luminosity-metallicity relation, in the context of detailed simulation of galaxy formation in the standard cosmological model. Then, we make predictions on the evolution of GRB rate with redshift and highlight a possible dramatic difference between overall star formation history and GRB rate history, if GRBs are not an unbiased tracer of star formation. In particular, if GRBs are predominantly produced by stars with metallicity $\\le 0.1\\zsun$, the GRB rate is expected in our model to rise obstinately from $z=0$ to $z\\sim 5$ by a factor of $\\sim 100$, when it flattens out towards higher redshift, whereas the overall star formation rate rises rapidly only from $z=0$ to $z\\sim 3$ and is roughly flat from $z\\ge 3$ until $z\\sim 7$. The evolution of GRB rate with redshift is thus expected to be stronger than that of star formation. ",
        "conclusions": "We utilize a simulation of galaxy formation in the standard cosmological model that has been shown to produce results consistent with extant observations of galaxy formation (e.g., Nagamine \\etal 2006) to shed light on the relation between GRB rate and star formation rate. We find that a consistent picture for Gamma-Ray Bursts and star formation that is in broad agreement with observations would emerge, {\\it if GRBs preferentially form in low metallicity environments and in galaxies substantially less luminous that $L_*$.} Because of the increase of metallicity with cosmic time, GRB rate consequently evolves more strongly with redshift than the overall star formation rate. We find that the observed evolution of GRB rate from $z=0$ to $z=4$ can be explained, {\\it if GRBs are primarily produced by massive stars with metallicity less than $0.1\\zsun$}, whereas an inclusion of stars with metallicity as high as $0.3\\zsun$ yields GRB rate evolution from $z=0$ to $z=4$ inconsistent with observations.  Therefore, we reach the conclusion that GRBs may not be a good tracer of cosmic star formation, especially over a long timeline. As a result, a simple inference of star formation rate or its derived quantities such as the ionizing photon production rate at high redshifts, based on the observed GRB rate, should be done with caution and may require careful calibrations.   \\smallskip"
    },
    "0710/0710.4146_arXiv.txt": {
        "abstract": "Motivated by the increasing use of the Kennicutt-Schmidt (K-S) star formation law to interpret observations of high redshift galaxies, the importance of gas accretion to galaxy formation, and the recent observations of chemical abundances in galaxies at \\ztwo--3, I use simple analytical models to assess the consistency of these processes of galaxy evolution with observations and with each other.  I derive the time dependence of star formation implied by the K-S law, and show that the sustained high star formation rates observed in galaxies at \\ztwo--3 require the accretion of additional gas.  A model in which the gas accretion rate is approximately equal to the combined star formation and outflow rates broadly reproduces the observed trends of star formation rate with galaxy age. Using an analytical description of chemical evolution, I also show that this model, further constrained to have an outflow rate roughly equal to the star formation rate, reproduces the observed mass-metallicity relation at \\ztwo. ",
        "introduction": "The motivations of this paper are several.  First, the empirical Kennicutt-Schmidt (K-S) law \\citep{s63,k98schmidt}, which states that the surface density of star formation is proportional to the surface density of gas, is widely used to interpret and describe star formation in galaxies, though its origins are not fully understood and it is just beginning to be tested at high redshift \\citep{btg+04,csn+07,bcd+07}.  The K-S law is a valuable tool when gas masses are not directly measurable; it has been used to estimate the gas masses of both high redshift galaxies \\citep{ess+06mass} and local galaxies in the distant past \\citep{cjp+07}.  The consequences of the K-S law for the evolution of star formation at high redshift are therefore worth considering in more detail, as is the consistency of these consequences with observations.  Second, the fueling of star formation by gas accretion is an essential element of models of galaxy formation, but has largely been neglected by observers of galaxies at high redshifts, mostly because of the lack of evidence for inflow in the spectra of these galaxies. In contrast, the evidence for strong outflows in galaxies at \\ztwo--3 is well-known, most notably in the form of offsets between the redshifts of nebular emission lines, rest-frame UV absorption lines, and \\lya\\ emission (\\citealt{pss+01}, Steidel et al.\\ 2007, in prep).  In spite of the lack of observations of inflow, however, its effects should be considered in the context of other observations, since a significant inflow rate would affect other, measurable properties.  Finally, in recent years metallicity measurements of increasingly large samples of galaxies at $z>1$ have become possible (e.g.\\ \\citealt{kk00,pss+01,sga+04,sep+04,mlc+06}), including the detection of a mass-metallicity relation at $z>2$ \\citep{esp+06}. These measurements still suffer considerably from limitations on the methods that can be used and from calibration uncertainties, but nevertheless they offer a unique opportunity to place constraints on star formation histories and gas flows at high redshift, provided that the effects of inflow, outflow and star formation can be disentangled. Many recent studies have addressed this issue in some detail, including \\citet{ke99,kh05,d07}; and \\citet{fd07}.  The goal of this paper is to formulate simple, analytical models including star formation according to the K-S law, gas inflows and outflows, and chemical evolution.  We would like to assess whether or not these models are consistent with each other and with our current observational knowledge of high redshift galaxies, and to see if they might give a general picture of how gas flows, star formation and metal enrichment may proceed at high redshift.  In \\S\\ref{sec:ks} I derive the explicit time dependence of star formation implied by the K-S law and test its consistency with observations of galaxies at \\ztwo.  In \\S\\ref{sec:metals} I consider simple models of chemical evolution which incorporate both inflows and outflows of gas, and again test their consistency with observations of high redshift galaxies and with the results of the previous section.  Some implications of the results are discussed in \\S\\ref{sec:discuss}.  I adopt a cosmology with $H_0=70\\;{\\rm km}\\;{\\rm s}^{-1}\\;{\\rm Mpc}^{-1}$, $\\Omega_m=0.3$, and $\\Omega_{\\Lambda}=0.7$, and use the \\citet{c03} initial mass function (IMF). ",
        "conclusions": "\\label{sec:discuss}  The above results provide a coherent picture in which strong star formation is sustained by the accretion of gas at approximately the gas processing rate, the outflow rate is roughly equal to the SFR, and metal enrichment is modulated by both outflows and inflows.  This is not a new result; \\citet{fd07} reached many of the same conclusions using cosmological hydrodynamic simulations to reproduce the \\ztwo\\ mass-metallicity relation, and the idea of a system in which star formation is balanced by inflow dates to work by \\citet{l72} and \\citet{tl78}.  Whatever the methods used to reach these conclusions, however, many questions remain about the mechanisms of gas flows and chemical enrichment at high redshift.  The only quantity not yet tied to observations is the gas accretion rate, which the models require to be approximately equal to the gas processing rate.  If the outflow rate is roughly equal to the SFR, the required accretion rate is $\\sim60$ \\msunyr\\ (assuming the average SFR of the UV-selected sample; \\citealt{ess+06stars}), and as much as several hundred \\msunyr\\ or higher for the most rapidly star forming galaxies.  These values are in general agreement with theoretical expectations.  For example, the predicted average gas accretion rates for galaxies in $10^{12}$ \\msun\\ haloes\\footnote{Clustering results indicate that the \\ztwo\\ UV-selected galaxies are typically associated with $\\sim10^{12}$ \\msun\\ haloes \\citep{asp+05}.}  given by \\citet{kkwd05} are $\\sim50$ \\msunyr\\ at the relevant redshifts, rising to several hundred \\msunyr\\ for the $10^{13}$ \\msun\\ haloes expected to host the most massive galaxies (though more recent simulations indicate rates a factor of $\\sim2$ or more lower; D. Kere{\\v s}, private communication).  Gas accretion and cooling rates from the semi-analytic models of \\citet{csw+06} are also of the right order of magnitude.  The question remains as to why such high accretion rates have not yet been observed.  One difficulty is that the velocity range of inflowing gas is likely to be much narrower than the several hundred \\kms\\ spread observed in the outflows.  Another suggestion is that cold, filamentary accretion may dominate at high redshifts \\citep{kkwd05,db06}, in which case detection would depend strongly on projection effects; alternatively, the accretion could occur largely in the form of minor mergers.  Another possibility is that the accreting gas may be too hot to produce signatures in the observed wavebands, or such signatures may simply be too weak to detect.  Even if the hot accreting gas does produce \\ion{C}{4} emission, for example, the line would likely be weak because of the low metallicity of the gas, and it would be superposed on the already complicated \\ion{C}{4} profile.  Detailed modeling of the likely line strengths would help to place limits on detectable accretion rates.  Gas heated to the virial temperature of $\\sim10^6$ K must also be sufficiently cooled in order to fuel star formation; work by \\citet{yssw02} and \\citet{csw+06}, among others, discusses the mechanisms by which this might proceed.  The strong star formation and gas accretion discussed herein will not continue indefinitely.  Observations suggest that most of the star-forming galaxies currently detected at \\ztwo--3 will become largely passively evolving by $z\\sim1$ \\citep{asp+05,pmd+06}.  Because the high observed star formation rates require accretion of new gas to sustain, a decline in the SFR is a natural consequence of the drop in accretion rates at lower redshifts predicted by theoretical models. Many theorists and observers have also proposed that AGN feedback may be responsible for the termination of star formation (e.g.\\ \\citealt{hhc+06} and references therein).  If an additional mechanism to shut off star formation is required, this is a plausible candidate, as the ubiquity of outflows suggests that starburst-driven winds may regulate star formation but do not usually terminate it, and this work implies that strong accretion and outflows may operate simultaneously, or at least alternate in relatively quick succession.  Finally, these results underscore the importance of metallicity measurements for understanding gas flows.  Until the flows can be observed and quantified directly, measurements of gas phase abundances offer the best hope for constraining the outflow and inflow rates of galaxies at high redshift.  There are still considerable difficulties associated with the measurements of these metallicities, but we look forward to improved constraints from new IR spectra and photoionization modeling, and to more direct estimates of outflow rates from detailed spectra (e.g.\\ \\citealt{psa+00}).  Such measurements will give a far more robust picture of star formation, gas flows and metallicity at high redshift than these simple models can provide."
    },
    "0710/0710.3416_arXiv.txt": {
        "abstract": "The \\wire\\ satellite was launched in March 1999 and was the first space mission to do asteroseismology from space on a large number of stars. \\wire\\ has produced very high-precision photometry of a few hundred bright stars ($V<6$) with temporal coverage of several weeks, including K~giants, solar-like stars, \\dss\\ stars, and $\\beta$~Cepheids. In the current work we will describe the status of science done on seven detached eclipsing binary systems. Our results emphasize some of the challenges and exciting results expected from coming satellite missions like COROT and Kepler. Unfortunately, on 23 October 2006, communication with \\wire\\ failed after almost eight years in space. Because of this sad news we will give a brief history of \\wire\\ at the end of this paper. ",
        "introduction": "The failure of the main mission of the \\wire\\ satellite shifted focus to the star tracker, which has a small aperture of $52$\\,mm and is equipped with a 512$\\times$512 pixel SITe CCD. During observations the main target is positioned near the middle of the CCD and the four brightest stars in the 8$\\times$8 degree field of view are also monitored. These four secondary stars are chosen automatically by the on-board software. Each observation comprises a time stamp and an 8$\\times$8 pixel window centred on the target. Two images per second are collected for each star, resulting in typically one million CCD windows in three weeks. Due to pointing restrictions two fields are observed during each \\wire\\ orbit. The duty cycle for one star per orbit is optimally 40\\%, but can be as low as 20\\%. The orbital period has decreased from 96 to 93 minutes over the cause of the mission. The filter reponse of the star tracker is not well defined but Johnson $V+R$ has been suggested \\citep{buzasi00}. For more details about \\wire\\ see Sect.~\\ref{sec:epitaph}.  \\begin{figure*}[h] \\centering \\includegraphics[width=14cm]{wire_vancouver_hr_markbin.ps} \\caption{\\label{fig:hr} Hertzsprung-Russell diagram of around 100 stars observed with \\wire. The locations of the seven dEBs and the Sun are marked.} \\end{figure*} ",
        "conclusions": "We have presented on-going work on the observations and modelling of several eclipsing binary systems observed with the \\wire\\ satellite. The week-long temporal coverage of the targets has made it possible to study dEB systems that are notoriously difficult to observe from the ground, due to their periods being close to an integer number of days. The high precision of the photometry from \\wire\\ is a huge improvement compared to even the best photometric observations from the ground. The two important advantages of space based photometry are: \\begin{itemize} \\item No atmospheric scintillation noise \\item High stability of the \\wire\\ star tracker over periods of several weeks. \\end{itemize}  These data have made it possible to push the limits on the constraints we can put on the theoretical models of these stars, covering the HR diagram from early B to late F type main sequence stars. However, at this level of precision the oscillations of the component stars, although amplitudes are small, need to be taken into account as part of the light curve analysis.  In a broader context, the results presented here give an idea of the potential of secondary science on dEBs that can be done with data from the future satellite photometry missions. Many new systems will be discovered with COROT (see \\citep{michel05}) and Kepler as discussed by \\cite{koch07}. Although the new systems will be much fainter, follow-up spectroscopy and multi-band photometry can still be done from the ground."
    },
    "0710/0710.0553_arXiv.txt": {
        "abstract": "The DAMA Collaboration has recently analyzed its data of the extensive  WIMP direct search (DAMA/NaI) which detected an annual modulation, by taking into account the channelling effect which occurs when an ion  traverses a detector with a crystalline structure. Among possible implications, this Collaboration has considered the case of a coherent WIMP-nucleus interaction and then derived  the form of the annual modulation region in the plane of the WIMP-nucleon cross section versus the WIMP mass, using a specific modelling for the channelling effect. In the present paper we first show that light neutralinos fit the annual modulation region also when channelling  is taken into account. To discuss the connection with indirect signals consisting in galactic antimatter, in our analysis we pick up a specific galactic model, the cored isothermal-sphere. In this scheme we determine the sets of supersymmetric models selected by the annual modulation regions and  then  prove that these sets are compatible with the available data on galactic antiprotons. We comment on implications when other galactic distribution functions are employed. Finally, we show that future measurements on galactic antiprotons and antideuterons will be able to shed further light on the populations of light neutralinos singled out by the annual modulation data. ",
        "introduction": "\\label{sec:intro}  \\begin{figure}[t] \\centering \\vspace{-20pt} \\includegraphics[width=1.0\\columnwidth]{A1_global.ps} \\vspace{-30pt} \\caption{WIMP--nucleon scattering cross-section as a function of the WIMP mass. The solid (dashed) line denotes the annual modulation region derived by the DAMA Collaboration with (without) the inclusion of the channeling effect. The two regions contain points where the likelihood- function values differ more than 4$\\sigma$ from the null hypothesis (absence of modulation). These regions are obtained by varying the WIMP galactic distribution function (DF) over the set considered in Ref. \\cite{bcfs} and by taking into account other uncertainties of different origins \\cite{damalast}. The scatter plot represents supersymmetric configurations calculated with the supersymmetric model summarized in the Appendix. The (red) crosses denote configurations with a neutralino relic abundance which matches the WMAP cold dark matter amount ($0.092 \\leq \\Omega_{\\chi} h^2 \\leq 0.124$), while the (blue) dots refer to configurations where the neutralino is subdominant ($\\Omega_{\\chi} h^2 < 0.092$).} \\label{fig:00} \\end{figure}   In a recent paper  \\cite{damalast} the DAMA Collaboration has analyzed the data of its extensive  WIMP direct search (DAMA/NaI) \\cite{dama} which measured an annual modulation effect at 6.3 $\\sigma$ C. L., by taking into account the channelling effect.  This effect occurs when an ion  traverses a detector with a crystalline structure  \\cite{drob}. In Ref. \\cite{damalast} implications of channelling  have been  discussed in terms of a specific modelling of this effect for the case of the DAMA NaI(Tl) detector; it is shown that the occurrence of  channelling  makes the response of this detector to WIMP-nucleus interactions more sensitive than in the case in which channelling is not included. Therefore, when applied to a WIMP with a coherent interaction with nuclei, the inclusion of the channelling effect implies that the annual modulation region, in the plane of the WIMP-nucleon cross section versus the WIMP mass, is considerably modified as compared to the one derived without including channelling. The extent of the modification depends on the specific model--dependent procedure employed in the evaluation of the channeling effect \\cite{damalast}.  These properties are shown in Fig. \\ref{fig:00}, where the quantity $\\sigma^{\\rm nucleon}_{\\rm scalar}$ denotes the WIMP-nucleon scalar cross-section, $\\xi = \\rho_{\\rm WIMP}/\\rho_0$ is the WIMP local fractional matter density and $m_{\\chi}$ is the WIMP mass. The dashed line denotes the annual modulation region derived by the DAMA Collaboration without including the channeling effect \\cite{dama}. The solid line shows the annual modulation region derived by the same Collaboration when the channeling effect is included as explained in Ref. \\cite{damalast}.  The regions displayed in Fig. \\ref{fig:00} are derived by varying the WIMP galactic distribution function (DF) over the set considered in Ref.\\cite{bcfs} and by taking into account other uncertainties of different origins \\cite{damalast,nota1}. Fig. \\ref{fig:00} shows that   the  effect of taking channelling into account is  that the annual modulation region modifies its contour with an extension towards lighter WIMP masses. Most remarkably, for WIMP masses $\\lsim $ 30 GeV, the WIMP-nucleon cross section involved in the annual modulation effect decreases sizeably, up to  more than an order of magnitude. As mentioned before, the specific shape of the annual modulation region depends on the way in which channelling is modelled \\cite{damalast}.  These features are  of great importance for a specific dark matter candidate, the light neutralino, which was extensively investigated in Refs. \\cite{lowneu,lowdir,ind}. Actually, in these papers  we analyzed  light neutralinos, {\\it i.e.} neutralinos with a mass $m_{\\chi} \\lsim 50$ GeV, which arise naturally in supersymmetric models where gaugino mass parameters are not related by a GUT--scale unification condition. In Refs. \\cite{lowneu,lowdir} it is proved that, when R-parity conservation is assumed, these neutralinos are of great relevance for the DAMA/NaI annual modulation effect. In these papers it is also shown that in MSSM without gaugino mass unification the lower limit of the neutralino mass is $m_{\\chi} \\gsim$ 7 GeV \\cite{hp}.  In Fig. \\ref{fig:00}, superimposed to the annual modulation regions is  the scatter plot of the supersymmetric configurations of our model, whose features are summarized in the Appendix. One sees that, also when the channeling effect is taken into account, the light neutralinos of our supersymmetric model  fit quite well the annual modulation region.  In the present paper we consider the phenomenological consequences for light neutralinos when the annual modulation region is the one indicated by the solid line in Fig.\\ref{fig:00}. More specifically we examine the properties of our supersymmetric population of light relic neutralinos in terms of the possible antimatter components generated by their pair annihilation in the galactic halo.  To do this, we have to resort to a specific form for the WIMP DF. We take as our representative DF a standard cored isothermal-sphere model, though we do not  mean to associate to this model prominent physical motivations as compared to other forms of DFs. Analyses similar to the one we present here for the cored isothermal-sphere can be developed for  other galactic models. We will comment about some of them, selected among those considered in Ref. \\cite{bcfs} (we will follow the denominations of this Reference to classify our DFs).  The scheme of the present paper is the following. In Sect. II, we show how  the model presented in Refs.  \\cite{lowneu,lowdir,ind} fits the DAMA/NaI annual modulation regions of Ref. \\cite{damalast} for the case of the cored isothermal-sphere model. In Sect. III we combine these results with constraints derivable from available data on cosmic antiprotons; we also discuss the sensitivity of upcoming measurements on cosmic antiprotons for investigating the neutralino populations selected by the annual modulation regions. Complementary investigations by measurements of galactic antideuterons are presented in Sect. IV. Conclusions are drawn in Sect.V. The main features of the supersymmetric scheme adopted here are summarized in the Appendix. ",
        "conclusions": "In the present paper we have considered the annual modulation regions which the DAMA Collaboration has recently determined, by including also the channelling effect which occurs when an ion traverses a detector with a crystalline structure, such the detector of the DAMA/NaI experiment. The inclusion of the channelling effect implies that the annual modulation region is considerably modified as compared to the one derived without including channelling. The extent of the modification depends on the specific model--dependent procedure employed in the evaluation of the channeling effect.  In the present paper we have considered the phenomenological consequences for light neutralinos when the annual modulation region includes the channelling effect as modelled in Ref.\\cite{damalast}. We have proved that these annual modulation data are fitted by light neutralinos which arise naturally in supersymmetric models where gaugino mass parameters are not related by the a GUT--scale unification condition.  The precise range of the neutralino mass which fits the annual modulation data depends on how the WIMP galactic distribution function is modelled and on a number of other assumptions, such as those mentioned in Sect. II. As an example, we have worked out in detail the case of a cored isothermal sphere DF. For this instance, the neutralino mass stays in the range $m_{\\chi} \\simeq (7 - 30)$ GeV, for values of the local rotational velocity, $v_0$, and of the local dark matter density, $\\rho_0$,  in the low-medium side of their own physical ranges, {\\it i.e.} $v_0 \\simeq$ (170 - 220) km sec$^{-1}$ and $\\rho_0 \\simeq (0.2 - 0.4)$ GeV cm$^{-3}$. Similar ranges are found also in the case of a Navarro--Frenk--White profile or for an isothermal model with a non-isotropic velocity dispersion.  We have then shown that the populations of light neutralinos selected by the annual modulation regions are consistent with present data on galactic antiprotons. We have also derived the intervals of the diffusion parameters which provide this agreement in correlation with the specific galactic halo model. For instance, for neutralinos with a mass of 20 GeV and a cored isothermal model with $v_0 = 170$ km s$^{-1}$ we have $0.55 \\lsim \\delta \\lsim 0.85$ and $L \\lsim 3$ kpc when $ \\rho_0 = \\rho_0^{\\rm max} = 0.42$ GeV cm$^{-3}$; instead when $ \\rho_0 = \\rho_0^{\\rm min} = 0.20$ GeV cm$^{-3}$, $L$ may go up to 15 kpc with a range of $\\delta$ which progressively shrinks to $\\delta \\sim 0.70 - 0.75$, when $L$ increases.  We have also shown that future measurements of galactic antiprotons and antideuterons will offer, together with the upcoming data from DAMA/LIBRA,  very interesting perspectives for further investigating the light neutralino populations selected by the annual modulation data. In case of models with a corotating halo or with triaxial spatial distributions, not investigated in the present paper, also heavier neutralinos can be involved.  Finally, a word of caution should be said concerning the fact that the distribution of WIMPs in the Galaxy could deviate from the models mentioned above, mainly because of the presence of streams and/or clumpiness. In such instances, the analysis should be appropriately adapted, along the lines discussed in the present paper."
    },
    "0710/0710.0765_arXiv.txt": {
        "abstract": "CCD $VRI$ photometry is presented for SN 2002hh from 14 days after the outburst till day 347. SN 2002hh appears to be normal type IIP supernova regarding both luminosity and the shape of the light curve, which is similar to SN 1999gi. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.0003_arXiv.txt": {
        "abstract": "We present a study of elemental abundances in a sample of thirteen Blue Compact Dwarf (BCD) galaxies, using the $\\sim$10--37$\\mu$m high resolution spectra obtained with Spitzer/IRS. We derive the abundances of neon and sulfur for our sample using the infrared fine-structure lines probing regions which may be obscured by dust in the optical and compare our results with similar infrared studies of starburst galaxies from ISO.  We find a good correlation between the neon and sulfur abundances, though sulfur is under-abundant relative to neon with respect to the solar value. A comparison of the elemental abundances (neon, sulfur) measured from the infrared data with those derived from the optical (neon, sulfur, oxygen) studies reveals a good overall agreement for sulfur, while the infrared derived neon abundances are slightly higher than the optical values.  This indicates that either the metallicities of dust enshrouded regions in BCDs are similar to the optically accessible regions, or that if they are different they do not contribute substantially to the total infrared emission of the host galaxy. ",
        "introduction": "Blue Compact Dwarf Galaxies (BCDs) are dwarf galaxies with blue optical colors resulting from one or more intense bursts of star-formation, low luminosities (M$_B>-$18) and small sizes. The first BCD discovered was I\\,Zw\\,18 by \\citet{Zwicky66}, which had the lowest oxygen abundance observed in a galaxy \\citep{Searle72}, until the recent study of the western component of SBS0335-052 \\citep{Izotov05}.  Although BCDs are defined mostly by their morphological parameters, they are globally found to have low heavy element abundances as measured from their HII regions (1/30\\,Z$_\\odot$$\\sim$1/2\\,Z$_\\odot$). The low metallicity of BCDs is suggestive of a young age since their interstellar medium is chemically unevolved. However, some BCDs do display an older stellar population and have formed a large fraction of their stars more than 1Gyr ago \\citep[see][]{Loose85,Aloisi07}. The plausible scenario that BCDs are young is intriguing within the context of Cold Dark Matter models which predict that low-mass dwarf galaxies, originating from density perturbations much less massive than those producing the larger structures, can still be forming at the current epoch. However, despite the great success in detecting galaxies at high redshift over the past few years, bona fide young galaxies still remain extremely difficult to find in the local universe \\citep{Kunth86,Kunth00,Madden06}. This is likely due to the observational bias of sampling mostly luminous more evolved galaxies at high redshifts. If some BCDs are truly young galaxies, they would provide an ideal local laboratory to understand the galaxy formation processes in the early universe.  Over the past two decades, BCDs have been studied extensively in many wavelengths using ground-based and space-born instruments \\citep [for a review see][]{Kunth00}.  In the FUV, the Far Ultraviolet Spectroscopic Explorer (FUSE) has been used to study the chemical abundances in the neutral gas in several BCDs \\citep{Thuan02, Aloisi03, Lebouteiller04}. Optical spectra have been obtained for a large number of BCDs and display strong narrow emission lines resulting from the intensive star-formation processes that take place in these systems \\citep{Izotov97, Izotov99b, Pustilnik05, Salzer05}. The Infrared Space Observatory (ISO) revealed unexpectedly that despite their low metallicities, BCDs, such as SBS\\,0335-052E, could still have copious emission from dust grains \\citep{Thuan99, Madden00, Madden06, Plante02}. More recently, the Spitzer Space Telescope \\citep{Werner04} has been used to observe these metal-poor dwarf systems in order to study their dust continuum properties and the polycyclic aromatic hydrocarbon (PAH) features \\citep{Houck04b, Hogg05, Engelbracht05, Rosenberg06, Wu06, OHalloran06, Hunt06, Wu07}. Finally, radio observations have also been performed for several BCDs to study their HI kinematics and distribution \\citep{Thuan04} and thermal/non-thermal continuum emission properties \\citep{Hunt05}.  \\begin{deluxetable*}{lrllllcc} \\tabletypesize{\\scriptsize} \\setlength{\\tabcolsep}{0.02in} \\tablecaption{Observing Parameters of the Sample\\label{tab1}} \\tablewidth{0pc} \\tablehead{ \\colhead{Object Name} & \\colhead{RA} & \\colhead{Dec} & \\colhead{AORKEY} & \\colhead{Observation} & \\colhead{Redshift} & \\multicolumn{2}{c}{On-source Time (Sen)}\\\\ \\colhead{} & \\colhead{(J2000)} & \\colhead{(J2000)} & \\colhead{} & \\colhead{Date} &  \\colhead{} & \\colhead{SH} & \\colhead{LR} \\\\ } \\startdata  Haro11              & 00h36m52.5s  & -33d33m19s &  9007104 & 2004-07-17  &  0.0206 & 480     & 240      \\\\ NGC1140             & 02h54m33.6s  & -10d01m40s &  4830976 & 2004-01-07  &  0.0050 & 480     & 240      \\\\ SBS0335-052E        & 03h37m44.0s  & -05d02m40s & 11769856 & 2004-09-01  &  0.0135 & 1440    & 960      \\\\ NGC1569             & 04h30m47.0s  & +64d50m59s &  9001984 & 2004-03-01  & $\\sim$ 0& 480     & 240      \\\\ IIZw40              & 05h55m42.6s  & +03d23m32s &  9007616 & 2004-03-01  &  0.0026 & 480     & 240      \\\\ UGC4274             & 08h13m13.0s  & +45d59m39s & 12076032 & 2004-10-23  &  0.0015 & 120     &  56      \\\\ &              &            & 12626688 & 2004-11-11  &         & 120     &  56      \\\\ IZw18               & 09h34m02.0s  & +55d14m28s &  9008640 & 2004-03-27  &  0.0025 & 480     & 240      \\\\ &              &            & 16205568 & 2005-12-16  &         & 2880    & 1440     \\\\ VIIZw403            & 11h27m59.9s  & +78d59m39s &  9005824 & 2004-12-09  & $\\sim$ 0& 480     & 240      \\\\ Mrk1450             & 11h38m35.6s  & +57d52m27s &  9011712 & 2004-12-12  &  0.0032 & 480     & 240      \\\\ UM461               & 11h51m33.3s  & -02d22m22s &  9006336 & 2005-01-03  &  0.0035 & 480     & 240      \\\\ &              &            & 16204032 & 2006-01-14  &         & 1440    & \\nodata  \\\\ SBS1210+537A        & 12h12m55.9s  & +53d27m38s &  8989952 & 2004-06-06  & \\nodata & 480     & 240      \\\\ Tol1214-277         & 12h17m17.1s  & -28d02m33s &  9008128 & 2004-06-28  &  0.0260 & 480     & 240      \\\\ Tol65               & 12h25m46.9s  & -36d14m01s &  4829696 & 2004-01-07  &  0.0090 & 480     & 240      \\\\ UGCA292             & 12h38m40.0s  & +32d46m01s &  4831232 & 2004-01-07  &  0.0010 & 480     & 240      \\\\ Tol1304-353         & 13h07m37.5s  & -35d38m19s &  9006848 & 2004-06-25  &  0.0140 & 480     & 240      \\\\ Pox186              & 13h25m48.6s  & -11d37m38s &  9007360 & 2004-07-14  &  0.0039 & 480     & 240      \\\\ CG0563              & 14h52m05.7s  & +38d10m59s &  8992512 & 2005-05-30  &  0.0324 & 240     & 120      \\\\ CG0598              & 14h59m20.6s  & +42d16m10s &  8992256 & 2005-03-19  &  0.0575 & 480     & 240      \\\\ CG0752              & 15h31m21.3s  & +47d01m24s &  8991744 & 2005-03-19  &  0.0211 & 480     & 240      \\\\ Mrk1499             & 16h35m21.1s  & +52d12m53s &  9011456 & 2004-06-05  &  0.0090 & 480     & 240      \\\\ {\\rm [RC2]}A2228-00 & 22h30m33.9s  & -00d07m35s &  9006080 & 2004-06-24  &  0.0052 & 480     & 240      \\\\  \\enddata  \\tablecomments{The coordinates and redshifts of the objects are cited from the NASA/IPAC Extragalactic Database (NED). In this paper, we only include the analysis of thirteen out of twenty-two sources which have SNRs sufficient for our abundance study. CG0563, CG0598 and CG0752 are included in the original sample as BCD candidates, however, they appear to be more starburst like \\citep[see][]{Hao07}. Thus even though they have high SNR, we do not include these three sources in this study.} \\end{deluxetable*}    Metallicity is a key parameter that influences the formation and evolution of both stars and galaxies. Detailed studies of the elemental abundances of BCDs have already been carried out by several groups \\citep{Izotov99b, Kniazev03, Shi05} and the well known metallicity-luminosity relation has also been studied in detail in the environment of dwarf galaxies \\citep{Skillman89, Hunter99, Melbourne02}. However, because these studies were performed in the optical, they were limited by the fact that the properties of some of the deeply obscured regions in the star-forming galaxies may remain inaccessible due to dust extinction at these wavelengths. In fact, \\citet{Thuan99}, using ISO, have shown that the eastern component of SBS\\,0335-052 does have an embedded super star cluster (SSC) that is invisible in the optical while contributing $\\sim$75\\% to the bolometric luminosity \\citep[see also][]{Plante02,Houck04b}, even though it has very low metallicity (12+log(O/H)=7.33), which would in principle imply a low dust content. In addition to probing the dust enshrouded regions, emission in the infrared also has the advantage that the lines accessible at these wavelengths are less sensitive to the electron temperature fluctuations than the corresponding optical lines of the same ion. In the infrared, more ionization stages of an element become available as well. The improved sensitivity of the Infrared Spectrograph (IRS\\footnote{The IRS was a collaborative venture between Cornell University and Ball Aerospace Corporation funded by NASA through the Jet Propulsion Laboratory and the Ames Research Center.})  \\citep{Houck04a} on Spitzer has enabled us to obtain for the first time infrared spectra for a much larger sample of BCDs than was previously possible \\citep{Thuan99, Madden00, Verma03, Martin06}, thus motivating this study to probe the heavy element abundances in BCDs.  In this paper, we analyze Spitze/IRS spectra of thirteen BCDs and present elemental abundances of neon and sulfur, which are both primary elements produced by the same massive stars in the nuclear synthesis processes. In section 2, we describe the sample selection, observations and data reduction. We present our results on the chemical abundances in section 3, followed by a comparison of the optical and infrared derived abundances in section 4. We show the interplay between the abundances and PAH emission in section 5. Finally, we summarize our conclusions in section 6. ",
        "conclusions": "We have studied the neon and sulfur abundances of thirteen BCDs using Spitzer/IRS high-resolution spectroscopy.  Our analysis was based on the fine-structure lines and the hydrogen recombination line detected in the SH spectra of the {\\em IRS}, combined with the radio continuum, H$\\alpha$ images and integrated optical spectral data in some cases. We find a positive correlation between the neon and sulfur abundances, though sulfur appears to be more under-abundant than neon (with respect to solar). The ratio of Ne/S for our sources is on average 11.4$\\pm$2.9, which is consistent with what has been found in other HII regions using infrared data. However, this average ratio appears to be higher than the corresponding optical value of 6.5$\\pm$1.8 (in BCDs), which could be due to the adopted ICFs in the optical studies. When comparing the newly derived neon and sulfur abundances with the oxygen abundances measured from the optical lines, we find a good overall agreement.  This indicates that there are few completely dust enshrouded HII regions in our BCDs, or if such HII regions are present, they have similar metallicities to the ones probed in the optical. Finally, the infrared derived neon and sulfur abundances also correlate, with some scatter, with the corresponding elemental abundances derived from the optical data."
    },
    "0710/0710.0185.txt": {
        "abstract": "Observational and theoretical evidence suggests that coronal heating is impulsive and occurs on very small cross-field spatial scales.  A single coronal loop could contain a hundred or more individual strands that are heated quasi-independently by nanoflares.  It is therefore an enormous undertaking to model an entire active region or the global corona. %therefore requires simulating the %time-dependent plasma evolution of enormous numbers of %elemental strands. Three-dimensional MHD codes have inadequate spatial resolution, and 1D hydro codes are too slow to simulate the many thousands of elemental strands that must be treated in a reasonable representation.  Fortunately, thermal conduction and flows tend to smooth out plasma gradients along the magnetic field, so ``0D models\" are an acceptable alternative.  We have developed a highly efficient model called Enthalpy-Based Thermal Evolution of Loops (EBTEL) that accurately describes the evolution of the average temperature, pressure, and density along a coronal strand.  It improves significantly upon earlier models of this type---in accuracy, flexibility, and capability. It treats both slowly varying and highly impulsive coronal heating; it provides the time-dependent differential emission measure distribution, $D\\!E\\!M(T)$, at the transition region footpoints; and there are options for heat flux saturation and nonthermal electron beam heating.  EBTEL gives excellent agreement with far more sophisticated 1D hydro simulations despite using four orders of magnitude less computing time.  It promises to be a powerful new tool for solar and stellar studies. ",
        "introduction": "An abundance of observational and theoretical evidence indicates that much of the corona is highly dynamic and evolves in response to heating that is strongly time-dependent.  The evidence further suggests that the cross-field spatial scale of the heating is very small, so that unresolved structure is ubiquitous.  In particular, many if not all coronal loops are bundles of thin strands that are heated impulsively and quasi-randomly by nanoflares.  It is estimated that a single loop contains several tens to several hundreds of such strands.  See \\citet{k06} for a detailed justification of these ideas and references to relevant work.  Three-dimensional (3D) magnetohydrodynamic simulations are extremely useful for studying the source of coronal heating (instabilities of electric current sheets, reconnection, turbulence, etc.), but they cannot adequately address the complexity that is present in a single coronal loop, much less an entire active region.  A more feasible approach is to treat the magnetic field as static and to solve the one-dimensional (1D) hydrodynamic equations along many representative flux strands using an assumed heating rate. The individual strands {\\it must} be treated separately.  It is not valid to approximate a loop as a monolithic structure with uniform heating corresponding to the average for the component strands.  This gives a completely different and incorrect result.  There is reason to believe that the diffuse corona that lies between distinct bright loops is also comprised of elemental strands (e.g., \\citealt{anb07}). If roughly 100 strands are present in a single loop, then the numbers present in active regions and the global Sun are truly staggering. While it is possible to construct a limited number of model active regions with time-dependent 1D simulations \\citep{ww07}, it is not possible to investigate a wide range of values for the coronal heating parameters that must be assumed, such as the dependence on magnetic field strength, loop length, etc. \\citep{mdk00}.  This is a major limitation, since we are still struggling to identify the properties and physical origin of the heating mechanism. Progress in the foreseeable future must therefore rely on simplified solutions to the hydro equations that treat field-aligned averages and are much less computationally intensive. These are sometimes called ``0D models\" because there is only one value of temperature, pressure, and density at any given time in the simulation.  0D models have been developed previously by \\citet{fh90} and \\citet{kp93}, but the best known is that of \\citet{c94}.  It has been used to study a variety of topics, including coronal loops \\citep{ck97, ck06, kc01,petal06}, flares \\citep{rw02,pak02}, post-eruption arcades \\citep{rf05}, and active stellar coronae \\citep{ck06}. We have learned a great deal with the Cargill model, and our understanding has now advanced to the point where a more accurate and flexible model is required.  This article presents an improved 0D model called Enthalpy-Based Thermal Evolution of Loops (EBTEL).  As the name suggests, a key aspect of the model is an explicit recognition of the important role that enthalpy plays in the energy budget.  EBTEL improves upon the Cargill model in several important ways.  First, whereas the Cargill model is limited to an instantaneous heat pulse, EBTEL accommodates any time-dependent heating profile and can include a low-level background heating if desired.  Second, EBTEL accounts for thermal conduction cooling and radiation cooling at all times during the evolution.  The Cargill model assumes that only one or the other operates at any given time.  Third, EBTEL has options for heat flux saturation and nonthermal electron beam heating. Finally, EBTEL is unique among 0D models in that it provides the time-dependent differential emission measure distribution of the transition region footpoints. Emission from the transition region plays a critical role in spatially unresolved observations, such as stellar observations and observations of the solar spectral irradiance, which is important for space weather \\citep{lean97}. Note that footpoint emission is not limited to the cooler ($<1$ MK) plasma traditionally associated with the transition region.  It can also include hot emissions that originate from the base of very hot loops.  The so-called moss seen in the ``coronal\" channels of the {\\it Transition Region and Coronal Explorer (TRACE)} is an example \\citep{betal99,mkb00}.   We describe the coronal and transition region parts of EBTEL in the next two sections.  We then present example simulations and compare them with corresponding simulations from a 1D model and, in one case, the Cargill model. We conclude with a discussion of EBTEL and the possible significance of the example simulations. ",
        "conclusions": "As evidenced by these examples, our simple 0D model is an excellent proxy for more sophisticated and far more computationally intensive 1D hydro simulations.  It improves substantially on the 0D models of \\citet{c94}, \\citet{fh90}, and \\citet{kp93}. The Cargill model assumes that heating is instantaneous and that cooling occurs either by thermal conduction or by radiation, but not by both at the same time. The Fisher-Hawley model: (1) predicts abrupt evolutionary changes as the strand evolves between three distinct regimes; (2) does not account for the evaporation that continues well beyond the end of an impulsive heating event; and (3) cannot return to the pre-event state due to unphysical catastrophic cooling.  The Kopp-Poletto model shares some similarities with EBTEL, but it treats the flows in a fundamentally different way.  Like EBTEL, it equates the enthalpy carried by evaporative upflows with an excess heat flux, but the excess is determined relative to the pre-event state, rather than to the time-varying radiative losses from the transition region.  Condensation downflows in the model are given by a density-dependent fraction of the free-fall velocity.  In actuality, gravity plays no direct role in condensation, since the downflows are driven by pressure gradient deficits relative to hydrostatic equilibrium, in the same way that evaporative upflows are driven by pressure gradient excesses. Gravity sets the value of the hydrostatic gradient, but it is only the deficit or excess relative to this value that is important for the flows. Inclined strands experience essentially the same condensation and evaporation as do upright strands of the same length.  Finally, EBTEL has advantages over all three of the other models in that it provides the $D\\!E\\!M(T)$ of the transition region and treats nonthermal electron beams and heat flux saturation.  One obvious application of EBTEL is to investigate the idea that the basic structural elements of the corona are very thin, spatially unresolved magnetic strands that are heated impulsively.  Loops may be bundles of such strands, as reviewed in \\citet{k06}, and the diffuse corona may be similarly structured.  Differential emission measure distributions are one important test of this idea.  Observed $D\\!E\\!M(T)$ curves from active regions and the quiet Sun tend to be peaked near $10^{6.5}$ and $10^{6.1}$ K, respectively, and to have a slope (temperature power law index) $\\ge 0.5$ coolward of the peak \\citep{rd81,dm93,betal96}. This is consistent with the coronal $D\\!E\\!M(T)$ curves of examples 1 and 4 (Figures \\ref{fig:dem1tr} and \\ref{fig:dem4}).  The full loop curves are discrepant, on the other hand, due to the strong contribution from the transition region.  The cited observations were made on the disk and should in principle include the transition region component.  However, it is possible that absorption from chromospheric material such as spicules significantly attenuates the intensities of transition region lines used to construct the $D\\!E\\!M(T)$ curves (e.g., \\citealt{ddg95}; \\citealt{detal99}; \\citealt{df82}; \\citealt{so79}). We are currently investigating the magnitude of this effect.  One of the great mysteries of coronal physics that has come to light in the last few years is the discovery that warm ($\\sim 1$ MK) coronal loops are much denser than expected for quasi-static equilibrium and live for much longer than a cooling time.  The loops are therefore neither steadily heated nor cooling as monolithic structures.  It has been shown that the observed densities and timescales can be explained by bundles of nanoflare heated strands, as long as nanoflares do not all occur at the same time (see \\citealt{k06} and references cited therein).  Neighboring strands will therefore have different temperatures, and loops are predicted to have multi-thermal cross sections.  In particular, emission should be produced at temperatures hotter than 3 MK.  Hot loops are sometimes observed at the locations of warm loops, but not always. Example 5 suggests that nonthermal electron beams are a possible explanation for the lack of hot emission.  As we have discussed, beams can produce excess densities through evaporation without the need for high temperatures.  We have just begun to explore this possibility. For now, we note that the coronal $D\\!E\\!M(T)$ curve of example 5 (Figure \\ref{fig:dem5}) bears a close resemblance to the observed curves reported by \\citet{setal01} for a loop seen above the limb.  %Other problems now being pursued with EBTEL include modeling the %emission characteristics of coronal arcades \\citep{pk07b}, modeling the %light curves of solar flares \\citep{rgm07}, investigating the %resolvability of elemental loop strands \\citep{pk07a}, and modeling coronal %loops as self-organized critical systems \\citep{lk05,kld06}.  In this %last study, loop strands are assumed to become tangled by turbulent %photospheric %convection and to release magnetic energy when the misalignment between %adjacent strands reaches a critical angle, as expected for the secondary %instability \\citep{dka05}.  EBTEL is used to follow the %plasma evolution of the strands.  Preliminary results suggest that %this model can reproduce the light curves of loops observed by the %Soft X-ray Imager (SXI) on the GOES-12 satellite \\citep{lkm07}.  However, %many more simulations with different combinations of model parameters are %necessary before any definitive conclusions can be drawn. This would %be very difficult with a computationally intensive 1D hydro code, but is %no problem with EBTEL.  %Another major program we have begun under the partial sponsorship of %NASA\u0092s Living With a Star Program is to build realistic %models of active regions and the global Sun.  Our approach is to %construct a ``magnetic skeleton\" %by extrapolating photospheric magnetograms and to populate many representative %field lines with plasma using EBTEL.  Models of this kind have been %published before, but they use static equilibrium strand solutions and cannot %adequately reproduce imaging observations over %a wide range of temperatures \\citep{letal04, setal04, metal05, ww06, bw07}.  %When the models are adjusted to %resemble soft X-ray images, they %they predict too little warm ($\\sim 1$ MK) emission at coronal altitudes %(EUV loops) and too much warm emission at the transition %region footpoints of hot coronal structures (EUV moss).   There is %reason to believe that impulsive heating can improve the situation %\\citep{kld06, pk07b}.  Indeed, \\citet{ww07} have built an impressive %model of an active region using the NRLFTM 1D hydro code and %find that impulsive heating gives a significant improvement over %steady heating.  The agreement with observations is still not adequate, %however.  %Since 1D simulations are computationally very intensive, %Warren and Winebarger were only able to consider one parameterized form for the %coronal heating function.  Using EBTEL, we will examine a wide range of %heating parameters, including %the magnitude, duration, and frequency of the nanoflares, and their dependence %on magnetic field strength and field line length %(see \\citealt{mdk00}).  We will also consider different proportions of %direct heating and nonthermal electron beams. %We estimate that 100 million hydro simulations %but is a manageable task with EBTEL.  Identifying the coronal heating %parameters will provide valuable constraints for testing competing %theories of the heating mechanism.  It should ultimately lead to a %physics-based operational model for nowcasting and %forecasting the X-ray and UV spectral irradiance.  In conclusion, EBTEL is a powerful new tool that can be applied to a variety of problems in which large numbers of evolving strands must be computed. For example, it is now feasible to construct multiple models of nanoflare-heated active regions or entire stars and therefore to examine a wide array of nanoflare parameters (magnitude, lifetime, occurrence rate, dependence on field strength and strand length, etc.). By determining which parameters best reproduce the observations, we can place important constraints on the heating and thereby gain insight into the physical mechanism (e.g., \\citealt{mdk00}; \\citealt{setal04}; \\citealt{ww06}).  EBTEL is currently being used to study the emission characteristics of coronal arcades \\citep{pk07b}, to explain the light curves of solar flares \\citep{rgm07}, and to model coronal loops as self-organized critical systems \\citep{lk05,kld06}.  Interested users are invited to contact us for a copy of our IDL code.  %We end by noting possible future improvements: %the inclusion of kinetic energy, cross-sectional area variation, and %gravitational stratification.  %% Included in this acknowledgments section are examples of the %% AASTeX hypertext markup commands. Use \\url without the optional [HREF] %% argument when you want to print the url directly in the text. Otherwise, %% use either \\url or \\anchor, with the HREF as the first argument and the %% text to be printed in the second."
    },
    "0710/0710.0373_arXiv.txt": {
        "abstract": "The physics behind the acceleration of the cosmic expansion can be elucidated through comparison of the predictions of dark energy equations of state to observational data.  In seeking to optimize this, we investigate the advantages and disadvantages of using principal component analysis, uncorrelated bandpowers, and the equation of state within redshift bins.  We demonstrate that no one technique is a panacea, with tension between clear physical interpretation from localization and from decorrelated errors, as well as model dependence and form dependence. Specific lessons include the critical role of proper treatment of the high redshift expansion history and the lack of a unique, well defined signal-to-noise or figure of merit. ",
        "introduction": "}  The acceleration of the universe poses a fundamental mystery to cosmology, gravitation, and quantum physics.  Understanding the nature of the dark energy responsible for the acceleration relies on careful, robust measurements of the dark energy properties, in particular its equation of state (EOS), or pressure to energy density, ratio that directly enters the Friedmann equation for cosmic acceleration.  As scientists design the next generation of dark energy experiments they seek to optimize the measurements for the clearest insight into this unknown physics.  Two critical pieces of information will be the value of the EOS at some epoch, such as the present, and a measure of its time variation, in much the way that early universe inflation theories are classified by the value of the spectral index and its running.  The best parametrized EOS are physics based and model independent, i.e.\\ able to describe dark energy dynamics globally, or at least over a wide range of behaviors.  Such EOS are very successful at fitting to data and projecting the results of future experiments, and can be robust to bias against inexact parametrization.  Other approaches seek to remove one drawback of parametrized EOS by not assuming a functional form for the time variation, lest the true dark energy model lie outside the apparently wide range of validity of the form, i.e.\\ they aim for form independence.  Two major avenues for achieving this are decomposition into basis functions or principal components (e.g.\\ \\cite{HutStark03}, also see \\cite{CritPog05,ShapTurn06,SimpBrid06,DickKnoxChu06,Ste06,HutPeir07}) and individual values of the EOS $w(z)$ over finite redshift bins, which become more general as the number of elements increases.  However uncertainties in estimation of the EOS properties also grow as the number of principal components or bins increases.  This article begins by examining general properties of the cosmological data and its dependence on the EOS in \\S\\ref{sec:cosdep}.  Many of the later, detailed results will already be foreshadowed by this straightforward and general analysis.  In \\S\\ref{sec:pca} we examine principal component analysis of the EOS and in \\S\\ref{sec:band} uncorrelated bandpowers. Bins of EOS in redshift is investigated in \\S\\ref{sec:bin}, including figures of merit for quantifying the uncertainties.  Further concentration on the crucial role of the high redshift EOS, and the risk of biasing parameter estimation, occurs in \\S\\ref{sec:whi}.  We consider physical constraints on EOS properties in \\S\\ref{sec:wlimit} and summarize our results and conclude in \\S\\ref{sec:concl}. ",
        "conclusions": "}  The dark energy equation of state properties contain clues crucial to understanding the nature of the acceleration of the cosmic expansion. Deciphering those properties from observational data involves a combination of robust analysis and clear interpretation.  We considered three approaches -- principal components, uncorrelated bandpowers, and binning; none of the approaches provides a panacea.  In particular, we identify issues of dependence on basis functions, binning variables, and baseline models.  The three approaches are not truly nonparametric and physical interpretation (not merely the values) of the results in the two decorrelated basis techniques depends on model, priors, and data, indeed even on an implicitly assumed functional form. Nevertheless, principal components can give a useful guide to the qualitative sensitivity, the best constrained aspects, of the data.  The uncorrelated bin approach unfortunately does not truly deliver uncorrelated bandpowers for the equation of state.  This approach using the square root of the Fisher matrix does not tightly localize the information (without a strong prior), making the interpretation nontrivial. This property of nonlocality is inherent in the cosmological characteristics.  One might prefer to stay with the original binned equations of state used as the initial step for this technique, which are readily interpreted.  Conversely, if the modes can be localized, the interpretation is easy, but in that case the original Fisher matrix is close to diagonal and thus the original bins almost uncorrelated.  Hence, again, one might as well stay with the bin parameters which have a clear meaning.  Indeed the goal is understanding the physics, not obtaining particular statistical properties.  Decorrelated parameters that are not readily interpretable physically are of limited use; for example one still prefers to analyze the cosmic microwave background in terms of physical quantities such as physical matter density and spectral tilt rather than the principal axes of the eigenvectors. Note that the uncertainty on the EOS behavior $\\sigma(w(z))$ is the same whether calculated by PCA (if all modes are kept), uncorrelated bands, or binned EOS, since the same information is in the data. We also emphasize that the modes most clearly determine the effect on the equation of state, not the weights, which are often the only quantity displayed.  Moderately localized, even all positive, weights do not guarantee a localized physical effect.  A further caution is that locality and positivity of weights can owe more to prior restrictions, especially the treatment of the high redshift equation of state, than to the data itself.  Assuming a fixed value for the high redshift equation of state has major, widespread impacts on the results, ranging from strongly misestimated uncertainties to spurious localization to bias in the derived cosmology. We emphasize that it is essential to fit for the high redshift behavior in order not to be misled.  Adding CMB data and marginalizing over a new, high redshift bin removes the ill effects of bias but ``cancels out'', providing no new constraints; multiple data points for $z>2$ are required, such as from high redshift distances or weak lensing measurements of the mass growth behavior. Assuming that dark energy is negligible at $z>2$ is also effectively assuming a functional form -- precisely what the use of eigenmodes was supposed to avoid.  Indeed, functional forms do not have many of the basis, model, binning, etc.\\ dependences of eigenmodes, while principal components are in turn not fully form independent. If one assumes a functional form to obtain informative constraints on the equation of state,  one must indeed choose the form to represent robustly the physical behavior (as has been shown to be widely the case for $w(a)=w_0+w_a(1-a)$ by \\cite{Linder03,Linder0708}), and carefully check the range of validity of the conclusions by examining other forms.  A good complementary analysis tool would be the binned equation of state approach examined here.  Regardless of the form of analysis, only a finite amount of information can be extracted from even next generation data.  As has been concluded for functional equations of state and principal component analysis \\cite{LinHut05}, the analysis here in terms of binned equation of state indicates that only two physically informative parameters can be fit with realistic accuracy.  However, we identify several issues in the PCA and uncorrelated bin approaches that cause accuracy or signal to noise criteria to be ill defined.  Similar difficulties arise in condensing the physical information on dark energy to a single figure of merit; the number is quite sensitive to cosmologically irrelevant aspects like the binning used (as well as very dependent on the treatment of the high redshift dark energy behavior).  In conclusion, physically motivated fitting of the equation of state such as the $w_0$-$w_a$ parametrization in complement with a binned equation of state approach (perhaps with physical constraints such as outlined in \\S\\ref{sec:wlimit}) have the best defined, clearest to interpret, and robust insights of the approaches we considered. With any method, one must use caution regarding the influence of priors and fit the dark energy physics over the entire expansion history."
    },
    "0710/0710.0780_arXiv.txt": {
        "abstract": "$Chandra$ and $XMM-Newton$ observations of the Cartwheel galaxy show $\\sim{}17$ bright X-ray sources ($\\gtrsim{}5\\times{}10^{38}$ erg s$^{-1}$), all within the gas-rich outer ring. We explore the hypothesis that these X-ray sources are powered by intermediate-mass black holes (IMBHs) accreting gas or undergoing mass transfer from a stellar companion. To this purpose, we run $N$-body/SPH simulations of the galaxy interaction which might have led to the formation of Cartwheel, tracking the dynamical evolution of two different IMBH populations: halo and disc IMBHs. Halo IMBHs cannot account for the observed X-ray sources, as only a few of them cross the outer ring. Instead, more than half of the disc IMBHs are pulled in the outer ring as a consequence of the galaxy collision. However, also in the case of disc IMBHs, accretion from surrounding gas clouds cannot account for the high luminosities of the observed sources. Finally, more than 500 disc IMBHs are required to produce $\\lesssim{}15$ X-ray sources via mass transfer from very young stellar companions. Such number of IMBHs is very large and implies extreme assumptions. Thus, the hypothesis that all the observed X-ray sources in Cartwheel are associated with IMBHs is hardly consistent with our simulations, even if it is still possible that IMBHs account for the few ($\\lesssim{}1-5$) brightest ultraluminous X-ray sources (ULXs). ",
        "introduction": "Ring galaxies are among the most fascinating objects in the Universe. The Cartwheel galaxy is surely the most famous of them, and also the biggest (with its optical diameter of $\\sim{}$50-60 kpc) and the most studied. Cartwheel has been thoroughly observed in almost every band: H$\\alpha{}$ (Higdon 1995) and optical (Theys \\& Spiegel 1976, Fosbury \\& Hawarden 1977) images, red continuum (Higdon 1995), radio line (Higdon 1996) and continuum (Higdon 1996; Mayya et al. 2005), near- (Marcum, Appleton \\& Higdon 1992) and far-infrared (Appleton \\& Struck-Marcell 1987a), line spectroscopy (Fosbury \\& Hawarden 1977) and X-ray (Wolter, Trinchieri \\&{} Iovino 1999; Gao et al. 2003; Wolter \\& Trinchieri 2004; Wolter, Trinchieri \\& Colpi 2006).    Cartwheel exhibits a double ringed shape, with some transversal 'spokes' (Theys \\& Spiegel 1976, Fosbury \\& Hawarden 1977; Higdon 1995), which have been detected only in few of the $\\lesssim{}300$ known ring galaxies (Arp \\& Madore 1987; Higdon 1996). The outer ring is rich of HII regions, especially in its southern quadrant, and dominates the H$\\alpha{}$ emission (Higdon 1995). This implies that Cartwheel is currently  undergoing an intense epoch of star formation (SF, with SF rate $\\sim{}20-30\\,{}M_\\odot{}$ yr$^{-1}$; Marston \\& Appleton 1995; Mayya et al. 2005), almost entirely confined to the outer ring. Inner ring, nucleus and spokes are nearly devoid of gas and dominated by red continuum emission, indicating a  relatively old, low- and intermediate-mass stellar population (Higdon 1995, 1996; Mayya et al. 2005). Cartwheel is located in a small group of 4 galaxies. All of its 3 companions (known as G1, G2 and G3) are less massive than Cartwheel, and host less than 20 per cent of the total gas mass in the group.   The analysis of X-ray data shows another intriguing peculiarity: most of the point sources detected by $Chandra$ are in the outer ring, and particularly concentrated in the southern quadrant (Gao et al. 2003; Wolter \\& Trinchieri 2004). According to Wolter \\& Trinchieri (2004) 13 out of 17 sources associated with Cartwheel are in the outer ring, the remaining 4 being close to the inner rim of the ring or to the optical spokes\\footnote{The total number of point X-ray sources in the Cartwheel group is 24 (Wolter \\& Trinchieri 2004), but 6 of them are associated with the companion galaxies G1 and G2, whereas 1 of them is probably background/foreground contamination.}. Gao et al. (2003) noted that all the five strongest H$\\alpha{}$ knots in the ring are associated with an X-ray source, indicating a possible correlation between X-ray sources and young star-forming regions. Furthermore, most of the observed sources have isotropic X-ray luminosity $L_X\\gtrsim{}10^{39}$ erg s$^{-1}$, belonging to the category of ultraluminous X-ray sources (ULXs). Given the distance of Cartwheel ($\\sim{}124$ Mpc for an Hubble constant $H_0=73$ km s$^{-1}$ Mpc$^{-1}$), the data are incomplete for $L_X\\lesssim{}5\\times{}10^{38}$ erg s$^{-1}$, and the faintest sources in the sample have $L_X\\sim{}10^{38}$ erg s$^{-1}$. Then, almost all the X-ray sources detected in Cartwheel are close to the ULX range.     Many theoretical studies, both analytical (Lynds \\& Toomre 1976; Theys \\& Spiegel 1976; Appleton \\& Struck-Marcell 1996) and numerical (Theys \\& Spiegel 1976; Appleton \\& Struck-Marcell 1987a, 1987b; Hernquist \\& Weil 1993; Mihos \\& Hernquist 1994; Struck 1997; Horellou \\& Combes 2001; Griv 2005) were aimed to explain the origin of Cartwheel and its observational features. In the light of these papers, the origin of propagating rings in Cartwheel and similar galaxies can be explained by galaxy collisions with small impact parameter (Theys \\& Spiegel 1976; Appleton \\& Struck-Marcell 1987a, 1987b; Hernquist \\& Weil 1993; Mihos \\& Hernquist 1994; Struck 1997; Horellou \\& Combes 2001). Among the companions of Cartwheel, both G1 and G3 are good candidates for this interaction, having a small projected impact parameter, being not too far from Cartwheel, and showing a disturbed distribution of neutral hydrogen (Higdon 1996). An alternative, less investigated model explains the rings with disc instabilities (Griv 2005).  Models of galaxy collisions explain quite well most of Cartwheel properties. However, none of the previous theoretical studies has investigated the nature of the X-ray sources, and especially of the ULXs, observed in Cartwheel.  The nature of the ULXs is still not understood. It has been suggested that they are  high-mass X-ray binaries (HMXBs) powered by  stellar mass black holes (BHs) with anisotropic X-ray emission (King et al. 2001; Grimm, Gilfanov \\& Sunyaev 2003; King 2006)  or with super-Eddington accretion rate (e.g. King \\&{} Pounds 2003; Socrates \\& Davis 2006; Poutanen et al. 2007). However, some ULXs, especially the brightest ones ($L_X\\gtrsim{}10^{40}$ erg s$^{-1}$), show characteristics which are difficult to reconcile with  the hypothesis of beamed emission, such as the presence of an isotropically ionized nebula (e.g. Pakull \\& Mirioni 2003; Kaaret, Ward \\& Zezas 2004) or of quasi periodic oscillations (e.g. Strohmayer \\& Mushotzky 2003). Then, it has been proposed that  some ULXs (or at least the brightest among them; Miller, Fabian \\& Miller 2004) might be associated with intermediate-mass black holes (IMBHs), i.e. BHs with mass $20\\lesssim{}m_{\\rm BH}/M_\\odot{}\\lesssim{}10^{5}$, accreting either gas or companion stars (Miller et al. 2004; see Mushotzky 2004 and Colbert \\& Miller 2005 for a review). On the other hand, most of ULXs can be explained with the properties of stellar mass  BHs (see Roberts 2007 for a review and references therein).  In this paper we present a new, refined $N-$body/SPH model of the Cartwheel galaxy, which reproduces the main features of Cartwheel. The aim of this paper is to use the $N-$body/SPH model in order to check whether the IMBH  hypothesis is viable to explain all or a part of the X-ray sources observed in Cartwheel. In fact, in the last few years the hypothesis that most of the ULXs are powered by  IMBHs has became increasingly difficult to support, as the observational features of the majority of ULXs are consistent with accreting stellar mass BHs (Roberts 2007). It would be interesting to see whether also the dynamical simulations agree with this conclusion\\footnote{We will not consider other possible scenarios (e.g. the production of ULXs by beamed emission in HMXBs), due to intrinsic limits of $N-$body methods.}.  In Section 2 we describe our simulations. In Section 3 we discuss the evolution of our models and the dynamics of either halo or disc IMBHs. In Section 4 we investigate the possibility that the X-ray sources in Cartwheel are powered by IMBHs accreting gas or stars. In the last case we consider both the accretion from old stars (which produce only transient sources) and from young stars, formed after the galaxy collision (which can lead also to persistent sources). We also investigate the hypothesis that new disc IMBHs are formed in very young clusters after the galaxy collision. Our findings are summarized in Section 5. ",
        "conclusions": "In this paper we investigated the possible connection among IMBHs and the $\\sim{}17$ (Gao et al. 2003; Wolter \\& Trinchieri 2004) bright X-ray sources detected in the outer ring of Cartwheel.  Recent observations show that models based on beamed emission or super-Eddington accretion in HMXBs including stellar mass BHs can explain most of ULXs, apart from the brightest ones (Roberts 2007 and references therein). However, the observations cannot definitely exclude that all the ULXs are powered by IMBHs. Thus, in our paper we checked whether the IMBHs can account for all or only for a part of the ULXs observed in Cartwheel.   We simulated the formation of a Cartwheel-like ring galaxy via dynamical interaction with an intruder galaxy. In this simulation we also integrated the evolution of 100 IMBHs particles.  We considered two different models of IMBH formation, i.e. IMBHs born as relics of population III stars (and distributed as a concentrated halo population) and IMBHs formed via runaway collapse of stars (and distributed as an exponential disc). For these models, we investigated both gas accretion from surrounding molecular clouds and mass transfer from a stellar companion. The main results of this study are the following. \\subsection{Halo IMBHs} IMBHs born as the relics of population III stars, if they are distributed as a halo population, cannot contribute to the X-ray sources, neither via gas accretion  nor via mass transfer in binaries. In particular, the luminosity produced by halo IMBHs accreting gas is always many orders of magnitude smaller than that of the observed sources, even if we assume that halo IMBHs have a large mass ($m_{\\rm BH}=10^3\\,{}M_\\odot{}$) and a high radiative efficiency ($\\eta{}=0.1$). This is due to the fact that only a small fraction of halo IMBHs passes through the disc (only for a short lapse of time), and even these IMBHs have a high ($v\\gtrsim{}100$ km s$^{-1}$) relative velocity with respect to gas particles. Similarly, halo IMBHs cannot accrete mass from stars, as the probability that they acquire a companion is very low. \\subsection{Disc IMBHs} IMBHs born from the runaway collapse of stars should be a disc population. Under overoptimistic assumptions ($n_g=10^2\\,{}{\\rm cm}^{-3}$, $m_{\\rm BH}=10^3\\,{}M_\\odot{}$ and $\\eta{}=0.1$) these IMBHs can produce, via gas accretion, X-ray sources with $L_X\\lesssim{}10^{39}$ erg s$^{-1}$,  a factor of $\\sim{}10$ fainter than the brightest ULXs observed in Cartwheel. Thus, also disc IMBHs accreting gas cannot explain the observed X-ray sources in Cartwheel. Our model of gas accreting IMBHs contains many rough assumptions; but most of them are upper limits, strengthening our result. However, a more realistic treatment of the local properties of the gas would be helpful, to understand the physical mechanisms of gas accretion onto IMBHs.  On the other hand, runaway collapse born IMBHs are hosted in dense young clusters. In such environment, it is easy for the IMBH to capture a stellar companion. Blecha et al. (2006) estimated that a $\\sim{}100\\,{}M_\\odot{}$ IMBH undergoes mass transfer from a companion star for about the $3$ per cent of the cluster lifetime. Previous papers (Portegies Zwart et al. 2004; Patruno et al. 2005) have shown that IMBHs accreting from low-mass ($\\lesssim{}10\\,{}M_\\odot{}$) and high-mass ($\\gtrsim{}10\\,{}M_\\odot{}$) companions  generate transient (with a bright phase of only a few days) and persistent bright X-ray sources, respectively. As $10\\,{}M_\\odot{}$ stars have a lifetime of $\\sim{}30-40\\,{}$ Myr, only IMBHs hosted in sufficiently young star clusters can generate persistent X-ray sources. Then, IMBHs hosted in clusters born before the dynamical encounter with the intruder (i.e. more than 100 Myr ago) can produce only transient sources.  We estimated that, out of 100 IMBHs which were present before the dynamical encounter with the intruder galaxy, only $\\lesssim{}2-3$ are expected to undergo mass transfer from low-mass companions at present, producing a comparable number of transient X-ray sources. As observations show that at least 4 X-ray sources in the Cartwheel ring are persistent over a time-scale of 6 months (Wolter et al. 2006), we conclude that pre-encounter formed IMBH binaries are not sufficient to explain the data.   We considered the possibility that pre-encounter IMBHs capture massive stars produced after the encounter with the intruder. In this case, under overoptimistic assumptions, 100 pre-encounter IMBHs can produce $\\lesssim{}1$ X-ray source, either persistent or not.  Finally, we hypothesized that very young ($<40$ Myr) star clusters, formed after the encounter, generate IMBHs at their centre. Under this hypothesis, $500-1000$ IMBHs are required to produce $\\sim{}15-30$ bright ($10^{36}-10^{41}$ erg s$^{-1}$) X-ray sources, some of them persistent and  some transient% . This scenario might account for the $\\sim{}17$  observed X-ray sources in the Cartwheel ring. It is also in agreement with the fact that many ULXs observed in Cartwheel are associated with bright $H_\\alpha{}$ spots, i.e. active star-forming regions (Gao et al. 2003).   The birth of $\\sim{}1000$ IMBHs (each one of 100 $M_\\odot{}$) in $\\sim{}$40 Myr  implies an IMBH formation rate of $2.5\\times{}10^{-3}\\,{}M_\\odot{}\\,{}{\\rm yr}^{-1}$, that is a factor of $\\sim{}10^4$ lower than the SF rate. This rate is acceptable for runaway collapse scenarios, as Portegies Zwart \\& McMillan (2002) show that  $\\sim{}$0.1 per cent of the mass of the parent young cluster merges to form the IMBH.  However, we stress that only the runaway collapse scenario, under extreme assumptions, can explain the formation of such a huge number of IMBHs in the disc. Thus, our simulations suggest that IMBHs can hardly account for all the ULXs observed in Cartwheel.    On the other hand, it is  possible that only the few brightest sources in Cartwheel  are powered by IMBHs, while the other ones are either beamed HMXBs, super-Eddington accreting stellar mass BHs  or a blending of multiple fainter sources. For example, $\\sim{}30$ IMBHs are expected to form in the ring, in order to produce only 1  very bright ULX, such as the source N.10 in Cartwheel.  These results agree with the semi-analytical model by King (2004), who showed that IMBHs cannot explain all the ULXs in Cartwheel. However, King (2004) concludes that $>3\\times{}10^4$ IMBHs are required to produce the observed number of ULXs, $\\sim{}30-60$ times more than in our analysis. This apparent discrepancy is due to the fact that  King (2004) assumes that the IMBHs power only transient sources (Kalogera et al. 2004), and thus he has to introduce a $\\sim{}10^{-2}$ duty-cycle. However, Patruno et al. (2005) showed that IMBHs accreting from young massive stars ($\\gtrsim10\\,{}M_\\odot{}$) produce non-transient sources, increasing the expected duty-cycle.    In conclusion, new $Chandra$ and $XMM-Newton$ observations of Cartwheel could partially  solve the mystery of Cartwheel X-ray sources, investigating which sources are transient, which variable, and which persistent. Deeper observations are also needed  to resolve possible blended sources.  In the future, it would be interesting to search whether other ring galaxies host  as many bright X-ray sources as Cartwheel and whether these sources are similarly concentrated in the outer ring."
    },
    "0710/0710.0749_arXiv.txt": {
        "abstract": "{}{We investigate the structure of a field around the position of V838~Mon as seen in the lowest CO rotational transitions. We also measure and analyse emission in the same lines at the position of V838~Mon.}{Observations have primarily been done in the $^{12}$CO $J = 2$$\\to$$1$ and $J=3$$\\to$$2$ lines using the KOSMA telescope. A field of 3.4 squared degrees has been mapped in the on-the-fly mode in these transitions. Longer integration spectra in the on-off mode have been obtained to study the emission at the position of V838~Mon. Selected positions in the field have also been observed in the  $^{12}$CO $J = 1$$\\to$$0$ transition using the Delingha telescope.}{In the observed field we have identified many molecular clouds. They can be divided into two groups from the point of view of their observed radial velocities. One, having $V_{\\rm LSR}$ in the range $18-32\\,{\\rm km\\,s}^{-1}$, can be identified with the Perseus Galactic arm. The other one, having $V_{\\rm LSR}$ between $44-57\\,{\\rm km\\,s}^{-1}$, probably belongs to the Norma-Cygnus arm. The radial velocity of V838~Mon is within the second range but the object does not seem to be related to any of the observed clouds. We did not find any molecular buble of a $1\\degr$ dimension around the position of V838~Mon claimed in van~Loon et~al. An emission has been detected at the position of the object in the $^{12}{\\rm CO}\\ J=2$$\\to$$1$ and $J=3$$\\to$$2$ transitions. The emission is very narrow (FWHM~$\\simeq$~1.2~km\\,s$^{-1}$) and at $V_{\\rm LSR} = 53.3$~km\\,s$^{-1}$. Our analysis of the data suggests that the emission is probably extended.}{} ",
        "introduction": "}  The outburst of V838 Mon was discovered in the beginning of January~2002. Initially thought to be a nova, the object appeared unusual and enigmatic in its nature. The eruption, as observed in the optical, lasted about three months (Munari et~al. \\cite{muna02}, Kimeswenger et~al. \\cite{kimes02}, Crause et~al. \\cite{crause03}). After developing an A--F supergiant spectrum at the maximum at the beginning of February~2002, the object evolved to lower effective temperatures and in April~2002 it practically disappeared from the optical, remaining very bright in the infrared. A detailed analysis of the evolution of the object in the outburst and decline can be found in Tylenda (\\cite{tyl05}).  Different outburst mechanisms have been proposed to explain the eruption of V838~Mon. They include an unusual nova (Iben \\& Tutukov \\cite{it92}), a late He-shell flash (Lawlor \\cite{law05}) and a stellar merger (Soker \\& Tylenda \\cite{soktyl03}). These models have critically been discussed in Tylenda \\& Soker (\\cite{tylsok06}) and the authors conclude that the only mechanism that can satisfactorily account for the observational data is a collision and merger of a low-mass pre-main-sequence star with an $\\sim$$8\\,M_\\odot$ main-sequence star.  V838 Mon is surrounded by diffuse matter which gave rise to a spectacular light-echo phenomenon (e.g. Bond et~al. \\cite{bond03}). Bond et~al. claim that the matter comes from previous eruptions of the object. However Tylenda (\\cite{tyl04}), Crause et~al. (\\cite{crause05}), as well as Tylenda, Soker \\& Szczerba (\\cite{tss05}) argue that the echoing matter is of interstellar origin. The latter is consistent with recent findings, namely that V838~Mon is a member of a young open cluster (Afsar \\& Bond \\cite{afsar}) and that the total mass of the echoing matter is probably of the order of $100\\ M_\\odot$ (Banerjee et~al. \\cite{baner06}).  van Loon et al. (\\cite{loon04}) have analyzed archive infrared and radio data on the sky around V838~Mon and claimed discovery of multiple shells ejected by the object in the past. In particular, from a compilation of CO galactic surveys done in Dame et~al. (\\cite{dame01}) van~Loon et~al. have suggested that V838~Mon is situated within a bubble of CO emission with a diameter of $\\sim$$ 1\\degr$. These results have been critically discussed in Tylenda et~al. (\\cite{tss05}), who have concluded that the shells of van~Loon et~al. are not realistic.  Deguchi et~al. (\\cite{degu05}) have discovered an SiO maser emission from V838~Mon. The main component is narrow and centered at $V_{\\rm LSR} \\simeq 54\\ {\\rm km}\\,{\\rm s}^{-1}$, which is thought to be a radial velocity of the object itself.  Further observations of Claussen et~al. (\\cite{clauss06}) have shown that the SiO maser is variable and that most of the emission comes from a region smaller than a milliarcsecond.  In the present paper we report on results of our observations of V838~Mon and its nearby vicinity in the $^{12}$CO $J = 1$$\\to$0, 2 $\\to$1 and 3$\\to$2 transitions. We describe the observational material and discuss results on the CO emission from a field of 3.4 squared degrees around the position of V838~Mon. One of the goal of this survey is to verify the existence of the CO bubble around V838~Mon claimed in van~Loon et~al. (\\cite{loon04}). Measurements of the CO emission obtained at the position of V838~Mon are also presented, analysed and discussed. A preliminary analysis of the data has been done in Kami\\'{n}ski et~al. (\\cite{kmst}). ",
        "conclusions": "}  Our on-the-fly maps in the CO rotational transitions, especially those done in the (2$\\to$1) line with the KOSMA telescope, have allowed us to identify 29 molecular clouds around the position of V838~Mon (see Appendix~\\ref{append}). As far as we know, the region mapped with the KOSMA telescope has not been observed before with a sensitivity and a spatial resolution comparable to our survey. The data presented in  Appendix~\\ref{append} can be used in studies of molecular matter in the outer Galaxy.  We do not confirm the existence of a molecular bubble of a 1 degree dimension claimed in van~Loon et~al. (\\cite{loon04}), which was probably an artefact resulted from merging two surveys of different resolution and sensitivity in the data used by van~Loon et~al.  Our maps did not reveal any molecular cloud in the near vicinity of V838~Mon. The nearest CO emission has been detected at $\\sim$$40$~arcmin from the position of the object (see Fig.~\\ref{map_2}). Thus the nearest dense molecular cloud is located at least at a distance of $\\sim$$80$~pc from V838~Mon (assuming a 7~kpc distance to V838~Mon -- see below). From the noise level in the (2$\\to$1) map (see Table~\\ref{tech_tab}) we can put an upper limit to $I_{\\rm CO}=\\int T_{\\rm mb}\\,dV$ of $1.26~{\\rm K~km~s}^{-1}$ ($3 \\sigma_{\\rm rms}$) at the position of V838~Mon. Assuming the $N$(H$_2$) to $I_{\\rm CO}$ conversion factor of $X_{\\rm CO} = 5.1$ (in units $10^{20}\\,{\\rm molecules}\\ {\\rm cm}^{-2}\\,{\\rm K}^{-1}\\,{\\rm km}^{-1}\\,{\\rm s}$, typical value for the NC clouds discussed in Appendix~\\ref{append}) we obtain an upper limit to the column density of $N({\\rm H}_2) = 6.4 \\times 10^{20}~{\\rm cm}^{-2}$. This upper limit is an order of magnitude lower than typical column densities observed for molecular clouds in the vicinity of the Sun (e.g. Blitz \\cite{blitz}). Thus we can conclude that V838~Mon is not located in a typical molecular cloud.  However, as presented in Sect.~\\ref{kosma}, long integrations in the on-off mode have allowed us to detect an emission in the CO (2$\\to$1) and (3$\\to$2) transitions at the position of V838~Mon. The question that arises is: what is the origin of this emission? Certainly it cannot come from the matter ejected in the 2002 outburst. That matter was expanding with large velocities, a few hundred km~s$^{-1}$, while our profiles have a FWHM of $\\sim$$1\\,{\\rm km\\,s}^{-1}$.  The $V_{\\rm LSR}$ position and width of the CO lines similar to those of the SiO maser (Deguchi et~al. \\cite{degu05}, Claussen et~al. \\cite{clauss06}) suggest that both emissions (CO and SiO) originate from the same place, e.g from the remnant of the 2002 outburst. However, in this case the KOSMA telescope would see a point source, which does not seem to be the case. The April 2006 observations show detectable emissions at the 1~arcmin offsets from the V838~Mon position (see Table~\\ref{res2_tab}). The HPBW of the KOSMA beam at 230~GHz is 130~arcsec. If the CO emission source were a point source at the (0,0) position, than the intensity measured at the (0,1) and ($0,-1$) offsets would be about twice fainter than that at the central position. Table~\\ref{res2_tab} shows that this ratio is larger, i.e. $0.6-0.8$. Unfortunantely the accuracy of the measurements at the 1~arcmin offsets was not high, so we cannot say that the results in Table~\\ref{res2_tab} are definitively inconsistent with a point source emission.  However, there are other arguments in favour of the idea that the CO emission is extended and/or situated outside the V838~Mon position. Our Delingha observations in the (1$\\to$0) transition (see Sect.~\\ref{delingha}) have not detected any emission at the V838~Mon position and the upper limit was 0.39~K. Assuming that the intensity in the (1$\\to$0) line is comparable to that in the (2$\\to$1) line (which is usually the case in molecular clouds) and that we observe a point source at the V838~Mon position than from the measured $T_{\\rm mb}$ in the (2$\\to$1) transition in December~2005 (see Table~\\ref{res_tab}), taking into account different beamsizes of the two telescopes (55~arcsec in Delingha versus 130~arcsec in KOSMA), the expected value of $T_{\\rm mb}$ in the (1$\\to$0) Delingha observations would be $\\sim$$1.8$\\,K. This is 4.5 times higher than the observed upper limit. Thus either the (2$\\to$1)/(1$\\to$0) ratio is exceptionally large ($> 4.5$) or the CO source is situated outside the Delingha beam but still inside the KOSMA beam. The later interpretation is supported by our possible detection of an emission with Delingha at three positions $\\sim$$1$\\,arcmin from the object (see Table~\\ref{resd_tab}). It is also supported by a finding of Deguchi et~al. (\\cite{degu06}) who, using the Nobeyama telescope, tried to measure the CO (1$\\to$0) emission from a few points in a field around V838~Mon. A signal was detected from a position 30~arcsec north of V838~Mon at a velocity very close to that of the (2$\\to$1) line in our KOSMA observations. No emission was however recorded at the position of the object.  The above discussion allows us to conclude that the CO emission, clearly seen in our KOSMA observations, most probably originates not from V838~Mon itself but from a region (regions) situated, typically, $\\sim$$1$\\,arcmin from the V838~Mon position. There are two possible ways of explaining the emission in this case. One is that the observed emission is a faint part of a larger CO structure belonging, for instance, to a molecular cloud. However, as discussed above, our CO maps have not revealed any dense CO cloud of similar $V_{\\rm LSR}$ closer than $\\sim$$40$\\,arcmin from the position of V838~Mon.  All the detected and possibly detected CO emission at and near the position of V838~Mon lie well within the optical light echo of V838~Mon. Hence the second possibility, namely that the CO emission comes from the same matter that gave rise to the light echo and that it was the light flash from the 2002 eruption which induced the emission. Then the observed narrowness of the line profile would be a strong argument that the matter is of interstellar origin rather than being ejected by V838~Mon in previous eruptions. The observed $V_{\\rm LSR} = 53.3\\,{\\rm km\\,s}^{-1}$ of the CO emission would imply a kinematical distance of $\\sim$$7.0$\\,kpc (using the Galactic rotational curve of Brandt \\& Blitz \\cite{bb93}). This can be compared to $6.1 - 6.2$~kpc found by Bond et~al. (\\cite{bond06}) and $\\sim$$8$~kpc advocated in Tylenda (\\cite{tyl05}).  Rushton et al. (\\cite{rush03}) searched for CO emission from V838~Mon about a year after the 2002 eruption. No measurable signal was detected in the three lowest rotational transitions and the upper limit was $T_A^* \\la 25-40$~mK. It is not straightforward to compare this result with our findings as the observations of Rushton et~al. were done about 3~years before ours and the object probably evolved significantly during this time span. Nevertheless, given the beamwidth of the telescopes used by Rushton et~al. (HPBW~$\\le 45\\,\\arcsec$), their negative result at the position of the object does not preclude a possibility that a significant emission was present in a near vicinity of the object, but outside the telescope beam.  More sensitive observations with a higher spatial resolution are required to distinguish between the above discussed possibilities and to futher investigate the problem of the CO emission from V838~Mon."
    },
    "0710/0710.5393_arXiv.txt": {
        "abstract": "The Zeeman pattern of Mn~{\\sc i} lines is sensitive to  hyperfine structure (HFS) and, because of this reason, they respond to hG magnetic field strengths differently from the lines commonly used in solar magnetometry. This peculiarity has been employed to measure magnetic field strengths in quiet Sun regions. However, the methods applied so far assume the magnetic field to be constant in the resolution element. The assumption is clearly insufficient to describe the complex quiet Sun magnetic fields, biasing the results of the measurements. The diagnostic potential of Mn~{\\sc i} lines can be fully exploited only after understanding the sense and the magnitude of such bias. We present the first syntheses of Mn~{\\sc i} lines in realistic quiet Sun model atmospheres. Plasmas varying in magnetic field strength, magnetic field direction, and velocity, contribute to the synthetic polarization signals. The syntheses show how the Mn~{\\sc i} lines weaken with increasing field strength. In particular, kG magnetic concentrations produce \\mni{5538} circular polarization signals (Stokes~$V$) which can be up to two orders of magnitude smaller than what the weak magnetic field approximation predicts. As corollaries of this result, (1) the polarization emerging from an atmosphere having weak and strong fields is biased towards the weak fields, and (2) HFS features characteristic of weak fields show up even when the magnetic flux and energy are dominated by kG fields. For the HFS feature of \\mni{5538} to disappear the filling factor of kG fields has to be larger than the filling factor of sub-kG fields. Since the Mn~{\\sc i} lines are usually weak, Stokes~$V$ depends on magnetic field inclination according to the simple consine law, a scaling that holds independently of the magnetic field strength. Atmospheres with unresolved velocities produce very asymmetric line profiles, which cannot be reproduced by simple one-component model atmospheres. Inversion techniques incorporating complex magnetic atmospheres must be implemented for a proper diagnosis. Using the HFS constants available in the literature we reproduce the observed line profiles of nine lines with varied HFS patterns. Consequently, the uncertainty of the HFS constants do not seem to limit the use of Mn~{\\sc i} lines for magnetometry. ",
        "introduction": "\\label{intro}  When the polarimetric sensitivity and the angular resolution exceed the required threshold, magnetic signals show up almost everywhere on the solar photosphere. The signals in supergranulation cell interiors are particularly weak, however, most of the solar surface is in the form of cell interiors and, therefore, these weak signals probably set the total unsigned magnetic flux and magnetic energy existing in the photosphere at any given time \\citep[e.g.,][]{unn59,ste82,yi93,san98c,san04,sch03b}. The importance of these ubiquitous quiet Sun magnetic fields depends, to a large extent, on the magnetic field strengths characterizing them. For example, the magnetic flux and the magnetic energy scale as powers of the field strength, and the connectivity between photosphere and corona is probably provided by the magnetic concentrations with the largest field strengths \\citep[e.g.,][]{dom06}. Unfortunately, measuring quiet Sun magnetic field strengths is not a trivial task. All measurements are based on the residual polarization left when a magnetic field of complex topology is observed with finite angular resolution \\citep[e.g.,][]{emo01,san03,tru04}. Measuring implies assuming many properties of the unresolved complex magnetic field and, in doing so, the measurements become model dependent and prone to bias. Depending on the technique used for measuring, the real quiet Sun exhibits weak daG fields \\citep[e.g.,][]{ste82,fau93,bom05}, intermediate hG fields \\citep[e.g.,][]{lin99,kho02,lop06}, or strong kG fields \\citep[e.g.,][]{san00,lit02,dom03b}. Such discrepancies can be naturally understood if the true quiet Sun contains a continuous distribution of field strengths going all the way from zero to two kG. Different techniques are biased differently and, therefore, they tend to pick out a particular part of the distribution. This scenario is very much consistent with realistic numerical simulations of magneto-convection \\citep[][]{cat99a,stei06,vog05,vog07}. In order to provide a comprehensive observational description of the quiet Sun magnetic field strengths, one has to combine different methods carefully chosen to have complementary biases. Then the full distribution can be assambled taking the biases into account \\cite[][]{dom06}.    In an effort to complement the existing magnetic field strength diagnostic techniques, \\citet{lop02} proposed the use of spectral lines whose Zeeman patterns are sensitive to hyperfine structure (HFS). The formalism to deal with the HFS of spectral lines in magnetic atmospheres was developed more than thirty years ago by \\citet{lan75}. According to such formalism, the polarization of the HFS lines vary with magnetic field strength  very differently from the lines commonly used in solar magnetometry. This unusual behavior was invoked by \\citet{lop02} when proposing the use of HFS Mn~{\\sc i}  lines as a  diagnostic tool for magnetic field strengths. L\\'opez Ariste and coworkers have applied the idea to measure magnetic field strengths in quiet Sun regions \\citep{lop06,lop07,ase07}. Since the number of observables is limited, they minimize the statistical error of the measuremet by  minimizing the number of free parameters to be tuned. Milne-Eddington atmospheres (ME) are used to synthesize the polarization of the lines. The magnetic field is assumed to be constant and, therefore, the measurements provide some kind of weighted average of the true field strengths existing in the resolution  elements. As we pointed out above, the topology of the quiet sun magnetic field is complex, with a distribution of field strengths and polarities coexisting in a typical resolution element. Then the ill-defined average provided by the Mn~{\\sc i} lines is expected to be biased toward a particular range of field strengths, as it happens with the rest of the quiet Sun measurements (e.g., the NIR Fe~{\\sc i} lines overlook kG magnetic concentrations, \\citealt{san00}, \\citealt{soc03}; the traditional visible Fe~{\\sc i} lines exaggerate the contribution of kG fields, \\citealt{bel03}, \\citealt{mar06}; the Hanle scattering polarization signals are not sensitive to hG and kG, \\citealt{fau01}, \\citealt{san05}). The bias presented by the HFS Mn~{\\sc i} lines is so far unknown, and the existing and forthcoming measurements based on those lines will be fully appreciated only when the sources of systematic effects are properly acknowledged and quantified. In order to explore the magnitude and the sense of the expected effects, we undertake the synthesis of  Mn~{\\sc i} lines in a number of realistic quiet Sun atmospheres with complex magnetic field distributions. The main trends are presented here.   The paper is organized as follows. Section~\\ref{code} describes the software developed to carry out the syntheses. First, a ME code is needed to compare our syntheses with the results in the literature. Then, a plane parallel one-dimensional (1D) code allows us to explore the influence of realistic thermodynamic conditions on the polarization of lines with HFS. Finally, a MIcro-Structured Magnetic Atmosphere (MISMA) code provides additional realism to the modeling since it includes  coupling between magnetic field strengths and thermodynamic conditions, magnetic fields varying along and across the line-of-sight (LOS), mixed polarities in the resolution element, etc. All these codes are based on the original FORTRAN routine written by \\citet{lan78}. The analysis is focused on the line most often used in observations, namely, \\mni{5538}. We discuss its intensity and circular polarization, since the linear polarization signals are very weak and they remain  undetected so far. Single component and multi-component syntheses of this line are described in \\S~\\ref{scs} and \\S~\\ref{mcs}, respectively. An exploratory attempt to consider unresolved mixed polarities and velocity fields is carried out in \\S~\\ref{comp_obs}. The polarization of other Mn~{\\sc i} lines is also considered in \\S~\\ref{other}. The implications of these syntheses are discussed in \\S~\\ref{conclusions}. ",
        "conclusions": "The Zeeman pattern of the Mn~{\\sc i} lines depends on hyperfine structure (HFS), which confers them a sensitivity to hG magnetic field strengths different from the  lines traditionally used in solar magnetometry. This peculiarity has been used to measure magnetic field strengths in quiet Sun regions (see \\S~\\ref{intro}). The methods applied so far assume the magnetic field strength to be constant in the resolution element, an approximation driven by feasibility rather than based on physical or observational arguments. Actually, it is not a good approximation since the magnetic fields of the quiet Sun are expected to vary on very small spatial scales, with field strengths spanning from zero to 2kG. Under these extreme conditions, all diagnostic techniques employed in quiet Sun magnetometry are strongly biased towards a particular range of field strengths, and the Mn~{\\sc i} signals are not expected to be the exception. Consequently, the  diagnostic content of the Mn~{\\sc i} lines cannot be fully exploited unless their biases are properly understood. Such task is undertaken in the paper by exploring the response of Mn~{\\sc i} lines in a number of realistic quiet Sun scenarios.  Three complementary LTE synthesis codes have been written and tested (ME, 1D and MISMA; \\S~\\ref{code}). They provide a number of relevant results, the first one being the ability to reproduce all observed Mn~{\\sc i} HFS patterns. We reproduce the observed unpolarized line profile of nine assorted lines corresponding to  all HFS sensitivities (\\S~\\ref{other}). The study is focused on the line most often used in magnetometry,  \\mni{5538}, however its behavior should be representative of the other lines. According to the weak magnetic field approximation, the Stokes $V$ signals scale with the longitudinal component of the magnetic field, i.e., the magnetic field strength times the cosine of the magnetic field inclination with respect to the LOS. We verify that the scaling on cosine holds very tightly. The scaling on the magnetic field strength, however, breaks down soon  (when $B \\ge  400$~G for \\mni{5538}). Even for ME  atmospheres, the weak field approximation predicts a  Stokes~$V$ signal twice as large as the synthetic one for $B\\simeq$1.5~kG. When the expected coupling between the thermodynamic conditions and the magnetic field strengths is taken into account, the dimming of the kG Stokes~$V$ signals can be as large as two orders of magnitude (see Figs.~\\ref{maxv} and \\ref{phi}). The dimming of the polarization signals formed in kG magnetic concentrations affects all Mn~{\\sc i} lines (\\S~\\ref{scs}),  and it has significant observational implications.  If the resolution elements contains both weak and strong fields, then the kG fields tend to be under-represented in the average profile. We have modelled the bias assuming a multi component atmosphere, where the synthetic signals are weighted means of the Stokes profiles corresponding to each single field strength. The weight is given by the fraction of atmosphere filled by each field strength, i.e., by the magnetic field strength probability density function PDF (equations~[\\ref{pdfIV}] and [\\ref{pdfIVb}]). According to our modeling, even when the (unsigned) magnetic flux and the magnetic energy are dominated by kG magnetic fields, the Stokes~$V$ profile of \\mni{5538} can show HFS reversal at line core characteristic of hG fields (Fig.~\\ref{vsme}). A pure morphological inspection of the Stokes profiles  does not suffice to infer which is the dominant magnetic field strength in the resolution element. For the HFS hump of \\mni{5538} to disappear the kG filling factor has to be larger than the sub-kG filling factor and, consequently, when the HFS hump disappears the magnetic flux and magnetic energy of the atmosphere are completely dominated by kG fields (\\S~\\ref{proxy}).     Detecting Stokes~$V$ profiles with HFS features indicates the presence of hG fields in the resolution element. However, this sole observation does not tell whether the hG field strengths dominate. There seem to be two extreme alternatives to exploit the diagnostic potential of these lines. First, improving the spatial resolution of the observations to a point where the quiet Sun magnetic structures can be regarded as spatially resolved. Unfortunately, this possibility does not seem to feasible at present. Realistic simulations of magneto-convection indicate that quiet Sun magnetic fields are uniform only at the diffusive length scales \\citep[e.g.][]{cat99a,vog07}, which are of the order of a few  km in the photosphere \\citep[e.g.][]{sch86}. These scales are very far from the angular resolution of the present measurements, and even much smaller than the length-scale for the radiative transfer average along the LOS \\citep{san96}. We prefer the alternative possibility, namely, developing inversion techniques where complex magnetic atmospheres are included into the diagnostics. Using appropriate constraints, the number of free parameters required to describe such atmospheres can be maintained within reasonable limits \\citep[e.g., the model MISMAs in ][]{san97b}. Dealing with unresolved velocities also favors detailled inversion codes (\\S~\\ref{comp_obs}). We are presently working on these improvements needed to develop the diagnostic technique pioneered by \\citeauthor{lop02}."
    },
    "0710/0710.3326_arXiv.txt": {
        "abstract": "{I provide a short review of the properties of Narrow-line Seyfert 1 (NLS1) galaxies across the electromagnetic spectrum and of the models to explain them. Their continuum and emission-line properties manifest one extreme form of Seyfert activity. As such, NLS1 galaxies may hold important clues to the key parameters that drive nuclear activity. Their high accretion rates close to the Eddington rate provide new insight into accretion physics, their low black hole masses and perhaps young ages allow us to address issues of black hole growth, their strong optical FeII emission places strong constraints on FeII and perhaps metal formation models and physical conditions in these emission-line clouds, and their enhanced radio quiteness permits a fresh look at causes of radio loudness and the radio-loud radio-quiet bimodality in AGN. }  \\resumen{Favor de proporcionar un resumen en espa\\~nol. \\\\ }  \\addkeyword{Galaxies: Active} \\addkeyword{Galaxies: Seyfert}   \\begin{document} ",
        "introduction": "\\label{sec:intro}  Narrow-line Seyfert 1 galaxies are a subclass of active galactic nuclei (AGN). Their spectra  exhibit exceptional emission-line and continuum properties. The most common NLS1 defining criterion is the width of the broad component of their optical Balmer emission lines in combination with the relative weakness of the [OIII]$\\lambda$5007 emission (FWHM$_{\\rm H\\beta} < 2000$ km/s and [OIII]/H$\\beta_{\\rm totl}$  $<$ 3; Osterbrock \\& Pogge 1985, Goodrich 1989){\\footnote{ While it is clear that a strict cutoff in line width (FWHM$_{\\rm H\\beta} < 2000$ km/s) is a gross simplification of any classification scheme, this historical value is still most commonly adopted for practical purposes. Suggestions have been made that more advanced NLS1 classification schemes would, for instance, incorporate the source luminosity (e.g., Laor 2000, Veron-Cetty et al. 2001). According to Sulentic et al. (2008 and references therein), AGN properties appear to change more significantly at a broad line width of FWHM$_{\\rm H\\beta} \\approx 4000$ km/s.}}. NLS1 galaxies typically show strong FeII emission which anticorrelates in strength with the [OIII] emission, and with the width of the broad Balmer lines. Often the presence of  FeII emission is added as further NLS1 classification criterion and Veron et al. (2001) suggest the use of an intensity ratio FeII/H$\\beta_{\\rm totl} > 0.5$.  NLS1 galaxies as AGN with the smallest Balmer lines from the Broad Line Region (BLR) and the strongest FeII emission, cluster at one extreme end of AGN correlation space. It is expected that such correlations provide some of the strongest constraints on, and new insights in, the physical conditions in the centers of AGN and the prime drivers of activity, and the study of NLS1 galaxies is therefore of particular interest. For instance, observations and interpretations hint at smaller black hole masses in NLS1 galaxies, and as such their black holes represent an important link with the elusive intermediate mass black holes, which have been little studied so far. Accreting likely at very close to the maximum allowed values, NLS1 galaxies are important test-beds of accretion models.  This paper provides a short overview of the multi-wavelength properties of NLS1 galaxies and major models to explain them. ",
        "conclusions": "-- {\\sl{It seems hard to resist the feeling that nature is telling us something important here, but we do not yet know what it is.}} ~~~~~~~~~~~~~~Lawrence et al. (1997) \\vskip0.3cm  While our knowledge has increased significantly in the last decade, important questions are still open. For instance, what are sufficient, what are necessary conditions for the onset of NLS1 activity ? For instance, the question is raised whether there are two types of Seyfert galaxies with low black hole masses: there is the unavoidable low-black-hole mass extension of BLS1 galaxies. Such systems would have their FWHM$_{\\rm H\\beta}$ fall below the formal cutoff value of 2000 km/s. Does low black hole mass already imply the emergence of some or all of the typical observed NLS1 characteristics ? Or is there a separate class of NLS1 galaxies ?  While many individual NLS1 galaxies have been studied in great detail, we still need larger samples free of selection biases and well-suited BLS1 comparison samples in order to identify robust trends. Correlation space will ultimately have to be expanded to include the radio and infrared properties of NLS1 galaxies, as well as the properties of their host galaxies. On the theoretical/modeling side, interesting questions that persist or have emerged are related to mechanisms of super-Eddington accretion, the simultaneous presence of (TeV) blazar activity and high accretion rate in extreme radio loud NLS1 galaxies, and mechanisms of fueling and feedback in NLS1 galaxies. The study of NLS1 galaxies will continue to provide important contributions to our understanding of AGN and their cosmic evolution."
    },
    "0710/0710.1609_arXiv.txt": {
        "abstract": "We examine the early phases of two near-limb filament destabilizations involved in coronal mass ejections on 16 June and 27 July 2005, using high-resolution, high-cadence observations made with the Transition Region and Coronal Explorer (TRACE), complemented by coronagraphic observations by Mauna Loa and the SOlar and Heliospheric Observatory (SOHO).  The filaments' heights above the solar limb in their rapid-acceleration phases are best characterized by a height dependence $h(t)\\propto t^m$ with $m$ near, or slightly above, 3 for both events. Such profiles are incompatible with published results for breakout, MHD-instability, and catastrophe models.  We show numerical simulations of the torus instability that approximate this height evolution in case a substantial initial velocity perturbation is applied to the developing instability. We argue that the sensitivity of magnetic instabilities to initial and boundary conditions requires higher fidelity modeling of all proposed mechanisms if observations of rise profiles are to be used to differentiate between them.  The observations show no significant delays between the motions of the filament and of overlying loops: the filaments seem to move as part of the overall coronal field until several minutes after the onset of the rapid-acceleration phase. ",
        "introduction": "Observations of the early rise phase of filaments and their overlying fields can in principle help constrain the mechanisms involved in the destabilization of the magnetic configuration through comparison with numerical simulations (e.g., \\citep{fan2005}; \\citep{torok+kliem2005}; \\citep{williams+etal2005}; and references therein), because the detailed evolution depends sensitively on the model details.  For example, a power-law rise with an exponent $m=2.5$ was obtained for a slender flux tube in the two-dimensional version of the catastrophe model (\\citep{Priest&Forbes02}). An MHD instability triggered by an infinitesimal perturbation implies an exponential rise, as was verified, for example, for a three-dimensional flux rope subject to a helical kink instability (\\citep{Torok&al04}, \\citep{torok+kliem2005}). The same holds for the torus (expansion) instability (TI), which starts as a $\\sinh(t)$ function (\\citep{Kliem&Torok06}) that is very similar to a pure exponential early on.  The CME rise in a breakout model simulation was well described by a parabolic profile (\\citep{Lynch&al04}).  The early rise phase of erupting filaments is best observed near the solar limb using high-resolution data, both in space and in time. Such data can be obtained by, for example, Big Bear Solar Observatory H$\\alpha$ observations (e.g., \\citep{kahler+etal1988}), the Mauna Loa K-coronameter (e.g., \\citep{gilbert+etal2000}), the Nobeyama Radioheliograph (e.g., \\citep{gopalswamy+etal2003}; \\citep{kundu+etal2004}), and the Transition Region and Coronal Explorer, TRACE (e.g., \\citep{vrsnak2001}; \\citep{gallagher+etal2003}; \\citep{goff+etal2005}; \\citep{sterling+moore2004}; \\citep{sterling+moore2005}; \\citep{williams+etal2005}). In those few cases where observers had the field of view for an appropriate diagnostic to attempt to establish whether the high loops or the filaments were accelerated first, the temporal resolution often was not adequate (see, e.g., \\citep{sterling+moore2004}, who use the standard 12-min.\\ cadence of SOHO/EIT).  These studies show that filaments that are about to erupt often --~but not always~-- exhibit a slow initial rise during which both the filament and the overlying field expand with velocities in the range of $1-15$\\,km/s. Then follows a rapid-acceleration phase during which velocities increase to a range of $100$\\,km/s up to over $1000$\\,km/s. The rapid-acceleration phase finally transitions into a phase with a nearly constant velocity or even a deceleration into the heliosphere.  The height evolution immediately following the onset of the rapid acceleration phase is often approximated by either an exponential curve (e.g., \\citep{gallagher+etal2003}; \\citep{goff+etal2005}; \\citep{williams+etal2005} --~who also show systematic deviations from that fit up to 2$\\sigma$ in position~-- ) or by a constant-acceleration curve (e.g., \\citep{kundu+etal2004}; and \\citep{gilbert+etal2000} --~ who show one case in which a third-order curve improves the fit to the earliest evolution, and leave others for future analysis); \\citet{kahler+etal1988} fit curves for the acceleration $a=ct^b$ to the first $10-50$\\,Mm for four erupting filaments, but do not list the best-fit values. \\cite{alexander&al02} find a best fit for the height of the early phase of a CME observed in X-rays by YOHKOH's SXT of the form $h_0+v_0t+ct^{3.7\\pm0.3}$. For 184 prominence events observed by the Nobeyama Radioheliograph, \\citet{gopalswamy+etal2003} show that higher in the corona velocity profiles include decelerating, constant velocity, and accelerating ones for heights from $\\sim 50$\\,Mm to 700\\,Mm above the solar surface.  In many cases, the detailed study of the evolution of the early phase is hampered by insufficient temporal coverage or by gaps between the fields of view of two complementing instruments that can be as large as a few hundred Mm. This results in substantial uncertainties in the height evolution. \\citet{vrsnak2001}, for example, concludes that ``[t]he main acceleration phase \\ldots\\ is most often characterized by an exponential-like increase of the velocity'', but notes that polynomial or power-law functions fit at comparable confidence levels.  In this study, we examine two events displaying the early destabilization and acceleration of ring filaments leading to coronal mass ejections. The high cadence down to 20\\,s, and the high spatial resolution of 1\\,arcsec, for the early evolution result in relatively small uncertainties in the height profiles. This enables a sensitive test of the height evolution against exponential, parabolic, and power-law fits. We find that a power-law with exponent near $3$, or slightly higher, is statistically preferred in both cases. As no published model matches that profile, we experiment with a numerical model for the torus instability, and find that this model can indeed approximate the observations provided that a sufficiently large initial velocity perturbation is applied (without which an exponential-like profile would be found). This finding reminds us of the sensitivity of developing instabilities to both initial and boundary conditions, and shows that the models, particularly their parametric dependencies, need to be worked out in greater detail in order to use observations of the height-time observations to differentiate successfully between competing models. ",
        "conclusions": "\\label{sec:conclusions} We study two well-observed filament eruptions, and find that their rapid acceleration phases are well fit by a cubic height-time curve that implies a nearly constant jerk for $10-15$\\,minutes, followed by a transition to a terminal velocity of $\\sim 750$\\,km/s and $\\sim 1250$\\,km/s, respectively. Simulations of a torus instability (TI) can reproduce such a behavior, provided that a substantial initial velocity perturbation is introduced. Without that perturbation, an exponential rise profile would be found.  We note that the initial slow rise and the onset of the subsequent rapid acceleration phase are shared between the filament and overlying loop structures: neither \\referee{leads the other to within the temporal resolution. For characteristic Alfv{\\'e}n speeds over active regions of $\\sim 1,000$\\,km/s, the propagation of a perturbation over the separation of $\\sim 75,000$\\,km would require only $\\sim 1.2$\\,min., which corresponds to only one or two exposures. Thus the observations allow for Alfv{\\'e}nic propagation of a signal between filament and overlying loops, but suggest no longer-term differential evolution.}  We observe no significant changes in the separation of erupting filament and overlying loops within that interval (Figs.~\\ref{fig:1c} and~\\ref{fig:2c}). After that, the distance increases in the 2005/06/16 eruption, suggesting the overlying field moves to the side for some time faster than the filament rises.  For the 2007/07/27 eruption, the distance stays the same for one loop and decreases for another for up to 10~min after the start of the rapid acceleration phase, which reflects the significant sideways motion component of the rising filament. The observed configuration of the filament and high loops may be part of a larger overall destabilizing field configuration. Our numerical modeling has assumed that, in the rapid acceleration phase, the overlying field starts to move rapidly only as a consequence of the flux rope's destabilization. This is consistent with the data. However, we cannot exclude that the filament and the overlying field were destabilized simultaneously by a process different from the one considered here. More study is needed to establish whether the common evolution of the filament and high loops has a significant diagnostic value as to the cause of the instability.  Comparison with other model studies in the literature leads us to conclude that the catastrophe model and the TI model are both marginally consistent with the observations of the two erupting filaments. The catastrophe model predicts a power-law exponent near the lower edge of the range of acceptable fits, but we have to allow for the possibility that changing that model's details may change the acceleration profile. \\referee{In order to yield the observed nearly cubic power-law rise (with $m$ slightly exceeding 3), our TI model requires an initial perturbation velocity that is in agreement with the observed rise velocity at the onset of the rapid-acceleration phase.  If a nearly exact cubic rise were to be matched, however, initial velocities moderately exceeding the observed ones, by a factor $\\approx1.5$, were required. In any case, our modeling is consistent with the observed velocities after the first few minutes of the eruption.}  Having established that the model for the TI instability is very sensitive to the initial conditions, we should of course also acknowledge that it depends sensitively on the model details itself. These include the details of the external field and of the rates and locations of the reconnection that occurs behind the erupting filament.  That such reconnection occurs in reality is suggested for both events by the occurrence of brightenings mainly at the bottom side of the filaments at the onset of the rapid-acceleration phase. These brightenings develop later into the streaks used for position determination in Sect.~\\ref{sec:Observations}.  The onset of reconnection even before the rapid-acceleration phase of the filament eruption on 27 July 2005 is strongly suggested by precursor soft and hard X-ray emission during about 04:00--04:30~UT, whose analysis revealed heating to 15~MK and the acceleration of non-thermal electrons to energies $>10$~keV (\\citep{Chifor&al06}).  The observed rise velocity early in the filament eruption may be an underestimate of the true expansion velocity of the hoop formed by the flux rope: the filament channel in the pre-eruption phase of AR\\,10775 is strongly curved, and one of the two possible channels in AR\\,10792 is too (ambiguity exists here because the eruptions occurred very near the limb, so that the configurations of the filament channels can only be observed some days before and after the events, respectively). If the initial expansion of the flux rope would have a strong component in the general direction of the inclined plane of the curved filament channel rather than be purely normal to the solar surface, projection effects could cause us to underestimate the expansion velocity in particular early in the evolution.  In addition to that, we must realize that the TI model assumes a flux rope that stands normal to the solar surface and that erupts radially. Future more detailed modeling will have to show how deviations from that affect the evolution of the eruption.  The fact that the torus-instability model yields qualitatively different rise profiles (exponential vs.\\ power law) in different parts of parameter space, cautions against expectations that precise measurements of the rise profile of filament eruptions by themselves permit a determination of the driving process: the non-linearities in the eruption models clearly require high-fidelity modeling if such observations are to be used to differentiate successfully between competing models. Our initial modeling discussed here suggests that the torus instability is a viable candidate mechanism for at least some filament eruptions in coronal mass ejections. \\referee{Given the dependence of nonlinear models on the details of boundary and initial conditions, it will be necessary to investigate how other models for erupting filaments compare to the data, as well as how the fidelity of our modeling of the torus instability can be improved before we can reach definitive conclusions about the mechanism(s) responsible for filament eruptions in general.}"
    },
    "0710/0710.4506_arXiv.txt": {
        "abstract": "The hydrogen-deficiency in extremely hot post-AGB stars of spectral class PG1159 is probably caused by a (very) late helium-shell flash or a AGB final thermal pulse that consumes the hydrogen envelope, exposing the usually-hidden intershell region. Thus, the photospheric elemental abundances of these stars allow to draw conclusions about details of nuclear burning and mixing processes in the precursor AGB stars. We compare predicted elemental abundances to those determined by quantitative spectral analyses performed with advanced non-LTE model atmospheres. A good qualitative and quantitative agreement is found for many species (He, C, N, O, Ne, F, Si, Ar) but discrepancies for others (P, S, Fe) point at shortcomings in stellar evolution models for AGB stars. PG1159 stars appear to be the direct progeny of [WC] stars. ",
        "introduction": "The PG1159 stars are a group of 40 extremely hot hydrogen-deficient post-AGB stars.  Their effective temperatures ($T_{\\rm eff}$) range between 75\\,000 and 200\\,000~K. Many of them are still heating up along the constant-luminosity part of their post-AGB evolutionary path in the HR diagram ($L \\approx 10^4$L$_\\odot$) but most of them are already fading along the hot end of the white dwarf cooling sequence (with $L$ {\\raisebox{-0.4ex}{${\\stackrel{>}{\\scriptstyle \\sim}} $}} 10\\,L$_\\odot$). Luminosities and masses are inferred from spectroscopically determined $T_{\\rm eff}$ and surface gravity ($\\log g$) by comparison with theoretical evolutionary tracks. The position of analysed PG1159 stars in the ``observational HR diagram'', i.e., the $T_{\\rm eff}$--$g$ diagram, are displayed in Fig.\\,\\ref{fighrd}. The high-luminosity stars have low $\\log g$ ($\\approx$\\,5.5) while the low-luminosity stars have a high surface gravity ($\\approx$\\,7.5) that is typical for white dwarf (WD) stars. The derived mean mass is 0.57\\,M$_\\odot$, a value that is practically identical to the mean mass of WDs \\citep{be:07}. The PG1159 stars co-exist with hot central stars of planetary nebulae and the hottest hydrogen-rich (DA) white dwarfs in the same region of the HR diagram. About every other PG1159 star is surrounded by an old, extended planetary nebula. For a recent review with a detailed bibliography see \\citet{werner:06}.  What is the characteristic feature that discerns PG1159 stars from ``usual'' hot central stars and hot WDs? Spectroscopically, it is the lack of hydrogen Balmer lines, pointing at a H-deficient surface chemistry. The proof of H-deficiency, however, is not easy: The stars are very hot, H is strongly ionized and the lack of Balmer lines could simply be an ionisation effect. In addition, every Balmer line is blended by a Pickering line of ionized helium. Hence, only detailed modeling of the spectra can give reliable results on the photospheric composition. The high effective temperatures require non-LTE modeling of the atmospheres. Such models for H-deficient compositions have only become available in the early 1990s after new numerical techniques have been developed and computers became capable enough.  \\begin{figure*}[bth] \\vspace{-3.1cm} \\begin{center} \\epsfxsize=1.0\\textwidth \\epsffile{kwerner02.eps} \\end{center} \\vspace{-6.9cm} \\caption{\\label{fig:hrd} Complete stellar evolution track with an initial mass of 2\\,M$_\\odot$ from the main sequence through the RGB phase, the HB to the AGB phase, and finally through the post-AGB phase that includes the central stars of planetary nebulae to the final WD stage. The solid line represents the evolution of a H-normal post-AGB star. The dashed line shows a born-again evolution of the same mass, triggered by a very late thermal pulse, however, shifted by approximately $\\Delta \\log T_{\\rm eff} = -0.2$  and $\\Delta \\log~L/$L$_\\odot = - 0.5$ for clarity. The double-loop structure of the path is the consequence of a hydrogen-ingestion flash. The ``$\\star$'' symbol shows the position of PG1159$-$035 \\citep[from][]{werner:06}. } \\end{figure*}  The first quantitative spectral analyses of optical spectra from PG1159 stars indeed confirmed their H-deficient nature \\citep{werner:91}. It could be shown that the main atmospheric constituents are C, He, and O. The typical abundance pattern is C=0.50, He=0.35, O=0.15 (mass fractions). It was speculated that these stars exhibit intershell matter on their surface, however, the C and O abundances were much higher than predicted from stellar evolution models. It was further speculated that the H-deficiency is caused by a late He-shell flash, suffered by the star during post-AGB evolution, laying bare the intershell layers. The re-ignition of He-shell burning brings the star back onto the AGB, giving rise to the designation ``born-again'' AGB star \\citep{iben:83a}. If this scenario is true, then the intershell abundances in the models have to be brought into agreement with observations. By introducing a more effective overshoot prescription for the He-shell flash convection during thermal pulses on the AGB, dredge-up of carbon and oxygen into the intershell can achieve this agreement \\citep{herwig:99c}. Another strong support for the born-again scenario was the detection of neon lines in optical spectra of some PG1159 stars \\citep{werner:94}. The abundance analysis revealed Ne=0.02, which is in good agreement with the Ne intershell abundance in the improved stellar models.  If we do accept the hypothesis that PG1159 stars display former intershell matter on their surface, then we can in turn use these stars as a tool to investigate intershell abundances of other elements. Therefore these stars offer the unique possibility to directly see the outcome of nuclear reactions and mixing processes in the intershell of AGB stars. Usually the intershell is kept hidden below a thick H-rich stellar mantle and the only chance to obtain information about intershell processes is the occurrence of the third dredge-up. This indirect view onto intershell abundances makes the interpretation of the nuclear and mixing processes very difficult, because the abundances of the dredged-up elements may have been changed by additional burning and mixing processes in the H-envelope (e.g., hot-bottom burning). In addition, stars with an initial mass below 1.5~M$_\\odot$ do not experience a third dredge-up at all.  The central stars of planetary nebulae of spectral type [WC] are believed to be immediate progenitors of PG1159 stars, representing the evolutionary phase between the early post-AGB and PG1159 stages. This is based on spectral analyses of [WC] stars which yield very similar abundance results (see papers by Crowther, Todt, and Gr\\\"afener in these proceedings).  \\begin{figure*} \\begin{center} \\epsfxsize=1.0\\textwidth \\epsffile{kwerner03.eps} \\end{center} \\vspace{-0.5cm} \\caption{\\label{fig:psifn} Detail from FUSE spectra of two relatively cool PG1159 stars (see labels). Note the following features. The \\ion{F}{vi}~1139.5~\\AA\\ line which is the first detection of fluorine at all in a hot post-AGB star; the \\ion{P}{v} resonance doublet at 1118.0 and 1128.0~\\AA, the first discovery of phosphorus in PG1159 stars; the \\ion{N}{iv} multiplet at 1132~\\AA. Also detected are lines from \\ion{Si}{iv} and \\ion{S}{vi}. The broader features stem from \\ion{C}{iv} and \\ion{O}{vi} \\citep{reiff:07}.} \\end{figure*} ",
        "conclusions": "It has been realized that PG1159 stars exhibit intershell matter on their surface, which has probably been laid bare by a late final thermal pulse. This provides the unique opportunity to study directly the result of nucleosynthesis and mixing processes in AGB stars. Abundance determinations in PG1159 stars are in agreement with intershell abundances predicted by AGB star models for many elements (He, C, N, O, Ne, F, Si, Ar). For other elements, however, disagreement is found (Fe, P, S) that points at possible weaknesses in the evolutionary models. Generally, the abundance patterns clearly support the idea that [WC] stars are direct progenitors of PG1159 stars."
    },
    "0710/0710.1786_arXiv.txt": {
        "abstract": "{A large fraction of the stellar mass is found to be located in groups of the size of the Local Group. Evolutionary status of poor groups is not yet clear and many groups could still be at an early dynamical stage or even still forming, especially the groups containing spiral and irregular galaxies only.} {We carry out a photometric study of a poor group of late-type galaxies around IC 65, with the aim: (a) to search for new dwarf members and to measure their photometric characteristics; (b) to search for possible effects of mutual interactions on the morphology and star-formation characteristics of luminous and faint group members; (c) to evaluate the evolutionary status of this particular group.} {We make use of our $BRI$ CCD observations, DPOSS blue and red frames, and the 2MASS $JHK$ frames. In addition, we use the \\ion{H}{i} imaging data, the far-infrared and radio data from the literature. Search for dwarf galaxies is made using the SExtractor software. Detailed surface photometry is performed with the MIDAS package. } {Four LSB galaxies were classified as probable dwarf members of the group and the $BRI$ physical and model parameters were derived for the first time for all true and probable group members. Newly found dIrr galaxies around the IC 65 contain a number of \\ion{H}{ii} regions, which show a range of ages and propagating star-formation. Mildly disturbed gaseous and/or stellar morphology is found in several group members. } {Various structural, dynamical, and star-forming characteristics let us conclude that the IC 65 group is a typical poor assembly of late-type galaxies at an early stage of its dynamical evolution with some evidence of intragroup (tidal) interactions.} ",
        "introduction": " ",
        "conclusions": "The main contributions of this paper are as follows: \\begin{enumerate} \\item We have selected four LSB dwarf companion candidates of the IC 65 group of galaxies on deep DPOSS frames according to their surface brightnesses, colours and morphology. \\item The $B, R$ and $I$ band surface photometry is presented for the first time for all certain and probable members of the group. The bright group members were studied in the NIR $J, H$ and $K$ bands, too. An image gallery and the deduced $SB$ and colour profiles are shown, permitting the detailed morphological analysis of the galaxies studied. Their relevant physical and model characteristics are determined. \\item Dynamical masses and star-forming characteristics of the bright group members are estimated using the new optical photometry and the available NIR, FIR, \\ion{H}{i} and radio data. The probable evolutionary status of the group is discussed. \\end{enumerate}  An analysis of the available photometric and kinematic data of individual galaxies with emphasis to study the possible mutual interactions between the group members leads to the following results:  \\noindent $\\bullet$ The available \\ion{H}{i} imaging data show that all bright members and at least one dwarf companion have a nearly normal gaseous fraction with \\ion{H}{i} mass to blue luminosity ratios in the range of 0.3 -- 1.0, consistent with their morphological type. The outer \\ion{H}{i} isophotes of the IC 65, and especially of the UGC 622 appear disturbed, in agreement with perturbation estimates.\\\\ $\\bullet$ The optical morphology of the bright galaxies generally appears to be regular, with barely significant disturbances in isophotes of the outer stellar disk of the IC 65 and UGC 622.\\\\ $\\bullet$ All bright group members (except PGC~138291, which we could not study in such detail) consist of many blue star-forming knots and plumes, especially UGC~608. A comparison of the surface photometry with stellar population models of Bruzual \\& Charlot (\\cite{bruzual03}) indicates that these blue knots must have formed recently. % The available data do not allow to establish whether they formed simultaneously, e.g. in star-bursts possibly triggered by interactions.\\\\ $\\bullet$ Two dIrr galaxies around the IC 65 both contain a number of \\ion{H}{ii} regions, which show a range of stellar ages and provide an evidence of propagating star-formation. One of these galaxies - A~0101+4744 is a confirmed member of the group; the second one - A~0100+4734 appears to be located in front of the group.\\\\ $\\bullet$ The brightest galaxies in both subgroups can fuel their current star-forming rates of $\\sim$ 1 - 2 ${\\cal M}_\\odot$ yr$^{-1}$ for about the next 3 - 7 Gyr.  The IC~65 group of galaxies obviously belongs to the class of less evolved groups. It is composed of late type spiral and irregular galaxies arranged in two subgroups. No massive early type galaxies are present. No hot gas has been detected in it by the ROSAT survey. % Some morphological irregularities and signs of enhanced SF in its members could be indicative of recent/ongoing  mutual interactions. Yet, the individual group members have retained much of their initial gas component. A few available velocities point to a short crossing time of only $\\sim$ 0.1~$H_0^{-1}$. However, this hardly means that the group has already reached a stable (virialized) configuration. The evidence, discussed above, lets us conclude that the IC~65 group of galaxies is a dynamically young system at a still relatively early stage of its collapse."
    },
    "0710/0710.3783_arXiv.txt": {
        "abstract": "We report the first results of the GammeV experiment, a search for milli-eV mass particles with axion-like couplings to two photons.  The search is performed using a ``light shining through a wall'' technique where incident photons oscillate into new weakly interacting particles that are able to pass through the wall and subsequently regenerate back into detectable photons.  The oscillation baseline of the apparatus is variable, thus allowing probes of different values of particle mass.  We find no excess of events above background and are able to constrain the two-photon couplings of possible new scalar (pseudoscalar) particles to be less than $3.1\\times 10^{-7} \\mbox{ GeV}^{-1}$ ($3.5\\times 10^{-7} \\mbox{ GeV}^{-1}$) in the limit of massless particles. ",
        "introduction": "The key to this experiment is the short 5 ns, 160 mJ pulses of 532 nm light emitted with a repetition rate of 20 Hz by our light source, a frequency-doubled Continuum Surelite I-20 Nd:YAG laser.  As described below, the small duty cycle enables a large reduction in the detector noise via coincidence counting.  The laser light is vertically polarized and when needed, a halfwave plate is used to obtain horizontal polarization.  The laser pulses are sent through a vacuum system (diagrammed in Fig.~\\ref{F:apparatus}) designed around an insulating warm bore inserted into a 6 m Tevatron superconducting dipole magnet.  The magnet produces a 5 T vertical field uniform across the aperture of the $48 \\mbox{ mm}$  inner diameter warm bore.  A ``wall'' consisting of a high-power laser mirror on the end of a long (7 m) hollow stainless steel ``plunger'' is inserted into the warm bore.   The mirror may be placed at various positions within the magnet by sliding the plunger.  The plunger mirror projects the reflected wave into a photon state and the transmitted wave into a (pseudo-)scalar state, provided that the scalar is sufficiently weakly-interacting to pass through the material of the mirror.  The mirror also functions to reflect the incident laser power out of the magnet to prevent heating of the magnet coils.  The mirror is mounted on a welded stainless steel cap on the end of the plunger in order to prevent stray photons from passing through.  Thus, the beam passing through the end of the plunger is a pure scalar beam.  These scalars can then oscillate back into photons through the remaining magnetic field region inside the $35 \\mbox{ mm}$ inner diameter hollow plunger.  Upon exiting the magnetic field region, the interaction ceases and the photon-scalar wavefunction is frozen.  This combined wavefunction then propagates $\\sim 6\\mbox{ m}$ into a dark box where a Hamamatsu H7422P-40 photomultiplier tube (PMT) module is used to detect single photons in coincidence with the laser pulses.  As described below, a high signal to noise ratio is achieved due to the very short pulses emitted by the laser.  The photon-scalar transition probability may be written in convenient units as: \\begin{eqnarray} \\label{E:conversionprob1} P_{\\gamma \\rightarrow \\phi} &=& \\frac{4 B^2\\omega^2}{M^2 (\\Delta m^2)^2} \\sin^2 \\left( \\frac{\\Delta m^2 L}{4\\omega} \\right) \\\\ \\label{E:conversionprob2} &\\approx&  \\frac{4 B^2\\omega^2}{M^2 m_\\phi^4} \\sin^2 \\left( \\frac{m_\\phi^2 L}{4\\omega} \\right) \\\\ \\label{E:conversionprob3} &=& 1.5\\times 10^{-11} \\frac{(B/\\mbox{Tesla})^2(\\omega/\\mbox{eV})^2}{(M/10^5 \\mbox{ GeV})^2 (m_\\phi/10^{-3} \\mbox{ eV})^4} \\nonumber  \\\\ & & \\times \\sin^2 \\left( 1.267 \\frac{(m_\\phi/10^{-3} \\mbox{ eV})^2 (L/\\mbox{m})}{(\\omega/\\mbox{eV})} \\right) \\end{eqnarray} where $B$ is the strength of the external magnetic field, $\\omega$ is the inital photon energy, $L$ is the magnetic oscillation baseline.   The mass-squared difference between the scalar mass and the effective photon mass, $\\Delta m^2=m_\\phi^2-m_\\gamma^2$, characterizes the mismatch of the phase velocities of the photon wave and the massive scalar wave and determines the characteristic oscillation length.  While the photon does not really gain a mass in a normal dielectric medium, the phase advance may be modelled with an effective imaginary mass $m_\\gamma^2 = -2 \\omega^2 (n-1)$, where $n$ is the index of refraction \\cite{vanBibber:1988ge}.  Both the warm bore and the interior of the plunger are pumped to moderate vacuum pressures of less than $10^{-4} \\mbox{ Torr}$, and a conservative estimate gives $\\sqrt{-m_\\gamma^2} <  10^{-4} \\mbox{ eV}$.  Therefore, starting with Eqn.~\\ref{E:conversionprob2} we assume that the contribution from the effective photon mass is negligible for larger values of $m_\\phi$ near the PVLAS region.  As can be seen from Eqn.~\\ref{E:conversionprob3}, the meter scale baseline provided by typical accelerator magnets is well-suited for probing the milli-eV range of possible particle masses.  This fact can be a curse as well as a boon because for a monochromatic laser beam, a fixed magnet length may accidentally coincide with a minimum in the oscillation rather than a maximum.  Indeed, this is a possible reason why the BFRT experiment \\cite{Ruoso:1992nx} did not see the PVLAS signal even though they had similar sensitivity.  GammeV's plunger design allows us to change the oscillation baseline and thus scan through all possible values of the scalar mass in the milli-eV range without any regions of diminished sensitivity.  The total conversion and regeneration probability contains two factors of Eqn.~\\ref{E:conversionprob2}, corresponding to the pre-mirror and post-mirror magnetic field regions of lengths $L_1$ and $L_2$.  The total probability varies as $\\sin^2 \\left(\\frac {m_\\phi^2 L_1}{4\\omega} \\right) \\sin^2 \\left(\\frac{m_\\phi^2 L_2}{4\\omega} \\right)$ where $L_1+L_2=6\\mbox{ m}$.   \\begin{figure}[t] \\includegraphics[width=0.48\\textwidth]{gammevcartoon.eps} \\caption{\\label{F:apparatus} Diagram (not to scale) of the experimental apparatus.  The initial vacuum chamber consists of a 10 m insulating warm bore which is offset by 1.6 m from the end of the 6 m magnetic field region, and is sealed to the sliding plunger via a double o-ring assembly.  The sliding plunger has a range of motion of 1.9 m, and contains an independent vacuum chamber.  The vacuum window at the far end slides within a stationary, long dark box.} \\end{figure}   To detect regenerated photons we use a $51\\mbox{ mm}$ diameter lens to focus the beam onto the $5 \\mbox{ mm}$ diameter GaAsP photocathode of the PMT.  The alignment is performed using a low power green helium-neon alignment laser and a mock target.  The alignment is verified both before and after each data-taking period by replacing the sealed plunger with an open-ended plunger, re-establishing the vacuum, and firing the Nd:YAG laser onto a flash paper target.  An optical transport efficiency of $92\\%$ is measured using the ratio of laser power transmitted through the open-ended plunger and through the various optics and vacuum windows, to the initial laser power, using the same power meter in both cases to remove systematic effects.  The quantum efficiency of the photocathode is factory-measured to be $38.7\\%$ while the collection efficiency of the metal package PMT is estimated to be $70\\%$.  The PMT pulses are amplified by 46 dB and then sent into a NIM discriminator.  Using a highly attenuated LED flasher as a single photon source, the discriminator threshold is optimized to give 99.4$\\%$ efficiency for triggering on single photo-electron pulses while also efficiently rejecting the lower amplitude noise.   By studying the trigger time distribution, the deadtime fraction due to possible multiple rapid PMT pulses is found to be negligible (0.001$\\%$).  Thus, we estimate the total photon transport and counting efficiency to be $(25\\pm 3)\\%$.  Using this threshold, and the built-in cooler to cool the photocathode to $0^\\circ$C, we measure a typical dark count rate of 130 Hz.  \\begin{figure}[t] \\includegraphics[width=0.48\\textwidth]{ttime_leaky_100ns.eps} \\caption{\\label{F:timing} PMT trigger times for the four run configurations, shown relative to the expected time distribution of photons as calibrated from the leaky mirror data.} \\end{figure}  To perform the coincidence counting we use two Quarknet boards \\cite{quarknet} \\cite{quarknetgps} with 1.25 ns timing precision, referenced to a GPS clock.  The Quarknet boards determine the absolute time of the leading edge of time-over-threshold triggers from the PMT and from a monitoring photodiode that is located inside the laser box.  The clocks on the laser board and on the PMT board are synchronized using an external trigger from a signal generator.  The absolute timing between the laser pulses and the PMT traces is established by removing the plunger with the mirror, and allowing the laser to shine on the PMT through several attenuation stages consisting of two partially reflective (``leaky'') mirrors, a pinhole, and multiple absorptive filters mounted directly on the aperture of the PMT module.  The $10^{19}$ photons per second emitted by the laser are thus attenuated to a corresponding PMT trigger rate of less than 0.1 Hz for this timing calibration and to provide an {\\it in situ} test of the data acquisition system.  The regenerated photons should arrive at the same time as the straight-through photons since milli-eV particles are also highly relativistic.  For coincidence counting between the laser pulses and the PMT, a 10 ns wide window is chosen and includes $99\\%$ of the measured photon time distribution shown in Fig.~\\ref{F:timing}.  The coincident dark count rate can be estimated to be $R_{\\rm{noise}}=20\\mbox{ Hz}\\times 130\\mbox{ Hz}\\times 10\\mbox{ ns} = 2.6\\times 10^{-5} \\mbox{ Hz}$.  This noise rate is negligible to the expected signal rate of $\\sim 2\\times 10^{-3} \\mbox{ Hz}$ estimated from the central values of the PVLAS parameters. \\begin{table} \\begin{tabular}{|l|c|c|c|c|} \\hline Configuration&$\\#$ photons&Est.Bkgd&Candidates&g[GeV$^{-1}$]\\\\ \\hline Horiz.,center&$6.3\\times 10^{23}$&$1.6$&1&$3.4\\times 10^{-7}$\\\\ Horiz.,edge&$6.4\\times 10^{23}$&$1.7$&0&$4.0\\times 10^{-7}$\\\\ Vert.,center&$6.6\\times 10^{23}$&$1.6$&1&$3.3\\times 10^{-7}$\\\\ Vert.,edge&$7.1\\times 10^{23}$&$1.5$&2&$4.8\\times 10^{-7}$\\\\ \\hline \\end{tabular} \\caption{\\label{F:table} Summary of data in each of the 4 configurations.} \\end{table} ",
        "conclusions": ""
    },
    "0710/0710.4450.txt": {
        "abstract": "In several studies of Low Luminosity Active Galactic Nuclei (LLAGNs), we have characterized the properties of the stellar populations in LINERs and LINER/HII Transition Objects (TOs).  We have found a numerous class of galactic nuclei which stand out because of their conspicuous 0.1--1 Gyr populations. These nuclei were called ''Young-TOs'' since they all have TO-like emission line ratios. To advance our knowledge of the nature of the central source in LLAGNs and its relation with stellar clusters, we are carrying out several imaging projects with the Hubble Space Telescope (HST) at near-UV, optical and near-IR wavelengths.  In this paper, we present the first results obtained with observations of the central regions of 57 LLAGNs imaged with the WFPC2 through any of the V (F555W, F547M, F614W) and I (F791W, F814W) filters that are available in the HST archive. The sample contains 34$\\%$ of the LINERs and 36$\\%$ of the TOs in the Palomar sample.  The mean spatial resolution of these images is 10 pc. With these data we have built an atlas that includes structural maps for all the galaxies, useful to identify compact nuclear sources and, additionally, to characterize the circumnuclear environment of LLAGNs, determining the frequency of dust and its morphology.  The main results obtained are: 1) We have not found any correlation between the presence of nuclear compact sources and emission-line type.  Thus, nucleated LINERs are as frequent as nucleated TOs.  2) The nuclei of \"Young-TOs\" are brighter than the nuclei of \"Old-TOs\" and LINERs. These results confirm our previous results that Young-TOs are separated from other LLAGNs classes in terms of their central stellar population properties and brightness. 3) Circumnuclear dust is detected in 88$\\%$ of the LLAGNs, being almost ubiquitous in TOs.  4) The dust morphology is complex and varied, from nuclear spiral lanes to chaotic filaments and nuclear disk-like structures. Chaotic filaments are as frequent as dust spirals; but nuclear disks are mainly seen in LINERs.  These results suggest an evolutionary sequence of the dust in LLAGNs, LINERs being the more evolved systems and Young-TOs the youngest. ",
        "introduction": "\\label{sec:Introduction}  %{\\bf\\ojo\\ojon ROSA: YOUR FILE CAME WITH FUNNY LINE BREAKS, SO CHECK %THAT I HAVE NOT UNDONE SOME OF THE PARAGRAPH BREAKS INADVERTEDLY...}  Low-luminosity active galactic nuclei (LLAGNs) comprise 30\\%\\ of all bright galaxies (B$\\leq$12.5) and are the most common type of AGN (Ho, Filippenko \\& Sargent 1997a, hereafter HFS97). These include LINERs, and transition-type objects (TOs, also called weak-[OI] LINERs). These two types of LLAGNs have similar emission line ratios in [OIII]/H$\\beta$, [NII]/H$\\alpha$, and [SII]/H$\\alpha$, but [OI]/H$\\alpha$ is lower in TOs than in LINERs. LLAGNs constitute a rather mixed class and different mechanisms have been proposed to explain the origin of the nuclear activity, including shocks, and photoionization by a non-stellar source, by hot stars or by intermediate age stars (e.g.  Ferland \\& Netzer 1983; Filippenko \\& Terlevich 1992; Binnette et al. 1994; Taniguchi, Shioya \\& Murayama 2000).  Because we do not know yet what powers them and how they are related to the Seyfert phenomenon, LLAGNs have been at the forefront of AGN research since they were first  systematically studied by Heckman (1980).  Are they all truly ``dwarf'' Seyfert nuclei powered by accretion onto nearly dormant supermassive black holes (BH), or can some of them be explained at least partly in terms of stellar processes? If LLAGNs were powered by a BH, they would represent the low end of the AGN luminosity function in the local universe and would also establish a lower limit to the fraction of galaxies containing massive BHs in their centers. If, on the contrary, LLAGNs were powered by nuclear stellar clusters, their presence would play an important role in the evolution of galaxy nuclei. Therefore, it is fundamental to unveil the nature of the central source in LLAGNs.  There is clear evidence that at least some LINERs harbor a bona fide AGN, and they may be considered the faint end of the luminosity function of Seyfert galaxies. It has been found that about 20\\%\\ of the nearby LINERs have a weak broad H$\\alpha$ emission component similar to those found in type 1 Seyferts (Ho, Filippenko \\& Sargent 1997b). In a few of these LINERs the H$\\alpha$ line shows a double peak component (e.g.  Storchi-Bergmann et al. 1997; Shields et al. 2000; Ho et al. 2000) suggesting that they are powered by an accreting black hole (BH). It has also been found that X-ray emission in LINERs has a nonthermal origin associated with an AGN (e.g. Terashima, Ho \\& Ptak 2000).  Recent works based on higher spatial resolution data taken with Chandra show that only in half of the LLAGNs observed the X-ray emission is associated with compact nuclear cores (Satyapal et al. 2004; Dudik et al. 2006; Gonz\\'alez-Mart\\'\\i n et al. 2006). This is consistent with VLA radio observations, which detect unresolved radio cores in half of the LINERs (Nagar et al. 2000, 2002). Finally, a monitoring study at near-UV wavelengths by Maoz et al. (2005) finds UV variability in a significant fraction of the 17 LLAGNs observed. The variation of the UV fluxes may be interpreted as the manifestation of low rate or low radiative efficiency accretion onto a supermassive BH.  On the other hand, detection of stellar wind absorption lines in the ultraviolet spectra of some TOs (Maoz et al. 1998; Colina et al. 2002) has proven unequivocally the presence of young stellar clusters in the nuclear region. Additional evidence comes from optical studies, in which we have focused on the study of the stellar population in the nuclei and circumnuclear region of LLAGNs to establish the role of stellar processes in their phenomenology. Ground-based (Cid Fernandes et al.\\ 2004; Cid Fernandes et al.\\ 2005) and HST+STIS spectra (Gonz\\'alez Delgado et al.\\ 2004) have shown that the contribution of an intermediate age stellar population is significant in a sizable fraction of the TO population.  These studies identified a class of objects, called ``Young-TOs'', which are clearly separated from LINERs in terms of the properties and spatial distribution of the stellar populations. They have stronger stellar population gradients, a luminous intermediate age stellar population concentrated toward the nucleus ($\\sim$100~pc) and much larger amounts of extinction than LINERs. These objects, which underwent a powerful star formation event $\\sim$ 1 Gyr ago, could correspond to post-Starburst nuclei or to evolved counterparts of the Seyfert 2 with a composite nucleus, characterized too by harboring nuclear starbursts (Gonz\\'alez Delgado et al. 2001; Cid Fernandes et al. 2001, 2004). HST imaging of the nuclei of these Seyfert 2 galaxies shows that the UV emission is resolved into stellar clusters (Gonz\\'alez Delgado et al. 1998) that are similar to those detected in starburst galaxies (Meurer et al.\\ 1995).  Nuclear stellar clusters are a common phenomenon in spirals, having been detected in 50-70\\% of these sources (Carollo et al.  1998, 2002; Boeker et al. 2002, 2004). Therefore stellar clusters are a natural consequence of the star formation processes in the central region of spirals. On the other hand, evidence has been accumulating during the past few years about the ubiquity of BH in the nuclei of galaxies. Furthermore, the tight correlation of the BH mass and stellar velocity dispersion (Ferrarese \\& Merrit 2000; Gebhardt et al.\\ 2000) implies that the creation and evolution of a BH is intimately connected to that of the galaxy bulge. Recently, in a HST survey in the Virgo Cluster, C\\^ot\\'e et al. (2006) have detected compact sources in a comparable fraction of elliptical galaxies. These compact stellar clusters, referred to as nuclei by the authors (see also Ferrarese et al 2006a), have masses that scale directly to the galaxy mass, in the same way as do the BH masses in high luminosity galaxies (Ferrarese et al. 2006b). Therefore, a natural consequence of the physical processes that formed present-day galaxies should be the creation of a compact massive object in the nucleus, either a BH and/or a massive stellar cluster.  To determine the nature of the nuclear source of active galaxies and their evolution we are carrying out several projects with HST+ACS imaging at the near-UV wavelengths a sample of Seyferts (ID. 9379, PI. Schmitt, Mu\\~noz-Mar\\'\\i n et al. 2007) and LLAGNs (ID. 10548, PI.  Gonz\\'alez Delgado). These observations are complemented with WFPC2 optical data retrieved from the HST archive. The high angular resolution provided by HST is crucial to determine the physical properties of the nuclei, the central structure of these galaxies, as well as to study the circumnuclear environment of AGNs. The main goals of these studies are to determine the frequency of nuclear and circumnuclear stellar clusters in AGNs, and whether they are more common in Seyferts, TOs or LINERs; to characterize the intrinsic properties of these clusters and to study whether there is evolution from Seyferts to TOs and LINERs. In addition, the frequency of dust and its morphology can also provide relevant information about the origin of nuclear activity. Dust is a valuable probe of the presence of cold interstellar gas in galaxies, and it is very sensitive to the perturbations that drive the gas toward the center and feed the AGN. Here, we present the initial results obtained for LLAGNs based on archival visible and red images obtained with the WFPC2.  The paper is organized as follows: section 2 presents the sample selection, and 3 the characteristics of the observations. Sections 4, 5 and 6 describe the imaging atlas, the dust morphology and the central properties of the galaxies. Finally, the summary and conclusions are presented in section 7.    %%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%% ",
        "conclusions": "LLAGNs, that include LINERs and TOs, are the most common type of AGN. What powers them is still at the forefront of AGN research.  To unveil the nature of the central source we are constructing a panchromatic atlas of the inner regions of these galaxies, which will be used to determine their nuclear stellar population.  To this end we have already carried out a near-UV snapshot survey of nearby LLAGNs with the ACS on board HST, that is complemented with optical and near-IR images available in the HST archive.  In this paper, the first of a series, we present observations of 57 LLAGNs imaged with the WFPC2 through any of the V (F555W, F547M, F606W) and/or I (F791W, F814W) bands. These objects comprise 36$\\%$ of the original HFS97 LLAGNs sample, and correspond to those for which there are WFPC2 images available in the HST archive and whose circumnuclear stellar population we have already studied spectroscopically (Papers I--III). The subset of objects studied here follows the same distance and morphological type distribution of the complete HFS97 LLAGN sample. Classifying the objects in strong-[OI] and weak-[OI] ([OI]/H$\\alpha$)$\\leq$ 0.25), this subset includes 34$\\%$ and 36$\\%$ of the strong- and weak-[OI], respectively, of the whole HFS97 LLAGN sample. Following our results obtained from the analysis of the circumnuclear stellar population, this sub-sample contains 17 Young-TOs, 20 Old-TOs, 18 Old-LINERs and 2 Young-LINERs. Young-TOs or Young-LINERs are LLAGNs in which intermediate age stars contribute significantly to the nuclear blue-optical continuum. The dearth of Young-LINERs in the sample can be understood from the result obtained in our spectroscopic studies, which show that the overwhelming majority of LINERs harbor an old stellar population.  With these data we have built an atlas that includes the structural map for all the images, and color maps for those galaxies for which images in two filters are available. We have identified those galaxies that have nuclear compact sources, and we have studied the circumnuclear environment of LLAGNs.  We have found circumnuclear dust in 88$\\%$ of the LLAGNs, but this fraction is somewhat larger (95$\\%$) in weak-[OI] LLAGNs. The dust morphology is quite complex, from nuclear spiral lanes, to chaotic filaments and to nuclear disk-like structures. Chaotic filaments are as frequent as the well organized dust spirals; but disks are mainly seen in strong-[OI] LLAGNs. The dust concentration (simply graded by its location relative to a radius of 100-200 pc, is similar in weak- and in strong-[OI] because the fraction of LLAGNs with dust located only in the inner part is larger in Old-LLAGNs than in Young-LLAGNs. These results suggest an evolutionary dust sequence from Young-TOs to Old-LLAGNs.  We have found that LINERs and TOs have both similar central magnitude and surface brightness, but LLAGNs with young and intermediate age populations are brighter than Old-TOs and LINERs. We have not found any correlation between the presence of nuclear compact sources and the emission line spectral type, ie., LINERs are as frequently nucleated as TOs.  However, the centers of Young-TOs are brighter than the centers of Old-TOs and LINERs. The difference in magnitude and surface brightness can be even larger if we account for internal extinction, since Young-TOs are dustier. This result indicates that Young-TOs are separated from other type of LLAGNs also in terms of their central brightness, in addition of the properties and spatial distribution of the stellar population.  These data have been very useful to study the circumnuclear environment of LLAGNs, and to identify which of these galaxies have a nuclear compact source. The fact that compact sources are as frequent in LINERs as in TOs, confirms again that LLAGNs are a mixed bag of objects. These results also suggest that the central morphology alone is not sufficient to elucidate the origin of their central source, and it cannot be used to ascertain whether LLAGNs are powered by AGNs or stellar clusters. These data will be complemented with near-UV (ACS) and near-IR (NICMOS) images to provide a panchromatic atlas of the inner regions of LLAGNs and to further investigate the origin of the nuclear sources and their relation with stellar clusters.  {\\bf Acknowledgements} We thanks the referee for her/his suggestions that helped to improve the paper. RGD and EP acknowledge support from the Spanish Ministerio de Educaci\\'on y Ciencia through the grant AYA2004-02703. The data used in this work come from observations made with NASA/ESA Hubble Space Telescope, obtained from the STScI data archive. Basic research in radio astronomy at the NRL is supported by 6.1 base funding. We also thank support from a joint CNPq-CSIC bilateral collaboration grant.   %%%REF%%%REF%%%REF%%%REF%%%REF%%%REF%%%REF%%%REF%%%REF%%%REF%%%REF%%%"
    },
    "0710/0710.4992_arXiv.txt": {
        "abstract": "We establish a nonminimal Einstein-Yang-Mills-Higgs model, which contains six coupling parameters. First three parameters relate to the nonminimal coupling of non-Abelian gauge field and gravity field, two parameters describe the so-called derivative nonminimal coupling of scalar multiplet with gravity field, and the sixth parameter introduces the standard coupling of scalar field with Ricci scalar. The formulated six-parameter nonminimal Einstein-Yang-Mills-Higgs model is applied to cosmology. We show that there exists a unique exact cosmological solution of the de Sitter type for a special choice of the coupling parameters. The nonminimally extended Yang-Mills and Higgs equations are satisfied for arbitrary gauge and scalar fields, when the coupling parameters are specifically related to the curvature constant of the isotropic spacetime. Basing on this special exact solution we discuss the problem of a hidden anisotropy of the Yang-Mills field, and give an explicit example, when the nonminimal coupling effectively screens the anisotropy induced by the Yang-Mills field and thus restores the isotropy of the model. ",
        "introduction": "The discussion of a nonminimal coupling (NMC) of gravity with fields and media has a long history. The most intensely this topic has been studied  in connection with the problem of nonminimal coupling of gravity and scalar field, which has numerous cosmological applications. The details of the investigations of this problem are discussed, e.g., in the review of Faraoni {\\it et al\\/}.\\cite{FaraR} The development of the theory of NMC of gravity and scalar field $\\phi$ has started by the introduction of the term $\\xi \\phi^2 R$ to the Lagrangian ($R$ is the Ricci scalar). In Ref.~\\refcite{Chernikov} the special choice $\\xi = 1/6$ has been motivated by the conformal invariance; in Ref.~\\refcite{Callan} this quantity was considered as an arbitrary parameter of the model. Such a model has been widely used for the cosmological applications, in which $\\xi$ played a role of extra parameter of inflation (see, e.g., Refs.~\\refcite{Abbott}--\\refcite{Fara4}). In Refs.~\\refcite{HDehnen1}--\\refcite{HDehnen4} the gauge-invariant term $\\alpha {\\bf \\Phi}^{+}{\\bf \\Phi}R $ has been introduced instead of $\\xi \\phi^2 R$ in the context of non-Abelian gauge theory (${\\bf \\Phi}$~is a multiplet of scalar complex Higgs fields interacting with gravity and spinor matter.) Subsequent generalizations have been related to the replacement of $\\xi \\phi^2$ by the function $f({\\Phi}^2)$ (see, e.g., Refs.~\\refcite{Bergmann}--\\refcite{Steinh}), as well as, to the inserting of the terms of the type $F({\\Phi}^2, {\\cal R})$ both linear and nonlinear in the Ricci scalar, Ricci and Riemann tensors (see, e.g., Refs.~\\refcite{Linde}--\\refcite{Inagaki}). The idea of nonminimal derivative coupling introduced in Ref.~\\refcite{Amen3} and developed further in Refs.~\\refcite{Capo1,Capo2} has enriched the NMC modeling by the terms $\\phi_{,ij..}$. Nonminimal cosmological models based on the formalism of derivative coupling are the multi-parameter ones and have supplementary abilities for a fitting of observational data. Let us note that the NMC of gravity and scalar field leads to the modifications of both the Klein-Gordon and the Einstein equations, and such modifications are of interest for various inflation scenarios. Thus, the modeling of nonminimal interactions of scalar and gravitational fields is one of the well established and physically motivated branch of modern cosmology. Natural extension of the nonminimal theory from the models with scalar fields coupled to curvature to the models describing scalar fields interacting with gauge fields has the same sound motivation and can disclose new aspects of cosmological dynamics.  The study of the nonminimal coupling of gravity with electromagnetic field has another motivation and another history. In 1971 Prasanna\\cite{Prasa1} introduced the invariant $R^{ikmn}F_{ik}F_{mn}$ ($R^{ikmn}$ is the Riemann tensor, $F_{ik}$ is the Maxwell tensor) as a possible element of a Lagrangian, and then in Ref.~\\refcite{Prasa2} obtained the corresponding nonminimal one-parameter modification of the Einstein-Maxwell equations. In 1979 Novello and Salim\\cite{Novello1}  proposed to insert the gauge non-invariant terms $R A^k A_k$ and $R^{ik}A_i A_k$ in the Lagrangian ($A_k$ is an electromagnetic potential four-vector). A qualitatively new step has been made by Drummond and Hathrell in Ref.~\\refcite{Drum}, where the one-loop corrections to the quantum electrodynamics (QED) are obtained, which take into account the nonminimal coupling of gravity and electromagnetism. The Lagrangian of such a theory happens to contain three fundamental $U(1)$-gauge-invariant scalars $R^{ikmn}F_{ik}F_{mn}$, $R^{ik}g^{mn}F_{im}F_{kn}$ and $RF_{mn}F^{mn}$ with  coefficients reciprocal to the square of the electron mass. This Lagrangian had no arbitrary parameters, but curvature induced modifications of the electrodynamic equations gave the impetus to wide discussions about the formal structure of the nonminimal Lagrangian, basic evolutionary equations, breaking the conformal invariance and the properties of the photons, coupled to curvature in  different gravitational backgrounds (see, e.g., Refs.~\\refcite{Acci1}--\\refcite{Lafrance}). The last paper revived, as well, the interest to the paradigm: curvature coupling and equivalence principle, various aspects of which are now discussed (see, e.g., Refs.~\\refcite{Prasa3,Solanki}). The QED-motivation of the use of the generalized Maxwell equations can also be found in the papers of Kosteleck\\'y and colleagues.\\cite{Kost1,Kost2} The effect of birefringence induced by curvature, first discussed in Ref.~\\refcite{Drum}, and some of its consequences for the electrodynamic systems have been investigated in Refs.~\\refcite{Balakin1}--\\refcite{Balakin5} for the case of pp-wave background. The generalization of the idea of nonminimal interactions to the case of torsion coupled to the electromagnetic field has been made in Refs.~\\refcite{Hehl1,Hehl2} (see, also, Ref.~\\refcite{Hehl3} for a review on the problem). To summarize we stress that the study of electrodynamic systems nonminimally coupled to the gravity field poses a natural question about curvature induced variations of photon velocity in the cosmological background. Since the interpretation of observational data in cosmology depends essentially on the velocity of photon propagation during different cosmological epochs, the modeling of nonminimal electrodynamic phenomena seems to be well motivated and interesting from physical point of view.  Concerning the nonminimal Einstein-Yang-Mills (EYM) theory, we can distinguish between two different ways to establish it. The first way is the direct nonminimal generalization of the Einstein-Yang-Mills (EYM) theory. In the framework of this approach Horndeski\\cite{Horn} and M\\\"uller-Hoissen\\cite{MH} obtained the nonminimal one-parameter EYM model from a dimensional reduction of the Gauss-Bonnet action. Now the Gauss-Bonnet models are of great interest in connection with the problem of dark energy (see, e.g., the Gauss-Bonnet model with nonminimal scalar field\\cite{OdinDE}). Thus, the non-Abelian multi-parameter extensions of nonminimal models are also well motivated, since they give a chance to explain the accelerated expansion of the Universe without addressing to exotic substance. We follow the alternative way, which is connected with a non-Abelian generalization of the nonminimal Einstein-Maxwell theory  along the lines proposed by Drummond and Hathrell\\cite{Drum} for the linear electrodynamics. Based on the results of Ref.~\\refcite{BL05} a three-parameter gauge-invariant nonminimal EYM model linear in curvature is considered.\\cite{1BZ06}\\cdash\\cite{BDZ07} Our goal is to formulate a nonminimal Einstein-Yang-Mills-Higgs (EYMH) theory, and this process, of course, also admits different approaches. In fact, the nonminimal EYMH theory should accumulate the ideas and methods both from the nonminimally extended EYM theory and from the nonminimally extended scalar field theory. Initial attempt to develop nonminimal EYMH theory can be found, for instance, in Ref.~\\refcite{Bij}, where the scalar Higgs field is nonminimally coupled with gravity via $\\xi {\\Phi}^2 R$ term, and the Higgs field ${\\bf \\Phi}$ is included into the Lagrangian of the Yang-Mills field in a composition with a square of the Yang-Mills potential: ${\\Phi}^2 A_k^{(a)} A^k_{(a)}$. Such a theory is not gauge-invariant.  In this paper we establish a new six-parameter nonminimal Einstein-Yang-Mills-Higgs model. First three coupling parameters, $q_1$, $q_2$ and $q_3$, describe a nonminimal interaction of Yang-Mills field and gravitational field. The fourth and fifth parameters, $q_4$ and $q_5$, describe the so-called gauge-invariant nonminimal derivative coupling of the Higgs field with gravity. Since the gauge-invariant derivative, $\\D_m {\\Phi}^{(a)}$, contains the potential of the Yang-Mills field, the corresponding nonminimal term is associated with ``triple'' interaction, namely, gravitational and scalar fields, gauge and scalar fields, and gauge and gravitational fields. The sixth parameter, $\\xi$, is the well-known coupling parameter nonminimally connecting gravitational and scalar fields via the term $\\xi R {\\Phi}^2$. Of course, this model is only one of a wide class of the nonminimal EYMH models. As for its motivation and possible physical applications, one can see that on the one hand, the interest to a six-parameter nonminimal EYMH model is based on the sound results obtained earlier in the framework of partial nonminimal models (Einstein-Maxwell, Einstein-Yang-Mills and scalar field theories), on the other hand, the six-parameter model under discussion shows new specific solutions of cosmological type, which can not appear in more simple models.  The paper is organized as follows. In Sec. \\ref{Formalism} we formulate the nonminimal EYMH model, which contains six phenomenological coupling parameters, and establish the nonminimally extended Yang-Mills, Higgs and Einstein equations. In Sec. \\ref{IsModel} we apply the introduced master equations to the spacetime with constant curvature and obtain the specific relationships between coupling constants, which turn the extended equations for the gauge field and scalar field into identities. In Subsec. \\ref{dSsptime} we discuss the exact solutions to the nonminimal EYMH equations attributed to the isotropic cosmological model with Yang-Mills field, characterized by hidden anisotropy. ",
        "conclusions": "\\label{Discussion} \\noindent \\par 1. The main mathematical result of the presented paper is the establishing of a new self-consistent nonminimal system of master equations for the coupled Yang-Mills, Higgs and  gravity fields from the gauge-invariant nonminimal Lagrangian (\\ref{1act}). The obtained mathematical model contains six arbitrary parameters, and, thus, admits a wide choice of special sub-models interesting for the applications to the nonminimal cosmology (isotropic and anisotropic) and nonminimal colored spherical symmetric objects. The applications require the phenomenological coupling constants $q_1$, $q_2,\\ \\dots,\\ q_5$ and $\\xi$ to be interpreted adequately. Following the idea, discussed in Ref.~\\refcite{HDehnen4}, we intend not to introduce ``new constants of Nature'', but to relate the phenomenological parameters with the constants well-known in the High Energy Particle Physics, on the one hand, and with the constants of cosmological origin, on the other hand. Indeed, in the specific cosmological model, established above, the sixth phenomenological parameter $\\xi$ is expressed in terms of the square of the effective mass of the Higgs bosons $\\mu$ and constant curvature $K$, $\\xi = \\frac{\\mu}{12K}$. Other parameters are expressed in terms of $K$ (see (\\ref{q})). Since in the de Sitter model the Hubble constant is $H=\\sqrt{K}$, one can say that $q_1$, $q_2,\\ \\dots,\\ q_5$ are connected with $H$. Analogously, one can consider the equality $H^2=K=\\frac{\\Lambda}{3}$ and thus, one can say that they are connected with the cosmological constant $\\Lambda$. In any case the parameters of nonminimal coupling $q_1$, $q_2,\\ \\dots,\\ q_5$ can be expressed in terms of cosmological parameters $K$, $H$ or $\\Lambda$, and define a specific radius of curvature coupling, $r_q \\equiv \\frac{1}{\\sqrt{K}}$ and the corresponding time parameter $t_q \\equiv r_q/c$.  2. The curvature coupling modifies the master equations for the Yang-Mills and Higgs fields. According to (\\ref{Heqs}) a new tensor ${H}^{ik}_{(a)}$ appears (see (\\ref{HikR})), which is an analog of the induction tensor in the Maxwell theory\\cite{Maugin}. This means that the curvature coupling of the non-Abelian gauge field with gravity acts as a sort of quasi-medium with a nonminimal susceptibility tensor ${\\cal R}^{ikmn}$ (see (\\ref{HikR})). As well, the curvature coupling modifies the master equations for the Higgs field, and the tensor $\\Re^{mn}$, according to (\\ref{Heq}), can be indicated as a simplest nonminimal susceptibility tensor for the Higgs field, and the vector ${\\Psi}^{m}_{(a)}$ (see (\\ref{21Heq})) can be defined as scalar induction. For the specific set of coupling constants (see (\\ref{AQU}), (\\ref{q})) the non-Abelian induction $H^{ik}_{(a)}$ and the scalar induction ${\\Psi}^{m}_{(a)}$ can turn into zero, despite the fact that the Yang-Mills field strength ${F}^{ik}_{(a)}$ and the Higgs field ${\\Phi}^{(a)}$ are non-vanishing. This means that, when (\\ref{AQU}) holds, the possibility exists to satisfy the nonminimally extended Yang-Mills and Higgs equations for arbitrary ${F}^{ik}_{(a)}$ and ${\\Phi}^{(a)}$. This possibility gives, in principle, a new option for modeling physical processes in Early Universe and shows very interesting analogy between this nonminimal model and resonance phenomena in plasma physics. Indeed, when we deal with plasma waves (for instance, with the longitudinal waves) one can see that electric induction $\\vec{D}$ is connected with the longitudinal electric field $\\vec{E}_{||}$ with the frequency $\\omega$ by the relation $\\vec{D}= \\varepsilon_{||}\\vec{E}_{||}$. Here $\\varepsilon_{||}$ is the longitudinal dielectric permittivity, the simplest expression for this quantity can be obtained in the limit of long waves and gives $\\varepsilon_{||}= 1-\\frac{\\Omega^2_{p}}{\\omega^2}$, where $\\Omega_{p}$ is the well-known plasma frequency. When $\\omega=\\Omega_{p}$, one obtains $\\vec{D}=0$ and electrodynamic equations are satisfied for arbitrary $\\vec{E}_{||}$. Analogous feature can be found in the nonminimal model described above (see Eq. (\\ref{1simpa})). Indeed, the quantity $K$ with the dimensionality of squared frequency ($c=1$) can be regarded as an analog of $\\Omega^2_{p}$, the quantity $2(6q_1+3q_2+q_3)$ can be indicated as $1/\\omega^2$, then the term $1-2K(6q_1+3q_2+q_3)$ plays a role of effective permittivity scalar $\\varepsilon_q$. When this effective permittivity scalar vanishes, i.e., when the constants of nonminimal coupling are connected with the constant curvature $K$ according to (\\ref{AQU}), we obtain the resonance case, for which the Yang-Mills and Higgs equations are satisfied identically for arbitrary strength field tensor $F^{ik}_{(a)}$ and Higgs multiplet $\\Phi^{(a)}$, the color induction $H^{ik}_{(a)}$ being equal to zero.   3. The vector potential of the Yang-Mills field $ A^{(a)}_i$ enters the master equations via the gauge covariant derivative $\\hat{D}_k$, thus, the gauge field generates an anisotropy in the spacetime. Such an anisotropy, in general case, breaks down the symmetry of the model and produces the isotropy violation. Nevertheless, as it was shown above, the nonminimal coupling can effectively screen the anisotropy and guarantee the symmetry conservation. In the framework of this model one can speak about hidden anisotropy of the Yang-Mills field, keeping in mind that non-Abelian gauge field enters the master equations for the gravity field in the isotropic combinations only."
    },
    "0710/0710.3867_arXiv.txt": {
        "abstract": "High speed collisions, although current in clusters of galaxies, have long been neglected, as they are believed to cause little damages to galaxies, except when they are repeated, a process called ``harassment\". In fact, they are able to produce faint but extended gaseous tails. Such low-mass, starless, tidal debris may become detached and appear as free floating clouds in the very deep HI surveys that are currently being carried out. We show in this paper that these debris possess the same apparent properties as the so-called ``Dark Galaxies\", objects originally detected in HI, with no optical counterpart, and presumably dark matter dominated. We present a numerical model of the prototype of such Dark Galaxies -- VirgoHI21 --, that is able to reproduce its main characteristics: the one-sided tail linking it to the spiral galaxy NGC 4254, the absence of stars, and above all the reversal of the velocity gradient along the tail originally attributed to rotation motions caused by a massive dark matter halo and which we find to be consistent with simple streaming motions plus projection effects. According to our numerical simulations, this tidal debris was expelled 750~Myr ago during a fly-by at 1100~km~s$^{-1}$ of NGC~4254 by a massive companion which should now lie at a projected distance of about 400~kpc. A candidate for the intruder is discussed. The existence of galaxies that have never been able to form stars had already been challenged based on theoretical and observational grounds. Tidal collisions, in particular those occurring at high speed, provide a much more simple explanation for the origin of such putative Dark Galaxies. ",
        "introduction": "With the availability of unprecedented deep HI blind surveys, a population of apparently free floating HI clouds without any detected stellar counterpart has become apparent \\citep{meyer04,davies04,things,alfalfa1,alfalfa2}. It has been suggested that a fraction of them could be ``dark galaxies\", a putative family of objects that would consist of a baryonic disk rotating in a dark matter halo, but that would differ from normal galaxies by being free of stars, having all their baryons under the form of gas. They would thus be ``dark'' in the optical and most other wavelengths, but visible through their HI emission, contrary to pure ``dark matter'' halos. Such dark galaxies would be extreme cases of Low Surface Brightness Galaxies (LSBs), a class of objects that have a particularly faint stellar content compared to their gaseous and dynamical masses \\citep[e.g.,][]{carignanfreeman88}. The formation of low mass dark galaxies is actually predicted by $\\Lambda$-CDM models \\citep[e.g.,][]{vandenbosch03,tully05}. \\citet{taylor05} provided theoretical arguments against the existence of galaxies that would have remained indefinitely stable against star formation, unless they are of very low mass, at least a factor ten below than that of classical dwarf galaxies.  If they exist, the dark galaxies are predicted to have a low dynamical mass and HI content. In the Local Group, some possibly rotating, high-velocity clouds were speculated to be dark galaxies \\citep{simon04,simon06}. Further away, \\citet{davies06} argued that most previous HI blind surveys were not sensitive enough to rule out the existence of dark galaxies. And indeed, while the HIPASS survey failed at detecting HI~clouds without optical counterparts \\citep{doyle05}, deeper recent HI observations, in particular with the Arecibo telescope, have revealed a number of dark galaxy candidates \\citep{alfalfa2}. Among these free-floating low-mass HI clouds, one object located in the outerskirts of the Virgo Cluster has attracted much attention and discussion: VirgoHI21 \\citep[][see Fig.~1]{davies04,minchin05}. Despite a HI mass of only $\\sim 10^8$~M$_\\sun$, this elongated gaseous structure, mapped with the Westerbork Synthesis Radio Telescope (WSRT) by \\citet[][hereafter M07]{minchin07}, exhibits a velocity gradient as large as 220~km~s$^{-1}$ (see Fig.~\\ref{fig:pv-vhi21}). {\\it Assuming} that the observed HI velocities trace rotation, the inferred dynamical mass would be as large as $\\sim 10^{11}$~M$_\\sun$. The object shows no optical counterpart, even on deep HST images (M07). With such extreme properties, VirgoHI21 has become the prototype for dark galaxies, although its high dynamical mass is atypical even in models predicting the existence of Dark Galaxies. If real, an object like VirgoHI21 could tidally disturb the galaxies in their neighborhood, as investigated by \\cite{karachentsev06}. Actually VirgoHI21 itself lies at about 150~kpc from the massive spiral galaxy NGC~4254 (M~99), to which it is connected by a faint HI~filament. This structure could in principle be a bridge linking the two galaxies and would then appear as a sign of a tidal interaction between them (M07).   \\begin{figure*} \\centering \\includegraphics[width=\\textwidth,angle=0]{f1.pdf} \\caption{The system VirgoHI21/NGC 4254. {\\it Left} The distribution of the atomic hydrogen is superimposed in blue on a true color optical SDSS image of the field acquired through the WIKISKY.org project. The HI maps actually combine two data set: the observations obtained with the Arecibo telescope as part of the ALFALFA project \\citep[courtesy of B. Kent,][]{haynes07}, which are sensitive enough to show the whole extent of the gaseous tail; the observations obtained at WSRT \\citep[courtesy of R. Minchin,][]{minchin07} which are not as deep, but have a much better spatial resolution. The HI maps were smoothed and the WRST data have been deconvolved by the elongated telescope beam. {\\it Right} Color coded first moment, velocity, field of the HI tail as mapped by Arecibo. The observed velocity range in km~s$^{-1}$ is indicated to the right.} \\label{fig:vhi21} \\end{figure*}  \\begin{figure} \\centering \\includegraphics[width=\\columnwidth,angle=0]{f2.pdf} \\caption{Observed Position (x-axis) Velocity (y-axis) diagram along the position angle 22 degrees corresponding to the main direction of the HI bridge. As for Figure~\\ref {fig:vhi21}, the WSRT (gray) and Arecibo (light blue) data set have been superimposed to show all available information.}\\label{fig:pv-vhi21} \\end{figure}  However, starless isolated gas clouds, showing a large velocity spread, are not necessarily genuine dark galaxies. Ram pressure can strip gas away from spirals in the vicinity of clusters, a process that does not affect stars. Interaction with an external field, for instance that of another galaxy, can expulse large amounts of material from the disk in the form of gas-rich tidal tails and debris. In that vein, \\citet[][hereafter B05]{bekki05}  have suggested that interactions between flying-by (i.e. interacting without merging) galaxies orbiting in a potential well similar as the one produced by the Virgo Cluster  form tidal tails that after some time can resemble isolated gas clouds containing little stars. Furthermore tidal tails are the place of large streaming motions, as shown for instance by the observations of the merger prototype NGC~7252 by \\citet{hibbard94} and the models of \\citet{bournaud04} and B05. The remaining HI structures of galaxy collisions can thus not only appear as isolated HI clouds, but also exhibit large velocity gradients that mimic those expected for rotating disks within dark matter haloes.  In these conditions, serious doubts have raised whether VirgoHI21 is really a dark galaxy and not simply the result of a tidal interaction or of an harassment process in the Virgo Cluster (B05). They have recently been reinforced by the publication by \\citet{haynes07} of a deep HI map of the field, obtained with the Arecibo Telescope as part of the ALFALFA survey \\citep{alfalfa1,alfalfa2}. It reveals that VirgoHI21 actually lies within an even larger HI structure that extends further to the North, in the opposite direction of NGC~4254. This further suggests that the thin and long HI feature is in fact a tidal tail emanating from the spiral and that VirgoHI21 is just a denser cloud within it and thus not a dark galaxy.  On the basis of HI spectra obtained with the  Effelsberg telescope and a numerical model including the cluster ram pressure, \\citet{vollmer05} suggested that NGC~4254 has indeed interacted recently with a companion. However this study mostly  focussed on the internal properties of the spiral and in particular the formation of  VirgoHI21 was not modeled. Nevertheless, several arguments against a tidal origin for the HI cloud have been raised  and addressed in detail in M07. The main ones regard the absence of a suitable interacting companion,  the nonexistence of a counter tail -- a feature that is generally present in tidal interactions --,  the total lack of stars in the HI tail/bridge,  and, above all, the remarkable 200~km~s$^{-1}$ velocity gradient associated to VirgoHI21, which seems to be reversed and amplified with respect to the large scale velocity field along the rest of the HI structure.  We will show in this paper that most of these criticisms actually apply to low-velocity encounters and not to the high-velocity ones, which are common in the cluster environment. High-velocity collisions have been neglected so far because they seem to cause little disturbances to stellar disks unless they are very numerous and participate to an harassment process \\citep[e.g.,][]{moore96}. To further investigate the role of high speed collisions in the formation of tidal debris, especially the gaseous ones, we have carried out a series of numerical simulations. We illustrate their impact showing a numerical model reproducing the morphology and kinematics of VirgoHI21.    The numerical simulations are presented in Section~2. In Section~3, we compare the properties of tidal tails formed in high- and low-velocity galaxy encounters. In Section~4, we present the model which best fits VirgoHI21. Its actual nature is discussed in Section~5. Our conclusions and the implications for the detection of dark galaxies are summarized in Section~6. ",
        "conclusions": "\\subsection{From tidal tails to fake dark galaxies} \\label{fakedark} As shown in Section~\\ref{sect:highvel}, tidal tails, especially those formed during high-velocity encounters, share many of the properties expected for dark galaxies: starless HI features, strong velocity gradients due in one case to streaming motions and in the other to the presence of a massive dark matter halo. If furthermore some gaseous condensations are present in the tail, they may resemble isolated dark galaxies: indeed the bridge to the parent galaxy can be very faint and hard or impossible to detect on moderately deep HI maps.  Tidal tails do not necessarily have uniform profiles; denser regions can even lie in their outermost regions \\citep[e.g.,][]{DBM04}. This can be because large pre-existing clouds from the parent disk are moved into the tail where they can form local overdensities \\citep{elmegreen93}, or because some parts of an initially uniform tail condense under the effect of gravity \\citep{BH92}. That these regions will lie far from the progenitor disk is a natural consequence of the extended flat rotation curves of spirals \\citep{DBM04}. Depending on their location and mass, these denser parts of tails can even be self-gravitating, form stars and become finally independent objects with the mass of dwarf galaxies: the so-called Tidal Dwarf Galaxies \\citep[TDGs,][for a review]{duc07}. In such cases, the internal motions within the young gravitationally bound object, in particular its rotation, will induce an additional velocity gradient. On the other hand, the less massive condensations that have not reached the critical HI column density threshold to form stars will appear as detached HI clouds without any stellar counterpart.    Many of the known free-floating HI clouds, which are considered as Dark Galaxy candidates, have been found in clusters. In this environment, high velocity encounters have a high probability to occur, due to the large velocity dispersion of the cluster galaxies. Thus intra-cluster low-mass HI clouds of tidal origin could then be common but more dedicated studies should check this. Of course, this mechanism applies only if the parent galaxy had before the collision an extended HI disk. This requires in particular that it has not already crossed the cluster core where ram pressure would have contributed to truncate its gaseous disk. Tidal material generally quickly falls back onto the parent spiral, but this can take more than 2~Gyr in the outer parts \\citep[e.g.,][]{BD06}. The cluster tidal field can even prevent tidal debris from falling back \\citep{Mihos04}. This leaves time for fake dark galaxies of tidal origin to be observed while the interloper galaxy can be far away: at 1000~km~s$^{-1}$, it can be at a projected distance up to 2~Mpc two billion years later. \\bigskip   \\subsection{The nature of VirgoHI21}  After having argued that dark galaxies may in general be mistaken with tidal features and thus be fake ones, we discuss more specifically the nature of VirgoHI21, presenting pro and con arguments for the different scenarios proposed so far for its origin.   \\subsubsection{Tidal debris?}  M07 argued that the VirgoHI21 + HI~bridge system cannot be a tidal tail from the spiral NGC~4254, because of the following reasons: \\begin{itemize} \\item Interacting galaxies generally have pair of tidal tails, and indeed the models in B05 have two-tailed morphologies, while NGC~4254 has no counter-tail. \\item The tidal HI clouds in B05 models are star-poor but not star-free, while HST optical observations show VirgoHI21 being completely dark. \\item The velocity gradient along the HI~bridge gets reversed around the VirgoHI21 cloud, which seems to rotate in a direction opposite to the rest of the HI bridge. This reversal is said by M07 not to be explained by interaction models as those of B05. \\item The velocity spread ($\\Delta V=200$~km~s$^{-1}$) of VirgoHI21 could only be explained by an encounter at a comparable velocity (according to M07), implying that the interloper should not be further away than a few arcminutes and would have been identified. Moreover, low relative velocities are rare in the Virgo Cluster. \\item High velocity encounters are more common in this environment but velocities as large as $\\sim$~1000~km~s$^{-1}$ are ``far too large to generate tidal features such as bridges and tails\" (M07) . \\end{itemize}  If indeed VirgoHI21 is a genuine dark galaxy, as claimed by M07, the HI~bridge would then be a tail expulsed from the dark galaxy by an interaction with NGC~4254 or the cluster field. We note here that our numerical model of VirgoHI21 which suggests that the tidal debris rather emanate from NGC~4254 addresses each of these above-mentioned concerns:  \\begin{itemize} \\item The interaction occurred about 750~Myr ago, so that the counter-tail from the spiral is not necessarily expected to be observed anymore. By that time the interloper might also be far away. \\item The HI cloud formed in our model during a high--velocity encounter is free of old stars pre-existing to the galaxy interaction, and is not dense enough to form new stars. \\item The velocity spread ($\\Delta V=200$~km~s$^{-1}$) of the HI structure does not imply that the galaxy interaction had a similar velocity. It is in fact accounted for by a high-velocity encounter. \\item The reversal of the velocity gradient along the tail is reproduced thanks to projection effects. \\end{itemize}  Whereas the global kinematics of VirgoHI21 and its bridge can be simply explained by streaming motions along a tidal structure, in detail, the model and the observations show some differences. They may be noted when comparing the observed (Fig.~\\ref{fig:pv-vhi21}) and simulated (Fig.~\\ref{fig:posvel}) Position--Velocity diagrams along the tail; in particular, as put forward by M07, the velocity gradient towards VirgoHI21 is locally larger within the HI cloud, while in our model it is not more enhanced at this location than further away near the tip of the tail.      This local difference is actually not a real concern for our scenario. First, part of the local amplification of the gradient can result from the self-gravity of the VirgoHI21 cloud itself, that could be somewhat denser than in our model. In particular, as recently shown by \\citet{B07}, tidal debris may contain a significant dark component that has not been included in the simulations.  Alternatively, the velocity field can be disturbed by objects in the neighborhood, for instance by the nearby dwarf galaxy, SDSS J121804.26+144510.4, also known as object~C (see Fig.~\\ref{fig:vhi21} and M07). The nature of this dwarf is actually unclear (see Appendix). We suppose here that it is a pre-existing object, physically interacting with the system. Object~C has a visible HI mass of $2\\times 10^7$~M$_\\sun$. We included its possible influence in a simple model. We describe it as a dwarf spheroid with a Plummer profile, total mass of $4\\times 10^8$~M$_\\sun$ (to include the dark matter mass) and scale-length of 5~kpc. It is located 8~kpc East from VirgoHI21 on the sky plane, and we assume it lies 20~kpc below VirgoHI21 in the radial direction. Simulating the whole trajectory of this dwarf in the simulation from $t=0$ would introduce too many parameters. Since we just want to illustrate its possible local effect, we simply add it as a fixed mass in the late stage, linearly increasing its mass from $0$ at $t=600$~Myr to the final value at $t=700$. The velocity perturbation induced by object~C is shown on Figure~\\ref{fig:posvel}. The model qualitatively reproduces the local amplification of the velocity gradient at the position of VirgoHI21, and in particular the ``S-shape\" of the velocity profile visible in the Position-Velocity diagram (see Fig.~\\ref{fig:pv-vhi21}). Tuning the parameters of the simulation may help to further reproduce the exact kinematical feature. Obviously other objects, such as NGC~4262 -- the gas--rich galaxy to the East in Figure~\\ref{fig:vhi21} -- might have also crossed the trajectory of the tidal tail and slightly interacted with it. In any case, whatever the real explication is, the ``S-shape\" of the velocity profile of VirgoHI21 is not inconsistent with a tidal origin.  \\bigskip  A second critical issue is the identification of the interloper responsible for the collision. In our high-velocity scenario, the interloper now lies at a projected distance of 400~kpc to the WNW of NGC~4254. A massive spiral is in fact present near this position: NGC~4192 (M~98). This galaxy, which is seen close to edge-on, has a maximum rotation velocity corrected for inclination  of 236~km~s$^{-1}$, significantly higher than that of NGC~4254 (193~km~s$^{-1}$, according to the HyperLeda database), making it compatible with the 1.5:1 mass ratio used in our simulation. In our model, the interloper is today approaching us with a radial velocity of 715~km~s$^{-1}$ w.r.t. the target spiral. In the real Universe, NGC~4192 is indeed approaching, but with a relative velocity of $\\sim 2000$~km~s$^{-1}$ with respect to NGC~4254, i.e. much larger than in our model. This difference could be explained by the tidal field of the cluster which is not taken into account in our model. Adding it would introduce too many additional free parameters since the radial position and tangential velocity of the studied objects with respect to the cluster are unknown. The cluster gravitational field can modify some details in the properties of tidal tails in the long term \\citep{Mihos04}, but not their fundamental characteristics. However the cluster tidal field can have a more dramatic effect on the relative orbit of the distant galaxy pair over the nearly 750~Myr period since the encounter. Indeed, a cluster with the mass typical of Virgo, $M_\\mathrm{C}=10^{15}$~M$_\\sun$, at a distance $R_\\mathrm{C}=1$~Mpc from the pair, and a typical separation between the two galaxies $d\\sim 300$~kpc (which is the average 3-D separation from the time of the encounter until today) would create a tidal acceleration \\begin{equation} a_\\mathrm{t} \\simeq \\frac{G M_\\mathrm{C} d} {{R_\\mathrm{C}} ^3} \\end{equation} and, over a timescale of $T=750$~Myr, a relative velocity impulsion of \\begin{equation} \\Delta V_\\mathrm{t} \\simeq \\frac{G M_\\mathrm{C} d} {{R_\\mathrm{C}} ^3} \\times T \\end{equation} which would come in addition to the relative velocity found in our model. The values above imply $\\Delta V \\simeq 1100$~km~s$^{-1}$. This additionnal velocity difference would account for the observed relative velocities of NGC~4192/4254, and goes in the right direction if NGC~4254 lies behind the Virgo Cluster center.  NGC~4192 is thus a fully possible interloper, in spite of its large distance and relative velocity. This galaxy does not show a strongly disturbed HI disk on moderately deep HI maps \\citep{Chung05}. One reason could be unfavorable internal/orbital parameters, for instance a retrograde orbit. Alternatively, faint tidal debris may be present, but only visible on deep HI observations, similar to those that revealed the existence of VirgoHI21.  However, we do not intent to claim that NGC 4192 is the only possible interloper. Other combination of orbits/projection/age can certainly reproduce the properties of VirogHI21, and a galaxy flying away at $\\sim 1000$~km~s$^{-1}$ during $\\sim 1$~Gyr can be at a projected distance of 1~Mpc today, possibly even at the center of the Virgo Cluster. The number of massive interloper candidates is then large, making hard to identify the real culprit. This is anyway not required for our demonstration that VirgoHI21 can be a tidal debris, since we have shown that possible interlopers do exist.  \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{f6.pdf} \\caption{Position--Velocity diagram of the gas for Model 7 at $t=$750 Myr when it best reproduces the actual morphology and velocity field of the system VirgoHI21/NGC 4254. The arrow indicates the axis of the ``Position\" direction. The PV diagram of a model where ``Object C\" has been artificially added during the simulation is also shown. }\\label{fig:posvel} \\end{figure}    \\subsubsection{A kinematically decoupled Tidal Dwarf Galaxy?} The presence of a strong velocity gradient in a tidal tail, if not due to streaming motions, may actually pinpoint the presence of a gravitationally bound object that need not be a pre-existing dark-matter dominated galaxy. Massive substructures in tidal tails often become kinematically decoupled, self-gravitating and form new stars, becoming rotating Tidal Dwarf Galaxies.  This is the case for VCC~2062, a  TDG candidate in Virgo \\citep{Duc07b}. VirgoHI21 appears as a gas condensation within a tidal tail; it is currently not a TDG since it is starless, but could be its gaseous progenitor. Whether stars will be formed in this structure later is questionable. If the surface column density has remained unusually very low during the first several hundreds of Myr after the formation of VirgoHI21, it is unlikely that star-formation is ignited later on. On the other hand, the dynamical collapse time of the cloud, $1/\\sqrt{G \\rho}$  \\citep{elmegreen02},  is as large as 400 -- 500 Myr for an initially resting system and an average volume mass density in our modeled cloud of $\\rho \\sim 10^{-3}$~M$_{\\sun}$~pc$^{-3}$ (this is also about the density of the real VirgoHI21 cloud assuming a vertical scale-height of $\\sim$ 300~pc). Comparing this time scale  to  the age of the cloud,   about 500~Myr (it appears in the model at $t=200-300$~Myr), one may conclude that  VirgoHI21 would barely have had the time to collapse and form stars even in the most favorable conditions and incidentally  that the absence of stars today is not dependent on an arbitrary choice of the threshold.  In other words, the system could still be contracting under the effect of its internal gravity today, and begin to form stars later-on {\\it if} its density comes to exceed the star formation threshold.  However, the main argument against VirgoHI21 being {\\it yet} a TDG fully responsible for the observed velocity gradient is the large dynamical mass inferred from the rotation curve: in galaxies made out of collisional debris, the dynamical mass should be of the same order as the luminous one, even if the presence of dark baryons may cause some differences between them \\citep{B07}. In the case of VirgoHI21, the dynamical mass inferred from the velocity curve is more than a factor of 3000 greater than the luminous one, i.e. that of the HI component. Clearly, streaming motions provide a much more reasonable explanation for the kinematical feature observed near VirgoHI21, if indeed this object is of tidal origin.      \\subsubsection{Harrasment by the cluster field?} In the group environment, \\cite{bekki05b} proposed a scenario in which  the group tidal field is able to strip gas from HI-rich galaxies,  explaining the presence of isolated intergalactic HI clouds.  Following this idea, B05 presented a model in which the combined action of galaxy-galaxy interactions and the cluster tidal field produce debris with properties similar to Dark Galaxies. \\citet{haynes07} even suggested that VirgoHI21 and the whole HI structure would result  from just the  long-term harassment by the large-scale cluster potential. However, the tidal field exerted by a structure of mass $M$ and typical scale $R$ scales as $M/R^3$. The tidal field of the Virgo Cluster ($10^{15}$~M$_\\sun$, 1~Mpc) at the present distance of NGC~4254 is then more than ten times smaller than that of the interloper galaxy in our interaction model ($2\\times 10^{12}$~M$_\\sun$, 50~kpc). It is just unlikely that the cluster field can develop a tail as long as the galaxy interaction can do. The harassment process has a longer timescale than the galaxy pair interaction, but over long timescales the orientation changes, which hardly accounts for the single, thin and long tail around NGC~4254. This structure is more typical of a short and violent interaction like a close galaxy encounter than a weaker and longer process like the harassment by the global cluster field.   \\subsubsection{Ram pressure stripping?} \\label{rampressure}  Ram pressure exerted by the cluster hot gas may also expulse gas from the outer regions of spiral disks and create isolated HI clouds without any optical counterpart. Yet, this scenario suffers fundamental concerns, also pointed out by M07, in particular:  \\begin{itemize} \\item structures known to result from ram-pressure stripping are rather short and thick as observed in Virgo \\citep{crowl05,lucero05,vollmer06,chung07} and suggested by hydrodynamical simulations \\citep[e.g.,][]{Roediger07}, while the HI bridge of VirgoHI21 is much thinner and longer. \\item the kinematics is not directly explained too, in particular the reversing velocity gradient in the HI~bridge, which in the context of a tidal interaction results from streaming motions along a curved tail (in 3-D) more or less seen edge-on. This may also be the case for ram pressure but should be demonstrated by a model. \\end{itemize}  \\cite{vollmer05} put forward the role of ram pressure, combined with a tidal interaction,  in shaping the internal gas  distribution, with its $m=1$ structure,  and velocity field of NGC~4254. This partly explains  why,  in the innermost regions, the detailed kinematics of the spiral  presents   some differences with that of our model which did not take into account the intracluster medium. More recently, \\cite{Kantharia07} presented a low radio frequency continuum map  of the galaxy which is best explained invoking a ram pressure scenario. However so far, its possible contribution on the properties of the HI bridge and VirgoHI21 has not yet been investigated. The two above-mentioned papers do not claim  that their origin is  ram-pressure, and indeed our pure tidal model is  able to reproduce these features provided that the HI disk was originally much more extended than the optical radius -- an hypothesis  which was not adopted in \\cite{vollmer05}.   \\subsubsection{A dark galaxy?}  Showing that virgoHI21 {\\it can} be a tidal debris taking the appearance of a dark galaxy does not directly rule out the possibility that it is a real dark galaxy. The dark galaxy hypothesis however suffers several difficulties unexplained so far:  \\begin{itemize} \\item the maximal velocity gradient is not centered on the peak of the HI emission, but actually lies on one side of the VirgoHI21 cloud. This is unexpected for a rotating HI disk within a massive dark halo. One may propose that the on-going interaction with NGC~4254 causes this asymmetry in the velocity field, but that NGC~4254 is massive and close enough to induce such major disturbances remains to be shown. \\item within the assumption that VirgoHI21 is a real dark galaxy, the HI bridge is a tidal tail expulsed from it and captured by the more massive galaxy NGC~4254 (M07). The gas in the HI bridge, falling onto NGC~4254 from 150~kpc away, is not expected to have the same velocity as the local gas settled in rotation: it could be on retrograde, polar or direct orbits but with different velocities. However, the base of the HI bridge has a radial velocity coherent with the outer disk of NGC~4254 to which it is morphologically connected: the velocity step between the base of the tidal tail and the outer disk is in fact less than $50$~km~$^{-1}$, much smaller than the circular velocity there \\citep[see data in M07 and also][]{phookun93}. This further suggests that the HI material in the bridge comes from NGC~4254. \\end{itemize}  These facts are naturally explained by the tidal scenario proposed for VirgoHI21. Whether they can also be addressed with the Dark Galaxy hypothesis remains to be demonstrated, in particular with a numerical model.   Although challenged, the putative existence of Dark Galaxies as massive as VirgoHI21, has fostered a number of follow-up works. As discussed by \\citet{karachentsev06} such invisible ghost objects should tidally perturb galaxies in their neighborhood, explaining why a fraction of apparently isolated spiral stellar disks seem to show signs of an external perturbation. We note however that other mechanisms may account for them, such as accretion of diffuse gas \\citep{bournaud05m1}."
    },
    "0710/0710.0675_arXiv.txt": {
        "abstract": "We investigate the number and type of pulsars that will be discovered with the low-frequency radio telescope LOFAR. We consider different search strategies for the Galaxy, for globular clusters and for galaxies other than our own. We show an all-sky Galactic survey can be optimally carried out by {\\em incoherently} combining the LOFAR stations. In a 60-day all-sky Galactic survey LOFAR can find over a thousand pulsars, probing the local pulsar population to a very deep luminosity limit. For targets of smaller angular size, globular clusters and galaxies, the LOFAR stations can be combined coherently, making use of the full sensitivity. Searches of nearby northern-sky globular clusters can find large numbers of low luminosity millisecond pulsars (eg.\\ over 10 new millisecond pulsars in a 10-hour observation of M15). If the pulsar population in nearby galaxies is similar to that of the Milky Way, a 10-hour observation can find the 10 brightest pulsars in M33, or pulsars in other galaxies out to a distance of 1.2Mpc. ",
        "introduction": "Since the discovery of the first four pulsars with the Cambridge radio telescope, an ongoing evolution of telescope systems has doubled the number of known radio pulsars roughly every 4 years.  The next step in radio telescope evolution will be the use of large numbers of low-cost receivers that are combined to form an interferometer or a single dish. These telescopes, LOFAR \\citep{ls07}, the Allen Telescope Array \\citep{bow07} and the SKA \\citep{kram04}, create new possibilities for pulsar research.  Here we outline and compare strategies for targeting normal and millisecond pulsars, both in the disk and globular clusters of our Galaxy, and in other galaxies. ",
        "conclusions": "Because of its large area and wide beam on the sky, LOFAR probes the local population of pulsars to a very deep luminosity limit. A 60-day Galactic survey at 140MHz can find over a thousand new pulsars, disclosing the local low-luminosity population. If we add all antennae coherently the sensitivity increases even further; with this setup, millisecond pulsars in nearby globular clusters can be detected to much lower flux limits than previously possible. Assuming the pulsar population in other galaxies is similar to that in ours, we can detect periodicities or giant pulses from extragalactic pulsars up to several Mpcs away."
    },
    "0710/0710.3730_arXiv.txt": {
        "abstract": "Models of terrestrial planet formation in the presence of a migrating giant planet have challenged the notion that hot-Jupiter systems lack terrestrial planets. We briefly review this issue and suggest that hot-Jupiter systems should be prime targets for future observational missions designed to detect Earth-sized and potentially habitable worlds. ",
        "introduction": "Since the discovery of the first extrasolar planets \\citep{wolszczan,mayor}, astronomical techniques and observational baselines have advanced to the point where over 200 extrasolar planetary systems have been identified \\citep{butler}. Most detected exoplanets are in the giant planet mass range and it is now clear that our solar system is but one variant within a great diversity of planetary system architectures. One of the most surprising discoveries has been of a population of giant planets, the so-called \\emph{hot-Jupiters}, found orbiting in a region of extreme insolation very close ($r < 0.1~\\mathrm{AU}$) to their central stars and well within the radius of the original nebular snowline ($r \\approx 3 - 5~\\mathrm{AU}$) where giant planets are thought to form \\citep{pollack}. Hot-Jupiters are not uncommon: they amount to about a quarter of exoplanet discoveries, and are thought to provide evidence that protoplanets can migrate over large radial distances via tidal interactions with the protoplanetary disk \\citep[e.g.][]{lin1,lin2,ward2,nelson1}. Since the disk gas is observed to disperse within the first few Myr of the system's existence \\citep{haisch}, giant planets must form and migrate through the inner system within this period, which is considerably less than the $\\sim$~10--100~Myr thought to be required to complete terrestrial planet formation \\citep{chambers2,kleine,halliday,obrien}.  Test particle studies have shown that terrestrial planets external to a hot-Jupiter would have stable orbits \\citep{jones}, and \\citet{raymond1} have shown that they should be able to form, in the presence of a giant planet already at $\\sim 0.1$~AU, from any available pre-planetary material with a period ratio roughly $>$~3. However, until recently it has been a common assumption that terrestrial planets could not have formed in hot-Jupiter systems due to the disruptive effect of the giant planet's migration which is deemed to have cleared the inner system of planet-forming material \\citep[e.g.][]{armitage1}, prompting claims that the observed abundance of hot-Jupiters could be used to constrain the general abundance of habitable planets \\citep{ward1}, and even their galactic location \\citep{lineweaver1}.  This picture has been challenged by the work of two groups who have modeled terrestrial planet formation concurrently with, and following, an episode of giant planet migration \\citep{fogg1,fogg2,fogg3,fogg4,raymond2,mandell}. Their findings suggest that inner system solids disks are \\emph{not} destroyed by the intrusion of a migrating giant planet and that terrestrial planet formation can resume in the aftermath and run to completion.  In this paper, we briefly describe our model of terrestrial planet formation and show some typical results. ",
        "conclusions": "Our models predict that terrestrial planets might be routinely expected in hot-Jupiter systems, including within their habitable zones, and may be detectable by forthcoming missions such as Kepler, Darwin, SIM PlanetQuest and TPF."
    },
    "0710/0710.1503_arXiv.txt": {
        "abstract": "{Observations of water lines are a sensitive probe of the geometry, dynamics and chemical structure of dense molecular gas. The launch of Herschel with on board HIFI and PACS allow to probe the behaviour of multiple water lines with unprecedented sensitivity and resolution.} {We investigate the diagnostic value of specific water transitions in high-mass star-forming regions. As a test case, we apply our models to the AFGL\\,2591 region.} {A multi-zone escape probability method is used in two dimensions to calculate the radiative transfer. Similarities and differences of constant and jump abundance models are displayed, as well as when an outflow is incorporated.} {In general, for models with a constant water abundance, the ground state lines, i.e., $\\mathrm{1_{10}}$-$\\mathrm{1_{01}}$, $\\mathrm{1_{11}}$-$\\mathrm{0_{00}}$, and $\\mathrm{2_{12}}$-$\\mathrm{1_{01}}$, are predicted in absorption, all the others in emission. This behaviour changes for models with a water abundance jump profile in that the line profiles for jumps by a factor of $\\sim$\\,10-100 are similar to the line shapes in the constant abundance models, whereas larger jumps lead to emission profiles. Asymmetric line profiles are found for models with a cavity outflow and depend on the inclination angle. Models with an outflow cavity are favoured to reproduce the SWAS observations of the $\\mathrm{1_{10}}$-$\\mathrm{1_{01}}$ ground-state transition. PACS spectra will tell us about the geometry of these regions, both through the continuum and through the lines.} { It is found that the low-lying transitions of water are sensitive to outflow features, and represent the excitation conditions in the outer regions. High-lying transitions are more sensitive to the adopted density and temperature distribution which probe the inner excitation conditions. The Herschel mission will thus be very helpful to constrain the physical and chemical structure of high-mass star-forming regions such as AFGL\\,2591.} ",
        "introduction": " ",
        "conclusions": "We have constructed models to examine the excitation of water in the high-mass star-forming region AFGL\\,2591. Depending on the adopted density, temperature and abundance structure, a completely different set of line profiles and strengths is found. Hence, the line profiles are very sensitive to the adopted physical and chemical structure. We have found that ({\\it i}) the ground-state transitions 1$_\\mathrm{10}$-1$_\\mathrm{01}$, 2$_\\mathrm{12}$-1$_\\mathrm{01}$ and 1$_\\mathrm{11}$-0$_\\mathrm{00}$, with relatively low upper energy levels ($\\lesssim$\\,110\\,K), become highly optically thick in the outer regions. The line profiles for these transitions, are mainly dominated by the emission from the outer regions, and are therefore not useful to put constraints on the water abundance in the inner regions. However, ({\\it ii}) the emission from lines with higher upper energy levels is dominated by the emission originating in the inner regions, and are therefore useful to probe the water abundance in the warm inner regions. ({\\it iii}) For models with an outflow cavity, the outflow feature (blue peak less strong than the red peak) is best seen in the ground-state transitions of o- and p-H$_\\mathrm{2}$O. ({\\it iv}) The influence of a moderate disk (few 100 AU in size) in the centre of the AFGL\\,2591 region does not change the water line profiles and strengths within the Herschel beam.  The Herschel mission will thus greatly help to understand the structure of high-mass protostellar objects, and consequently the formation process of high-mass stars."
    },
    "0710/0710.1359_arXiv.txt": {
        "abstract": "In a systematic study, we compare the density statistics in high-resolution numerical experiments of supersonic isothermal turbulence, driven by the usually adopted solenoidal (divergence-free) forcing and by compressive (curl-free) forcing. We find that for the same rms Mach number, compressive forcing produces much stronger density enhancements and larger voids compared to solenoidal forcing. Consequently, the Fourier spectra of density fluctuations are significantly steeper. This result is confirmed using the $\\Delta$-variance analysis, which yields power-law exponents $\\beta\\!\\sim\\!3.4$ for compressive forcing and $\\beta\\!\\sim\\!2.8$ for solenoidal forcing. We obtain fractal dimension estimates from the density spectra and $\\Delta$-variance scaling, and by using the box counting, mass size and perimeter area methods applied to the volumetric data, projections and slices of our turbulent density fields. Our results suggest that compressive forcing yields fractal dimensions significantly smaller compared to solenoidal forcing. However, the actual values depend sensitively on the adopted method, with the most reliable estimates based on the $\\Delta$-variance, or equivalently, on Fourier spectra. Using these methods, we obtain $D\\!\\sim\\!2.3$ for compressive and $D\\!\\sim\\!2.6$ for solenoidal forcing, which is within the range of fractal dimension estimates inferred from observations ($D\\!\\sim\\!2.0\\dots2.7$). The velocity dispersion to size relations for both solenoidal and compressive forcings obtained from velocity spectra follow a power law with exponents in the range $0.4\\dots0.5$, in good agreement with previous studies. ",
        "introduction": "Observations provide velocity dispersion to size relations for various molecular clouds (MCs), which document the existence of supersonic random motions on scales larger than $\\sim\\!0.1\\,\\mathrm{pc}$ \\citep[e.g.,][]{Larson1981,Myers1983,PeraultFalgaronePuget1986,SolomonEtAl1987,FalgaronePugetPerault1992,HeyerBrunt2004}. These motions are associated with compressible turbulence \\citep[e.g.,][]{ElmegreenScalo2004,ScaloElmegreen2004,MacLowKlessen2004} in the interstellar medium \\citep{Ferriere2001} and exhibit a single turbulent cascade or spatially separated coexisting inertial ranges \\citep{PassotPouquetWoodward1988} similar to the kinetic energy cascade of incompressible \\citet{Kolmogorov1941c} turbulence. However, there are various physical processes (e.g., self-gravity, magnetic fields, nonequilibrium chemistry) and especially the compressibility of the gas, that alter the scaling laws \\citep[e.g.,][]{Fleck1996} and statistics \\citep[e.g., intermittency corrections measured by][]{HilyBlantFalgaronePety2008} established for incompressible turbulence.  The physical origin and characteristics of the turbulent fluctuations are still a matter of debate. To advance on the question of how turbulence isdriven in the interstellar medium, we present results of high-resolution numerical experiments of supersonic isothermal turbulence comparing two distinct and extreme ways of driving the turbulence in a systematic study: 1) solenoidal forcing (divergence-free or rotational forcing), and 2) compressive forcing (curl-free or dilatational forcing).  Various numerical and analytical studies have provided important insight into the statistics of supersonic isothermal turbulence \\citep[e.g.,][]{PorterPouquetWoodward1992,Vazquez1994,PadoanNordlundJones1997,PassotVazquez1998,StoneOstrikerGammie1998,MacLow1999,Klessen2000,OstrikerStoneGammie2001,BoldyrevNordlundPadoan2002,LiKlessenMacLow2003,PadoanJimenezNordlundBoldyrev2004,JappsenEtAl2005,BallesterosEtAl2006,KritsukEtAl2007,LemasterStone2008}. Most of these studies use purely solenoidal or weakly compressive kinetic energy injection mechanisms (forcing) to excite turbulent motions. In the present study, we aim at comparing the usual case of solenoidal (divergence-free) forcing with the case of fully compressive (curl-free) forcing. The actual way of turbulence production in real MCs is expected to be far more complex compared to what we can model with the present simulations, probably consisting of a convolution of various agents producing turbulence, and mixtures of solenoidal and compressive modes \\citep[e.g.,][]{ElmegreenScalo2004,MacLowKlessen2004}. Here, we systematically investigate the extreme cases of purely solenoidal versus purely compressive energy injection.  Analyzing the density correlation statistics and fractal structure obtained in our hydrodynamic simulations, we show that compressive forcing leads to significantly steeper density fluctuation spectra and consequently to fractal dimensions of the turbulent gas structures, that are significantly smaller compared to the usually adopted solenoidal forcing. We use Fourier analysis, $\\Delta$-variance analysis, structure functions, the fractal mass size, box counting, and perimeter area methods to obtain fractal dimension estimates. We apply the $\\Delta$-variance analysis to both our 3-dimensional data and to 2-dimensional projections, and the perimeter area method to projections and slices through the turbulent density structures supporting the result of a significantly smaller fractal dimension for compressive forcing compared to solenoidal forcing. Although compressive forcing yields significantly smaller fractal dimensions than solenoidal forcing, the estimated fractal dimensions are in the range $2.0\\dots2.7$ consistent with observational estimates \\citep[e.g.,][]{ElmegreenFalgarone1996,SanchezEtAl2007}  We explain our numerical method, construction of solenoidal and compressive forcing fields and fractal analysis techniques in Section~\\ref{sec:methods}. In Section~\\ref{sec:results}, we show that our results are consistent with previous studies using solenoidal forcing, whereas compressive forcing yields much stronger density contrasts and consequently leads to significantly smaller fractal dimensions. In Section~\\ref{sec:conclusions}, we summarize our conclusions. ",
        "conclusions": "\\label{sec:conclusions}  We have presented results of two high-resolution ($1024^3$ grid cells) hydrodynamic simulations of supersonic isothermal turbulence driven to rms Mach numbers $\\mathcal{M}\\!\\sim\\!5.5$. The first simulation uses the typically adopted solenoidal (divergence-free) forcing to excite turbulent motions, whereas the second one uses compressive (curl-free) forcing. We have shown that compressive forcing yields much stronger density contrasts compared to solenoidal forcing for the same rms Mach number. This implies that the turbulence production mechanism leaves a strong imprint on compressible turbulence statistics, especially altering the density statistics. Our results particularly suggests that the mixture of solenoidal and compressive modes of the turbulence forcing must be taken into account. We summarize our results as follows:  \\begin{itemize}  \\item The velocity Fourier spectra exhibit power laws in the inertial range for solenoidal and compressive forcing. The slopes obtained for both forcings are significantly steeper ($\\sim\\!1.9$) compared to the Kolmogorov slope ($5/3$), in agreement with previous studies \\citep[e.g.,][]{KritsukEtAl2007,SchmidtEtAl2008} and in agreement with velocity dispersion to size relations inferred from observations \\citep[e.g.,][]{Larson1981,FalgaronePugetPerault1992,HeyerBrunt2004,PadoanEtAl2006}.  \\item From the integral of the density fluctuation Fourier spectra and from the asymptotic behavior of the 2nd order density structure function, we obtained the standard deviation of the density distribution $\\sigma_\\rho$. Compressive forcing yields a standard deviation $\\sim\\!$ three times larger compared to solenoidal forcing, in agreement with the results found in our previous study analyzing density probability distribution functions \\citep{FederrathKlessenSchmidt2008} and in agreement with the studies by \\citet{PassotVazquez1998}, \\citet{KritsukEtAl2007}, \\citet{BeetzEtAl2008} and \\citet{SchmidtEtAl2008}.  \\item The density fluctuation Fourier spectra are significantly steeper for compressive forcing in the inertial range compared to solenoidal forcing. Consistent results were obtained using complementary analysis methods, i.e., by comparing the $\\Delta$-variances \\citep{OssenkopfKripsStutzki2008a} and the 2nd order structure functions of the density field. Our estimates of density spectra for solenoidal forcing are in agreement with previous studies, e.g., the weakly magnetized super-Alfv\\'enic supersonic MHD models by \\citet{PadoanEtAl2004} and \\citet{KowalLazarianBeresnyak2007}, and consistent with the hydrodynamic estimates by \\citet{KritsukNormanPadoan2006}. Although a comparison with observational results must be regarded with caution due to systematic uncertainties, our results for solenoidal and compressive forcing are in the range of inferred scaling exponents by observations \\citep[e.g.,][]{BenschStutzkiOssenkopf2001}.  \\item From the scaling of the density fluctuation Fourier spectra and the $\\Delta$-variance applied to the 3-dimensional data and applied to 2-dimensional projections, we obtained fractal Hurst exponents following the analysis by \\citet{StutzkiEtAl1998}. This implies fractal box counting and fractal perimeter area dimensions significantly smaller for compressive forcing compared to solenoidal forcing (see Table~1).  \\item We analyzed the density structure using the fractal mass size method as introduced by \\citet{KritsukEtAl2007}. Compressive forcing yields a smaller fractal mass dimension compared to solenoidal forcing. The mass size method is, however, particularly sensitive to the temporal fluctuations of density peaks. Given the large uncertainties, our results using this method are roughly consistent with the estimates by \\citet{KritsukEtAl2007} and \\citet{KowalLazarian2007}.  \\item We analyzed the fractal density structure using the box counting method described in Section~\\ref{sec:box-counting} and the perimeter area method (Section~\\ref{sec:perimeter-area}) applied to projections and slices. We recover the significant differences between solenoidal and compressive forcing inferred from the density spectra and $\\Delta$-variance analysis. However, the box counting dimension varies strongly with the defining density threshold. The perimeter area dimensions obtained from slices are roughly consistent with the computed perimeter area dimensions from the $\\Delta$-variance given the systematic uncertainties (of order $\\sim\\!0.1$ for fractal dimension estimates) comparing different methods. The range of fractal dimensions obtained is consistent with the observations analyzed by \\citet{ElmegreenFalgarone1996} suggesting an overall fractal dimension of interstellar clouds in the range $D\\!\\sim\\!2.3\\pm0.3$.  \\end{itemize}"
    },
    "0710/0710.3454_arXiv.txt": {
        "abstract": "{There are now four low mass X-ray binaries with black holes which show twin resonant-like HFQPOs. Similar QPOs might have been found in Sgr A$\\sp *$. I review the power spectral density distributions of the three X-ray flares and the six NIR flares published for  Sgr A$\\sp *$ so far, in order to look for more similarities than just the frequencies between the microquasar black holes and Sgr A$\\sp *$. The three X-ray flares of Sgr A$\\sp *$ are re-analysed in an identical way and white noise probabilities from their power density distributions are given for the periods reported around $\\sim$ 1100 s. Progress of the resonant theory using the anomalous orbital velocity effect is summarized. ",
        "introduction": "% \\label{sect:intro} Quasi-periodic oscillations (QPOs) are believed to arise from variations of the accretion flow around compact objects, i.e. white dwarfs, neutron stars and black holes. As far as black holes are concerned, Remillard and McClintock (2006) have compiled a list of a total of 40 sources which show up in galactic X-ray binaries, 20 of which show a dark companion with a mass largely exceeding the mass of a neutron star but apparently limited to about 18 solar masses. Out of the 20 sources with established dynamically measured masses 16 sources show low frequency QPOs, whereas high frequency QPOs (HFQPOs, $\\nu$ $>$ 10 Hz) have been detected in 8 sources. ",
        "conclusions": ""
    },
    "0710/0710.4274_arXiv.txt": {
        "abstract": "We performed MHD simulations of very light bipolar jets with density contrasts down to $10^{-4}$ in axisymmetry, which were injected into a medium of constant density and evolved up to $200$ kpc ($200\\: r_{\\mathrm{j}}$) full length. These jets show weak and roundish bow shocks as well as broad cocoons and thermalize their kinetic energy very efficiently. We argue that very light jets are necessary to match low-frequency radio observations of radio lobes as well as the bow shocks seen in X-rays. Due to the slow propagation, the backflows and their turbulent interaction in the midplane are important for a realistic global appearance. ",
        "introduction": "During the last years, simulations of extragalactic jets with reasonable resolution and realistic sizes became computationally feasible, which makes comparisons between simulated and observed properties possible \\citep{Saxton2002,Zanni2003,Carvalho2005,KrauseVLJ2,ONeill2005}. Unfortunately, the direct physical variables and the observed properties are rather hard to link, which leaves simulations with a wide range of parameters. Simulations are mainly governed by the initial setup of the density ratio between jet and ambient gas, the Mach number and the magnetic field. If the magnetic field is not dynamically dominant (though important), the density contrast is the most dominant parameter, but may be one of the hardest to measure. The thermal jet pressure has turned out to be of little importance in the very light jet limit \\citep{KrauseVLJ1}. As the (kinetic) power of a jet can be estimated from energies in X-ray bubbles, typical values of velocity, lifetime, jet radius and cluster gas densities indicate that density contrasts of $10^{-2}$ to $10^{-4}$ (or even lower) are necessary to describe real sources. Parameter studies support this further, if the global jet/cocoon/bow shock properties are compared. Thus, we concentrate on the very light jets with magnetic fields as another important ingredient. ",
        "conclusions": ""
    },
    "0710/0710.3381_arXiv.txt": {
        "abstract": "We present optical and X-ray data for the first object showing strong evidence for being a black hole in a globular cluster.  We show the initial X-ray light curve and X-ray spectrum which led to the discovery that this is an extremely bright, highly variable source, and thus must be a black hole.  We present the optical spectrum which unambiguously identifies the optical counterpart as a globular cluster, and which shows a strong, broad [O III] emission line, most likely coming from an outflow driven by the accreting source. ",
        "introduction": "Since the early days of X-ray astronomy, there has been considerable debate over whether globular clusters contained black holes.  With the discovery of Type I X-ray bursts from all globular clusters in the Milky Way with bright X-ray sources (starting with Grindlay et al. 1976), and their subsequent explanation as episodes of thermonuclear burning on the surfaces of neutron stars (Woosley \\& Taam 1976; Swank et al. 1977), it became clear that there was no evidence for any accreting black holes in the Milky Way's globular cluster system.  Intepretations of the observations have been taken in two directions. One is simply that given only 13 bright X-ray sources in the Milky Way's globular cluster system, it is not so unlikely for them all to have neutron star accretors, especially in light of the fact that about 10 times as many neutron stars as black holes are expected to be produced for most stellar initial mass functions.  The alternative is that dynamical effects are responsible for ejecting black holes from globular clusters.  Severe mass segregation is likely to take place for globular cluster black holes, as they should be many times heavier than all the other stars in the cluster.  This can lead to the formation of a ``cluster within a cluster'' where the heaviest stars (i.e. the black holes) feel negligble effects from the other stars in the cluster, which in turn leads to a cluster with a short evaporation timescale (Spitzer 1969).  Numerical calculations have found that this evaporation can be accelerated further due to binary processes (e.g Portegies Zwart \\& McMillan 2000).  Early results from the Chandra X-ray Observatory gave new hope that globular cluster black holes might be detectable, by opening up the window of looking in other galaxies.  Previously, only ROSAT could resolve point sources in other galaxies, and its localization of sources was generally not good enough to allow for unique identification of optical counterparts.  The first few years of Chandra observations revealed several extragalactic globular cluster X-ray sources brighter than the Eddington limit for a neutron star (e.g. Angelini et al. 2001; Kundu et al. 2002), but a globular cluster may contain multiple bright neutron stars (as does, for example M~15 in our own galaxy -- White \\& Angelini 2001), and that the quality of X-ray spectra available from Chandra for even the brightest extragalactic sources is insufficient to make phenomenological determinations that a source has a black hole accretor.  It was pointed out that only large amplitude variability could prove that we were seeing the emission from a single source, rather than multiple sources (Kalogera, King \\& Rasio 2003).   Furthermore, the optical catalogs used to identify globular cluster counterparts to X-ray sources have been predominantly photometric catalogs, sometimes made even without color selections being used to ensure that that the optical source in question truly is a globular cluster.  Most studies done to date have focused on HST images of the central regions of elliptical galaxies with high specific frequencies of globular clusters.  In these regions, and with the angular resolution of HST, contamination will be rare, especially if color cuts are used to ensure that the contribution of background quasars is minimized.  In the halos of galaxies, the surface density of real globular clusters will drop, and contamination will be a more serious problem.  In either case, when one is looking for conclusive proof that an object is a globular cluster black hole, spectroscopic confirmation that the object is a globular cluster is essential -- the fractional contamination of the X-ray sources due to background AGN will be more serious at very high fluxes, corresponding to luminosities above $10^{39}$ ergs/sec than it will at lower levels consistent bright neutron star accretors.  Furthermore, the optical to X-ray ratios for globular cluster black holes near the Eddington limit and background quasars are quite similar. ",
        "conclusions": "We have shown for the first time clear evidence of a globular cluster black hole, on the basis of strong, highly variable X-ray emission from a source in a spectroscopically confirmed globular cluster. Based on the X-ray spectrum, the characteristic variability, and the [O III] emission in the optical spectrum, the black hole is most likely a stellar mass object accreting far faster than its Eddington rate. Given that it is difficult to develop a scenario in which a globular cluster could have both a stellar mass black hole in a binary and an intermediate mass black hole, this argues against the idea that all globular clusters contain intermediate mass black holes.  This discovery also motivates future searches for quiescent stellar mass black holes in the Milky Way's globular clusters; these may be hiding among the X-ray sources currently classified as cataclysmic variable stars.  Radio emission should be detectable only from quiescent stellar mass black holes, and should be the one feasible discriminant between the two classes of systems."
    },
    "0710/0710.4935_arXiv.txt": {
        "abstract": "Intercluster filaments negligibly contribute to the weak lensing signal in general relativity (GR), $\\gamma_{N}\\sim 10^{-4}-10^{-3}$. In the context of relativistic modified Newtonian dynamics (MOND) introduced by Bekenstein, however, a single filament inclined by $\\approx 45^\\circ$ from the line of sight can cause substantial distortion of background sources pointing towards the filament's axis ($\\kappa=\\gamma=(1-A^{-1})/2\\sim 0.01$); this is rigorous for infinitely long uniform filaments, but also qualitatively true for short filaments ($\\sim 30$Mpc), and even in regions where the projected matter density of the filament is equal to zero. Since galaxies and galaxy clusters are generally embedded in filaments or are projected on such structures, this contribution complicates the interpretation of the weak lensing shear map in the context of MOND. While our analysis is of mainly theoretical interest providing order-of-magnitude estimates only, it seems safe to conclude that when modeling systems with anomalous weak lensing signals, e.g. the ``bullet cluster\" of Clowe et al., the ``cosmic train wreck\" of Abell 520 from Mahdavi et al., and the ``dark clusters\" of Erben et al., {\\it  filamentary structures might contribute} in a significant and likely complex fashion. On the other hand, {\\it our predictions of a (conceptual) difference in the weak lensing signal could, in principle, be used to falsify MOND/TeVeS} and its variations. ",
        "introduction": "\\label{intro} Without resorting to cold dark matter (CDM), the modified Newtonian dynamics (MOND) paradigm \\citep{Mond3, mondnew} is known to reproduce galaxy scaling relations like the Tully-Fisher relation \\citep{tully}, the Faber-Jackson law \\citep{faber} and the fundamental plane \\citep{fundamental}) as well as the rotation curves of individual galaxies over five decades in mass \\citep{spiral1,mondref1,mondref2,mondref3,mondref4,mondref5,escape}. In particular, the recent kinematic analysis of tidal dwarf galaxies by \\cite{debris} is very hard to explain within the classical CDM framework while it is in accordance with MOND \\citep{tidal1,tidal2}. In addition, observations of a tight correlation between the mass profiles of baryonic matter and dark matter in relatively isolated (field) galaxies at all radii \\citep{insight2,insight} are most often interpreted as supporting MOND. Nevertheless, in rich clusters of galaxies, the MOND prescription is not enough to explain the observed discrepancy between visible and dynamical mass \\citep{neutrinos2,tevesfit,asymmetric}. At very large radii, the discrepancy is about a factor of $2$, meaning that there should be as much dark matter (mainly in the central parts) as observed baryons in MOND clusters. One solution is that neutrinos have a mass at the limit of detection, i.e. $\\sim2$ eV, which can solve the bulk of the problem of the missing mass in galaxy clusters, but other issues remain \\citep{group}. These $2$ eV neutrinos have also been invoked to fit the angular power spectrum of the cosmic microwave background (CMB) in relativistic MOND \\citep{tevesneutrinocosmo}, and are thus part of the only consistent MOND cosmology presented so far. In the following, we will refer to this model as the MOND hot dark matter ($\\mu$HDM) cosmology \\citep{tevesfit}\\footnotemark\\footnotetext[6]{Note, however, that one could also switch to sterile neutrinos with masses of a few eV \\citep[e.g.][]{sterile,maltoni1,maltoni2} and that massive (sterile) neutrinos are not indispensable within certain covariant formulations of modified gravity, e.g. the $V\\Lambda$ model, which can mimic the effects of neutrinos in clusters and cosmology as well as the behavior of a cosmological constant \\citep{vector}.}.  On the other hand, strange features have recently been discovered in galaxy clusters, which are hard to explain, such as the ``dark matter core\" devoid of galaxies at the center of the ``cosmic train wreck\" cluster Abell 520 \\citep{abell520} and others \\citep{darkcluster,bullet}. Here, we consider the possibility that this kind of features could be due to the gravitational lensing effects generated by an intercluster filament in a universe based on tensor-vector-scalar gravity \\citep[TeVeS;][]{teves}, one possible relativistic extension of MOND \\citep[cf.][]{tv1,tv2,vector}. However, we are not performing a detailed lensing analysis of any particular cluster in the presence of filaments, but rather provide a proof of concept that the influence of filaments could be much less negligible in a MONDian universe than within the framework of general relativity (GR).  Filaments are among the most prominent large-scale structures of the universe. From simulations in $\\Lambda$CDM cosmology, we know that almost every two neighboring clusters are connected by a straight filament with a length of approximately $20-30$ Mpc \\citep{LCDMfilament}. For instance, the dynamics of field galaxies, which are generally embedded in such filaments, as well as their weak lensing properties are persistently influenced by this kind of structures, generally encountering accelerations of about $0.01-0.1\\times 10^{-10}$ m s$^{-2}$. Filaments also cover a fair fraction of the sky, much larger than the covering factor of galaxy clusters. Thus, there is a good chance that filaments might be superimposed with other objects on a given line of sight, hence affecting the analysis of observational data like, for example, weak lensing shear measurements. Such recent studies prompted us to investigate the possibility that, in the context of MOND, end-on filamentary structures could be responsible for creating anomalous features in reconstructions of weak lensing convergence maps such as the peculiar ``dark matter core\" devoid of galaxies in Abell 520 \\citep{abell520}.  Short straight filaments are structures which, at the best, are partially virialized in two directions perpendicular to their axis. According to \\cite{LCDMfilament}, a filament generally corresponds to an overdensity of about $10-30$, having a cigar-like shape. Furthermore, filamentary structures tend to have a low-density gradient along their axis and, in the perpendicular directions, they have a nearly uniform core which tapers to zero at larger radii, usually about $2-5$ times their core radius. Since filaments are typically much longer than their diameter, we shall approximately treat them as infinite uniform cylinders of radius $R_{f}=2.5$ $h^{-1}$ Mpc.  Lacking a MOND/TeVeS structure formation $N$-body simulation (with or without substantially massive neutrinos), we shall adopt the naive assumption that filamentary structures have roughly the same properties in MOND and in CDM, which will be justified in \\S \\ref{app}. Deriving expressions for the TeVeS deflection angle and setting up a cosmological background, we conclude that the order of magnitude of the TeVeS lensing signal caused by filaments is compatible with that of the previously mentioned observed anomalous systems. In addition, we find that there is fundamental difference between GR and MOND/TeVeS for cylindrically symmetric lens geometries (see Fig. \\ref{fig1}); in contrast to GR, the framework of MOND/TeVeS allows us  to have image distortion and amplification effects where the projected matter density is equal to zero. As for a more realistic approach, we also consider a model where the filament has a fluctuating density profile perpendicular to its axis. Compared to the uniform model, we find that the lensing signal in this case is smaller, but still of the same order, taking into account that the filamentary structures may be inclined to the line of sight by rather small angles ($\\theta\\lesssim 20^{\\circ}$). Finally, we demonstrate the impact of filaments onto the convergence map of other objects by considering superposition of such structures with a toy cluster along the line of sight. Again, our results show an additional contribution comparable to that of a single isolated filament. \\begin{figure} \\centering \\includegraphics[trim=-20 0 0 0,width=\\linewidth]{figures/Fig1.pdf} \\caption{Light deflection by an infinitely elongated cylinder of constant mass density; the unperturbed photon traveling along the $z$-direction passes the filament at the distance $y$ (impact parameter) from the filament's axis and is deflected by the angle $\\hat\\alpha$. The line density of the filament is assumed to be constant, $\\lambda = M/L = \\rho \\pi R_f^2$, where $\\rho$ is the volume density and $R_f$ is the cylinder's radius.} \\label{fig1} \\end{figure} ",
        "conclusions": "\\label{discussion} In this work, we have analyzed the gravitational lensing effect by filamentary structures in TeVeS, a relativistic formulation of the MOND paradigm.  For this purpose, we have set up two different cosmological models in TeVeS: the so-called $\\mu$HDM cosmology including massive neutrinos on the order of $2$ eV which have already been proposed as a remedy for the discrepancies between dynamical and visible mass on cluster scales \\citep{neutrinos2} as well as for the CMB \\citep{tevesneutrinocosmo}, and a simple minimal-matter cosmology accounting for a universe which is made up of baryons alone. Encouraged by several HDM simulations and the fact that filamentary structures are generic, we have assumed that the properties of such structures, i.e. their shape and relative densities, are similar in CDM and MOND/TeVeS scenarios independent of the particularly used cosmological background.  Modeling these filaments as infinite uniform mass cylinders, we have derived analytic expressions for their lensing properties in MOND/TeVeS and Newtonian/GR gravity. Regardless of the actual used cosmological background, we have shown that TeVeS filaments can account for quite a substantial contribution to the weak lensing convergence and shear field, $\\kappa \\sim \\gamma \\sim 0.01$, as well as to the amplification bias, $A^{-1}\\sim 1.02$, which is even true outside but close $(y\\sim 2R_{f})$ to the projected ``edges\" of the filament's matter density. Exploring a simple oscillating density model of a filament and its surrounding area, we have found that the lensing signal in this case is generally smaller, but can still be of the same order, taking into account that the filamentary structures may be inclined to the line of sight by rather small angles ($\\theta\\lesssim 20^{\\circ}$). In addition, we find that there is fundamental difference between GR and MOND/TeVeS considering idealized cylindrically symmetric lens geometries: wherever the projected matter density is zero, there will be no distortion as well as no amplification effects, i.e. image and source will look exactly the same. In the context of MOND/TeVeS, however, this changes as one can have such effects in these regions. Finally, we have demonstrated the impact of filaments onto the convergence map of other objects by considering superposition with a toy cluster along the line of sight. Again, our results have shown an additional contribution comparable to that of a single isolated filament and that the contribution pattern of filaments can be generally quite complex.  Here we have considered the lensing signal generated by single filaments alone. Simulating the cosmic web in a standard $\\Lambda$CDM cosmology, \\cite{dolag} have found a shear signal $\\gamma\\sim 0.01-0.02$ along filamentary structures, which seems quite similar to what MOND/TeVeS can do. Note, however, that this signal is entirely dominated by the simulation's galaxy clusters, with the filament's signal being much smaller, approximately on the order of $10^{-4}-10^{-3}$.  Although our analysis is mainly of theoretical interest, the above result points to an interesting possibility concerning recent measurements of weak lensing shear maps. For instance, the weak shear signal in the ``dark matter peak\" of Abell 520 \\citep{abell520} is roughly at a level of $0.02$, % which is comparable to what filaments could produce in MOND/TeVeS, but not in Newtonian gravity (also cf. \\cite{wedding}).% Therefore, we conclude that filamentary structures might actually be able to cause such anomalous lensing signals within the framework of MOND/TeVeS.  In principle, the predicted difference in the weak lensing signal could also be used to test the validity of modified gravity. As several attempts to detect filaments by means of weak lensing methods have failed so far, e.g. the analysis of Abell $220$ and $223$ by \\cite{dietrich}, this might already be a first hint to possible problems within MOND/TeVeS gravity. On the other hand, shear signals around $\\gamma\\sim 0.01$ are still rather small to be certainly detected by today's weak lensing observations, and lacking $N$-body structure formation simulations in the framework of MOND/TeVeS, we cannot even be sure about how filaments form and how they look like in a MONDian universe compared to the CDM case. Clearly, more investigation is needed to gain a better understanding about the impact of filamentary structures."
    },
    "0710/0710.3809_arXiv.txt": {
        "abstract": "The bipolar morphology of the planetary nebula (PN) K 3-35 observed in radio-continuum images was modelled with 3D hydrodynamic simulations with the adaptive grid code {\\sc yguaz\\'u-a}. We find that the observed morphology of this PN can be reproduced considering a precessing jet evolving in a dense AGB circumstellar medium, given by a mass loss rate $\\dot{M}_{csm}=5\\times 10^{-5}~\\mathrm{M_{\\odot} yr^{-1}}$ and a terminal velocity $v_\\mathrm{w}=10~\\mathrm{km s^{-1}}$. Synthetic thermal radio-continuum maps were generated from numerical results for several frequencies. Comparing the maps and the total fluxes obtained from the simulations with the observational results, we find that a model of precessing dense jets, where each jet injects material into the surrounding CSM at a rate $\\dot{M}_j=2.8\\times 10^{-4}~\\mathrm{M_{\\odot}\\ yr^{-1}}$ (equivalent to a density of $8 \\times 10^{4}~\\mathrm{cm}^{-3}$), a velocity of $1\\,500~\\mathrm{km~s^{-1}}$, a precession period of 100~yr, and a semi-aperture precession angle of 20\\degr\\ agrees well with the observations. ",
        "introduction": "\\label{sec:intro} The evolution between the end of the asymptotic giant branch (AGB) and the planetary nebula (PN) phases has for a long time been a poorly understood link in the late stages of intermediate-mass stars (1 -- 8 M$_\\odot$). It is in this transition phase where the fast stellar wind of the emerging PN interacts with the slow wind from its precursor AGB star (Kwok, Purton, \\& Fitzgerald 1978), shaping the final morphology of the PN.  Other ingredients that can contribute to the final shape of a PN are the presence of interacting binary stars (e.g.  Morris 1987; Balick \\& Frank 2002; Soker \\& Bisker 2006) or magnetic fields (e.g. Garc\\'\\i a-Segura \\& L\\'opez 2000; Blackman et al. 2001; Garc\\'\\i a-Segura, L\\'opez \\& Franco 2005). Thus, the study of objects in this transition phase can give us important clues about the physical mechanisms responsible for the different morphologies observed in PNe.  K 3-35 is a very young PN with a characteristic S-shaped emission morphology that suggests the presence of precessing bipolar jets (Aaquist \\& Kwok 1989; Aaquist 1993; Miranda et al. 2000; 2001).  The detection of OH and H$_2$O maser emission as well as the presence of CO and HCO$^+$ emission suggest that K~3-35 departed from the proto-PN phase only a few decades ago (Miranda et al. 2001; Tafoya et al. 2007).  Miranda et al. (2001) estimate a dynamical age of $\\leq$ 15 years for the ionised inner core, which expands at $\\sim$25 km~s$^{-1}$. For the jets, assuming a modest jet velocity of $\\sim$100 km~s$^{-1}$, an age of $\\sim$800 years is obtained. Therefore, the jet formation in K~3-35 occurred during the proto-PN phase. In this paper we show that the radio morphology of K 3-35 can be explained by a precessing jet model. ",
        "conclusions": "Several 3D hydrodynamical simulations were carried out employing the adaptive grid code {\\sc yguaz\\'u-a}, in order to  model both the morphology and the thermal radio emission of the planetary nebula  K 3-35.  After analysing our results, we found that the bipolar structure of this PN can be described as the result of the interaction of a dense jet (with an initial number density of $8\\sim 10^4$~cm$^{-3}$ or $\\dot{M}_j=2.8\\times 10^{-4} M_{\\odot} yr^{-1}$) moving into a dense environment, previously swept up by the AGB wind of the central star.  The `S' morphology shown by  K~3-35 in 8.3 GHz radio-continuum images \\citep{lfm98,lfm01,gomez03} can be reproduced if the modelled jet precesses with a period of 100~yrs on a cone with a half-opening angle of 20\\degr.  For an integration time of 40~yrs, the simulated jet has a total length of $1.8\\times 10^{17}$~cm, which is equivalent to 2.4\\arcsec \\citep{lfm01,gomez03}considering an estimated distance of 5 kpc to K 3-35 and also the orientation of this object \\citep{uscanga07}. This time is almost 20 times smaller than the one given by \\citet{lfm01} for the jet, where a slower velocity was assumed. However, it is only $2.7$ times larger than the dynamical age of the inner core \\citep[also in][]{lfm01}, suggesting that the two events are more related than previously thought, and further supporting the idea that K 3-35 is a young object.  Synthetic radio-continuum maps were generated from our numerical results.  These maps show that the predicted morphologies and fluxes are in reasonable agreement with the observations. At an integration time of 40~yrs, the obtained spectral index is the one of optically thin emission. For an integration time of 50~yrs, the observed change of the spectral index with frequency (Aaquist 1993) is also reproduced by our simulation.  A direct comparison between the observational and numerical results is given in Fig. \\ref{fcomp}, where we show the observed and synthetic (for an integration time of 40~yrs) 8.3 GHz radio maps.  We must note that the values for the velocity and mass loss rate employed in the simulation for the jet and the CSM seem to be rather high. It is difficult for AGB and post-AGB starts to launch jets at velocities of the order of $1\\,500\\,\\mathrm{km~s}^{-1}$ (although Riera et al. 2003 have reported this kind of velocity for the outflows of the PPN 3-1475).  Besides, the mass injection is also a bit extreme, in 40 years both jets have injected $0.02~\\mathrm{M}_{\\odot}$ into the surrounding CSM. This scenario might be explained in terms of a binary system, if the jet is produced by a companion accreting material (at a rate about ten times higher than $\\dot{M}_j$) from a massive star (the AGB progenitor that produced the density distribution of the circumstellar material). The primary in the last 40 years has lost $0.2~\\mathrm{M}_{\\odot}$. This means that it started with an envelope containing this mass, which appears to correspond to an AGB star rather than a post-AGB star (even though at the present time the star has already evolved to the planetary nebula stage).  Furthermore, the CSM is very dense. With the parameters employed, in a radius of $10^{17}$~cm, a mass of 1 M$_{\\odot}$ is contained, implying a massive PN. There is observational evidence that favours a very dense surrounding CSM. However, the total mass derived from HCO$^+$ observations (Tafoya et al. 2007) is quite low ($\\sim 0.017~$M$_\\odot$), so that the molecular emission from this molecule would be confined to a small, possibly shock excited volume, in order to be consistent with our much higher mass CSM. These values imply that this scenario would be plausible if it is a short lived event. Clearly, a better determination of these parameters could be done with proper motion studies and better distance determinations.   Notwithstanding the extreme values for the  employed parameters, it is important to note that a simple model of a precessing dense jet moving into an also dense CSM, successfully reproduces the observed morphology of the PN K 3-35, obtaining a qualitatively and quantitatively good agreement between the model predictions and the observations, although this scenario would be a short lived event.   \\begin{figure} \\includegraphics[width=\\columnwidth]{fig4.eps} \\caption{Comparison between observational and numerical results. The top panel shows the 8.3~GHz radio continuum image of K 3-35 \\citep{lfm01,gomez03}.  The bottom panel shows the tilted synthetic radio emission map at the same frequency, obtained for an integration time of 40~yr.  An inclination of 36\\degr between the precession axis and the plane of the sky was considered. The simulated map was smoothed with the observed beam of $0.2\\arcsec\\times 0.2\\arcsec$ and it is shown in the same scale that the observed one.  The contours corresponds to the levels 0.18, 0.31, 0.56, 1.0, 1.78, 3.16, 5.62, and 10.00 mJy beam$^{-1}$. The arrows indicate the jet direction (projected on the plane of the sky) for integrations times of 0 and 40~yr. The last one is coincident with the P.A. measured by \\citet{lfm01}.} \\label{fcomp} \\end{figure}"
    },
    "0710/0710.2722_arXiv.txt": {
        "abstract": "We present the results of the multi-frequency observations of radio outburst of the microquasar Cyg X-3 in February and March 2006 with the Nobeyama 45-m telescope, the Nobeyama Millimeter Array, and the Yamaguchi 32-m telescope.  Since the prediction of a flare by RATAN-600, the source has been monitored from Jan 27 (UT) with these radio telescopes.   At the eighteenth day after the quench of the activity, successive flares exceeding 1 Jy were observed successfully. The time scale of the variability in the active phase is presumably shorter in higher frequency bands.  We also present the result of a follow-up VLBI observation at 8.4 GHz with the Japanese VLBI Network (JVN) 2.6 days after the first rise. The VLBI image exhibits a single core with a size of $<8$ mas (80 AU). The observed image was almost stable, although the core showed rapid variation in flux density.  No jet structure was seen at a sensitivity of $T_b = 7.5\\times 10^5$ K\\@. ",
        "introduction": "Cyg X-3 is a famous X-ray binary including a black hole candidate (e.g., \\cite{schalinski1998}). This object is classified as a microquasar due to its bipolar relativistic jet accompanied by radio flares.  Because it is located on the Galactic plane at a distance of about 10 kpc (e.g., \\cite{predehl2000}) and obscured by intervening interstellar matter, it has been observed mainly in radio and X-ray.  Its giant radio flares have been observed once every several years since its initial discovery \\citep{gregory1972,braes1972}.  The peak flux densities in the radio flares have often increased up to levels of 10 Jy or more at centimeter wave (e.g., \\cite{waltman1994}). The radio emission seems to be correlated with hard X-ray emission, and not with soft X-ray emission \\citep{mccollough1999}. Although the radio emission arises through synchrotron process of relativistic electrons in the jet \\citep{hjellming1988}, the millimeter behavior during the flares is not yet established.  An observation at a shorter wavelength and with a higher time resolution is desirable to understand the mechanism of the flares.  The quenched state of Cyg X-3, in which the radio emission is suppressed below 1 mJy, is a possible precursor of flares (e.g., \\cite{waltman1994}).  In January 2006, this quenched state was detected in monitoring observations with the RATAN-600 radio telescope \\citep{trushkin2006}.  The source has been monitored from MJD$=53762$ (Jan 27 2006 in UT) with the Nobeyama 45-m radio telescope (NRO45), the Nobeyama Millimeter Array (NMA), and the Yamaguchi 32-m radio telescope (YT32).  We detected the initial state, or  rising phase, of the radio flare of Cyg X-3 at MJD$= 53768$ (February 2 2006) and observed successive flares exceeding 1 Jy \\citep{tsuboi2006}, which turned out to be the beginning of an active phase lasting more than 40 days .   In this paper, radio observations with NRO45, NMA, YT32, and the Japanese VLBI Network (JVN) are reported.  The observation procedures are summarized in section 2. The light curves and the spectral evolution observed with NRO45, NMA and YT32 are shown in section 3, together with the result of JVN\\@.  Detailed discussion based on these observations will be published as separate papers. ",
        "conclusions": ""
    },
    "0710/0710.0727_arXiv.txt": {
        "abstract": "{ Gamma-ray burst are thought to be produced by highly relativistic outflows. Although upper and lower limits for the outflow initial Lorentz factor $\\Gamma_0$ are available, observational efforts to derive a direct determination of $\\Gamma_0$ have so far failed or provided ambiguous results. As a matter of fact, the shape of the early-time afterglow light curve is strongly sensitive on $\\Gamma_0$ which determines the time of the afterglow peak, i.e. when the outflow and the shocked circumburst material share a comparable amount of energy. We now comment early-time observations of the near-infrared afterglows of GRB\\,060418 and GRB\\,060607A performed by the REM robotic telescope. For both events, the afterglow peak was singled out and allowed us to determine the initial fireball Lorentz, $\\Gamma_0\\sim400$. ",
        "introduction": "The early stages of gamma-ray burst (GRB) afterglow light curves display a rich variety of phenomena at all wavelengths and contain significant information which may allow determining the physical properties of the emitting fireball. The launch of the \\textit{Swift} satellite \\citep{Neil04}, combined with the development of fast-slewing ground-based telescopes, has hugely improved the sampling of early GRB afterglow light curves.  Since many processes work in the early afterglow, it is often difficult to model them well enough to be able to determine the fireball characteristics. The simplest case is a light curve shaped by the forward shock only. This case is particularly interesting because, while the late-time light curve is independent of the initial conditions (the so-called self-similar solution), the time at which the afterglow peaks depends on the original fireball Lorentz factor $\\Gamma$, thus allowing a direct measurement of this fundamental parameter \\citep{SP99}. The short variability timescales, coupled with the nonthermal GRB spectra, indeed imply that the sources emitting GRBs have a highly relativistic motion \\citep{Ruderman75,Fenimore93,Piran00,Lith01}, to avoid suppression of the high-energy photons due to pair production. This argument, however, can only set a lower limit to the fireball Lorentz factor. Late-time measurements (weeks to months after the GRB) have shown $\\Gamma \\sim \\mbox{a few}$ \\citep{Frail97,Taylor05}, but a direct measure of the initial value (when $\\Gamma$ is expected to be $\\sim 100$ or more) is still lacking.  We present here the NIR early light curves of the GRB\\,060418 and GRB\\,060607A afterglows observed with the REM robotic telescope\\footnote{\\url{http://www.rem.inaf.it}} \\citep{Zerb01,Chinc03} located in La Silla (Chile). These light curves show the onset of the afterglow and its decay at NIR wavelengths as simply predicted by the fireball forward shock model, without the presence of flares or other peculiar features. A detailed discussion of these data has also been reported by \\citet{Mol07} and \\citet{Mal07}.   \\begin{figure} \\centering \\includegraphics[width=11cm]{CovinoS_2007_01_fig01.ps} \\\\ \\centering\\includegraphics[width=11cm]{CovinoS_2007_01_fig02.ps} \\caption{NIR and X-ray light curves of GRB\\,060418 (upper panel) and GRB\\,060607A (lower panel). The dotted lines show the models of the NIR data using the smoothly broken power law (see Sect. \\ref{modelling}). For GRB\\,060418 the dashed line shows the best-fit to the X-ray data.\\label{fig:lcs}} \\end{figure} ",
        "conclusions": "The REM discovery of the afterglow onset has demonstrated once again the richness and variety of physical processes occurring in the early afterglow stages. The very fast response observations presented here provide crucial information on the GRB fireball parameters, most importantly its initial Lorentz factor. This is the first time that $\\Gamma(t_{\\rm peak})$ is directly measured from the observations of a GRB. The measured $\\Gamma_0$ value is well within the range ($50 \\la \\Gamma_{0} \\la 1000$) envisaged by the standard fireball model \\citep{Piran00,Guetta01,Alicia02,Meszaros}. It is also in agreement with existing measured lower limits \\citep{Lith01,zhang06}.  Using $\\Gamma_0=400$ we can also derive other fundamental quantities characterising the fireball of the two bursts. In particular, the isotropic-equivalent baryonic load of the fireball is $M_{\\rm fb} = E/(\\Gamma_0 c^2) \\approx 7\\times 10^{-4}\\,M_\\odot$, and the deceleration radius is $R_{\\rm dec} \\approx 2 c t_{\\rm peak} [\\Gamma(t_{\\rm peak})]^2/(1+z) \\approx 10^{17}$~cm. This is much larger than the scale of $\\sim 10^{15}$~cm where the internal shocks are believed to power the prompt emission \\citep{MesRees97,Rees94}, thus providing further evidence for a different origin of the prompt and afterglow stages of the GRB."
    },
    "0710/0710.0382_arXiv.txt": {
        "abstract": "% We present results from a 150 ksec {\\it Suzaku} observation of the Seyfert 1.5 NGC 3516 in October 2005. The source was in a relatively highly absorbed state. Our best-fit model is consistent with the presence of a low-ionization absorber which has a column density near 5 $\\times$ 10$^{22}$ cm$^{-2}$ and covers most of the X-ray continuum source (covering fraction 96--100$\\%$). A high-ionization absorbing component, which yields a narrow absorption feature consistent with Fe K {\\sc XXVI}, is confirmed. A relativistically broadened Fe K$\\alpha$ line is required in all fits, even after the complex absorption is taken into account; an additional partial-covering component is an inadequate substitute for the continuum curvature associated with the broad Fe line. A narrow Fe K$\\alpha$ emission line has a velocity width consistent with the Broad Line Region. The low-ionization absorber may be responsible for producing the narrow Fe K$\\alpha$ line, though a contribution from additional material out of the line of sight is possible. We include in our model soft band emission lines from He- and H-like ions of N, O, Ne and Mg, consistent with photo-ionization, though a small contribution from collisionally-ionized emission is possible. ",
        "introduction": "The 6.4 keV Fe K$\\alpha$ emission line has long been known to be an important diagnostic of the material accreting onto supermassive black holes. The Compton reflection hump, frequently seen in Seyfert spectra above $\\sim$7~keV and peaking near 20--30 keV (Pounds et al.\\ 1990), indicates that Seyferts' Fe K lines may have an origin in optically thick material, such as the accretion disk. Observations with {\\it ASCA} indicated that many Fe K$\\alpha$ lines were broad (FWHM velocities up to $\\sim$0.3$c$) and asymmetrically skewed towards lower energies, implying an origin in the inner accretion disk; the line profile is sculpted by gravitational redshifting and relativistic Doppler effects (e.g., Tanaka et al.\\ 1995, Fabian et al.\\ 2002). However, {\\it XMM-Newton} and {\\it Chandra} observations have been revealing a more complex picture. A narrow Fe K component (FWHM velocities $\\sim$5000 km s$^{-1}$ or less) appears to be quite common; these lines' widths suggest emission from distant material, such as the outer accretion disk, the optical/UV Broad Line Region (BLR) or the molecular torus. Spectral observations in which the broad and narrow components are deconvolved are thus a prerequisite for using the Fe K line as a tracer of the geometry of the emitting gas.  At the same time, there is strong evidence from X-ray and UV grating observations for the presence of ionized material in the inner regions of a large fraction of AGN (e.g., Blustin et al.\\ 2005; McKernan et al.\\ 2007). High-resolution spectroscopy shows the gas is usually outflowing from the nucleus; typical velocities are $\\sim$ a few hundred km s$^{-1}$. Absorption due to a broad range of ionic species is commonly seen; and for many sources, there is evidence for several different photo-ionized absorbing components, as opposed to a single absorber, along the line of sight. In the Fe K bandpass, some Seyferts show evidence for absorption by H- or He-like Fe, indicating a zone of highly-ionized absorbing material (e.g., NGC 3783, Reeves et al.\\ 2004).  Cold absorbing gas, with line of sight columns in excess of the Galactic value, are routinely observed in Seyfert 2 AGN in accordance with unification schemes (Urry, Padovani 1995), and have also been reported in some Seyfert 1 AGN. Importantly, variations in column density on timescales from hours to years have been observed in both Seyfert 1 AGN (e.g., in I Zw 1, Gallo et al.\\ 2007; see also Lamer et al.\\ 2003, Puccetti et al.\\ 2004) and Seyfert 2 AGN (Risaliti et al.\\ 2002; Risaliti et al.\\ 2005), suggesting that the absorbing circumnuclear material is not homogeneous in either Seyfert type, has a high tranverse velocity and occurs over a range of length scales.  NGC 3516 is a well-studied, nearby (z = 0.008836; Keel 1996) Seyfert 1.5 AGN that can display strong 2--10 keV flux variability on timescales of hours to years (e.g., Markowitz, Edelson 2004). Previous X-ray spectroscopic observations of NGC 3516 e.g., Nandra et al.\\ (1997) using {\\it ASCA}, have indicated the presence of a broad Fe K line, and this source is known to also contain complex and ionized absorption. Numerous UV absorption lines, including N {\\sc V}, C {\\sc IV} and Si {\\sc IV}, were observed with the {\\it International Ultraviolet Explorer} (Ulrich, Boisson 1983); absorption line strengths vary on timescales as short as weeks as the absorber responds to variations in the ionizing flux (e.g., Voit et al.\\ 1987). {\\it Hubble Space Telescope} observations have revealed that this component of absorbing gas (henceforth called the ``UV absorber'') may consist of several distinct kinematic components (Crenshaw et al.\\ 1998).  X-ray spectra of NGC 3516 can exhibit evidence for large columns ($\\gtrsim$10$^{22}$ cm$^{-2}$) of absorbing gas (e.g., Kolman et al.\\ 1993), and the X-ray absorbers can display variability on timescales of years (e.g., Mathur et al.\\ 1997). Using {\\it Chandra} gratings data from observations in April 2001 and November 2001, Turner et al.\\ (2005) observed K-shell absorption lines due to H- like Mg, Si and S, and He-like Si, evidence for a highly-ionized absorber, likely with column density $\\gtrsim$10$^{22}$ cm$^{-2}$, outflowing at $\\sim$1100 km s$^{-1}$. Simultaneous with these {\\it Chandra} observations in 2001 were two {\\it XMM-Newton} observations. Turner et al.\\ (2005) modeled the continuum curvature of the two {\\it XMM-Newton} EPIC spectra by including a partial covering, mildly-ionized absorber; the column density was $\\sim$2.5 $\\times$ 10$^{23}$ cm$^{-2}$, with a covering fraction of $\\sim$50$\\%$. However, the formal requirement for the broad Fe line was reduced, leading to uncertainty as to whether the broad Fe line really existed in NGC 3516. Spectral fitting using an instrument with a wide bandpass is thus necessary to remove such model degeneracies.  In this paper, we report on an observation of NGC 3516 made with the {\\it Suzaku} observatory in October 2005. The combination of the X-ray Imaging Spectrometer (XIS) CCD and the Hard X-ray Detector (HXD) instruments have yielded a broadband spectrum covering 0.3 to 76 keV, allowing us to deconvolve the various broadband emitting and absorbing components. Furthermore, the exceptional response of the XIS CCD and high signal-to-noise ratio of this observation have allowed us to study narrow emission lines in great detail. $\\S$2 gives a brief overview of the {\\it Suzaku} observatory, and describes the observation and data reduction. $\\S$3 describes fits to the time-averaged spectrum. Variability analysis is briefly discussed in $\\S$4. Flux-resolved spectral fits are discussed in $\\S$5. In $\\S$6, we describe a search for narrow red- or blue-shifted lines in the Fe K bandpass. The results are discussed in $\\S$7, and a brief summary is given in $\\S$8. ",
        "conclusions": "During the late 1990's, NGC 3516 typically displayed a 2--10 keV flux of $\\sim$4--6 $\\times$ 10$^{-11}$ erg cm$^{-2}$ s$^{-1}$ (e.g., Markowitz, Edelson 2004). During the 2001 {\\it XMM-Newton}/{\\it Chandra} observations, however, the observed 2--10 keV flux was much lower: 1.6--2.3 $\\times$ 10$^{-11}$ erg cm$^{-2}$ s$^{-1}$ (Turner et al.\\ 2005). Table 4 lists the inferred absorption-corrected 2--10 keV nuclear fluxes from the {\\it XMM-Newton} observations, as well as during the 2005 {\\it Suzaku} observation. In addition, Figure 10 shows the unfolded observed spectra for the {\\it Suzaku} XIS and the 2001 {\\it XMM-Newton} EPIC-pn data (see Turner et al.\\ 2005 for details regarding the {\\it XMM-Newton} data). The {\\it Suzaku} observation apparently caught the source in a similar low level of nuclear flux as the 2001 observations.  The observed 0.5--2.0 keV flux during the {\\it Suzaku} observation, however, was $\\sim$2--3 times lower than during the {\\it XMM-Newton} observations, indicating that the source was still heavily obscured, and confirming that the complex absorption in this source cannot be ignored when fitting the broadband spectrum and modeling diskline components.   \\subsection{Power-law Components}   The primary power-law observed in the hard X-rays is likely emission from a hot corona very close to the supermassive black hole, as seen in all Seyferts. The nature of the soft power-law component, however, is not as clear. It could represent nuclear emission scattered off optically-thin material, e.g., in the optical/UV Narrow Line Region (NLR). In the baseline model, the normalization of the soft power-law relative to that of the primary power-law was 4.2$\\pm$0.4$\\%$. Assuming a covering fraction of unity, this ratio is equal to the optical depth of the scattering material, indicating a column density of roughly 5 $\\times$ 10$^{22}$ cm$^{-2}$, consistent with this notion, though the column density is somewhat too high to likely be associated with the NLR. It is interesting to note that this column density is similar to that obtained for the high-ionization absorber, suggesting the possibility that this absorbing component could be associated with a zone of scattering. Using {\\it Chandra}-ACIS, George et al.\\ (2002) found the extended circumnuclear gas to have a 0.4--2.0 keV flux of roughly 10$^{-14}$ erg cm$^{-2}$ s$^{-1}$. However, that emission was studied over an annular extraction region 3$\\arcsec$ to 10$\\arcsec$ (0.6--1.8 kpc), and so that flux value is likely a lower limit to the 0.4--2.0 keV flux that Suzaku would observe (given the XRTs' 2$\\arcmin$ HPD). In our baseline model, we found an unabsorbed 0.4--2.0 keV flux of 1 $\\times$ 10$^{-12}$ erg cm$^{-2}$ s$^{-1}$, consistent with the notion that the soft power-law is scattered emission.  In this case, the decrease in observed 0.5--2.0 keV flux from 2001 to 2005 could potentially be explained by the scattered emission responding to a recent decrease in nuclear continuum flux. However, this scenario would require the bulk of the scattered emission to lie within at most a few light years of the black hole, and the nuclear flux would have had to decrease between 2001 and 2005 (when the source was not observed by any major X-ray mission in the 2--10 keV band) then return to 2001 levels by the {\\it Suzaku} observation.    Alternatively, the soft power-law could be unobscured, ``leaked'' nuclear emission as part of a partial covering scenario. In this case, the primary absorber would obscure 96$\\%$ of the sky as seen from the nuclear continuum source. The lack of significant variability in the 0.3--1.0 keV band could argue for the soft power-law to originate in scattered emission, since we might expect to observed variability of the same amplitude as the 2--10 keV band only if the soft power-law were leaked nuclear emission. However, this is far from certain, as the 0.3--1.0 keV band had a low count rate and the presence of the soft emission lines in the XIS spectrum could contribute to dilution of intrinsic variability in the soft power-law. A broadband observation spanning a larger observed flux range is needed to clarify this issue. The soft power-law could of course represent a blend of scattered emission plus leaked nuclear emission. We therefore conclude that the primary absorber has a covering fraction between 96--100$\\%$.    \\subsection{Complex Absorption}  We detect two zones of absorption: in addition to the primary absorber, which has a covering fraction of 96--100$\\%$, there is the high-ionization absorber, which is assumed here to have a covering fraction of unity. The high-ionization absorber is potentially the same as that reported by Turner et al.\\ (2005); we find a column density $N_{\\rm H}$ of 4.0$^{+4.6}_{-3.1}$ $\\times$ 10$^{22}$ cm$^{-2}$, consistent with the column density of 2 $\\times$ 10$^{22}$ cm$^{-2}$ used by Turner et al.\\ (2005), although we use a slightly higher ionization parameter (see Turner et al.\\ 2007). Previous studies of NGC 3516, such as Netzer et al.\\ (2002), have discussed in detail the UV absorber, responsible for H Ly$\\alpha$, C {\\sc IV} and N {\\sc V} absorption features in {\\it Hubble Space Telescope} spectra (Kraemer et al.\\ 2002). In the X-ray band, discrete features associated with Mg {\\sc VII--IX} and Si {\\sc VII--IX} are expected from this component, but with the CCD resolution and with calibration-related artifacts near 1.7--1.8 keV in the XIS, such features are not detected by {\\it Suzaku}.  {\\it Suzaku} has found the primary absorber of the hard X-ray continuum to be lowly-ionized (log($\\xi$) = 0.3$\\pm$0.1 erg cm s$^{-1}$), with a column density $N_{\\rm H}$ of 5.5$\\pm$0.2 $\\times$ 10$^{22}$ cm$^{-2}$. It is possible that it is the same absorber that Turner et al.\\ (2005) designated as the ``heavy'' partial-covering absorber, though we use a somewhat lower ionization parameter (see Turner et al.\\ 2007). A new observation of NGC 3516 with {\\it XMM-Newton} in 2006 October showed that the source had returned to similar $>$6 keV brightness and similar obscuration levels ($N_{\\rm H}$ $\\sim$ 2 $\\times$ 10$^{23}$ cm$^{-2}$; covering fraction $\\sim$ 45$\\%$) as during the 2001 observations (Turner et al.\\ 2007). Long-term changes in the covering fraction of the heavy absorber could explain the bulk of the spectral variability changes between the 2001, 2005 and 2006 observations. In this case, the column density has decreased by a factor of 4.5, while the covering fraction has increased from $\\sim$40--60$\\%$ to 96--100$\\%$, from 2001 to 2005, and subsequently returned to approximately the same levels as in 2001 within the next 12 months. However, because we have not actually observed the entire eclipse associated with a specific, discrete blob of absorbing gas traversing the line of sight, it is not clear whether the covering fractions derived are associated with single, large blobs partially blocking the line of sight to the X-ray continuum source, or if the absorber consists of numerous, discrete blobs or has a filamentary or patchy structure. On the other hand, given the gaps between the 2001 and 2006 {\\it XMM-Newton} and 2005 {\\it Suzaku} observations, it is certainly plausible that these observations could have caught independent, discrete blobs or filaments with differing column densities and differing physical sizes and/or radial distances lying on the line of sight.  To estimate the distance $r$ between the central black hole and the absorbing gas, we can use a definition of the ionization parameter $\\xi$ = $L_{1-1000 {\\rm Ryd}}$/($n$$r^2$), where $n$ is the number density. $L_{1-1000 {\\rm Ryd}}$ is the 1--1000 Ryd illuminating continuum luminosity, and the value of the ionization parameter is taken to be the value of 2 erg cm s$^{-1}$ measured above. We estimate the maximum possible distance to the material by assuming that the thickness $\\Delta$$r$ must be less than the distance $r$. The column density $N_{\\rm H}$ = $n$$\\Delta$$r$, yielding the upper limit $r$ $<$ $L_{1-1000 {\\rm Ryd}}$/($N_{\\rm H}$$\\xi$). We estimate the 1-1000 Ryd flux from the baseline model to be 9.9 $\\times$ 10$^{-11}$ erg cm$^{-2}$ s$^{-1}$,  which corresponds to  $L_{1-1000 {\\rm Ryd}} = 1.7 \\times 10^{43}$ erg s$^{-1}$ (assuming $H_{\\rm o}$ = 70 km s$^{-1}$ Mpc$^{-1}$ and $\\Lambda_{\\rm o}$ = 0.73). $r$ is thus $<$2 $\\times$ 10$^{20}$ cm  (180 light-years), a very loose upper limit encompassing both distances associated with the BLR ($\\sim$10 light-days; Peterson et al.\\ 2004) and a possible cold molecular torus at a 1 pc radius. Variability in the absorber properties between 2001--2005 and 2005--2006 thus imply a radial distance of at most a few light years in the case of clouds traversing the line of sight to the nucleus. In addition, in the case of a partial covering scenario, it is plausible that the absorber's size could be of the same order as that of the X-ray continuum source. The absorbing material could thus be e.g., associated with the base of an outflow or from dense clouds associated with magnetohydrodynamic disk turbulence (e.g., Emmering et al.\\ 1992).    NGC 3516's transition from an unobscured source to a moderately-obscured source in a 4 year span presents a challenge to standard Seyfert 1/2 unification schemes. If the obscuration in NGC 3516 is associated with an equatorial molecular torus usually invoked in Seyfert 1/2 unification schemes, then it is possible that during the {\\it Suzaku} observation, the inner edge of the torus could have intersected the line of sight, but given NGC 3516's classification as a Seyfert 1, this could only occur if the torus opening angle were extremely small and the torus were not azimuthally symmetric. Alternatively, variations in column density and/or covering fraction could be due to fine structure in large-distance (tens of pc), non-equatorial filaments that traverse the line of sight (e.g., Malkan et al.\\ 1998).    \\subsection{Fe K Emission Components and Compton Reflection}  We have deconvolved the broadband emitting components, and determined that 1) the existence of the broad Fe line is robust in that it was required in all models for an adequate fit, and 2) a partial covering component could not mimic the curvature associated with a relativistic broad line. We note, for instance, that if we remove the diskline components from the baseline model and refit, not only is the fit worse ($\\chi^2$ increases by over 170), but the value of $R$ becomes $\\sim$ 3.2.  This value is incompatible with the observed $EW$ of the narrow line unless the Fe abundance is extremely sub-solar ($\\lesssim$0.3; see below). The best-fit disk inclination was typically $\\lesssim$25$\\arcdeg$. The inner radius was typically $\\lesssim$5$R_{\\rm g}$. The line energy was seen to be consistent with neutral to mildly-ionized Fe (up to Fe $\\sim$ {\\sc XX}; Kallman et al.\\ 2004). The equivalent width with respect the primary continuum was 287$^{+49}_{-24}$ eV, consistent with the value of 431$^{+193}_{-172}$ eV obtained by Turner et al.\\ (2005) for the April 2001 {\\it XMM-Newton} observation, where the spectrum could be fit with a diskline component in addition to the complex absorbing components.     The line energy of the narrow Fe K$\\alpha$ line was also consistent with emission from neutral Fe. The intensity of the narrow line during the {\\it Suzaku} observation is roughly 40$\\%$ higher than that measured during the 2001 {\\it XMM-Newton} and {\\it Chandra} observations (Turner et al.\\ 2002), possibly indicating that a substantial fraction of the Fe K line photons originate in a region $\\lesssim$5 lt.-yr.\\ in size. We measured a FWHM velocity line width for the narrow Fe K$\\alpha$ line of $<$ 4900 km s$^{-1}$ (99$\\%$ confidence level for two interesting parameters). This velocity does not rule out an origin in the BLR; Peterson et al.\\ (2004) reported FWHM velocities for the H$\\alpha$ and H$\\beta$ lines of 4770$\\pm$893 and 3353$\\pm$310 km s$^{-1}$, respectively.  However, we also cannot exclude a contribution from an origin in the putative molecular torus; there could potentially be a very narrow line component with FWHM velocity $\\sim$ a few hundred km s$^{-1}$, but the XIS would be unable to separate it from the relatively broader line component.   It is possible that the same material that absorbs the hard X-rays along the line of sight is responsible for producing the narrow Fe line. The material producing the Fe line cannot have a column substantially less than 10$^{\\sim 22}$ cm$^{-2}$ or else there would be insufficient optical depth to produce a prominent Fe K line. The primary absorber, with its column density of 5.5 $\\times$ 10$^{22}$ cm$^{-2}$ and low ionization state, is thus a plausible candidate for the narrow line origin. As an estimate of the Fe K$\\alpha$ equivalent width expected in this case, we can assume an origin in optically-thin gas which completely surrounds a single X-ray continuum source and is uniform in column density, and use the following equation: \\begin{equation} EW_{\\rm calc} = f_{\\rm c} \\omega f_{\\rm K\\alpha} A \\frac{\\int^{\\infty}_{E_{\\rm K edge}}P(E) \\sigma_{\\rm ph}(E) N_{\\rm H} dE}{P(E_{\\rm line})} \\end{equation}  Emission is assumed to be isotropic. Here, $f_{\\rm c}$ is the covering fraction, initially assumed to be 1.0. $\\omega$ is the fluorescent yield, 0.34 (Kallman et al.\\ 2004). $f_{\\rm K\\alpha}$ is the fraction of photons that go into the K$\\alpha$ line as opposed to the K$\\beta$ line; this is 0.89 for Fe {\\sc I}. $A$ is the number abundance relative to hydrogen. We assumed solar abundances, using Lodders (2003). $P(E)$ is the spectrum of the illuminating continuum at energy $E$; $E_{\\rm line}$ is the K$\\alpha$ emission line energy. $\\sigma_{\\rm ph}(E)$ is the photo-ionization cross section assuming absorption by K-shell electrons only (Veigele 1973\\footnote{http://www.pa.uky.edu/$\\sim$verner/photo.html}).  For $N_{\\rm H}$ = 5.5 $\\times$ 10$^{22}$ cm$^{-2}$, $EW_{\\rm calc}$ = 29 eV, substantially lower than the observed $EW$ of 123$\\pm$7 eV. We conclude that it is possible for the primary absorber to contribute to the observed line $EW$, but there is also likely a contribution from some other (non-continuum absorbing) material lying out of the line of sight, likely with column densities 10$^{\\sim 23}$ cm$^{-2}$ (e.g., Matt et al.\\ 2002). For instance, if the putative cold molecular torus does not intersect the line of sight, it could contribute to the observed $EW$.  The 13$\\%$ upper limit to ratio of the Compton shoulder/ narrow Fe K$\\alpha$ core intensity was a 90$\\%$ confidence limit only, and does not exclude at high confidence the possibility of Compton-thick material out of the line of sight.    An additional possibility is that the material emitting the bulk of the line photons could be responding to a continuum flux that was higher in the past. For instance, if the putative molecular torus is located $\\sim$ a pc or so from the black hole, the torus will yield a line $EW$ corresponding to the continuum flux averaged over the past few years. This situation is plausible for NGC 3516, as the 2--10 keV flux of NGC 3516 during $\\sim$1998--2001 (Markowitz, Edelson 2004) was a factor of $\\sim$1.5--2 times brighter than during 2005.   We now turn our attention to properties of the Compton reflection continuum. {\\it Suzaku} has observed other Seyferts to display reduced levels of variability in the PIN band compared to the 2--10 keV band, e.g., in MCG--6-30-15 (Miniutti et al.\\ 2007). This behavior is thought to be caused by the presence of the relatively non-varying Compton reflection hump, which dilutes the observed $>$10 keV variability of the power-law component. Gravitational light-bending in the region of strong gravity has been invoked to explain the relative constancy of the reflection spectrum (Compton hump and Fe K diskline component) despite large variations in the coronal power-law flux in MCG--6-30-15, for instance (Miniutti et al.\\ 2007). In the case of NGC 3516, the observed fractional variability amplitudes for the 2--10 and 12--76 keV bands were $F_{\\rm var,2-10}$  = 9.2$\\pm$0.3$\\%$ and $F_{\\rm var,12-76}$ $<$ 4.4$\\%$, respectively. These measurements allow us to rule out the possibility that the Compton hump varies in concert with the power-law, since the variability amplitudes would be consistent in that case. The primary power-law and Compton hump contribute 44$\\%$ and 56$\\%$, respectively, of the total 12--76 keV flux. In the case of a constant Compton hump and variable power-law, $F_{\\rm var,12-76}$ would then be equal to $F_{\\rm var,2-10}$ / 2.25, or roughly 4.1$\\%$. This is roughly consistent with the observed upper limit on $F_{\\rm var,12-76}$, suggesting that the reflection component varies less strongly than the primary power-law over the course of the observation. To verify this, however, we would need to observe NGC 3516 over a larger % X-ray flux range than in the current {\\it Suzaku} observation to potentially observe any significant variability in the PIN band.   Finally, we discuss the origin of the material that gives rise to the observed Compton reflection hump. The primary and high-ionization absorbers lack the necessary column density, and are excluded. We next consider an origin in the same material that yields either the broad or narrow Fe lines. George, Fabian (1991) calculated that $R$ = 1 corresponds to an observed Fe K$\\alpha$ line $EW$ (relative to a primary continuum with a photon index of 1.9) of 140 eV for neutral Fe, assuming an inclination angle of 25$\\arcdeg$. However, George, Fabian (1991) used the elemental abundances of Morrison, McCammon (1983), where the Fe number abundance relative to hydrogen was $A_{\\rm Fe}$ = 3.3 $\\times$ 10$^{-5}$. More recent papers have slightly lower values of $A_{\\rm Fe}$, 2.7--3.0 $\\times$ 10$^{-5}$ (Lodders 2003; Wilms et al.\\ 2000). The expected equivalent width corresponding to $R$ = 1 is thus 115--125 eV. In our baseline model, we found a best-fit value of $R$ = 1.7, which corresponds to an expected line $EW$ (relative to the primary continuum) of 200--215 eV, a value in between the observed $EW$s of the broad line (287 eV in the baseline model) and the narrow line (123 eV). It is thus not clear from this measurement alone whether the total Compton reflection continuum is associated with the broad line (disk), narrow line (a distant origin), or both. That is, while is it a possibility that at least some portion of the Compton reflection component is associated with the broad Fe K component, we cannot exclude the possibility that the narrow line contributes as well and that there is reflection off cold, distant material. For example, in $\\S$3.3, we demonstrated that a model wherein there existed both blurred reflection from an ionized disk plus reflection from cold, distant material, such as the molecular torus, provided a good fit to the data. In addition, we demonstrated in this section that the observed $EW$ of the narrow Fe K line means we cannot rule out a contribution to the narrow Fe line, and to the reflection continuum as well, from Compton-thick material out of the line of sight."
    },
    "0710/0710.2343.txt": {
        "abstract": "This paper analyses the radio properties of a subsample of optically obscured ($R\\ge 25.5$) galaxies observed at 24$\\mu$m by the {\\it Spitzer Space Telescope} within the First Look Survey. 96 $F_{24\\mu\\rm m}\\ge 0.35$~mJy  objects out of 510 are found to have a radio counterpart at 1.4~GHz, 610~MHz or at both frequencies respectively down to $\\sim 40\\mu$Jy and $\\sim 200\\mu$Jy. IRAC photometry sets the majority of them in the redshift interval $z\\simeq [1-3]$ and allows for a broad distinction between AGN-dominated galaxies ($\\sim 47$\\% of the radio-identified sample) and systems powered by intense star-formation ($\\sim 13$\\%), the remaining objects being impossible to classify. The percentage of radio identifications is a strong function of 24$\\mu$m flux: almost all sources brighter than $F_{24\\mu\\rm m}\\sim 2$~mJy are endowed with a radio flux at both 1.4~GHz and 610~MHz, while this fraction drastically decreases by lowering the 24$\\mu$m flux level. The radio number counts at both radio frequencies suggest that the physical process(es) responsible for radio activity in these objects have a common origin regardless of whether the source shows mid-IR emission compatible with being an obscured AGN or a star-forming galaxy. We also find that both candidate AGN and star-forming systems follow (although with a large scatter) the relationship between 1.4~GHz and 24$\\mu$m fluxes reported by Appleton et al. (2004) which identifies sources undergoing intense star formation activity. However, a more scattered relation is observed between 24$\\mu$m and 610~MHz fluxes. On the other hand, the inferred radio spectral indices $\\alpha$ indicate that a large fraction of objects in our sample ($\\sim 60$\\% of all galaxies with estimated $\\alpha$) may belong to the population of Ultra Steep Spectrum (USS) Sources, typically 'frustrated' radio-loud AGN. We interpret our findings as a strong indication for concurrent AGN and star-forming activity, whereby the 1.4~GHz flux is of thermal origin, while that at 610~GHz mainly stems from the nuclear source. ",
        "introduction": "The advent of the {\\it Spitzer Space Telescope} has marked a fundamental milestone in our understanding of the assembly history of massive spheroidal galaxies, one of the major issues for galaxy formation models. The unprecedented sensitivity of the Multiband Imaging Photometer for {\\it Spitzer} (MIPS) at 24$\\mu$m has in fact for the first time allowed the detection at high redshifts of a population of Luminous and UltraLuminous Infrared Galaxies (LIRGs; ULIRGs) with huge infrared luminosities ($L_{\\rm IR}>10^{11}L_{\\odot}$). Such sources are underluminous at rest frame optical and UV wavelengths because they are reprocessing and radiating much of their energy in the IR (e.g. Sanders \\& Mirabel 1996). As a consequence, LIRGs and ULIRGs at high redshifts had been missed so far either due to their extreme optical faintness or because previous infrared missions such as the InfraRed Astronomical Satellite (IRAS) or the Infrared Space Observatory (ISO) did not have enough sensitivity to push the observations beyond $z \\sim 1$.  Recent studies have shown that these objects, while relatively rare in the local universe (e.g. Sanders \\& Mirabel 1996), become an increasingly significant population at higher redshifts (e.g. Le Floc'h et al. 2004; 2005; Lonsdale et al. 2004; Caputi et al. 2006; 2007) and likely dominate the luminosity density at $z>1$ (see e.g. Dole et al. 2006). Their space density, found to range between $10^{-4}$ and a few $10^{-5}$~Mpc$^{-3}$ according to the selection criteria adopted by different studies (see e.g. Caputi et al 2007; Daddi et al. 2007a, Magliocchetti et al. 2007a), is a factor of 10 to 100 higher than that of optically selected quasars in the same redshift range (e.g. Porciani, Magliocchetti \\& Norberg 2004). Furthermore, clustering studies (e.g. Magliocchetti et al. 2007; 2007a; Farrah et al. 2006) prove that, at variance with their local counterparts, LIRGs and ULIRGs at $z\\sim 2$ are associated with extremely massive ($M\\simgt 10^{13} M_{\\odot}$, where $M$ here refers to the dark matter) structures, only second to those which locally host very rich clusters of galaxies. Given their properties, it then appears clear that these sources represent a fundamental phase in the build up of massive galactic bulges, and in the growth of their supermassive black holes.  Emission line diagnostics for bright ($F_{24\\mu \\rm m}\\simgt 1$~mJy) mid-IR samples of LIRGs and ULIRGs at $z\\sim 2$ in the near and mid-IR spectral regimes (see e.g. Yan et al. 2005; 2007; Weedman et al. 2006; 2006a; Brand et al. 2007; Martinez-Sansigre et al. 2006a) have shown these sources to be a mixture of obscured type1-type2 AGN and systems undergoing intense star formation activity. These findings are confirmed by photometric follow up mainly undertaken in the mid-IR and X-ray (both soft and hard) bands which also prove that the fraction of galaxies dominated by a contribution of AGN origin is drastically reduced at faint mid-IR fluxes (e.g. Brand et al. 2006; Weedman et al. 2006;  Treister et al. 2006; Magliocchetti et al. 2007). Unfortunately, the exact proportion of AGN vs starforming dominated galaxies is still undetermined. This separation is further complicated by the existence of a noticeable number of mixed systems where both star formation and AGN activity significantly contribute to the IR emission. For instance, Daddi et al. (2007a) find that about 20\\% of 24$\\mu$m-selected galaxies in the GOODS sample show a mid-IR excess which is not possible to reconcile with pure star-forming activity. Such a fraction increases to $\\sim 50-60$\\% at the highest (stellar) masses probed by their study.   Clearly, understanding how the mid-IR sources divide between starbursts, AGN and composite systems is now the next essential step in order understand the relationships amongst the formation and evolution of stars, galaxies and massive black holes powering AGNs within dusty environments and more generally within massive systems observed at the peak of their activity.  This paper approaches the study of the population of optically faint luminous infrared galaxies from the point of view of their multifrequency radio emission. Diagnostics based on the radio signal steming from these sources can in fact provide precious information on the process(es) which are actively taking place within such systems. Enhanced radio activity can stem from supernova remnants associated with regions which are vigorously forming stars, or originate from nuclear activity (AGN-dominated sources). These two processes determine a rather different spectral behaviour at radio wavelengths: star-forming systems are in general characterized by radio spectra which feature power-law shapes with slope (hereafter called {\\it radio spectral index} $\\alpha$, defined as $F\\propto \\nu^{-\\alpha}$, with $F$ radio flux and $\\nu$ radio frequency) of the order of $\\sim$0.7-0.8, while typical radio-loud quasars exhibit values for $\\alpha$ between 0 and 0.5 even though, especially at high redshifts, there is a non negligible population of radio-loud sources with very high, $\\alpha>1$, values (see e.g. De Breuck et al. 2000).  Radio counterparts to the 510 optically faint mid-IR selected sources drawn from the whole First Look Survey sample (Fadda et al. 2006) have been searched in the overlapping region between MIPS and IRAC observations (Magliocchetti et al. 2007). Despite not having direct redshift estimates except for a handful of cases (e.g. Weedman et al. 2006; Yan et al. 2005; 2007; Martinez-Sansigre et al. 2006a), mid-IR photometry indicates that the overwhelming majority of such sources reside at redshifts $1.7\\simlt z\\simlt 2.5$ (\\S~4; see also Brand et al. 2007; Houck et al. 2005; Weedman et al. 2006a).\\\\ The First Look Survey region provides an excellent laboratory for investigations of the radio properties of dusty galaxies set at redshifts $z\\sim 2$ as its area has been observed at a number of radio frequencies down to very low flux densities (Condon et al. 2003; Morganti et al. 2004; Garn et al. 2007), therefore maximizing our chances of finding radio emitting galaxies.  The layout of the paper is as follows. In \\S2 we introduce the parent (mid-IR and radio) catalogues, while in \\S3 we present the matching procedure leading to the sample of optically obscured {\\it Spitzer}-selected sources with a radio counterpart at 1.4~GHz and/or 610~MHz. \\S4 uses IRAC photometry to provide some information on the typology of such objects (i.e. whether mainly powered by an obscured AGN or by a starburst) and also -- where possible -- to assign them to a redshift interval. \\S5 presents the results on the radio number counts both at 1.4~GHz and 610~MHz (\\S5.1) and on the relationship between radio and 24$\\mu$m emission (\\S5.2). \\S6 discusses our findings on the 1.4~GHz vs 610~MHz radio spectral indices for the objects in our sample, while \\S7 summarizes our conclusions. ",
        "conclusions": "This paper has presented an analysis of the radio properties of a subsample of optically faint ($R\\simgt 25.5$), 24$\\mu$m-selected galaxies observed by {\\it Spitzer} in the FLS (Magliocchetti et al. 2007). These objects have been cross-correlated with a number of radio catalogues which cover the same region of the sky, namely that of Condon et al. (2003) which probes 1.4~GHz fluxes brighter than $\\sim 100 \\mu$Jy, that of Garn et al. (2007) -- which probes 610~MHz fluxes brighter than $\\sim 200 \\mu$Jy -- and, on a smaller portion of the sky, that of Morganti et al. (2004) which reaches 1.4~GHz fluxes as faint as $\\sim 40\\mu$Jy.  70 optically faint {\\it Spitzer} sources have been identified in the Condon et al. (2003) catalogue, 33 in the Morganti et al. (2004) dataset, while 52 are found in the survey performed by Garn et al. (2007). After performing a number of corrections to account for multiple identifications, sources erroneously split in the original {\\it Spitzer} catalogue into different components and mid-IR objects with real radio counterparts at one of the two radio frequencies which were further away than the allowed ($10^{\\prime\\prime}$) matching radius, we end up with a sample of 96 radio-identified, optically faint, mid-IR emitting sources, 45 of which have an identification at both 1.4~GHz and 610~MHz. The fraction of radio identifications is a strong function of 24$\\mu$m flux: almost all sources brighter than $F_{24\\mu\\rm m}\\sim 2$~mJy are endowed with a radio flux at both 1.4~GHz and 610~MHz, while this fraction drastically decreases by lowering the flux level.  IRAC photometry for all those sources which also have detected fluxes in at least one of the four 8$\\mu$m, 5.8$\\mu$m, 4.5$\\mu$m and 3.6$\\mu$m channels (64 out of 96), allows to classify them into two categories: obscured AGN (45 sources) and systems mainly powered by starformation activity, (SB, 13 objects). We also find five low-z (i.e. $z\\simlt 0.5$ m51-type) interlopers, while the remaining 33 sources are unclassified. Furthermore, with the help of IRAC photometry it was possible to assign broad redshift intervals to all those sources (mostly AGN) which presented in the lowest 3.6$\\mu$m and $4.5\\mu$m wavelength channels of IRAC a 'bump' compatible with being produced by an evolved (old) stellar population. The majority ($\\sim 66$\\%) of these galaxies reside at redshifts $z\\simgt 1$, in agreement with other studies mainly based on mid-IR and near-IR spectroscopy of optically faint, 24$\\mu$m-selected galaxies (i.e. Weedman et al. 2006; Yan et al. 2005; 2007; Brand et al. 2007). We stress that this inferred fraction can only be considered as a lower limit to the real portion of faint {\\it Spitzer} sources set at high redshifts. In fact, because of their characteristic spectral properties in the mid-IR regime, we expect the majority of star forming systems (too faint at the IRAC frequencies to have measurable 3.6$\\mu$m and/or 4.5$\\mu$m fluxes) to be indeed located in the $z\\sim [1.6-2.7]$ redshift range.  A small fraction of objects in our sample present radio morphologies such as jets and/or lobes compatible with them being identified as radio-loud AGNs. Interestingly enough, we find that for most of these few extended radio sources the mid-IR emission is associated to such peripheral regions rather than steming from the centre of radio activity, generally coinciding with the location of the AGN. However, the majority of the objects in our sample present unresolved radio images.  A compared analysis of the radio number counts for optically obscured {\\it Spitzer} sources indicates that the $\\Delta N/\\Delta F$ as estimated at 1.4~GHz can only be reconciled with what found at 610~MHz if the population under investigation is endowed with an average value for the radio spectral index (defined as $F_R\\propto \\nu_R^{-\\alpha}$, where $F_R$ is the radio flux and $\\nu_R$ the generic radio frequency) $\\left<\\alpha\\right>\\simgt 1$. Classical, 'flat spectrum' radio sources can confidently be excluded as the typical objects constituting our sample. Furthermore, we have found that the radio number counts of sources classified as AGN and of those identified as starburst galaxies are quite similar, evidence which suggests that radio emission in $z\\sim 2$ {\\it Spitzer} galaxies originates from similar process(es), despite of the different mid-IR emission.  Direct investigations of the relation between 24$\\mu$m and 1.4~GHz fluxes show that the overwhelming majority of those galaxies not excluded from our analysis because morphologically classified as radio-loud AGN follow the relationship identified for $0\\simlt z\\simlt 1$ star-forming objects by Appleton et al. (2004), although with a large scatter. This happens regardless of whether the galaxy has been classified as an AGN or a starforming system on the basis of its mid-IR colours. The distribution of 24$\\mu$m vs 610~MHz fluxes is instead found to be more scattered.  The majority of these radio-identified objects (26, a figure which rises to 34 if one also includes sources with estimated lower limits on $\\alpha$) present very steep, $\\alpha> 1$ radio spectral indices, some galaxies being endowed with $\\alpha$'s as high as 2.5. This excess of galaxies with large $\\alpha$ values is statistically significant as it corresponds to 60 per cent of our sample, to be compared with the 43 per cent found by considering {\\it the whole population} of faint radio objects with an identification at both 610~MHz and 1.4~GHz.\\\\ Such very high figures for $\\alpha$ would identify the corresponding sources as Ultra Steep Spectrum galaxies, generally high redshift radio-loud AGN. However this is in striking disagreement with what found for the relation between 24$\\mu$m and 1.4~GHz fluxes for the very same objects, relation which would explain their (1.4~GHz) radio emission as mainly due to processes connected with star forming activity.  A natural explanation to the above issues could be found by assuming that AGN and star-formation activity are concomitant in the majority of $z\\sim 2$, {\\it Spitzer} sources, at least in those which present enhanced radio emission. The radio signal steming from these systems would then simply be the combination of two components: a shallower one -- dominating the spectrum at 1.4~GHz -- due to processes connected with star formation, and a steeper one -- being responsible for most of the 610~MHz signal -- connected with AGN activity. This framework could then also explain why the 1.4~GHz emission in our sources follows that of star forming systems, while the same does not seem to happen in the case of 610~MHz fluxes.  Various explanations can be found in the literature to account for the presence of USS: from slow adiabatic expansion losses in high-density environments (e.g. Klamer et al. 2006) to the scattering between CMB photons and relativistic electrons at $z\\sim 2$ (e.g. Martinez-Sansigre et al. 2006). However, the results presented in this work might provide an alternative scenario. In fact, they suggest that high values for $\\alpha$ might be due to the concomitant presence within the same systems of an AGN and of a star forming region: the AGN expansion would then be halted by the encounter with the cooler/denser sites in which star formation takes place. This would determine the 'strangling' of the AGN, causing its radio spectrum at low radio frequencies to steepen. Clearly, more theoretical work is needed in order to quantify the above issue and we are planning to present it in a forthcoming paper. For the time being, we note that intense star forming activity within a high redshift galaxy host of a USS has been recently reported by Hatch et al. (2007).  From a more observational point of view, the 'ultimate truth' on these sources could only come from very high resolution (and, given their faintness, sensitivity) measurements, capable to clearly disentangle the emission related to the AGN to that associated to star forming regions. The advent of instruments such as ALMA will then provide the answers we need."
    },
    "0710/0710.2452_arXiv.txt": {
        "abstract": "We present  highlights and an overview of 20 {\\em FUSE} and {\\em HST}/STIS observations of the bright symbiotic binary \\object{EG And}. The main motivation behind this work is to obtain spatially-resolved information on an evolved giant star in order to understand the mass-loss processes at work in these objects. The system consists of a low-luminosity white dwarf and a mass-losing, non-dusty M2 giant. The ultraviolet observations follow the white dwarf continuum through periodic gradual occultations by the wind and chromosphere of the giant, providing a unique diagnosis of the circumstellar gas in absorption.  Unocculted spectra display high ionization features, such as the \\ion{O}{6} resonance doublet which is present as a variable ({\\em hourly} time-scales), broad wind profile, which diagnose the hot gas close to the dwarf component.  Spectra observed at  stages of partial occultation display a host of low-ionization, narrow, absorption lines, with transitions observed from lower energy levels up to $\\sim$5 eV above ground. This absorption is due to chromospheric/wind material, with most lines due to transitions of \\ion{Si}{2}, \\ion{P}{2}, \\ion{N}{1}, \\ion{Fe}{2} and \\ion{Ni}{2}, as well as heavily damped \\ion{H}{1} Lyman series features. No molecular features are observed in the wind acceleration region despite the sensitivity of {\\em FUSE} to H$_2$. From analysis of the ultraviolet dataset, as well as optical data, we find that the  dwarf radiation does not dominate the wind acceleration region of the giant, and that observed thermal and dynamic wind properties are most likely representative of isolated red giants. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.1876_arXiv.txt": {
        "abstract": "{During the period 1966.5--2006.2 the 15GHz and 8GHz lightcurves of 3C454.3 (z=0.859) show a quasi-periodicity of $\\sim$12.8 yr ($\\sim$6.9 yr in the rest frame of the source) with a double-bump structure. This periodic behaviour is interpreted in terms of a rotating double-jet model in which the two jets are created from the black holes in a binary system and  rotate with the period of the orbital motion. The periodic variations in the radio fluxes of 3C454.3 are suggested to be mainly due to the lighthouse effects (or the variation in Doppler boosting) of the precessing jets which are caused by the orbital motion. In addition, variations in the mass-flow rates accreting onto the black holes may be also involved.\\\\ ",
        "introduction": "In recent years many works have been devoted to the study of the periodic properties discovered in extragalactic sources, especially in Blazars. These include the periodic (or quasi-periodic) variations of flux density (or luminosity) in optical bands (for example, for OJ287: Sillanp{\\\"a}{\\\"a} et al. \\cite{Si88}, Lehto \\& Valtonen \\cite{Le96}, Valtonen et al. \\cite{Va99}, Villata et al. \\cite{Vi98}, Valtaoja et al. \\cite{Va00}, Valtonen \\cite{Va06a}, \\cite{Va06b}; AO\\,0235+16: Raiteri et al. \\cite{Ra01}); periodic variations of flux density in radio bands (for example, for AO0235+16: Raiteri et al. \\cite{Ra01}, 3C454.3: Ciaramella et al. \\cite{Ci04}, Kudryavtseva \\& Pyatunina \\cite{Ku06}); the periodic changes of the position angle of the ejection of superluminal knots (for example, Qian et al. \\cite{Qi01}, Britzen et al. \\cite{Br01}, Abraham \\& Carrara \\cite{Ab00}, Lobanov \\& Roland \\cite{Lo05}); periodic changes of the trajectory of superluminal knots (for example, for 3C345: Qian et al. \\cite{Qi91}, Lobanov \\& Roland \\cite{Lo05}, Klare et al. \\cite{Kl05}), etc. Up to now, the most reliable periodic phenomenon observed in blazars may be the 12 year quasi-periodic optical variations of the BL Lac object OJ287 which has been convincingly determined in the more than 100 year long records of optical observations since 1890.\\\\ In the interpretations of these periodic phenomena various binary balck hole models have been invoked. For the periodic optical variations in OJ287  all the proposed models invoke a binary system with two supermassive black holes. But the mechanisms for the production of the optical outbursts can be divided into two classes: accretion models and lighthouse models. \\begin{itemize} \\item Accretion models\\\\ These models (Lehto \\& Valtonen \\cite{Le96}, Valtonen \\cite{Va06a}, \\cite{Va06b}) propose that the optical periodicity is caused by a precessing binary black hole system of two supermassive black holes with a large eccentric orbit. The observed periodic 'superflares'  with double peaks are suggested to be due to the crossings of the secondary black hole into the accretion disk of the primary black hole during the pericenter passage. Thus in this class of models the optical emission is thermal (bremsstrahlung) and the observed flux increase reflects the enhanced accretion related to the disk crossings and  the general enhancement of the accretion rate during the pericenter passage. Doppler boosting is not taken into account. Detailed models for a binary black hole system have been proposed to fit the past  optical records for OJ287 and predict the future optical events. If the timing of the optical activity (superflares) in OJ287 can be accurately predicted, then it is possible to calculate or determine the parameters of the binary black hole system (such as the masses of the black holes, the orbital parameters etc.).  \\item Lighthouse models\\\\ The periodic optical flares observed in OJ287 are suggested to be caused by the change of the orientation of the realtivistically moving emission regions with respect to the line of sight, resulting in an increase of the Doppler boosting factor (Camenzind \\& Krockenberg \\cite{Ca92}, Katz \\cite{Ka97}, Villata et al. et al. \\cite{Vi98}, Qian et al. \\cite{Qi01}) and the apparent optical flux density (${\\propto}{{\\delta}^3}$). Katz (\\cite{Ka97}) proposed that in a bianry black hole system the companion exerts a torque on the accretion disk of the primary black hole and results in the precessing of the relativistic jet from the primary sweeping across the line of sight and the periodic variation of the Doppler factor and thus the apparent flux density. Villata et al. (\\cite{Vi98}) suggests that the two black holes of the binary system both create relativistic jets which are bent significantly in different directions. In the course of the binary's orbit motion, the directions of the bent parts of the jets from the two black holes rotate with the orbital period, resulting in periodic double-peak flares. In this class of models only the variation of the Doppler boosting factor plays the role, not considering any change in accretion in the central disk-hole systems. \\item Alternative models\\\\ Valtaoja et al. (\\cite{Va00}) have argued against both the accretion models and lighthouse models for OJ287. They found that in OJ287 the first optical flare of the double-structure is thermal (non-polarized) and lacks radio counterpart, but the second one is (polarized) synchrotron emission, having a radio counterpart. They proposed an alternative model in which the first optical flare is caused by the disk-crossing during the pericenter passage, having a regular period. And the second one occurs about one year later in the relativistic jet with an association with a radio outburst. For checking this model optical polarization measurements and VLBI studies of the optical-radio association are significant. \\end{itemize} In this paper we will discuss the possible existence of a $\\sim$13 year periodicity in the radio lightcurve of the optically violently variable (OVV) quasar 3C454.3 at 8GHz and 14.5GHz in which two broad peaks (or bumps) are observed during one period. We propose that the periodicity can be interpreted in the frame of a binary black hole model in which two jets from the two black holes rotate with the period ($\\sim$13 year) of the orbital motion. Model-fits to the lightcurves are given. It is shown that both the periodic Doppler boosting effect (lighthouse effect) and the variability of the accretion activity (i.e. the mass-energy inflow into the jets) should be taken into account in order to fully explain the lightcurves. We also discuss some alternative models and future observations for further studying the periodic phenomenon in 3C454.3. ",
        "conclusions": "\\begin{itemize} \\item (1) the supposed  model is not unique. In addition to the double-jet models proposed in this paper, models with a single jet are also plausible. But in this case, the two emitting components could be assumed to be situated at different positio ns in the curved jet and their rotation would produce the Doppler boosting profiles. In this paper we do not invoke a specific model for a binary black hole system. \\item (2) In our model-fittings, multi-parameters can be chosen. When  parameters were chosen we consider mainly three ingredients: (a) the width of the two radio bumps; (2) the difference of Doppler factor at flaring and quiescent epochs should not be too large in order to avoid a huge difference in timescale of varibility at different epochs (Valtaoja et al. \\cite{Va00}); \\item (3)  We should consider both effects due to Doppler boosting and the energy transfer into the jets (or accretion rates). In our models, the Doppler boosting determine the 'profiles' of the possible activity in the radio and optical bands. If the intrinsic (in the rest frame of the flows) strengths are stable, then the apparent activity follows the pattern of the Doppler boosting. This seems correct for the observed radio variations, for example during 1978--1995. \\item (4) As for the optical variations, they do not follow the Doppler boosting profiles derived from our models. This could imply we should adopt different parameters for the optical variations, i.e., the optically emitting regions could have orientations different from the radio regions. Moreover, optical flux variations could be largely determined by the variation in accretion rate and their much shorter timescales may be due to mechanisms for acceleartion of electrons and radiative losses. The optical peak at 2005.4 and radio peak at 2006.2, which both deviate the Doppler boosting pattern maximum, further indicate the necessity to consider the  effects both from Doppler boosting, accretion and transfer of enery into the jets. \\item (5) The radio periodic behaviour still needs to be tested, because the 2005 optical flare has a quite peculiar mm-radio outbursts, which have much narrower profiles than before, very different from the 1981 and 1994 radio flare's profiles. This may imply that the radio periodic behaviour could be a short-term phenomenon. The optical-radio relation is still important to be tested and thus could obtain some information about the combination of the black hole disk system and the Doppler boosting effect. \\item (6) Relationship between mm- and cm-outburtsts needs to be further studied in order to study the evolution of the outbursts. \\item  (7) At present,  we cannot predict whether the periodic behaviour in 3C454.3 in radio and optical bands can hold long, and this should be observationally tested during the next period (2007--2020). Both optical and radio monitoring (intensity and polarization) are required. \\end{itemize}"
    },
    "0710/0710.3042_arXiv.txt": {
        "abstract": "{We show that future Ultra-High Energy Cosmic Ray samples should be able to distinguish whether the sources of UHECRs are hosted by galaxy clusters or ordinary galaxies, or whether the sources are uncorrelated with the large-scale structure of the universe.  Moreover, this is true independently of arrival direction uncertainty due to magnetic deflection or measurement error.  The reason for this is the simple property that the strength of large-scale clustering for extragalactic sources depends on their mass, with more massive objects, such as galaxy clusters, clustering more strongly than lower mass objects, such as ordinary galaxies. }  \\begin{document} ",
        "introduction": "\\label{sec:intro} Identifying the sources of ultrahigh energy cosmic rays (UHECRs, here $E>10^{19}$eV $\\equiv 10$ EeV) is complicated by the deflection they presumably experience in Galactic and extragalactic magnetic fields, as well as their relatively poor arrival direction determinations, typically $\\sim 1^{\\circ}$.  Arrival directions of most UHECRs are thus not known well enough to match their positions with specific astrophysical objects. However, there is also useful information in the clustering of UHECRs on large scales, where $\\sim$ few degree uncertainties in position become unimportant.  The clustering of galaxies in the universe is typically quantified by the two-point correlation function or its analog in Fourier space, the power spectrum.  The two-point correlation function $\\xi(r)$ of any class of objects (e.g., galaxies of a certain luminosity or color) is defined as the excess number of pairs of such objects at physical separation $r$ over that expected for a random (Poisson) distribution. In Cold Dark Matter models, the large-scale amplitude of $\\xi(r)$ (usually referred to as the bias) of a population of objects depends only on their mass, with more massive objects, such as clusters of galaxies, clustering more strongly than less massive objects, such as ordinary galaxies \\cite{mo_white_96,sheth_tormen_99,berlind_etal_06b}. The large-scale bias of a UHECR sample is therefore a robust and informative measure of the clustering properties of the source.  We cannot measure physical separations for pairs involving UHECRs because they do not have measured redshifts.  However, we can measure the angular correlation function $\\omega(\\theta)$.  As is the case for $\\xi(r)$, the large-scale amplitude of $\\omega(\\theta)$ for a UHECR sample depends on the nature of the astrophysical source.  However, it also depends on the depth of the sample because deeper samples mix more physically uncorrelated pairs and thus show weaker angular clustering.  In order to access the information in the large-scale angular clustering of UHECRs, we must therefore know the depth of our UHECR sample.  In this paper, we demonstrate what can be learned from the large-scale angular clustering of UHECRs, we estimate what kind of sample is needed to do this analysis, and we show how to deal with the unknown depth of a UHECR sample, using the GZK effect. ",
        "conclusions": ""
    },
    "0710/0710.4986_arXiv.txt": {
        "abstract": "We calculate axisymmetric toroidal modes of magnetized neutron stars with a solid crust. We assume the interior of the star is threaded by a poloidal magnetic field that is continuous at the surface with the outside dipole field whose strength $B_p$ at the magnetic pole is $B_p\\sim10^{16}$G. Since separation of variables is not possible for oscillations of magnetized stars, we employ finite series expansions of the perturbations using spherical harmonic functions to represent the angular dependence of the oscillation modes. For $B_p\\sim 10^{16}$G, we find distinct mode sequences, in each of which the oscillation frequency of the toroidal mode slowly increases as the number of radial nodes of the eigenfunction increases. The frequency spectrum of the toroidal modes for $B_p\\sim10^{16}$G is largely different from that of the crustal toroidal modes of the non-magnetized model, although the frequency ranges are overlapped each other. This suggests that an interpretation of the observed QPOs based on the magnetic toroidal modes may be possible if the field strength of the star is as strong as $B_p\\sim10^{16}$G. ",
        "introduction": "Recent discovery of quasi-periodic oscillations (QPOs) of magnetar candidates is one of the observational manifestations of global oscillations of neutron stars. Israel et al (2005) detected QPOs of frequencies $\\sim18$, $\\sim 30$ and $\\sim$92.5Hz in the tail of the SGR 1806-20 hyperflare observed December 2004, and suggested that the 30Hz and 92.5Hz QPOs could be caused by seismic vibrations of the neutron star crust (see, e.g., Duncan 1998). Later on, in the hyperflare of SGR 1900+14 detected August 1998, Strohmayer \\& Watts (2005) found QPOs of frequencies 28, 53.5, 84, and 155 Hz, and claimed that the QPOs could be identified with the low $l$ fundamental toroidal torsional modes of the solid crust of the neutron star. These recent discoveries of QPOs in the giant flares of Soft Gamma-Ray Repeaters SGR 1806-20 (Israel et al 2005, Watts \\& Strohmayer 2006, Strohmayer \\& Watts 2006) and SGR 1900+14 (Strohmayer \\& Watts 2005) have made promising asteroseismology for magnetars, neutron stars with an extremely strong magnetic field (see, e.g., Woods \\& Thompson 2006 for a review of SGRs).  It is currently common to identify these QPOs with seismic vibrations caused by crustal toroidal modes of the neutron stars, since the frequency range of the modes overlap that of the observed QPOs and from the energetics point of view the crustal toroidal modes would be most easily excited to observable amplitudes by spending a least amount of available energies, for example, those released in magnetic field restructuring (e.g., Duncan 1998). Although the interpretation based on crustal torsional modes looks promising, we need detailed theoretical analyses of oscillations of magnetized neutron stars so that we could get information of physical conditions of the stars through the confrontation between theoretical modelings and observations. This is particularly true for high frequency QPOs (e.g., 625Hz QPO, Watts \\& Strohmayer 2006; 1835Hz QPO and less significant QPOs at 720 and 2384 Hz in SGR 1806-20, Strohmayer \\& Watts 2006), since there exist classes of modes other than the crustal toroidal modes that can generate the frequencies observed.  The presence of a magnetic field makes it possible for toroidal modes to exist in a fluid star even without rotation as does the presence of the shear modulus in the solid crust. In this paper we are interested in axisymmetric toroidal modes since axisymmetirc toroidal and spheroidal modes are decoupled for a poloidal field when the star is non-rotating. For non-axisymmetric modes, the toroidal and spheroidal components are coupled even without rotation and hence the modal analyses of magnetized stars would be much more complicated.  Theoretical calculations of toroidal modes of strongly magnetized neutron stars have been carried out by several authors, including Piro (2005), Glampedakis et al (2006), Sotani et al (2006, 2007), and Lee (2007). The analyses by Piro (2005), Glampedakis et al (2006), and Lee (2007) assume Newtonian gravity, while those by Sotani et al (2006, 2007) use general relativistic formulation. Although the studies by Piro (2006) and Lee (2007) ignore the effects of magnetic fields in the fluid core, those by Glampedakis et al (2006) and Sotani et al (2006, 2007) consider magnetic waves propagating in the fluid core, assuming the core is threaded by a magnetic field of substantial strength. Besides the differences mentioned above, most of the authors except Lee (2007) represent the angular dependence of the oscillations by a single spherical harmonic function $Y_l^m(\\theta,\\phi)$. Since the shear modulus in the crust dominates the magnetic pressure in most parts of the crustal regions for a dipole field of strength $B_p\\ltsim 10^{15}$G, this treatment may be justified so long as the crustal toroidal modes are well decoupled from the fluid core. But, if the torsional waves in the crust are strongly coupled with magnetic waves in the core, the treatment may not be justified because the angular dependence of the magnetic waves in the fluid core cannot be correctly represented by a single spherical harmonics. For example, the analysis by Reese, Fincon, \\& Rieutord (2004) of toroidal modes in a fluid shell have employed finite series expansions of long length for the perturbations.  In this paper, using the method of series expansions of perturbations we calculate toroidal modes of a strongly magnetized neutron star having a fluid core and a solid crust, where the entire interior is assumed to be threaded by a poloidal magnetic field. We employ two different sets of oscillation equations, one for fluid regions and the other for the solid crust, and solutions in the solid and fluid regions are matched at the interfaces between them to obtain an entire solution of a mode. The method of calculation we employ is presented in \\S 2, and the numerical results are given in \\S 3, and conclusions are in \\S 4. ",
        "conclusions": "In this paper, we have calculated toroidal modes of magnetized stars with a solid crust, where the entire interior of the star is assumed to be threaded by a poloidal magnetic field that is continuous at the stellar surface to the outside dipole field. We find distinct mode sequences of the toroidal modes, in each of which the mode frequency remains rather constant and only slowly increases as the radial order of the modes increases. In the presence of a solid crust, the frequency separation between the lowest and second lowest frequency mode sequences is approximately given by $\\Delta\\omega$, but that between the second and third lowest frequency mode sequences by $\\Delta\\omega/2$, where the frequency separation $\\Delta\\omega$ is roughly proportional to the field strength $B_p$. This frequency pattern of the low frequency sequences is different from that found for the model without the crust, for which the frequency separation between the sequential sequences of low frequency modes is given by $\\Delta\\omega/2$. We also find that for the equation of state we use, $\\Delta\\omega$ is larger for smaller $M$.  The eigenfunction $\\xi_\\phi$ of the modes belonging to the lowest frequency sequence have much larger amplitudes than $B^\\prime_\\phi$ and can penetrate into the solid crust. On the other hand, $\\xi_\\phi$ of the modes belonging to the higher frequency sequences is much smaller than $B^\\prime_\\phi$, which is well confined into the fluid core and does not have any substantial amplitudes in the solid crust.   The frequency ranges of the toroidal modes we find for the magnetized neutron star with $B_p\\sim10^{16}$G overlap the QPO frequencies found for the magnetar candidates, SGR 1806-20 and SGR 1900+14. This suggests that we may interpret  the observed QPOs based on the magnetic toroidal modes, and that detailed comparisons between observed frequency spectra and theoretical calculations make it relevant to infer physical parameters of the magnetar candidates, such as the equation of state and the strength of the magnetic field. But, we think it worth pointing out that except for the modes belonging to the lowest frequency sequence, the magnetic perturbations, which have much larger amplitude than $\\xi_\\phi$, are well confined in the fluid core and do not have substantial amplitudes in the solid crust, which make it difficult for the modes to be directly observable. If the magnetar candidates do no have a crust, the problem of observability could be avoided. In this case, however, we have to use more realistic surface boundary conditions than $\\pmb{B}^\\prime=0$ used in the present calculations. Note that, although the existence of a solid crust affects the frequency pattern of the low frequency mode sequences, the frequency range of the magnetic modes itself is not very much dependent on the presence or absence of a solid crust unless the crust is extremely thick.   We have tried to find toroidal modes well confined in the solid crust or core toroidal modes that are in resonance with the crustal toroidal modes, but failed. We also find it extremely difficult to identify distinct toroidal mode sequences for magnetic fields weaker than $B_p\\ltsim10^{15}$G when a solid crust is included in the models. It is therefore not clear whether the toroidal mode sequences of the kind we find in the present paper can survive also for weakly magnetized neutron stars with a solid crust. If the field strength is much weaker than $B_p\\sim10^{15}$G, the magnetic fields may have only minor effects on the crustal toroidal modes (e.g., Lee 2007), and it will be justified to use frequency spectra of the crust modes theoretically obtained in the weak field limit to interpret observed QPOs. As briefly noted in the last section, it is not theoretically clear how the frequency spectra of the toroidal modes of magnetized stars should look like, and how the toroidal modes behave in the limit of $n\\ge k\\rightarrow\\infty$. We may even speculate that the frequency spectra we obtained reflect the existence of continuous spectra (see, e.g., Goedbloed \\& Poedts 2004, see also Levin 2007), but the detailed numerical analysis of continuous frequency spectra is beyond the scope of this paper. It will be worthwhile to examine the effects of an interior toroidal field on magnetic modes. It is also needed to extend the present analysis to a general relativistic formulation (e.g., Sotani et al 2006, 2007)."
    },
    "0710/0710.1271_arXiv.txt": {
        "abstract": "We present results of 3-neutrino flavor evolution simulations for the neutronization burst from an O-Ne-Mg core-collapse supernova. We find that nonlinear neutrino self-coupling engineers a single spectral feature of stepwise conversion in the inverted neutrino mass hierarchy case and in the normal mass hierarchy case, a superposition of two such features corresponding to the vacuum neutrino mass-squared differences associated with solar and atmospheric neutrino oscillations. These neutrino spectral features offer a unique potential probe of the conditions in the supernova environment and may allow us to distinguish between O-Ne-Mg and Fe core-collapse supernovae. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.3104_arXiv.txt": {
        "abstract": "The Milky Way is the only galaxy for which we can resolve individual stars at all evolutionary phases, from the Galactic center to the outskirt. The last decade, thanks to the advent of near IR detectors and 8 meter class telescopes, has seen a great progress in the understanding of the Milky Way central region: the bulge. Here we review the most recent results regarding the bulge structure, age, kinematics and chemical composition. These results have profound implications for the formation and evolution of the Milky Way and of galaxies in general. This paper provides a summary on our current understanding of the Milky Way bulge, intended mainly for workers on other fields. ",
        "introduction": "How did the Milky Way form? Decades ago, the Galactic formation scenarios focused on the disk and the halo populations, because these were the components that astronomers knew something about. For example, in the classic works of Eggen et al. (1962) or Searle \\& Zinn (1978) the bulge was not mentioned. All-sky optical maps (Fig.~\\ref{maps}) did not show a clear, separate component in the inner Milky Way.  Yet upon looking at the new DIRBE or 2MASS near-infrared maps of the whole sky, it is evident that the Milky Way is a spiral galaxy with a peanut-shape bulge. The simultaneous observation that bulge stars were mainly old made it clear that to answer this question the attention had to shift towards what seems to be the first massive component to be formed in the Milky Way. This is a Copernican-like revolution on Galactic scales, and it is happening now! We see our Galaxy as if it were an external galaxy for the first time, and are fortunate to have such new perspective, and also to be able to provide specific answers to the many questions regarding its formation.  \\begin{figure} \\includegraphics[height=5.5in,width=5.3in,angle=00]{minniti_fig1.ps} \\caption{ Top: Near-IR COBE-DIRBE map of the whole sky (Dwek et al. 1995). Bottom: Optical map for comparison (Copyright Axel Mallenhoff 2001). These sky maps illustrate why astronomers did not realize before the importance of the bulge.} \\label{maps} \\end{figure}   Note that we have only one object to study: our Galaxy. This is a fundamental limitation because we have to get the answers right. The advantage in the case of our Galaxy, of course, is that we can study the subject in detail: in no other galaxies the fundamental problems of Astronomy can be surpassed. Basically, these fundamental problems are: (1) The distance problem: we can resolve the components of the Milky Way into stars, and study them in situ, obtaining 3-D positions and motions, and detailed chemical compositions of individual stars; and (2) The timescale problem: we see an instant snapshot, and must infer histories from that. In our galaxy we can separate and date the components, using the main-sequence turn-offs of different stellar populations and clusters.  In spite of a small community working on the Galactic bulge, a revolution has occurred in this field in the last decade, with great progress in comparison with other fields of Astronomy. For the first time, we have the answers to important questions such as:  {\\bf 1.} Is the bulge a different Galactic component?  {\\bf 2.} When did the Galactic bulge form?  {\\bf 3.} How did it form?  {\\bf 4.} Is there a radial gradient in the bulge?  {\\bf 5.} Are there globular clusters associated with the bulge?  {\\bf 6.} Are there planets in the bulge?  The 8m class telescopes were built last decade, aiming to answer these questions. The proposed options were endless (bulge formation from secular evolution of the disk, from the halo, in a single burst, in several episodes, by slow accretion of smaller sub-units, etc). The evidence and the answers described below serve  as a  basis for understanding   more    distant galaxies    which  cannot  be  studied (dissected) in similar detail. ",
        "conclusions": "For the first time, we have the answers to the following basic questions:  {\\bf 1.} Is the bulge a different component? Yes, based on all the evidence available, the bulge is a distinct Galactic component, with different kinematics and compositions from the thin disk, the thick disk and the halo.  {\\bf 2.} How did the Galactic bulge form? It formed on a short timescale ($\\sim 1$ Gyr), as demonstrated by the $\\alpha$-element enrichment. Despite the presence of the bar, models like a bulge formation via secular evolution of the disk can be firmly excluded.  {\\bf 3.} What is the age of the bulge? The bulk of the stellar population is $\\sim 10$ Gyr old. However, there are traces of a small fractions of intermediate-age stars, and of metal-poor stars. The latter might well be the oldest population in the Galaxy.  {\\bf 4.} Is there a gradient in the bulge? Yes, there is a stellar population gradient as shown for example by the CMDs. Now it is found that this gradient is mostly due to metallicity, which decreases along the Galactic minor axis.  {\\bf 5.} Are there globular clusters associated with the bulge? Yes, there is a population of metal-rich globular clusters in the central regions that share the kinematics, spatial distribution, and composition of the bulge field stars.  {\\bf 6.} Are there planets in the bulge? Even though this question seems to belong to another field, another of the recent advances was the discovery of planets in the bulge by HST. The available data suggests that giant planets are as numerous in the bulge as they are in the Solar neighborhood.  These revolutionary advances that impact the whole of extragalactic Astronomy cannot be attributed to the success of a single group, but to the combined contributions of different teams. Where controversy was present before, today similar answers are given. Progress has occurred!"
    },
    "0710/0710.3618_arXiv.txt": {
        "abstract": " ",
        "introduction": "In the course of carrying out a wide-field CCD imaging survey, two new methods for correlating the images to star catalogues have been developed, motivated by the need to efficiently handle the large number of stellar sources present on the images. Most previously published algorithms successfully cater for small lists ($\\le$ 50 stars), but do not scale well to wide-fields containing $10^3$ or more stellar sources.  The problem of matching coordinate lists of point sources is a necessary prerequisite for deriving an astrometric plate solution. The objective is to match a subset of stars found on an image to their corresponding entries in a stellar catalogue in order to determine the transformation between detector coordinates and sky coordinates. The algorithm must handle translation and rotation, and small changes in scale caused by temperature related changes in focal length. In addition, it must cope with additional and missing stars. That is, the two lists may only partially overlap.  The efficiency of the algorithm is of paramount concern, since it is embodied within the closed-loop pointing system of the telescope and therefore affects the duty-cycle time, and ultimately constrains the number of images that can be acquired each night. Surveys that require very high photometric precision typically seek to accurately align their fields on the same detector pixels each night to overcome residual flat-fielding errors \\citep{Everett}, and would benefit from the efficiency gains of a fast matching algorithm. Similarly, high cadence surveys, such as the Southern Sky Survey \\citep{Keller} could improve precision and reduce its duty-cycle by utilizing a fast closed-loop pointing algorithm. Moreover, real-time attitude adjustments on spacecraft might be possible with the aid of an efficient matching algorithm to analyze on-board star camera images \\citep[see for example][]{Fraser}.  A number of algorithms have been proposed to solve this problem. \\citet{Groth} describes an algorithm that matches geometrically similar shapes (triangles) in the two lists. By limiting the number of triangles constructed, and by only matching those triangles whose ratio of longest to shortest side are within a defined limit, his matching phase has a computational complexity of $O(n^{4.5})$ where $n$ is the number of stars in each list. \\cite{Stetson} describes a very similar algorithm that he developed independently at around the same time.  \\citet{Murtagh} reviews a number of approaches and proposes his own, based upon characterization of a set of coordinates couples, with matching based on the proximity of feature vectors in the two lists. His method's matching phase has a computational complexity of $O(n^2)$.  Nevertheless, Groth's algorithm appears to be the most widely accepted, with the methods applied across disciplines. For example, \\citet{Arz} discuss its application to the problem of computer-aided identification of whale sharks, while \\citet{Mars}, building upon the work of \\citet{Groth}, describe an optimization to the voting phase of the algorithm, concluding that their method reduces the need for complicated filtering methods while successfully reducing the number of false matches.  More recently, \\citet{PB} describe another variation of triangle matching, optimized to handle large lists of objects extracted from wide field images. Large fields contain thousands of stars and pose a severe test for matching algorithms, requiring efficient methods to accommodate the large number of point sources.  The following sections discuss two new methods for pattern matching that have a matching phase with a complexity that is nearly $O(1)$, at the cost of a slight loss in generality. They are collectively referred to as \\emph{Optimistic Pattern Matching (OPM)} because they assume that (i) a good match is likely to be found, and (ii) the scale of the image is approximately known, thus permitting the use of an early exit strategy whereby only a small percentage of the candidate list is examined. By contrast, previous methods assumed an unknown scale which required the entire candidate list to be processed to determine the most likely match using a statistical approach. This required additional phases and complexity. In practice, an \\emph{a priori} knowledge of an instrument's focal length is common place, and the use of a more general algorithm that assumes it is unknown mandates strategies that unnecessarily degrade performance.  Section 2 describes the algorithms in detail. $OPM_A$ is based upon a new definition of triangle space, while $OPM_B$ uses an alternative shape characterization method. Section 3 tests their performance using a large sample of survey images and compares them to earlier methods. Conclusions are summarized in Section 4. ",
        "conclusions": "Two new techniques for matching two-dimensional coordinate lists in nearly constant time have been presented. The matching phase of $OPM_B$ is nearly $O(1)$, being independent of list size. These algorithms have a significant performance advantage over previous techniques, at a slight loss in generality, caused by the requirement that the approximate focal length of the optical system is known \\emph{a priori}. This requirement permits the determination of the image scale from the physical dimensions of the detector, allowing $OPM$ algorithms to directly compare a subset of triangles (or shapes) to their counterparts derived from a reference catalogue, without having to process the entire set, as is the case when the scale is unknown. By employing early exit strategies, postponing work until absolutely necessary, testing candidates in the order most likely to yield success, and combining these with and an efficient mechanism for rejecting false positives, a highly efficient search, in nearly constant time is possible.  Small uncertainties in the focal length, such as caused by temperature related changes, are accommodated by selecting an appropriate matching tolerance. The actual focal-length is determined and reported as part of the astrometric solution.  The $OPM$ algorithms are particularly suited to processing large lists or in situations where pattern matching must be performed as quickly as possible. The performance of these algorithms makes it practical to search thousands of fields very quickly, if for example, the coordinates of the field center were unknown. Similarly, when only an approximate focal-length is known, it is perfectly reasonable to attempt to iteratively match the field using a range of focal-lengths."
    },
    "0710/0710.2102_arXiv.txt": {
        "abstract": "We use hydrodynamical simulations of disk galaxies to study relations between star formation and properties of the molecular interstellar medium (ISM). We implement a model for the ISM that includes low-temperature ($T<10^{4}$~K) cooling, directly ties the star formation rate to the molecular gas density, and accounts for the destruction of $\\Hmol$ by an interstellar radiation field from young stars.  We demonstrate that the ISM and star formation model simultaneously produces a spatially-resolved molecular-gas surface density Schmidt-Kennicutt relation of the form $\\SigmaSFR \\propto \\SigmaHmol^{\\nmol}$ with $\\nmol\\approx1.4$ independent of galaxy mass, and a total gas surface density -- star formation rate relation $\\SigmaSFR \\propto \\Sigmagas^{\\ntot}$ with a power-law index that steepens from $\\ntot\\sim2$ for large galaxies to $\\ntot\\gtrsim4$ for small dwarf galaxies.  We show that deviations from the disk-averaged $\\SigmaSFR \\propto \\Sigmagas^{1.4}$ correlation determined by \\cite{kennicutt1998a} owe primarily to spatial trends in the molecular fraction $\\fHmol$ and may explain observed deviations from the global Schmidt-Kennicutt relation. In our model, such deviations occur in regions of the ISM where the fraction of gas mass in molecular form is declining or significantly less than unity. Long gas consumption time scales in low-mass and low surface brightness galaxies may owe to their small fractions of molecular gas rather than mediation by strong supernovae-driven winds.  Our simulations also reproduce the observed relations between ISM pressure and molecular fraction and between star formation rate, gas surface density, and disk angular frequency. We show that the Toomre criterion that accounts for both gas and stellar densities correctly predicts the onset of star formation in our simulated disks. We examine the density and temperature distributions of the ISM in simulated galaxies and show that the density probability distribution function (PDF) generally exhibits a complicated structure with multiple peaks corresponding to different temperature phases of the gas. The overall density PDF can be well-modeled as a sum of lognormal PDFs corresponding to individual, approximately isothermal phases. We also present a simple method to mitigate numerical Jeans fragmentation of dense, cold gas in Smoothed Particle Hydrodynamics codes through the adoption of a density-dependent pressure floor. ",
        "introduction": "\\label{section:introduction}    Galaxy formation presents some of the most important and challenging problems in modern astrophysics.  A basic paradigm for the dissipational formation of galaxies from primordial fluctuations in the density field has been developed \\citep[e.g.][]{white1978a,blumenthal1984a,white1991a}, but many of the processes accompanying galaxy formation are still poorly understood.  In particular, star formation shapes the observable properties of galaxies but involves a variety of complicated dynamical, thermal, radiative, and chemical processes on a wide range of scales \\citep[see][for a review]{mckee2007a}. Observed galaxies exhibit large-scale correlations between their global star formation rate (SFR) surface density $\\SigmaSFR$ and average gas surface density $\\Sigmagas$ \\citep{kennicutt1989a,kennicutt1998a}, and these global correlations serve as the basis for treatments of star formation in many models of galaxy formation. While such models have supplied important insights, detailed observations of galaxies have recently provided evidence that the molecular, rather than the total, gas surface density  is the primary driver of global star formation in galaxies \\citep[e.g.,][]{wong2002a,boissier2003a,heyer2004a,boissier2007a,calzetti2007a,kennicutt2007a}.  In this study, we adopt an approach in which empirical and theoretical knowledge of the star formation efficiency (SFE) in dense, molecular gas is used as the basis for a star formation model in hydrodynamical simulations of disk galaxy evolution.  This approach requires modeling processes that shape properties of the dense phase of the interstellar medium (ISM) in galaxies. The purpose of this paper is to present such a model.  Stellar populations in galaxies exhibit salient trends of colors and metallicities with galaxy luminosity \\citep[e.g.,][]{kauffmann2003c,blanton2005a,cooper2007a}. In the hierarchical structure formation scenario these trends should emerge through the processes of star formation and/or stellar feedback in the progenitors of present-day galaxies. Observationally, ample evidence suggests that the efficiency of the conversion of gas into stars depends strongly and non-monotonically on mass of the system. For example, the faint-end of the galaxy luminosity function has a shallow slope \\citep[$\\alphaL\\approx1.0-1.3$, e.g.,][]{blanton2001a,blanton2003a} compared to the steeper mass function of dark matter halos \\citep[$\\alphaDM\\approx2$, e.g.,][]{press1974a,sheth1999a}, indicating a decrease in SFE in low-mass galaxies.  At the same time, the neutral hydrogen (HI) and baryonic mass functions may be steeper than the luminosity function \\citep[$\\alphaHI\\approx1.3-1.5$, e.g.,][]{rosenberg2002a,zwaan2003a}. The baryonic \\cite{tully1977a} relation is continuous down to extremely low-mass dwarf galaxies \\citep[e.g.,][]{mcgaugh2005a,geha2006a}, indicating that the fractional baryonic content of galaxies of different mass is similar. Hence, low-mass galaxies that are unaffected by environmental processes are gas-rich, yet often form stars inefficiently.   While feedback processes from supernovae and AGN \\citep[e.g.,][]{brooks2007a,sijacki2007a}, or the efficiency of gas cooling and accretion \\citep{dekel2006a,dekel2008a}, may account for part of these trends, the SFE as a function of galaxy mass may also owe to intrinsic ISM processes \\citep[e.g.,][]{tassis2008a,kaufmann2007a}.  To adequately explore the latter possibility, a realistic model for the conversion of gas into stars in galaxies is needed.  Traditionally, star formation in numerical simulations of galaxy formation is based on the empirical Schmidt-Kennicutt (SK) relation \\citep{schmidt1959a,kennicutt1989a,kennicutt1998a}, in which star formation rate is a {\\it universal} power-law function of the total disk-averaged or global gas surface density: $\\SigmaSFR\\propto \\Sigmagas^{\\ntot}$ with $\\ntot\\approx 1.4$ describing the correlation for the entire population of normal and starburst galaxies.  However, growing observational evidence indicates that this relation may not be universal on smaller scales within galaxies, especially at low surface densities.  Estimates of the slope of the SK relation within individual galaxies exhibits significant variations. For example, while \\citet{schuster2007a} and \\citet{kennicutt2007a} find $\\ntot\\approx 1.4$ for the molecular-rich galaxy M51a, similar estimates in other large, nearby galaxies \\citep[including the Milky Way (MW);][]{misiriotis2006a} range from $\\ntot\\approx1.2$ to $\\ntot\\approx 3.5$ \\citep{wong2002a,boissier2003a} depending on dust-corrections and fitting methods. The disk-averaged total gas SK relation for normal (non-starburst) galaxies also has a comparably steep slope of $\\ntot\\approx2.4$, with significant scatter \\citep{kennicutt1998a}. While the variations in the $\\SigmaSFR-\\Sigmagas$ correlation may indicate systematic uncertainties in observational measurements, intrinsic variations or trends in galaxy properties may also induce differences between the global relation determined by \\cite{kennicutt1998a} and the $\\SigmaSFR-\\Sigmagas$ correlation in individual galaxies.  Galaxies with low gas surface densities, like dwarfs or bulgeless spirals, display an even wider variation in their star formation relations. \\citet{heyer2004a} and \\citet{boissier2003a} show that in low-mass galaxies the SFR dependence on the total gas surface density exhibits a power-law slope $\\ntot\\approx 2-3$ that is considerably steeper than the global \\citet{kennicutt1998a} relation slope of $\\ntot\\approx 1.4$. Further, star formation in the low-surface density outskirts of galaxies also may not be universal. Average SFRs appear to drop rapidly at gas surface densities of $\\Sigmagas\\lesssim 5-10{\\rm\\, \\Msun\\,pc^{-2}}$ \\citep{hunter1998a,martin2001a}, indicating that star formation may be truncated or exhibit a steep dependence on the gas surface density. The existence of such threshold surface densities have been proposed on theoretical grounds \\citep{kennicutt1998a,schaye2004a}, although recent GALEX results using a UV indicator of star formation suggest that star formation may continue at even lower surface densities \\citep{boissier2007a}. Star formation rates probed by damped Lyman alpha absorption (DLA) systems also appear to lie below the \\cite{kennicutt1998a} relation, by an order of magnitude \\citep{wolfe2006a,wild2007a}, which may indicate that the relation between SFR and gas surface density in DLA systems differs from the local relation measured at high $\\Sigmagas$.   In contrast, observations generally show that star formation in galaxies correlates strongly with {\\it molecular\\/} gas, especially with the highly dense gas traced by HCN emission \\citep{gao2004a,wu2005a}.  The power-law index of the SK relation connecting the SFR to the surface density of molecular hydrogen consistently displays a value of $\\nmol\\approx1.4$ and exhibits considerably less galaxy-to-galaxy variation \\citep{wong2002a,murgia2002a,boissier2003a,heyer2004a,matthews2005a,leroy2005,leroy2006,gardan2007a}. Molecular gas, in turn, is expected to form in the high-pressure regions of the ISM \\citep{elmegreen1993a,elmegreen1994a}, as indicated by observations \\citep{blitz2004a,blitz2006a,gardan2007a}.  Analytical models and numerical simulations that tie star formation to the fraction of gas in the dense ISM are successful in reproducing many observational trends \\citep[e.g.,][]{elmegreen2002a,kravtsov2003a,krumholz2005a,li2005a,li2006a,tasker2006a,tasker2007a,krumholz2007a,wada2007a,tassis2007a}. Recently, several studies have explored star formation recipes based on molecular hydrogen.  \\citet{pelupessy2006a} and \\citet{booth2007a} implemented models for $\\Hmol$ formation in gaseous disks and used them to study the molecular content and star formation in galaxies. However, these studies focused on the evolution of galaxies of a single mass and did not address the origin of the SK relation, its dependence on galaxy mass or structure, or its connection to trends in the local molecular fraction.  Our study examines the SK relation critically, including its dependence on the structural and ISM properties of galaxies of different masses, to explain the observed deviations from the global SK relation, to investigate other connections between star formation and disk galaxy properties such as rotation or gravitational instability, and to explore how the temperature and density structure of the ISM pertains to the star formation attributes of galaxies.  To these ends, we develop a model for the ISM and star formation whose key premise is that star formation on the scales of molecular clouds ($\\sim 10$~pc) is a function of molecular hydrogen density with a universal SFE per free-fall time \\citep[e.g.,][]{krumholz2007b}.  Molecular hydrogen, which we assume to be a proxy for dense, star forming gas, is accounted for by calculating the local $\\Hmol$ fraction of gas as a function of density, temperature, and metallicity using the photoionization code Cloudy \\citep{ferland1998a} to incorporate $\\Hmol$-destruction by the UV radiation of local young stellar populations.  We devise a numerical implementation of the star formation and ISM model, and perform hydrodynamical simulations to study the role of molecular gas in shaping global star formation relations of self-consistent galaxy models over a representative mass range.  The results of our study show that many of the observed global star formation correlations and trends can be understood in terms of the dependence of molecular hydrogen abundance on the local gas volume density. We show that the physics controlling the abundance of molecular hydrogen and its destruction by the interstellar radiation field (ISRF) play a key role in shaping these correlations, in agreement with earlier calculations based on more idealized models of the ISM \\citep{elmegreen1993a,elmegreen1994a}. While our simulations focus on the connection between the molecular ISM phase and star formation on galactic scales, the formation of molecular hydrogen has also been recently studied in simulations of the ISM on smaller, subgalactic scales \\citep{glover2007a,dobbs2007a}. These simulations are complementary to the calculations presented in our study and could be used as input to improve the molecular ISM model we present.   The paper is organized as follows.  The simulation methodology, including our numerical models for the ISM, interstellar radiation field, and simulated galaxies, is presented in \\S \\ref{section:methodology}.  The results of the simulations are presented in \\S \\ref{section:results}, where the simulated star formation relations in galactic disks and correlations of the molecular fraction with the structure of the ISM are examined.  We discuss our results in \\S \\ref{section:discussion} and conclude with a summary in \\S \\ref{section:summary}. Details of our tests of numerical fragmentation in disk simulations and calculations of the model scaling between star formation, gas density, and orbital frequency are presented in the Appendices. Throughout, we work in the context of a dark-energy dominated cold dark matter cosmology with a Hubble constant $H_{0} \\approx 70\\kms\\Mpc^{-1}$.   \\begin{figure*} \\figurenum{1} \\epsscale{1} \\plotone{fig1.eps} \\caption{\\label{fig:cooling} Cooling ($\\Lambda$) and heating ($\\Heating$) rates for interstellar and intergalactic gas as a function of gas density ($\\nH$), temperature ($T$), metallicity ($Z$), and interstellar radiation field (ISRF) strength ($\\Uisrf$), in units of the ISRF strength in the MW at the solar circle, as calculated by the code Cloudy \\citep{ferland1998a}.  Shown are the cooling (dotted lines), heating (dashed lines) and net cooling (solid) functions over the temperature range $T=10^{2}-10^{9}~\\K$.  Dense gas can efficiently cool via atomic and molecular coolants below $T=10^{4}\\K$, depending on the gas density and the strength of the ISRF.  A strong ISRF can enable the destruction of $\\Hmol$ gas and thereby reduce the SFR. } \\end{figure*} ",
        "conclusions": "\\label{section:discussion}  Results presented in the previous sections indicate that star formation prescriptions based on the local abundance of molecular hydrogen lead to interesting features of the global star formation relations.  We show that the inclusion of an interstellar radiation field is critical to control the amount of diffuse $\\Hmol$ at low gas densities.  For instance, without a dissociating ISRF the low mass dwarf galaxy eventually becomes almost fully molecular, in stark contrast with observations. We also show that without the dissociating effect of the ISRF our model galaxies produce a much flatter relation between molecular fraction $\\fHmol$ and pressure $\\Pext$, as can be expected from the results of \\cite{elmegreen1993a}.  Including $\\Hmol$-destruction by an ISRF results in a $\\fHmol-\\Pext$ relation in excellent agreement with the observations of \\citet{wong2002a} and \\citet{blitz2004a,blitz2006a}.  Our model also predicts that the relation between $\\Sigma_{\\rm SFR}$ and $\\Sigma_{\\rm gas}$ should not be universal and can be considerably steeper than the canonical value of $n_{\\rm tot}=1.4$, even if the three-dimensional Schmidt relation in molecular clouds is universal. The slope of the relation is controlled by the dependence of molecular fraction (i.e., fraction of star forming gas) on the total local gas surface density. This relation is non-trivial because the molecular fraction is controlled by pressure and ISRF strength, and can thus vary between different regions with the same total gas surface density. The relation can also be different in the regions where the disk scale-height changes rapidly (e.g., in flaring outer regions of disks), as can be seen from equation~\\ref{equation:structural_schmidt_law} \\citep[see also][]{schaye2007a}.   We show that the effect of radial variations in the molecular fraction $\\fHmol$ and gas scale heights ($\\hSFR$ and $\\hgas$) on the SFR can be accounted for in terms of a structural, SK-like correlation, $\\SigmaSFR\\propto\\fHmol \\hSFR \\hgas^{-1.5} \\Sigmagas^{1.5}$, that trivially relates the local SFR to the consumption of molecular gas with an efficiency that scales with the local dynamical time.  A generic testable prediction of our model is that deviations from the SK $\\SigmaSFR-\\Sigmagas$ relation are expected in galaxies or regions of galaxies where the molecular fraction is declining or much below unity.  As we discussed above, star formation in the molecule-poor ($\\fHmol\\sim0.1$) galaxy M33 supports this view as its total gas SK power-law index is $\\ntot\\approx 3.3$ \\citep{heyer2004a}.  While our simulations of an M33 analogue produce a steep SK power-law, the overall SFE in our model galaxies lies below that observed for M33.  However, as recently emphasized by \\cite{gardan2007a}, the highest surface density regions of M33 have unusually efficient star formation compared with the normalization of the \\cite{kennicutt1998a} relation and so the discrepancy is not surprising. Given that dwarf galaxies generally have low surface densities and are poor in molecular gas, it will be interesting to examine SK relation in other small-mass galaxies.  Another example of a low-molecular fraction galaxy close to home is M31, which has only a fraction $\\fHmol\\approx0.07$ of its gas in molecular form within 18 kpc of the galactic center \\citep{nieten2006a}. Our model would predict that this galaxy should deviate from the total gas SK relation found for molecular-rich galaxies. Observationally, the SK relation of M31 measured by \\citet{boissier2003a} is rather complicated and even has an increasing SFR with decreasing gas density over parts of the galaxy.  Low molecular fractions can also be expected in the outskirts of normal galaxies and in the disks of low surface brightness galaxies. The latter have molecular fractions of only $\\fHmol\\lesssim 0.10$ \\citep{matthews2005a,das2006a}, and we therefore predict that they will not follow the total gas SK relation obeyed by molecule-rich, higher surface density galaxies. At the same time, LSBs do lie on the same relation between $\\Hmol$ mass and far infrared luminosity as higher surface brightness (HSB) galaxies \\citep{matthews2005a}, which suggests that the dependence of star formation on molecular gas may be the same in both types of galaxies.  An alternative formulation of the global star formation relation is based on the angular frequency of disk rotation: $\\Sigma_{\\rm SFR}\\propto \\Sigmagas\\Omega$. That this relation works in real galaxies is not trivial, because star formation and dynamical time-scale depend on the local gas density, while $\\Omega$ depends on the total mass distribution {\\it within} a given radius. Although several models were proposed to explain such a correlation \\citep[see, e.g.,][for reviews]{kennicutt1998a,elmegreen2002a}, we show in \\S \\ref{section:results:sfr_rotation} and the second Appendix that the star formation correlation with $\\Omega$ can be understood as a fortuitous correlation of $\\Omega$ with gas density of $\\Omega\\propto \\rho_{\\rm gas}^{\\alpha}$, where $\\alpha\\approx 0.5$, for self-gravitating exponential disks or exponential disks embedded in realistic halo potentials. Moreover, we find that $\\SigmaSFR\\propto \\Sigmagas\\Omega$ breaks down at low values of $\\Sigmagas\\Omega$ where the molecular fraction declines, similarly to the steepening of the SK relation. Our models therefore predict that the $\\SigmaSFR\\propto\\SigmaHmol\\Omega$ relation is more robust than the $\\SigmaSFR\\propto\\Sigmagas\\Omega$ relation.  An important issue related to the global star formation in galaxies is the possible existence of star formation thresholds \\citep{kennicutt1998a,martin2001a}.  Such thresholds are expected to exist on theoretical grounds, because the formation of dense, star-forming gas is thought to be facilitated by either dynamical instabilities \\citep[see, e.g.,][for comprehensive reviews]{elmegreen2002a,mckee2007a} or gravithermal instabilities \\citep{schaye2004a}. We find that the two-component Toomre instability threshold that accounts for both stars and gas, $Q_{\\rm sg}<1$, works well in predicting the transition from atomic gas, inert to star formation, to the regions where molecular gas and star formation occur in our simulations.  Our results are in general agreement with \\citet{li2005a,li2005b,li2006a}, who used sink-particle simulations of the dense ISM in isolated galaxies to study the relation between star formation and the development of gravitational instabilities, and with observations of star formation in the LMC \\citep{yang2007a}.  Our simulations further demonstrate the importance of accounting for all mass components in the disk to predict correctly which regions galactic disks are gravitationally unstable.  Given that our star formation prescription is based on molecular hydrogen, the fact that $Q_{\\rm sg}$ is a good threshold indicator may imply that gravitational instabilities strongly influence the abundance of dense, molecular gas in the disk. Conversely, the gas at radii where the disk is stable remains at low density and has a low molecular fraction. We find that in our model galaxies, the shear instability criterion of \\cite{elmegreen1993a} does not work as well as the Toomre $Q_{\\rm sg}$-based criterion. Almost all of the star formation in our model galaxies occurs at surface densities $\\Sigmagas\\gtrsim 3{\\rm\\ M_{\\odot}\\,pc^{-2}}$, which is formally consistent with the \\citet{schaye2004a} constant surface density criterion for gravithermal instability. However, as Figure~\\ref{fig:kennicutt.gas} shows, we do not see a clear indication of threshold at a particular surface density and our GD-SF models that have an effectively isothermal ISM with $T\\approx10^{4}\\K$ (and hence do not have a gravithermal instability) still show a good correlation between regions where $\\Qsg<1$ and regions where star formation operates.   Our results have several interesting implications for interpretation of galaxy observations at different epochs. First, low molecular fractions in dwarf galaxies mean that only a small fraction of gas is participating in star formation at any given time. This connection between SFE and molecular hydrogen abundance may explain why dwarf galaxies are still gas rich today compared to larger mass galaxies \\citep{geha2006a}, without relying on mediation of star formation or gas blowout by supernovae. Note that a similar reasoning may also explain why large LSBs at low redshifts are gas rich but anemic in their star formation. Understanding the star formation and evolution of dwarf galaxies is critical because they serve as the building blocks of larger galaxies at high redshifts. Such small-mass galaxies are also expected to be the first objects to form large masses of stars and should therefore play an important role in enrichment of primordial gas and the cosmic star formation rate at high redshifts \\citep{hopkins2004a,hopkins2006am}.  The star-forming disks at $z\\sim2$ that may be progenitors of low-redshift spiral galaxies are observed to lack centrally-concentrated bulge components \\citep{elmegreen2007a}.  Given that galaxies are expected to undergo frequent mergers at $z>2$, bulges should have formed if a significant fraction of baryons are converted into stars during such mergers \\citep[e.g.,][]{gnedin2000a}. The absence of the bulge may indicate that star formation in the gas rich progenitors of these $z\\sim 2$ systems was too slow to convert a significant fraction of gas into stars. This low SFE can be understood if the high cosmic UV background, low-metallicities, and low dust content of high-$z$ gas disks keep their molecular fractions low \\citep{pelupessy2006a}, thereby inhibiting star formation over most of gas mass and keeping the progenitors of the star-forming $z\\sim 2$ disks mostly gaseous.  Gas-rich progenitors may also help explain the prevalence of extended disks in low-redshift galaxies despite the violent early merger histories characteristic of $\\Lambda$CDM universes, as gas-rich mergers can help build high-angular momentum disk galaxies \\citep[][]{robertson2006c}. Mergers of mostly stellar disks, on the other hand, would form spheroidal systems.   Our results may also provide insight into the interpretation of the results of \\cite{wolfe2006a}, who find that the SFR associated with neutral atomic gas in damped Lyman alpha (DLA) systems is an order of magnitude lower than predicted by the local \\citet{kennicutt1998a} relation.  The DLAs in their study sample regions with column densities $N_{\\mathrm{H}}\\approx 2\\times 10^{21}\\rm\\ cm^{-2}$, or surface gas densities of $\\Sigma_{\\rm DLA}\\approx 20\\rm\\ M_{\\odot}\\,pc^{-2}$, assuming a gas disk with a thickness of $h\\approx100$~pc. Suppose the local SK relation steepens from the local relation with the slope $n_0=1.4$ to a steeper slope $n_1$ below some surface density $\\Sigma_{\\rm b}>\\Sigma_{\\rm DLA}$. Then for $\\Sigmagas<\\Sigma_{\\rm b}$, the SFR density will be lower than predicted by the local relation by a factor of $(\\Sigmagas/\\Sigma_{\\rm b})^{n_1-n_0}$. For $n_0=1.4$ and $n_1=3$ the SFR will be suppressed by a factor of $>10$ for $\\Sigmagas/\\Sigma_{\\rm b}<0.25$. Thus, the results \\cite{wolfe2006a} can be explained if the total gas SK relation at $z\\sim3$ steepens below $\\Sigmagas\\lesssim 100\\rm\\ M_{\\odot}\\, pc^{-2}$. We suggest that if the majority of the molecular hydrogen at these redshifts resides in rare, compact, and dense systems \\citep[e.g.,][]{zwaan2006a}, then both the lack of star formation and the rarity of molecular hydrogen in damped Ly$\\alpha$ absorbers may be explained simultaneously.  Our results also indicate that the thermodynamics of the ISM can leave an important imprint on its density probability distribution. Each thermal phase in our model galaxies has its own log-normal density distribution.   Our results thus imply that using a single lognormal PDF to build a model of global star formation in galaxies \\citep[e.g.,][]{elmegreen2002a,wada2007a} is likely an oversimplification. Instead, the global star formation relation may vary depending on the dynamical and thermodynamical properties of the ISM.  We can thus expect differences in the SFE between the low-density and low-metallicity environments of dwarf and high-redshift galaxies and the higher-metallicity, denser gas of many large nearby spirals.  Note that many of the results and effects we discuss above may not be reproduced with a simple 3D density threshold for star formation, as commonly implemented in galaxy formation simulations.  Such a threshold can reproduce the atomic-to-molecular transition only crudely and would not include effects of the local interstellar radiation field, metallicity and dust content, etc.  A clear caveat for our work is that the simulation resolution limits the densities we can model correctly. At high densities, the gas in our simulations is over-pressurized to avoid numerical Jeans instability. The equilibrium density and temperature structure of the ISM and the molecular fraction are therefore not correct in detail. Note, however, that our pressurization prescription is designed to scale with the resolution, and should converge to the ``correct'' result as the resolution improves. In any event, the simulations likely do not include all the relevant physics shaping density and temperature PDFs of the ISM in real galaxies. The results may of course depend on other microphysics of the ISM as they influence both the temperature PDF and the fraction of gas in a high-density, molecular form. Future simulations of the molecular ISM may need to account for new microphysics as they resolve scales where such processes become important.     Using hydrodynamical simulations of the ISM and star formation in cosmologically motivated disk galaxies over a range of representative masses, we examine the connection between molecular hydrogen abundance and destruction, observed star formation relations, and the thermodynamical structure of the interstellar medium.  Our simulations provide a variety of new insights into the mass dependence of star formation efficiency in galaxies. A summary of our methodology and results follows.   \\begin{itemize} \\item[1.]  A model of heating and cooling processes in the interstellar medium (ISM), including low-temperature coolants, dust heating and cooling processes, and heating by the cosmic UV background, cosmic rays, and the local interstellar radiation field (ISRF), is calculated using the photoionization code Cloudy \\citep{ferland1998a}. Calculating the molecular fraction of the ISM enables us to implement a prescription for the star formation rate (SFR) that ties the SFR directly to the molecular gas density.  The ISM and star formation model is implemented in the SPH/N-body code GADGET2 \\citep{springel2005c} and used to simulate the evolution of isolated disk galaxies.  \\item[2.]  We study the correlations between gas surface density ($\\Sigmagas$), molecular gas surface density ($\\SigmaHmol$), and SFR surface density ($\\SigmaSFR$).  We find that in our most realistic model that includes heating and destruction of $\\Hmol$ by the interstellar radiation field, the power law index of the SK relation, $\\SigmaSFR\\propto\\Sigmagas^{\\ntot}$, (measured in annuli) varies from $\\ntot\\sim2$ in massive galaxies to $\\ntot\\gtrsim4$ in small mass dwarfs.  The corresponding slope of the $\\SigmaSFR\\propto\\SigmaHmol^{\\nmol}$ molecular-gas Schmidt-Kennicutt relation is approximately the same for all galaxies, with $\\nmol\\approx 1.3$.  These results are consistent with observations of star formation in different galaxies \\citep[e.g.,][]{kennicutt1998a,wong2002a,boissier2003a,heyer2004a,boissier2007a,kennicutt2007a}.  \\item[3.] In our models, the SFR density scales as $\\SigmaSFR \\propto \\fHmol h_{\\SFR} h_{\\gas}^{-1.5} \\Sigmagas^{1.5}$, where $h_{\\gas}$ is the scale-height of the ISM and $h_{\\SFR}$ is the scale-height of star-forming gas. The different $\\SigmaSFR-\\Sigmagas$ relations in galaxies of different mass and in regions of different surface density in our models therefore owe to the dependence of molecular fraction $\\fHmol$ and scale height of gas on the gas surface density.  \\item[4.] We show that the $\\SigmaSFR\\propto\\Sigmagas\\Omega$ and $\\SigmaSFR\\propto\\SigmaHmol\\Omega$ correlations describe the simulations results well where the molecular gas and total gas densities are comparable, while the simulations deviate from $\\SigmaSFR\\propto\\Sigmagas\\Omega$ \\citep[e.g.,][]{kennicutt1998a} at low $\\Sigmag$ owing to a declining molecular fraction. We demonstrate that these relations may owe to the fact that the angular frequency and the disk-plane gas density are generally related as $\\Omega\\propto\\sqrt{\\rho}$ for exponential disks if the potential is dominated by either the disk, a \\cite{navarro1996a} halo, \\cite{hernquist1990a} halo, or an isothermal sphere. The correlation of $\\SigmaSFR$ with $\\Omega$ is thus a secondary correlation in the sense that  $\\Omega\\propto\\sqrt{\\rho}$ is set during galaxy formation and $\\Omega$ does not directly influence star formation.  \\item[5.] The role of critical surface densities for shear instabilities ($\\SigmaA$) and \\cite{toomre1964a} instabilities ($\\SigmaQ$) in star formation \\citep[e.g.,][]{martin2001a} is examined in the context of the presented simulations.  We find that the two-component Toomre instability criterion $\\Qsg<1$ is an accurate indicator of the star-forming regions of disks, and that gravitational instability and star formation are closely related in our simulations. Further, the $\\Qsg$ criterion works even in simulations in which cooling is restricted to $T>10^4$~K where gravithermal instability cannot operate.  \\item[6.] Our simulations that include $\\Hmol$-destruction by an ISRF naturally reproduce the observed scaling $\\fHmol\\propto\\Pext^{0.9}$ between molecular fraction and external pressure \\citep[e.g.,][]{wong2002a,blitz2004a,blitz2006a}, but we find that simulations without an ISRF have a weaker scaling $\\fHmol\\propto \\Pext^{0.4}$.  We calculate how the connection between the scalings of the gas surface density, the stellar surface density, and the ISRF strength influence the $\\fHmol-\\Pext$ relation in the ISM, and show how the simulated scalings reproduce the $\\fHmol-\\Pext$ relation even as the power-law index of the total gas Schmidt-Kennicutt relation varies dramatically from galaxy to galaxy.  \\item[7.] We present a method for mitigating numerical Jeans fragmentation in Smoothed Particle Hydrodynamics simulations that uses a density-dependent pressurization of gas on small scales to ensure that the Jeans mass is properly resolved, similar to techniques used in grid-based simulations \\citep[e.g.,][]{truelove1997a,machacek2001a}.  The gas internal energy $u$ at the Jeans scale is scaled as $u\\propto\\mJeans^{-2/3}$, where $\\mJeans$ is the local Jeans mass, to ensure the Jeans mass is resolved by some $\\NJeans$ number of SPH kernel masses $2\\Nneigh m_{\\mathrm{SPH}}$, where $\\Nneigh$ is the number of SPH neighbor particles and $m_{\\mathrm{SPH}}$ is the gas particle mass. For the simulations presented here, we find the \\cite{bate1997a} criterion of $\\NJeans=1$ to be insufficient to avoid numerical fragmentation and that $\\NJeans\\sim15$ provides sufficient stability against numerical fragmentation over the time evolution of our simulations. Other simulations may have more stringent resolution requirements \\citep[e.g.,][]{commercon2008a}. We also demonstrate that isothermal galactic disks with temperatures of $T=10^{4}\\K$ may be susceptible to numerical Jeans instabilities at resolutions common in cosmological simulations of disk galaxy formation, and connect this numerical effect to possible angular momentum deficiencies in cosmologically simulated disk galaxies. \\end{itemize}  The results of our study indicate that star formation may deviate significantly from the relations commonly assumed in models of galaxy formation in some regimes and that these deviations can be important for the overall galaxy evolution.  Our findings provide strong motivation for exploring the consequences of such deviations and for developing further improvements in the treatment of star formation in galaxy formation simulations.\\\\[2mm]"
    },
    "0710/0710.2428_arXiv.txt": {
        "abstract": "{ A recent analysis of cosmic-ray data from a space borne experiment by the AMS collaboration supports the observation of an excess in the cosmic-ray positron spectrum by previous balloon experiments. The combination of the various experimental data establishes a deviation from the expected background with a significance of more than four standard deviations. The observed change in the spectral index cannot be explained without introducing a new source of positrons. When interpreted within the MSSM a consistent description of the antiproton spectrum, the diffuse gamma-ray flux and the positron fraction is obtained which is compatible with all other experimental data, including recent WMAP data. \\PACS{ {98.70.Sa}{Cosmic rays} \\and {95.35.+d}{Dark Matter} \\and {11.30.Pb}{Supersymmetry} } % } % ",
        "introduction": "Among the cosmic-ray species, antiparticles and diffuse $\\gamma$-rays are of particular interest because they are produced secondarily in hadronic interactions of protons and nuclei with the interstellar medium at low rates. Their small abundance makes them a sensitive probe for the existence of additional -- and possibly exotic -- cosmic-ray sources which would be visible as an excess of particles above conventional expectations.  One of the most important unsolved questions in modern cosmology is the nature of dark matter. The most promising dark matter candidate is the weakly interacting lightest neutralino, $\\chi_1^0$, predicted by supersymmetric extensions to the standard model of particle physics. The annihilation of neutralinos might constitute an additional primary source of particles with a unique spectral shape which would be determined by the parameters of supersymmetry, allowing to put constraints on new physics beyond the standard model.  A recent reanalysis of the data from the \\AMS{} spectrometer~\\cite{aguilar07a} supports the observation of an excess of cosmic-ray positrons by the HEAT experiments~\\cite{beatty04a}. In this work, we discuss the combined results on the cosmic-ray positron fraction $e^+ / (e^+ + e^-)$. Assuming that dark matter is largely constituted by neutralinos, we determine the cosmic-ray preferred parameter space of the minimal supersymmetric standard model (MSSM) from a simultaneous fit to the cosmic-ray positron, antiproton and diffuse $\\gamma$-ray data. ",
        "conclusions": "\\label{sec:conclusions} In this work, the combined recent experimental results on the cosmic-ray positron fraction have been presented. The data exhibit an excess of positrons above energies of 6\\+\\GeV{} which cannot be explained by purely secondary positron production alone and thus requires an additional primary source of positrons. In this work, we interpret this source to be the annihilation of supersymmetric neutralinos constituting dark matter. A simultaneous fit to the cosmic-ray positron, antiproton and $\\gamma$-ray data shows that, for particular sets of the MSSM parameters, this hypothesis gives a fully consistent description of the cosmic-ray spectra which is compatible with all other experimental data. We find that the cosmic-ray data clearly prefer the focus point region of the MSSM parameter space but reveal almost no sensitivity to $\\tan\\beta$."
    },
    "0710/0710.5452_arXiv.txt": {
        "abstract": "We use a series of ray-tracing experiments to determine the magnification distribution of high-redshift sources by gravitational lensing. We determine empirically the relation between magnification and redshift, for various cosmological models. We then use this relation to estimate the effect of lensing on the determination of the cosmological parameters from observations of high-$z$ supernovae. We found that, for supernovae at redshifts $z<1.8$, the effect of lensing is negligible compared to the intrinsic uncertainty in the measurements. Using mock data in the range $1.8<z<8$, we show that the effect of lensing can become significant. Hence, if a population of very-high-$z$ supernovae was ever discovered, it would be crucial to fully understand the effect of lensing, before these SNe could be used to constrain cosmological models. We show that the distance moduli $m-M$ for an open CDM universe and a $\\Lambda$CDM universe are comparable at $z>2$. Therefore if supernovae up to these redshifts were ever discovered, it is still the ones in the range $0.3<z<1$ that would distinguish these two models. ",
        "introduction": "High-redshift supernovae have become a major tool in modern cosmology. By measuring their apparent magnitudes, we can estimate their luminosity distances $d_L$ (see \\citealt{tonryetal03,barrisetal04,riessetal04}, and references therein). Since the relationship between $d_L$ and the redshift $z$ depends on the cosmological parameters, observations of distant SNe can constrain the cosmological model. Prior to the announcement of the {\\sl WMAP} results \\citep{bennettetal03}, observations of high-$z$ SNe provided the most compelling evidence of the existence of a nonzero cosmological constant. Since then, they have been used in combination with the {\\sl WMAP} data to refine the determination of the cosmological parameters.  The luminosity distances $d_L$ are determined by combining the observed fluxes $F$ with estimates of the SNe luminosities $L$. Uncertainties in $d_L$ are caused by uncertainties in $L$, because SNe are not perfect standard candles. The flux $F$ is much easier to measure, but for distant sources the value of $F$ might be altered by gravitational lensing caused by the intervening distribution of matter. For instance, a magnification $\\mu>1$ would result in a increase in $F$, and an underestimation of $d_L$.  Estimating the effect of lensing on the statistics of high-$z$ supernovae is a complex problem. Using either an analytical model or ray-tracing simulations, we can estimate the effect of lensing of a large number of sources in a statistical sense. We would then need to redo the error analysis on the SNe data to include in a consistent way the effect of lensing. This would be a very complex task, and in this paper we have chosen a much simpler approach. {\\it Our goal is not to obtain a precise estimate of the error introduced by lensing, but rather to assess the importance of this effect: is it dominant, important, or negligible, and for what range of redshift? and how does it affect the discrimination between different cosmological models?\\/} To answer these questions, we take at face value the published results of Type~Ia SNe, including their error bars which account for every source of uncertainty but gravitational lensing. Then, we include {\\it a posteriori\\/} the effect of lensing, to estimate the change in the errors. This approach is not rigorous at all, and does not constitute a substitute for a rigorous treatment of the errors. But it has the great advantage of simplicity. We do not have to redo the detailed error analysis performed by the high-redshift SNe groups, and, more importantly, our conclusions will not be tied to any particular sample or particular data reduction and error analysis technique used by any particular group. We are seeking to make generic statements about the importance of lensing (or lack of) that are relevant to any current or future sample of high-$z$ SNe.  The lensing of distant supernovae has been the focus of several recent studies. In an early study, \\citet{wambsganssetal97} used ray-tracing experiments to estimate the effect of weak lensing on the determination of the deceleration parameter $q_0$. \\citet{md05}, \\citet{dv05}, and \\citet{mv06} focused on SNe as a mean to study the nature of weak lensing. The issue of determining the cosmological parameters for distant SNe, and how this determination is affected by lensing, was addressed by \\citet{wang05} who used semi-analytical models to determine the magnification distribution function, \\citet{hl05} who used Monte Carlo ray-tracing simulations to study the effect of weak and strong lensing, and \\citet{gunnarssonetal06} and \\citet{jonssonetal06}, who estimated the effect on lensing along individual lines of sight by considering the properties of foreground galaxies in the same direction. These various studies concluded that the effect of lensing on current determinations of the cosmological parameters is small. \\citet{alderingetal06} discussed the effect of gravitational lensing on a population of SNe at $z>1.7$.  What distinguishes our approach is mostly its simplicity. Our calculations depend on very few assumptions, and this implies a certain amount of robustness to our results. Even though we rely on numerical simulations, this work should be regarded as a back-of-the-envelope calculation, whose purpose is to obtain a qualitative estimate of the effect of lensing on the determination of cosmological parameters by distant SNe. Using ray-tracing experiments, rather than a semi-analytical approach, enables us to extend our study to redshifts much higher than the ones considered by \\citet{wang05} and \\citet{hl05}.  This paper is organized as follow: In \\S2, we describe our calculation of the magnification distribution $P(\\mu)$, and how to estimate that distribution at any redshift $z$. In \\S3, we describe the real and mock samples of supernovae we use for our calculations. Results are presented in \\S4. In \\S5, we address various observational issues. Summary and conclusion are presented in \\S6. ",
        "conclusions": "We have performed a series of ray-tracing experiments using a multiple lens-plane algorithm. We have determined the distributions of  magnifications $P(\\mu)$ for sources in the redshift range $0<z<8$, for three different cosmological models. We have used these distributions to estimate the effect of gravitational lensing on the determination of the cosmological parameters with high-redshift Type~Ia supernovae. We used a generic, {\\it a posteriori} approach which is not tied to any particular sample.  We found that errors introduced by lensing are unimportant for SNe with redshift $z<1.8$. These errors are negligible compared to the intrinsic errors already present in the supernovae data. Since those intrinsic errors do not prevent us from determining the cosmological parameters, the additional errors introduced by lensing have no consequences. A similar conclusion was reached by \\citet{alderingetal06}.  Using a mock catalog of high-$z$ SNe, extending to $z=8.1$, we showed that the effect of lensing on a hypothetical population of SNe at redshifts $z>2$ could be quite significant, and must be understood before such SNe could be used to constrain cosmological models. Furthermore, the open CDM and $\\Lambda$CDM are difficult to distinguish at that redshift. We showed that, even if SNe at redshift $z\\sim8$ were ever discovered, it is the SNe in the range $z=0.3-1$ that would still provide the best discriminant between these two models. The data at that redshift already exist, and they support the $\\Lambda$CDM model."
    },
    "0710/0710.3369_arXiv.txt": {
        "abstract": " ",
        "introduction": "It is thought that the decay of the magnetic field in \\emph{magnetars} is the main source of its X-ray and $\\gamma$-ray luminosity since these objects appear to radiate substantially more power than available from the energy rotational loss \\cite{TD-96}. The spontaneous decay of the magnetic field could occur through \\emph{ambipolar diffusion}, \\emph{Hall drift} and \\emph{ohmic diffusion}, which are non-ideal magnetohydrodynamics processes that occur in thousands of years, compared to dynamical sound and Alfv\\'en time scales of milliseconds to seconds. Ambipolar diffusion promote a dissipative magnetic field advection, through the movement of a bulk of charged particles relative to the neutrons, the \\emph{Hall drift} is a non-dissipative advection of the magnetic field caused by the electrical current associated with it, and the magnetic field ohmic diffusion is a dissipative process caused by the electrical resistivity. The time scales of these process were estimated in Ref.~2. For classical pulsar magnetic field strengths, these were found to be longer than the lifetimes of these stars, and therefore unlikely to be observationally important. For the case of magnetars, it was found that magnetic field decay by these processes might be occurring \\cite{TD-96,ACT-04}. However, a full understanding of these processes, their interactions, and their effectiveness in neutron stars is still lacking. The work of \\cite{GR-92} was analytical, and therefore useful to identify general processes and identify the relevant time scales, but not to address the action of the identified processes in their full nonlinear development, and their interaction with each other. The full evolution of the magnetic field can only be addressed by numerical simulations. Recent ideal three-dimensional single-fluid magnetohydrodynamics simulations showed that, in a stably stratified star, a complicated, random, initial field generally evolves on a short, Alfv\\'en-like time scale to a relatively simple, roughly axisymmetric, large-scale configuration containing a toroidal and a poloidal component of comparable strength, both of which are required in order to stabilize each other \\cite{BS-04}. It is interesting to study the effects of the non-ideal processes studied in Ref.~2 on the evolution of this configuration. As a first step in this direction, here we simulate the decay of a magnetic field, including the effects studied in  Ref.~2, in a system where the magnetic field points in one Cartesian direction but varies only along an orthogonal direction. It is shown that the magnetic field evolves through different quasi-equilibrium states and we estimate the characteristic time scales where these quasi-equilibria occur. ",
        "conclusions": "Using numerical simulations in one dimension, we studied the effects of some non-ideal MHD processes on the magnetic field evolution in neutron stars. We found that the system evolves through succesive quasi-equilibria, and we estimated the characteristic time scales on which these quasi-equilibria occur.   \\begin{center} Acknowledgments \\end{center} We acknowledge financial support through FONDECYT postdoctoral project 3060103, Gemini project 32070014, and regular FONDECYT projects 1060644 and 1070854. We also thank the ESO-Chile Mixed Committee."
    },
    "0710/0710.4549_arXiv.txt": {
        "abstract": "We investigate the chemical abundances of NGC\\,3603 in the Milky Way, of 30\\,Doradus in the Large Magellanic Cloud, and of N\\,66 in the Small Magellanic Cloud. Mid-infrared observations with the Infrared Spectrograph onboard the Spitzer Space Telescope allow us to probe the properties of distinct physical regions within each object: the central ionizing cluster, the surrounding ionized gas, photodissociation regions, and buried stellar clusters. We detect [S\\3], [S\\4], [Ar\\3], [Ne\\2], [Ne\\3], [Fe\\2], and [Fe\\3] lines and derive the ionic abundances. Based on the ionic abundance ratio (Ne\\3/H)/(S\\3/H), we find that the gas observed in the MIR is characterized by a higher degree of ionization than the gas observed in the optical spectra. We compute the elemental abundances of Ne, S, Ar, and Fe. We find that the $\\alpha$-elements Ne, S, and Ar scale with each other. Our determinations agree well with the abundances derived from the optical. The Ne/S ratio is higher than the solar value in the three giant H\\2\\ regions and points toward a moderate depletion of sulfur on dust grains. We find that the neon and sulfur abundances display a remarkably small dispersion (0.11\\,dex in 15 positions in 30\\,Doradus), suggesting a relatively homogeneous ISM, even though small-scale mixing cannot be ruled out. ",
        "introduction": "\\label{sec:intro}  Giant H\\2\\ regions are ideal laboratories to understand the feedback of star-formation on the dynamics and energetics of the interstellar medium (ISM). Supernov{\\ae} and stellar winds arising in such regions are reponsible for producing shocks, destroying dust grains and molecules, while compressing molecular clouds and triggering subsequent star-formation. They also allow the release of newly synthetized elements into the ISM, altering its metallicity.  In order to study the star-formation properties as a function of the environment, we observed three giant H\\2\\ regions spanning a wide range of physical conditions (gas density, mass, age) and chemical properties (metallicity) with the Spitzer Space Telescope (Werner et al.\\ 2004). Observations are part of the GTO program PID\\#63. The regions are NGC\\,3603 in the Milky Way, 30\\,Doradus (hereafter 30\\,Dor) in the Large Magellanic Cloud (LMC), and N\\,66 in the Small Magellanic Cloud (SMC). The scope of this program is to address crucial issues such as the destruction of complex molecules by energetic photons arising from massive stars, the polycyclic aromatic hydrocarbon (PAH) abundance dependence on metallicity, or conditions that lead to the formation/disruption of massive stellar clusters. Photometry with Spitzer/IRAC (Fazio et al.\\ 2004) has been performed and will be discussed in Brandl et al.\\ (in preparation). The brightest mid-infrared (MIR) regions (knots, stellar clusters, shockfronts, ...) were followed spectroscopically with the Infrared Spectrograph (IRS; Houck et al.\\ 2004). In Lebouteiller et al.\\ (2007), we analyzed the spatial variations of the PAH and fine-structure line emission across individual photodissociation regions (PDRs) in NGC\\,3603. The two other regions will be investigated the same way in follow-up papers (Bernard-Salas et al.\\ in preparation; Whelan et al.\\ in preparation). In this paper, we introduce the full IRS dataset (low- and high-resolution) of the giant H\\2\\ regions and we derive their chemical abundances. A subsequent paper will be focused on the study of molecules and dust properties (Lebouteiller et al.\\ in preparation).  Elemental abundances in H\\2\\ regions are historically derived from optical emission-lines. Large optical telescopes, together with sensitive detectors makes it possible to determine the chemical composition of very faint H\\2\\ regions. Because of dust extinction, optical spectra only observe ionized gas toward sighlines with low dust content. In this view, the MIR range allows analyzing denser lines of sight, with possibly different chemical properties because of small-scale mixing and/or differential depletion on dust grains. MIR emission-lines constitute the only way to measure abundances in more obscure regions, and these abundances ought to be compared to abundances from the optical range. Although the optical domain gives access to some of the most important elements to constrain nucleosynthethic and stellar yields (C, N, O, Ne, S, Ar, Fe), it does not include some essential ionization stages necessary for abundance determinations of certain elements, such as S\\4\\ or Ne\\2. The MIR range enables the abundance determination of Ne, S, and Ar, with the most important ionization stages observed. Iron abundance can be also determined from MIR forbidden emission-lines, but with considerably larger uncertainty due to ionization corrections. Finally, it must be stressed that abundance determinations in the optical are more sensitive to the electronic temperature ($T_e$) determination as compared to the MIR range. The effect of $T_e$ on abundances determinations is a significant source of error in optical abundance results.  Wu et al.\\ (2007) recently studied a sample of blue compact dwarf galaxies (BCDs) with the IRS and found a global agreement between abundances derived from the optical and those derived from the MIR. This suggests that the dense lines of sight probed in the MIR have a similar chemical composition as unextincted lines of sight and/or dense regions with possibly peculiar abundances do not contribute significantly to the integrated MIR emission-line spectrum. MIR abundances of the BCDs were calculated using mostly H$\\beta$ or H$\\alpha$ lines from the optical as tracer of the hydrogen content, with significant uncertainties from aperture corrections, or different observed regions because of extinction. The present sample of giant H\\2\\ regions provides the unique opportunity to measure accurate abundances, with a signal-to-noise ratio sufficiently high to observe directly the H\\1\\ recombination line at 12.37\\mic. We provide abundances of Ne, S, Ar, and Fe toward lines of sight with different physical properties (PDRs, ionized gas, embedded source, stellar cluster, ...) within each giant H\\2\\ region.  We first present the sample of the three giant H\\2\\ regions in \\S\\ref{sec:pres}. The data reduction and analysis are discussed in \\S\\ref{sec:observations}. We infer the ion abundances in \\S\\ref{sec:ionicab}. Elemental abundances are determined in \\S\\ref{sec:eleab} and are discussed in \\S\\ref{sec:discussion}. ",
        "conclusions": "We analyzed the chemical abundances in the ISM of three giant H\\2\\ regions, NGC\\,3603 in the Milky Way, 30\\,Dor in the LMC, and N\\,66 in the SMC using the MIR lines observed with the IRS onboard Spitzer. \\begin{itemize}  \\item Our observations probe the ISM toward various physical regions, such as stellar clusters, ionized gas, photodissociation regions, and deeply embedded MIR bright sources. The spectra show the main ionization stages of neon, sulfur, and argon in the ionized gas. We also detect [Fe\\2] and [Fe\\3] lines.  \\item Ionic abundances of Ne\\2, Ne\\3, S\\3, S\\4, Ar\\2, Ar\\3, Fe\\2, and Fe\\3\\ were derived. The internal variation of electron density across a region has no impact on the ionic abundance determination. On the other hand, we find that electron temperature uncertainties and/or intrinsic variations could be responsible for an error of 20\\% at most on the abundance determinations. Based on the (Ne\\3/H)/(S\\3/H) ionic abundance ratio, we find that the optical spectra probe a gas with a degree of ionization equal to or higher than the gas probed in the MIR.  \\item Elemental abundances were determined from the ionic abundances. No ionization corrections were needed, except for iron. We find that neon, sulfur, and argon scale with each other, which is expected from stellar yields. Abundances do not show any dependence on the physical region (PDR, stellar cluster, embedded region, ...).  \\item The Ne/S ratio is larger than the solar value, and suggests that sulfur could be depleted onto dust grains. The sulfur abundance in the MIR agrees best with the lowest optical determinations, which is likely due to uncertainties in the abundance determinations.  \\item Iron abundance shows a larger uncertainty than Ne/H, S/H, and Ar/H. The comparison of iron and neon abundances hints at significant depletion of iron onto dust grains at large metallicities. The agreement with the optical determination of Fe/H indicates however that there is no differential depletion on dust grains between the gas probed in the MIR and in the optical.  \\item  Fe/H is found to be spectacularly large in one position, corresponding to a supernova remnant. This strongly suggest that iron atoms have been released from dust grains due to schocks from the SN.  \\item The metallicity of NGC\\,3603 agrees with the Galactic abundance gradient. The metallicities of 30\\,Doradus and N\\,66 agree well with those of the PNe in their respective host galaxies. These findings suggest that the giant H\\2\\ regions did not experience a significant metal enrichment for at least 1\\,Gyr. If enrichement occured, the metallicity was altered by less than a factor of two.  \\item Neon and sulfur abundances show remarkably little dispersion in the three H\\2\\ regions (e.g., 0.11\\,dex dispersion in 15 positions in 30\\,Dor). Small-scale mixing is apparently effective, abundance fluctuations are smaller than $\\sim$55\\%. However, internal variations of the abundances are likely to be on the order of $\\lesssim$5\\%, and determining their existence would require a significant improvement of the data quality and of the method to be evidenced. \\end{itemize}"
    },
    "0710/0710.1193_arXiv.txt": {
        "abstract": "\\small Small perturbations in spherical and thin disk stellar clusters surrounding massive a black hole are studied. Due to the black hole, stars with sufficiently low angular momentum escape from the system through the loss cone. We show that stability properties of spherical clusters crucially depend on whether the distribution of stars is monotonic or non-monotonic in angular momentum. It turns out that only non-monotonic distributions can be unstable. At the same time the instability in disk clusters is possible for both types of distributions. ",
        "introduction": "The study of the gravitational loss-cone instability, a far analog of the plasma cone instability, has begun with the work of V. Polyachenko (1991),  in which a simplest analytical model of thin disk stellar cluster has been treated. The interest to the problem of stability of stellar clusters has been revived recently by detailed investigation by Tremaine (2005) and Polyachenko, Polyachenko, Shukhman, (2007; henceforth, Paper I) of low mass clusters around massive black holes. The both papers have considered stability of small amplitude perturbations of stellar clusters of disk-like and spherical geometry.  Tremaine (2005) has shown using Goodman's (1988) criterion that thin disks with symmetric DFs over angular momentum and empty loss cone are generally unstable. By contrast, analyzing perturbations with spherical numbers $l=1$ and $l=2$, he deduced that spherical clusters with monotonically increasing DF of angular momentum should be generally stable.  Later we demonstrated (see Paper I) that spherical systems with non-monotonic distributions may be unstable for sufficiently small-scale perturbations $l \\ge 3 $, while the harmonics $l=1,2$ are always stable. For the sake of convenience, we have used two assumptions. The first one is that the Keplerian potential of the massive black hole dominates over a self-gravitating potential of the stellar cluster (which does not mean that one can neglect the latter). Then the characteristic time of system evolution is of the order of the orbit precessing time, which is slow, compared to typical dynamical (free fall) time. Since a star makes many revolutions in its almost unaltered orbit, we can regard it as to be ``smeared out'' along the orbit in accordance with passing time, and study evolution of systems made of these extended objects.   The second assumption is a so called {\\it spoke approximation}, in which a system consists of near-radial orbits only. This approximation was earlier suggested by one of the authors (Polyachenko 1989, 1991). The spoke approximation reduces the problem to a study of rather simple analytical characteristic equations controlling small perturbations of stellar clusters.   There are two questions that naturally arise in this context. First: Does the instability remain when abandoning the assumption of strong radial elongation of orbits? Second: Does the instability occur in spheres with monotonically increasing distributions in angular momentum if one consider smaller-scale perturbations with $l \\ge 3$? The aim of the paper is to provide answers to these questions.  To achieve the task we use semi-analytical approach based on analysis of integral equations for slow modes elaborated recently in Polyachenko (2004, 2005) for thin disks, and in Paper I for spherical geometry. Following Paper I, we shall restrict ourselves to studying monoenergetic models with DFs in the form \\begin{align}\\label{eq:1.1} F(E,L)=A\\,\\delta(E-E_0)\\,f(L). \\end{align} The models specified by function $f(L)$ are suitable for studying the effects of angular momentum distribution on gravitational loss-cone instability. On the other hand, the Dirac $\\delta$-function permits one to reduce the integral equations for slow modes to one-dimensional integral equations, and to advance substantially in analytical calculations.  Several arguments can be brought in favour of our simplified approach. First of all, the Lynden-Bell derivative (see Paper I, eq. 4.7) of the DF with respect to angular momentum $L$, keeping $J = L + I_1$ constant (here $I_1$ is the radial action) in the limit where the slow mode approximation is applicable, can be replaced by a derivative, keeping energy $E$ constant: $$ \\left(\\frac{\\p F}{\\p L} \\right)_{LB} = \\Omega_\\textrm{pr} \\left(\\frac{\\p  F}{\\p E} \\right)_L  + \\left(\\frac{\\p F}{\\p L} \\right)_E \\approx \\left(\\frac{\\p F}{\\p L} \\right)_{E}, $$ because $\\Omega_\\textrm{pr}$ is small. Thus, the derivative over energy is not included into the slow integral equation, and one can loosely say, that dependence on energy is only parametric. Another argument is that the results of independent study by Tremaine (2005), who used a non-monoenergetic DF, are in agreement with our conclusions.  Section 2 is devoted to spheres, Section 3  -- to thin disks with symmetric DFs. The sections are organized alike. In the beginning we derive integral equations for initial distribution functions in the form (\\ref{eq:1.1}). Then follow analytical and numerical investigations of these equations. We demonstrate that by contrast to the case of near-Keplerian sphere, the loss-cone instability in disks takes place even for the monotonic DF, $df/d|L|>0$, provided the precession is retrograde and the loss cone is empty: $f(0)=0$. Sec. 2 is complimented by stability analysis of models with circular orbits, which of course doesn't belong to the class of monoenergetic models of (\\ref{eq:1.1}) type.  In the last,  Section 4, we discuss the results and some perspectives of further studies. ",
        "conclusions": "We have studied the stability of the spherically-symmetric and thin disk stellar clusters around a massive black hole. We conclude that stability properties of spherical clusters depend crucially on monotonity of initial distribution functions, while thin disk clusters are almost always unstable.  If the initial distribution of the spherical cluster is monotonic, the cluster is most likely to be stable. This conclusion was first made in Tremaine (2005), where stability of $l=1$ mode was generally proved, and $l=2$ was tested numerically. We confirm this conclusion by considering a number of monotonic distributions for modes with arbitrary $l$. Besides, we have checked distributions obtained from monotonic ones by making them vanish quickly but smoothly at circular orbits. These models were also stable. However, a general proof of stability for any monotonic distributions was not yet found.  Spherical clusters with the non-monotonic DFs should be generally affected by the gravitational loss-cone instability. The instability was first found in our Paper I using a simplification of systems with near-radial orbits. In the Sec. 2 we show that this instability is due to just non-monotony of distributions over angular momentum, the orbits may not necessary be near-radial.  In our opinion, both monotonic and non-monotonic distributions are important for possible applications to real stellar clusters around black holes. The DFs monotonically increasing from the loss cone radius up to circular orbits are formed naturally due to two-body collisions of stars. It follows from numerical experiments (see, e.g., Cohn and Kulsrud, 1978), which predict establishment of such distributions after a characteristic time for collisional relaxation. These distributions may be approximated by the formula $F\\propto \\ln \\bigl(L/L_{\\rm min}\\bigr)$.  Such a slowly increasing function is, in fact, predetermined by the boundary conditions imposed in the cited numerical study and some other investigations. Indeed, the vanishing condition at $L=L_{\\rm min}$, and the matching condition to isotropic (Maxwellian) distribution, $F=F(E)$, at the boundary $E=E_{\\rm bound}=0$ of the phase space $(E,L)$ (boundary separates stars which is gravitationally coupled to the black hole from the others) is required. The last condition means the asymptotic (when $E \\to E_{\\rm bound}$) independence of the function $F(E,L)$ on the momentum $L$. So monotonic, or logarithmic, dependence of type of (\\ref{eq:3.2}) is quite reasonable.  The non-monotonic distributions are also real. If the cluster, is formed, for example, as a result of the collisionless collapse (several free fall times), then it remains collisionless for a long timescale of collisional relaxation (see, e.g., Merritt \\& Wang, 2005). In principle, the system can have almost arbitrary DF both in the energy and in the angular momentum. During the collapse, a typical non-monotonic distribution of stars over the angular momentum, with empty loss cone and maximum at some value $L=L_{\\ast}$, is formed.  In Paper I we argued that stability properties of such a distribution is effectively analogous to one of typical plasma distributions of the ``beam-like'' type. But they can readily become unstable, as it is well-known in plasma physics (and also confirmed by direct stability study of corresponding stellar systems in Paper I). It is possible (as it is often so in plasma) that for the time of collisionless behavior, DF can undergo a dramatic change from its initial form. In particular, the collective flux of stars into the loss cone caused by the instability could, in principle, lead to the formation of a considerable part of the black hole. Checking of such possibilities is the most urgent task for future studies of unstable {\\it non-monotonic} models.  Since spherically-symmetric models with the {\\it monotonic} DF are apparently stable, but analogous disk systems are unstable (see Tremaine 2005 and Sec. 3), a critical flatness of ellipsoid models at which the instability begins is expected. Study of such systems, as well as systems with more complex triaxial ellipsoids can be performed using numerical simulations."
    },
    "0710/0710.1700_arXiv.txt": {
        "abstract": "The general world model for homogeneous and isotropic universe has been proposed. For this purpose, we introduce a global and fiducial system of reference (world reference frame) constructed on a $5$-dimensional space-time that is embedding the universe, and define the line element as the separation between two neighboring events that are distinct in space and time, as viewed in the world reference frame. The effect of cosmic expansion on the measurement of physical distance has been correctly included in the new metric, which differs from the Friedmann-Robertson-Walker metric where the spatial separation is measured for events on the hypersurface at a constant time while the temporal separation is measured for events at different time epochs. The Einstein's field equations with the new metric imply that closed, flat, and open universes are filled with positive, zero, and negative energy, respectively. The curvature of the universe is determined by the sign of mean energy density. We have demonstrated that the flat universe is empty and stationary, equivalent to the Minkowski space-time, and that the universe with positive energy density is always spatially closed and finite. In the closed universe, the proper time of a comoving observer does not elapse uniformly as judged in the world reference frame, in which both cosmic expansion and time-varying light speeds cannot exceed the limiting speed of the special relativity. We have also reconstructed cosmic evolution histories of the closed world models that are consistent with recent astronomical observations, and derived useful formulas such as energy-momentum relation of particles, redshift, total energy in the universe, cosmic distance and time scales, and so forth. It has also been shown that the inflation with positive acceleration at the earliest epoch is improbable. ",
        "introduction": "\\label{sec:intro}  The main goal of modern cosmology is to build a cosmological model that is consistent with astronomical observations. To achieve this goal, tremendous efforts have been made both on theories and on observations since the general theory of relativity was developed. So far the most successful model of the universe is the Friedmann-Robertson-Walker (FRW) world model \\cite{fri22,fri24,rob29,wal35}. The FRW world model predicts reasonably well the current observations of the cosmic microwave background (CMB) radiation and the large-scale structures in the universe. The precisely determined cosmological parameters of the FRW world model imply that our universe is consistent with the spatially flat world model dominated by dark energy and cold dark matter ($\\Lambda\\textrm{CDM}$) with adiabatic initial condition driven by inflation \\cite{spergel07,tegmark06}.  Although the flat FRW world model is currently the most reliable physical world model, one may have the following fundamental questions on the nature of the FRW world model. First, mathematically, if a space-time manifold is flat, then the Riemann curvature tensor should vanish, and vice versa. However, the Riemann curvature tensor of the flat FRW world model does not vanish unless the cosmic expansion speed and acceleration are zeros, which implies that the physical space-time of the flat FRW world is not geometrically flat but curved. Only its spatial section at a constant time is flat.  Secondly, the cosmic evolution equations of the FRW world model can be derived from an application of the Newton's gravitation and the local energy conservation laws to the dynamical motion of an expanding sphere with finite mass density \\cite{milne34,mccrea34}. Besides, the Newton's gravitation theory has been widely used to mimic the non-linear clustering of large-scale structures in the universe even on the horizon-sized $N$-body simulations \\cite{colberg2000,park05}. On large scales, the close connection between the FRW world model and the Newton's gravitation law is usually attributed to the fact that the linear evolution of large-scale density perturbations satisfies the weak gravitational field condition. Recently, Hwang and Noh \\cite{hwang06} show that the relativistic fluid equations perturbed to second order in a flat FRW background world coincide exactly with the Newtonian results, and prove that the Newtonian numerical simulation is valid in all cosmological scales up to the second order. However, one may have a different point of view that the Newton's gravitational action at a distance appears to be valid even on the super-horizon scales in the FRW world just because the world model does not reflect the full nature of the relativistic theory of gravitation.  Thirdly, according to the FRW world model, the universe at sufficiently early epoch ($z \\gtrsim 1000$) is usually regarded as flat since the curvature parameter contributes negligibly to the total density. The present non-flat universe should have had the density parameter approaching to $\\Omega = 1$ with infinitely high precision just after the big-bang (flatness problem). On the other hand, if we imagine the surface of an expanding balloon with positive curvature, then the curvature of the surface is always positive and becomes even higher as the balloon is traced back to the earlier epoch when it was smaller. This prediction from the common sense contradicts the FRW world model.  Observationally, the flat $\\Lambda\\textrm{CDM}$ universe is favored by the recent joint cosmological parameter estimation using the Wilkinson Microwave Anisotropy (WMAP) CMB \\cite{hinshaw07,page07}, large-scale structures \\cite{cole05,tegmark04}, type Ia supernovae (SNIa; \\cite{riess07,wood07}), Hubble constant \\cite{freedman01,macri06,sandage06}, baryonic oscillation data \\cite{eisen05}, and so on. However, the WMAP CMB data alone is more compatible with the non-flat FRW world model ($\\S7.3$ of \\cite{spergel07} and Table III of \\cite{tegmark06}). Besides, some parameter estimations using SNIa data or angular size-redshift data of distant radio sources alone suggest a possibility of the closed universe \\cite{clocchi06,jackson06}. The combinations of the WMAP plus the SNIa data or the Hubble constant data also imply the possibility of the closed universe, giving curvature parameters $\\Omega_k = -0.011\\pm 0.012$ and $\\Omega_k = -0.014\\pm 0.017$, respectively \\cite{spergel07}, although the estimated values are still consistent with the flat FRW world model.  The questions above and the observational constraints on the cosmological model may bring about possibilities of non-flat or non-FRW world models. Interestingly, Einstein claimed that our universe is spatially bounded or closed \\cite{ein22}. The primary reason for his preference to the closed universe is because Mach's idea \\cite{mach93,misner73} that the inertia depends upon the mutual action of bodies is compatible only with the finite universe, not with a quasi-Euclidean, infinite universe. According to Einstein's argument, an infinite universe is possible only if the mean density of matter in the universe vanishes, which is unlikely due to the fact that there is a positive mean density of matter in the universe \\footnote{However, in the appendix to the second edition of his book \\cite{ein22}, Einstein summarized Friedmann's world models and discussed a universe with vanishing spatial curvature and non-vanishing mean matter density, which differs from his original argument.}.  In this paper, we propose the general world model for homogeneous and isotropic universe which supports Einstein's perspective on the physical universe. The outline of this paper is as follows. In Sec. \\ref{sec:metric}, we consider the effect of cosmic expansion on the physical space-time distance between neighboring events and describe how to define the line element for homogeneous and isotropic universes of various spatial curvature types. The metric and the cosmic evolution equations for flat, closed, and open universes are derived in Sec. \\ref{sec:nspace}. It will be shown that our universe is spatially closed. In Sec. \\ref{sec:some}, we reconstruct cosmic evolution histories of the closed world models, and derive interesting properties of the closed universe. In Sec. \\ref{sec:inflation}, we discuss whether the inflation theory is compatible with the closed world model or not. Conclusion follows in Sec. \\ref{sec:conc}.  Throughout this paper, we adopt a sign convention $(+,-,-,-)$ for the metric tensor $g_{ik}$, and denote a 4-vector in space-time as $p^{i}$ ($i=0,1,2,3$) and a 3-vector in space as $p^{\\alpha}$ ($\\alpha=1,2,3$) or $\\mathbf{p}$. The Einstein's field equations are \\begin{equation} R_{ik} - \\frac{1}{2} g_{ik}R = 8\\pi G T_{ik}+\\Lambda g_{ik}, \\label{eq:einstein} \\end{equation} where $R_{ik}=R^{a}_{~iak}$ is the Ricci tensor, $R=R^i_{~i}$ the Ricci scalar, $T_{ik}$ the energy-momentum tensor, $G$ the Newton's gravitational constant, and $\\Lambda$ the cosmological constant. The Riemann curvature tensor is given by $R^{a}_{~ibk} =\\partial_b \\Gamma^{a}_{ki}-\\partial_{k}\\Gamma^{a}_{bi} +\\Gamma^{a}_{bn} \\Gamma^{n}_{ki}-\\Gamma^{a}_{kn}\\Gamma^{n}_{bi}$, with the Christoffel symbol $\\Gamma^{a}_{ik}={1\\over 2} g^{ab} (\\partial_i g_{kb}+\\partial_k g_{ib} -\\partial_b g_{ik})$. The energy-momentum tensor for perfect fluid is \\begin{equation} T_{ik} = (\\varepsilon_\\textrm{b} + P_\\textrm{b})u_i u_k - P_\\textrm{b} g_{ik}, \\label{eq:emtensor} \\end{equation} where $\\varepsilon_\\textrm{b}$ and $P_\\textrm{b}$ are background energy density and pressure of ordinary matter and radiation, and $u_{i}$ is the 4-velocity of a fundamental observer. We assume that the cosmological constant acts like a fluid with effective energy density $\\varepsilon_\\Lambda = \\Lambda/8\\pi G$ and pressure $P_\\Lambda = -\\varepsilon_\\Lambda$. The limiting speed in the special theory of relativity is set to unity ($c\\equiv 1$). ",
        "conclusions": "\\label{sec:conc}  In this paper, the general world model for homogeneous and isotropic universe has been proposed. By introducing the world reference frame as a global and fiducial system of reference, we have defined the line element so that the effect of cosmic expansion on the physical space-time separation can be correctly included in the metric. With this framework, we have demonstrated theoretically that the flat universe is equivalent to the Minkowski space-time and that the universe with positive energy density is always spatially closed and finite. The open universe is unrealistic because it cannot accommodate positive energy density. Therefore, in the world of ordinary materials, only the spatially closed universe is possible to exist.  The naturalness of the finite world with positive energy density comes from the Mach's principle that the motion of a mass particle depends on the mass distribution of the entire world. The principle is consistent only with the finite world because the dynamics of a reference frame cannot be defined in the infinite, empty world. The closed world model satisfies the Mach's principle and supports Einstein's perspective on the physical universe.  We have reconstructed evolution histories of the closed world models that are consistent with the recent astronomical observations, based on the nearly flat FRW world models (Model I and II; Sec. \\ref{sec:some}). The present curvature radius of the universe is $a_0 = 25.7$ Gpc ($a_0=40.2$ Gpc) for Model I (Model II). The expansion histories of both models imply that the closed universe dominated by dark energy expands eternally. However, the currently favored flat FRW world exists only as a limiting case of the closed universe with infinite curvature radius that is expanding with the maximum speed ($\\dot{a}_0=1$, $\\ddot{a}_0=0$).  From the local nature of the FRW metric (Sec. \\ref{sec:metric}) and of the proper time of a comoving observer (Sec. \\ref{sec:rel_friedmann}), it is clear that the FRW world model describes the local universe as observed by the comoving observer. Since the Newton's gravitation law can be derived from the Einstein's field equations in the weak field and the small velocity limits, the gravitational action at a distance usually holds at a local region of space on scales far smaller than the Hubble horizon size (e.g., \\cite{peebles80}). The proper Hubble radius $d_\\textrm{H}$ (Fig. \\ref{fig:horizon}, top) may provide a reasonable estimate of the characteristic distance scale where the Newton's gravity applies. The cosmic structures simulated by the Newton's gravity-based $N$-body method will significantly deviate from the real structures on scales comparable to $d_\\textrm{H}$. The variation of the comoving Hubble radius $\\chi_\\textrm{H} = d_\\textrm{H} /a$ also implies that in the past (future) the Newtonian dynamics was (will be) applicable on smaller region of space compared to the size of the universe (Fig. \\ref{fig:horizon}, bottom).  In this paper, the history of the universe has been tentatively reconstructed based on cosmological parameters of non-flat FRW world models. The more general cosmological perturbation theory and parameter estimation are essential for accurate reconstruction of the cosmic history."
    },
    "0710/0710.4914_arXiv.txt": {
        "abstract": "We present 2D simulations of the cooling of neutron stars with strong magnetic fields ($B \\geq 10^{13}$ G). We solve the diffusion equation in axial symmetry including the state of the art  microphysics that controls the cooling such as slow/fast neutrino processes, superfluidity, as well as possible heating mechanisms. We study how the cooling curves depend on the the magnetic field strength and geometry. Special attention is given to discuss the influence of magnetic field decay. We show that Joule heating effects are very large and in some cases control the thermal evolution. We characterize the temperature anisotropy induced by the magnetic field for the early and late stages of the evolution of isolated neutron stars. ",
        "introduction": "The observed thermal emission of neutron stars (NSs) can provide information about the matter in their interior. Comparing the theoretical cooling curves with observational data\\cite{Yakovlev2004,Page2006} one can infer not only the physical conditions of the outer region (atmosphere) where the spectrum is formed but also of the poorly known interior (crust, core) where high densities are expected.  There is increasing evidence that most of nearby NSs whose thermal emission is visible in the X-ray band have a non uniform temperature distribution\\cite{Zavlin2007,Haberl2007}~. There is a mismatch between the extrapolation to low energy of the fits to X-ray spectra, and the observed Rayleigh Jeans tail in the optical band ({\\it optical excess flux}), that cannot be addressed with a unique temperature (e.g. \\rxdieciocho\\cite{Pons2002}~, \\rbdoce\\cite{Schwope2007}~,  and \\rxcerosiete\\cite{Perez2006}~).  A non uniform temperature distribution may be produced not only in the low density regions\\cite{Greenstein1983}~, but also in intermediate density regions, such as the solid crust. Recently, it has been proposed that crustal confined magnetic fields with strengths larger than $10^{13}$ G   could be responsible for the surface thermal anisotropy \\cite{Geppert2004,Azorin2006}~. In the crust, the magnetic field limits the movement of electrons (main responsible for the heat transport) in the direction perpendicular to the field and the thermal conductivity in this direction is highly suppressed, while remains almost unaffected along the field lines.  Moreover, the observational fact that most thermally emitting isolated NSs have magnetic fields larger than $10^{13}$ G implies that a realistic cooling model must include magnetic field effects. In a recent work\\cite{Aguilera2007}~, first 2D simulations of the cooling of magnetized NSs have been presented. In particular, it has been stated that magnetic field decay, as a heat source, could strongly affect the thermal evolution and the observations should be reinterpreted in the light of these new results. We present the main conclusions of this work next. ",
        "conclusions": "The main result of this work is that, in magnetized NSs with $B> 10^{13}$~G, the decay of the magnetic field affects strongly their cooling. In particular,  there is a huge effect of Joule heating on the thermal evolution. In NSs born as magnetars, this effect plays a key role in maintaining them warm for a long time. Moreover,  it can also be important in high magnetic field radio pulsars and in radio--quiet isolated NSs. As a conclusion, the thermal and magnetic field evolution of a NS is at least a two parameter space (Fig~\\ref{fig_coupled}), and a first step towards a coupled magneto-thermal evolution has been given in this work. \\begin{figure}[htb] \\begin{center} \\psfig{file=BT_v1_1.ps,width=0.5\\textwidth,angle=-90} \\end{center} \\caption{Coupled magneto-thermal evolution of isolated neutron stars\\cite{Aguilera2007a}~: $T_s$ for the hot component as a function of $B$ and $t$. Observations: squares for magnetars ($B>10^{14}$~G), triangles for intermediate-field isolated NSs ($10^{13}$~G$<B<10^{14}$~G) and circles for radio pulsars ($B<10^{12}$~G). Corresponding cooling curves in solid, dashed and dotted lines, respectively. } \\label{fig_coupled} \\end{figure}"
    },
    "0710/0710.5123_arXiv.txt": {
        "abstract": "We investigate the effect of spiral structure on the Galactic disk as viewed by pencil beams centered on the Sun, relevant to upcoming surveys such as ARGOS, SEGUE, and GAIA. We create synthetic Galactic maps which we call Pencil Beam Maps (PBMs) of the following observables: line-of-sight velocities, the corresponding velocity dispersion, and the stellar number density that are functions of distance from the observer. We show that such maps can be used to infer spiral structure parameters, such as pattern speed, solar phase angle, and number of arms. The mean line-of-sight velocity and velocity dispersion are affected by up to $\\sim35$ km/s which is well within the detectable limit for forthcoming radial velocity surveys. One can measure the pattern speed by searching for imprints of resonances. In the case of a two-armed spiral structure it can be inferred from the radius of a high velocity dispersion ring situated at the 2:1 ILR. This information, however, must be combined with information related to the velocities and stellar number density in order to distinguish from a four-armed structure. If the pattern speed is such that the 2:1 ILR is hidden inside the Galactic bulge the 2:1 OLR will be present in the outer Galaxy and thus can equivalently be used to estimate the pattern speed. Once the pattern speed is known the solar angle can be estimated from the line-of-sight velocities and the number density PBMs. Forthcoming radial velocity surveys are likely to provide powerful constraints of the structure of the Milky Way disk. ",
        "introduction": "It has been well established by now that the Milky way is not axisymmetric with both a central bar and spiral structure perturbing its disk. Due to our location in the Galactic plane both spiral and bar structure is impossible to observe directly. Galactic bar parameters such as orientation and pattern speed have been inferred indirectly from both asymmetries around the Galactic center (e.g., \\citealt{blitz91,weinberg92}) and its effect on the local velocity distribution of old stars, i.e., the Hercules stream \\citep{dehnen99,dehnen00,fux01,mq07b}.  Spiral structure parameters, however, are much more uncertain. Current spiral density wave models \\citep{fux01,lepine01,desimone04,qm05} strongly disagree on the strength of the spiral structure, the number of arms, and the pattern speed. These models differ in their predictions of the induced velocity streaming at different angular positions in the Galaxy. For example, a four-armed density wave with velocity perturbations of $\\sim20$ km/s will exhibit rapidly varying radial and tangential velocity components with azimuth across distances of a few kpc, and we could expect to detect $\\sim20-50$ km/s variations in the mean line-of-sight stellar velocity as a function of the distance from the Sun. However, the strength of the spiral arm perturbation remains controversial. Based on the OGLE number counts, \\cite{paczynski94} estimated that the Sagittarius-Carina arm has a factor of two increase in density compared to the underlying disk. This model is inconsistent with COBE studies which find a much smaller contrast ($\\sim15\\%$) and show that the Perseus and Scu-Cru arms are more dominant \\citep{drimmel01}.  HI, CO, Cepheid, and far-infrared observations suggest that the Galactic disk contains a four-armed tightly wound structure. On the other hand, \\cite{drimmel01} have shown that the near-infrared observations are consistent with a dominant two-armed structure. \\citet{lepine01} suggest that locally the Milky Way can be modeled by the superposition of a two- and four-armed structure moving at the same pattern speed. By studying the nearby spiral arms, \\cite{naoz07} find that the Sagittarius-Carina arm is a superposition of two features, moving at different pattern speeds. The effect of a two- and four-armed structure, moving at different angular velocities, on the velocity dispersion of a galactic disk has been explored numerically by \\cite{mq06}.  Estimates for the pattern speed of the Milky Way spiral structure, or equivalently, the Sun's position with respect to resonances associated with spiral structure, span a large range of values. Reviewing previous work, \\citet{shaviv03} finds a clustering of estimates for the pattern speed of local spiral structure near $\\Omega_s \\sim  20 {\\rm  km s^{-1} kpc}^{-1}$, though other studies suggest $\\Omega_s \\sim 13 {\\rm  km s^{-1} kpc}^{-1}$. The model by \\citet{lepine01} places the Sun near the corotation resonance $\\Omega_s \\sim  28 {\\rm  km s^{-1} kpc}^{-1}$), and was fit to Cepheid kinematics. The recent gas dynamical studies \\citep{martos04, bissantz03} match the properties of the gas in nearby arms with a spiral pattern speed of $\\sim 20 {\\rm km s^{-1} kpc}^{-1}$. \\citet{martos04} propose that a two-armed stellar structure consistent with the stellar distribution inferred from COBE could cause four-arms in the gas distribution near the Sun. The gas dynamical model proposed by \\citet{bissantz03} with a similar spiral pattern speed matches HI and CO kinematics. The pattern speed of a spiral density wave can be tightly constrained from the location of its resonances. For example, \\cite{qm05} associated stellar streams in the solar neighborhood with the 4:1 ILR resonance of a two-armed pattern and were then able to tightly constrain the pattern speed of the driving spiral density wave to within 5\\%. Independent constraints on the pattern speed come from recent surveys of nearby open clusters (e.g. \\citealt{dias05}) where the older clusters are found to have drifted further from their original density wave location. These authors concluded that the Sun is located near the CR. A solar circle near the CR is also favored by \\cite{lepine01} and \\cite{naoz07}.  In this paper we investigate how spiral structure parameters can be inferred from velocity and density maps resulting from pencil-beam and large-scale surveys of the Galaxy. At present the influence of spiral arms on the observed kinematic properties of the Galactic disk is very poorly understood. With the advent of future Galactic all-sky (GAIA, SEGUE) and pencil-beam (ARGOS, BRAVA) radial velocity surveys, large amounts of kinematic data will be collected. The types of dynamical constraints made possible with these new data sets is not currently known. We address that issue here with synthetic models for the purpose of exploring how spiral structure might be constraint from these data. ",
        "conclusions": "Upcoming Galactic disk surveys will reveal the age, composition and phase space distribution of stars within the various Galactic components. These stellar excavations will provide essential clues for understanding the structure, formation and evolution of our Galaxy. To facilitate the interpretation of the huge amounts of data resulting from these surveys, Galactic disk models, such as the one presented here, are needed to interpret the observations.  We have investigated how the Milky Was spiral structure parameters, such as pattern speed and solar phase angle, can be estimated in a deep all-sky survey. We performed a series of test-particle simulations of a warm galactic disk approximating the disk kinematics of the Milky Way. We considered both two- and four-armed spiral structure and suggested a way to distinguish between the two using velocity and number density maps.  We found that the axisymmetric potential needs to be known to $\\sim10\\%$, line-of-sight velocities to $\\sim20$ km/s, and distance uncertainties need to be less than $\\sim30\\%$. The mean line-of-sight velocity and the velocity dispersion are affected by up to $\\sim35$ km/s which is well within the detectable limit for forthcoming radial velocity surveys. Pattern speed can be constrained by a hot ring at the 2:1 ILR in both two- and four-armed spiral structure. To distinguish between the two, however, we also need information related to the velocities and stellar number density. If the pattern speed is such that the 2:1 ILR is hidden inside the Galactic bulge the 2:1 OLR would be present in the outer Galaxy and thus can equivalently be used to estimate the pattern speed. Once the pattern speed is known the solar angle can be estimated from the number density variation with heliocentric distance; $\\phi_0$ is also reflected in the $v_d$ PBMs.  Future work needs to address the issue of how to obtain the axisymmetric background potential needed to subtract from the observational data as discussed at the end of Section \\ref{sec:pbm}. Also, it is important to know what type of tracer stars are needed that would allow the estimation of photometric parallaxes with errors less than $\\sim30\\%$, and the distribution of those stars.  While here we only considered steady state spiral stricture, other theories of spiral structure, such as transient and swing-amplified spirals, need to be investigated as well. It has also been suggested that the Galaxy harbors two sets of spiral structure moving at the same \\citep{lepine01} or different \\citep{naoz07,mq06} pattern speeds. We expect in all those cases it will be again resonant features to relate to the pattern speed and solar angle. In the case of non-steady state spirals, however, the structure in the PBMs will vary with integration time and interpretation will become more complicated. We refer these cases to a future study.  It is also known that the Milky Way is a barred galaxy. The simulations performed here do not include the influence of the bar. This is not necessarily a shortcoming since most of the features in the PBMs we use to infer spiral structure parameters are caused by resonances and, unless a resonance overlap with the central bar exists in the same location, those would not be different when a bar is included in the simulations. Future work should also look at this problem."
    },
    "0710/0710.2190_arXiv.txt": {
        "abstract": "The long-slit spectra obtained along the minor axis, offset major axis and diagonal axis are presented for 12 E and S0 galaxies of the Coma cluster drawn from a magnitude-limited sample studied before. The rotation curves, velocity dispersion profiles and the $H_3$ and $H_4$ coefficients of the Hermite decomposition of the line of sight velocity distribution are derived. The radial profiles of the \\Hb , Mg, and Fe line strength indices are measured too. In addition, the surface photometry of the central regions of a subsample of 4 galaxies recently obtained with Hubble Space Telescope is presented. The data will be used to construct dynamical models of the galaxies and study their stellar populations. ",
        "introduction": "This is the fourth of a series of papers aimed at investigating the stellar populations and the kinematics of early-type galaxies in the Coma Cluster. Spanning about 4 dex in the observed radial variation of the surface density of cluster members \\citep[e.g.][]{kent1982}, the Coma Cluster is the ideal place to investigate these galaxy properties as a function of the environmental density in order to test the theories for galaxy formation and evolution.  The sample of 35 E and S0 galaxies of the Coma Cluster is presented in the first paper of the series \\citep[hereafter Paper I]{mehlert2000} along with the photometry and long-slit spectroscopy along their major axis. From these spectra the rotation curves, velocity dispersion profiles and the $H_3$ and $H_4$ coefficients of the Hermite decomposition of the line-of-sight velocity distribution (LOSVD) were measured out to 1--3 effective radii with high signal-to-noise ratio ($S/N$). Moreover, the radial profiles of the \\Hb\\, Mg, and Fe line strength indices were measured too. Subsequently, the spectroscopic database was complemented with the long-slit spectra obtained along the minor axis, an offset axis parallel to the major one and one diagonal axis for 10 objects \\citep[hereafter Paper II]{wegner2002}.  The central values and major-axis logarithmic gradients for the line strength indices were derived by \\citet[][hereafter Paper III]{mehlert2003}. This allowed the estimation of the average ages, metallicities and \\aFe\\ ratios in the center and at the effective radius by using stellar population models with variable element abundance ratios from \\citet{dthomas2003}. There is a dichotomy among the population of S0 galaxies. Some of them are dominated by old stellar populations and are indistinguishable from E galaxies. The remaining ones host very young stellar populations; hence they must have experienced relatively recent star formation episodes. Most massive galaxies had the shortest star formation timescales and were the first to form. The absence of age gradients implies that the stellar populations at different radii formed at a common epoch. The \\aFe\\ enhancement is not restricted to galaxy centers but it is a global phenomenon. Finally, negative metallicity gradients were measured to be significantly flatter than what is expected from gaseous monolithic collapse models. This suggests the importance of mergers in the galaxy formation history.  Here the spectroscopic database of Paper I and II is completed with the long-slit spectra obtained along the minor axis, offset major axis and one diagonal axis for another 12 objects. As done in Paper II, these galaxies were selected from the sample of Paper I as the objects with the most extended and precise major-axis kinematics and therefore best-suited for dynamical modelling, balancing between the number of E and S0 galaxies. Moreover, the surface photometry of the central regions of a subsample of 4 galaxies recently obtained with Hubble Space Telescope (HST) is presented.  The data shown here and in Paper I will allow the study of the stellar population gradients for a large number of early-type galaxies \\citep{dthomas2008} in order to investigate possible systematic differences between the disk and bulge components of S0 galaxies. The photometric and kinematic data of the combined dataset allowed the construction of dynamical models of the objects to study the properties of the dark matter halos of flattened and rotating E and S0 galaxies \\citep{jthomas2005,jthomas2007}. In fact, the implementation of Schwarzschild's orbit superposition technique for axisymmetric potentials by \\citet{jthomas2004} was used to derive the stellar mass-to-light ratios and dark matter halo parameters for a subsample of 17 galaxies. About 10--50 percent of the mass inside the effective radius is dark with a central density which is at least one order of magnitude lower than the luminous mass density. The orbital system of the stars is reasonably close to isotropy, but the distribution function shows a lot of fine structure. This study was complementary to the one presented by \\citet{gerhard2001} focusing on round and non-rotating ellipticals.  The HST photometry is described in Sect. \\ref{sec:photometry}.  The spectroscopic galaxy sample, relative observations and data reduction are described in Sect. \\ref{sec:spectroscopy}. The measured stellar kinematics, and line indices are given in Sect. \\ref{sec:results}. Conclusions are drawn in Sect. \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  New radially resolved spectroscopy of 12 E and S0 galaxies of the Coma cluster was presented. The rotation curves, velocity dispersion profiles and the $H_3$ and $H_4$ coefficients of the Hermite decomposition of the line-of-sight velocity distribution were derived along the minor axis, offset major axis and one diagonal direction. Moreover, the line strength index profiles of Mg, Fe and \\Hb\\ line indices were measured too. In addition, the surface photometry of the central regions of a subsample of 4 galaxies recently obtained with HST/WFPC2 was presented.  The data complement the existing set (Paper I, Paper II) and have a precision and radial extent sufficient to construct flattened and rotating dynamical models of the galaxies and study their radially resolved stellar populations. Other papers address these issues (Paper III, \\citealt{jthomas2005,jthomas2007})."
    },
    "0710/0710.2473_arXiv.txt": {
        "abstract": "In order to study the process of cooling in dark-matter (DM) halos and assess how well simple models can represent it, we run a set of radiative SPH hydrodynamical simulations of isolated halos, with gas sitting initially in hydrostatic equilibrium within Navarro-Frenk-White (NFW) potential wells. Simulations include radiative cooling and a scheme to convert high density cold gas particles into collisionless stars, neglecting any astrophysical source of energy feedback. After having assessed the numerical stability of the simulations, we compare the resulting evolution of the cooled mass with the predictions of the classical cooling model of White \\& Frenk and of the cooling model proposed in the {\\sc morgana} code of galaxy formation. We find that the classical model predicts fractions of cooled mass which, after about two central cooling times, are about one order of magnitude smaller than those found in simulations. Although this difference decreases with time, after 8 central cooling times, when simulations are stopped, the difference still amounts to a factor of 2--3.  We ascribe this difference to the lack of validity of the assumption that a mass shell takes one cooling time, as computed on the initial conditions, to cool to very low temperature. Indeed, we find from simulations that cooling SPH particles take most time in traveling, at roughly constant temperature and increasing density, from their initial position to a central cooling region, where they quickly cool down to $\\sim10^4$ K. We show that in this case the total cooling time is shorter than that computed on the initial conditions, as a consequence of the stronger radiative losses associated to the higher density experienced by these particles. As a consequence the mass cooling flow is stronger than that predicted by the classical model.  The {\\sc morgana} model, which computes the cooling rate as an integral over the contribution of cooling shells and does not make assumptions on the time needed by shells to reach very low temperature, better agrees with the cooled mass fraction found in the simulations, especially at early times, when the density profile of the cooling gas is shallow. With the addition of the simple assumption that the increase of the radius of the cooling region is counteracted by a shrinking at the sound speed, the {\\sc morgana} model is also able to reproduce for all simulations the evolution of the cooled mass fraction to within 20--50 per cent, thereby providing a substantial improvement with respect to the classical model.  Finally, we provide a very simple fitting function which accurately reproduces the cooling flow for the first $\\sim10$ central cooling times. ",
        "introduction": "Understanding galaxy formation is one of the most important challenges of modern cosmology. The rather stringent constraints on cosmological parameters now placed by a number of independent observations \\citep[e.g.,][ for a recent review]{Springel06} allows us to precisely set the initial conditions from which the formation of cosmic structures has started. As a consequence, understanding the complex astrophysical processes, related to the evolution of the baryonic component, represents now the missing link toward a successful description of galaxy formation and evolution.  So far, two alternative approaches have been pursued to make quantitative predictions on the observational properties of the galaxy population and their evolution in the cosmological context. The first one is based on the so-called semi-analytical models (SAMs, hereafter; e.g., \\citealt{Kauffmann93,Somerville99,Cole00,Menci05,Monaco07}). In this approach, the background cosmological model predicts the hierarchical build-up of the Dark Matter (DM) halos, where gas flows in, cools and gives rise to the formation of galaxies, while the complex interplay between gas cooling, star formation, chemical enrichment and release of energy from supernova (SN) explosions and Active Galactic Nuclei (AGN) is modeled through a set of simplified or phenomenological models, which are specified by a number of free parameters. A posteriori, the values of the relevant parameters should then be constrained by comparing SAM predictions to observational data. The rather low computational cost of this approach makes it a quite flexible tool to explore the model parameter space.  The second approach is based on resorting to full hydrodynamical simulations, which include the processes of gas cooling and suitable sub-resolution recipes for star formation and feedback. The obvious advantage of this method, with respect to SAMs, is that galaxy formation can be described by following in detail the gas-dynamical processes which determine the evolution of the cosmic baryons during the shaping of the large-scale cosmic structures. However, its limitation lies in the high computational cost, which makes it difficult to explore in detail the parameter space describing the process of galaxy formation and evolution. For this reason, following galaxy formation with full hydrodynamical simulations in a cosmological environment of several tens of Mpc is a challenging task for simulations of the present generation \\citep[e.g., ][]{Nagamine04,Nagai05,Romeo05,Saro06}.  This discussion shows that SAMs and hydrodynamical simulations provide complementary approaches to the cosmological study of galaxy formation. The ability of hydrodynamical simulations of accurately following gas dynamics calls for the need of a close comparison between these two approaches, in order to test the basic assumptions of the SAMs. The best regime to perform this comparison is when one excludes the effect of all those processes, like feedback in energy and metals, whose different modeling in SAMs and simulation codes would make the comparison scarcely telling. Since gas cooling is the most basic ingredient in any model of galaxy formation \\citep[e.g., ][]{White91}, an interesting comparison would be performed when cooling is the only process turned on. In this spirit \\cite{Benson01}, using a hydrodynamical simulation of a cosmological box and a stripped-down version of SAM, compared the statistical properties of ``galaxies'' found in the two cases. They discovered that SPH simulation and SAM give similar results for the thermodynamical evolution of gas and that there is a very good agreement in terms of final fractions of hot, cold and uncollapsed gas.  Similar conclusions were reached by \\cite{Helly03} and \\cite{Cattaneo07}. They improved the comparison performed by \\cite{Benson01} by giving to the down-stripped SAM the same halo merger histories extracted from the cosmological simulation. In this way they were able to compare cooling in DM halos not only statistically but on an object-by-object basis.  The result was again that the two methods provide comparable ``galaxy'' populations. \\cite{Yoshida02} performed a similar comparison for a simulation of a single galaxy cluster, obtaining similarly good results.  While the general agreement between the two methods is encouraging, still all the above analyses generally concentrated on comparing the statistical properties of the galaxy populations. Furthermore, if one wants to test the reliability of the cooling model implemented in the SAMs, the cleanest approach would be that of turning off the complications associated to the hierarchical merging of halos, thereby allowing gas to cool in isolated halos.  The purpose of this paper is to present a detailed comparison between the predictions of cooling models, as implemented in SAMs, and results of hydrodynamical simulations in which gas is allowed to cool inside isolated halos. Our controlled numerical experiments will be run for halos having the density profile (for DM particles) of \\cite{NFW} (NFW hereafter), with a range of masses, concentration parameters and average densities (related to the halos' redshift). As a baseline model for gas cooling, we consider the classical one, as originally proposed by \\cite{White91}.  In this model, the cooling radius is defined as the radius at which the cooling time equals the time elapsed since radiative cooling is turned on.  As a result, the growth rate of the cooled gas mass is simply related to the growth rate of the cooling radius. This model was claimed by \\cite{White91} to be close the exact self-similar solutions of cooling flows presented by \\cite{Bert89}.  Simulation results will also be compared to another model of gas cooling, which has been recently proposed by \\cite{Monaco07} in the context of the {\\sc morgana} SAM, and is based on a ``dynamical'' definition of cooling radius. As a result of our analysis, we will show that the gas cooling rate in the simulations is initially faster than predicted by the classical cooling model. When the simulations are stopped, after about 8 central cooling times, this initial transient causes the classical cooling model to underestimate the cooled mass by an amount which can be as large as a factor of three, depending on the halo concentration, density and mass. A much better agreement with simulations is achieved by the alternative {\\sc morgana} model.  The plan of the paper is as follows. In section 2 we describe first the ``classical'' analytic model for cooling \\citep{White91}, and the alternative {\\sc morgana} cooling model \\citep{Monaco07}. In section 3 we present numerical simulations, performed with the {\\tt GADGET-2} code and in section 4 we discuss the results obtained by comparing simulations to analytical cooling models. We discuss our results in section 5 and draw our final conclusions in section 6.  A more technical discussion on the differences between the classical and the {\\sc morgana} cooling models is provided in Appendix A, while Appendix B gives a very simple fitting formula for the cooling flows. ",
        "conclusions": "We have presented a detailed analysis of cooling of hot gas in DM halos, comparing the predictions of semi-analytic models with the results of controlled numerical experiments of isolated NFW halos with hot gas in hydrostatic equilibrium. Simulations have been performed spanning a range of masses (from galaxy- to cluster-sized halos), concentrations and redshift (from 0 to 2). Smaller halos at higher redshift have not been simulated because the validity of the assumption of a hydrostatic atmosphere is doubtful when the cooling time is much shorter than the dynamical time.  We have considered the ``classical'' cooling model of \\cite{White91}, used in most SAMs, and the model recently proposed by \\cite{Monaco07} within the {\\sc morgana} code for the evolution of galaxies and AGNs.  The main features of these models can be summarized as follows.  The density and pressure profiles of the gas are computed by solving the equation of hydrostatic equilibrium in an NFW halo \\citep{Suto98}.  The classical cooling model assumes that each mass shell cools to low temperature exactly after one cooling time $t_{\\rm cool}(r)$, computed on the initial conditions.  The cooling radius $r_{\\rm C}$ is then the inverse of the $t_{\\rm cool}(r)$ function, and the cooled fraction is the fraction of gas mass within $r_{\\rm C}$.  The ``unclosed {\\sc morgana}'' cooling model computes the cooling rate of each mass shell, then integrates over the contribution of all mass shells and follows the evolution of the cooling radius assuming that the transition from hot to cold phases is quick enough so that a sharp border in the density profile of hot gas is always present.  This determines the evolution of the cooling radius $r_{\\rm M}$.  Moreover, to mimic the closure of the ``cooling hole'' due to the lack of pressure support at $r_{\\rm M}$, the cooling radius (now called $r_{\\rm M,ch}$) is induced to close at the sound speed. This defines the ``closed {\\sc morgana}'' model.  Our main results can be summarized as follows.  \\noindent {\\bf (i)} The classical cooling model systematically underestimates the fractions of cooled mass. After about two central cooling times, they are predicted to be about one order of magnitude smaller than those found in simulations. Although this difference decreases with time, after 8 central cooling times, when simulations are stopped, the difference still amounts to a factor 2--3. This disagreement is ascribed to the lack of validity of the assumption that each mass shell takes one cooling time, computed on the initial conditions, to cool to low temperature. Seen from the point of view of a mass element, the time required by it to cool to low temperature is shorter than the initial cooling time when density increases and temperature is constant during cooling.  This is what happens to gas particles in the simulations: they take most of time to travel from their initial position toward the cooling region, at roughly constant temperature and increasing density.  The disagreement is stronger when the cooling gas comes from the shallow central region, in which case the cooling flow is markedly not self-similar.  \\noindent {\\bf (ii)} The unclosed {\\sc morgana} model gives a much better fit of the cooled mass fraction.  This is mostly due to the relaxation of the assumption on the cooling time, mentioned in point (i).  This model correctly predicts cooling flows which are stronger than the classical model, by a larger amount for flatter gas density profiles.  In the Appendix we show that the solution is not self-similar if the slope of the density profile is shallower than $r^{-3/2}$. In this case cooling is not dominated by the the shells just beyond the cooling radius but the whole region for which the density profile is shallow contributes.  \\noindent {\\bf (iii)} The closed {\\sc morgana} model further improves the fit to the simulation results on the evolution of the cooled mass fraction, giving accurate results to within 20--50 per cent in all the considered cases, after about 8 central cooling times.  This agreement is a good reward for the increase of physical motivation of this model, obtained at the modest price of letting the cooling radius close at the local sound speed.  However, the closure of the cooling radius must be halted at later times for the model to give realistic results. In general, we consider the closed {\\sc morgana} as a successful effective model of cooling, rather than as a rigorous physical model.  \\noindent {\\bf (iv)} The cooling flow is well approximated by a constant flow, for which we give a fitting formula in Appendix B, and which is valid up to $\\sim10$ central cooling times.  In the context of models of galaxy formation, cooling of hot virialized gas is the starting point for all the astrophysical processes involved in the formation of stars (and supermassive black holes) and their feedback on the interstellar and intra-cluster media. We find that the classical model, used in most SAMs, leads to a significant underestimate of the cooled mass at early times. This result is apparently at variance with previous claims, discussed in the Introduction, of an agreement of models and simulations in predicting the cooled mass, (\\citealt{Benson01,Helly03,Cattaneo07,Yoshida02}). Given the much higher complexity of the cosmological initial conditions used in such analyses, it is rather difficult to perform a direct comparison with the results of our simulations of isolated halos. Here we only want to stress the advantage of performing simple and controlled numerical tests in order to study in detail how the process of gas cooling takes place.  It is well possible that the different behaviour of simulations and the classical cooling model is less apparent when the more complex cosmological evolution is considered.  However, there is no doubt that the results of our analysis are quite relevant for the comparison between SAM predictions and observations. For instance, \\citep{Fontanot07} have recently shown that the {\\sc morgana} model of galaxy formation is able to reproduce the observed number counts of sources in the sub-mm band by using the standard Initial Mass Function (IMF) by \\cite{Salpeter55}, without any need to resort to a top-heavier IMF \\citep{Baugh06}.  As argued by these authors, the bulk of starbursts are driven by massive cooling flows, so this difference is mostly due to the different cooling models."
    },
    "0710/0710.5253_arXiv.txt": {
        "abstract": "Nine supergiant shells (SGSs) have been identified in the Large Magellanic Cloud (LMC) based on H$\\alpha$ images, and twenty-three SGSs have been reported based on \\ion{H}{1} 21-cm line observations, but these sets do not always identify the same structures. We have examined the physical structure of the optically identified SGSs using \\ion{H}{1} channel maps and P-V diagrams to analyze the gas kinematics. There is good evidence for seven of the nine optically identified SGSs to be true shells. Of these seven H$\\alpha$ SGSs, four are the ionized inner walls of \\ion{H}{1} SGSs, while three are an ionized portion of a larger and more complex \\ion{H}{1} structure. All of the H$\\alpha$ SGSs are identified as such because they have OB associations along the periphery or in the center, with younger OB associations more often found along the periphery. After roughly 12 Myrs, if no new OB associations have been formed a SGS will cease to be identifiable at visible wavelengths. Thus, the presence and location of ionizing sources is the main distinction between shells seen only in \\ion{H}{1} and those also seen in H$\\alpha$. Based on our analysis, H$\\alpha$ observations alone cannot unambiguously identify SGSs, especially in distant galaxies. ",
        "introduction": "\\label{sec:intro}  Observations of the interstellar medium (ISM) of spiral and irregular galaxies have revealed complex filamentary structures indicative of a violent ISM \\citep{MS79}. The largest of these structures are the supergiant shells (SGSs), with diameters approaching 1 kpc \\citep{gm}. SGSs are thought to be formed by the fast stellar winds and supernova explosions of multiple OB associations; they require $10^{52}$--$10^{53}$ ergs for their creation, the equivalent of tens to hundreds of supernova explosions \\citep{M80}. The diameter of these shells often exceeds the scale height of the galactic gas disk, allowing them to puncture the gas disk and vent their hot interior gas into the galactic halo. The expansion of SGSs may also cause further star formation.  It has been suggested that the compression of the ISM on the rims of these shells and its subsequent gravitational collapse may be a significant cause of self-propagating star formation in galactic disks, thereby providing important insight into galactic evolution.  To examine the impact of SGSs on their host galaxy, it is necessary to determine the physical structure of SGSs that have been identified based on their morphology alone, since some may be merely chance superpositions of unrelated filaments.  The Large Magellanic Cloud (LMC) is an excellent site to explore the nature of SGSs.  With its low inclination there is little line-of-sight confusion, and its proximity (50 kpc; Feast 1999) allows us to probe its ISM in a detail impossible for most other galaxies. SGSs in the LMC were first identified by \\citet{gm}, based on visual inspection of H$\\alpha$ images of the ionized ISM. This original paper tabulated four SGSs, a number that was later expanded to nine by \\citet{M80}. Later work on LMC SGSs, however, has also been done based on observations of \\ion{H}{1}, the neutral atomic ISM. \\citet{Kim99} examined the LMC in \\ion{H}{1} and tabulated 23 SGSs that have diameters greater than 360 pc. Curiously, there is not a one-to-one correspondence between the largest \\ion{H}{1}-selected SGSs and the nine H$\\alpha$-selected SGSs.  To investigate the physical structure of the H$\\alpha$-selected SGSs and their relationship with the \\ion{H}{1}-selected SGSs, we have examined the \\ion{H}{1} environment of the H$\\alpha$-selected SGSs using the H$\\alpha$ images from the Magellanic Cloud Emission Line Survey \\citep[MCELS;][]{MCELS05} and the \\ion{H}{1} synthesis maps made with combined observations from the Australia Telescope Compact Array and Parkes Telescope \\citep{Kim03}.  In this paper we describe these data sets in \\S 2 and our methodology in \\S 3, and critically examine the physical structure of the nine H$\\alpha$-selected SGSs in \\S 4.  We discuss the nature of the SGSs and compare the \\ion{H}{1}-selected and  H$\\alpha$-selected shells in \\S 5.  A summary is given in \\S 6. ",
        "conclusions": "\\label{sec:disc}  Comparison of neutral and ionized hydrogen images has allowed us to gain an understanding of the structure and kinematics of the optically identified SGSs. We have determined that these shells can be separated into three categories:  \\begin{itemize} \\item {\\bf Simple Coherent Shells:} These are neutral shells whose inner walls are photoionized. While the stars in the interior of the SGSs are responsible for the expanding structure, most of the ionizing sources are distributed along the periphery. Of the optically identified LMC SGSs, LMC 1, 4, 5 and 6 fall within this category.  \\item {\\bf Complex Shells:} These are structures in which the ionized shell delineates only part of a larger or more complex neutral shell. The complex shells include LMC 2, 3 and 8. In the cases of LMC 2 and LMC 3, the optical shell illuminates only the northern portion of an larger neutral structure, while the single optically identified shell LMC 8 appears to contain multiple expanding neutral structures, of which only the eastern part is associated with ionized gas.  \\item {\\bf False Shells:} These are those optically identified structures whose neutral \\ion{H}{1} counterpart does not show any of the following characteristics to indicate a shell structure: (1) expansion within the shell boundary, as expected from an expanding shell; (2) raised column densities along shell rims at the systemic velocity, showing a shell morphology; and (3) central cavity at the systemic velocity, or for an \\ion{H}{1} hole. A physical shell, even if it has been stalled, can still be recognized by the latter two characteristics. We found no evidence in neutral hydrogen to confirm the shell structure of LMC 7 and LMC 9, and characterize them as false shells.  LMC 7 lacked any corroborating \\ion{H}{1} structure, while the optically identified shell LMC 9 was seen to be a chance superposition of ionized filaments and OB associations, with a neutral hydrogen structure in no way indicating a shell. \\end{itemize}  The nature of the optically identified SGSs can be understood as the result of the interplay between OB associations and the ISM. The fast stellar winds and supernova explosions from OB associations clear away the interstellar gas and form a large expanding shell structure.  The initial OB associations will lose their ionizing power after $\\sim$12 Myr, the lifetime of a 15 $M_{\\odot}$ star, the least massive star whose UV flux still produces detectable \\ion{H}{2} gas.  The recombination timescale of ionized gas is $7.6\\times10^4 N_{\\rm e}^{-1}$ yr, where $N_{\\rm e}$ is the electron density in cm$^{-3}$.  The interstellar gas density in a SGS is most likely in the range of 0.1--1 cm$^{-3}$, and the recombination timescale is much less than 1 Myr.  Thus, a SGS with no new OB associations will cease to be observable in H$\\alpha$ soon after 12 Myrs, and only an \\ion{H}{1} shell will be detected.  A summary of the diameters ($D$), average expansion velocities ($V_{\\rm exp}$), and maximum expansion velocities ($V_{\\rm max}$) of the optically identified SGSs in the LMC is given in Table \\ref{tab:SGS}. We estimate their ages from these data using the formula $t \\simeq \\eta (D/2) V_{\\rm exp}^{-1}$, where $\\eta$ is 0.4 for a momentum-conserving bubble \\citep{S75}, and 0.6 for an energy-conserving bubble \\citep{C75, W77}. We adopt an intermediate value $\\eta=0.5$. Examining the ages derived in Table \\ref{tab:SGS}, we find that all but one, SGS 8, are less than 12 Myr old. SGS 8, however, contains the OB associations LH18, LH24 and LH26 (see section \\ref{sec:sgs8}), of which LH24 and LH26 are still young enough to contain \\ion{H}{2} gas. Those that are younger than 12 Myr all have young OB associations along their peripheries.  These young OB associations explain the presence of an optical counterpart to this \\ion{H}{1} shell, regardless of their ages.  Of the \\ion{H}{1}-identified SGSs \\citep{Kim99}, 15 have average diameters $\\geq$ 600 pc, the threshold used for the optical SGSs: SGS 1, 3, 4, 5, 6, 7, 9, 10, 11, 12, 17, 18, 19, 20, and 23. Some of these \\ion{H}{1}-identified SGSs are associated with the H$\\alpha$-identified SGSs: SGS 3 with LMC 1, SGS 4 with LMC 8, SGS 7 with LMC 5, SGS 9 with LMC 9, SGS 11 with LMC 4, SGS 12 with LMC 3, and SGS 19 and 20 with LMC 2.  As described in \\S 4 and discussed earlier in this section, the H$\\alpha$-identified SGSs may correspond to partial or complete walls of \\ion{H}{1}-identified SGSs, depending on the distribution of OB associations that provide ionizing fluxes. Among the other \\ion{H}{1} SGSs that are not associated with H$\\alpha$ shells, SGS 1 and 6 do not contain any OB associations, and hence have no ionized counterparts; SGS 5 and 23 have no OB associations or \\ion{H}{2} regions along the most well-defined shell rims; SGS 10 and 18 do not show obvious shell structures in either column density or channel maps, so they may not be physical shells; SGS 17's northern rim is coincident with a long H$\\alpha$ arc that is likely to be ionized by massive stars in the 30 Doradus complex in the south \\citep[the H$\\alpha$ arc is sketched in dashes in Fig.\\ 1 of][]{M80}.  These comparisons adequately illustrate the importance of the presence of young OB associations to photoionize the gas in \\ion{H}{1} SGSs to produce H$\\alpha$-identified shell morphology.  Supergiant shells have been identified in distant galaxies based on H$\\alpha$ images \\citep[e.g.,][]{HG97,M98}. However, from our analysis of the LMC H$\\alpha$ SGSs, it is clear that the H$\\alpha$ morphology alone is not sufficient to identify physically expanding shells and that the H$\\alpha$ SGSs represent only the SGS population that possess young OB associations. High-resolution \\ion{H}{1} position-velocity datacubes are needed to verify the shells' structure and to reveal the entire population of SGSs."
    },
    "0710/0710.3063_arXiv.txt": {
        "abstract": "{The high energy neutrino telescope NT200+ is currently in operation in Lake Baikal. We review the status of the Baikal Neutrino Telescope, and describe recent progress on key components of the next generation kilometer-cube (km3) Lake Baikal detector, like investigation of new large area phototubes, integrated into the telescope. }  \\begin{document} ",
        "introduction": "The Baikal Neutrino Telescope  is operated in Lake Baikal, Siberia,  at a depth of {1.1~km}. Deep Baikal water is characterized by an absorption length of $L_{abs}(480 $nm$) =20\\div 24$ m, a (geometric) scattering length of $L_s =30\\div 70$ m and a strongly anisotropic scattering function with a mean cosine of scattering angle $0.85\\div 0.9$ \\cite{APP1}, and by a level of bioluminescence and other natural backgrounds that are well below seawater sites.   The first stage telescope, NT200, started full operation in spring 1998 and contained 192 Optical Modules (OMs). The favorable water properties, and a relatively simple and reliable design led to the physics success of this comparably small telescope. Low light scattering allows for a sensitive volume of a few Mtons at PeV shower energy scale, well beyond the geometric detector limits. For a review of the high sensitivity limits on UHE astrophysical neutrino's as well as best so far obtained limits on relativistic magnetic monopoles and other results, see \\cite{ICRC07_B}. The upgrade to NT200+ was a logical consequence of the large external sensitive volume, now to be fenced by sparsely instrumented external strings of OMs.    In this paper, we review the current status of the Baikal Neutrino Telescope as of 2007, and the activities towards the km3-scale detector. Results on a prototype device for acoustic neutrino detection, obtained with a stationary setup in 2006/2007, are reported elsewhere in these proceedings \\cite{ICRC07_C}. ",
        "conclusions": "The Baikal Neutrino Telescope is taking data currently in it's NT200+ configuration - an upgrade of the original NT200 telescope for improved high energy shower sensitivity.     For a km3-detector in Lake Baikal, R\\&D-activities have been started. The NT200+ detector is, beyond its better physics sensitivity, used as an ideal testbed for critical new components. Modernization of the NT200+  DAQ allowed to install a prototype FADC PM readout. Six large area hemispherical PMs have been integrated into NT200+ (2 Photonis XP1807/12\" and 4 Hamamatsu R8055/13\"), to facilitate an optimal PM choice. A prototype new technology string will be installed in spring 2008; and a km3-detector Technical Design Report is planned for fall 2008.               {\\it This work was supported by the Russian Ministry of Education and Science, the German Ministry of Education and Research and the Russian Fund of Basic Research (grants 05-02-17476, 05-02-16593, 07-02-10013, 07-02-00791), by the Grant of the President of Russia NSh-4580.2006.2 and by NATO-Grant NIG-9811707 (2005).}    \\vspace*{-3mm}"
    },
    "0710/0710.3313_arXiv.txt": {
        "abstract": "A multi-dimension, time-dependent Monte Carlo code is used to compute sample $\\gamma$-ray spectra to explore whether unambiguous constraints could be obtained from $\\gamma$-ray observations of type~Ia supernovae. Both spherical and aspherical geometries are considered and it is shown that moderate departures from sphericity can produce viewing-angle effects that are at least as significant as those caused by the variation of key parameters in one-dimensional models. Thus $\\gamma$-ray data could in principle carry some geometrical information, and caution should be applied when discussing the value of $\\gamma$-ray data based only on one-dimensional explosion models. In light of the limited sensitivity of current $\\gamma$-ray observatories, the computed theoretical spectra are studied to revisit the issue of whether useful constraints could be obtained for moderately nearby objects. The most useful $\\gamma$-ray measurements are likely to be of the light curve and time-dependent hardness ratios, but sensitivity higher than currently available, particularly at relatively hard energies ($\\sim 2$ -- $3$~MeV), is desirable. ",
        "introduction": "\\label{sect:intro}  Although the paradigm that Type Ia supernovae (SNe~Ia) result from the explosions of carbon-oxygen white dwarfs is well established, many issues regarding the nature of the progenitors and the explosion mechanism remain unclear (see e.g. \\citealt{hillebrandt00}). Achieving a clearer understanding of SNe~Ia is important because of their role in the chemical evolution of galaxies and as cosmological distance indicators.  Considerable theoretical effort has gone into modelling of SNe~Ia explosions. Recently, fully three-dimensional (3D) modelling of the explosion \\citep{reinecke02,gamezo03,roepke05,roepke06,jordan07} has become feasible, allowing a detailed treatment of the hydrodynamical instabilities and turbulence which play a pivotal role. The 3D structure predicted by these models has been shown to affect observables such as optical/ultra-violet/infrared ({\\sc uvoir}) light curves and spectra (e.g. \\citealt{kasen06c, sim07, sim07b}).  Although they are primarily detected through their optical emission, SNe~Ia are also expected to be $\\gamma$-ray sources owing to the large masses of radioactive isotopes synthesised in the explosion (e.g. \\citealt{travaglio04}). In principle, measurements of $\\gamma$-ray emission from SNe~Ia could provide important diagnostics since $\\gamma$ rays trace almost directly the mass and velocity distribution of the products of nuclear burning. Therefore, there has been considerable theoretical work on SN~Ia $\\gamma$-ray emission (e.g. \\citealt{ambwani88, burrows90, mueller91, hoeflich92, kumagai97, hoeflich98, gomez98, milne04}). Unfortunately, owing to the low sensitivity of $\\gamma$-ray observatories, to date only one SN~Ia (SN1991T) has been detected (see discussion by \\citealt{milne04}). However, with current instrumentation such as the {\\it SPI} \\citep{vedrenne03} on-board {\\it Integral} \\citep{winkler03}, detection of SN~Ia within several Mpc would be possible (\\citealt{gomez98}). Moreover, increased sensitivity in future missions should make detection feasible for more distant SNe~Ia.  Most studies of SN~Ia $\\gamma$-ray emission focus on making predictions from individual models representing specific explosion mechanisms. In general, these studies have been restricted to one-dimension (although see \\citealt{hoeflich02}) and to Chandrasekhar-mass models (some sub-Chandrasekhar models have been considered, e.g. \\citealt{hoeflich98}; \\citealt{gomez98}). Such studies demonstrated that, for nearby SN~Ia, the prospects of obtaining useful data are fairly good. Here we adopt a different but complementary approach, motivated by the increasing variety of explosion conditions suggested (both in theoretical and semi-empirical studies): rather than determining the quality of $\\gamma$-ray data that would be needed to distinguish specific models, we attempt to show what might be unambiguously determined solely from data and the physics of $\\gamma$-ray radiation transport.  We use a multi-dimensional, time-dependent code to compute $\\gamma$-ray spectra for a set of parameterised geometries spanning a broad range in the relevant physical conditions. Using these reference spectra, we highlight the quantities that would be most useful diagnostics and therefore most worthy of consideration in the design of future instrumentation.  In Section~\\ref{sect:code}, we briefly describe the code used to compute the reference $\\gamma$-ray spectra. The reference models we employ are motivated in Section~\\ref{sect:models} and their spectra are discussed in Section~\\ref{sect:spectra}. Guided by our reference spectra, in Section~\\ref{sect:diagnostics} we examine observational diagnostics. Finally, in Section~\\ref{sect:summ}, we summarise our results. ",
        "conclusions": "\\label{sect:spectra}  \\begin{figure} \\epsfig{file=fig3.eps, width=8.5cm} \\caption{ Representative time series of $\\gamma$-ray spectra computed with four spherically symmetric models (identified in the second panel) at times after explosion as indicated on each panel. The $^{56}$Ni and $^{56}$Co emission lines are identified in the top panel. The four flux-bands discussed in Section 5.3 are indicated in the bottom panel. Fluxes are for a source distance of 1~Mpc. \\label{fig:spec} } \\end{figure}  Fig.~\\ref{fig:spec} shows a time series of spectra computed from Model~SC and spectra from three other spherical models (SS, SFeR and SM) are over-plotted. (For clarity, spectra from the remaining four models are not plotted. They are, however, included in all the discussions of observable diagnostics in subsequent sections.)  At times close to maximum light the $\\gamma$-ray spectrum consists of strong emission lines, mainly due to $^{56}$Co, with significant continuum arising from Compton scattering of line photons. Since optical depths decrease with time in the expanding ejecta, the strength of the lines relative to the continuum increases with time. At early times, the ``mixed'' model (Model~SM) shows a harder continuum than the standard model and stronger Ni emission lines (particularly 0.158~MeV and 0.275~MeV). These lines are also fairly strong in Model~SNiS (not plotted) where they originate in the surface layer of $^{56}$Ni. As noted by \\citet{gomez98}, these Ni lines can only form when the source Ni lies at small $\\tau_C$; since $\\sigma_C$ decreases with increasing photon energy, the soft-energy lines are degraded most easily, becoming swamped by photons down-scattered from harder energies.  In Model~SS, the 0.158 and 0.275~MeV lines are almost completely buried by the strong continuum which persists until well after maximum light as a consequence of the high $\\tau_C$'s in this model.  Model~SFeR shows the effect of composition as intended. Above about 0.3~MeV, its spectra are indistinguishable from those of Model~SC. At soft energies, however, the Model~SFeR flux is lower by up to an order of magnitude. This is due entirely to the difference in the photoabsorption cross-section arising from the choice of composition. Although weaker, a similar effect appears in Model~SNiS because of the relatively large amounts of iron-group material in that model, particularly in the surface layer.  Should energy-resolved data of sufficient sensitivity to measure both the line and continuum emission across the entire 0.1 -- 3.5~MeV spectrum be obtained, it would be used to evaluate SN models by direct comparison. However, given the rarity of very nearby SNe~Ia and the limited sensitivity of $\\gamma$-ray observatories, in the next section we will consider what information could be extracted from low quality $\\gamma$-ray data in the form of simple diagnostics. We will focus on line and continuum fluxes. Although the energy resolution of $\\gamma$-ray instruments can be high enough to resolve spectral lines (e.g. with {\\it SPI Integral}, see \\citealt{roques03}), sensitivity limits make it practically impossible to measure detailed line shapes except for extremely nearby events \\citep{gomez98}.   \\label{sect:summ}  Using a Monte Carlo code we computed $\\gamma$-ray spectra for a variety of models to explore whether unambiguous constraints could be obtained from $\\gamma$-ray observations of SNe~Ia. Two aspherical toy geometries (a lop-sided distribution of Ni and an ellipsoidal ejecta) show that moderate departures from sphericity can produce viewing-angle effects at least as significant as those due to variations of key parameters in 1D models. Thus $\\gamma$-ray data could carry some useful constraints on possible geometries, but caution must be applied when evaluating its potential usefulness in distinguishing specific explosion scenarios.  Given the limited sensitivity of current $\\gamma$-ray missions, we conclude, in agreement with previous studies (e.g. \\citealt{gomez98}), that there is little prospect for obtaining useful constraints from line-ratio diagnostics except for fortuitously nearby objects. Instead, we suggest that the best prospects are offered by hardness ratios. In particular, owing to the simplicity of the physics underlying the $\\gamma$-ray spectrum, a simple ratio of the total emission in a hard energy, line-dominated part of the spectrum to a soft energy, continuum-dominated region could discriminate between our more extreme models. We also emphasise the value of obtaining multiple observations over a wide time period given the diagnostic power of the light curve shape. In planning future $\\gamma$-ray missions, greater sensitivity at harder energies ($\\simgt 2$~MeV) should be given high priority since several of the most potentially useful diagnostics (e.g. our $R_{1}$- and $H_{4}$-ratios) require measuring the hardest energies in the spectrum."
    },
    "0710/0710.4019_arXiv.txt": {
        "abstract": "Principal Component Analysis (PCA) is a well--known technique used to decorrelate a set of vectors. It has been applied to explore the star formation history of galaxies or to determine distances of mass--lossing stars. Here we apply PCA to the optical data of Planetary Nebulae (PNe) with the aim of extracting information about their morphological differences. Preliminary analysis of a sample of 55 PNe with known abundances and morphology shows that the second component (PC2), which results from a relation produced by the parameters log(N/O), initial and final mass of PNe, is depending on the morphology of PNe. It has been found that when log(N/O) $< -0.18$ the PNe's nitrogen is low independently on the oxygen abundance for either Bipolar ({\\it B}), Elliptical ({\\it E}) or Round ({\\it R}) PNe. An interesting result is that both {\\it E} and {\\it R} PNe have log(N/O) $< 0$ while only {\\it B} PNe show negative and positive values. Consequently, {\\it B} PNe are expected to have higher nitrogen values than the {\\it E} and {\\it R} PNe. Following that and a second sample of 35 PNe, n$_{\\rm e}$ is also found to be higher in {\\it B} PNe. Also, in all PNe morphologies PC2 appears to have a minimum at 0.89 and PNe's initial mass at 2.6 M$_{\\odot}$. 5--D diagrams between PCAs components and physical parameters are also presented. More results will follow while simple models will be applied in order to try to give a physical meaning to the components. ",
        "introduction": "\\label{sec:1} Planetary Nebulae (PNe) are powerful tools in the study of the evolutionary scenario of intermediate mass stars. They play an important role in the chemical enrichment history of the interstellar medium and many efforts have been devoted to determine the physical parameters of Galactic PNe (like T$_{\\rm e}$, n$_{\\rm e}$, T$^{\\star}$, L$^{\\star}$, distance, abundances; \\cite{journal1} and references therein). For this study a number of methods have been used either by using the observational results directly or by developing simulation models like CLOUDY \\cite{journal2}. Using the latter, the values of the physical parameters can be determined (i.e. \\cite{journal3}) but not any possible correlation among them. So far, only statistical methods had been used in order to find correlations between the parameters. However, there is a well--known technique (PCA) which can be used in order to study the possible correlation between these parameters. PCA technique uses a sample of observed parameters and creates a new sample of independent components which are linear combination from the previous parameters. ",
        "conclusions": ""
    },
    "0710/0710.3712.txt": {
        "abstract": "We apply a wind model, driven by combined cosmic-ray and thermal-gas pressure, to the Milky Way, and show that the observed Galactic diffuse soft X-ray emission can be better explained by a wind than by previous static gas models.  We find that cosmic-ray pressure is essential to driving the observed wind.  Having thus defined a ``best-fit'' model for a Galactic wind, we explore variations in the base parameters and show how the wind's properties vary with changes in gas pressure, cosmic-ray pressure and density.  We demonstrate the importance of cosmic rays in launching winds, and the effect cosmic rays have on wind dynamics.  In addition, this model adds support to the hypothesis of Breitschwerdt and collaborators that such a wind may help explain the relatively small gradient observed in $\\gamma$-ray emission as a function of galactocentric radius. ",
        "introduction": "Large-scale galactic outflows are usually considered in the context of starburst galaxies or Active Galactic Nuclei \\citep{VeCeBH2005}. These outflows are interesting not only intrinsically (what drives the outflow?) but for the interstellar and intergalactic media (how is the host galaxy affected, and what metals are ejected from the galaxy?).  To examine these questions, we have built a thermal and cosmic-ray driven wind model.  Our investigation into such models was first inspired by observational hints that the Milky Way may possess a kiloparsec-scale wind; this paper further explores that possibility. To motivate this study, we first introduce the observational evidence for a Galactic wind (\\S\\S\\ref{diffuseXrays} and \\ref{CRdensity}) and then introduce the cosmic-ray and thermally driven wind model (\\S\\ref{modelDef}).  In \\S\\ref{compareToObs}, we calculate the X-ray emission from this wind, and then compare it to the observations. After finding the best-fit wind model, we then explore the parameter space around that model (\\S\\ref{ParameterSurvey}) to understand more about how the wind is modified by varying the input and fit parameters.  Our conclusions are given in \\S\\ref{Conclusions}.  But first, observational hints for an outflow from our own Galaxy.  %To start applying this model, we look to a perhaps %unlikely first suspect: the Milky Way.  \\subsection{X-ray Observations}\\label{diffuseXrays}  The Milky Way does exhibit clues that it might drive a large-scale wind.  The first of these is an enhancement in the diffuse soft X-ray emission, stretching over the longitude range $-20^\\circ \\la l \\la 35^\\circ$ with an emission scale height in the southern Galactic hemisphere of $b \\sim -17^\\circ$ (see Fig.~\\ref{r45Intensity}).  This emission was first noted by \\citet{SnowdenEtAl95}, who modeled it with an isothermal plasma with a temperature $T = 4 \\times 10^6$~K, a midplane electron density of $n_{\\rm e, midplane} \\sim 3.5 \\times 10^{-3}$~cm$^{-3}$, and a midplane thermal pressure of $P_{\\rm g, midplane}/k \\sim 2.8 \\times 10^4$~cm$^{-3}$~K.  At approximately the same time, \\citet{BreitschwerdtSchmutzler94} suggested that the average \\textit{all-sky} X-ray emission (not only that emission in the region defined above) in all \\textit{ROSAT} bands might be explained by delayed-recombination in a large-scale cosmic-ray and thermally driven wind \\citep[see also][]{BreitschwerdtSchmutzler99}.  Later, \\citet{AlmyEtAl00} used intervening absorption to show that at least half of the central, enhanced X-ray emission lies more than 2~kpc from the sun \\citep[see also][]{ParkEtAl97, ParkEtAl98}.  Since that measurement was made in the Galactic plane, where the absorption is strongest, it was inferred that most of the emission observed at higher latitudes lies beyond that 2~kpc distance.  \\citet{AlmyEtAl00} also improved on previous modeling efforts: that work presents a model of the emission due to a static polytropic gas (with $\\gamma = 5/3$), and very importantly, includes the effects of known background components, such as the stellar background, extragalactic background, and an additional isotropic background (to fit high-latitude emission).  For comparison, their model had a central temperature of $T_0 = 8.2 \\times 10^{6}$~K, a central electron density of $n_{\\rm e} = 1.1 \\times 10^{-2}$~cm$^{-3}$, and a central pressure of $P_{\\rm g,0} = 1.8 \\times 10^5$~cm$^{-3}$~K.  We will compare our results with this static polytrope model to investigate whether a wind model for this emission is feasible. \\begin{figure*}[] \\begin{center} \\includegraphics[width=14cm]{f1_color.ps} \\caption{X-ray emission at $3/4$~keV (the ``R45 band'') as seen by \\textit{ROSAT} \\citep{SnowdenEtAl97}.  These observations suggest a ``Galactic X-ray Bulge'', seen most clearly in the southern Galactic Hemisphere, and stretching over the Galactic longitude range, $l$, from $|l| \\la 30^\\circ$ and down to approximately $-15^\\circ$ in Galactic latitude.  This paper asks whether the X-ray bulge in the southern Galactic Hemisphere can be explained with a combined thermal and cosmic-ray driven wind. \\label{r45Intensity}} \\end{center} \\end{figure*}   \\subsection{Cosmic Ray Source Density}\\label{CRdensity}  Another indicator of a Galactic wind comes from measurements of the density of cosmic rays as a function of Galactocentric radius, $R$. The source density of cosmic rays can be determined via $\\gamma$-ray emission: the production of $\\gamma$-ray photons with energies exceeding about 50~MeV is dominated by collisions of cosmic rays with gas in the interstellar medium \\citep{BlBlHe84}.  Since the galaxy is largely transparent to such high-energy photons, the $\\gamma$-ray emissivity at those energies yields the cosmic ray source density.  If cosmic rays are produced in supernovae remnants, then since the source density of supernovae remnants seems to increase with decreasing $R$, the cosmic-ray source density should increase as well. However, it has been known for some time \\citep[e.g,][]{Bloemen89} that the inferred cosmic-ray source density is relatively flat, compared to the supernova density, as a function of $R$.  There has, however, been some debate about whether supernovae remnants are an accurate tracer \\citep[since those surveys are subject to various selection effects; see, e.g.,][and references therein]{StrongEtAl04}.  Recent surveys of the pulsar population \\citep{LorimerEtAl06} also show that the pulsar source density increases towards the center of the Galaxy, as shown in Figure~\\ref{crSourceDensity}.  This is true irrespective of the model of how $n_{\\rm e}$ varies in the disk, although the magnitude of the pulsar population gradient with $R$ depends strongly on the $n_{\\rm e}$ model.  So, there remains a mismatch between the observed source density of cosmic ray ``producers'' and the cosmic rays themselves.  It has already been pointed out that the observed slow rise in cosmic rays may be due to a wind emerging from the disk, advecting cosmic rays outwards \\citep{BloemenEtAl93, BrDoVo02}.  In the case of \\citet{BloemenEtAl93}, a wind model was applied to the entire Galactic disk; as a result, only a very slow wind was found to be compatible with the inferred cosmic-ray source density.  In contrast, \\citet{BrDoVo02} applied their cosmic-ray and thermally driven wind model, where the wind velocity varied as a function of radius and height; they also took into account anisotropic diffusion.  With this model, a small radial gradient in the cosmic ray source density could be explained.  An alternate explanation for this slow change in the cosmic ray population with $R$ was proposed by \\citet{StrongEtAl04}, who found that a radial variation in the $W_{\\rm CO}$-to-$N(H_2)$ ratio by a factor of 5 to 10 could explain the $\\gamma$-ray observations.  In this paper we primarily address the question of the origin of the diffuse, soft X-ray background emission; we will, however, concentrate on a large-scale wind model, keeping in mind its possible application to the cosmic ray source density. \\begin{figure}[h] \\begin{center} \\includegraphics[width=6cm,angle=-90]{f2.ps} \\caption{Comparison of two different calculations of the pulsar population as a function of Galactocentric radius \\citep{LorimerEtAl06} vs. the cosmic ray source density implied from the observed $\\gamma$-ray emissivity \\citep{StrongEtAl04}.  The two different curves for the pulsar distribution result from assuming a smooth distribution of $n_{\\rm e}$ in the Galaxy \\citep{LyMaTa85}, or a clumped distribution, using \\citet{CordesLazio02} and \\citet{FaKa06}; for details, see \\citet{LorimerEtAl06}.  The fact that the cosmic-ray distribution does not seem to follow the pulsar population has been known for some time \\citep{Bloemen89}, but there is no consensus on the reason.  A cosmic-ray and thermal pressure-driven wind may help explain the cosmic-ray source population. \\label{crSourceDensity}} \\end{center} \\end{figure}  %Both the diffuse X-ray emission and the slow change in cosmic-ray %population towards the Galactic center hint at the possibility of a %large-scale wind driven by both thermal and cosmic-ray pressure.  We %next outline a wind model including these components. ",
        "conclusions": "We have implemented a simplified cosmic ray- and thermally-driven wind and have used it to try to explain the soft, diffuse X-ray emission seen towards the Galactic Center.  We find that such a wind can indeed match the observed averaged X-ray emission quite well, and in fact fits demonstrably better than the static polytrope model of \\citet{AlmyEtAl00}.  It is important to note that this wind is approximately equally powered by both cosmic rays and thermal pressure: cosmic rays are important in helping this relatively cool wind escape from the Galactic potential.  It is also quite interesting that this wind does not require excessive thermal or cosmic-ray pressures (both pressures are not extreme compared to what has been estimated for the inner Milky Way), nor does this simple model require much more energy than the standard inferred supernova rate implies. Taking this result at face value, such a wind would be very important to the ``ecology'' of the Milky Way due to the high mass loss rate of 2~$M_{\\odot}$~yr$^{-1}$.  In addition, such a wind would also play an important role in removing angular momentum from matter in the Galactic disk and allowing matter to move radially inward \\citep{Zirakashvili1996}.  At the least, this shows that such wind models should be considered further for the Milky Way; they may be able to explain at least a substantial fraction of the observed soft X-ray emission.  Further, as other researchers have already shown \\citep{BrDoVo02}, such winds can also be used to explain the unexpectedly slow rise in $\\gamma$-ray emission towards the center of the Galaxy; this work therefore gives independent support to the Galactic wind hypothesized in \\citet{BrDoVo02}.  We have not yet calculated the effect of the best-fit wind model on the cosmic-ray distribution, but a simple approximation will investigate if this wind is removing cosmic rays at too high a rate.  We reason as follows.  The best-fit wind model shown here has a cosmic-ray advective timescale of $\\sim 4.3 \\times 10^6$~years.  Therefore, supernova in the disk must resupply the cosmic-ray pressure on that timescale.  We know that the approximate total supernova energy in the disk (from observations, see \\S\\ref{totalEnergyFlux}) is $\\sim 1.7 \\times 10^{41}$~ergs~s$^{-1}$. So, if some fraction of this energy, $\\epsilon_{\\rm CR}$, is given to cosmic rays, and distributed over the volume occupied by hot gas where the wind is launched (over the Galactocentric radius range of $1.5$ to $4.5$~kpc, and a height range of $\\pm$2~kpc, to be conservative), that energy density should be similar to the cosmic-ray pressure required to launch the wind.  Calculating the resultant buildup of $P_{\\rm c}$ over $4.3 \\times 10^6$ years at the observed SN rate, we find $P_{\\rm c} \\sim \\epsilon_{\\rm CR} \\cdot 6.9 \\times 10^{-12}$~dyne~cm$^{-2}$ or $\\sim \\epsilon_{\\rm CR} \\cdot 5.0 \\times 10^4$~K~cm$^{-3}$.  This is actually of order the $P_{\\rm c}$ that the best-fit model requires (see Table~\\ref{bestFitParams}), although it would require $\\epsilon_{CR} \\sim 0.6$ to duplicate the best-fit $P_{\\rm c}$, which is relatively high.  Still, this simple, rather conservative calculation shows that the high $\\dot{M}$ wind shown here does not remove cosmic rays much more quickly than they can be replenished by the normal SN rate in the Galaxy, although the removal rate of cosmic rays is certainly non-negligible, and would affect the density of cosmic rays towards the Galactic center.  Of course, a more detailed calculation is required (with a more detailed wind model), but this again shows that the best-fit wind model is at least feasible, and would have a significant but not destructive effect on the Galaxy's cosmic-ray density.  \\subsection{Future Improvements}  More detailed models are clearly needed; there are a few concerns about the current model that could be addressed with more realistic wind models.  For instance, we have assumed uniform density at the base of the wind over the area of the disk from $R=1.5$~kpc to 4.5~kpc, which leads to a mass outflow rate of order 2~$M_{\\odot}$~yr$^{-1}$.  This seems quite high, but in the context of a more detailed model with variations in density within the wind, we might expect that the $n^2$ weighting of emissivity would favor overdensities, and allow an inhomogeneous wind to better reproduce the observations with a smaller mass outflow rate.  In addition, it is possible that other effects limit the gas to velocities below those in this simple model; drag effects may slow down the wind \\citep[e.g.,][]{EM07}, and could lead to some of the gas forming part of the Galactic fountain \\citep{Bregman99}.  On the other hand, turbulence may be an important additional source of energy for the wind, but we have not included such an input in this work.  Also, the effects of distributed mass loading, which could be relatively easily incorporated into this model, have been ignored so far; such mass loading would be very important to the emission properties of the wind, and the potential observability of such winds in many galaxies. Finally, note that we have used \\citet{AG89} abundances; if the wind starts out with super-solar abundances (particularly in oxygen), a smaller mass outflow rate would be required.  The cosmic ray physics in this wind is still quite simple.  As mentioned previously, we have assumed zero diffusivity of the cosmic rays throughout the wind \\citep[c.f.,][]{BrMcVo1993}.  We have also neglected thermal conductivity (our calculations show that conductivity may be important for $z \\la 350$~pc).  Including both of these effects would be essential to future progress for this wind model: as mentioned previously (see \\S\\ref{totalEnergyFlux}), including the effects of conduction may help lower the energy requirements in the wind, bringing the total power of the wind closer to that of the total inferred supernovae power in the Galactic disk below the wind.  We also note that we are using a model of the gravitational potential that we know overestimates the potential in the Milky Way (see \\S\\ref{modelDef}).  Adopting the more detailed potential model of \\citet{DehnenBinney98} may allow a lower total-energy wind to duplicate the soft X-ray observations.  Finally, we have also assumed the hydrodynamic model of \\citet{McKenzieWebb84} and \\citet{BrMcVo1991} for the interaction of cosmic rays with \\alfven waves and the gas.  To further examine this model, we will next consider the effects of higher cosmic-ray fluxes \\citep{Zweibel2003} and apply it in other settings.  \\subsection{Future Tests}  How can we further test this model?  In analogy to early studies of the solar wind, this outflow may impact clouds in the vicinity of the galaxy, perhaps causing ``comet-tail'' extensions to high velocity clouds above the plane of the Milky Way.  The formation of such ``tails'' would depend on the velocity of the wind.  This has been studied in some detail before \\citep{BenjaminCox02}, but should be reconsidered in the context of the predictions of these winds.  In addition, it may be possible to study the kinematic impact of this wind on, for instance, the Magellanic Stream (A. Burkert, personal communication).  Another way to test the model would be to compare the wind with absorption columns and emission spectra towards the center of the Galaxy.  Concerning absorption measurements, recent \\textit{Chandra} observations \\citep{FutamotoEtAl04} towards the low-mass X-ray binary 4U 1820-303 show significant columns in \\al{O VII}, \\al{O VIII}, and \\al{Ne IX}.  As this X-ray binary is located within ten degrees of the Galactic center, at a distance of approximately 7.6~kpc, it seems an ideal target.  However, similar absorption columns are found towards objects on sightlines that do not intercept the base of the wind; for instance, Mrk~421 shows similar absorption columns \\citep{FutamotoEtAl04}, but is far out of the plane.  This leads us to conclude that the observed absorption is somewhat local to the Sun's position, and, as such, this absorption does not constrain the wind model.  However, recent Suzaku measurements of emission at various latitudes along different sightlines towards the Galactic center (Rocks, 2008, in preparation) may help constrain the properties of the wind.  Finally, these winds, including their important cosmic-ray component, will also emit synchrotron radiation.  We are now calculating the synchrotron emission expected from these wind models (Schiller et al., 2008, in preparation).  This will allow exploration of the wind's synchrotron emission as compared to recent models of Galactic synchrotron which begin to map the three-dimensional cosmic-ray emissivity \\citep{NordEtAl06}.  Looking at a wider field of application, such wind models may also be quite important in application to starburst galaxies \\citep[e.g.,][]{GallagherSmith05, SDR06} and dwarf galaxies.  The general applicability of these kinds of models to starbursting galaxies has been shown by \\citet{Breitschwerdt03} in fitting cosmic-ray and thermally driven wind models to NGC 3079."
    },
    "0710/0710.5535_arXiv.txt": {
        "abstract": "We present new \\emph{Chandra} observations of a high redshift ($z$$\\sim$1) galaxy cluster discovered in the Red-Sequence Cluster Survey (RCS): \\rcsB.  X-ray luminosity measurements and mass estimates are consistent with $L_X$--$T_X$ and $M_\\delta$--$T_X$ relationships obtained from low-redshift data.  Assuming a single cluster, X-ray mass estimates are a factor of $\\sim$10--100 below the red-sequence optical richness mass estimate.  Optical spectroscopy reveals that this cluster comprises two components which are close enough to perhaps be physically associated.  We present simple modeling of this two-component system which then yields an X-ray mass and optical richness consistent with expectations from statistical samples of lower redshift clusters.  An unexpectedly high gas mass fraction is measured assuming a single cluster, which independently supports this interpretation.  Additional observations will be necessary to confirm the excess gas mass fraction and to constrain the mass distribution. ",
        "introduction": "There is considerable interest in measuring the evolution of the galaxy cluster mass function with redshift.  In principle, such a measurement can test structure formation and constrain the cosmic expansion history.  One of the observational challenges in such an undertaking is to identify all clusters in a particular volume; the other is to reliably infer cluster mass from observable cluster properties, over a range of cluster redshifts.  Recent optical surveys, such as the Red-Sequence Cluster Survey \\citep[RCS,][]{GY05}; and the Sloan Digital Sky Survey \\citep[e.g.,][]{Kea07} have found large numbers of clusters at moderate to high redshift, with $0.2\\lesssim z\\lesssim 1$ and $0< z \\lesssim 0.3$, respectively.  At the high redshift extent, where spectroscopic information can be difficult to obtain on a survey scale, various optical properties, such as richness or optical luminosity, can be taken to be mass proxies, and have been shown to be strongly correlated with observables understood to be related to cluster mass, such as galaxy velocity dispersion and X-ray temperature \\citep[e.g.,][]{YE03, Lea03, Pea04, Lea06}.   In order to use optical richness (or any other optical property) as a mass proxy, understanding the evolution of the mass-richness relation is critical to extracting cosmological information, since any evolution in the relation must be separated from the evolution of the mass function.  In particular, some care must be taken when attempting to extend the mass/observable relation to high redshift as many of these relations are computed using local data.  Previous work has shown that optically-selected, high-redshift clusters at a fixed optical richness are typically lower in temperature and underluminous in X-rays when compared to local clusters \\citep{Hea05,Gea04,Lea02,Dea01}.  Since these clusters generally follow the local luminosity-temperature relationship, the low temperatures and luminosities are interpreted as an indication that the cluster galaxies are not in virial equilibrium with the X-ray emitting gas.  However, it is important to account for the significant scatter in the relation between optical richness and X-ray luminosity, as an Eddington-like bias will scatter points of lower than expected X-ray luminosity into an optically-selected sample \\citep[e.g.,][]{Bea94}.  Observing these high redshift clusters is thus an avenue to study how clusters approach equilibrium.  \\rcsB\\ is part of a larger sample of high redshift, high richness RCS clusters selected for X-ray follow up observations.  X-ray properties of the sample as a whole are described in a forthcoming paper by \\citet{Hea07}.  Here we present \\emph{Chandra} X-ray observations and optical spectroscopy of this apparently rich, X-ray underluminous object.  In \\S 2 we describe the X-ray data reduction and analysis.  In \\S 3 we compare mass estimates from the X-ray data to the known mass-temperature and mass-richness relations in literature.  In \\S 4 we describe the optical spectroscopy, in \\S 5 we discuss our results and we summarize in \\S 6.  Throughout the paper we take a standard $\\Lambda$CDM cosmology with $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m = 0.3$ and $\\Omega_\\Lambda = 0.7$.  At the cluster redshift, $z=0.9558$, 1\\arcsec\\ corresponds to a metric distance of 6.64 kpc.  Errors are indicated with 90\\% confidence intervals unless otherwise noted. ",
        "conclusions": "\\rcsB\\ has an inferred X-ray mass which is one to two orders of magnitude lower than the mass indicated by its red-sequence richness, assuming that the X-ray emission is from a single, spherical, isothermal gas distribution and that the richness is associated with a single cluster ($M_{200,X}=4.6\\pm^{6.0}_{1.7}\\times10^{13}$ M$_{\\odot}$ versus $M_{200,B_{gcR}}=1.1\\times10^{15}\\pm0.46$ dex M$_{\\odot}$).  This is a significant discrepancy for two reasons.  First, as mentioned before, richness is well correlated on average with other observable properties.  Secondly, clusters with masses above $10^{15}$ M$_{\\odot}$ are quite rare, especially at $z\\sim1$.  This means that if the richness mass estimate of \\rcsB\\ were truly indicative of its size, it would be among the largest overdensities in the observable universe, while also being very underluminous.  Yet the X-ray properties, which suggest a much less massive object, are consistent with the expected $M_{\\delta}$--$T_X$ and $L_X$--$T_X$ relations, meaning that the temperature of the plasma is consistent with its observed distribution.  One possible interpretation of these results is that the red-sequence galaxies used to measure the optical richness and the X-ray emitting gas trace different volumes.  Specifically, while the X-ray luminous plasma is confined within one or more deep gravitational potential wells, the galaxy population may occupy a more extended region which is not yet dynamically relaxed.  The existence of such a structure in \\rcsB\\ is supported by its two-peaked radial velocity distribution.  Merging and dynamically active clusters are expected and observed to be increasingly common at higher redshifts \\citep[e.g.,][]{Jea05}.  Thus any cluster sample will include an increasing fraction of dynamically young systems at higher redshift.  Furthermore, N-body simulations indicate that a significant fraction of cluster samples selected using broadband color discrimination will be systems whose galaxy members are not predominantly associated with a single, large potential well, but rather are spread amongst a number of smaller, but still physically associated, dark matter halos distributed along the line-of-sight \\citep{Cea07}.  Simulations of the RCS technique on mock catalogs tuned to reproduce observables such as the observed galaxy color distribution and the two-point correlation function show that we expect false-positive cluster detections to occur at a frequency of $\\sim$5\\% in the RCS \\citep{G02}.  This agrees with initial results based on a small number of clusters \\citep{Gilbank_ea07,Bea07}.  Since the red-sequence galaxies used in the RCS richness measurements form very early ($z\\gtrsim2$) in high density regions, overdensities of red-sequence galaxies can be associated with very large structures which may not be relaxed during the observed epoch.  The X-ray luminous plasma, on the other hand, will be confined to one or more gravitational potential wells whose size is limited by the collapse timescale.  In the hierarchical collapse paradigm of a $\\Lambda$CDM universe, the large, virialized wells seen in local clusters do not develop until long after the galaxies have formed.  In this scenario, for any evolving cluster, there is likely to be an epoch at which the apparent richness overestimates the virialized mass.  Moreover, the misinterpretation of the X-ray emission from a complex, dynamically young cluster as a single, virialized structure can lead to an overestimate of the inferred gas mass fraction and an underestimate of the total mass, which can be seen as follows.  Let us consider two models for the X-ray emission seen in \\rcsB.  Our one-component model (Model I) is the set of assumptions used in our X-ray analysis thus far:  X-ray emission is from a single, spherically-symmetric distribution of gas in hydrostatic equilibrium with a gravitational potential which is well-described by a $\\beta$-model.  The two-component model (Model II) assumes that there are two nominally identical (i.e.~with the same values of $\\beta$, $r_0$ and $T_X$), virialized gas distributions, separated along the line of sight, and that each of the two gas distributions is responsible for half of the measured X-ray flux.  Each component is assumed to obey the local $M_{\\delta}$--$T_X$ and $L_X$--$T_X$ relation.  Because it assumes the flux is halved between the two components, the luminosity inferred per cluster in Model II is half that of Model I.  The luminosity scales as the square of the gas mass, so since the spatial distribution is the same for each cluster, the inferred gas mass and thus also the electron density will be reduced by a factor of $\\sqrt{2}$ in each of the clusters relative to Model I.  Since our total mass estimate for each cluster depends only on the temperature of the X-ray emitting gas and its distribution in the plane of the sky, each of the two components in Model II will have the same total mass as the single component in Model I.  This means that the inferred gas mass fraction will be a factor of $\\sqrt{2}$ lower using Model II versus Model I.  The total mass, total gas mass and gas mass fraction results using each of the two models can be found in Table \\ref{masstable}.  Similarly, $n$ identical components along the line of sight would reduce the inferred gas mass per component and the overall gas mass fraction by a factor of $\\sqrt{n}$ in addition to reducing the luminosity per component by a factor of $n$.  In principle, luminosity measurements and the $L_X$--$T_X$ relation constrain the number of components allowed, but our luminosity errors are too large to distinguish between Model I and Model II.  \\citet{ME01}, using a suite of hydrodynamic cluster simulations, found that during merger events the ICM spectral fit temperature will underestimate the mass-weighted ICM temperature by $\\sim20\\%$, because the cool, denser inflowing gas will dominate the emission over the gas already heated by the merger.  This means that for a dynamically young system, such as \\rcsB, where even the most virialized components are likely to have undergone recent mergers, X-ray mass estimates for those clusters may still be below the virialized masses by several tens of percent.  In a recent detailed study of the cluster Cl 0024+17, \\citet{Jea07} found a similar underestimation of the true mass distribution due to the assumption of a single virialized mass structure rather than two components extended along the line of sight and in a state of ongoing dynamic interaction.  In summary, the assumption of spherical symmetry and virial equilibrium for a cluster system containing extended line-of-sight structure and dynamic evolution can lead to overestimation of the gas mass fraction and underestimation of the total mass.  The baryon mass fraction in relaxed galaxy clusters is expected to be a universal quantity \\citep[e.g.,][]{Wea93,Vea03}, suggesting that the comparison of cluster gas mass fractions inferred by assuming spherical symmetry to canonical cluster values might be used to infer the presence of line-of-sight structure.  As noted by \\citet{Wea93}, the cosmological baryon mass ratio, $\\Omega_b/\\Omega_m$, provides an upper limit to the total gas mass fraction.  This ratio from the WMAP three-year data results is $\\Omega_b/\\Omega_m=0.175\\pm0.012$ \\citep{Sea07}.  Galaxy cluster samples at high redshift are expected to have lower cluster gas mass fractions than local samples \\citep[e.g.,]{Hea07, Sea05, Eea04}.  Also, a positive correlation between temperature and gas mass fraction has been observed \\citep[e.g.,][]{Vea06, Sea03}.  This suggests that a high-redshift cluster with a low X-ray temperature, such as \\rcsB, ought to have a correspondingly low gas mass fraction.  Instead, we observe a very high gas mass fraction assuming that it is a single, spherical matter distribution, especially measured within a large radius.  At $r_{500}$,  we find $f_{gas,500}=0.17\\pm^{0.07}_{0.08}$.  We observe cluster emission to approximately this radius, so this gas mass fraction estimate is unlikely to have the extrapolation errors that the measurement at $r_{200}$ might.  Assuming that Model II is more appropriate for the physical state of \\rcsB\\ lowers the gas mass fraction towards what is expected.  Would this pair of clusters be physically associated?  As an order-of-magnitude argument, we note that the richness mass ($\\sim$1$\\times10^{15}$ M$_{\\odot}$) is roughly equal to that of a sphere of radius $\\sim$13 Mpc with the cosmic density of $z=0.96$.  Optical spectroscopy shows that the redshift separation of the two components in \\rcsB\\ is $\\Delta z\\approx 0.005$, or a physical separation of about 12 Mpc in the Hubble flow.  This rough agreement supports the conclusion that \\rcsB\\  is an incompletely-virialized system which is still approaching equilibrium, and that the galaxy distribution traces unvirialized matter extended along the line-of-sight, as well as the virialized matter traced by the X-ray gas.  The free-fall collapse time for the two components given by the richness mass is approximately equal to the lookback time, indicating that this cluster would be nearly virialized by about the present epoch.  Additional observations are required to determine the physical state of \\rcsB.  Due to the low number of source photons from this cluster, there are large errors in the X-ray luminosity and temperature.  Improvements on both would increase the precision of the X-ray mass estimates and the gas mass fractions.  An independent mass estimate, such as might be obtained from weak lensing, could also help distinguish which mass estimate (richness or X-ray) is more appropriate, as well as giving additional insight into the distribution of matter.  A weak lensing measurement, though difficult, would be particularly interesting for this system if it could trace the spatial distribution of cluster mass perpendicular to the line of sight, possibly revealing more substructure.  Additional spectroscopy, currently underway, will help to further constrain line-of-sight substructure."
    },
    "0710/0710.0898_arXiv.txt": {
        "abstract": "{ We report the detection of transits by the $3.1 M_{\\rm Jup}$ companion to the V=8.17 G0V star \\object{HD 17156}. The transit was observed by three independant observers on Sept. 9/10, 2007 (two in central Italy and one in the Canary Islands), who obtained detections at confidence levels of 3.0~$\\sigma$, 5.3~$\\sigma$, and 7.9~$\\sigma$, respectively. The observations were carried out under the auspices of the Transitsearch.org network, which organizes follow-up photometric transit searches of known planet-bearing stars during the time intervals when transits are expected to possibly occur. Analyses of the 7.9~$\\sigma$ data set indicates a transit depth $d=0.0062 \\pm 0.0004$, and a transit duration $t=186 \\pm 5$ min. These values are consistent with the transit of a Jupiter-sized planet with an impact parameter $b=a \\cos{i}/R_{\\star} \\sim 0.8$. This planet occupies a unique regime among known transiting extrasolar planets, both as a result of its large orbital eccentricity ($e=0.67$) and long orbital period ($P=21.2 {\\rm d}$). The planet receives a 26-fold variation in insolation during the course of its orbit, which will make it a useful object for characterization of exoplanetary atmospheric dynamics. } ",
        "introduction": "During the past several years, the discovery rate of transiting planets has begun to increase rapidly, and twenty transiting planets with secure characterizations are currently known \\footnote{ Extrasolar Planets Encyclopedia {\\tt http://exoplanet.eu}}. This aggregate consists mostly of short-period hot-Jupiter type planets, with prototypical examples being HD~209458b \\citep{2000ApJ...529L..45C,2000ApJ...529L..41H} and HD~189733b \\citep{2005A&A...444L..15B}. These planets tend to have $M\\sim 1M_{\\rm Jup}$, $2 {\\rm d}<P<5{\\rm d}$, and tidally circularized orbits.  In the past year, two remarkable discoveries have significantly extended the parameter space occupied by known transiting planets. HD~147506b \\citep{2007arXiv0705.0126B} with $M=8.04 \\, M_{\\rm Jup}$ is by far the most massive planet known to exhibit transits. It also has the longest orbital period (5.63 days) and a startlingly large orbital eccentricity, $e\\sim 0.5$. At the other end of the mass scale, Gl~436b \\citep{2004ApJ...617..580B,2007A&A...472L..13G} has $M=0.07 M_{\\rm Jup}$, a 2.64 day orbital period, and an eccentricity $e=0.15\\pm0.01$ \\citep{2007arXiv0707.2778D}. These two planets straddle more than a hundred-fold difference in mass, and their significant non-zero eccentricities are also capable of imparting important information.  At present, infrared observations of transiting extrasolar planets by Spitzer present an incomplete and somewhat contradictory overall picture. It is not understood how the wind vectors and temperature distributions on the observed planets behave as a function of pressure depth, and planetary longitude and latitude. Most importantly, the effective radiative time constant in the atmospheres of short-period planets remains unmeasured, and as a result, dynamical calculations of the expected planet-wide flow patterns \\citep{2003ApJ...587L.117C,2005ApJ...629L..45C,2005ApJ...618..512B,2007ApJ...657L.113L,2007arXiv0704.3269D} have come to no consensus regarding how the surface flow should appear. This lack of agreement between the models stems in large part from the paucity of unambiguous measurements of the radiative time constant in the atmosphere. What is needed, is a transiting planet with both a long-period orbit and a large orbital eccentricity. If such a planet were known, then one could use Spitzer to obtain infrared time-series photometry of the planet during the periastron passage. The transit guarantees knowledge of both the geometric phase function and the planetary mass. This information would in turn allow a measured rate of increase in flux to inform us of the planet's atmospheric radiative time constant in the observed wavelength regime.  The orbital periods of the known transiting planets are all significantly shorter than 6 days. This bias is due both to the intrinsically lower geometric probability of transit as one moves to longer periods, and also to the fact that ground-based wide-field transit surveys that rely on photometry folding become very significantly incomplete for planets with orbital periods longer than 5 days. If one wants to detect longer-period transiting planets from the ground, a more productive strategy is to monitor known RV-detected planet-bearing stars at the times when the radial velocity solution suggests that transits may occur. This strategy has the further advantage of producing transits around stars that tend to be both bright and well-suited for follow-up observations.  Long-period transiting planets present an ideal observing opportunity for small telescope observers. \\cite{2003PASP..115.1355S} have suggested that a global network of telescopes, all capable of $\\sim1\\%$ photometry can easily outperform a single large telescope in terms of efficiency of transit recovery. Since inception in 2002, the Transitsearch.org network has conducted follow-up searches on a number of intermediate-period planets (see e.g. \\citealt{2006ApJ...653..700S}).  The Doppler-based discovery of HD~17156b was recently published by the N2K consortium \\citep{2007arXiv0704.1191F}. The planet has M$\\sin i = 3.12{\\rm M_{Jup}}$, with $P=21.22$ days and $e\\sim0.67$. \\cite{2007arXiv0704.1191F} report that the V=8.17 G0V host star has $M= 1.2 M_\\odot$ and $R= 1.47 R_\\odot$. The planet's semi-major axis $a=0.15 {\\rm AU}$ thus indicates a periastron distance of $a_{min}=0.0495 {\\rm AU}=7.2 R_{\\star}$. A best fit to the radial velocities indicates longitude of periastron $\\omega=121 \\pm 11^{\\circ}$. The orbital orientation is favorable, yielding an a-priori geometric transit probability of $P\\sim13$\\%.  In their discovery paper, \\cite{2007arXiv0704.1191F} reported 241 individual photometric measurements obtained over a 179 day interval, and with a mean dispersion $\\sigma=0.0024$ mag. No significant rotation-induced periodicity was seen. Together, the observations sampled approximately 25\\% of the $1-\\sigma$ transit window, and no evidence for a transit was observed. After the \\cite{2007arXiv0704.1191F} discovery paper was made public, the star was added to the Transitsearch.org candidates list\\footnote {http://207.111.201.70/transitsearch/dynamiccontent/candidates.html} and observers throughout the Northern Hemisphere were repeatedly encouraged to obtain photometry of the star\\footnote{see www.oklo.org}. The first available window of opportunity occurred on 9/10 Sept., 2007, with the transit midpoint predicted to occur at HJD $2454353.65 \\pm 0.30$. ",
        "conclusions": "The detection of transits by a planet with a three-week orbital period demonstrates the utility of ad-hoc networks of small telescopes for obtaining photometric follow-up of planets whose orbital parameters have been determined via Doppler radial velocities. Indeed, the transits of HD~17156b offer a plethora of interesting opportunities for follow-up observations.  With its high orbital eccentricity and small periastron distance, HD~17156b appears to bear a curious kinship to HD~80606b, HD~147506b, and HD~108147b. All three of these planets occupy a locus of the $a$--$e$ plane where they should actively be undergoing tidal dissipation, and therefore they should be generating significant quantities of excess interior heat. Our measurement indicates that tidal heating is not significantly inflating the planetary radius. The nominal $R=1.1 R_{\\rm Jup}$ radius predicted by baseline models (e.g. those of \\citealt{2003ApJ...592..555B}) is confirmed by our observations.  Follow-up photometric measurements during future transits will allow a more accurate determination of the orbital inclination of HD~17156b. An improved value for $i$, in turn, will generate an accurate assessment of likelihood that the planet can be observed by Spitzer in secondary transit, and will enable a much-improved constraint on the still-uncertain radius of the parent star. In the event that secondary transits can be observed, a direct measurement of the excess tidally generated luminosity from the planet is a distinct possibility (see e.g. \\citealt{2007arXiv0707.2778D}).  As a consequence of its highly eccentric orbit, HD~17156b experiences a 26-fold variation in insolation during the 10.6 day interval between periastron and apoastron. This extreme radiative forcing may drive interesting, and potentially observable dynamical atmospheric flows on the planet \\citep{2007ApJ...657L.113L}. The large tidal forces experienced during periastron have almost certainly forced the planet into pseudo-synchronous rotation (e.g. \\citealt{1968ARA&A...6..287G,1981A&A....99..126H,2005MNRAS.364L..66P}). Rotationally induced modulation in the infrared light curve following periastron is potentially observable, and may be of great utility in selecting between the current divergent predictions for the actual value of the pseudo-synchronous spin frequency.  HD~17156b is quite massive, as is often the case for planets orbiting one member of a binary pair \\citep{2007A&A...462..345D}, and the eccentricity is large. These characteristics favor a formation scenario involving migration and/or dynamical evolution in the presence of a sufficiently close external perturber. Such a perturber could be either a companion star or an additional planet(s) in the system.  For some of the close-in planets with $m \\sin i > 1.5~M_{\\rm Jup}$ orbiting single stars, there are already indications of a history of significant dynamical perturbations. For example, HD~118203b, HD~68988b and HIP~14810b all have anomalously high eccentricities that may be indicative of additional perturbing companions, perhaps with masses below (or periods longer than) the threshold of immediate radial velocity detection (this is certainly the case for HD~68988b and HIP~14810b, which both have long-period planetary companions \\citealt{2006ApJ...646..505B}). Due to mutual perturbations, the eccentricities of the bodies in the precursor system may have grown to the point where crossing orbits were achieved. Repeated close encounters among the planets would have then generated a period of chaotic evolution that typically terminates with the ejection of one planet on a hyperbolic trajectory \\citep{2002Icar..156..570M}.  Alternately, a stellar companion could also effectively trigger dynamical evolution or instability in a precursor system, eventually leading to the current configuration (for some examples, see \\citealt{2007A&A...467..347M,2007A&A...472..643M,2003ApJ...589..605W}). With reference to a stellar companion, a quick inspection of POSSI, POSSII, and 2MASS images do not reveal any clear association between faint field stars and HD~17156. The only potentially interesting source is 2MASS~02494068+7144583. It shows an appreciable proper motion, but in the opposite direction of the proper motion of HD~17156. The star lies 22.2\\arcsec~from HD~17156; if they are at the same distance, the apparent separation is $\\sim$1740 AU.  A combination of continued radial velocity monitoring of HD~17156, in conjunction with accurate measurements of successive transit midpoints, gives hope for the detection and accurate characterization of additional bodies in the system via a novel set of constraints."
    },
    "0710/0710.5473_arXiv.txt": {
        "abstract": "We identify a large sample of isolated bright galaxies and their fainter satellites in the 2dF Galaxy Redshift Survey (2dFGRS). We analyse the dynamics of ensembles of these galaxies selected according to luminosity and morphological type by stacking the positions of their satellites and estimating the velocity dispersion of the combined set. We test our methodology using realistic mock catalogues constructed from cosmological simulations. The method returns an unbiased estimate of the velocity dispersion provided that the isolation criterion is strict enough to avoid contamination and that the scatter in halo mass at fixed primary luminosity is small. Using a maximum likelihood estimator that accounts for interlopers, we determine the satellite velocity dispersion within a projected radius of 175~\\hkpc. The dispersion increases with the luminosity of the primary and is larger for elliptical galaxies than for spiral galaxies of similar b$_{\\rm J}$ luminosity. Calibrating the mass-velocity dispersion relation using our mock catalogues, we find a dynamical mass within 175~\\hkpc\\ of $\\rm M_{175}/h^{-1}M_\\odot \\simeq 4.0^{+2.3}_{-1.5} \\times 10^{12}\\,(L_{\\bjm}/L_*)$ for elliptical galaxies and $\\rm M_{175}/h^{-1}M_\\odot \\simeq 6.3^{+6.3}_{-3.1} \\times 10^{11}\\,(L_{\\bjm}/L_*)^{1.6}$ for spiral galaxies. Finally, we compare our results with recent studies and investigate their limitations using our mock catalogues. ",
        "introduction": "The view that galaxies are surrounded by large dark matter halos dates back more than 30 years to the pioneering study of the rotation curve of M32 by Rubin \\& Ford (1970). Extended galactic halos are, in fact, a generic feature of the cold dark matter model of galaxy formation (Blumenthal et al 1984, Frenk et al 1985), but this fundamental theoretical prediction has limited observational support.  Zaritsky \\etal~(1993) attempted to measure the mass and extent of dark matter halos by analysing the dynamics of satellite galaxies found around `isolated' galaxies. Since galaxies generally have only a few detectable satellites, they used a method that consists of stacking satellites in a sample of primaries of similar luminosity. In spite of the small size of their relatively inhomogeneous sample, Zaritsky \\etal~(1993) were able to detect massive halos around isolated spiral galaxies extending to many optical radii. Having nearly doubled their satellite sample to 115 members, Zaritsky \\etal (1997b) confirmed their earlier claims including a puzzling lack of correlation between the velocity dispersion of the (stacked) satellite system and the luminosity of the primary.  More recently, McKay \\etal~(2002) performed a similar analysis on data from the Sloan Digital Sky Survey (SDSS; York \\etal~2000). They compared mass estimates derived from satellite dynamics to those derived from weak lensing analyses of the same data (McKay \\etal 2001). With a much larger sample than that of Zaritsky et al (1997), they were able to detect a correlation between satellite velocity dispersion and primary luminosity. This trend was confirmed by Prada \\etal~(2003) who also used SDSS data. Although they are both based on SDSS data, these two studies find results that, while consistent at first sight, are in, fact, somewhat contradictory. For example, although Prada et al (2003) found a strong dependence of satellite velocity dispersion on galactrocentric distance, their measured velocity dispersion within a radius of 125~\\hkpc\\ is similar to the values obtained by McKay \\etal~(2002) at a radius of 275~\\hkpc. Discrepant results were also found by Brainerd \\& Specian (2003) who applied the same technique to the early, ``100k'' data release of the 2degree-Field Galaxy Redshift Survey (2dFGRS; Colless 2001) and derived satellite velocity dispersions which are in qualitative and quantitative disagreement with those of Zaritsky \\etal (1997), McKay \\etal~(2002) and Prada \\etal~(2003). A more extensive analysis of the complete 2dFGRS (Colless \\etal~2003) by Brainerd (2005) also led to disagreements with the results of earlier work. This somewhat confused picture of satellite dynamics is due in large part to different choices of primary galaxy samples and to differences in the modelling and analysis methods.  This paper has multiple aims. Firstly, we carry out a new analysis of the dynamics of satellites around bright galaxies of different morphological types selected from the full 2dFGRS. The goal is to constrain the velocity dispersion and mass of their dark matter halos and we therefore select a sample of isolated galaxies chosen according to strict criteria. Secondly, we investigate the reliability and accuracy of commonly used dynamical analysis methods. For this, we make extensive use of realistic mock catalogues constructed from large cosmological N-body simulations and different semi-analytic galaxy formation models (Cole \\etal~2000; \\munichsambrak). A similar approach, but in a different context, was adopted by \\vdbpapa. Finally, we attempt to understand the root cause of the differences found in previous work, again relying on the use of realistic mock catalogues.  The paper is organised as follows. In Section~\\ref{sec:data}, we briefly present some of the characteristics of the 2dFGRS data and simulations used for our analysis. In Section~\\ref{sec:sample}, we describe the satellite sample selection scheme, together with its robustness to changes in the selection parameters. The analysis of the 'stacked' satellite velocity distribution is carried in Section~\\ref{sec:anal} while, in Section~\\ref{sec:vel_disp_est}, we present velocity dispersion estimates for our mock catalogues and for 2dFGRS primaries split according to luminosity and morphological type. Using a model for the relationship between dark halo mass and satellite velocity dispersion, we give, in Section~\\ref{sec:mass_est}, an estimate of the mass of the halos around 2dFGRS galaxies. In Section~\\ref{sec:discussion}, we compare our results with those of previous studies, and we conclude in Section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have developed, tested and applied a method to probe the properties of extended dark matter haloes around bright galaxies. We do this by carefully selecting isolated galaxies in the 2dFGRS and using their faint satellites as tracers of the gravitational potential. By stacking systems of similar primary luminosity to improve the signal-to-noise, we estimate the satellite velocity dispersion. Realistic mock galaxy catalogues, created from cosmological N-body simulations populated with a semi-analytic galaxy formation scheme, enable us to relate the measured velocity dispersion of the satellite system to the velocity dispersion and mass of the underlying dark halo.  Fig.~\\ref{fig:sigma_amag_morph} shows evidence for the existence of dark matter halos around typical galaxies. Our sample of satellites probes the potential well of the primaries out to several hundred kiloparsecs and demonstrates that the dark halos extend many times beyond the optical radius of the primary. This is in agreement with the current theoretical picture of galaxy formation in a cold dark matter universe (e.g. White \\& Frenk 1991, Kauffmann \\etal~1993, Cole \\etal~2000). The satellite velocity dispersion increases with the luminosity of the primary and is much larger for elliptical galaxies than for spiral galaxies of similar b$_{\\rm J}$ luminosity.  The total extent of the dark halo is not constrained by our data. Although, the satellite distribution extends to $r_p \\sim$~375~\\hkpc, most of the signal comes from within $r_p \\sim$~175~\\hkpc. Within the errors, the velocity dispersion appears to be constant within $r_p \\sim$~175~\\hkpc\\ and $r_p \\sim$~375~\\hkpc. In this range, the satellite velocity dispersion does not depend strongly on the luminosity of the primary.  Our mock catalogues allow us to calibrate the velocity dispersion-mass relation for galaxies selected according to the isolation criterion of our 2dFGRS sample. Fig.~\\ref{fig:amag_mass} then indicates that elliptical galaxies reside in haloes which are at least 4 times more massive than spiral galaxies of similar b$_{\\rm J}$ brightness. Galaxy like the Milky-Way typical reside in dark matter halos of mass $\\rm \\sim 3.5^{+4.0}_{-2.1}\\,10^{11}$~\\hmsol\\ within 175~\\hkpc.  A key assumption in our analysis is that isolated galaxies of similar luminosity reside in halos of similar mass. It is this that justifies the stacking procedure. Our semi-analytic models allow us to test the validity of this assumption. We find that in one of two semi-analytic models that we have investigated, the \\munichsam\\ model, there is very little scatter in the relation between central galaxy luminosity and halo mass. A different semi-analytic model, however, that by Cole \\etal~(2000), predicts considerable scatter in this relation and this introduces large errors in the halo properties inferred from a stacking analysis. Mocks constructed from this model return an increasing satellite velocity dispersion as a function of primary luminosity which, however, deviates systematically from the velocity dispersion of the host dark halos. The reasons behind this difference in the galaxy formation models are not investigated in detail here but it serves to illustrate that significant theoretical uncertainties remain in the kind of analysis that we have presented here."
    },
    "0710/0710.5645_arXiv.txt": {
        "abstract": "We describe the development and implementation of the SEGUE (Sloan Extension for Galactic Exploration and Understanding) Stellar Parameter Pipeline (SSPP). The SSPP derives, using multiple techniques, radial velocities and the fundamental stellar atmospheric parameters (effective temperature, surface gravity, and metallicity) for AFGK-type stars, based on medium-resolution spectroscopy and $ugriz$ photometry obtained during the course of the original Sloan Digital Sky Survey (SDSS-I) and its Galactic extension (SDSS-II/SEGUE). The SSPP also provides spectral classification for a much wider range of stars, including stars with temperatures outside of the window where atmospheric parameters can be estimated with the current approaches. This is Paper I in a series of papers on the SSPP; it provides an overview of the SSPP, and initial tests of its performance using multiple data sets. Random and systematic errors are critically examined for the current version of the SSPP, which has been used for the sixth public data release of the SDSS (DR-6). ",
        "introduction": "The Sloan Extension for Galactic Understanding and Exploration (SEGUE) is one of three surveys that are being executed as part of the current extension of the Sloan Digital Sky Survey (SDSS-II), which consists of LEGACY, SUPERNOVA SURVEY, and SEGUE. The SEGUE program is designed to obtain $ugriz$ imaging of some 3500 square degrees of sky outside of the original SDSS-I footprint (Fukugita et al. 1996; Gunn et al. 1998, 2006; York et al. 2000; Stoughton et al. 2002; Abazajian et al. 2003, 2004, 2005; Pier et al. 2003; Adelman-McCarthy et al. 2006, 2007a). The regions of sky targeted are primarily at lower Galactic latitudes ($|b|~< ~35^{\\circ}$), in order to better sample the disk/halo interface of the Milky Way. As of Data Release 6 (DR-6, Adelman-McCarthy et al. 2007b), about 85\\% of the planned additional imaging has already been completed. SEGUE is also obtaining $R$ $\\simeq$ 2000 spectroscopy, over the wavelength range $3800-9200$\\, {\\AA}, for some 250,000 stars in 200 selected areas over the sky available from Apache Point, New Mexico. The spectroscopic candidates are selected on the basis of $ugriz$ photometry to populate roughly 16 target categories, chosen to explore the nature of the Galactic stellar populations at distances from 0.5 kpc to over 100 kpc from the Sun. Spectroscopic observations have been obtained for roughly half of the planned targets thus far, a total of about 120,000 spectra. The SEGUE data clearly require automated analysis tools in order to efficiently extract the maximum amount of useful astrophysical information for the targeted stars, in particular their stellar atmospheric parameters, over wide ranges of effective temperature ($T_{\\rm eff}$), surface gravity (log $g$), and metallicity ([Fe/H]).  Numerous methods have been developed in the past in order to extract atmospheric-parameter estimates from medium-resolution stellar spectra in a fast, efficient, and automated fashion. These approaches include techniques for finding the minimum distance (parameterized in various ways) between observed spectra and grids of synthetic spectra (e.g., Allende Prieto et al. 2006), non-linear regression models (e.g., Re Fiorentin et al. 2007, and references therein), correlations between broadband colors and the strength of prominent metallic lines, such as the Ca~II K line (Beers et al. 1999), auto-correlation analysis of a stellar spectrum (Beers et al. 1999, and references therein), obtaining fits of spectral lines (or summed line indices) as a function of broadband colors (Wilhelm et al. 1999), or the behavior of the Ca~II triplet lines as a function of broadband color (Cenarro et al. 2001a,b). However, each of these approaches exhibits optimal behavior only over restricted temperature and metallicity ranges; outside of these regions they are often un-calibrated, suffer from saturation of the metallic lines used in their estimates at high metallicity or low temperatures, or lose efficacy due to the weakening of metallic species at low metallicity or high temperatures. The methods that make use of specific spectral features are susceptible to other problems, e.g., the presence of emission in the core of the Ca~II K line for chromospherically active stars, or poor telluric line subtraction in the region of the Ca~II triplet. Because SDSS stellar spectra cover most of the entire optical wavelength range, one can apply several approaches, using different wavelength regions, in order to glean optimal information on stellar parameters. The combination of multiple techniques results in estimates of stellar parameters that are more robust over a much wider range of $T_{\\rm eff}$, log $g$, and [Fe/H] than those that might be produced by individual methods. In this first paper of a series, we describe the SEGUE Stellar Parameter Pipeline (hereafter, SSPP), which implements this ``multi-method'' philosophy. We also carry out a number of tests to assess the range of stellar atmospheric parameter space over which the estimates obtained by the SSPP remain valid. The second paper in the SSPP series (Lee et al. 2007b; hereafter Paper II) seeks to validate the radial velocities and stellar parameters determined by the SSPP by comparison with member stars of Galactic three globular clusters (M~15, M~13, and M~2) and two open clusters (NGC~2420 and M~67). A comparison with an analysis of high-resolution spectra for SDSS-I/SEGUE stars is presented in the third paper in this series (Allende Prieto et al. 2007; hereafter Paper III).  Section 2 describes determinations of radial velocity used by the SSPP. The procedures used to obtain an estimate of the appropriate continuum, and the determination of line indices, is explained in \\S 3. The methods that the SSPP employs for determining stellar parameters are presented in \\S 4. Section 5 addresses the determinations of auxiliary estimates of effective temperature, based on both theoretical and empirical approaches; these methods are used for stars where adequate estimates of $T_{\\rm eff}$ are not measured by our primary techniques. A decision tree that gathers the optimal set of parameter estimates based on the multiple methods is introduced in \\S 6. Section 7 summarizes the conditions for raising various flags to warn potential users where uncertainties in parameter determinations remain. In \\S 8, we discuss validation of the parameters determined by the SSPP based on SDSS-I/SEGUE stars for which we have obtained higher dispersion spectroscopy on various large telescopes, and also compare the parameters with those of likely member stars of Galactic open and globular clusters. Assignments of approximate MK spectral classifications are described in \\S 9. In \\S 10 we describe several methods (still under testing) for the determination of distance estimates used by the SSPP. Section 11 presents a summary and conclusions.  In the following, the colors ($u-g$, $g-r$, $r-i$, $i-z$, and $B-V$) and magnitudes ($u,~g, ~r,~i,~z,~B,~\\rm and~V$) are understood to be de-reddened and corrected for absorption (using the dust maps of Schlegel et al. 1998), unless stated specifically otherwise. ",
        "conclusions": "We have described the development and execution of the SEGUE Stellar Parameter Pipeline (SSPP), which makes use of multiple approaches in order to estimate the fundamental stellar atmospheric parameters (effective temperature, $T_{\\rm eff}$, surface gravity, log $g$, and metallicity, parameterized by [Fe/H]) for stars with spectra and photometry obtained during the course of the original Sloan Digital Sky Survey (SDSS-I) and its current extension (SDSS-II/SEGUE).  The use of multiple approaches allows for an empirical determination of the internal errors for each derived parameter, based on the range of the reported values from each method. From consideration of about 140,000 spectra of stars obtained during SDSS-I and SEGUE that have derived stellar parameters available in the range 4500~K $\\le$ $T_{\\rm eff}$ $\\le$ 7500~K, typical internal errors obtained by the SSPP are $\\sigma(T_{\\rm eff})$ = 73~K (s.e.m), $\\sigma(\\log~g)$ = 0.19 (s.e.m), and $\\sigma([\\rm Fe/\\rm H])$ = 0.10 (s.e.m). Paper III points out that the internal scatter estimates obtained from averaging the multiple estimates of the parameters produced by the SSPP underestimate the external errors, owing to the fact that several methods in the SSPP use similar same parameter indicators and atmospheric models.  The results of a comparison with an average of two different high-resolution spectroscopic analyses of 124 SDSS-I/SEGUE stars suggests that the SSPP is able to determine $T_{\\rm eff}$, log $g$, and [Fe/H] to precisions of 135~K, 0.26 dex, and 0.21 dex, respectively, after combining small systematic offsets quadratically for stars with 4500~K $\\leq T_{\\rm eff} \\leq$ 7500~K. These errors differ slightly from the those obtained by Paper III ($\\sigma(T_{\\rm eff})$ = 130~K, $\\sigma(\\log g)$ = 0.21 dex, and $\\sigma([\\rm Fe/\\rm H])$ = 0.11 dex), even though they share a common set of high-resolution calibration observations. This arises because Paper III derived the external uncertainties of the SSPP only taking into account the stars observed with the HET (on the grounds of internal consistency). The sample referred to as OTHERS in Paper III exhibits somewhat larger scatter in its parameters, when compared with those determined by the SSPP. Observation of several hundred additional stars from SDSS-I/SEGUE with HET is now underway. Thus, in the future, we will be able to use a homogeneous sample gathered by HET in our tests. Also, additional high-resolution data for stars outside of our adopted temperature range will enable tests for both cooler and warmer stars.  Considering the internal scatter from the multiple approaches and the external uncertainty from the comparisons with the high-resolution analysis together, the typical uncertainty in the stellar parameters delivered by the SSPP are $\\sigma(T_{\\rm eff})$ = 154~K, $\\sigma(\\log~g)$ = 0.31 dex, and $\\sigma([\\rm Fe/\\rm H])$ = 0.23 dex, over the temperature range 4500~K $\\leq T_{\\rm eff} \\leq$ 7500~K.  However, it should be kept in mind that the errors stated above apply for the very highest $S/N$ spectra obtained from SDSS ($S/N >$ 50/1), as only quite bright stars were targeted for high-resolution observations. In addition, outside of the quoted temperature range (4500~K $\\le$ $T_{\\rm eff}$ $\\le$ 7500~K), we presently do not have sufficient high-resolution spectra to fully test the parameters obtained by the SSPP.  The results of a comparison with likely member stars of a sample of Galactic open and globular clusters suggest that SSPP may slightly overestimate [Fe/H] (by $\\sim 0.15$ dex) for stars with [Fe/H] $< -$2.0, and underestimate [Fe/H] (by $\\sim 0.30$ dex) for stars with near-solar metallicities. Slight trends of [Fe/H] with $g-r$ are noticed for the higher metallicity clusters as well, although further data will be needed in order to verify this.  Approximate spectral types are assigned for stars, based on two methods, with differing limitations. A preliminary set of distance determinations for each star is also obtained, although future work will be required in order to identify the optimal method.  We conclude that the SSPP determines sufficiently accurate and precise radial velocities and atmospheric parameter estimates, at least for stars in the effective temperature range from 4500~K to 7500~K, to enable detailed explorations of the chemical compositions and kinematics of the thick-disk and halo populations of the Galaxy."
    },
    "0710/0710.2247_arXiv.txt": {
        "abstract": "{The deeper and more extended survey of the central parts of the Galactic Plane by H.E.S.S. during 2005-2007 has revealed a number of new point-like, as well as, extended sources. Two point-like sources can be associated to two remarkable objects around ``Crab-like'' young and energetic pulsars in our Galaxy : G21.5-0.9 and Kes~75. The characteristics of each of the sources are presented and possible interpretations are briefly discussed.}     \\begin{document} ",
        "introduction": "\\begin{figure*}[!t]% \\begin{center} \\includegraphics*[width=0.46\\textwidth,angle=0,clip]{g215exc008Fig1.eps} \\includegraphics*[width=0.46\\textwidth,angle=0,clip]{J1846m029exc008Fig1.eps} \\end{center} \\caption{ Smoothed excess maps ($\\sigma=0.08^{\\circ}$) of the 0.5$^{\\circ}\\times0.5^{\\circ}$ field of view around the positions of HESS~J1833-105 (left) and HESS~J1846-029 (right). The white contours show the pre-trials significance levels for 4, 5, 6 $\\sigma$, and 7, 8, 9 $\\sigma$, respectively. The black triangle marks the position of the pulsars. The best-fit positions of the two sources are marked with an error cross (for HESS~J1846-029 the latter overlaps with the triangle).} \\label{skymaps} \\end{figure*}  The standard candle of VHE astronomy, the Crab Nebula, has served for decades as a yardstick in almost all wavelengths, and yet it is a very peculiar object, harbouring the most energetic and one of the youngest pulsars of our Galaxy. Since the early days, where the similarities of the historical trio Crab/3C~58/G~21.5-0.9 were under debate~\\cite{WilsonWeiler76}, radio and X-ray astronomy have provided a wealth of information by detecting and characterizing nebulae around rotation-powered pulsars. In the VHE domain, H.E.S.S. has revealed more than a dozen pulsar wind nebulae (PWN), either firmly established as such or compelling candidates~\\cite{HESSpwnICRC07}, almost all of which are middle-aged (at least few kyrs up to $\\sim$100 kyrs, except MSH~15-52) and exhibit an offset between the pulsar position and the nebula center. We report here on the VHE emission discovery of two remarkable objects, G~21.5-0.9 and Kes75, which also harbor very young and energetic pulsars and which on some aspects, especially their plerionic nebular emission due to an energetic pulsar, can be considered as Crab-like.  \\begin{figure*}[tb] \\begin{center} \\includegraphics[width=0.48\\textwidth]{g21.5m09on011Pwl1Z.eps} \\includegraphics[width=0.48\\textwidth]{J1846on011Pwl1Zdec.eps} \\end{center} \\caption{Differential energy spectra above for HESS~J1833-105 (left) and HESS~J1846-029 (right). The shaded area shows the 1 $\\sigma$ confidence region for the fit parameters.} \\label{spectra} \\end{figure*}  G21.5-0.9 \\cite{Altenhoff70}, recently revealed as a composite SNR consisting of a centrally peaked PWN and a 4{\\arcmin} shell~\\cite{Bocchino2005,MathesonSafi-Harb2005}, was previously classified as one of the about ten Crab-like SNR~\\cite{Green2004}. Its flat spectrum PWN is polarised in radio~\\cite{BeckerSzymkowiak1981} with a spectral break above 500 GHz~\\cite{GallantTuffs1998}. The non-thermal X-ray PWN with radius $\\sim$40{\\arcsec} shows significant evidence of cooling ~\\cite{Slane2000}, with the power-law photon index steepening from 1.43$\\pm$0.02 near the pulsar to 2.13$\\pm$0.06 at the edge of the PWN. There appears to be a synchrotron X-ray halo at a radius of 140{\\arcsec} from the pulsar which could originate in the shell~\\cite{Bocchino2005,MathesonSafi-Harb2005}, with a contribution of scattering off dust grains as proposed by Bocchino et al.~\\cite{Bocchino2005}. The 61.8 ms pulsar PSR~J1833-1034, with a spin-down power of ${\\dot E} = 3.3\\times 10^{37}$erg/s and a characteristic age of 4.9 kyr was discovered only recently through its faint radio pulsed emission~\\cite{Gupta2005,Camilo2006}. Given the derived distance of 4.7$\\pm$0.4 kpc, the age of G~21.5-0.9 was revised downwards by a factor of $\\sim$10 to force consistency with the freely expanding SNR shell~\\cite{Camilo2006}.  PSR~J1833-1034 in G~21.5-0.9 is the second most energetic pulsar known in the Galaxy.   Kes 75 (SNR G29.70.3) is also a prototypical example of a composite remnant for which the distance of 19 kpc was estimated through neutral hydrogen absorption measurements~\\cite{BeckerHelfand1984}. Its 3.5{\\arcmin} radio shell surrounds a flat-spectrum highly polarized radio core, and harbors, at its center, the 325 ms X-ray Pulsar, PSR~J1846-258~\\cite{Gotthelf2000}. The latter has the shortest known characteristic age $\\tau_c= 723$ yr and a large inferred magnetar-like magnetic field of B$=4.9\\times 10^{13}$G. The pulsar lies within a 25{\\arcsec}$\\times$20{\\arcsec} X-ray nebula which exhibits an photon index of 1.92$\\pm$0.04, but no evidence of cooling as a function of the distance to the pulsar. Like in G~21.5-0.9 there is an X-ray halo, in this case due mostly to dust scattering, but a non-thermal contribution from electrons accelerated in the shell remains possible~\\cite{Helfand2003}. ",
        "conclusions": "It is remarkable that de Jager et al.~\\cite{deJager1995} predicted that plerionic VHE $\\gamma$-rays from G21.5-0.9 would be detectable at a level of $4\\times10^{-13}$ cm$^{-2}$~s$^{-1}$ at 1~TeV with an electron spectral index of $\\sim$2.8, which would give a photon index near 2.0 at VHE energies (after including KN effects given the contributions from dust and CMBR). Their prediction was based on an assumed equipartition field strength of 22 $\\mu$G which is close to the value of $\\sim$15$\\mu$G implied from $\\gamma$-ray observations reported here (assuming IC scattering on CMB photons only, and using the ratio of the X-ray to the $\\gamma$-ray luminosities: $L_{\\rm X}/L_{\\gamma}\\sim30$). The equipartition field strength was afterwards increased to 0.3 mG following the revision of the maximum spectral range of the radio PWN to 500 GHz \\cite{GallantTuffs1998, Camilo2006}. However, the detection of VHE $\\gamma$-rays by H.E.S.S. from PWN tends to confirm the suggestion of Chevalier \\cite{Chevalier2004} that some PWN may be particle dominated, so that the true PWN field strength may be significantly lower than equipartition for some objects. In the case of Kes~75, $L_{\\rm X}/L_{\\gamma}\\sim$10 yields also a lower than equipartition nebular magnetic field strength of $\\sim$10 $\\mu$G. It should be noted that Kes~75 shows the highest conversion efficiency in X-rays ($\\sim15$\\%) as compared to other ``Crab-like'' pulsars ($\\sim3$\\% and $\\sim0.6$\\% for the Crab and G~21.5-0.9, respectively) and a 100 times larger $\\gamma$-ray efficiency ($\\sim$2\\%) than the Crab and G~21.5-0.9 which are similar in that respect ($\\sim$0.02\\%). However, the latter object's $L_{\\rm X}/L_{\\gamma}\\sim30$ is 4 times smaller than that of the Crab Nebula $L_{\\rm X}/L_{\\gamma}\\sim120$. These numbers together with the spin parameters and high surface magnetic field in the case of PSR~J1846-0258, show that these objects, although ``Crab-like'' in some aspects, do possess peculiar properities.  Given the evidence for synchrotron emission in the SNR shell, an alternative interpretation of the VHE emissions of G~21.5-0.9 and Kes75 would be radiation from particles accelerated at the non-relativistic forward shock of the freely expanding SNR. However the required field strength in the shell to explain the H.E.S.S. detection in terms of IC scattering should be much lower than 10 $\\mu$G, value which may be unreasonably low for typical expanding SNR shells. Deeper observations of both sources could help to constrain the size of the VHE emission region and to ascertain whether it is compatible with this scenario."
    },
    "0710/0710.0418_arXiv.txt": {
        "abstract": "We are investigating mass fractions on the crust of a neutron star which would remain after one year of cooling. We use cooling curves corresponding with various densities, or depths, of the neutron star just after its formation. We assume the modified Urca process dominates the energy budget of the outer layers of the star in order to calculate the temperature of the neutron star as a function of time. Using a nuclear reaction network up to technetium, we calculate how the distribution of nuclei quenches at various depths of the neutron star crust. The initial results indicate that $^{28}$Si is the lightest isotope to be optically thick on the surface after one year of cooling. ",
        "introduction": "The observed  emission from a neutron star passes through a crustal layer of the neutron star. In order to fully interpret the observed emission from a neutron star, we need to understand what comprises this crustal layer of the neutron star \\cite{Hern84b,heylmnras01}. Here, the mass fractions on the surface are calculated for what would exist on the crust after one year of cooling of the neutron star. These mass fractions are calculated using a 489 isotope reaction network which burns up to technetium written by F.X. Timmes \\cite{timmesapjs99} and is available on Timmes's webpage \\footnote{\\url{http://www.cococubed.com/code_pages/burn.shtml}}. ",
        "conclusions": "We have examined the mass fractions that will result in an optically thick surface layer by looking at two cases: a density of $10^{12}\\;$g/cm$^3$ and another of $10^{7}\\;$g/cm$^3$. The mass fractions at these densities are calculated using a 489 isotope reaction network which burns up to technetium. The neutron star is cooled for a year, starting with a temperature of $10^{10}$K, assuming the modified Urca process dominates. Upon calculating the mass fractions the minimum mass fraction abundance required for enough of an isotope to form an optically thick layer such that it has a surface density of 1g/cm$^2$ is calculated. From the two cases we found that the deeper density resulted in a larger variety of lighter isotopes can form an optically thick layer. We found $^{28}$Si, $^{30}$Si, $^{31}$P, $^{32}$S, $^{33}$S, $^{34}$S, and $^{36}$Ar rise to the surface. Future work on this includes exploring more densities, comparing to analytic calculations and estimating the timescale for the light impurities to float to the surface.                   \\begin{theacknowledgments} We would like to thank Ed Brown for helpful discussions during the conference. The Natural Sciences and Engineering Research Council of Canada, Canadian Foundation for Innovation and the British Columbia Knowledge Development Fund supported this work. This research has made use of NASA's Astrophysics Data System Bibliographic Services. \\end{theacknowledgments}"
    },
    "0710/0710.5521_arXiv.txt": {
        "abstract": "I calculate the linear stability of a stratified low collisionality plasma in the presence of a weak magnetic field.  Heat is assumed to flow only along magnetic field lines.  In the absence of a heat flux in the background plasma, Balbus (2000) demonstrated that plasmas in which the temperature {\\it increases} in the direction of gravity are buoyantly unstable to convective-like motions (the ``magnetothermal instability'').  I show that in the presence of a background heat flux, an analogous instability is present when the temperature {\\it decreases} in the direction of gravity.  The instability is driven by the background heat flux and the fastest growing mode has a growth time of order the local dynamical time.  Thus, independent of the sign of the temperature gradient, weakly magnetized low collisionality plasmas are unstable on a dynamical time to magnetically-mediated buoyancy instabilities.  The instability described in this paper is predicted to be present in clusters of galaxies at radii $\\sim 0.1-100$ kpc, where the observed temperature increases outwards. Possible consequences for the origin of cluster magnetic fields, ``cooling flows,'' and the thermodynamics of the intercluster medium are briefly discussed. ",
        "introduction": "\\label{sec:intro}  Thermally stratified fluids are buoyantly unstable when the entropy increases in the direction of gravity, a result of considerable importance to the theory of stellar structure (Schwarzschild 1958). Remarkably, however, this well-known result changes in a low collisionality plasma in which i) the collisional mean free path of electrons is larger than the electron Larmor radius and ii) thermal conduction is the dominant mode of heat transport (Balbus 2000).  In such a plasma, heat is transported primarily along magnetic field lines.  For the simple problem of a horizontal magnetic field in a vertically stratified plasma, Balbus (2000) showed that the condition for the plasma to be buoyantly unstable becomes that the temperature (not entropy) increase in the direction of gravity.  The resulting ``magnetothermal instability'' (MTI) has been studied with nonlinear simulations by Parrish \\& Stone (2005; 2007).  In a subsequent paper, Balbus (2001) generalized his initial result to rotating flows and magnetic fields of arbitrary orientation, but still under the assumption that there is no heat flux in the background plasma (i.e., that the field lines are initially isothermal).  This latter assumption is unlikely to hold in many low collisionality astrophysical plasmas such as clusters of galaxies and hot accretion flows onto black holes.  In this paper I extend Balbus's calculation and study the stability of weakly magnetized plasmas in the presence of a background heat flux. I show that the presence of a heat flux drives a buoyancy instability analogous to the MTI when the temperature {\\it decreases} in the direction of gravity (a situation that is MTI stable according to Balbus's analysis).  This instability is distinct from the heat flux driven overstabilities described in Socrates, Parrish, \\& Stone (2007).\\footnote{In my analysis below, I utilize the Boussinesq approximation to focus on nearly incompressible perturbations.  In this limit, Socrates et al. predict that the slow mode is stable while the fast mode is unstable on a dynamical time.  Because the Boussinesq approximation filters out fast waves, I do not expect any version of their overstabilities to be present in my analysis.}  In the next two sections I summarize the equations and assumptions used in my analysis (\\S \\ref{sec:eqns}) and the results of the linear stability calculation (\\S \\ref{sec:results}).  I then discuss possible applications of the heat flux-driven version of the MTI, in particular to the intercluster plasma in clusters of galaxies (\\S \\ref{sec:disc}). ",
        "conclusions": "\\label{sec:disc}  The above analysis shows that, regardless of the sign of the temperature gradient, a weakly magnetized low collisionality plasma in which heat flows primarily along magnetic field lines is buoyantly unstable.  For $dT/dz < 0$, this instability is the magnetothermal instability (MTI) derived by Balbus (2000; 2001) and simulated by Parrish \\& Stone (2005; 2007). Although a plasma with $dT/dz > 0$ is MTI stable according to Balbus (2001), I have shown that an analogous buoyancy instability in fact exists for $dT/dz > 0$ in the presence of a vertical magnetic field and a background heat flux. Physically, this new instability arises because perturbed fluid elements are heated/cooled by the background heat flux in such a way as to become buoyantly unstable (\\S \\ref{sec:new}).  In many astrophysical plasmas, the sign of the temperature gradient is fixed to be $dT/dz < 0$ by basic principles.  These systems may be MTI unstable, but they will be stable to the new buoyancy instability discussed in this paper.  This is typically the case in cooling white dwarfs and neutron stars, where the flow of heat outwards requires $dT/dz < 0$.  It also also the case in hot accretion flows onto compact objects, because the inflow of matter and the release of gravitational potential energy drives $dT/dz < 0$.  The heat flux driven instability described in this paper may act in the transition region between cool dense gas and hot low density plasma in stellar coronae, accretion disks, and the multi-phase interstellar medium.  However, these regions tend to be strongly magnetized, which will inhibit the instability (\\S \\ref{sec:strongB}). A more promising application is to the hot intercluster plasma in galaxy clusters.  Plasma in hydrostatic equilibrium in a Navarro, Frank, \\& White (1997) dark matter potential well has a temperature profile which is locally isothermal ($\\rho \\propto r^{-2}$) at a scale radius $R_s \\simeq 100-400$ kpc.  The temperature is predicted to decrease for radii both smaller and larger than $\\sim R_s$.  Such a radial variation in the temperature of the intercluster plasma is directly observed in many systems (e.g., Piffaretti et al. 2005).  At radii larger than $\\sim R_s$, the intercluster plasma is MTI unstable (e.g., Parrish \\& Stone 2007), but it is MTI stable in the cores of clusters where the temperature decreases inwards.  However, it is precisely these radii that are unstable to the buoyancy instability discussed in this paper.  Thus, provided that the field is not too strong (see \\S \\ref{sec:strongB}), I conclude that the entire intercluster plasma in galaxy clusters is unstable to magnetically mediated buoyancy instabilities.  The implications of this instability for the intercluster medium will be investigated in future papers using nonlinear simulations.  Here I briefly comment on the possible consequences. Given the presence of exponentially growing instabilities that amplify magnetic fields {\\it at all radii} in galaxy clusters, it is natural to suspect that these instabilities play a significant role in generating the observed (e.g., Federica \\& Feretti 2004) $\\mu G$ magnetic fields in clusters from smaller cosmological ``seed'' fields.  In addition, just as the MTI is found to re-orient the magnetic field to be largely radial, allowing heat to flow down the temperature gradient (Parrish \\& Stone 2007; Sharma \\& Quataert, in preparation), I suspect that the heat flux driven buoyancy instability discussed here will generate a significant horizontal magnetic field if one was not present originally.  This will act so as to decrease the net heat flux through the plasma (which is the origin of the instability in the first place).  Given the close connection between the heat flux and magnetic field described in this paper, it is unclear whether current calculations of the effective conductivity and heat flux in cluster plasmas (e.g., Narayan \\& Medvedev 2001; Chandran \\& Maron 2004) are correct, since they do not capture this {dynamical} coupling.  It may thus turn out that thermal conduction from large radii will prove to be less effective than previous authors have suspected (e.g., Bertschinger \\& Meiksin 1986; Ruszkowski \\& Begelman 2002; Zakamska \\& Narayan 2003) at heating ``cooling flow'' cores in clusters. Regardless of the accuracy of this speculation, the results of this paper highlight the need for a proper treatment of the combined effects of thermal conduction and magnetically-mediated buoyancy instabilities on the plasma in galaxy clusters.  In future work it will also be interesting to study the dynamics of the heat flux driven instability in the presence of cosmic rays, which may be energetically significant in cluster cores because of a central black hole, and which are known to modify the MTI (Chandran \\& Dennis 2006).   \\"
    },
    "0710/0710.1302_arXiv.txt": {
        "abstract": "{We use the techniques of effective field theory in an expanding universe to examine the effect of choosing an excited inflationary initial state built over the Bunch-Davies state on the CMB bi-spectrum. We find that even for Hadamard states, there are unexpected enhancements in the bi-spectrum for certain configurations in momentum space due to interactions of modes in the early stages of inflation. These enhancements can be parametrically larger than the standard ones and are potentially observable in future data. These initial state effects have a characteristic signature in $l$-space which distinguishes them from the usual contributions, with the enhancement being most pronounced for configurations corresponding to flattened triangles for which two momenta are collinear.}  \\preprint{PI-COSMO-64}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.5071_arXiv.txt": {
        "abstract": "The Laser Interferometer Space Antenna (LISA) mission will use advanced technologies to achieve its science goals: the direct detection of gravitational waves, the observation of signals from compact (small and dense) stars as they spiral into black holes, the study of the role of massive black holes in galaxy evolution, the search for gravitational wave emission from the early Universe. The gravitational red-shift, the advance of the perihelion of Mercury, deflection of light and the time delay of radar signals  are the classical tests in the first order of General Relativity (GR). However, LISA can possibly test Einstein's theories in the second order and perhaps, it will show some particular feature of non-linearity of gravitational interaction. In the present work we are seeking a method to construct theoretical templates that limit in the first order the tensorial structure of some metric fields, thus the non-linear terms are given by exponential functions of gravitational strength. The Newtonian limit obtained here, in the first order, is equivalent to GR. ",
        "introduction": "The extreme difficulties which arise if one tries to draw physically important conclusions from the basic assumptions of Einstein's theory are mainly due to the non-linearity of the field equations. Moreover, the fact that the spacetime topology is not given a priori, and the impossibility to integrate tensors over finite regions cause difficulties unknown in other branches of mathematical physics. Actually in this respect they are not so different from others fields, for example the electromagnetic field, the scalar field, etc., by themselves obey linear equations in a given spacetime, they form a non-linear system when their mutual interactions are taken into account. The distinctive feature of the gravitational field is that it is self-interacting (as the Yang-Mills field): it is non-linear even in the absence of other fields. This is because it defines the spacetime over which it propagates \\citep{Hawking}.   Linearized gravity is any approximation to General Relativity obtained from $g_{\\mu\\nu}=g_{\\mu\\nu}^{(0)}+h_{\\mu\\nu}$ (where $g_{\\mu\\nu}^{(0)}$ is any curved background spacetime) in Einstein's equation  and retaining only the terms linear in $h_{\\mu\\nu}$ \\citep{Wald}. The weakness of the gravitational field means in the context of general relativity that the spacetime is nearly flat. Small gravitational perturbations in Minkowski space can be treated in the simplest linearized version of General Relativity, \\begin{equation} \\label{linear} g_{\\mu\\nu}=\\eta_{\\mu\\nu}+h_{\\mu\\nu}, \\end{equation} as describing a theory of a symmetric tensor field $h_{\\mu\\nu}$ propagating on at background spacetime. This theory is Lorentz invariant in the sense of Special Relativity. If one wants to obtain a solution of the non-linear equations, it is necessary to employ an iterative method on approximate linear equations whose solutions are shown to converge in a certain neighbourhood of initial surface. It should be possible to avoid some of these difficulties of non-linearity by working in some spacetimes that shall be described in this paper. The proposal is that some metric fields can be separated into the parts carrying the dynamical information and those parts characterizing the coordinates system. In this proposal,  terms of the coordinates system will have tensorial structure limited  only in the first order. The tensors that will describe $g_{\\mu\\nu}$ have linear behavior. Naturally, there is a price that must be paid for the linear tensors. The dynamic terms that carry the gravitational strength have exponential structure. The principal idea came from the basic principle that one should interpret (\\ref{linear}) as separation between pure mathematical and physical terms in metric field tensor. In spite of the fact that $\\eta_{\\mu\\nu}$ plays a key role as empty flat and background spacetime of the Standard Model in the description of fundamental interactions, this background tensor metric $\\eta_{\\mu\\nu}$ is an object wholly mathematical and entirely geometrical, while $h_{\\mu\\nu}$ contains the physical information. The strength of gravity is tied in the components of $h_{\\mu\\nu}$. The proposal of this paper is a working hypothesis to untie the strength of gravity from geometrical tensors. This proposal is valid for a family of metric field tensors $g_{\\mu\\nu}$, and some basic examples such Newtonian limit and gravitational plane waves of low amplitudes are described.  This paper is outlined as follows: in Sec. II, we present the basic mathematic concepts of the (quasi-)idempotent tensors that compose the structure of metric fields approached in this work. In Sec. III, we propose how to link the strength of gravity with the tensors from Sec. II, then is defined a family of exponential metrics. In Sec. IV, we present some examples of these exponential metrics, such as: Yilmaz metric, circularly polarized wave and rotating bodies. In Sec. V, we present exponential metrics (`adjoint metric field') that are non-physical, but which help us to compute Christoffel symbols, and consequently the curvature tensors, Ricci tensors and determinant of metric field. In Sec. VI, we verify the Newtonian limit and also we obtain gravitational waves. In Sec. VII, we present a general conclusion.   We assume spacetime $({\\cal M},{\\bf g})$ to be a ${C^{\\infty}4-}$dimen\\-sional, globally hyperbolic, pseudo-Riemannian manifold ${\\cal M}$ with Lorentzian metric tensor ${\\bf g}$ (whose components are $g_{\\mu\\nu}$) associated with the line element $$ds^2=g_{\\mu\\nu}(x)dx^{\\mu}dx^{\\nu},$$ assumed to have  signature $(+---)$ \\citep{Landau}. Lower case Greek indices refer to coordinates on ${\\cal M}$ and take the values $0,1,2,3.$ The relation between the metric field $g_{\\mu\\nu}$ and the material contents of spacetime is expressed by Einstein's field equation, \\begin{equation} \\label{field equation} R_{\\mu\\nu}-\\frac{1}{2}g_{\\mu\\nu}R=\\frac{8\\pi G}{c^4}T_{\\mu\\nu}, \\end{equation} $T_{\\mu\\nu}$ being the stress-energy-momentum tensor, $R_{\\mu\\nu}$ the contract curvature tensor (Ricci tensor) and $R$ its trace. In an empty region of spacetime we have $R_{\\mu\\nu}=0$, such a region is called  vacuum field. ",
        "conclusions": "This paper deals with a tensorial structure that assumes a (quasi-)idempotent feature to be able to improve at least the linear tensorial template of some tensor metric fields. It is clear that Einstein's field equations are non-linear, however, with these  (quasi-)idempotent tensorial structure, without quadratic tensorial values, the non-linearity becomes more moderate although there is a price to pay. The part that carries  the dynamical information, the strength of gravity is tied to the tensorial structure by exponential functions. In this approach the metric field can be characterized by a background spacetime conformally flat affected by a disturbance. We have approached some examples in this tensorial structure that results in exponential metric fields, we can point out as the main exponential metric obtained in this paper which has been extensively explored: the Yilmaz exponential metric \\citep{Yilmaz1,Yilmaz2,Yilmaz3,Yilmaz4,Yilmaz5,Yilmaz6,Yilmaz7,Clapp,Robertson1,Robertson2,Ibison}. H. Yilmaz has argued that in his theory, the gravitational field can be quantized via Feynman's method \\citep{Yilmaz9,Alley}. Further, it has been found that the quantized theory is finite. Incidentally in the exponential metric fields approached in this work just as in the Yilmaz theory there are no black holes in the sense of event horizons, but there can be stellar collapse \\citep{Robertson1,Robertson2}. However, there are no point singularities.  Interesting results obtained in this work from exponential metric fields are: circularly polarized wave; rotating bodies that in the first order is a deformation of Kerr metric and also we have a deformed static spherically symmetric spacetime. Many discussions around massive stellar objects have suggested, for example, that Kerr metric should  be slightly deviated from Kerr. The possibility of discovering a non-Kerr object should be taken into account when constructing waveform templates for LISA's data analysis tools \\citep{Glampedakis,White}. The technological development is ripe enough so much so in the years to come we might be able to test the second order relativist-gravity effects and may lead to answers to some important questions of gravity.  In this work, we have obtained a simple and general expression for the volume element of a manifold in coordinates $(t,x,y,z)$ given in terms of strength of gravity and of traces of tensors $\\bm \\eta$ and $\\bm \\Upsilon$. It is possible that an analysis of any Lagrangian of field interacting with gravity will become easer.  An interesting observation is the spacetime of circularly polarized plane wave, in this spacetime the volume element $\\sqrt{-g}\\,\\,d^4x$ is the same of Minkowski spacetime, in this sense this  gravitational radiation obtained from exponential metric field does not modify the volume element of background Minkowski spacetime where this plane wave travels onto. Moreover, it was purposed and verified the Newtonian limit as solution for Einstein's equation, since we can assume that the trace of stress-energy tensor is $T\\approx \\rho c^2$. Other important solution of Einstein's equation analysed in this paper was the plane gravitational wave for the empty space since we have considered the vanished stress-energy tensor to the first order in $\\Phi$. Both solved Einstein's equations for Newtonian limit and plane gravitational wave propagating in the vacuum are cases that the strength of gravity is small, $\\Phi\\ll 1$. We have analysed  the Newtonian limit  in the case that $\\partial_{\\alpha}\\bm\\Upsilon=0$, and analysed the plane gravitational wave considering  the strength of gravity as a constant term, thus we had two independent Ricci tensors, $R_{\\mu\\nu}^{(1)}$ which $\\partial_{\\alpha}\\bm\\Upsilon=0$ (for Newtonian limit) and $R_{\\mu\\nu}^{(2)}$ which $\\partial_{\\alpha}\\Phi=0$ (for plane gravitational wave). In a forthcoming work, an analysis of Einstein's equations with both non-vanished $\\partial_{\\alpha}\\bm\\Upsilon$ and $\\partial_{\\alpha}\\Phi$, will be considered.   It is missing a discussion about quantities of physical interest in the solutions of Einstein's equations which describe the exterior and interior gravitational field. Yilmaz has argued the existence of the matter part in the right-hand side of the field equations correspondent to field energy in the exterior. This paper lacks a discussion about the interior and the exterior field energies denoted by a total stress-energy tensor. An analysis about the total stress-energy tensor will be the object of a forthcoming study, where the physical consequences of terms of deformity in Kerr and Schwarzschild solutions could be analysed.  We know that the dark energy and the dark matter problems are challenges to modern astrophysics and cosmology; as a typical example, we could mention the galactic rotation curves of spiral galaxies, that probably, indicates the possible failure of both Newtonian gravity and General Relativity on galactic and intergalactic scales. To explain astrophysical and cosmological problems with arguments against dark energy and dark matter many works have been devoted to the possibility that the Einstein-Hilbert Lagrangian, linear in the Ricci scalar R, should be generalized. In this sense, the choice of a generic function $f(R)$ can be derived by matching the data and the requirement that no exotic ingredient have to be added \\citep{Allemandi,Barrow,Capozziello1,Capozziello2,Capozziello3,Carroll1,Carroll2,Faraoni,Flanagan,Koivisto,Nojiri1,Nojiri2,Nojiri3,Sotiriou}. This class of theories when linearized exhibits others polarization modes for the gravitational waves, of which two correspond to the massless graviton and others such as massive scalar and ghost modes in $f(R)$ gravity \\citep{Bellucci,Bogdanos}. In this way, analyses in any order to $f(R)$ gravity with `exponential metrics' proposed in the present work could give a positive contribution to the debate of astrophysical and cosmological questions."
    },
    "0710/0710.0979_arXiv.txt": {
        "abstract": "We have obtained velocity-resolved spectra of the $\\hoz ~ (\\lambda = 2.1218 \\micron)$ emission line at $2\\arcsec$ angular resolution (or $\\sim 0.08$~pc spatial resolution) in four regions within the central 10 pc of the Galaxy where the supernova-like remnant Sgr A East is colliding with molecular clouds. To investigate the kinematic, physical, and positional relationships between the important gaseous components in the center, we compared the H$_2$ data cube with previously published NH$_3$ data. The projected interaction-boundary of Sgr A East is determined to be an ellipse with its center offset $\\sim1.5$~pc from Sgr A* and dimensions of 10.8~pc $\\times$ 7.6~pc. This H$_2$ boundary is larger than the synchrotron emission shell but consistent with the dust ring which is believed to trace the shock front of Sgr A East. Since Sgr A East is driving shocks into its nearby molecular clouds, we can determine their positional relationships using the shock directions as indicators. As a result, we suggest a revised model for the three-dimensional structure of the central 10~pc. The actual contact between Sgr A East and all of the surrounding molecular material, including the circum-nuclear disk and the southern streamer, makes the hypothesis of infall into the nucleus and feeding of Sgr A* very likely. ",
        "introduction": "In the central 10 pc of our Galaxy, the Sgr A region contains several characteristic objects; a candidate for super-massive black hole (Sgr A*) of about $4 \\times 10^6~\\msun$ (see \\citealt*{ghe03,sch03} and references therein), a surrounding cluster of stars (the Central cluster), molecular and ionized gas clouds (the circum-nuclear disk (CND) and Sgr A West), supernova remnants (SNR G~359.92-0.09 and Sgr A East). They are surrounded by molecular structures including two giant molecular clouds (GMCs) M-0.02-0.07 and M-0.13-0.08 (also known as the `$50~\\kms$ cloud' and the `$20~\\kms$ cloud', respectively).  In addition to the two GMCs, recent accurate radio observations have resolved several dense and filamentary molecular features around the Sgr A complex; the `molecular ridge', the `southern streamer', the `northern ridge', and the `western streamer' (see Figure~\\ref{3D_model} of this paper and Figures~3, 9, and 10 of \\citealt*{mcg01} and Figure~14 of \\citealt*{her05}). These molecular features are believed to play important roles in feeding the central massive black hole \\citep{ho91,coi99,coi00,mcg01}. The interaction between these various components is responsible for many of the phenomena occurring in this complicated and unique portion of the Galaxy. Developing a comprehensive picture of the primary interactions between the components at the Galactic center will also improve our understanding of the nature of galactic nuclei in general.  As the complicated morphology of the central 10 pc is being unveiled thanks to the dramatic progress of radio technology, effort is being made to understand whether these features are really associated with the Galactic center or just seen along the line-of-sight in that direction, and to determine the relative positions of them along the line-of-sight, i.e., the three-dimensional (3-D) spatial structure of the Galactic center.  Observations of 327~MHz absorption toward Sgr A West definitely place Sgr A East behind Sgr A West (\\citealt*{yus87}; see also \\citealt*{ped89}). \\citet*{mez89} observed ring-shaped 1.3~mm dust emission surrounding Sgr A East across the $50~\\kms$ cloud and the $20~\\kms$ cloud, and suggested that Sgr A East has expanded into these molecular clouds. Based on these observational arguments, \\citet{mez89} proposed a 3-D structure of the Sgr A complex and concluded that the event which created Sgr A East and the associated dust shell did not occur deep within the GMCs but close to their surfaces facing the sun.  However, \\citet*{geb89} found CO absorption toward a few Galactic center infrared (IR) sources and suggested that the $20~\\kms$ cloud may be located in front of Sgr A West. They also found some evidence that the $50~\\kms$ cloud lies partly in front of Sgr A West.  Based on the NH$_3$ morphology and kinematics observed using the Very Large Array (VLA), \\citet*{coi99,coi00} located the Galactic nucleus (defined to include Sgr A*, Sgr A West, and the CND throughout this paper) behind the southern streamer (and the $20~\\kms$ cloud; see also \\citealt*{gus80}) and the $50~\\kms$ cloud (or the northern part of the molecular ridge) slightly behind Sgr A East. They also argued that the distance between Sgr A East and the $20~\\kms$ cloud along the line-of-sight should be smaller than 8.4 pc, which is the size of the SNR G~359.92-0.09 in 20~cm radio continuum images \\citep*{yus87,ped89}.  \\citet*{her05} updated and modified the 3-D model of \\citet{coi00} based on their additional NH$_3$ line data and more recently published results (\\citealt*{mae02} and references therein; \\citealt{par04}), as follows.  The nuclear region is placed just inside the leading edge of Sgr A East. Only some part of the $50~\\kms$ cloud is located in front of the nucleus. The western streamer seen in NH$_3$ emission is highly inclined to the line-of-sight and is expanding outward with Sgr A East.  The northern ridge is placed along the northern edge of Sgr A East and is expanding perpendicular to the line-of-sight.  The southern streamer passes over the nucleus in projection but probably does not interact with it.  Together, the 3-D models above agree on the following features:  \\begin{enumerate} \\item The Galactic nucleus lies in front of Sgr A East but behind the southern streamer and part of the $20~\\kms$ cloud along the line-of-sight. \\item Sgr A East is expanding into the $50~\\kms$ cloud, the northern ridge, and the western streamer. \\item SNR G~359.92-0.09 is colliding with the southern part of the molecular ridge, the eastern edge of the $20~\\kms$ cloud, and the southern edge of Sgr A East. \\end{enumerate}  On the other hand, contradictions among the models raise the following questions.  \\begin{enumerate} \\item Is the nucleus in contact with or contained within Sgr A East? \\item Is the southern streamer falling into the nucleus? \\item Has Sgr A East expanded into the $50~\\kms$ cloud significantly, or just started to contact it? \\item Is Sgr A East colliding with the northern part of the molecular ridge? \\item Is Sgr A East in contact with the $20~\\kms$ cloud? \\item Is the $20~\\kms$ cloud located only in front of Sgr A East, or also extended further to the backside of it along the line-of-sight? \\end{enumerate}  It should be noted that the models above are all based on indirect evidence like morphology, kinematics of molecular clouds, or absorption of background radiation by these clouds, rather than on direct, physical interactions between the objects.  To answer some of the above questions directly, we have observed molecular hydrogen (H$_2$) emission and constructed a 3-D picture of the Galactic center. H$_2$ emission is an excellent tracer of interactions between dense molecular clouds and other hot and powerful objects, like Sgr A East.  In this paper we report observations of H$_2$ line emission from regions of interaction between Sgr A East and other gaseous components within the central 10 pc.  Our observations were almost entirely of the H$_2$ 1-0 S(1) line. Unlike most previous work, we observed this line at sufficiently high spectral resolution to resolve the velocity profiles and at high enough angular resolution to obtain detailed information on the spatial structure of the emission. We also obtained measurements of the H$_2$ 2-1 S(1) line at 2.2477~$\\mu$m at one location in order to investigate the excitation mechanism of the H$_2$.  We describe the observations in Section~\\ref{observation} and the reduction of the spectroscopic data in Section~\\ref{reduction}. Based on the directions of the shocks derived from the direct comparison of radial velocities with those from the NH$_3$(3,3) data of \\citet{mcg01}, we construct a 3-D model for the structure of the central 10~pc in Section~\\ref{structure}. In a forthcoming paper we discuss the properties of the shocks and estimate the explosion energy and age of Sgr A East, from which we constrain its origin. ",
        "conclusions": "Based on the H$_2$ emission map, we determine the outer boundary of Sgr A East where it is driving shocks into the surrounding molecular clouds, to be approximately an ellipse with the center at ($+32\\arcsec$, $+18\\arcsec$) or $\\sim1.5$~pc offset from Sgr A*, a major axis of length 10.8~pc, which is nearly parallel to the Galactic plane, and a minor axis of length 7.6~pc. This boundary is significantly larger than the synchrotron emission shell \\citep*{eke83,yus87,ped89} but is closely consistent with the dust ring suggested by \\citet*{mez89}.  Since Sgr A East is in physical contact with all of its nearby molecular clouds (the $50~\\kms$ cloud, the northern ridge, the molecular ridge, the southern streamer, the CND, and the western streamer),  we are able to determine the positional relationships between Sgr A East and the molecular clouds along the line-of-sight using the shock directions as indicators.  Based on the determined relationships and the strong evidence that Sgr A East is in contact with the nucleus, we suggest a revised model for the 3-D spatial structure of the central 10 pc of our Galaxy modifying the previous models of \\citet*{mez89}, \\citet*{coi00}, and \\citet*{her05}.  Our conclusions on the 3-D structure resolve most of the debates in previous studies as follows: \\begin{enumerate} \\item Is the nucleus in contact with or contained within Sgr A East? -- The Galactic nucleus is in physical contact with Sgr A East since the CND is pushed toward us by the expanding hot cavity of Sgr A East. \\item Is the southern streamer falling into the nucleus? -- It is highly probable that the southern streamer is falling into the nuclear region and feeding the CND. Sgr A East is driving shocks into the northern-most part of this cloud where it meets the CND in projection. \\item Has Sgr A East expanded into the $50~\\kms$ cloud significantly, or just started to contact it? -- In the H$_2$ data of the northeastern field, we can see that most of this region is filled with shocked gas from the $50~\\kms$ cloud. The area corresponds to at least one third of that of the entire cloud. Thus Sgr A East has significantly penetrated the cloud. If the cloud is wrapping around the Sgr A East shell with a casual morphological coincidence, however, the shocked layer might still be thin and on the surface. \\item Is Sgr A East colliding with the northern part of the molecular ridge? -- Yes, we detected shocked H$_2$ emission from the northern-most part of this cloud. \\end{enumerate}  However, we cannot answer the questions related to the $20~\\kms$ cloud. We know the branches (the southern streamer and the western streamer) from this GMC are interacting with Sgr A East but the main body of this cloud is located far to the south, beyond the scope of our observations. Therefore the position and extent along the line-of-sight of this GMC in Figure~\\ref{3D_model} is uncertain. According to \\citet*{coi99,coi00}, SNR G~359.92-0.09 is expanding into the molecular ridge, the $20~\\kms$ cloud, and Sgr A East. Thus, if we observe the H$_2$ emission around SNR G~359.92-0.09 in a similar way to the 3-D observations for Sgr A East, the questions about the $20~\\kms$ cloud could also be answered."
    },
    "0710/0710.5767.txt": {
        "abstract": "In studies of star-forming regions, near-infrared excess (NIRX) sources---objects with intrinsic colors redder than normal stars---constitute both signal (young stars) and noise (e.g. background galaxies). We hunt down (identify) galaxies using near-infrared observations in the Perseus star-forming region by combining structural information, colors, and number density estimates. Galaxies at moderate redshifts (z = 0.1 - 0.5) have colors similar to young stellar objects (YSOs) at both near- and mid-infrared (e.g. Spitzer) wavelengths, which limits our ability to identify YSOs from colors alone. Structural information from high-quality near-infrared observations allows us to better separate YSOs from galaxies, rejecting 2/5 of the YSO candidates identified from Spitzer observations of our regions and potentially extending the YSO luminosity function below K of 15 magnitudes where galaxy contamination dominates. Once they are identified we use galaxies as valuable extra signal for making extinction maps of molecular clouds. Our new iterative procedure: the Galaxies Near Infrared Color Excess method Revisited (GNICER), uses the mean colors of galaxies as a function of magnitude to include them in extinction maps in an unbiased way. GNICER increases the number of background sources used to probe the structure of a cloud, decreasing the noise and increasing the resolution of extinction maps made far from the galactic plane. ",
        "introduction": "The near-infrared (NIR) has been a valuable window into star-forming regions, providing us with a variety of tools to study the process by which dark clouds coalesce into stars: luminosity functions in the $K$ band (2.2$\\mu$m) can provide an accurate estimate of the initial mass function \\citep[e.g.][]{Muench:2002,Stolte:2005};  studies of embedded clusters reveal important details of the processes by which stars form either in groups or in isolation \\citep[e.g.][]{Lada:1991,Roman-Zuniga:2007}; and extinction mapping in the near-infrared allows precise determination of the column density structure of the cloud \\citep[e.g.][]{Alves:2001,Cambresy:2002}.  Of particular use is the narrow range of intrinsic near-infrared colors of main-sequence stars (typically $H-K$ = 0 - 0.4 and $J-H$ = 0 - 1.0). This narrow range is due both to the near-infrared bands lying on the Wien tail of all stellar (hydrogen-fusing)-temperature blackbodies and to the relative paucity and uniformity of absorption feature\\footnote{Principally H$^{-}$, but also CO features which are sensitive to surface gravity and produce a small split between dwarf and giant colors}. By assuming that all stars have the same intrinsic color we can do purely photometric extinction mapping. This is the heart of the Near Infrared Color Excess method (NICE) \\citep{Lada:1994} and the Near Infrared Color Excess method Revisited (NICER) \\citep{Lombardi:2001}.  Additionally, we can use the narrow range of stellar near-infrared colors to identify young stars by their near-infrared excess. The thermal contribution from their disk or envelope changes their color from that of a plain photosphere and places them in a certain region of a near-infrared color-color ($J-H$ versus $H-K$) diagram. This is the CTTS locus defined by \\citet{Meyer:1997}.  Unfortunately, a number of other astronomical objects also have intrinsic red colors, similar to CTTS. We refer to all intrinsically red objects as near-infrared excess (NIRX) sources. Of particular concern for studies of star forming regions are objects which are not either young-stellar objects (YSOs) or T-Tauri type young stars (TTS). These objects---brown dwarfs, a variety of evolved stars, galaxies, and AGN, are often studied in the near-IR in their own right, but in studies of star forming regions where we have limited information about their nature (perhaps only $J$,$H$, and $K$ photometry) they are contaminants which must be understood and identified in order to avoid common biases in these studies. For example, including intrinsically red objects in NICER produces an overestimate of the extinction.  In this paper we identify as galaxies a large number of NIRX sources at moderate redshifts (z = 0.1-0.5) in high-quality near-infrared images outside of the Perseus molecular cloud complex. After describing our observations and reduction in \\S\\ref{Obs} we present a detailed case that our NIRX sources are galaxies in \\S\\ref{ColorAndNumber}. In the heart of the paper we show how deep NIR images can help identify genuine YSOs in Spitzer observations (\\S\\ref{Clusters}) and how we can make use of these contaminating galaxies as additional probes for extinction mapping (\\S\\ref{GNICERSection}). ",
        "conclusions": ""
    },
    "0710/0710.0712_arXiv.txt": {
        "abstract": "We investigate molecular evolution in a star-forming core that is initially a hydrostatic starless core and collapses to form a low-mass protostar. The results of a one-dimensional radiation-hydrodynamics calculation are adopted as a physical model of the core. We first derive radii at which CO and large organic species sublimate. CO sublimation in the central region starts shortly before the formation of the first hydrostatic core. When the protostar is born, the CO sublimation radius extends to 100 AU, and the region inside $\\lesssim 10$ AU is hotter than 100 K, at which some large organic species evaporate. We calculate the temporal variation of physical parameters in infalling shells, in which the molecular evolution is solved using an updated gas-grain chemical model to derive the spatial distribution of molecules in a protostellar core. The shells pass through the warm region of $10 -100$ K in several $\\times$ $10^4$ yr, and fall into the central star $\\sim 100$ yr after they enter the region where $T \\gtrsim 100$ K. We find that large organic species are formed mainly via grain-surface reactions at temperatures of $20 -40$ K and then desorbed into the gas-phase at their sublimation temperatures. Carbon-chain species can be formed by a combination of gas-phase reactions and grain-surface reactions following the sublimation of CH$_4$. Our model also predicts that CO$_2$ is more abundant in isolated cores, while gas-phase large organic species are more abundant in cores embedded in ambient clouds. ",
        "introduction": "In a last decade, great progress has been made in our understanding of the chemical evolution of low-mass star-forming cores \\citep[][and references therein]{fra07, bt06}.  Among the observational advances are the detection of chemical fractionation in several prestellar cores; emission lines of the rare isotopes of CO are weaker at the core center than at outer radii, while emission lines of nitrogen-bearing species (e.g., N$_2$H$^+$) show relatively good correlation with the centrally-peaked dust continuum \\citep[e.g.][]{cas99, taf02}.  Theoretical models show that the CO depletion is caused by adsorption onto grains \\citep{bl97,aik01}, and that the CO depletion helps to temporarily maintain the N$_2$H$^+$ abundance at the core center. Theoretical models also predict that a fraction of the adsorbed CO will be hydrogenated to form H$_2$CO and CH$_3$OH on grain surfaces \\citep{ar77,hhl92}, which is confirmed by laboratory experiments \\citep{wkl02,fuchs07}. The existence of solid methanol in low-mass star formation regions has been confirmed observationally; \\citet{ppp03} detected a high abundance of CH$_3$OH ice (15-25 \\% relative to water ice) towards three low-mass protostars  among $\\sim 40$ observed protostars. Radio observations find gaseous CH$_3$OH to be abundant in the central regions of protostars \\citep{sjv02}. Since the formation of CH$_3$OH is inefficient in the gas phase \\citep{gpc06,gwh06}, it must be formed by the hydrogenation of CO on grain surfaces in the prestellar core stage, and then sublimated to the gas-phase as the core is heated by the protostar. Although CH$_3$OH ice is not detected towards the majority of low-mass protostars \\citep{ppp03} and the background star Elias 16 \\citep{chi96}, the upper limit for the CH$_3$OH ice abundance is a few $\\%$ relative to water ice, which is not low enough to contradict the idea that gaseous CH$_3$OH around protostars is originally formed by grain-surface reactions and then sublimated.  Two major questions, however, remain to be answered, the first being at what stage CO returns to the gas phase. The core remains nearly isothermal as long as  cooling by radiation is more efficient than heating by contraction (compression). Eventually, though, the heating overwhelms the cooling, so that the core center becomes warmer. A newly-born protostar further heats the surrounding core. Laboratory experiments show although some fraction of CO can be entrapped in water ice \\citep{cdf03}, a significant amount of CO sublimates at around 20 K \\citep{sa88}. In order to predict if an observable amount of CO returns to the gas-phase during the prestellar core stage or after the birth of a protostar, the temperature distribution in a core should be calculated by detailed energy transfer. Such a prediction is important in order to observe very young protostars.  Once it is sublimated, CO again becomes a useful observational probe. In addition, CO sublimation significantly affects the gas-phase chemistry; for example, it destroys N$_2$H$^+$.  The second question concerns how large organic molecules are formed in protostellar cores. In recent years, diverse organic molecules, including methanol (CH$_3$OH), dimethyl ether (CH$_3$OCH$_3$), acetonitrile (CH$_3$CN), and formic acid (HCOOH), have been detected towards low-mass protostars \\citep[][and references therein]{cec07}. They are still referred to as \"hot-core species\", because it was once thought that they are only characteristic of hot ($T\\sim 200$ K) cores in high-mass star forming regions. A large number of modeling studies on hot-core chemistry show that sublimed formaldehyde (H$_2$CO) and CH$_3$OH are transformed to other organic species by gas-phase reactions within a typical timescale of $10^4-10^5$ yr \\citep[e.g.][]{mh98}. In low-mass cores, however, the time scale for the cloud material to cross the hot ($T\\sim 100$ K) region should be smaller than $10^4$ yr, considering the infall velocity and temperature distribution in the core \\citep{bot04a}. Furthermore, theoretical calculations and laboratory experiments recently showed that  gas-phase reactions are much less efficient in producing some hot-core species, such as methyl formate (HCOOCH$_{3}$) and dimethyl ether, than assumed in previous models \\citep{hor04, gep06,gh06}.   Several model calculations have been performed on the chemistry that occurs in low-mass protostellar cores. \\citet{dot04} solved a detailed gas-phase reaction network assuming a core model for IRAS 16293-2422, and succeeded in reproducing many of the observed lines within 50 \\%. The physical structure of the core; i.e., its density and temperature distribution, was fixed with time. They assumed gas-phase initial molecular abundances that  pertain to the high-mass hot-core AFGL 2591. \\citet{lbe04}, on the other hand, constructed a core model that evolves from a cold hydrostatic sphere to a protostellar core by combining a sequence of Bonnor-Ebert spheres with the inside-out collapse model by \\citet{shu77}. They solved a chemical reaction network that includes gas-phase reactions along with adsorption and desorption of gas-phase/ice-mantle species. The resulting molecular distributions and line profiles are significantly different from those of the static core models \\citep{leb05}. The chemical network of \\citet{lbe04}, however, does not include large organic species.  In the present paper, we re-investigate molecular evolution in star-forming cores with the partial goal of answering the two questions posed above. We adopt a core model by \\citet{mi00}; it accurately calculates the radial distribution of temperature, which determines when and where the ice components are sublimated. The model also enables us to follow molecular evolution smoothly from a prestellar core to a protostellar core. The chemistry includes both gas-phase and grain-surface reactions according to \\citet{gh06}; the surface reactions in particular are important for producing organic molecules in a warming environment.  Here, we report a solution of the reaction network following the temporal variation of density and temperature in infalling shells to derive a spatial distribution of molecules, including complex organic ones, in a protostellar core. ",
        "conclusions": "\\subsection{Physical structure of the core and sublimation radius} We have re-analyzed and adopted the results of \\citet{mi00} to derive the sublimation temperatures of CO and large organic species, and to investigate molecular evolution in a star-forming core. The chosen conditions in the envelope can be different from those of other models because the temperature distribution in the envelope depends not only on the evolutionary state and mass of the central object (either the first or second core), but also on the mass distribution in the envelope, which should depend on the initial conditions.  Recently, \\citet{omu07} investigated the temperature distribution in the first-core envelope assuming mass distributions from the L-P model \\citep{lar69} and Shu model \\citep{shu77} similarity solutions. The former has a more massive envelope than the latter. When the first core has a mass of $0.05 M_{\\odot}$, for example, the hydrogen number density at $r=10$ AU is several $10^{10}$ cm$^{-3}$ and $\\sim 10^9$ cm$^{-3}$ with the L-P model and the Shu model, respectively. The L-P model gives a core luminosity of $10^{-1} L_{\\odot}$, and sublimation radii of $r_{\\rm 20K} \\sim 100$ AU and $r_{\\rm 100 K}\\sim 10$ AU, while the Shu model yields a  core luminosity of $\\sim 10^{-3} L_{\\odot}$ and an $r_{\\rm 20 K}$ of $\\sim 20$ AU. The latter model does not exceed 100 K at any radius. Our core model is warmer than the Shu model but colder than the L-P model.  In the second core stage, on the other hand, the envelope structure can deviate from spherical symmetry and be accompanied by a circumstellar disk. Considering a typical angular velocity for  molecular cloud cores \\citep[$\\sim 10^{-14}$ s$^{-1}$; e.g.,][]{ga85}, the centrifugal radius (i.e. the initial disk radius) is $\\sim 100$ AU. Our results at $r\\gtrsim$ several hundred AU are thus relatively robust, while the core structure would be significantly different from our model at smaller radii. For example, the large organic species could be sublimated by the accretion shock onto the forming disk rather than in the envelope, since the centrifugal radius $\\sim 100$ AU coincides with the sublimation radius of large organic species in our model.  \\subsection{Simple molecules} \\cite{lbe04} investigated the evolution of relatively simple molecules, such as HCN and N$_2$H$^+$, in a star-forming core by combining a sequence of Bonnor-Ebert spheres for the prestellar stage with the inside-out collapse model of \\cite{shu77} for the core after the first-core formation. These simple molecules are often observed in star-forming cores and are of importance as observational probes, while we mainly discussed  large organic species in \\S 3. Figure \\ref{cfLee} shows the radial distribution of simple molecules in our model at $t_{\\rm final}$.  Comparison with \\cite{lbe04} is not easy because there are many differences in the physical core models and chemical reaction networks. First, the inside-out core model, which is adopted in \\cite{lbe04}, yields lower temperatures than our core model, as discussed above. Secondly, we adopt different binding energies for molecules to the grain surface. While the difference is on the order of only a few hundred K for many species, the binding energy of NH$_3$ is significantly higher in our model (5534 K) than in \\cite{lbe04} (1082 K); the latter value seems to originate from \\cite{hh93}, in which the hydrogen bonding of NH$_3$ is not taken into account. Thirdly, grain-surface reactions, which are the main formation processes of saturated species such as NH$_3$ and H$_2$CO in our model, are not included in \\cite{lbe04}. Hence, here we compare our results with \\cite{lbe04} qualitatively rather than quantitatively.  A main conclusion of \\cite{lbe04} is that the molecular abundances vary significantly near and inside the CO sublimation radius. For instance, N$_2$H$^+$ is destroyed by CO, and thus declines steeply inwards across the CO sublimation radius, an effect which happens in our model as well. \\cite{lbe04} also found that some molecules, such as H$_2$CO, have their peak  abundance at the sublimation radius; the gas-phase abundance first increases inward via sublimation, but then decreases due to gas-phase reactions at smaller radii. Similar phenomena can be seen in our model, but the spatial variation of the gas-phase abundances is more moderate than in \\cite{lbe04}, for which we can think of a few reasons. First, \\cite{lbe04} calculated the molecular evolution in 512 shells, while we calculated the chemistry in only 13 shells. A larger number of shells are needed to resolve abundance fluctuations on a smaller radial scale. Secondly, higher infall speeds, in general, make the abundance distributions more uniform. Since the infall speed is higher at inner radii in the infalling envelope, it would be natural that the molecular abundances remain relatively uniform at $r\\lesssim 100$ AU, which is not calculated in \\cite{lbe04}. Thirdly, our chemical network includes a much larger number of species and reactions. \\cite{lbe04} used a reduced reaction network with $\\sim 80$ species and $\\sim 800$ reactions to save computational time, while our model includes 655 species and 6309 reactions. In a small reaction network, a sudden increase of a species (e.g., as caused by sublimation) can easily change the abundances of other species though chemical reactions. In a large reaction network, on the other hand, a larger number of reactions contribute to the formation and destruction of each species, and thus the sudden abundance change of one species does not necessararily propagate to other species. Another noticeable difference in our results from those of  \\cite{lbe04} is that HCO$^+$ decreases more steeply inwards at $r\\lesssim 1000$ AU in our model; H$_3$CO$^+$ and C$_3$H$_5^+$ are the dominant positive ions rather than HCO$^+$ at the inner radii.  The oxygen-bearing species atomic oxygen (O), molecular oxygen (O$_2$), and gaseous H$_2$O are also of special interest because O  is very reactive and because the  two molecules have been intensively observed by {\\it the Submillimeter Wave Astronomy Satellite (SWAS)} and {\\it Odin} Satellite in recent years.  Our model predicts that O  reaches an abundance of  $4\\times 10^{-5}$ relative to hydrogen nuclei at $r=8000$ AU, while it steeply decreases inwards from $\\sim 1000$ AU to 100 AU.  Even at $r=8000$ AU, however, H$_2$O ice is the most abundant O-bearing species (Figure \\ref{dist}). \\cite{ber00} summarized the {\\it SWAS} observations of cold molecular clouds by stating that the fractional abundance of O$_2$ lies under $10^{-6}$ and that of gaseous water lies in the range $10^{-9}$ to  a few $\\times$ $10^{-8}$. Molecular oxygen has recently been detected towards $\\rho$ Oph by Odin with an abundance of $5\\times 10^{-8}$ relative to hydrogen \\citep{lar07}. These abundances are consistent with our predictions for the outermost radius $r=8000$ AU, where the density $n_{\\rm H}\\sim 10^4$ cm$^{-3}$ and temperature $T<20$ K are similar to those in molecular clouds.   \\subsection{Dependence on visual extinction at the core edge}  So far we have assumed that the model core is embedded in ambient clouds of $A_{\\rm v}= 3$ mag. In reality, some cores (e.g. Bok globules) are isolated, while others are embedded in clouds. Isolated cores are directly irradiated by interstellar UV radiation, which causes photo-reactions (photodissociation and ionization) in the gas-phase and ice mantles \\citep[e.g.][]{lee96,rh01}. In order to evaluate the effect of ambient UV radiation, we recalculated molecular abundances in a core that is directly irradiated by interstellar UV radiation; i.e. $A_{\\rm v}=0$ mag at the outer edge of the core ($r=4\\times 10^4$ AU). The photodissociation rates of H$_2$ and CO were calculated by following \\citet{lee96}, which gives the shielding factors as a function of $A_{\\rm v}$ and column densities of CO and H$_2$ in the outer radii. The outermost shell for which we calculate molecular evolution is initially located at $r\\sim 1.4\\times 10^4$ AU and declines to $r=8000$ AU at $t_{\\rm final}$. Column densities of CO and H$_2$ outside of this shell were estimated by assuming $n$(CO)/$n_{\\rm H}=5\\times 10^{-5}$ and $n$(H$_2$)/$n_{\\rm H}=0.5$.  Figure \\ref{av0} shows the resultant distribution of molecular abundances in a protostellar core at $t_{\\rm final}$. Compared with the embedded core model (Figure \\ref{dist}), the fractional abundances of CH$_3$OH and H$_2$CO are lower by more than one order of magnitude, while the CO$_2$ abundance is higher. When the shells are still at $r\\gtrsim$ several thousand AU and have relatively low $A_{\\rm v}$ ($\\lesssim 4$ mag), the photodissociation of H$_2$O ice efficiently produces OH, which reacts with CO to produce CO$_2$ ice. Methanol in the ice mantle is dissociated to H$_2$CO, which is further dissociated to CO. Species such as CH$_3$CHO, HCOOCH$_3$ and CH$_3$OCH$_3$ are also less abundant in the isolated model, because their formation path includes H$_2$CO in the ice mantle. Formic acid in the gas phase extends only up to $\\sim 100$ AU, while it extends to several hundred AU in the embedded model. In the isolated model, it is formed mainly by OH + HCO on the grain surface. Our results may indicate that molecular abundances in hot cores and corinos depend on whether the core is isolated or embedded in clouds.  The importance of photo-reactions on ice mantle abundances has also been investigated in  a number of laboratory experiments.   For example, \\citet{wk02} and \\citet{wm07} found that carbon dioxide is efficiently produced by UV irradiation on  a binary ice mixture of H$_2$O and CO, a result that is consistent with our model. However, \\cite{wm07}  found that the UV irradiation also produces CH$_3$OH with a relative abundance of $n$(CH$_3$OH)/$n$(CO) $\\sim 10$ \\% in the ice mixture. Comparison of our model with their experiment is not straightforward because of  differences in temperature and included reactions,  but we may have underestimated the CH$_3$OH ice abundance in the irradiated core model. The discrepancy can arise from the H atom desorption rate in our model. We calculated that UV radiation penetrates into the ice mantle and dissociates H$_2$O to produce H atoms embedded in ice. Although we do not discriminate between H atoms on the ice surface and those embedded in the ice mantle, the embedded H atoms in reality would have a lower desorption rate and a better chance of reacting with neighboring CO, a reaction that initiates the formation of methanol in the mantles. Discrimination between surface and embedded hydrogen atoms should be included in future work.  \\subsection{Comparison with observation}  Our model results can be compared with observational results of low-mass protostars. Comparison of the physical structure of the core has  already been discussed in detail by \\citet{mi00}.  Here we concentrate on molecular abundances. First, we compare molecules in ice mantles. The observation of ices towards the low-mass protostar Elias 29 is summarized in \\citet{es00}; the abundances of CO, CO$_2$, CH$_3$OH and CH$_4$ relative to water ice are 5.6 \\%, 22 \\%, $< 4$ \\%, and $< 1.6$ \\%, respectively. On the other hand, \\citet{ppp03} detected a high abundance (15-25 \\% relative to water ice) of CH$_3$OH ice towards 3 low-mass protostars among 40 observed protostars. Although the observation of ice features is difficult, it is probable that the composition of ice mantles depends on their environment and the history of the observed regions. From a theoretical point of view, the abundances of molecules on grains and their fractional abundances compared with H$_2$O ice depend on time and radius from the protostar. In addition, an embedded core and an isolated core lead to significantly different calculated abundances for solid CH$_3$OH, H$_2$CO and CO$_2$.  Table~\\ref{solid_theo} lists ratios of ice species with respect to water ice at $t_{\\rm final}$ for local abundances at radii of 1000 AU and 8000 AU and for column densities towards the core center. Since most of the core material exists at $r\\le 8000$ AU along the line of sight, and since the ices are present at $r>10$ AU, the column density is calculated by integrating the number density of ice species from 10 AU to 8000 AU. It is interesting to note that regardless of the distance from the protostar, the isolated core model gives smaller surface abundances of CH$_3$OH and H$_2$CO and a higher surface abundance of CO$_2$ than the embedded model. In both models, the surface abundance of CO at 1000 AU is much smaller than observed, but it increases towards larger radii up to 10 and 20\\% in the embedded and isolated models respectively. Comparison with the observations is best done using our column density ratios, which are in reasonable agreement with Elias 29 for both models. Considering the variation among cores, our models show reasonable agreement with observation.  The disagreement with CO doubtless results from our assumption concerning desorption rates. In the present work, we use a desorption energy for each species mainly referring to laboratory experiments on pure ice sublimation or theoretical estimates that sum up the van der Waals forces between adsorbed atoms and grain surface. But in reality, interstellar ice is a mixture, and hence volatile species can be entrapped by less volatile species; recent temperature-programmed desorption results show that in a mixture rich in water ice, much CO desorbs at considerably higher temperatures \\citep{cdf03}.  Table \\ref{obs} lists the gas-phase abundances of large molecules towards the low-mass protostar IRAS 16293-2422 and in the central region ($r=30.6$ AU) of our core model. The physical and chemical structures of IRAS 16293-2422 have been intensively studied \\citep{cec00a,cec00b, sch02,cau03, cha05}. The physical parameters derived by the model at $t_{\\rm final}$ ($\\sim 23 L_{\\odot}$ and $r_{\\rm 100K}\\sim1.2\\times 10^2$ AU) are very close to the ones of IRAS16293-2422; the observed luminosity of the source is 27 $L_{\\odot}$ \\citep{wal86} and the physical structure has been constrained by multi-line analysis of H$_2$O and H$_2$CO observations through a detailed radiative transfer code \\citep{cec00b}. It should be noted that the emission lines of large organic molecules observed in this source are not spatially resolved, except for a few lines investigated by interferometric observations \\citep{bot04b, kua04, rem06}. The density and temperature of the gas should vary both within the beam and along the line of sight. Hence the derived molecular abundances vary significantly depending on the assumptions concerning the physical structure of the core and  the emitting region, which explains the difference in abundances determined by different investigators (see Table~ \\ref{obs}).  Although it is not obvious which core model, embedded or isolated, should be compared with IRAS16293-2422, we would prefer the embedded core model. While IRAS16293-2422 is in the Ophiuchus molecular cloud, which harbors several UV sources in the form of OB stars, the $^{13}$CO and C$^{18}$O observations indicate that the core is embedded in dense gas \\citep{lor89,tac00}. The high molecular D/H ratios observed towards IRAS 16293-2422 and its neighboring starless core IRAS16293E \\citep{cec07,vas04} indicate that these cores have been very cold and thus well-shielded from the ambient stellar radiation. Considering the uncertainties in observationally-estimated molecular abundances, HCOOCH$_3$, HCOOH, and CH$_3$CN in our embedded core model show reasonable agreement with the observations. The embedded model, however, underestimates CH$_3$OCH$_3$ and overestimates H$_2$CO and CH$_3$OH. The isolated core model, on the other hand, reproduces observed abundances of CH$_3$OH, HCOOH and CH$_3$CN, but strongly underestimates H$_2$CO, HCOOCH$_3$, and CH$_3$OCH$_3$.  We can think of several possible solutions to improve the agreement with observation. First, it is noteworthy that the gaseous CH$_3$OH abundance estimated in IRAS 16293-2422 is much lower than the abundance of CH$_3$OH ice ($n$(CH$_3$OH ice)/$n_{\\rm H} \\sim 10^{-5}$, assuming n(H$_2$O ice)/ $n_{\\rm H}\\sim 10^{-4}$) detected by \\citet{ppp03} towards some low-mass protostars. We may have missed or underestimated the reactions which transform CH$_3$OH to other large organic species. Secondly, the abundances of large organic species in the central region can vary with time. In the present work, we have concentrated on the distribution of molecules only at $t_{\\rm final}$. Shells that reach the central region at different times should experience different temporal variations of physical conditions. Some shells may experience longer periods at $T\\sim 20-40$ K, where large organic species start to be efficiently formed. This possibility will be pursued in a future work. Thirdly, core models with rotation could produce higher abundances of large organic species; because of the centrifugal force, core material migrates more slowly and stays longer in the temperature range preferable for the formation of large organic molecules (e.g. in a forming disk).  Because of the beam size of the current radio observations, little is known concerning the spatial distribution of the large organic species; they are mostly confined within a few arc seconds from the core center \\citep[e.g.][]{kua04}. But there are exceptions; \\citet{rem06} found that HCOOH and HCOOCH$_3$ show extended emission of $\\sim 5$ arcsec. Our embedded core model reproduces the extended emission of HCOOH, while HCOOCH$_3$ is confined to $r<100$ AU.    \\subsection{Carbon chains in a protostellar core}  Recently, \\cite{sshk07} detected strong emission lines of carbon chain species such as C$_4$H and C$_4$H$_2$ towards the low-mass protostar L1527, which is considered to be in a transient phase from class 0 to class I. In general, carbon-chain species are associated with  the early stages of a cold cloud core when the dominant form of carbon changes from atomic carbon to CO. Hence, the existence of carbon-chain species towards L1527 is a surprise.  Figure \\ref{carbon_chain} (a-b) shows the temporal variation of CH$_4$ and carbon-chain abundances in the shell that reaches $r=2.5$ AU at $t_{\\rm final}$.  When the core starts to contract, methane (CH$_4$) and C$_3$H$_4$ are already abundant in ice mantles.  Methane has been formed by the hydrogenation of carbon on grain surfaces, while C$_3$H$_4$ has been formed by a combination of gas-phase reactions (to form unsaturated carbon chains such as C$_{3}$H and C$_{3}$H$_{2}$) and grain-surface reactions (to hydrogenate them). When the CH$_4$ sublimates, some fraction reacts with C$^+$ to form C$_2$H$_3^+$, which is a precursor for ion-molecule reactions that lead to the production of larger unsaturated hydrocarbons. While gas-phase reactions tend to make the carbon chains longer,  adsorbed species experience hydrogenation and dissociation by cosmic-ray induced UV radiation. Figure \\ref{carbon_chain} (c-d) shows the distribution of carbon-chain species in our embedded core model. Methane and C$_3$H$_4$ are abundant in the ice mantle even at $r\\gtrsim 10^3$ AU, while other hydrocarbons are abundant at a few 100 AU $\\lesssim r \\lesssim 10^3$ AU. Most of the carbon-chain species  desorb at $r\\lesssim$ a few 100 AU.  In summary, our model indicates that carbon-chain species can be re-generated in the protostellar core by the combination of gas-phase reactions and grain-surface reactions. It should be noted, however, that large oxygen-containing organic species are not detected in L1527 \\citep{sak07}, while they are abundant in the central region of our model. In future work, the temporal variation of the molecular distribution in model cores should be investigated to see if at certain evolutionary stages carbon-chain species are abundant while large organic species are not.  For example, in their recent study of the chemistry of cold cores, \\cite{gwh06} found that generation of gas-phase hydrocarbons from precursor surface methane occurs at very late times, after the abundance of surface methanol has essentially vanished. More detailed  observations are also needed for a quantitative comparison with models. While our model predicts the column density of C$_4$H to be $2\\times 10^{16}$ cm$^{-2}$ towards the central star, the observation gives $2\\times 10^{14}$ cm$^{-2}$. The actual column density towards the central star can, however, be larger, because the emission is averaged within the beam (a few 10 arcsec) in the current observation. Methane ice, the precursor of the carbon-chain species in our model, is as abundant as 23 \\% relative to water ice at the outer edge of our core model. A deep integration of ice features towards field stars is needed to constrain the CH$_4$ ice abundances at the outer edge of the core and/or quiescent clouds, while the column density ratio of CH$_4$ ice to H$_2$O ice (1 \\%, see Table 2) in our model is consistent with the observation towards YSO's (\\S 4.3), because of the relatively large sublimation radius of CH$_4$."
    },
    "0710/0710.1839_arXiv.txt": {
        "abstract": "Black hole mass (\\mbh) is a fundamental property of active galactic nuclei (AGNs). In the distant universe, \\mbh\\ is commonly estimated using the MgII, \\hb, or \\ha\\ emission line widths and the optical/UV continuum or line luminosities, as proxies for the characteristic velocity and size of the broad-line region. Although they all have a common calibration in the local universe, a number of different recipes are currently used in the literature. It is important to verify the relative accuracy and consistency of the recipes, as systematic changes could mimic evolutionary trends when comparing various samples.  At $z=0.36$, all three lines can be observed at optical wavelengths, providing a unique opportunity to compare different empirical recipes.  We use spectra from the Keck Telescope and the Sloan Digital Sky Survey to compare \\mbh\\ estimators for a sample of nineteen AGNs at this redshift. We compare popular recipes available from the literature, finding that \\mbh\\ estimates can differ up to $0.38\\pm0.05$ dex in the mean (or $0.13\\pm0.05$ dex, if the same virial coefficient is adopted). Finally, we provide a set of 30 internally self consistent recipes for determining \\mbh\\ from a variety of observables. The intrinsic scatter between cross-calibrated recipes is in the range $0.1-0.3$ dex. This should be considered as a lower limit to the uncertainty of the \\mbh\\ estimators. ",
        "introduction": "Understanding the growth of supermassive black holes along with their host galaxies is one of the fundamental questions in current astrophysics \\citep[e.g.][]{DMS05,Cro++06b}.  Black hole mass (\\mbh) is a key parameter in revealing the nature of black hole-galaxy coevolution as well as the physics of active galactic nuclei (AGNs). However, direct mass measurements using the motions of gas and stars in the sphere of influence of a central black hole is limited to very nearby galaxies \\citep[e.g.][]{K+G01,F+F05}.  Beyond the very local universe, the so-called ``virial'' or ``empirically calibrated photo-ionization'' method based on the reverberation sample is popularly used for active galaxies \\citep[e.g.][]{WPM99, Kas++00, Kas++05, Ben++06a}.  This method utilizes broad line widths as velocity indicators and monochromatic continuum or line luminosities as indicators of broad-line region size, hence estimating virial \\mbh. A combination of the MgII, \\hb, or \\ha\\ broad emission line widths and the 3000\\AA, 5100\\AA, \\hb, or \\ha\\ luminosities is typically used, depending on the redshift of the source and the observational setup. Several equations have been presented in the literature to estimate \\mbh\\ using various combinations of these indicators \\citep[e.g.,][]{W+U02b,W+U02a,M+J02,TMB04,Kol++06,G+H05,V+P06,Woo++06,Sal++07,N+T07,Tre++07}.  Although all three emission lines have a common calibration based on the reverberation sample in the local universe, it is important to verify that different recipes give consistent results; any systematic changes could mimic evolutionary trends given that different recipes are often used in various studies.  At $z=0.36$, all three lines can be observed at optical wavelengths, providing a unique opportunity to cross-calibrate the different methods of \\mbh\\ estimation. Using data from the Keck Telescope and the Sloan Digital Sky Survey for a sample of nineteen AGNs at $z=0.36$, we compare the different methods of estimating \\mbh, and derive a set of self-consistent equations for \\mbh\\ estimates using every combination of velocity scale (FWHM and line dispersion \\sline\\, of MgII, \\hb, or \\ha) and luminosity (3000\\AA, 5100\\AA - nuclear and total - \\hb, or \\ha).  The paper is organized as follows.  In \\S~\\ref{sec:data} we describe the sample selection, observations, and data reduction.  In \\S~\\ref{sec:meas} we describe our line fitting process, based on expansion in Gauss-Hermite series, and the resulting luminosity and width measurements.  In \\S~\\ref{sec:form} we review the various formulae adopted in the literature, and compare the various \\mbh\\ estimators. In \\S~\\ref{sec:res} we present our self-consistent recipes. Section~\\ref{sec:sum} summarizes our results.  Throughout this paper magnitudes are given in the AB scale. We assume a concordance cosmology with matter and dark energy density $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$, and Hubble constant H$_0$=70 kms$^{-1}$Mpc$^{-1}$.  \\begin{figure*} \\epsscale{0.8} \\plotone{f1.eps} \\caption{Flux-calibrated spectra.  The SDSS spectra are shown in black, and the Keck spectra are shown in blue and red.  The MgII line can be seen on the far left of each wavelength range, while \\hb\\ is located in the center and \\ha\\ to the far right. } \\label{spectra1} \\end{figure*}  \\begin{figure*} \\epsscale{0.8} \\plotone{f2.eps} \\caption{As in Figure~\\ref{spectra1} for objects S11 to S29.} \\label{spectra2} \\end{figure*} ",
        "conclusions": "\\label{sec:sum}  In this paper we have used Keck and SDSS spectra of nineteen Seyferts at $z=0.36$, to perform a comprehensive study of ``virial'' black hole mass estimators for broad line AGNs.  The main results can be summarized as follows:  \\begin{enumerate}  \\item We have fit Gauss-Hermite series to the data in order to measure the FWHM and \\sline\\, of MgII, \\hb\\,, and \\ha\\,, as well as \\ha\\, and \\hb\\, luminosities and continuum luminosities at 3000\\AA\\ and 5100\\AA.  Measurement errors are approximately 0.02 dex on the MgII and \\hb\\ line widths, 0.04-0.05 on the \\ha\\ line widths, 0.01 dex on the continuum luminosity, 0.01 dex on the \\hb\\ luminosity, and 0.06 dex on the \\ha\\ luminosity.  \\item We have compared twelve formulae taken from the literature, showing that \\mbh\\ estimates can differ systematically by as much as $0.38\\pm0.05$ dex (or $0.13\\pm0.05$ dex, if the same virial coefficient is adopted).  Such differences should be taken into account when comparing data obtained with different methods.  \\item We have cross-calibrated a set of 30 empirical recipes based on all combinations of the velocity and luminosity indicators corresponding to the \\ion{Mg}{2}, \\hb, and \\ha\\ broad lines.  Taking the masses measured by Treu et al. (2007) as our fiducial black hole masses, we find that: the absolute scale of the different indicators is calibrated to within $\\sim$0.05 dex; the best agreement is found when using the line dispersion of \\hb\\, as a velocity estimator, with the residual 0.1 dex r.m.s. scatter resulting from the various continuum luminosity estimators; adopting the line dispersion of \\ion{Mg}{2} raises the scatter to 0.2 dex; for the other estimators the intrinsic scatter is in the range 0.2-0.38 dex. This implies a lower limit of 0.1-0.2 dex on the validity of each estimator for each individual case.  \\end{enumerate}  The newly calibrated recipes should be useful to reduce the sources of systematic uncertainties when comparing different studies."
    },
    "0710/0710.3461_arXiv.txt": {
        "abstract": "We present the results of a deep, wide-field search for transiting `Hot Jupiter (HJ)' planets in the globular cluster $\\omega$ Centauri. As a result of a 25-night observing run with the ANU 40-inch telescope at Siding Spring Observatory, a total of 109,726 stellar time series composed of 787 independent data points were produced with differential photometry in a 52$'$$\\times$52$'$ (0.75 deg$^2$) field centered on the cluster core, but extending well beyond. Taking into account the size of transit signals as a function of stellar radius, 45,406 stars have suitable photometric accuracy ($\\le$0.045 mag to V$=$19.5) to search for transits. Of this sample, 31,000 stars are expected to be main sequence cluster members. All stars, both cluster and foreground, were subjected to a rigorous search for transit signatures; none were found. Extensive Monte Carlo simulations based on our actual data set allows us to determine the sensitivity of our survey to planets with radii $\\sim$1.5R$\\rm{_{Jup}}$, and thus place statistical upper limits on their occurrence frequency $F$.  Smaller planets are undetectable in our data. At 95$\\%$ confidence, the frequency of Very Hot Jupiters (VHJs) with periods P satisfying 1d~$<$~P~$<$~3d can be no more than $F_{\\rm VHJ} < 1/1040$ in $\\omega$ Cen. For HJ and VHJ distributed uniformly over the orbital period range 1d~$<$~P~$<$5d, $F_{\\rm VHJ+HJ} < 1/600$. Our limits on large, short-period planets are comparable to those recently reported for other Galactic fields, despite being derived with less telescope time. ",
        "introduction": "The identification and subsequent study of extrasolar planets has become a subject of intense interest in recent years. To date, $\\sim$250 giant planets have been discovered\\footnote{http://exoplanet.eu/}, ranging in mass from Neptune to greater than Jupiter. These new worlds are altering our understanding of the formation and evolutionary processes of giant planets in the immediate solar neighborhood. Currently, the majority of them have been found using radial velocity (RV) techniques, which favours the detection of close-in, massive planets.  Statistical analysis of current RV detections indicates that 1.2$\\%$$\\pm$0.3$\\%$ of nearby F,G and K stars are orbited by `Hot Jupiter' (HJ) planets, those with orbital periods of only a few days ($\\le$0.1 AU) and minimum masses approximately equal to that of Jupiter \\citep{M2005}. Such discoveries have challenged traditional ideas of planetary evolution, implying that a rapid migration of the planet takes place soon after formation.  Planet frequency from radial velocities appear to depend on the metallicity of the host star \\citep{G1997,L2000,S2001,FV2005}. However, there is very little observational evidence for a lower planetary frequency at quite low metallicities, due to the bias of radial velocity detections in association with bright nearby high metallicity stars. Hence dedicated surveys in low metallicity environments, like globular clusters, can provide information to help understand this relationship in a more robust manner. Indeed, does low metallicity halt planet formation or just affect the planetary migration process?  Of the currently known planets, $\\sim$90$\\%$ have only a minimum mass assigned to them, due to the unknown inclination of the planetary orbit, and unknown radii and densities. These quantities can be measured for transiting planets. Due to its short orbital period, each HJ has a non-negligible probability of transiting its host star that depends on the orbital separation and the ratio between the stellar and planetary radii. Typical transit depths and durations are $\\sim$1.5$\\%$ and $\\sim$2 hours respectively. With precise photometry these transiting systems can be identified in the field, leading to direct measurements of the planetary radius (provided the stellar radius is known) from the depth of the transit dip. Studies of the planetary atmosphere may be attempted if transits occur, and when coupled with RV measurements, accurate mass and density determinations can be made.  Currently, 28 transiting exoplanets are known \\citep{C2000,He2000,K2003,K2005,B2004,B2005,A2004,P2004,S2005,MCC2006,Bakos2006,Bakos2007,Bakos2007c,OD2006,OD2007,CC2006,Sahu2006,Bu2007,G2007,Mand07,Barb2007,Kov2007,N2007}, only a handful of which were first discovered through radial velocity searches. Despite an increasing number, the detection rate of planets from transit searches is significantly lower than initially expected eg, \\citet{H2003}. This lack of detections is due, in part, to observing strategy: a long observing window ($\\ge$1 month equivalent) with a dedicated telescope coupled with a wide field and high temporal resolution are needed to sample enough stars frequently enough to allow the detection of a transit. Also the production of a large enough number ($\\sim10^{5}\\rightarrow10^{6}$) of high-quality ($\\le$0.02 mag) lightcurves of dwarf stars coupled with a sensitive detection algorithm with low false alarm rates is non-trivial.  Recently, \\citet{Gouldetal2006} analyzed OGLE~III transit surveys in Galactic fields and concluded that the occurrence frequencies of the detected planets in these surveys is not statistically different from that found in RV surveys of nearby stars. However, they did conclude that the frequency of HJ planets with periods P satisfying 3$<$P$<$5 days ($F =$~1/320(1$^{+1.37}_{-0.59}$) at a 90$\\%$ confidence) was statistically different from and larger than that of Very Hot Jupiter (VHJ) planets (1$<$P$<$3 days), which have $F = $1/710(1$^{+1.10}_{-0.54}$) at 90$\\%$ confidence.  Since the OGLE~III surveys detected no planets with radius larger than 1.3R$_{\\rm{Jup}}$,they placed upper limits on the occurrence frequency of larger worlds.  Transit searches, unlike RV, are not limited to the immediate solar neighborhood and can be used to measure relative planet frequencies in various regions of the Galaxy, providing information on the role environmental effects play on HJ planet formation. Early predictions for the success of transit searches in open clusters was presented by \\citet{J1996}, indicating that with a large amount of telescope time planets could be detected via the transit technique in nearby clusters. More recently, \\citet{PG2005} presented an analysis of the prospective harvest of cluster transit surveys by discussing the observational techniques and methods to maximise the chances of a detection. They concluded that due to their mass functions, the most populous, nearby and bright clusters have the greatest chance of yielding a planet. \\citet{PG2006} then went on to discuss specifically the detection of short period `Hot Earth' and `Hot Neptune' planets. The detection yields for various nearby clusters with various instruments was estimated. \\citet{Ai2007} agree that small-aperture wide field surveys targetting nearby clusters have the potential to discover transiting Hot Neptune planets.  Transit searches have been undertaken in bright metal-rich open clusters including STEPPS \\citep{Bu2005}, UStAPS \\citep{S2003,Ho2005}, PISCES \\citep{M2003,Mo2005,Moch2006}, $\\it{Monitor}$ \\citep{A2007}, EXPLORE-OC \\citep{V2005} and in the Praesepe cluster (M44) with KELT \\citep{Pepper07}. Searches have also been performed in the general Galactic field \\citep{U2002,U2004,H2005,Ka2005,W2007a} and toward the Galactic Bulge \\citep{Sahu2006}. If cluster candidates are confirmed as planetary in nature, difficult if fainter than V$\\sim$17.0, they can provide information on the timescales of HJ formation and subsequent migration. Null results of high significance allow planet frequency upper limits to be estimated.  Globular clusters provide an excellent opportunity to study the effects of environment on planetary frequency. Two bright, nearby southern clusters, 47 Tuc and $\\omega$ Centauri, have stars in sufficient numbers ($\\sim$10$^{5}$) and brightness (V$\\leqslant$17) for meaningful statistics to be gained using ground-based telescopes of moderate aperture. 47 Tuc was previously sampled for planetary transits, resulting in a high significance ($>$3$\\sigma$) null result in both the cluster core \\citep{G2000} and in the outer halo \\citep{W2005}. These two results strongly indicate that system metallicity - not crowding - is the dominant factor determining HJ frequencies in this cluster.  This paper presents the results of a dedicated transit search in the second cluster, $\\omega$ Centauri, in order to test further the dependence of planetary frequency on stellar metallicity and crowding. Omega Centauri has only 1/10$^{\\rm{th}}$ the core density of 47 Tuc yet contains five times the total mass (5.1$\\times$10$^6$M$_{\\odot}$, \\citet{Meylan1995}). Due to its low stellar density compared to other globular clusters and long stellar interaction timescale, a null result for $\\omega$ Centauri can be used to test the relative importance of stellar metallicity over density in the formation of giant planets.  Omega Centauri ($\\omega$ Cen, NGC 5139) has been subjected to intense research over the years. The cluster is unique among globular clusters in that it displays a distinct spread of metallicity among its stars \\citep{DW1967,NB1975,Lee1999,P2000,Sol2005}, due to an extended period of star formation and chemical enrichment. Using He abundances, \\citet{N2004} has shown that the cluster has three distinct stellar populations, with metallicities of $-$1.7, $-$1.2 and $-$0.6 dex, corresponding to 0.80, 0.15 and 0.05 of the total population respectively.  The cluster has a highly-bound, retrograde orbit \\citep{D1999} and is by far the most massive of the globular clusters \\citep{Meylan1995}. Indeed, these vagaries have led to the theory that the cluster had an external origin, being the left-over remains of a tidally disrupted dwarf galaxy \\citep{BF2003,IM2004,BN2005}. With its relative proximity, $\\omega$ Cen presents a statistically significant number of upper main sequence stars that can be searched for transiting HJ planets.  Here we present the result from a vigorous search for the transit signatures of large planets on 45,406 lightcurves in a 0.75deg$^2$ field centered on $\\omega$ Cen. The same set of observations yielded a total of 187 variable stars in the field, 81 of which are new discoveries, and are presented in a companion paper \\citep{W2006}. Furthermore, we observed a control field in the Lupus Galactic Plane to test the data reduction and transit identification strategies. Analysis for this field is ongoing, but has led to the identification of several transit candidates, of which none similar were seen in the $\\omega$ Cen dataset. One candidate in particular has excellent prospects for being a new Hot Jupiter planet \\citep{W2007a}.  Section 2 of this paper describes our observational strategy and data reduction details. Section 3 details how the photometry was obtained for both the crowded core regions and outer halo parts of the dataset. The cluster Color Magnitude Diagram dataset (along with astrometry) are also briefly discussed. Section 4 describes the stellar parameters of the cluster stars that were searched for transits, and the expected characteristics of the transits themselves. The total number of stars in the field (both in the cluster and the foreground galactic disk) is also calculated. Section 5 describes the transit detection algorithms used in our search and our removal of systematics in the photometry. Our Monte Carlo simulations to derive expected recovery of real transits and false alarms are described in Section 6, and their application to estimate our HJ sensitivity outlined in Section 7. The results of our transit search in $\\omega$ Cen are presented in Section 8, with discussion, comparison to the literature and interpretation in Section 9. We conclude in Section 10. ",
        "conclusions": "We have presented the results of a wide-field, deep photometric search for transiting short-perod planets in the globular cluster $\\omega$ Centauri, a region previously un-sampled for planetary transits. The cluster was observed with a 52$'$$\\times$52$'$ field of view for 25 contiguous nights with the ANU 40-inch telescope at Siding Spring Observatory. From application of difference imaging analysis, a total of 109,726 time series were produced across the field, each being composed of 787 independent data points.  A total database of 45,406 stars have photometric accuracy suitable for the search ($\\le$0.045 mag scatter down to V$=$19.5), including 31,000 cluster stars extending 2.5 magnitudes down the main sequence. All of these were subjected to a rigorous (and vigorous) search for transit-like events; none were detected after variable stars and clear false-positives were removed.  Simulations have shown that if large Hot Jupiters (HJs) formed in the cluster then dynamically speaking they would survive to be detectable in our data.  Extensive Monte Carlo simulations via injection of transit signals into actual light curves were used to determine the sensitivity of the survey to R$\\le$1.5R${\\rm_{Jup}}$ planets over a range of orbital periods.  Coupled with our null result, we are thus able to place strict, statistically significant upper limits on the occurrence frequency $F$ of large (R~$=$~1.5R), short-period planets in $\\omega$ Centauri.  We determine a limit of $F_{\\rm VHJ} < 1/1040$ at 95$\\%$ confidence for Very Hot Jupiter (VHJ) planets with periods distributed uniformly over 1d~$<$~P~$<$~3d. This upper limit for the cluster is less than that determined by \\citet{Gouldetal2006} for smaller (1.3R${\\rm_{Jup}} < $~R~$< 1.5$R${\\rm_{Jup}}$) planets with the same period distribution in the Galactic fields surveyed by OGLE~III.  The two results are consistent at the 90$\\%$ confidence level, and more understandable if the low metallicity of $\\omega$ Cen suppresses planet formation or planetary migration.  Under the assumption that there is no difference in occurrence frequency for VHJ and HJ across the orbital period range 1d~$<$~P~$<$5d, we derive an upper limit of $F_{\\rm VHJ+HJ} < 1/600$ in $\\omega$ Cen. The corresponding result in the Galactic OGLE~III fields for comparably-sized planets is an upper limit of $F < 1/640$.  Both results are quoted at 95$\\%$ confidence.  Our results are less dependent on model assumptions about the distance to the target population since the vast majority of stars in our fields are members of, and thus at the distance of, $\\omega$~Cen.  It is noteworthy that despite the fact that the OGLE~III campaigns monitored considerably more stars in better median seeing with more frames per field, our $\\omega$ Cen study produces a comparable upper limit on the frequency of large, short-period planets. While part of the reason may lie in the longer exposure times, denser sampling, and the use of different cleaning and detection algorithms in our survey, a large part of the difference is due to the small fraction of the more distant and obscured Galactic bulge stars (which constituted about a third of the total OGLE~III sample) that can be meaningfully probed for transiting planets.  This null result for VHJ and HJ planets in $\\omega$ Cen, coupled with the null result of 47 Tucanae \\citep{W2005} strengthens the evidence for the dominance of system metallicity over stellar interactions in determining short period planetary frequencies in globular clusters. At longer orbital periods stellar encounters may play a role in determining planetary frequencies. This is a result aligned with current work on the metallicity trend of planet-bearing host stars in the Solar Neighborhood and N-body simulations of planets in dense environments. Such a metallicity dependence is one of the main predictions of the core accretion model of planet formation."
    },
    "0710/0710.2128_arXiv.txt": {
        "abstract": "We have resolved the classical nova V1663 Aql  using long-baseline near-IR interferometry  covering the period from $\\sim$5--18 days after peak brightness. We directly measure the shape and size of the fireball, which we find to be asymmetric. In addition we measure an apparent expansion rate of  $0.21 \\pm 0.03$ ${\\rm mas\\,day^{-1}}$. Assuming a linear expansion model we infer a time of initial outburst approximately 4 days prior to peak brightness. When combined with published spectroscopic expansion velocities our angular expansion rate implies a distance of $8.9\\pm3.6$ kpc. This distance measurement is independent of, but consistent with, determinations made using widely available photometric relations for novae. ",
        "introduction": "Novae are violent stellar explosions exceeded in energy output only by $\\gamma$-ray bursts and supernovae. They are erratic outbursts that occur in systems containing a white dwarf accreting mass from a late-type stellar companion (e.g Prialnik \\& Kovetz\\nocite{pk05} 2005). When the amount of accreted material on the surface of the white dwarf reaches some critical value a thermonuclear-runaway is ignited, giving rise to the observed nova outburst in which material enriched in heavy elements is ejected into the surrounding medium at high velocities. For certain elements, this ejected material may influence observed abundances in the ISM \\citep{gehrz98,hernanz05}. Direct observations of the expansion of the nova shell provide an opportunity to accurately determine the distance to the nova. Such observations are usually only possible many months or years after the outburst, when the expanding shell can be resolved \\citep{bode02}. However, several optical/IR interferometers are now capable of resolving novae and other explosive variables, allowing detailed studies of the initial \"fireball\", \\citep{ches07,mon06}  as well as later stages of development \\citep{lane05,lane07}.   Nova Aquilae 2005 (ASAS190512+0514.2, V1663 Aql) was discovered on 9 June 2005 by \\citet{iauc8540}.  At the time of discovery the magnitude was $m_V$ = 11.05; the source reached $m_V \\sim 10.8$ the following day, and declined in brightness thereafter. A possible progenitor near the source coordinates (sep. $\\sim$ 4.5 arcseconds) is seen on Palomar Optical Sky Survey plates (USNO-B1.0 0952-00410569, \\citet{usno-b1}), with magnitudes $m_R \\sim 18.1$ and $m_I \\sim 16.45$.  Soon after discovery \\citet{iauc8544} obtained an optical spectrum with features indicating a heavily reddened, peculiar nova. H-$\\alpha$ emission lines exhibited P Cygni line profiles and indicated an expansion velocity of $700 \\pm 150 {\\rm km\\,s^{-1}}$ \\nocite{iauc8544} (Dennefeld et al. 2005, confirmed via personal communication), somehwhat slow for a classical nova, but not outside the range of observed values. Recently, \\citet{pog06} published spectra and analysis of published light-curves of this nova, deriving a distance in the range 7.3--11.3 kpc, and an expansion velocity of $\\sim2000$~${\\rm km\\,s^{-1}}$.  We have used the Palomar Testbed Interferometer (PTI) to resolve the $2.2 \\mu$m emission from V1663 Aql and measure its apparent angular diameter as a function of time. We are able to follow the expansion starting $\\sim 9$ days after the initial explosion; when combined with radial velocities derived from spectroscopy we are able to infer a distance to, and luminosity of, the object. We compare this result with values inferred by a maximum magnitude-rate of decline (MMRD) relation in the literature (see Poggiani 2006 for a summary).  The Palomar Testbed Interferometer (PTI) was built by NASA/JPL as a testbed for developing ground and space-based interferometry and is located on Palomar Mountain near San Diego, CA \\citep{col99}. It combines starlight from two out of three available 40-cm apertures and measures the resulting interference fringes. The high angular resolution provided by this long-baseline (85-110 m), near infrared ($2.2 \\mu$m) interferometer is sufficient to resolve emission on the milli-arcsecond scale. ",
        "conclusions": "We have used long-baseline near-IR interferometry to resolve the classical nova V1663 Aql, starting $\\sim 9$ days after outburst.  We measure an apparent expansion rate of $0.21 \\pm 0.03$ ${\\rm mas day^{-1}}$, which can be combined with previously determined expansion velocities to produce a distance estimate to the nova of $8.9\\pm3.6$ kpc; the precision is limited by the precision of the available spectroscopic radial velocities. Such a large distance is consistent with the large reddening ($E(B-V) \\sim 1.1$) determined from photometry, as well as with distances found from MMRD relations. This represents only the third time a nova has been resolved using optical/IR interferometry, the previous cases being Nova V1974 Cyg 1992 \\citep{q93} and more recently RS Oph \\citep{mon06,lane07}. We anticipate that further instrumental improvements, in particular high spectral resolution interferometry such as that recently deployed on the Keck Interferometer \\citep{eis07}, may break the inclination degeneracy and thus yield very precise expansion distances to these interesting objects."
    },
    "0710/0710.0407_arXiv.txt": {
        "abstract": "The ``galactic shocks'' \\citep{fujimoto68,roberts69} is investigated using a full three-dimensional hydrodynamic simulations, taking into account self-gravity of the ISM, radiative cooling, and star formation followed by energy feedback from supernovae. This is an essential progress from the previous numerical models, in which 2-D isothermal, non-self-gravitating gas is assumed.  We find that the classic galactic shocks appears is unstable and transient, and it shifts to a globally quasi-steady, inhomogeneous pattern due to non-linear development of instabilities in the disk. The spiral patterns consists of many GMC-like dense condensations, but those local structures are not steady, and they evolves into irregular spurs in the inter-arm regions. Energy feedback from supernovae do not destroy the quasi-steady spiral arms, but it mainly contributes to vertical motion and structures of the ISM. The results and methods presented here are a starting point for more consistent treatment of the ISM in spiral galaxies, in which effects of magnetic fields, radiative transfer, chemistry, and dynamical evolution of a stellar disk are taken into account. ",
        "introduction": "A standing spiral shock solution was first discovered in the 60s numerically in a rotating gas disk in the galactic potential with a tightly wound spiral perturbation (i.e. the ``pitch angle'' is very small). Since then in the most of theoretical studies on the galactic ``shock'' \\citep[e.g.][]{fujimoto68,roberts69,wood75,johns86,lubow86}, the interstellar medium (ISM) is treated as a isothermal and homogeneous fluid with two-dimensional approximation, or it is modeled as a $N-$body system of small cloudlets \\citep[e.g.][]{tomisaka86}. Global evolution of the spiral shock was studied in the 80s \\citep{johns86} using time-dependent, two-dimensional (2D) hydrodynamic simulations show that spiral shocks are stable and long-standing for various pitch angles.  However, the steady, smoothed galactic shock does not consistently explain the complicated distribution of the ISM around spiral arms and the inter-arm substructures akin to so-called `spurs' or `feathers'\\footnote{In this paper, we refer the terms 'spurs' or 'feathers' as inter-arm gas structures, which often show quasi-periodic features associated with main spiral arms. See Paper I, \\citet{shetty06}, and \\citet{kim02} for numerical examples, and \\citet{vigne06} for observations. On the other hand, `blanches' are sub-structures bifurcated from main spiral arms.  They are smoother and longer azimuthal extent than spurs, which may be caused by resonances \\citep{chark03}.} observed in real spiral galaxies \\citep{elm80,scov01,calz05,vigne06}. Moreover, observed molecular clouds in galactic disks do not match the picture of hydrodynamic shocks in a uniform media. It was however shown by full 2D global simulations of a non-self-gravitating, isothermal gas disk that the spurs are in fact natural consequences of ``wiggle'' instability, which is caused by a purely hydrodynamic phenomenon, i.e. Kelvin-Helmholtz instability \\citep{wadakoda04} (hereafter Paper I).  This phenomenon was also found in \\citet{shetty06}. They also pointed out that the features only grow in the inner-most several kpc regions if self-gravity and magnetic fields are ignored.   More recently, three-dimensional (3D) response of the gas to the spiral potential was modeled using a local shearing box approximation in isothermal, MHD simulations taking into account self-gravity \\citep{kim06}, and found that the wiggle instability is suppressed by radial flapping motion of the shock.    These previous results suggest that hydrodynamic effects, self-gravity of the gas and magnetic fields play some important roles on the gas structures in a spiral potential. However, an important feature of the real ISM has been ignored; its inhomogeneous multi-phase structures. Apparently the ISM is not `isothermal fluid', nevertheless it is assumed in most HD and MHD simulations\\footnote{\\citet{dobbs07} recently studied gas dynamics in a spiral potential, taking into account multi-phase nature of the ISM.  However, in their non-self-gravitating SPH simulations, an energy equation with realistic cooling and heating processes is not solved, alternatively warm ($10^4$ K) and cold ($100$ K) components are treated separately as isothermal gases without phase exchange.}. Effects of self-gravity and magnetic fields highly depend on gaseous temperatures and phases. In this sense, an energy equation with a realistic cooling and heating processes should be solved before taking into account those effects.  Local box-approximation with a shearing periodic boundary \\citep[e.g.][]{kim06} is not necessarily relevant for representing dynamics of the multi-phase ISM in galactic disks, because typical scales of the inhomogeneous structures of the ISM are not small enough compared to the disk size \\citep[cf.][]{wada01, wada07, tasker06}.  Moreover, since gravity is a long-range force, global, 3D simulations for the whole disk are essential if self-gravity of the gas is considered.  This is also necessary for the multi-phase ISM, because the scale height and velocity dispersion are different for cold and hot gases.  In 3D hydrodynamic simulations of self-gravitating gas disks with the radiative cooling, we should take into account energy feedback from stars, otherwise the cold gas disk becomes unrealistically thin and the gravitational instability is strongly affected.   Here we show, for the first time, 3D evolution of the ISM in a galactic spiral potential, taking into account realistic radiative cooling and energy feedback from stars, especially from supernovae explosions.  The simulations are performed for the whole gas disks without any assumptions for symmetry.  In order to ensure a high spatial resolution (10 pc), which is essential to reproduce the multi-phase ISM, we here focus on a central part of a relatively small galaxy (radius is 2.56 kpc and the maximum circular velocity is 150 km s$^{-1}$). However, this is large enough to see the global effect of the spiral potential on the inhomogeneous ISM. Our simulations clarify apparent discrepancy between the steady solution of ``galactic shock'' and the complicated, non-steady structures of the ISM in real spiral galaxies.  This is a major progress from the previous simulations, and it will be a starting point for more realistic numerical models of the ISM, taking into account `lived' stellar potential, magnetic fields, UV radiation from massive stars, and chemistry of molecules and atoms. ",
        "conclusions": "The high resolution hydrodynamic simulations, taking into account self-gravity of the gas and realistic cooling and heating processes for the ISM, reveal for the first time that the galactic spiral arms of the ISM are neither hydrodynamic shock waves nor an assembly of long-lived, bullet-like cloudlets.  The global spiral arms are consist of complicated time-dependent substructures from which stars can be formed, but over a long time (at least 5-6 rotational periods), they stably exist under the influence of the spiral potential.  The pseudo-spiral in the multi-phase ISM is robust for energy input from the supernovae, which mainly cause the vertical non-uniform structure of cold and hot gases.  The pattern speed of the spiral potential and its strength are not key parameters to alter above features.  The ISM in galactic disks has often been represented by isothermal, non-self-gravitating fluid or inelastic particles in many astrophysical simulations.  Moreover, introducing a periodic boundary condition and reducing spatial dimensions were usual. However, those kinds of simplification do not necessarily represent the nature of spiral galaxies.  Magnetic fields are not considered in the present simulations, but the present results suggest that even if the magnetic fields are weak, the spiral patterns can steadily exist on a global scale. Effect of magnetic fields and other important physical processes, such as UV radiation and chemistry, should be investigated based on 3D global models, as presented in this paper.   The present treatment of the multi-phase ISM could be used for direct comparison with observations, coupled with radiative transfer calculations for various observational probes \\citep[e.g.][]{wada00,wada05}. Fine structures of molecular gas associated with the spiral arms and in the inter-arm regions in external galaxies, which could be compared with the present numerical model, will be revealed by high resolution observations, e.g. by the Atacama Large Millimeter/Submillimeter Array."
    },
    "0710/0710.0588_arXiv.txt": {
        "abstract": "Outflows from active galactic nuclei (AGNs) seem to be common and are thought to be important from a variety of perspectives: as an agent of chemical enhancement of the interstellar and intergalactic media, as an agent of angular momentum removal from the accreting central engine, and as an agent limiting star formation in starbursting systems by blowing out gas and dust from the host galaxy.  To understand these processes, we must determine what fraction of AGNs feature outflows and understand what forms they take.  We examine recent surveys of quasar absorption lines, reviewing the best means to determine if systems are intrinsic and result from outflowing material, and the limitations of approaches taken to date. The surveys reveal that, while the fraction of specific forms of outflows depends on AGN properties, the overall fraction displaying outflows is fairly constant, approximately 60$\\%$, over many orders of magnitude in luminosity.  We emphasize some issues concerning classification of outflows driven by data type rather than necessarily the physical nature of outflows, and illustrate how understanding outflows probably requires more a comprehensive approach than has usually been taken in the past. ",
        "introduction": "The role of outflows from quasars and active galactic nuclei (AGN) has recently become an important feature in the overall framework of how galaxies and star formation processes evolve over cosmic time. Mergers and other interactions triggering AGN seem to provide feedback affecting the larger scale environment. Recent efforts to include the effects of this so-called AGN feedback focus on two modes: a ``radio'' mode whereby a relativistic jet heats the surrounding interstellar and intercluster media \\citep[e.g.,][]{best07}, and a ``quasar'' mode whereby a lower velocity but higher mass outflow also helps to clear out post-merger shrouding gas and quenches star formation \\citep*[e.g.,][]{tdm05}. We focus on this second mode in this paper. For this mode, a number of questions require addressing. How common are outflows? Do all AGN have outflows? What drives outflows? Is there a single all-governing structure of AGN? Answering these questions will help us to understand the role AGN outflows with respect to issue of feedback, and other important issues like chemical enrichment and accretion.  In the ensuing sections we aim to achieve several goals: (1) to review the ways in which outflows are detected in AGN over all luminosity scales; (2) to comment on the merits of various catalogs of outflows; and (3) to arrive at the true (possibly property-dependent) observed frequency of outflows. In its most basic interpretation, the observed frequency of outflows can be equated with the fraction of solid angle (from the view point of the central black hole) subtended by outflowing gas. This interpretation assumes that all AGN feature outflows and that not all sight-lines to the emitting regions are occulted by the outflow. Alternatively (and equally simplistic), the frequency can be interpreted as the fraction of the duty cycle over which AGN feature outflows (assuming the outflow subtends 4$\\pi$\\ steradians). The actual conversion of the fraction of AGN featuring spectroscopic evidence of outflow to the solid angle subtend by such outflows has been treated by \\citet{cren99} and \\citet*{cren03b}. This computation involves further knowledge of the line-of-sight covering factor (that is, the fraction of lines-of-sight that reach the observer that are occulted by the outflow) as well as an understanding of range of solid angle sampled by the AGN used (e.g., Type 1 versus Type 2 AGN). The true situation is likely in between these two extremes, and may depend also on properties we can not currently constrain, such as the time since the AGN was triggered.  Additionally, we strive here to build a case that more effort should be made to consider outflows of all types together.  Often data limitations of one sort or another have led to the study of limited parts of parameter space (e.g., outflow velocity or velocity dispersion), creating artificial or at least biased divisions.  There appears to be a continuous range in properties of outflows and these should only be regarded as fundamentally different when there is clear evidence to reach such a conclusion.  Below we discuss the identification of outflows (\\S 2) and the data-driven subcategories (\\S 3).  We show an illustrative example of how combining the different outflow subclasses may lead to a more unified physical understanding of outflows (\\S 4).  Finally, we bring together the different survey methodologies to determine an overall fraction of AGN displaying the signatures of outflows (\\S 5) and summarize the case for more global studies of the outflow phenomenon.  We adopt a cosmology with $\\Omega_M$ = 0.3, $\\Omega_{\\Lambda}$ = 0.7, and H$_0$ = 70 km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": ""
    },
    "0710/0710.0893_arXiv.txt": {
        "abstract": "We use two-band imaging data from the Advanced Camera for Surveys on board the Hubble Space Telescope for a detailed study of NGC\\,1533, an SB0 galaxy in the Dorado group surrounded by a ring of \\HI.  NGC\\,1533 appears to be completing a transition from late to early type: it is red, but not quite dead.  Faint spiral structure becomes visible following galaxy subtraction, and luminous blue stars can be seen in isolated areas of the disk.  Dust is visible in the color map in the region around the bar, and there is a linear color gradient throughout the disk.  We determine an accurate distance from the surface brightness fluctuations (SBF) method, finding $\\mM=31.44\\pm0.12$ mag, or $d = 19.4\\pm1.1$ Mpc. We then study the globular cluster (GC) colors, sizes, and luminosity function (GCLF).  Estimates of the distance from the median of the GC half-light radii and from the peak of the GCLF both agree well with the SBF distance.  The GC specific frequency is $S_N=1.3\\pm0.2$, typical for an early-type disk galaxy. The color distribution is bimodal, as commonly observed for bright galaxies. There is a suggestion of the redder GCs having smaller sizes, but the trend is not significant.  The sizes do increase significantly with galactocentric radius, in a manner more similar to the Milky Way GC system than to those in Virgo.  This difference may be an effect of the steeper density gradients in loose groups as compared to galaxy clusters. Additional studies of early-type galaxies in low density regions can help determine if this is indeed a general environmental trend. ",
        "introduction": "The Hubble Space Telescope (\\hst) has opened the door to our understanding of extragalactic star cluster systems, revealing numerous globular clusters (GCs) in early-type galaxies (e.g.\\ Gebhardt \\&\\ Kissler-Patig, 1999; Peng \\etal\\ 2006) as well as ``super-star clusters'' (SSCs) in late-type galaxies (Larsen \\&\\ Richtler 2000), especially starbursts (e.g. Meurer \\etal\\ 1995, Maoz \\etal\\ 1996).  In early-type systems the color distribution of the GCs is often bimodal, consisting of a blue metal-poor component and a red metal-rich component (e.g. West \\etal\\ 2004; Peng \\etal\\ 2006).  This observation gives a hint to the connection between the early- and late-type systems.  Mergers often have particularly strong starbursts and rich populations of SSCs, as seen for example in NGC4038/39 - ``the Antennae'' system (Whitmore \\&\\ Schweizer 1995; Whitmore \\etal\\ 1999), and hierarchical merging is one possible origin for the redder population of GCs in early-type galaxies (e.g., Ashman \\& Zepf 1998; Beasley \\etal\\ 2002; Kravtsov \\& Gnedin 2005).  Much of the research on GC systems has concentrated on galaxy clusters which are rich in early-type galaxies (e.g., the ACS Virgo Cluster Survey, C{\\^o}t{\\'e} \\etal\\ 2004; the ACS Fornax Cluster Survey, Jord{\\'a}n \\etal\\ 2007).  Early-type galaxies in groups and the field are somewhat less studied, particularly with \\hst\\ and the Wide Field Channel (WFC) of its Advanced Camera for Surveys (ACS).  The relatively wide ($3\\farcm4$) field of view of the ACS WFC combined with its fine pixel sampling make it an exceptional tool for imaging GCs out to a few tens of Mpc where they have measurable angular sizes (Jord{\\'a}n \\etal\\ 2005).  Here we report \\hst\\ ACS/WFC imaging of the SB0 (barred lenticular) galaxy NGC~1533 in the Dorado group.  This group is in the ``Fornax wall'' (Kilborn \\etal\\ 2005) and hence at a similar distance to the Fornax  cluster (e.g., Tonry \\etal\\ 2001).  Dorado is interesting in that it is richer than the Local Group but still dominated by disk galaxies (its brightest members being the spiral NGC~1566 and the S0 NGC~1553), and its members have \\HI\\ masses similar to non-interacting galaxies with the same morphology (Kilborn \\etal\\ 2005).  While the apparent crossing time of the group is only $\\sim 13$\\%\\ of the age of the universe (Firth \\etal\\ 2006; see also Ferguson \\&\\ Sandage 1990), the most recent analyses conclude the group is unvirialized (Kilborn \\etal\\ 2005; Firth \\etal\\ 2006), which may explain the richness in spirals and \\HI.  \\begin{figure*} \\plottwo{f1a.eps}{f1b.eps} \\caption{\\textit{Left panel:} \\HI\\ contours from Ryan-Weber \\etal\\ (2003) are overlaid on a ground-based $R$-band image of the NGC\\,1533 field from the SINGG survey (Meurer \\etal\\ 2006). The outermost contour (bold) is at a column density of $10^{20}$ cm$^{-2}$, and the contours increase in steps of $0.5{\\times}10^{20}$ cm$^{-2}$.  The small companion galaxies IC~2038/2039 are in the upper right corner of the image.  NGC\\,1533 itself is in an \\HI\\ ``hole'' (the galaxy center is not detected), and this distribution has been described as a ring. \\textit{Right panel:} ACS HRC (green) and WFC (blue) fields of view for the two HST roll angles described in the text (labeled 1 and 2). The outlines of the camera fields are overlaid on a $\\sim\\,$9\\arcmin\\ portion of the SINGG $R$-band image. North is up and East is to the left in both panels. } \\label{fig:ACSfields} \\end{figure*}  NGC~1533 is the seventh brightest member of the Dorado group, with $M_V{\\,\\approx\\,}{-}20.7$. It lies within the virial radius, but is a $\\sim\\,$2-$\\sigma$ velocity outlier (Kilborn \\etal, 2005; Firth \\etal\\ 2006) so that it is moving at high speed through the intra-group medium.   A vast \\HI\\ arc is seen in the outskirts of NGC~1533 connected to the Sdm galaxy IC~2038 and the small S0 galaxy IC~2039  (Ryan-Weber \\etal\\ 2004). This suggests that NGC~1533 is ``stealing'' ISM from its companions or has cannibalized another gas-rich satellite (see Figure~\\ref{fig:ACSfields}, left panel). As is typically seen in S0 galaxies, star formation is weak in NGC~1533.  Observations of this galaxy in spectroscopic surveys note the presence of emission lines (Jorgensen \\etal\\ 1997; Bernardi \\etal\\ 2002); the nuclear spectrum available from the 6dF survey (Jones \\etal\\ 2005) shows [{\\sc N ii}]6584 and weak \\Halpha.  \\Halpha\\ imaging from the Survey of Ionization in Neutral Gas Galaxies (SINGG, an \\Halpha\\ imaging survey of \\HI\\ selected galaxies; Meurer \\etal\\ 2006) shows a few weak \\HII\\ regions beyond the end of its bar (the nucleus is too bright to allow faint nuclear \\HII\\ regions to be detected in the SINGG images), as well as a scattering of very faint ``intergalactic \\HII\\ regions''.  These are discussed in more detail by Ryan-Weber \\etal\\ (2004) who show that they are so faint that it would only take one to a few O stars to ionize each one. Although its current rate is low, the star formation in NGC~1533 illustrates another possible channel for building up cluster systems in early-type galaxies: slow re-ignited star formation in ISM stripped from companions.  The ACS WFC images of NGC~1533 used in the present study were obtained with \\hst\\ as ``internal parallel images'' while the ACS High Resolution Channel (HRC) was pointed at the intergalactic \\HII\\ regions (\\hst\\ GO Program 10438; M. Putman, PI).  The HRC observations are discussed elsewhere (Werk \\etal\\ 2007, in preparation). Here we use the WFC observations to measure the structural properties of the galaxy, characterize its GC population, and use the GC luminosity function (GCLF), GC half-light radii, and surface brightness fluctuations (SBF) to provide accurate distance estimates.  The contrast between NGC~1533, a (weakly) star-forming gas-rich barred S0 in a loose group environment, and galaxies in the richer environments of the Virgo and Fornax cluster, provides a useful test of the ubiquity of the various relations found in the denser environments.  The following section describes the observations and data reductions in more detail. Sec.~\\ref{sec:props} discusses  the galaxy morphology, structure, color profile, and isophotal parameters.  Sec.~\\ref{sec:sbf} presents the SBF analysis and galaxy distance, while Sec.~\\ref{sec:gccolors}--\\ref{sec:gclf} discuss the GC colors, effective radii, luminosity function, and specific frequency. The final section summarizes our conclusions.    \\begin{figure*} \\epsscale{1.0} \\plotone{f2.eps} \\caption{\\textit{Upper left:} Combined F814W ACS/WFC image of NGC~1533. \\textit{Lower left:} Contour map of a 3\\farcm6$\\times$4\\farcm0 portion of the image.  Contours are plotted in steps of a factor of two in intensity, with the faintest being at $\\mu_I = 20.7$ mag~arcsec$^{-2}$. \\textit{Upper right:} The image following galaxy model subtraction, showing the faint spiral structure (the ``plume'' 2\\arcmin\\ north of the galaxy is a ghost image).  The same 3\\farcm6$\\times$4\\farcm0 field is shown; the box marks the central~1\\arcmin. \\textit{Lower right:} Colormap of the central 1\\arcmin\\ region of NGC1533, with dark indicating red  areas and white indicating blue. The dark spot at center marks the center of the galaxy.  Dust can be seen as faint, dark, wispy features.  A compact blue star-forming region is visible to the left of the galaxy center, near the center-left of the map. } \\label{fig:cmb4} \\end{figure*} ",
        "conclusions": "We have analyzed deep F606W and F814W images of the galaxy NGC\\,1533 and its GC population taken at two roll angles with the ACS/WFC on \\hst.  Although it is classified as an early-type barred lenticular galaxy, we found faint spiral structure once a smooth fit to the galaxy isophotes was subtracted.  The color map shows faint dust features in the area around the bar and inner disk. Previous ground-based \\Halpha\\ imaging had shown that the galaxy disk contains several faint, compact \\HII\\ regions.  We find that all of these regions have some blue stars associated with them. Four of the \\HII\\ regions lie within one of the faint spiral arms, and several other blue stars are spread out within the arm. These observations suggest that NGC\\,1533 is in the late stages of a transition in morphology from type SBa to~SB0.  Transition objects such as NGC\\,1533 may be the key to understanding the evolution of the morphology-density relation in galaxy clusters, which is often explained as infalling spirals being transformed into S0s by the harsh cluster environment (e.g., Dressler \\etal\\ 1997; Postman \\etal\\ 2005).  However, recent evidence at intermediate redshift indicates that the transitions begin outside the clusters in small group environments through galaxy-galaxy interactions (Moran \\etal\\ 2007).  Following infall, the intra-cluster medium then serves to expedite the process by removing the remaining cool gas.  With NGC\\,1533, we have a close-up view of this transition in the Dorado group, including interaction with the neighboring galaxies IC~2038/2039 (Ryan-Weber \\etal\\ 2004).   From two-dimensional two-component parametric modeling of the galaxy surface brightness, we find a bulge-to-total ratio $B/T\\approx0.42$.  The half-light radii of the bulge and disk are $\\sim\\,$7\\arcsec\\ and $\\sim\\,$46\\arcsec, respectively.  We find a best-fitting S\\'ersic index $n=2.0$ for the bulge, which can be reasonably approximated by an $r^{1/4}$ law in the 1-D profile. However, the disk has a relatively flat profile over a factor-of-three in radius, from $\\sim\\,$15\\arcsec\\ to $\\sim\\,$45\\arcsec, then steepens fairly abruptly beyond $\\sim\\,$50\\arcsec. This gives the disk a very low S\\'ersic index of $n\\approx0.4$, which might result from past high-speed interactions of NGC\\,1533 within the group environment.  Overall, the color of NGC\\,1533 is that of an evolved, red population, except in the few, small isolated regions where the blue stars occur.  The bulge color is $\\vi\\gta1.22$, similar to cluster ellipticals, and then there is a mild, but significant, linear color gradient throughout the disk.  There is a gradual isophotal twist and the isophotes increase in ellipticity out to a semi-major axis distance of 24\\arcsec, where $\\epsilon$ goes above 0.4 before falling sharply again towards the round outer disk.  The peak of the $A_4$ harmonic term, measuring ``diskiness,'' occurs at a smaller semi-major axis of 21\\arcsec. This is because the pointed lens-like isophotes occur inside of the bar.  We measured the SBF amplitude in four broad radial annuli for each of the two observations at different roll angles.  A gradient in the SBF amplitude is clearly detected and follows the color gradient (the bluer outer regions have relatively brighter SBF).  By matching our ACS photometry against ground-based \\vi\\ data for this galaxy, we have accurately calibrated the SBF measurements to obtain distance moduli.  We find excellent agreement among the different annuli but with an offset of 0.04 mag in distance between the two observations. However, the distance error is dominated by systematic uncertainty in the color and calibration zero point.  We find a final distance modulus $(m{-}M) = 31.44 \\pm 0.12$ mag, or $d = 19.4 \\pm 1.1$ Mpc.  Candidate globular clusters were selected according to color, magnitude, radial position, and FWHM.  Analysis of the color distribution of these objects with the KMM algorithm indicates with a very high degree of confidence that the distribution is bimodal with peaks at $\\vi\\approx0.92$ and 1.22. There is no evidence that the blue GCs become redder at bright magnitudes, the so-called ``blue tilt.''  The absence of this effect in NGC\\,1533, an intermediate luminosity galaxy with a small GC population, is consistent with a self-enrichment explanation, since the GCs in such systems do not reach the high masses that they do in richer systems.  The sizes of the GC candidates were measured using the Ishape software. By comparing the results from the two different roll angles, we found that the effective (half-light) radii \\reff\\ have an accuracy of about 13\\% down to $\\iacs=23$, but are not reliable beyond this.  We did not find a significant trend of \\reff\\ with GC color, although the red-peak GCs have a median \\reff\\ smaller by $11\\pm8$\\% than the blue-peak GCs.  However, we did find a significant (4\\,$\\sigma$) trend of \\reff\\ with galactocentric radius.  In this respect, NGC\\,1533 is more like the Milky Way than the Virgo early-type galaxies.  This may be an effect of the environment: since the sizes of the GCs are limited by the tidal field, and the density gradients will be steeper in small groups such as Dorado or the Local Group, GC sizes should have a stronger dependence on radius in such environments.  The dominance of this radial effect may weaken or obscure any relation between size and color.  More studies of size and color trends for the GCs of galaxies in loose groups are needed to verify this hypothesis, although this may be difficult because of the low GC populations in such systems.  We then used the median half-light GC radius to obtain another estimate of the distance to NGC\\,1533.  Following Jord\\'an \\etal\\ (2005), we find $d = 18.6\\pm2.0$ Mpc, in good agreement with the SBF distance.  We modeled the $I_{814}$-band GCLF of NGC\\,1533 as a Gaussian using a maximum likelihood fitting routine. The best-fit peak magnitude $m^0_I=22.84^{+0.18}_{-0.24}$ corresponds to $M^0_I \\approx -8.6$ for the measured SBF distance, in good agreement with expectations based on other galaxies.  The fitted Gaussian dispersion of $\\sigma_{\\rm LF} = 1.10 \\pm 0.15$ mag is in accord with the relation between $\\sigma_{\\rm LF}$ and galaxy luminosity found recently by Jord\\'an \\etal\\ (2006) for Virgo galaxies. Finally we estimate the GC specific frequency in the analysis region to be $S_N = 1.3\\pm0.2$, typical for a disk galaxy. We conclude that the GCs in NGC\\,1533 have the same average size, color, and luminosity within the errors as the Virgo early-type galaxies, but the stronger dependence of size on galactocentric distance is more reminiscent of the Milky Way.  NGC\\,1533 represents an interesting class of transitional objects, both in terms of morphology and environment.  A large, multi-band, systematic study of such systems with \\hst, similar to the ACS Virgo and Fornax Cluster surveys but focusing on group galaxies, has yet to be undertaken and must await either a revived ACS or Wide Field Camera~3.  Such an effort would be extremely valuable in piecing together a more complete picture of the interplay between galaxy structure, globular cluster system properties, and environment."
    },
    "0710/0710.1547_arXiv.txt": {
        "abstract": "We present our new advanced model for population synthesis of close-by cooling NSs. Detailed treatment of the initial spatial distribution of NS progenitors and a detailed ISM structure up to 3 kpc give us an opportunity to discuss the strategy to look for new isolated cooling NSs. Our main results in this respect are the following: new candidates are expected to be identified behind the Gould Belt, in directions to rich OB associations, in particular in the  Cygnus-Cepheus region; new candidates, on average, are expected to be hotter than the known population of cooling NS. Besides the usual approach (looking for soft X-ray sources), the search in 'empty' $\\gamma$-ray error boxes or among run-away OB stars may yield new X-ray thermally emitting  NS candidates. ",
        "introduction": "More than 10 years after the discovery of its brightest member RX J1856-3754 \\cite{Walter1996}, a group of seven radio-quiet isolated neutron stars (NSs) detected by ROSAT gained an important place in the rich zoo of compact objects. Together with Geminga and several close-by young radio pulsars, these objects form the local population of cooling NSs. Studies of this group of sources already provided a wealth of information on NSs physics (see e.g. \\cite{Haberl2007,Page2007,Zane2007} for recent reviews).  Since 2001 the number of known close-by radio-quiet NSs has not been growing despite all attempts to identify new candidates. Partly this is due to the fact that all these searches are blind. To advance the identification of new near-by cooling NSs it is necessary to perform a realistic modeling of this population.  To investigate the population of close-by young cooling NSs the method of population synthesis is used here. In this short note  an advanced population synthesis model is briefly discussed for the population of close-by ($< 3$~kpc) isolated NSs which can be observed via their thermal emission in soft X-rays (the detailed description and full analysis of new results will be presented elsewhere \\cite{Posselt2008}). Previously our models were applied to confirm the link between the seven radio quiet NSs (the Magnificent Seven) and the Gould Belt \\citep{p03} (Paper I), to study distribution of NSs in the Galaxy and in the solar vicinity \\cite{p04} (Paper II), and to test theories of thermal evolution of NSs \\citep{pgtb04} (Paper III). The major interest of the present study is to get a hint how to find more objects of this type. ",
        "conclusions": "The main aim of this study is to make some advances in the strategy for searching for new isolated cooling NSs. According to our results, new candidates expected to be identified at ROSAT count rates $<0.1$~cts~s$^{-1}$ should be young objects born in rich OB associations behind the Gould Belt. Most of the recent studies \\cite{Agueros2006,Chieregato2005,Rutledge2003} looked for new candidates far from the galactic plane. It seems that this is not very promising. Our results indicate that new cooling NSs should be searched in directions of OB associations such as Cyg OB7 and Cep OB3.  Considering sky coverage the ROSAT All Sky Survey is currently the best choice to look for new ``cowboys'' in the Cygnus-Cepheus region which is, according to our results, the most promising area. However, the relatively large positional error circle of ROSAT usually includes many possible optical counterparts, especially at these low galactic latitudes. Furthermore one has to exclude variable X-ray sources to find isolated cooling NSs. In this respect the recently published XMM-$Newton$ Slew Survey may become an important database. A major step can be expected from the planned eROSITA all sky survey which - compared with the ROSAT all sky survey --  will provide a factor of $\\sim$~10 in soft X-ray sensitivity and factor of $\\sim$~4 in energy resolution \\cite{Predehl2006}.   Some of unidentified $\\gamma$-ray sources (already observed by EGRET and forthcoming due to AGILE and GLAST) can be identified as cooling NSs as it was with Geminga and 3EG J1835+5918. In particular, GLAST observations of the 56 EGRET error boxes studied in \\cite{Crawford2006} and later cross-correlation with the ROSAT (or/and XMM) data can result in new identification of cooling NSs.  Another possibility to find new isolated coolers is to search for (un)bound compact companions of OB runaway stars. More than one hundred OB runaway stars are known in a 1 kpc region around the Sun \\citep{Zeeuw1999}. They are characterized by large spatial velocities or/and by large shifts from the galactic plane. Two main origins of these large velocities are currently discussed: dynamical interaction and explosion of a companion in a close binary system. The latter case is interesting for the discussion of search for new close-by cooling NSs.  A binary can survive after the first SN explosion in, roughly, 10-20\\% of cases. Then one expects to have a runaway system consisting of an OB star and a compact object (most probably a NS). A young NS can appear as a radio pulsar. In \\cite{sayer1996} and \\cite{philp1996} the authors searched for radio pulsar companions of $\\sim 40$ runaway OB stars. Nothing was found. This result is consistent with the assumption that in less than 20\\% cases OB stars have radio pulsar companions. Still, it is interesting to speculate that runaway massive stars can have cooling radio quiet NS as companions. Then a companion can be identified as a source of additional X-ray emission.   \\begin{theacknowledgments} We thank D. Blaschke, H. Grigorian, and D. Voskresensky for data on cooling curves and discussions; A. Mel'nik for discussion of properties of OB associations; R. Lallement for the sodium data; and A. Pires for discussions about the ISM model. S.B.P. was supported by INTAS and Dynasty foundations. \\end{theacknowledgments}"
    },
    "0710/0710.1637_arXiv.txt": {
        "abstract": "A simple speed-up cosmology model is proposed to account for the dark energy puzzle. We condense contributions from dark energy and curvature term into one effective parameter in order to reduce parameter degeneracies and to find any deviation from flat concordance $\\Lambda$CDM model, by considering that the discrimination between dynamical and non-dynamical sources of cosmic acceleration as the best starting point for analyzing dark energy data sets both at present and in future. We also combine recent Type Ia Supernova (SNIa), Cosmic Microwave Background (CMB) and Baryon Oscillation (BAO) to constrain model parameter space. Degeneracies between model parameters are discussed by using both degeneracy diagram and data analysis including high redshift information from Gamma Ray Bursts (GRBs) sample. The analysis results show that our model is consistent with cosmological observations. We try to distinct the curvature effects from the specially scaling dark energy component as parameterized. We study the linear growth of large scale structure, and finally show the effective dark energy equation of state in our model and how the matter component coincidences with the dark energy numerically. ",
        "introduction": "It is now well-established that the expansion of our universe is currently in an accelerating phase, supported by the most direct and robust evidence from the redshift - apparent magnitude measurements of the \"cosmic lighthouse\" type Ia supernova \\cite{Perlmutter}, and indirect others such as the observations of Cosmic Microwave Background (CMB) by the WMAP satellite \\cite{Spergel,jarosik,hinshaw,page,boomerang,cbi,kuo,sa,planck}, and large-scale galaxy surveys by 2dF and SDSS \\cite{cole,sdss1,sdss2,tegmark}. Under the assumption that general relativity is valid on cosmological scale, the combined analysis of different observation data sets indicates a spatially-flat universe with about 70\\% of the total energy content of the universe today as so called dark energy with effectively negative pressure responsible for the accelerating expansion (see Ref.~\\cite{Peebles} for reviews on this topic). Among multitudinous candidates of dark energy models, the \"simplest\" and theoretically attractive one might be the so called vacuum energy, i.e. $\\rho_{\\Lambda} = \\Lambda/8 \\pi G$, where $\\Lambda$ is the cosmological constant, which has been long considered as a leading candidate and works quite well on explaining observations through out the history of our universe at different scales. But the origin or mechanisms responsible for the cosmic accelerating expansion are not very clear. On the other hand, some authors suggest that maybe there does not exist such mysterious dark component, but instead the observed cosmic acceleration is a signal of our first real lack of understanding of gravitational physics \\cite{Lue} on cosmic scale. An example is the braneworld theory with the extra dimensions compactified or non-compactified \\cite{Dvali,rs,hw,ms2}. Consequently, finding the different cosmological implications to distinguish modified gravity models and dark energy scenario from observations is essentially fundamental to physically understanding of our universe\\cite{distinguish}.  Along with the matter (mainly cold dark matter) component and possible curvature term, the mysterious dark energy dominates the fate of our universe (we do not consider the radiation component contribution as it is supposed to be very tiny for the current universe evolution, at least for the present discussion interests). Ironically so far we do not know much to either of them, even full of puzzling to some extends. So any progress or reasonable understanding to each of them is undoubtedly valuable. Specifically, the quest to distinguish between dark energy and modified gravity scenario and further to differentiate cosmological constant and dynamical dark energy models from observations has become the focus of cosmology study since it holds the key to new fundamental physics.   Although we have built up a successful parametrization to describe the properties and evolution of our universe, and in principle distinct dark energy models live at different sub-space of fully descriptive multi-dimension parameter space, due to serious degeneracies among different parameters, we cannot get tight enough constraints from observations by global fitting various observational data sets. One way to extract useful information from observation data and get hints for fundamental physics from cosmology study is to reduce the dimension of parameter space (thus reduce the parameter degeneracies) with particular purpose in mind without apparently biased input to model parametrization. Since we have not found any evidence of inconsistency of standard $\\Lambda$CDM model, including more parameters which describe detailed properties of each component if the \"cosmic pie\" will complicate the situation to constrain model parameters.  In order to find deviation of dark energy equation of state parameter $w$ from $-1$ (e.g. evidence of non-cosmological constant dark energy), the assumption of a flat universe is widely accepted in the literature with claims that curvature is negligible from inflation predictions and with emphasis on combined analysis results \\emph{with} prior assumption $w=-1$. On the other hand, typically one looks for evidence of dynamical dark energy in the absence of spatial curvature to get better constraints (for an exception, see \\cite{curvature1}). It has been concluded in \\cite{curvature} that the non-curvature assumption can induce critically large errors in reconstructing the dark energy equation of state even if the true cosmic curvature is on sub-percent level. These claims motivate us proposing a parameterized dark component term to mimic the effective contributions from either dark energy or curvature term plus the dark energy (It is also possible that the parameterized term we postulate may be from a fundamental theory or reasonably modified gravity model we are seeking), besides the conventional matter term.   In the first step, it is reasonable to introduce only one parameter which stands for any kind of deviation from standard cosmology model. In some limit case, it should be reduced to the simple four dimensional (4D) $\\Lambda$CDM cosmology. The constraint on this parameter from observations should provide insightful hints to further explore fundamental physics.   In the next section, we propose a simple cosmic parametrization for the current universe, a parameterized model for the later evolution of our universe. In section 3 we give various cosmic probes to this model, with comparison to the DGP model Universe\\cite{Dvali} and the concordance model with a cosmological constant, i.e, the $\\Lambda$CDM model, with the hope to locate new features to this new model. Then in section 4, we discuss the new degeneracies between the parameters we introduced and dark matter content. The possible constraints from high redshift observations are also discussed. The last section devotes discussions and conclusions for the general framework studies to this present model. ",
        "conclusions": "We have presented a cosmic model parameterizing the late universe which collapses curvature and dark energy effects into one parameter $B$ that may indicate any deviation from standard flat $\\Lambda$CDM model and we find that we can not conclude that the cosmic curvature term is constantly zero, instead it may contribute rich phenomenological effects. In order to show the advantages of our parametrization, we study the degeneracy properties between $B$ and $\\Omega_m$, emphasizing the contribution from high redshift distance information from GRB or gravitational waves experiments on-going and up-coming. It is well-known that deducing the number of free parameter without significant physics lost is quite important to constrain cosmology models and to find new physics behind.  In this paper we also investigated the DGP cosmological model in the simplest flat geometry case with the extra dimension contribution as an effective \"cosmological constant\", compared with our parameterized model and the reduction to the power-law $\\Lambda$CDM model for the 4D real Universe. We find that the DGP model even in the simplest case is still an interesting candidate for the current cosmic speed-up expansion mechanism at long distances, while we know that in the short ranges the model behaves as 4D conventional gravity. We will exploit the non-compact extra dimension to see its possible existence signatures via cosmic effects in the general DGP model later as a promising model, while we do not intent to discuss the quantum aspects of this model as a basic theory\\cite{alu}.   As a generalization of the $\\Lambda$CDM model with naive cosmological constant as dark energy candidate we has parameterized a curvature like term with new phenomenological features via numerical fittings and show the term explicit physics meanings when we perform the parameter B reduction directly to zero or 2. It may be interesting also to study the general properties of the parameterized term as the matter-energy contents in our Universe continuous equation to see what kind of \"matter\" it may describe effectively, without specifying the form of the parameter. Besides, the phantom case can be realized too, for example, the equation of state parameter $w=p/\\rho<-1$ if we take $B>2$ and quintessence corresponds to $B<2$ with $w=p/\\rho>-1$ numerically. We think this picture is in conformity with other popular models and enlarges phenomenological dark energy study possibilities to explain the late-time accelerating expansion of our Universe, thus it is worth of further endeavors."
    },
    "0710/0710.5128_arXiv.txt": {
        "abstract": "{} {A number of recent works have suggested that the period-luminosity (PL) relation for the Large Magellanic Cloud (LMC) Cepheids exhibits a controversial nonlinear feature with a break period at 10 days. Therefore, the aim of this Research Note is to test the linearity/nonlinearity of the PL relations for the LMC Cepheids in $BVI_cJHK_s$ band, as well as in the Wesenheit functions.} {We show that simply comparing the long and short period slopes, together with their associated standard deviations, leads to a strictly larger error rate than applying rigorous statistical tests such as the $F$-test. We applied various statistical tests to the current published LMC Cepheid data. These statistical tests include the $F$-test, the testimator test, and the Schwarz information criterion (SIC) method.} {The results from these statistical tests strongly suggest that the LMC PL relation is nonlinear in $BVI_cJH$ band but linear in the $K_s$ band and in the Wesenheit functions. Using the properties of period-color relations at maximum light and multi-phase relations, we believe that the nonlinear PL relation is not caused by extinction errors. } {} ",
        "introduction": "Recently, Fouqu\\'{e} et al. (\\cite{fou07}) have derived the Galactic Cepheid period-luminosity (PL) relation with several different techniques, including parallax measurements (from {\\it Hipparcos} and {\\it HST}), variants of the Baade-Wesselink method, and distances inferred from open clusters. We point out that such an approach has been applied before in Ngeow \\& Kanbur (\\cite{nge04}) and Groenewegen et al. (\\cite{gre04}). In addition, Fouqu\\'{e} et al. (\\cite{fou07}) also derive Large Magellanic Cloud (LMC) PL relations in the $BVR_cI_cJHK_s$ band, and refer to the work of Sandage et al. (\\cite{san04}), which suggests a possible change of slope for the LMC PL relation at 10 days. In fact, there are several other papers on the topic of nonlinear\\footnote{By nonlinearity we mean that the PL relation can be broken into two relations, with a break period adopted at 10 days.} LMC PL relations (see Kanbur \\& Ngeow \\cite{kan04,kan06}; Kanbur et al. \\cite{kan07}; Ngeow et al. \\cite{nge05}; Ngeow \\& Kanbur \\cite{nge06a,nge06b}; Koen et al. \\cite{koe07}).  These previous works concentrate on the $VI_c$ band (Kanbur \\& Ngeow \\cite{kan04,kan06}; Ngeow \\& Kanbur \\cite{nge06b}) or $V$ band only (Ngeow \\& Kanbur \\cite{nge06a}; Kanbur et al. \\cite{kan07}), with data mostly from the OGLE (Optical Gravitational Lensing Survey, Udalski et al. \\cite{uda99}) database. For the $JHK_s$ band PL relations, Ngeow et al. (\\cite{nge05}) investigated possible nonlinearities using the 2MASS data from Nikolaev et al. (\\cite{nik04}) that cross-correlated with the MACHO LMC Cepheids, and a random-phase correction to derive the mean magnitudes of these 2MASS data.  Our motivation for this Research Note is to extend the previous work in $BVI_cJHK_s$ band, using the LMC Cepheid data from Fouqu\\'{e} et al. (\\cite{fou07}), with various rigorous statistical tests. The $JHK_s$ band data used in Fouqu\\'{e} et al. (\\cite{fou07}) are the 2MASS data matched to the OGLE Cepheids, and the mean magnitudes are derived using the method presented in Soszy\\'{n}ski et al. (\\cite{sos05}), which is different from the data used in Ngeow et al. (\\cite{nge05}). It is important to test the nonlinearity results in $JHK_s$ band results with different Cepheid samples and different methods deriving the $JHK_s$ mean magnitudes.  As emphasized in Ngeow \\& Kanbur (\\cite{nge06a}), statistical tests are needed to test and detect the existence of the nonlinear PL relation. We also point out that in searching for nonlinearity or a change of slope at 10 days, the method of comparing the short and long period slope with their associated standard deviations is more prone to error than applying a statistical test, such as the $F$-test as indicated by the following, purely analytical example. A statement such as the ``the slope is $x\\pm \\delta x$'' means that the probability that the slope is in the interval $(x-\\delta x,x+\\delta x)$ is $ 1 - \\alpha$, where $\\alpha$ is the desired significance level. Then if $A$ is the event that the calculated short period slope is wrong and $B$ is the event that the calculated long period slope is wrong, we have $P(A)=\\alpha$ and $P(B) = \\alpha$. Then in comparing the short and long period slopes using just their calculated standard deviations, the probability of at least one mistake is $P(A\\cup B) = P(A) + P(B) - P(A\\cap B) = 2{\\alpha} - {\\alpha}^2$. If $1 > {\\alpha} > 0$, then $2{\\alpha} - {\\alpha}^2 > {\\alpha}$. If the $F$-test or any other statistical test is carried out to the level of significance ${\\alpha}$, then this states that the probability of making an error in just comparing long and short period slopes through their standard deviations is greater than the probability that the $F$-test makes a mistake. In essence the $F$-test compares both short and long period slopes (as well as the zero-points) of the nonlinear PL relations simultaneously.   \\begin{table*} \\caption{$F$-test Results of the LMC PL Relations.} \\label{tab1} \\begin{center} \\begin{tabular}{lcccccccccc} \\hline\\hline Band & $a_S$ & $b_S$ & $\\sigma_S$ & $N_S$ & $a_L$ & $b_L$ & $\\sigma_L$ & $N_L$ & $F$ & $p(F)$ \\\\ \\hline $B$  & $-2.628\\pm0.072$ & $17.493\\pm0.046$ & 0.262 & 618 &  $-2.402\\pm0.192$ & $17.419\\pm0.238$ & 0.316 & 96 & 7.10 & 0.001 \\\\ $V$  & $-2.899\\pm0.052$ & $17.148\\pm0.033$ & 0.191 & 621 &  $-2.763\\pm0.141$ & $17.127\\pm0.176$ & 0.233 & 95 & 6.83 & 0.001 \\\\ $I_c$& $-3.073\\pm0.035$ & $16.657\\pm0.022$ & 0.126 & 604 &  $-2.951\\pm0.104$ & $16.609\\pm0.129$ & 0.162 & 88 & 7.15 & 0.001 \\\\ $J$  & $-3.237\\pm0.040$ & $16.330\\pm0.025$ & 0.126 & 481 &  $-3.035\\pm0.151$ & $16.184\\pm0.179$ & 0.134 & 48 & 5.00 & 0.007 \\\\ $H$  & $-3.347\\pm0.036$ & $16.116\\pm0.023$ & 0.114 & 481 &  $-3.099\\pm0.137$ & $15.925\\pm0.162$ & 0.122 & 48 & 7.69 & 0.001 \\\\ $K_s$& $-3.294\\pm0.043$ & $16.027\\pm0.028$ & 0.137 & 481 &  $-3.211\\pm0.144$ & $15.992\\pm0.171$ & 0.128 & 18 & 1.98 & 0.140 \\\\ \\hline $W_{bi}$  & $-3.463\\pm0.021$ & $15.933\\pm0.013$ & 0.074 & 598 &  $-3.507\\pm0.055$ & $15.999\\pm0.068$ & 0.086 & 88 & 0.886 & 0.413 \\\\ $W_{vi}$  & $-3.349\\pm0.019$ & $15.897\\pm0.012$ & 0.069 & 601 &  $-3.316\\pm0.050$ & $15.883\\pm0.062$ & 0.078 & 87 & 1.631 & 0.196 \\\\ \\hline \\end{tabular} \\\\ \\end{center} The subscripts $_S$ and $_L$ are for the short ($\\log P<1.0$) and long period Cepheids, respectively, while $a$, $b$ and $\\sigma$ are the slope, zero-point and dispersion of the fitted PL relations. \\end{table*} ",
        "conclusions": "Combining the results from the three statistical tests presented in the previous sections, we find that there is strong statistical evidence to suggest the LMC PL relation is nonlinear in the $BVI_cJH$ band but linear in the $K_sW_{bi}W_{vi}$ band. Including additional data from Persson et al. (\\cite{per04}) for the $JHK_s$ band does not alter the results as well. We have to emphasize that both of the testimator and SIC methods are applied to the $BI_cJHK_s$ band and the Wesenheit functions for the first time, in contrast to the $V$ band data that has been studied in Kanbur et al. (\\cite{kan07}).  The nonlinear LMC PL relation has been found from a Cepheid sample that consists of OGLE Cepheids only (Kanbur \\& Ngeow \\cite{kan04}). To extend the OGLE sample, mostly at the long period end, Ngeow \\& Kanbur (\\cite{nge06a}) included various additional data from literature (see table 1 of Ngeow \\& Kanbur \\cite{nge06a}), and again found strong evidence of nonlinearity of the LMC PL relation. In this Research Note, results using the Fouqu\\'{e} et al. (\\cite{fou07}) data alone, and with additional data from Persson et al. (\\cite{per04}), further supports the conclusion given in Ngeow \\& Kanbur (\\cite{nge06a}): that the sample selection does not play an important role in detecting the nonlinear LMC PL relation. However, Fouqu\\'{e} et al. (\\cite{fou07}) have suggested that the mixture of data used in previous work may lead to the nonlinearity seen in the statistical tests. This may certainly be the case and the analysis of a homogeneous sample, such as that provided by the ``LMC shallow survey'' (Fouqu\\'{e} 2007; Gieren 2007 -- private communication) is desirable.  The nonlinearity of the PL relation that is seen in the optical and $JH$ band but not in the reddening insensitive $K_s$ band and the Wesenheit function may suggest that extinction is the cause of the nonlinearity. However, extinction is not the only explanation and there is some evidence against the hypothesis of extinction errors as a cause for the apparent nonlinearity. The linearity of the $K_s$ band PL relation, as compared to other shorter wavelength PL relations, is expected from black-body arguments (Ngeow \\& Kanbur \\cite{nge06a}). Simply speaking, the temperature variation dominates the luminosity variation in the optical, and extends to $JH$ band for Cepheid-like temperatures. But in the $K_s$ band the luminosity variation is dominated by the radius variation of Cepheid variables. The proposed mechanism that may cause the nonlinear PL relation, the interaction between the hydrogen ionization front and the stellar photosphere (Kanbur \\& Ngeow \\cite{kan06}), will only affect the temperature variation and not the radius variation. The linearity of the Wesenheit functions is also not a surprise, and has been studied and discussed in Ngeow \\& Kanbur (\\cite{nge05w}) and in Koen et al. (\\cite{koe07}), and will not be repeated here.  Since the additional data used is mainly at the long period end, the possibility remains of systematic errors in reddening as a function of period. However we note that a reddening error as a function of long period LMC Cepheids would also change the observed properties of LMC Cepheids at other phases. A reddening error as a function of period such that LMC Cepheids obey a linear PL relation at mean light would force the LMC Cepheids to have a period-color relation such that they get bluer or hotter at maximum light as the period increases (see Figure \\ref{pcmax} for a schematic illustration). This is in stark contrast to the behavior of Galactic Cepheids and long period LMC Cepheids, which are known to have a flat period-color relation at maximum light (Code \\cite{cod47}; Simon et al \\cite{sim93}; Kanbur \\& Ngeow \\cite{kan04,kan06}; Kanbur et al \\cite{kan04a}). Moreover, it is difficult to explain, theoretically, how a Cepheid could get hotter at maximum light as the period increases.  The PL relation at phase $\\sim0.8$ described in Ngeow \\& Kanbur (\\cite{nge06b}) presents clearly the dramatic nature of the nonlinearity at 10 days and the dynamic nature of the PL relation as a function of phase. It is difficult to reconcile this behavior as being due to sampling errors and/or reddening errors. It is worth to point out that the mean light PL relation used in the literature is an average of the PL relations in all phases (Kanbur \\& Ngeow \\cite{kan04}; Kanbur et al. \\cite{kan04a}; Ngeow \\& Kanbur \\cite{nge06b}). nonlinearity of the PL relations at certain phases will certainly affect the linearity/nonlinearity of the mean light PL relation.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{8812fig5.eps}} \\caption{Schematic illustration for the argument with PC(max) relation. Top panels show the observed PC relations (after corrected for extinction) at mean (top-left panel) and maximum (top-right panel) light. Bottom panels show that if additional extinction as function of period to make the mean light PC relation linear, then the same extinction will cause the colors at maximum light get bluer as period increases, which are against observation and theoretical expectation.} \\label{pcmax} \\end{figure}  We have to remind that the data used in this study (and in most of our previous work) were published data that have been corrected for extinction using the ``state-of-the-art'' and well-developed methodology. If extinction error is believed to be the cause of nonlinear PL relations, it would imply that the extinction correction done previously in the literature is incorrect and/or incomplete. This would affect the previous work that using these extinction corrections, and those results need to be revised in future work."
    },
    "0710/0710.5544_arXiv.txt": {
        "abstract": "We use current theoretical estimates for the density of long cosmic strings to predict the number of strong gravitational lensing events in astronomical imaging surveys as a function of angular resolution and survey area. We show that angular resolution is the most important factor, and that interesting limits on the dimensionless string tension $\\G$ can be obtained by existing and planned surveys. At the resolution of the Hubble Space Telescope ($0\\farcs14$), it is sufficient to survey of order a few square degrees -- well within reach of the current HST archive -- to probe the regime $\\G\\sim10^{-7}$. If lensing by cosmic strings is not detected, such a survey would improve the limit on the string tension by a factor of two over that available from the cosmic microwave background. Future high resolution imaging surveys, covering a few hundred square degrees or more, either from space in the optical or from large-format radio telescopes on the ground, would be able to further lower this limit to $\\G\\lesssim10^{-8}$. ",
        "introduction": "Superstrings of cosmic size \\citep[introduced by ][]{Kib76} are a generic prediction of a number of string theory models \\citep[see, e.g., ][and references therein]{Pol04,D+K05}. Given their macroscopic nature, they are in principle detectable through astronomical observations. Therefore, they provide a perhaps unique opportunity for direct empirical tests of the physics of the very early Universe.  Considering strings whose only interactions are gravitational, all of their effects are controlled by the global constant dimensionless string tension~$\\G$.  Current limits on this parameter are given mainly by the properties of the cosmic microwave background (CMB) and by studies of pulsar timing. As far as the former is concerned, cosmic strings produce a smooth component in the CMB power spectrum and a non-gaussian signature in the CMB anisotropy map \\citep[e.g.,][]{L+W05,J+S07}. As far as the latter is concerned, strings produce a stochastic gravitational wave background, detectable in the time series of pulsars \\citep[e.g.,][]{KTR94,D+V05}. Current limits are $\\G \\lesssim 3 \\times 10^{-7}$ from the CMB~\\citep{PWW06}, and perhaps one or two orders of magnitude more stringent from pulsar timing depending on the details of the statistical analysis and the assumed string loop size distribution. An up-to-date review of current observational limits is given by~\\citet{Pol07}.  The idea of detecting cosmic strings by observing their gravitational lensing effect dates back to~\\citet{Vil84}. Briefly, cosmic strings produce a conical space time resulting in a very clear strong lensing signature, i.e.\\ they produce identical (neither parity-flipped, nor magnified or sheared) offset replica images of background objects, separated by an angle proportional to the string tension. Thus, strong gravitational lensing provides an opportunity for the {\\it direct} detection of cosmic strings, and even a single detected event would provide a measurement of the string tension, independent on the overall demographics of cosmic strings. A number of past studies have identified cosmic string lens candidates in optical surveys, but unfortunately none so far has withstood the test of higher resolution imaging \\citep{AHP06,Saz++07}.  In this paper we use current theoretical knowledge about the abundance of cosmic strings to predict the number of lensing events as a function of~$\\G$ for a realistic set of current and future imaging surveys. Our calculations show that, for sufficiently high angular resolution and sufficiently high (yet currently allowable) string tension, imaging surveys will either be able to detect cosmic string lenses or to at least set interesting limits on the dimensionless string tension parameter. For simplicity, we restrict our analysis to long strings, i.e. strings that are the size of the cosmic volume, and in particular straight with respect to the typical image separation, neglecting the contribution from string loops. \\citep[For a discussion of the lens statistics in future radio surveys from the loop population we refer the reader to the recent paper by ][]{MWK07} This assumption simplifies significantly the treatment and, since they do not depend on the detailed topology of the string network, nor on the timescales for gravitational decay, makes our predictions quite robust. ",
        "conclusions": "  \\begin{enumerate}  \\item Present-day high-resolution imaging surveys are capable of probing the putative cosmic string tension parameter to a factor of two lower than the current CMB limit, making lensing both competitive with, and complementary to, the pulsar timing methods. As has been noted before, in the event of a detection, gravitational lensing would provide a direct measurement of the tension, and perhaps the velocity, of this string.  \\item The main practical considerations in detecting string lensing events are two-fold: firstly, the expected faint pairs of images must first be carefully deblended and then understood in the context of neighbouring events; secondly, the survey geometry should be such that the fields are large enough to contain the characteristic multiple neighbouring events, but sparsely distributed to ensure that the global, not local, lensing rate is being probed.  \\item The upper bound on the tension, from the failure to detect a single string lensing event in a given survey, is principally determined by the available angular resolution. Relatively little is gained from studying an area of sky greater than some critical value: for typical ground-based optical image resolution, this critical survey size is a few tens of square degrees. Likewise, string tensions of $\\G\\sim10^{-8}$ should be able to be investigated in future surveys with image resolutions of 10-100~milliarcsec covering a few hundred square degrees.  \\end{enumerate}"
    },
    "0710/0710.2963_arXiv.txt": {
        "abstract": "{The identification of increasingly smaller signal from objects observed with a non-perfect instrument in a noisy environment poses a challenge for a statistically clean data analysis.} {We want to compute the probability of frequencies determined in various data sets to be related or not, which cannot be answered with a simple comparison of amplitudes. Our method provides a statistical estimator for a given signal with different strengths in a set of observations to be of instrumental origin or to be intrinsic.} {Based on the spectral significance as an unbiased statistical quantity in frequency analysis, Discrete Fourier Transforms (DFTs) of target and background light curves are comparatively examined. The individual False-Alarm Probabilities are used to deduce \\em conditional \\rm probabilities for a peak in a target spectrum to be real in spite of a corresponding peak in the spectrum of a background or of comparison stars. Alternatively, we can compute joint probabilities of frequencies to occur in the DFT spectra of several data sets simultaneously but with different amplitude, which leads to \\em composed \\rm spectral significances. These are useful to investigate a star observed in different filters or during several observing runs. The composed spectral significance is a measure for the probability that none of coinciding peaks in the DFT spectra under consideration are due to noise.} {\\sc Cinderella \\rm is a mathematical approach to a general statistical problem. Its potential reaches beyond photometry from ground or space: to all cases where a quantitative statistical comparison of periodicities in different data sets is desired. Examples for the composed and the conditional \\sc Cinderella \\rm mode for different observation setups are presented.} {} ",
        "introduction": "\\label{s1} %  The micromag precision, achieved by the MOST\\footnote{MOST is a Canadian Space Agency mission, jointly operated by Dynacon Inc., the University of Toronto Institute of Aerospace Studies, the University of British Columbia, and with the assistance of the University of Vienna, Austria.} (Microvariability \\& Oscillations of STars) mission (Walker et al.~2003; Matthews 2004), does not only provide exciting new results in asteroseismology, but reveals instrumental problems which challenge our data reduction techniques (see Sect.\\,\\ref{s1.1}). Cosmic ray impacts on the detector, stray light, positioning errors of the satellite, and thermal stability problems introduce periodic and, in the worst case, pseudo-periodic effects into photometric measurements. All this calls for new techniques in data reduction and analysis (see Sect.\\,\\ref{s1.2}).  Space observations in general can provide an unprecedented amount of measurements, requiring an enhanced degree of automatic data analysis without sacrificing accuracy and reliability. In this context, {\\sc SigSpec} (Reegen~2007) was developed to combine the Discrete Fourier Transform (DFT) -- a standard method to determine stellar pulsation frequencies -- with a clean statistical quantity: the spectral significance of a peak in an amplitude or power spectrum by comparison to white noise.  The basic idea of {\\sc Cinderella} is to use target and comparison data sets simultaneously for a cross-identification of artifacts in the frequency domain. It is the first technique permitting a statistically unbiased and quantitative comparison of different (not necessarily photometric) time series in the {\\em frequency} domain. Being applicable to practically all measurements of physical quantities over time, {\\sc Cinderella} has the potential to become a valuable tool beyond the scope of micromag space photometry.  \\subsection{The MOST mission}\\label{s1.1}  The first space telescope designed and built for photometric stellar seismology was EVRIS (Vuillemin et al.~1998), a 10-cm photoelectric telescope aboard the MARS-96 probe, but it unfortunately did not achieve the transfer orbit. An instrument providing photometric information on a large scale useful for asteroseismology was NASA's WIRE satellite, whose primary scientific goal of infrared mapping failed, but a 5-cm star tracker telescope with a CCD detector turned out to permit stellar photometry of remarkable quality (e.\\,g., Buzasi et al.~2000). The MOST satellite launched in June, 2003, assumed the role as a precursor to the CNES-led mission COROT (Baglin et al.~2004), which was successfully launched on December 27, 2006, and which is producing extremely useful space photometric data of hitherto unprecedented accuracy and volume.  MOST, WIRE and COROT are low-Earth-orbit (LEO) missions with comparable environmental effects (e.g., cosmic radiation, stray light scattered from the Earth's surface). A further commonality of all three missions is the requirement to extract asteroseismic information from a series of up to hundreds of thousands of CCD frames (or sub-rasters, respectively), each of which may consist of a few hundred to several million pixels. Hence, the present work may apply to other LEO space photometry missions and to ground-based multi-object photometry.  The MOST telescope is a 15-cm Maksutov optical telescope, supplied with a single broadband filter and initially with two identical CCD detectors: one used for science data acquisition, the other for the {\\em Attitude Control System} (ACS). Thanks to the low mass of $54$\\,kg and the ACS developed by Dynacon, Inc. (Groccott, Zee \\& Matthews~2003; Carroll, Rucinski \\& Zee~2004), a pointing stability to approximately $\\pm 1\\arcsec$ rms is achieved.  In {\\em Fabry Imaging} mode the telescope entrance pupil is imaged onto the CCD via a Fabry microlens as is shown by Figs.\\,7 and 8 of Walker et al.~(2003). Each Fabry Image is an annulus with an outer diameter of 44 pixels. The pixels in a square subraster outside the annulus are used to estimate the background. MOST also obtains {\\em Direct Imaging} photometry of typically $1-6$ stars, based on defocussed images (FWHM $\\sim$ 2.2 pixels; Rowe et al.~2006; Huber \\& Reegen~2008), and {\\em Guide Star} photometry of about $20-30$ stars (Aerts et al.~2006; Saio et al.~2006).  \\subsection{Data reduction}\\label{s1.2}  The data reduction described by Reegen et al.~(2006) applies linear correlations between pairs of target and background pixels for stray light correction. This so-called {\\em decorrelation technique} is also applicable to simultaneous photometry of several stars, in this case correlating variable vs.~constant stars.  Fig.\\,\\ref{fig1} illustrates the performance of the Fabry imaging photometry with MOST data of $\\beta$ CMi (Saio et al.~2007). The blue graph refers to the raw data and the red graph to the reduced light curve. The overall noise level decreased by an order of 10, and so did the harmonics of the orbital frequency of the spacecraft, ($\\approx 14.2$\\,d$^{-1}$ for $101.4$\\,min; Walker et al.~2003). However, instrumental peaks (dotted green lines) persisted on a lower level and their amplitudes still exceeded the stellar signal (main frequencies: $3.257$\\,d$^{-1}$ \\& $3.282$\\,d$^{-1}$; dotted black line).  \\begin{figure}\\includegraphics[width=256pt]{f1.eps} \\caption{The raw light curve ({\\em blue}) of the MOST Fabry target $\\beta$\\,CMi  and after data reduction ({\\em red}).  Harmonics of the satellite's orbital frequency ($\\approx 14.2$\\,d$^{-1}$; {\\em dotted green}), the detected stellar signal ($3.257$\\,d$^{-1}$ \\& $3.282$\\,d$^{-1}$; {\\em dotted black}) are indicated.}\\label{fig1} \\end{figure}  \\subsection{\\sc SigSpec}\\label{s1.3}  {\\sc SigSpec} (Reegen~2007), is based on DFT amplitude spectra and consecutive prewhitening of dominant peaks. But instead of considering the peak with the highest amplitude to be significant and estimating the reliability roughly in terms of signal-to-noise ratio, the {\\em Probability Density Function} (PDF) is employed. The PDF depends on the frequency and phase of the examined peak using white noise as a reference. The mean photometric magnitude in a time series is usually reduced to zero before evaluating the DFT. {\\sc SigSpec} the resulting statistical consequences into account, and is furthermore not restricted to Gaussian distributed residuals.  The {\\em False-Alarm Probability} is a frequently used statistical quantity in time series analysis. It is the probability of a peak at a given amplitude level to be generated by noise. Formally it is obtained through integration of the PDF. To avoid problems in computing extremely low numerical values, {\\sc SigSpec} returns a quantity called {\\em spectral significance} (hereafter abbreviated by ``sig''), which is the negative logarithm of the False-Alarm Probability. It gives  the number of uncorrelated data sets needed, containing pure noise, so that a peak in the Fourier domain appears which is comparable in amplitude and phase to the peak under consideration in the observed data.  Although {\\sc SigSpec} prevailed as a powerful tool for analyzing MOST photometry, it occasionally suffered from the weakness of having to refer to uncorrelated (i.e. white) noise.  \\subsection{The virtue of {\\sc Cinderella}}\\label{s1.4}  Frequencies with individual amplitudes and phases (``peaks'') in the DFT spectra of a target and comparison data sets are examined by {\\sc Cinderella} for compatibility. In other words, {\\sc Cinderella} allows us to investigate whether these data sets are related by any physical (deterministic) process. The procedure is the same if the comparison data represent sky background or a star with a different frequency spectrum as the target star, which -- in the best case -- is a constant star. Subsequently, the terms ``target star'' and ``comparison star'' will be used, keeping in mind that everything discussed here readily applies to sky readings instead of comparison stars as well. Obviously, all compared data sets have to be observed under similar circumstances. An extension of the method to handle more than one comparison data set is useful for multi-object environments, such as  photometry in a field.  In conditional mode, {\\sc Cinderella} establishes a quantitative comparison of significant frequencies occurring at the same time in at least two different data sets. It returns a statistically robust value, called {\\em conditional sig}, for the probability that a peak in the spectrum of one data set is not (deterministically) related to a peak in the other data set(s) within a given frequency resolution.  The alternative composed mode is dedicated to testing whether peaks in different DFT spectra with similar frequencies are ``real'', in the sense of not due to noise. The corresponding quantity, the {\\em composed sig} is a measure for the probability that none of the examined peaks is due to noise.  \\subsection{Frequency resolution}\\label{s1.6}  The question how to set the frequency difference acceptable for the consideration of peaks as coincidental is crucial to the examination of corresponding peaks in different DFT spectra. In this context, an alternative definition to the Rayleigh resolution, \\begin{equation}\\label{eqRayleigh} \\delta f_R := \\frac{1}{\\Delta t}\\: , \\end{equation} with $\\Delta t$ denoting the total time interval width of the time series is introduced by Kallinger, Reegen \\& Weiss~(2007). They suggest to additionally employ the sig for a peak amplitude according to \\begin{equation}\\label{eqKallinger} \\delta f_K := \\frac{1}{\\Delta t\\sqrt{\\mathrm{sig}\\left( A\\right)}}\\: , \\end{equation} for obtaining a more realistic criterion for matching peaks ({\\em frequency resolution}) than provided by Eq.\\,(\\ref{eqRayleigh}). Their numerical simulations show an excellent compatibility of this quantity, subsequently termed {\\em Kallinger resolution}, to the {\\em frequency error} derived by Montgomery \\& O'Donoghue~(1999).  For practical applications, it is useful to enhance the flexibility of {\\sc Cinderella} by introducing an exponent $z$ and to re-define the frequency resolution according to \\begin{equation}\\label{eqFRes} \\delta f := \\frac{1}{\\Delta t\\left[\\sqrt{\\mathrm{sig}\\left( A\\right)}\\right] ^z}\\: , \\end{equation} where $z$ usually attains values in the range $\\left[ 0,1\\right]$. The Rayleigh resolution is obtained for $z=0$, whereas $z=1$ yields the Kallinger resolution. ",
        "conclusions": "\\label{s5} %  This paper introduces a technique to interpret periodicities in an ensemble of data of common origin. {\\sc Cinderella} relies on {\\sc SigSpec} (Reegen~2007), thus benefitting from a correct employment of the complex phase information in Fourier Space on the one hand and a clean statistical description of interrelation of datasets on the other.  The conditional {\\sc Cinderella} mode is based on a quantitative comparison between one target and one or more comparison datasets and returns a measure of the probability (conditional sig) for periodicities identified in the target data to be deterministically related (to be `unique') to the target.  The composed {\\sc Cinderella} analysis returns a measure of the joint probability (composed sig) that a given periodicity observed in individual datasets -- but with different signal strengths -- is not due to noise. Such datasets could contain, e.g., measurements of the same target in different observing runs or with different instruments (e.g., different filters or simultaneous spectroscopy and photometry).  Our experience (as outlined in our examples in Sect.\\,\\ref{s3}) confirms that {\\sc Cinderella} reliably identifies residual instrumental signal in the MOST data even after a fairly sophisticated data reduction in the time-domain and also provides quantitative arguments to distinguish intrinsic from instrumental signal.  {\\sc Cinderella} is a statistically correct technique replacing what experienced observers achieve based on their ``good feelings'' when evaluating, for example, differential photometry, but, of course, the method is not limited to photometry. It quantitatively determines conditional and composed probabilities for matching peaks in DFT spectra of any kind of datasets containing periodicities.  \\subsection*{Acknowledgements}  PR received financial support from the Fonds zur F\\\"orderung der wis\\-sen\\-schaft\\-li\\-chen Forschung (FWF, projects P14546-PHY, P14984-PHY) Furthermore, it is a pleasure to thank D.\\,B.~Guenther (St.~Mary's Univ., Halifax), M.~Hareter, D.~Huber, T.~Kallinger (Univ.~of Vienna), R.~Kusch\\-nig, J.\\,M. Matthews (UBC, Vancouver), A.\\,F.\\,J.~Moffat (Univ.~de Montreal), D.~Punz (Univ.~of Vienna), S.\\,M. Rucinski (D.~Dunlap Obs., Toronto), D.~Sasselov (Harvard-Smithsonian Center, Cambridge, MA), G.\\,A.\\,H.~Walker (UBC, Vancouver), and K.~Zwintz (Univ.~of Vienna) for valuable discussion and support with extensive software tests.  Finally, we address our very special thanks to S.\\,M.~Mochnacki (University of Toronto) for his careful revision and valuable comments that substantially improved the presentation of this work."
    },
    "0710/0710.4068_arXiv.txt": {
        "abstract": "To optimise the science results of the asteroseismic part of the CoRoT satellite mission a complementary simultaneous ground-based observational campaign is organised for selected CoRoT targets. The observations include both high-resolution spectroscopic and multi-colour photometric data. We present the preliminary results of the analysis of the ground-based observations of three targets.  A line-profile analysis of 216 high-resolution FEROS spectra of the $\\delta$ Sct star HD~50844 reveals more than ten pulsation frequencies in the frequency range 5--18 d$^{-1}$, including possibly one radial fundamental mode (6.92 d$^{-1}$). Based on more than 600 multi-colour photometric datapoints of the $\\beta$ Cep star HD~180642, spanning about three years and obtained with different telescopes and different instruments, we confirm the presence of a dominant radial mode $\\nu_1=5.48695\\,$d$^{-1}$, and detect also its first two harmonics. We find evidence for a second mode $\\nu_2=0.3017$\\,d$^{-1}$, possibly a g-mode, and indications for two more frequencies in the 7--8 d$^{-1}$ domain. From Str\\\"omgren photometry we find evidence for the hybrid  $\\delta$ Sct/$\\gamma$ Dor character of the F0 star HD~44195, as frequencies near 3 d$^{-1}$ and 21 d$^{-1}$ are detected simultaneously in the different filters. ",
        "introduction": "The asteroseismic window of the CoRoT satellite mission aims at the monitoring of several types of pulsators along the Main Sequence. To optimise the science results, its targets are carefully chosen and selected. Preparatory observations from the ground have been a key stone in the selection process \\cite{ref1}. With the CoRoT satellite successfully launched (December 2006), simultaneous ground-based observations are very important and are complementary to the space data.  Multi-colour photometry provides colour information, which allows identification of the degree $\\ell$ by means of amplitude ratios and phase shifts, while high-resolution spectroscopy allows the detection of high-degree modes and the identification of both the degree $\\ell$ and the azimuthal order $m$.  In this framework a simultaneous ground-based observing campaign was organised by the CoRoT/SWG Ground-based Observations Working Group. Large Programme proposals (i.e. guaranteed observing time during four observing seasons) were submitted and accepted at ESO La Silla (FEROS/2.2m; 10+5 nights per semester), Observatoire de Haute Provence (SOPHIE/1.92m; 10+5 nights per semester) and Calar Alto Astronomical Observatory (FOCES/2.2m; 10+10 nights per semester), with the aim to obtain multi-site time-series of high-resolution spectra of a selection of $\\delta$ Sct, $\\gamma$ Dor, $\\beta$ Cep and Be CoRoT primary and secondary targets.  The first two observing seasons (winter 2006-2007 and summer 2007), (nearly) coinciding with CoRoT's IR01 (first Initial Run) and LRc1 (first Long Run in the center direction), have been succesfully completed. \\Tref{logbook} gives an overview of the targets and the amount of spectra obtained.  In addition to the Large Programmes, and in continuation of a project started three years ago, observing time has been awarded at smaller telescopes with multi-colour photometric instruments (Str\\\"omgren  photometry: 90cm@Sierra Nevada Observatory (SNO), 1.5m@San Pedro M\\'{a}rtir Observatory (SPMO); Geneva photometry: 1.2m Mercator@Observatorio Roque de los Muchachos (ORM); Johnson photometry: Konkoly Observatory (KO)). In particular, 18 and 14 consecutive nights have been awarded with the $uvby\\beta$ photometers at SPMO and SNO, respectively, in Nov-Dec 2006 and 2007.  \\begin{center} \\begin{table}[h] \\caption{\\label{logbook} Logbook of the spectroscopic observations, obtained in Jan-Feb 2007 and May-Jul 2007 with the FEROS, SOPHIE and FOCES instruments, dedicated to a selection of  targets of the CoRoT IR01 (Feb-Apr 2007), LRc1 (summer 2007) and LRa1 (winter 2007-2008) runs. The different columns give the target name, its V magnitude, spectral type, variable type, name of the CoRoT run, amount of spectra obtained with FEROS, SOPHIE and FOCES, respectively. The targets indicated in boldface are discussed in this paper. \\\\} \\centering \\begin{tabular}{llllllll} \\br target & V & Sp.T. & type & CoRoT run & FEROS & SOPHIE & FOCES  \\\\ \\mr {\\bf HD~50844} & 9.09 & A2 & $\\delta$ Sct & IR01 & 216 & & \\\\ HD~50747 & 5.45 & A4 & SB2 & IR01 & 17 & 14 & 6\\\\ HD~51106 & 7.35 & A3m & SB2 & IR01 & 15 & 14 & 4\\\\ HD~50846 & 8.43 & B5 & EB & IR01 & 16 & 12 & 4\\\\ \\hline HD~49434 & 5.74 & F1V & $\\gamma$ Dor & LRa1 & 71 & 444 & 75 \\\\ HD~50209 & 8.36 & B9Ve & Be star & LRa1 & 68 & & \\\\ \\hline HD~181555 & 7.52 & A5 & $\\delta$ Sct & LRc1 & 343 & 66 & 285 \\\\ HD~174966 & 7.72 & A3 & $\\delta$ Sct & LRc1 & & & 119 \\\\ {\\bf HD~180642} & 8.27 & B1.5III & $\\beta$ Cep & LRc1 & 213 & 35 & \\\\ HD~181231 & 9.69 & B9.0V & Be star & LRc1 & 72 & & \\\\ \\br \\end{tabular} \\small{SB2=double-lined spectroscopic binary; EB=eclipsing binary} \\end{table} \\end{center}  In the next sections we report on the preliminary results of the analysis of the ground-based datasets of the $\\delta$ Sct star HD~50844 and the $\\beta$ Cep star HD~180642, and give an outlook towards an interesting candidate CoRoT target, the  hybrid $\\delta$ Sct/$\\gamma$ Dor star HD~44195. ",
        "conclusions": "We obviously are in a challenging era of asteroseismology. The preparatory work and the simultaneous ground-based observations of the CoRoT mission require a huge effort of the asteroseismic community. Challenging objects are chosen (e.g. faint targets, fast rotating stars) and in more than one way are we pushing the limits. The ambitious choice of targets asks for new methodologies and analysis techniques, and requires a close collaboration between theoreticians and observers. With joined forces, with shared expertise and with complementary data from space and from the ground we have excellent prospects to probe the interiors of stars.    \\ack KU acknowledges financial support from a \\emph{European Community Marie Curie Intra-European Fellowship}, contract number MEIF-CT-2006-024476. This work is supported by the italian ESS project, contract ASI/INAF I/015/07/0, WP 03170, and by the Research Council of Leuven University under grant GOA/2003/04."
    },
    "0710/0710.4859_arXiv.txt": {
        "abstract": "{Many binary stellar systems in which the primary star is beyond the asymptotic giant branch (AGB) evolutionary phase show significant orbital eccentricities whereas current binary interaction models predict their orbits to be circularised.} {In the search for a mechanism to counteract the circularising effect of tidal interaction we analyse how the orbital parameters in a\tsystem are modified under mass loss and mass exchange among its binary components.} {We propose a model for enhanced mass-loss from the AGB star due to tidal interaction with its companion, which allows a smooth transition between the wind and Roche-lobe overflow mass-loss regimes. We explicitly follow its effect along the orbit on the change of eccentricity and orbital semi-major axis, as well as the effect of accretion by the companion. We calculate timescales for the variation of these orbital parameters and compare them to the tidal circularisation timescale.} {We find that in many cases, due to the enhanced mass loss of the AGB component at orbital phases closer to the periastron, the net eccentricity growth rate in one orbit is comparable to the rate of tidal circularisation. We show that with this eccentricity enhancing mechanism it is possible to reproduce the orbital period and eccentricity of Sirius system, which under the standard assumptions of binary interaction is expected to be circularised. We also show that this mechanism may provide an explanation for the eccentricities of most barium star systems, which are expected to be circularised due to tidal dissipation. } {By proposing a tidally enhanced model of mass loss from AGB stars we find a mechanism which efficiently works against the tidal circularisation of the orbit. This mechanism can explain the significant eccentricities observed in binary systems containing a white dwarf and a less evolved companion, which are predicted to be circularised due to their proximity, such as Sirius and systems with barium stars.} ",
        "introduction": "Detached binary systems containing a white dwarf and a relatively unevolved companion, i.e., a main-sequence star or a (sub)giant, are a useful tool to understand the binary evolution of systems with an asymptotic giant branch (AGB) star, given that their orbital and chemical properties do not change significantly from the moment that the primary finished its AGB evolution and became the current white dwarf. An example of such a system is Sirius, a 2.1 M$_{\\odot}$ main sequence star with a 1.05 M$_{\\odot}$ white dwarf companion in a 50-year orbit \\citep[e.g., ][]{1960JO.....43..145V,1978ApJ...225..191G}. Under the standard picture of binary evolution with tidal interaction this system should have circularised when the primary became an AGB star, as we show in $\\S$\\ref{sec:sirius} of this paper. However, this system has an eccentricity $e=0.59$ and is not an exception. A similar problem is faced when considering the barium-star systems, which are red giants with over-abundances of $s$-process elements (prominently barium) with white dwarf companions. The best current explanation for the enhancement of $s$-process elements in these stars is that they accreted mass from their companion when it was an AGB star. Due to the large size of AGB stars these systems are expected to be circularised by tidal interaction for periods smaller than about 3500-4000 days \\citep{2003ASPC..303..290P}. However, barium-star systems with period as short as 600 days are observed to be significantly eccentric \\citep{1998A&A...332..877J}. The problems stated above are indications that a mechanism must exist that counteracts the circularising effect of the tides by the time the primary star is on the AGB. \\citet{1995A&A...293L..25V} propose that enhanced mass loss at periastron could enhance the eccentricity and \\citet{2000A&A...357..557S} shows that this works when the AGB star fills its Roche lobe during periastron passages. It has also been proposed that the tidal interaction between a binary system and a circumbinary disk could account for an eccentricity enhancing mechanism \\citep{1996A&A...314L..17W,1998Natur.391..868W}. However, \\citet{2000A&A...357..557S} argues that the masses of observed circumbinary disks are too small to counteract the circularisation.  Mass loss from stars in binary systems is commonly treated as single-star wind mass-loss as long as the stars are inside their Roche lobes, while when one star fills its Roche lobe the mass-loss rate increases abruptly to the high values that correspond to Roche-lobe overflow. This approximation is fairly accurate for stars with a steep density gradient in their atmospheres, but it is not appropriate for the case of AGB stars, given their large atmospheric pressure scale height and their weakly bound envelope. The fact that AGB stars undergo dynamical pulsations further reduces density gradient in the layers above the photosphere \\citep[e.g.,][]{1988ApJ...329..299B}. We propose a prescription for enhanced mass loss which smoothly grows from the single star wind mass loss rate to the Roche lobe overflow mass loss rate as the radius of the star approaches its Roche lobe radius. This gives a variable mass loss rate along eccentric orbits which is higher at orbital phases closer to the periastron, even if the star does not fill its Roche lobe. This effect works in the same way as discussed above, but no filling of the Roche lobe is needed, so that (during the AGB phase) this effect is permanently competing against tidal circularisation.  \\citet{1988A&A...205..155B} carried out calculations of the variation of orbital parameters considering instantaneous mass transfer and only linear momentum conservation. Angular momentum conservation was not taken into account in this exploratory study. An extension including angular momentum conservation was made by \\citet{2000A&A...363..660L}, but they make use of the orbital average of the distance between the components and orbital angular velocity to carry out their calculations, which in the case of a variable mass-loss along the orbit does not give the correct results.  In $\\S$\\ref{sec1} we revise the variation of orbital parameters due to stellar mass loss and mass transfer by taking into account the conservation laws of linear and angular momentum with respect to the centre of mass of the actual binary system. We do not make use of the orbital averages a priori, but give the variations as a function of the orbital phase, allowing different rates of mass loss and mass transfer along the orbit. In $\\S$\\ref{sec2} we calculate the rate of change of the eccentricity according to different assumptions of mass loss and accretion. In $\\S$\\ref{sec3} we evaluate the competition between the tidal circularisation and the eccentricity pumping due to our proposed mass loss, and show that the latter can prove effective in counteracting the former for systems such as Sirius and barium stars in which their AGB component did not fill its Roche lobe. A summary is presented in $\\S$\\ref{sec4}. ",
        "conclusions": "\\label{sec4} We have revised the description of the evolution of orbital parameters due to mass loss and mass transfer in a binary system where the primary is an AGB star. We also propose a tidally enhanced mass-loss rate for AGB stars in binary systems which allows a smooth transition between wind mass loss and mass loss due to Roche lobe overflow. With our revised prescription for the evolution of the orbital eccentricity and the fact that our proposed mass loss is not constant along an eccentric orbit, we find an eccentricity enhancing mechanism which % counteracts the circularising effect of tidal dissipation. We have shown that the standard scenario of binary interaction cannot explain the observed orbital parameters of the eccentric binary system Sirius unless the strength of tidal dissipation in AGB stars is at least 3 orders of magnitude smaller than in normal red giants. On the other hand, our models with the eccentricity-enhancing mechanism can reproduce the orbital properties of the Sirius system with a reasonable choice of the tidal dissipation parameter ($f'=0.25$). Our eccentricity enhancing mechanism also allows binary systems containing barium stars to remain eccentric with periods as short as about 1000 days under reasonable parameter assumptions, while under the same assumptions in the standard scenario of tidal circularisation only systems with periods longer than about 2500 days are expected to show a significant eccentricity. We also show that we can explain the eccentricities of almost all barium star systems which do not fill their Roche lobes if the uncertain convective tidal dissipation strength in AGB stars is reduced by a factor of 10, compared to what is measured for red giant stars. Moreover, this assumption allows Roche-lobe overflow to occur in significantly eccentric systems, which may be the key to explain shorter-period barium-star systems which are still eccentric.   Whether our eccentricity-enhancing mechanism is a viable solution to the problem of explaining short-period eccentric binaries with (former) AGB stars, in particular barium stars, needs to be tested with a binary population synthesis model. Results of such a test will be given in a separate paper \\citep{BMP07}. Preliminary barium star population synthesis calculations indicate that when the full evolutionary history of a binary population is taken into account, a reduction of the tidal dissipation strength of only $f'=0.25$ is needed to reproduce the observed systems. This is consistent with our findings about the values of $f'$ needed to reproduce the Sirius system."
    },
    "0710/0710.1198_arXiv.txt": {
        "abstract": "We investigate the influence of an interaction between dark energy and dark matter upon the dynamics of galaxy clusters. We obtain the general Layser-Irvine equation in the presence of interactions, and find how, in that case, the virial theorem stands corrected. Using optical, X-ray and weak lensing data from 33 relaxed galaxy clusters, we put constraints on the strength of the coupling between the dark sectors.  Available data suggests that this coupling is small but positive, indicating that dark energy might be decaying into dark matter. Systematic effects between the several mass estimates, however, should be better known, before definitive conclusions on the magnitude and significance of this coupling could be established. ",
        "introduction": "Cosmological accelerated expansion is by now a well-established observational fact \\cite{cosmoaccel,wmap,bao}, leading either to an asymptotically de Sitter cosmology, plagued with an astonishingly small cosmological constant, or else to a universe filled up to 80\\% with a strange dynamical component with negative pressure -- dark energy \\cite{darkenergy}.  If dark energy contributes a significant fraction of the content of the Universe, it is natural, in the framework of field theory, to consider its interactions with the remaining fields of the Standard Model and well-motivated extensions thereof. For lack of evidence to the contrary, interactions of dark energy or dark matter with baryonic matter and radiation must be either inexistent or negligible. Nevertheless, some level of interaction between dark energy and the dark matter sector, which is present in most extensions of the Standard Model, is still allowed by observations.   The possibility that dark energy and dark matter can interact has been studied in \\cite{amendola}-\\cite{[14]}, among others. It has been shown that the coupling between a dark energy (or \ufffdquintessence\ufffd) field and the dark matter can provide a mechanism to alleviate the coincidence problem \\cite{amendola,[11]}. A suitable choice of the coupling, motivated by holographic arguments, can also lead to the crossing of the \ufffdphantom barrier\ufffd which separates models with equations of state $w = p/\\rho > -1$ from models with $w < -1$ \\cite{[12]} -- see also \\cite{[14],[15]}. In addition, it has been argued that an appropriate interaction between dark energy and dark matter can influence the perturbation dynamics and affect the lowest multipoles of the CMB spectrum, accounting for the observed suppression of the quadrupole \\cite{[13],[16]}.  The strength of the coupling could be as large as the fine structure constant \\cite{[13],[17]}. Recently, it was shown that such an interaction could be inferred from the expansion history of the universe, as manifested in, e.g., the supernova data together with CMB and large-scale structure\\cite{[18]}. Nevertheless, the observational limits on the strength of such an interaction remain weak \\cite{[19]}.    A complementary and fundamentally different way in which the coupling between dark energy and dark matter can be checked against the observations is through its impact on large-scale structure. If dark energy is not a cosmological constant, it must fluctuate in space and in time -- and, in particular, if dark energy couples to dark matter, then it must surely be dynamical. If that is the case, dark energy affects not only the expansion rate, but the process of structure formation as well, through density fluctuations, both in the linear \\cite{[11]}, \\cite{[20]}-\\cite{[23]} and the non-linear \\cite{[24],[25]} regimes. The growth of dark matter perturbations can in fact be enhanced due to the coupling between these two components \\cite{[13],[14],[26]}.   Recently, it was suggested that the dynamical equilibrium of collapsed structures would be affected by the coupling of dark energy to dark matter, in a way that could be observed in the galaxy cluster Abell A586 \\cite{[27]}. The basic idea is that the virial theorem is distorted by the non-conservation of mass caused by the coupling.    In this paper we show precisely how the Layser-Irvine equation, which describes the flow to virialization \\cite{[28]}, is changed by the presence of the coupling, in such a way that the final state of equilibrium violates the usual virial condition, $2K + U= 0$, where $K$ and $U$ are respectively the kinetic and the potential energies of the matter constituents in an isolated system. We show that this violation leads to a systematic bias in the estimation of masses of clusters if the usual virial conditions are employed. Although it is still possible that systematic errors from observations smear the results, the fact that some shift in the mean value of the coupling for two independent sets of observations (compared to the third set) may signalize some new physics.  Even though the uncertainties associated with any individual galaxy cluster are very large, by comparing the naive virial masses of a large sample of clusters with their masses estimated by X-ray and by weak lensing data, we may be able to constrain such a bias and to impose tighter limits on the strength of the coupling than has been achieved before. ",
        "conclusions": "We have estimated the effective coupling between dark energy and dark matter through the internal dynamics of galaxy clusters. In the presence of coupling, the flow of mass and energy between the components changes the virial condition in a way that can be tested by comparing different estimators for the mass of clusters. We searched for this signature in 33 galaxy clusters for which reliable X-ray, weak lensing and optical data were available.  Our results indicate a weak preference for a small but positive effective coupling constant $\\bar\\zeta$ -- in line with predictions made by some of us \\cite{[12],[13]}. Since the statistical significance is still low ($\\Delta \\chi^2/{\\rm d.o.f.} \\sim 0.2$), it is paramount that more clusters (with homogeneous mass determinations and good control of systematics) be tested. If a significant indication of such coupling is still found, this would open a tantalizing new window on the nature of the dark sector."
    },
    "0710/0710.3648_arXiv.txt": {
        "abstract": "{} {We present a qualitative analysis of key (but yet unappreciated) linear phenomena in stratified hydrodynamic Keplerian flows: \\emph{(i)} the occurrence of a vortex mode, as a consequence of strato-rotational balance, with its transient dynamics; \\emph{(ii)} the generation of spiral-density waves (also called inertia-gravity or $g\\Omega$ waves) by the vortex mode through linear mode coupling in shear flows.} {Non-modal analysis of linearized Boussinesq equations were written in the shearing sheet approximation of accretion disk flows.} {It is shown that the combined action of rotation and stratification introduces a new degree of freedom, vortex mode perturbation, which is in turn linearly coupled with the spiral-density waves. These two modes are jointly able to extract energy from the background flow, and they govern the disk dynamics in the small-scale range. The transient behavior of these modes is determined by the non-normality of the Keplerian shear flow. Tightly leading vortex mode perturbations undergo substantial transient growth, then, becoming trailing, inevitably generate trailing spiral-density waves by linear mode coupling. This course of events -- transient growth plus coupling -- is particularly pronounced for perturbation harmonics with comparable azimuthal and vertical scales, and it renders the energy dynamics similar to the 3D unbounded plane Couette flow case.} {Our investigation strongly suggests that the so-called bypass concept of turbulence, which has been recently developed by the hydrodynamic community for spectrally stable shear flows, can also be applied to Keplerian disks. This conjecture may be confirmed by appropriate numerical simulations that take the vertical stratification and consequent mode coupling into account in the high Reynolds number regime.} ",
        "introduction": "According to classical fluid dynamics, unmagnetized disk flows in Keplerian rotation (more generally: with angular momentum increasing outward and with no extremum of vorticity) are spectrally stable; however, there is irrefutable observational evidence that such disks have to be turbulent.  Due to this apparent contradiction, disk turbulence is often considered as some sort of mystery.  However, an analogous dilemma that existed in laboratory/engineering flows has been solved by the hydrodynamic community in the 90s of the last century, where a breakthrough was accomplished in the comprehension of turbulence in spectrally/asymptotically stable shear flows (e.g. in the plane Couette flow).  Let us briefly recall the essence of that breakthrough  (see Chagelishvili et al. 2003 for details). Traditional stability theory followed the approach of Rayleigh (1880) where the instability is determined by the presence of exponentially growing modes that are solutions of the linearized dynamic equations. Only recently has one become aware that operators involved in the the modal analysis of plane shear flows are not normal, hence that the corresponding eigenfunctions are non-orthogonal and would strongly interfere (Reddy et al. 1993). For this reason, the emphasis was shifted in the 90s from the analysis of long time asymptotic flow stability to the study of short time behavior. It was established that \\emph{asymptotically/Rayleigh stable flows allow for linear transient growth of vortex and/or wave mode perturbations} (cf. Gustavsson 1991; Butler and Farrell 1992; Reddy and Henningson 1993; Trefethen et al. 1993). This fact incited a number of fluid dynamicists to examine the possibility of a subcritical transition to turbulence, with the linear stable flow finding a way to bypass the usual route to turbulence (via linear classical/exponential instability). On closer examination, the perturbations reveal rich and complex behavior in the early transient phase, which leads to the expectation that they may become self-sustaining  when there is nonlinear positive feedback.  Based on the interplay of linear transient growth and nonlinear positive feedback, a new concept emerged in the hydrodynamic community for the onset of turbulence in spectrally stable shear flows and was named \\emph{bypass transition} (cf. Boberg \\& Brosa 1988; Butler \\& Farrell 1992; Farrell \\& Ioannou 1993; Reddy \\& Henningson 1993; Gebhardt \\& Grossmann 1994; Henningson \\& Reddy 1994; Baggett et al. 1995; Grossmann 2000; Reshotko 2001; Chagelishvili et al. 2002; Chapman 2002). The bypass scenario differs fundamentally from the classical scenario of turbulence. In the classical model, exponentially growing perturbations permanently supply energy to the turbulence and they do not need any nonlinear feedback for their self-sustenance, so the role of nonlinear interaction is just to reduce the scale of perturbations to that of viscous dissipation. In the bypass model, nonlinearity plays a key role.  The nonlinear processes are conservative, but in the case of positive feedback, they ensure the repopulation of perturbations that are able to extract energy transiently from the mean flow. The self-sustenance of turbulence is then the result of a subtle and balanced interplay of linear transient growth and nonlinear positive feedback. Consequently, thorough examination of the nonlinear interaction between perturbations is a problem of primary importance, and the first step is to search and to describe the linear perturbation modes that will participate in the nonlinear interactions.  Such linear transient growth is also at work in rotating hydrodynamic disk flows; however, the Coriolis force causes a quantitative reduction of the growth rate  there which delays the onset of turbulence.  Keplerian flows are therefore expected to become turbulent for Reynolds numbers a few order of magnitudes higher than for plane subcritical flows (see: Longaretti 2002; Tevzadze et al. 2003). The possibility of an alternate route to turbulence gave new impetus to the research on the dynamics of astrophysical disks (Lominadze et al. 1988; Richard \\& Zahn 1999; Richard 2001; Ioannou \\& Kakouris 2001; Tagger 2001; Longaretti 2002; Chagelishvili et al. 2003; Tevzadze et al. 2003; Klahr \\& Bodenheimer 2003; Yecko 2004; Afshordi et al. 2004 Umurhan \\& Regev 2004; Umurhan \\& Shaviv 2005; Klahr 2004; Bodo et al. 2005; Mukhopadhyay et al. 2005; Barraco \\& Marcus 2005; Johnson \\& Gammie 2005a, 2005b; Umurhan 2006).  By adapting the progress of the hydrodynamic community to the disks flow, this research is promising for solving the disks' hydrodynamic turbulence problem.  But it remains to be seen whether this route to turbulence actually applies to astrophysical disks. Compared to plane shear flows, these possess two additional properties: differential rotation and vertical stratification. Separate studies of these factors show that each exerts a stabilizing effect on the flow: these include numerical calculation of the stability of unstratified flows by Shen et al. (2006), experiments on Keplerian rotation without stratification by Ji et al. (2006), estimates of the growth rates with stratification by Brandenburg \\& Dintrans (2006). However, it appears that the combined action of differential rotation and stratification introduces a new degree of freedom that may influence the flow stability and lead to turbulence at a high enough Reynolds number. Indeed, it has been shown that strato-rotational flows may exhibit global instability in bounded domains (Dubrulle et al. 2005). However, it is probable that local disk dynamics will also lead to hydrodynamic turbulence. The study of the linear perturbations in strato-rotational flow in the local limit can be found in Tevzadze et al. (2003; hereafter T03); it is shown there that the combined action of rotation and stratification generates an aperiodic vortex mode, which undergoes nonmodal transient growth. We conjecture that this transient growth may be the main energy source for the turbulence in the bypass scenario.  Although the importance of the transient exchange of energy between perturbations and mean flow has now been realized by most working in the field, only a few seem aware that another linear process may play also a crucial role, namely the linear coupling of modes, which allows for transient exchange of energy between them. As shown by Chagelishvili et al. (1997a,b) and Gogoberidze et al. (2004), the energy exchange between modes is inherent to shear flows (as the transient exchange of energy between the mean flow and perturbations), and it determines in many respects the diversity of perturbation modes and, therefore, of nonlinear processes. Once one fully realizes the role of nonlinear processes in the bypass concept discussed above, it becomes evident that the neglect of linear mode coupling may lead to an incorrect picture of nonlinear (and, consequently, turbulent) phenomena.  There are signs of such linear mode coupling in the simulations performed by Klahr (2004), Barraco \\& Marcus (2005), and Brandenburg \\& Dintrans (2006), but apparently they were not identified as such. Compressive waves are present in the simulation by Johnson \\& Gammie (2005b), along with vortical perturbations, but their origin is not recognized, namely linear mode coupling. On this mode coupling, attention is focused in T03 and Bodo et al. (2005).  The latter paper studies the linear dynamics of an imposed two-dimensional pure vortex mode perturbation,  in a compressible Keplerian disk with constant mean pressure and density. (Two-dimensionality, i.e. the neglect of cross-disk variation, is only correct  for perturbations with characteristic scales close to or larger than the vertical stratification scale of the disk.) Two modes -- a vortex mode and a spiral-density wave mode -- exist in the system, and they are strongly coupled. This investigation points out the importance of mode coupling and the necessity of considering compressibility for dynamic processes with characteristic scales close to or larger than the disk thickness.  In T03 we studied the linear dynamics of three-dimensional small-scale perturbations (with characteristic length scales much shorter than the disk thickness) in compressible, vertically (stably) stratified Keplerian disks. \\emph{The first novelty} presented in that paper is the occurrence of an interplay between the disk rotation and stratification. Separately, each of these factors is stabilizing. However, their interplay gives rise to a new vortex/aperiodic mode that is able to extract the basic flow energy  transiently. \\emph{The second novelty} is the existence of a linear coupling of that vortex mode with spiral-density wave (SDW) modes, that makes SDWs valuable participants of the dynamical processes. Furthermore, we compared the linear dynamics of the small-scale perturbations in steady stratified disks with that of the perturbations in unbounded plane Couette flow. We showed that just the linear coupling of the disk modes makes the dynamics in the disk and plane flows similar to each other (provided the Reynolds number of the disk flow is chosen about three orders of magnitudes higher than in the plane flow). This similarity suggests that the bypass concept can also be applied to disk flow and it  further motivates investigation in this direction.  In this respect, the present paper is a sequel of T03: we focus again on the mathematical and physical aspects of mode coupling, while introducing a significant simplification by neglecting the rotational-acoustic waves. This is justified by the fact that the characteristic timescale of these modes is much shorter than for the two other modes, when the characteristic lengthscale of the perturbations is much shorter than the disk thickness. Then the rotational-acoustic waves are not coupled to the vortex and SDW modes, and they play a negligible role in the  slow, small-scale dynamics.  That is why we can cut out the rotational-acoustic waves (i.e., the flow compressibility) without detriment to the dynamic picture and confine ourselves to the Boussinesq approximation. Keeping just vortex and SDW modes, this approximation simplifies mathematical aspects of the problem and allows an advance in the analytical description and comprehension of mode coupling.  The paper is organized as follows. In Sect. 2 we introduce the physical approximations and the mathematical formalism, and describe the linear strato-rotational balance and the perturbation modes. In Sect. 3 we present the qualitative and quantitative analysis of the linear dynamics of perturbations. We summarize and discuss the results in Sect. 4. ",
        "conclusions": "Our motivation in studying mode coupling in the presence of vertical stratification is quite clear: in unstratified rotating flows the strato-rotational balance is absent (see Eq. \\ref{vort1}),  hence there will be no vortex mode, whose role is so powerful in extracting the shear flow energy. Moreover, in stratified rotating flows mode coupling causes the generation of SDW, which are able to conserve the extracted energy. Thus one expects astrophysical disks, which are vertically stratified, to demonstrate intrinsically different behavior compared to unstratified rotating flows. Therefore it is not possible to draw conclusions about the stability of astrophysical disks by considering unstratified rotating flows, as has been done often in the past and again in recent investigations: analytical (Balbus 2006), numerical (Shen et al. 2006), and experimental (Ji et al. 2006).  Our analysis shows that, in the local limit, spectrally stable stratified Keplerian disks allow for two modes of perturbations -- vortex and SDW -- that are jointly able to extract the background flow energy and determine the disk dynamical activity in the small-scale range. These modes are linearly coupled due to the non-normality of Keplerian/shear flow. The coupling is asymmetric: the vortex mode is able to generate the related  SDW, but the inverse is not true. This mode coupling is transient (like the energy exchange between perturbations and the basic flow): the SFH of the vortex mode generates the wave SFHs during the brief time interval where it switches from leading to trailing, thus rendering the dynamics non-adiabatic.  At first sight, the mode coupling described here seems similar to the phenomenon of geostrophic adjustment, since both lead to the generation of the spiral density waves. However, these processes are intrinsically different. Geostrophic adjustment is an initial value problem that describes the transition from an initially unbalanced state to that of geostrophic balance. In the general case, part of the initial conditions containing no potential vorticity (the wave component) will radiate away as inertial-gravity waves (see, e.g., Pedlosky 2003) during the geostrophic adjustment, whereas our initial conditions do not include zero potential vorticity corresponding to the wave component. Hence, we have eliminated the wave generation due to the process of the initial geostropic adjustment.  The waves generated in our case stem from linear mode coupling induced by the velocity shear. Moreover, geostrophic adjustment is mainly a nonlinear process that leads to equilibrium. In contrast, wave generation due to mode coupling is a linear process and does not describe the relaxation of the system, but hopefully the opposite: it promotes the transition to turbulence.  The linear dynamics of each leading SFH of vortex mode proceeds in the following way: initially, the SFH extracts energy from the basic flow and it grows. At the same time $~k_x(t)/k_y \\to -0$. Becoming trailing, the vortex SFH generates the related SFH of SDW. In what follows, while tilting (i.e., increasing $~k_x(t)/k_y$), the wave SFH keeps the energy, whereas the vortex SFH gives its energy back to the basic flow, so the energy gained  by the leading vortex SFH is conserved by the SDW. This course of events -- the transient growth plus coupling -- is strongly pronounced for SFHs with $~k_z/k_y \\sim 1$ and makes their energy dynamics similar to that of a 3D unbounded plane Couette flow.  We are aware that there are differences between them: in the Couette case, only the vortex mode participates in the dynamical process, whereas in the disk case, this role is played (in the local limit) by the symbiosis between vortex and SDW perturbations. However, given the similarities that we have discussed,  we conjecture that the bypass concept, which has been developed for the transition to turbulence in the Couette flow, is also applicable to rotating stratified disks. One question remains, of course: will the nonlinear interactions provide positive feedback that is efficient enough for that bypass transition? To answer it, more numerical simulations are needed, which must include the vertical stratification and grasp the mode coupling."
    },
    "0710/0710.3224_arXiv.txt": {
        "abstract": "We report the detection of OH~1667~MHz and wide \\hi~21cm absorption at $z \\sim 0.67$ towards the red quasar 1504+377, with the Green Bank Telescope and the Giant Metrewave Radio Telescope. The \\hi~21cm absorption extends over a velocity range of $\\sim 600$~km/s blueward of the quasar redshift ($z=0.674$), with the new OH~1667~MHz absorption component at $\\sim -430$~\\kms, nearly coincident with earlier detections of mm-wave absorption at $z \\sim 0.6715$. The atomic and molecular absorption appear to arise from a fast gas outflow from the quasar, with a mass outflow rate ${\\dot M} \\sim 12 M_\\odot$~yr$^{-1}$ and a molecular hydrogen fraction $f_{\\rm H_2} \\equiv (N_{\\rm H_2}/N_{\\rm HI}) \\sim 0.2$. The radio structure of 1504+377 is consistent with the outflow arising due to a jet-cloud interaction, followed by rapid cooling of the cloud material. The observed ratio of HCO$^+$ to OH column densities is $\\sim 20$~times higher than typical values in Galactic and high-$z$ absorbers. This could arise due to small-scale structure in the outflowing gas on sub-parsec scales, which would also explain the observed variability in the \\hi~21cm line. ",
        "introduction": "\\label{intro}  The quasar 1504+377 is a rare case of a radio-loud active galactic nucleus (AGN) hosted by a disk galaxy (e.g. \\citealt{perlman96,carilli97}). The flat-spectrum radio emission arises from a compact core and a one-sided jet to the southwest, with the jet axis aligned (within $\\sim 15^\\circ$) with the major axis of the host galaxy \\citep{polatidis95,fomalont00}. The AGN is heavily reddened ($r$-K~=~5.1) and was not detected in a deep R-band image, suggesting a high level of dust obscuration \\citep{stickel96}. Consistent with this, strong redshifted mm-wave molecular absorption has been detected towards the radio source \\citep{wiklind96}, with two absorption complexes at $z \\sim 0.6734$ (system~A) and $z \\sim 0.6715$ (system~B), close to the redshift of the host galaxy ($z = 0.674 \\pm 0.001$; \\citealt{stickel94}).  Besides the mm-wave transitions, \\hi~21cm, OH~1665~MHz and OH~1667~MHz absorption have all been detected from system~A, with strong, wide profiles extending over a velocity range of $\\gtrsim 100$~\\kms\\ \\citep{wiklind96,carilli97,kanekar02}. In contrast, the mm-wave absorption in system~B is quite narrow [full-width-at-half-maximum (FWHM)$ \\sim 15$~\\kms] and neither \\hi~21cm nor OH absorption have been detected at this redshift \\citep{carilli97,carilli98}. This is the only $z > 0.1$ mm-wave absorber that has not hitherto been detected in OH or \\hi~21cm absorption \\citep{wiklind94,wiklind95,wiklind96,wiklind96b,chengalur99,kanekar02,kanekar03c} and is thus an excellent candidate for a deep search in these transitions. Beside studying physical conditions in the interstellar medium (ISM) of the QSO host, the detection of these lines would, in principle, also allow one to test for changes in the fundamental constants from $z \\sim 0.67$ to the present epoch \\citep{darling03,chengalur03,kanekar04b}. Unfortunately, the OH~1665~MHz line from $z \\sim 0.6715$ lies at the same frequency as the known 1667~MHz absorption from $z \\sim 0.6734$ \\citep{kanekar02}, implying that it (and the latter 1667~MHz line) cannot be used to probe fundamental constant evolution. We report here a search for the other three redshifted OH ground-state lines (at rest frequencies of 1667.3590, 1612.2310 and 1720.5299~MHz) and the \\hi~21cm line towards 1504+377 with the Giant Metrewave Radio Telescope (GMRT) and the Green Bank Telescope (GBT), resulting in the detection of OH~1667~MHz and \\hi~21cm absorption at $z \\sim 0.6715$. ",
        "conclusions": "\\label{sec:discuss}  \\begin{figure} \\centering \\epsfig{file=fig3.eps,height=3.3truein,width=3.3truein} \\caption{Final difference spectrum between the \\hi~21cm optical depth spectra towards 1504+377 in 2003 and 2006, with optical depth difference (in units of $10 \\times \\tau_{diff}$) plotted against heliocentric velocity in \\kms\\, relative to $z = 0.674$. The original difference spectrum had a resolution of $\\sim 0.54$~\\kms; this was boxcar-smoothed to, and resampled at, a resolution of $\\sim 4.8$~\\kms\\ to produce this spectrum.} \\label{fig:variability} \\vskip -0.1in \\end{figure}  \\subsection{Variability in the \\hi~21cm profile} \\label{sec:21cm-change}  Fig.~\\ref{fig:variability} shows a plot of the difference between the \\hi~21cm optical depths in 2003 and 2006 versus heliocentric velocity, in \\kms, relative to $z = 0.674$. The strong features in the difference spectrum between $\\sim -150$~\\kms\\ and $\\sim -70$~\\kms\\ indicate significant changes ($\\sim 10\\%$ of the line depth) in the \\hi~21cm profile between 2003 and 2006. Note that the difference cannot be due to a simple scaling of one or both of the spectra, as different spectral components show changes of opposite sign. While the possibility that the observed change might be due to RFI cannot be ruled out, no evidence was seen for RFI at these frequencies, in either these or our other 850~MHz GBT datasets. The profile ``variability'' is coincident with the strongest spectral components, with the rest of the profile showing no evidence for changes within the noise.  Variability in redshifted \\hi~21cm profiles has been seen earlier in two damped Lyman-$\\alpha$ systems, at $z \\sim 0.524$ towards 0235+164 \\citep{wolfe82} and $z \\sim 0.3127$ towards 1127$-$145 \\citep{kanekar01c}. While changes in the latter two profiles have been detected on far shorter timescales (a few days) than in 1504+377, it is interesting that all three sources contain highly compact ($\\sim $~mas-scale) components. Possible explanations for the observed changes towards 1504+377 include refractive scintillation in the Galactic interstellar medium  (for which the background source need not be compact; \\citealt{macquart05}), or transverse motion of a source component on VLBI scales \\citep{briggs83c}. Both models require small-scale structure in the 21cm optical depth of the absorbing gas.   \\subsection{Physical conditions in the absorbing gas} \\label{sec:1504}  The radio core of 1504+377 and the nucleus of the host galaxy are coincident within the errors ($\\sim 1''$) in the R-band image of \\citet{stickel94}. At mm-wave frequencies, the core dominates the quasar flux density, with very little emission coming from the steep-spectrum jet \\citep{wiklind96}. The core is also likely to be extremely compact at these frequencies, implying that both mm-wave absorbers arise along a single line of sight, which must also pass extremely close to the centre of the host galaxy. \\citet{wiklind96} noted that it is impossible to produce two absorption components at very different velocities in such circumstances if the absorbing gas is in pure rotational motion. The large separation ($\\sim 330$~\\kms) between the two observed absorption velocities is thus suggestive of the presence of strong non-circular orbits; these authors argued in favour of a scenario in which the broad absorption from system~A originates close to the nucleus (in a nuclear ring or a bar), while the narrow absorption of system~A arises in a more-distant cloud in the disk of the host galaxy. In this picture, the systemic redshift is $z \\sim 0.6715$. On the other hand, \\citet{carilli97} used the fact that the optical emission redshift of the host galaxy ($z = 0.674 \\pm 0.001$) is in excellent agreement with that of the higher-redshift complex to argue that the latter is the systemic redshift. They also pointed out that the outflow velocity of system~A in this model ($\\sim 330$~\\kms) is too large for a cloud in the outer disk of the parent galaxy and suggested the possibility that it might arise in a high-velocity cloud, due to tidal interactions between the host galaxy and a nearby object seen in the R-band image of \\citet{stickel94}. Our new GBT \\hi~21cm spectrum of Fig.~\\ref{fig:21cm} clearly shows that the two absorption systems are, in fact, part of a continuous absorption complex, spanning $\\sim 600$~\\kms\\ and extending from the optical redshift of $z \\sim 0.674$ out to $z \\sim 0.6706$. The 21cm absorption lies entirely blueward of the optical redshift, implying that it must arise in gas that is outflowing from the quasar.  The large velocity spread of the \\hi\\ outflow in 1504+377 is similar to that seen in a number of low-redshift AGNs \\citep{morganti05}. These authors note that all known fast \\hi\\ outflows have been detected in radio galaxies in early or re-started phases of their radio activity. There is also evidence that the most likely mechanism to explain such fast \\hi\\ outflows is interaction between the radio jets and the surrounding interstellar medium (e.g. \\citealt{morganti05b}), with rapid cooling taking place in the gas after a jet-cloud interaction, as expected from numerical simulations (e.g. \\citealt{fragile04}). The fact that 1504+377 shows no extended radio structure (the outer jet extends to only $\\sim 55$~mas, i.e. $\\sim 387$~pc, from the nucleus; \\citealt{polatidis95}) suggests that it too is in a early phase of its radio activity. Recent 5~GHz VLBI observations \\citep{bolton06} have found a new north-eastern extension, which was not seen in earlier (deeper) images (e.g. \\citealt{fomalont00}), demonstrating that the source is currently in an active phase. Finally, the fact that the radio structure in 1504+377 is strongly one-sided (e.g. \\citealt{fomalont00}) indicates that the jet lies close to the line of sight towards the core. The above suggestion that jet-cloud interactions are responsible for local gas cooling is consistent with the fact that mm-wave absorption (which takes place in cold gas and, as noted earlier, must arise towards the core) is seen at multiple velocities along the line of sight.  It thus appears that the wide \\hi~21cm and molecular absorption towards 1504+377 arise in outflowing gas from the AGN that is cooling rapidly after an interaction with the south-western radio jet. This is the highest redshift at which such a high-velocity outflow has been observed (e.g. \\citealt{morganti05}) and, perhaps more interesting, the first case where molecular gas has been detected in the outflow. The ${\\rm H_2}$ fraction is $f_{\\rm H_2} = \\left[N_{\\rm H_2}/N_{\\rm HI}\\right] \\le 2 \\times (\\ts/100) (\\tx/10) (f_{\\rm OH}/f_{\\rm 21})$. \\citet{morganti05} assume $\\ts \\sim 1000$~K to estimate \\hi\\ column densities for sources in their sample due to the proximity of the gas to the AGN and the likely presence of shocks. Using this value for consistency gives a molecular fraction of $f_{\\rm H_2} \\sim 0.2$ in the outflowing gas.   We estimate the mass outflow rate ${\\dot M}$ using the model of \\citet{heckman00}, in which a constant-velocity, mass-conserving wind flows into a solid angle $\\Omega$ from a  minimum radius $r_*$, viz.  \\begin{equation} {\\dot M} = 30 \\left[ \\frac{\\Omega}{4\\Pi}\\right] \\left[ \\frac{r_*}{1 \\: {\\rm kpc}} \\right] \\left[ \\frac{N_{\\rm H}}{10^{21} \\: {\\rm cm^{-2}}} \\right] \\left[ \\frac{v}{300 \\:{\\rm km s^{-1}}} \\right] \\: M_\\odot \\: {\\rm yr}^{-1}, \\end{equation} \\noi where $v$ is the outflow velocity and $N_{\\rm H}$, the total hydrogen column density of the outflowing gas. We will assume that the minimum radius $r_*$ is $\\sim 10$~pc, the size of the radio core, and, following \\citet{morganti05}, that $\\Omega = \\Pi$~steradians and $v = \\:{\\rm FWBN}/2 \\sim 300$~\\kms.  The total hydrogen column density at $z \\sim 0.67$ is $N_{\\rm H} = \\left[ N_{\\rm HI} + 2 \\times N_{\\rm H_2} \\right] \\sim 1.6 \\times 10^{23}$~\\cm, again assuming $\\ts \\sim 1000$~K. This leads to an estimated mass outflow rate of ${\\dot M} \\sim 12 M_\\odot$~yr$^{-1}$, comparable to estimates in nearby fast \\hi\\ outflows \\citep{morganti05}.    \\citet{wiklind96} noted that HCO$^+$ is highly over-abundant in system~B, enhanced by at least an order of magnitude relative to expected abundances in chemical models. The ratios of HCO$^+$ to CO and HCO$^+$ to HCN column densities here are $3 - 5$ times larger than in system~A. While such large differences in relative abundances between HCO$^+$ and species such as CO, HCN, etc, have been observed in Galactic clouds \\citep{lucas98}, the ratio of OH to HCO$^+$ column densities in both the Galaxy and a sample of four redshifted HCO$^+$ and OH absorbers has been found to be fairly constant, with $N_{\\rm HCO^+}/N_{\\rm OH} \\sim 0.03$ \\citep{liszt96,kanekar02} over more than two orders of magnitude in HCO$^+$ column density. \\citet{liszt04} found this ratio to show a large spread (by a factor of $\\sim 4$) in the clouds towards Cen.A and NGC1052, with the HCO$^+$ and OH lines also showing very different kinematics, but argued that this could be explained by differing source structure and foreground free-free opacity at the OH and HCO$^+$ frequencies, source variability between observing epochs, and excitation effects at high OH column densities ($\\gtrsim 10^{15}$~\\cm; \\citealt{langevelde95}). Conversely, system~B has $N_{\\rm HCO^+}/N_{\\rm OH} \\sim 0.5 \\times (\\tx/10) (0.46/f_{\\rm OH})$, discrepant by more than an order of magnitude from the expected value. However, 1504+377 is highly compact at both mm-wave and cm-wave frequencies (with a cm-wave core-fraction of $\\sim 46$\\%; \\citealt{carilli97}) and the HCO$^+$ and OH lines have very similar FWHMs [$\\sim 16.5 \\pm 2.2$~\\kms\\ (OH) and $\\sim 15.2 \\pm 0.9$~\\kms\\ (HCO$^+$)], making it likely that they arise from similar gas. Increasing $\\tx$ by an order of magnitude could resolve this problem but such high $\\tx$ values have never been seen in the Galaxy (e.g. \\citealt{liszt96}). Similarly, the  ratio of peak optical depths in the HCO$^+$ and \\hi~21cm lines in system~A is  $R \\equiv \\tau_{HCO^+}/\\tau_{\\rm 21} \\sim 30$, far larger than that seen in Galactic clouds ($0.1 \\le R \\le 6$; \\citealt{lucas96,liszt96}). \\citet{carilli97} point out that high values of $R$ could result from either far warmer \\hi\\ or low molecular dissociation, but this would not explain the discrepancy in the ratio of OH and HCO$^+$ column densities. If the latter is not due to real chemical differences between OH and HCO$^+$ (which seems unlikely; \\citealt{liszt00}), a plausible explanation is extreme small-scale structure in the opacity of the absorbing gas on sub-parsec scales, smaller than the size of the radio core at cm wavelengths. This could arise due to internal shocks or turbulence in the rapidly outflowing gas. As noted earlier, the observed variability in the \\hi~21cm absorption at $z \\sim 0.674$ suggests similar small-scale structure at a different location in the outflow, which could also account for the large velocity difference ($\\sim 15$~\\kms) in peak OH and HCO$^+$ absorption in system~A \\citep{kanekar02}. Monitoring the mm-wave lines for variability would be one way of testing this hypothesis.  Finally, comparisons between the OH, \\hi~21cm and HCO$^+$ redshifts from an absorber can be used to test the evolution of fundamental constants \\citep{darling03,chengalur03}. However, the above possibility of small-scale structure in the absorbing gas makes it likely that any such comparisons in the absorbing gas towards 1504+377 will be dominated by local systematic velocity offsets. We conclude that this absorber is unlikely to be useful for the purpose of probing fundamental constant evolution."
    },
    "0710/0710.3189_arXiv.txt": {
        "abstract": "We present a simple and efficient method to set up spherical structure models for $N$-body simulations with a multimass technique. This technique reduces by a substantial factor the computer run time needed in order to resolve a given scale as compared to single-mass models. It therefore allows to resolve smaller scales in $N$-body simulations for a given computer run time. Here, we present several models with an effective resolution of up to $1.68 \\times 10^9$ particles within their virial radius which are stable over cosmologically relevant time-scales. As an application, we confirm the theoretical prediction by \\citet{2005MNRAS.360..892D} that in mergers of collisonless structures like dark matter haloes always the cusp of the steepest progenitor is preserved. We model each merger progenitor with an effective number of particles of approximately $10^{8}$ particles. We also find that in a core-core merger the central density approximately doubles whereas in the cusp-cusp case the central density only increases by approximately 50\\%. This may suggest that the central region of flat structures are better protected and get less energy input through the merger process. ",
        "introduction": "Resolution is a key issue in $N$-body simulations. In state-of-the-art cosmological $N$-body simulations today, structures can be fully resolved down to the scale of a fraction of a per cent of the virial radius \\citep{2007ApJ...657..262D}. But there are many problems where even higher resolution in central regions of structures is needed. Also, one often would like to study a certain problem without the cosmological context, i.e. one needs a possibility to set up isolated structures with high resolution in order to study dynamical effects in detail and with precise control of the initial condition, which can be very difficult within the framework of a cosmological $N$-body simulation.  For example, the question of whether the central dark matter density profile of haloes that form in cosmological $N$-body simulations is cuspy or cored needs at least a resolution of $\\approx 10^{-3}~\\rvir$ to be answered. Another example are centrally flat profiles: in order to resolve isolated structures with flat central profiles, a lot of particles are needed since in flat profiles the resolution scaling with the number of particles is the slowest (see below for more details about the scaling of resolution with the number of particles). Another problem we had in mind when developing the multimass technique presented in this paper, was the dynamics of super-massive black hole binaries in the centre of remnants of galaxy mergers. There, the relevant scales are of order of a few pc $\\approx 10^{-6}-10^{-5}~\\rvir$ for Milky Way size dark matter haloes.  \\begin{figure*} \\centering \\includegraphics[width=0.49\\textwidth]{g10.res.eps} \\includegraphics[width=0.49\\textwidth]{g15.res.eps} \\caption{Plot of $r_1$, $r_{100}$, $\\rrel(\\tdyn(\\rvir)/10)$, $\\rrel(\\tdyn(\\rvir))$ as a function of $\\Nvir$ for a spherical structure with $\\cvir = 10$, $\\alpha = 1$, $\\beta = 3$ but different central profiles: $\\gamma = 1$ (left) and $\\gamma = 3/2$ (right). One can nicely see that in general the constraint on $\\rres$ is set by the scale where relaxation becomes important since $\\rres = \\max(r_{100}(\\Nvir),\\rrel(t_0,\\Nvir))$ for a given simulation time $t_0$. The asymptotic scaling for $r_N$ and $\\rrel$ is given in the plots; the discrepancy between the sampled scale $r_{100}$ and the relaxation scale $\\rrel$ as a function of $\\Nvir$ is bigger for larger $\\gamma$.} \\label{fig:resolution} \\end{figure*}  We illustrate the resolution problem in more detail with a commonly used family of spherically symmetric density profiles: the so called $\\alpha\\beta\\gamma$-models family \\citep{1996MNRAS.278..488Z}. An $\\alpha\\beta\\gamma$-model density profile is given by \\begin{equation} \\label{eq:rho} \\rho(r) = \\frac{\\rho_0}{\\left(\\frac{r}{\\rs}\\right)^\\gamma \\left(1+\\left(\\frac{r}{\\rs}\\right)^\\alpha \\right)^{\\left(\\frac{\\beta-\\gamma}{\\alpha}\\right)}}~, \\end{equation} where $\\gamma$ determines the inner slope and $\\beta$ the outer slope of the density profile, whereas $\\alpha$ controls the transition between the inner and outer region. The normalisation is given by $\\rho_0$ and $\\rs$ is the scale radius defined by $\\rs \\equiv \\rvir/\\cvir$ ($\\cvir$ is the virial concentration). Many models used by observers as well as theorists to describe structures in the universe belong to this family, e.g. the Hernquist profile \\citep{1990ApJ...356..359H}, the NFW profile \\citep{1996ApJ...462..563N} or the Moore profile \\citep{1998ApJ...499L...5M} are just special cases of the general form (\\ref{eq:rho}).  An obvious minimal criterion for a given length scale to be resolved is, that it has to be populated with enough particles by the sampling procedure. For example, if one would like to set up an isolated structure given by the density profile (\\ref{eq:rho}), then the radius $r_{N}$ that contains $N$ particles is simply given by the solution of \\begin{equation} \\label{eq:rNdef} M(r_{N}) = N m~, \\end{equation} where $M(r)$ is the enclosed mass given by \\begin{equation} \\label{eq:Menc} M(r) \\equiv \\int_0^r 4 \\pi \\rho(x) x^2 \\d x~, \\end{equation} and $m$ is the mass of the particles which can be expressed in virial quantities for single-mass models as $m = \\Mvir / \\Nvir$. In the central asymptotic regime (i.e. $r_{N} \\ll \\rs$), equation (\\ref{eq:rNdef}) can be solved for $r_{N}$ to yield \\begin{equation} \\label{eq:rN} r_{N} = \\left(\\frac{3-\\gamma}{4 \\pi} \\frac{N m}{\\rho_0 \\rs^{\\gamma}}\\right)^{\\frac{1}{3-\\gamma}}~. \\end{equation} For $N = 1$ one obtains the scale of the location of the innermost particle, $\\rimp = r_1$.\\footnote{An alternative definition for the location of the innermost particle can be made with the mean particle separation given by $h(r) \\equiv \\sqrt[3]{\\frac{m}{\\rho(r)}}$. The two definitions are essentially equivalent and only differ by the geometric factor $\\left(\\frac{4 \\pi}{3-\\gamma}\\right)^{\\frac{1}{3-\\gamma}}$ which is of order unity for most useful profiles.} But of course one particle is not enough to resolve that scale. If we say at least 100 particles are needed in the innermost bin so that the bin is well resolved, then $r_{100}$ provides a better estimate of the resolved scale.  Another constraint to resolution comes from the fact that often the structures that one would like to simulate can be very well approximated as collisionless systems, i.e. the local relaxation time-scale is much longer than the age of the system. This property must be preserved in $N$-body simulations of such collisionless systems. Otherwise, purely artificial relaxation processes due to under-resolving the structures will have a dominant effect on the dynamics of the system \\citep{2003MNRAS.338...14P,2004MNRAS.348..977D,2004MNRAS.349.1117B}. This is not a numerical artefact since a real astrophysical system with a small particle number will also be subject to relaxation processes. The problem is, that in most cases, we can not simulate the astrophysical system under study with the number of particles the system would have in nature, e.g. a galaxy size dark matter halo would have of the order of $O(10^{67})$ dark matter particles if we assume that the dark matter particle has a mass of 100 GeV/c$^2$. Therefore, this effect due to under-resolving the system with not enough particles will always be a limitation of collisionless $N$-body simulations.  In principle, there are only few known dynamically stable systems in the universe, e.g. the Keplerian 2-body system or the figure-8 rotation of 3 bodies \\citep{1993PhRvL..70.3675M, 2000AM....152..881C}.\\footnote{This is only true in the Newtonian regime. In General Relativity, not even a 2-body orbit is dynamically stable and the orbit decays due to the emission of gravitational radiation \\citep{1979RvMP...51..447D, 1977QJRAS..18....3I, 1997RvMP...69..337A}.} Dynamical effects like relaxation or evaporation will sooner or later lead to the disruption of any system. The issue is always on what time-scale the system is actually stable. The time-scale of interest is normally of the order of the age of the universe and for most cases collisionless astrophysical systems can be regarded as perfectly stable within that time-scale. This is different in $N$-body simulations. Here, one just tries to generate models that show the desired stability behaviour during the time-scale of interest with the least amount of particles needed. Artificial $N$-body models will be subject to such disruption effects much sooner than the real astrophysical system one wants to study.  We illustrate this in more details now. The local relaxation time is defined by \\begin{equation} \\label{eq:trelax} \\trel(r) \\equiv \\frac{N(r)}{\\ln(N(r))}~\\tdyn(r)~, \\end{equation} where \\begin{equation} \\label{eq:tdyn} \\tdyn(r) \\equiv 2 \\pi \\sqrt{\\frac{r^3}{G M(r)}} \\end{equation} is the dynamical or orbital time at radius $r$ and $N(r) \\equiv M(r)/m$ denotes the number of particles within $r$. Here, we have a slightly different normalisation than in the usual expression for the local relaxation time since we dropped the factor of 8 in front of the logarithmic term in the denominator. We found better agreement with results from $N$-body simulations with this normalisation (see section \\ref{chap:tests} below for more details). Had we kept the factor of 8, then our definition (\\ref{eq:trelax}) for spherically symmetric structures would be equivalent to the empirical expression found by \\citet{2003MNRAS.338...14P}. Normally, the Coulomb logarithm is given by $\\ln(\\Lambda) = \\ln(b_{\\mathrm{max}}/b_{\\mathrm{min}})$, where $b_{\\mathrm{max}}$ and $b_{\\mathrm{min}}$ are the maximum and minimum impact parameters of the particles under consideration. Since the minimum impact parameter is related to the softening length and the latter scales with the number of particles, we prefer the direct formulation of the Coulomb logarithm as a function of the number of particles, $\\ln(N(r))$.  In a simulation run for a time $t_0$ relaxation processes become important on a scale $\\rrel$ given by the solution of \\begin{equation} \\label{eq:rrelaxdef} \\trel(\\rrel) = t_0~. \\end{equation} In the central asymptotic regime (i.e. $\\rrel \\ll \\rs$), equation (\\ref{eq:rrelaxdef}) can be inverted for $\\rrel$ to yield \\begin{equation} \\label{eq:rrelax} \\rrel(t_0) = \\left( \\frac{W_{-1}(X)~\\frac{(3 - \\gamma) m}{4 \\pi \\rho_0 \\rs^{\\gamma}}} {-\\frac{1}{2} \\frac{6-\\gamma}{3-\\gamma} \\frac{\\pi}{t_0} \\left( \\frac{3-\\gamma}{G \\pi \\rho_0 \\rs^{\\gamma}} \\right)^{\\frac{1}{2}}} \\right)^{\\frac{2}{6-\\gamma}} \\end{equation} where $W_{-1}$ denotes the $k = -1$ branch of the Lambert W function \\citep{1758AH......3..128L,1996ACM.....5..329C} and \\begin{equation} X \\equiv -\\frac{1}{2} \\frac{6-\\gamma}{3-\\gamma} \\frac{\\pi}{t_0} \\left(\\frac{3-\\gamma}{G \\pi \\rho_0 \\rs^{\\gamma}}\\right)^{\\frac{1}{2}} \\left( \\frac{(3 - \\gamma) m}{4 \\pi \\rho_0 \\rs^{\\gamma}} \\right)^{\\frac{\\gamma}{2 (3-\\gamma)}}~. \\end{equation}  In Fig. \\ref{fig:resolution} we plot $r_1$, $r_{100}$, $\\rrel(\\tdyn(\\rvir)/10)$ and $\\rrel(\\tdyn(\\rvir))$ as a function of $\\Nvir$ by setting $m = \\Mvir / \\Nvir$. The virial radius $\\rvir$ is defined so that the enclosed average density within $\\rvir$ is given by \\begin{equation} \\frac{\\Mvir}{4 \\pi \\rvir^3/3} = \\rhovir = \\Deltavir \\rho_{\\mathrm{crit},0} = 1.41\\times10^{4}~\\Mo~\\kpc^{-3} \\end{equation} where $\\Deltavir = 178~\\Omega_{\\mathrm{M},0}^{0.45} \\approx 104$ \\citep{1998ApJ...503..569E} for our choice of cosmology with $\\Omega_{\\mathrm{M},0} = 0.3$, $\\Omega_{\\Lambda,0} = 0.7$, $\\rho_{\\mathrm{crit},0} = {3 H_0^2}/{8 \\pi G}$ and $H_0 = 70~\\km~\\s^{-1}~\\Mpc^{-1}~(h_0 = 0.7)$, which we use throughout this paper. The dynamical time at the virial radius is then \\begin{equation} \\tdyn(\\rvir) = \\sqrt{\\frac{3 \\pi}{G \\Deltavir \\rho_{\\mathrm{crit},0}}} = 12.2~\\Gyr~. \\end{equation} For example for a galaxy size dark matter halo with $\\Mvir = 10^{12}/h_0~\\Mo = 1.43 \\times 10^{12}~\\Mo$ one would obtain $\\rvir = 289~\\kpc$ and the values for systems of different virial mass $\\Mvir$ can be obtained by the simple scaling relation \\begin{equation} \\rvir = 289~\\kpc \\sqrt[3]{\\frac{\\Mvir}{1.43 \\times 10^{12}~\\Mo}}~. \\end{equation}  Although the virial radius is a rather artificial definition of the size of a dark matter halo and the virialised region of haloes in cosmological $N$-body simulations is generally much larger \\citep{2006ApJ...645.1001P,2007ApJ...667..859D} it is a convenient normalisation and cut-off scale for isolated models since we are anyway mainly interested in the central dynamics of the structure. We set $\\cvir = 10$, $\\alpha = 1$, $\\beta = 3$ in both cases and present plots for $\\gamma = 1$ (left) respectively $\\gamma = 3/2$ (right).  We only plot these quantities for $\\Nvir \\geq 10^{6}$ since the expressions (\\ref{eq:rN}) and (\\ref{eq:rrelax}) are only valid in the central asymptotic regime. One can see that $r_N \\propto \\Nvir^{-\\frac{1}{3-\\gamma}}$ and $\\rrel \\sim \\Nvir^{-\\frac{2}{6-\\gamma}}$ since the Lambert W function only varies very slowly as a function of $\\Nvir$ which reflects the weak dependence of the logarithmic term of the local relaxation time. For a given particle resolution $\\Nvir$ and simulation time $t_0$, the radius $\\rres$ that we can still resolve with correct collisionless physics (i.e. this scale does not suffer from too much artificial relaxation) is given by \\begin{equation} \\rres = \\max(r_{100}(\\Nvir),\\rrel(t_0,\\Nvir))~. \\end{equation} As one can see from Fig. \\ref{fig:resolution}, this resolution scale $\\rres$ is in general set by $\\rrel(t_0,\\Nvir)$ for isolated high resolution structures. It is worth remarking here, that for structures assembled hierarchically in a cosmological $N$-body simulation, the amount of relaxation is significantly larger and $\\rrel$ scales slower as a function of $\\Nvir$ \\citep{2004MNRAS.348..977D}. Although a structure might be sampled with enough particles at a certain scale at the final time, the relaxation time at that scale would have been much smaller in the past since particles were in lower mass structures during the hierarchical growth.  By inspecting Fig. \\ref{fig:resolution}, we see that more than approximately of order $O(10^{12})$ particles in the centre of a structure with $\\gamma = 1$ are needed in order to resolve scales of $\\approx 10^{-5}~\\rvir$. It is worth remarking, that with the same number of particles $\\Nvir$ much smaller scales are populated in the $\\gamma = 3/2$ profile than in the $\\gamma = 1$ profile. Generally, the steeper the central profile, the more the particles are concentrated. But unfortunately, the relaxation scale $\\rrel$ does not scale equally fast so that the discrepancy as a function of $\\Nvir$ becomes bigger for larger values of $\\gamma$.  Nonetheless, such an enormous amount of particles per structure is hardly doable today - even with large supercomputers. But since we only need this high resolution at the very centre of the structure, our solution to this problem is to use models where we only populate regions of the phase space that are in the centre or will reach the centre in the future with high resolution particles.  In section \\ref{chap:method} we present the simple idea behind the multimass models and present stability tests in section \\ref{chap:tests}. In section \\ref{chap:preservation}, we test Dehnen's prediction with high resolution mergers and we summarize our results in section \\ref{chap:conclusions}. ",
        "conclusions": "\\label{chap:conclusions}  The multimass technique for modelling haloes is a simple method to perform high resolution $N$-body simulations. An early version of this technique without orbit dependent refinement has already been used successfully in several applications which also include cosmological structure formation simulations \\citep{2005MNRAS.364..665D,2006MNRAS.368.1073G}. With a careful choice of parameters it is possible to gain over an order of magnitude in computer run time for a given resolution scale.  As an application of this technique we confirm the earlier findings that core-core mergers lead to a cored merger remnant while cups-cusp mergers lead to a cuspy merger remnant with high resolution multimass $N$-body simulations. In cusp-core merges, the merger remnant has a final profile corresponding to the steepest progenitor which is in excellent agreement with the theoretical predictions. We find that in the core-core case the central density approximately doubles whereas in the cusp-cusp case the central density in the merger remnant only increases by approximately 50\\%. This may suggest that the central region of flat structures are better protected and get less energy input from the merger.  A software tool called \\textsc{halogen} (HALO GENerator) for generating multimass initial conditions is available from the author upon request."
    },
    "0710/0710.4556_arXiv.txt": {
        "abstract": "It has been known for over 30 years that Galactic globular clusters (GCs) are overabundant by orders of magnitude in bright X-ray sources per unit mass relative to the disk population.  Recently a quantitative understanding of this phenomenon has developed, with a clear correlation between the number of X-ray sources in a cluster, $N_X$, and the cluster's encounter frequency, $\\Gamma$, becoming apparent.  We derive a refined version of $\\Gamma$ that incorporates the finite lifetime of X-ray sources and the dynamical evolution of clusters.  With it we find we are able to explain the few clusters that lie off the $N_X$--$\\Gamma$ correlation, and resolve the discrepancy between observed GC core radii and the values predicted by theory. Our results suggest that most GCs are still in the process of core contraction and have not yet reached the thermal equilibrium phase driven by binary scattering interactions. ",
        "introduction": "\\label{sec:dyn} It was realized more than 30 years ago that Galactic globular clusters (GCs) are overabundant by orders of magnitude in bright X-ray sources per unit mass relative to the disk population \\citep{1975ApJ...199L.143C,1975Natur.253..698K}. It was quickly understood that strong dynamical scattering interactions of binaries in the dense cluster cores should be responsible for this overabundance \\citep{1987IAUS..125..187V}.  With the advances in X-ray astronomy made possible by observatories such as {\\em Chandra}, the relationship between X-ray sources and core cluster dynamics has recently been quantitatively studied.  \\citet{2003ApJ...591L.131P} performed {\\em Chandra} observations of many Galactic GCs down to a limiting luminosity of $4 \\times 10^{30}\\,{\\rm erg}/{\\rm s}$ in the 0.5--6 keV range (which includes low-mass X-ray binaries [LMXBs] in outburst and quiescence, cataclysmic variables [CVs], millisecond pulsars [MSPs], and magnetically active main sequence binaries [ABs]), and looked for correlations between the number of X-ray sources in each cluster and properties of the cluster itself.  They found the strongest correlation with the ``encounter frequency'' $\\Gamma$, a rough estimate of the current dynamical encounter rate in the cluster.  More recently, \\citet{2003ApJ...598..501H} and \\citet{2006ApJ...646L.143P} have isolated the quiescent LMXBs (qLMXBs) and CVs, respectively, from the X-ray source populations, and have shown that their numbers are indeed consistent with dynamical formation.  These results represent quantitative, empirical evidence that dynamical encounters are responsible for the formation of X-ray sources in clusters.  However, they suffer from at least a few drawbacks. First, the correlation between the number of X-ray sources, $N_X$, and the encounter frequency appears to be sub-linear, with $N_X \\propto \\Gamma^{0.74 \\pm 0.36}$, although for LMXBs the exponent is $0.97 \\pm 0.5$ \\citep{2003ApJ...591L.131P}. Second, there are three clusters for which $N_X$ is significantly larger than predicted by $\\Gamma$.  In the original work of \\citet{2003ApJ...591L.131P} it was already clear that NGC 6397 has an $N_X$ that is $\\sim 5$ times larger than predicted by the $N_X$--$\\Gamma$ correlation. Recent observations show that $N_X$ is factor of $\\sim 2$ times that predicted by $\\Gamma$ for NGC 7099 \\citep{2007ApJ...657..286L}, and a factor of $\\sim 20$ for Ter 1 \\citep{2006MNRAS.369..407C}. The common thread among these three clusters is that they are observationally ``core-collapsed,'' while all others in the \\citet{2003ApJ...591L.131P} sample are not \\citep{2006AdSpR..38.2923G}. (A possible exception is NGC 6752, whose collapsed core status is debated \\citep{2003ApJ...595..179F,1995ApJ...439..191L}.)  A cluster is observationally termed core-collapsed if its surface brightness profile is consistent with a cusp at the limit of resolution, making it more difficult to measure the core radius.   As described below, the collapsed core status of a cluster is linked to its dynamical state, implying that cluster evolution complicates the $N_X$--$\\Gamma$ correlation.  \\begin{figure} \\begin{center} \\includegraphics[width=0.9\\columnwidth]{f1.eps} \\caption{Number of observed cluster X-ray sources with $L_X\\gtrsim 4 \\times 10^{30}\\,{\\rm erg}/{\\rm s}$ for several Galactic GCs versus the encounter rate $\\Gamma$. The power-law fit and data points are from \\citet{2003ApJ...591L.131P} with the exception of NGC 7099 \\citep{2007ApJ...657..286L} and Ter 1 \\citep{2006MNRAS.369..407C}. The $N_X$ error bars for Ter 1, NGC 6397, and NGC 7099 represent source counting noise and background source uncertainty, but for the remaining clusters represent only background uncertainty. \\label{fig:gamma}} \\end{center} \\end{figure} ",
        "conclusions": "\\label{sec:disc}  This {\\em letter} presents the confluence of three suggestive observational and theoretical results into a self-consistent picture.  The first is that of the clusters that have been observed sufficiently to determine their XRB population, the three that are core-collapsed are the same three that have a significant X-ray source overabundance (of a factor of $\\sim 2$ to $\\sim 20$). The second is the semi-analytical result derived in this {\\em letter} that a cluster in the binary burning phase for the last few Gyr should have $\\sim 5$ times more dynamically formed X-ray sources than if it were in the core contraction phase for the same time.  The third is the recently confirmed discrepancy between observations and theory for the core radii of Galactic GCs, which suggests that only the observationally core-collapsed clusters are in the binary-burning phase. In light of these facts, the conclusion that seems strongly suggested is that most Galactic GCs are currently still in the core contraction phase, while only the $\\sim 20\\%$ of clusters that are core-collapsed are in the binary burning phase.  This goes counter to the widely held belief that most clusters are currently in the binary burning phase, and complicates the many existing studies that have assumed cluster core properties that are constant with time.  The implications of this result are manifold.  There are many studies of the dynamical production of interesting source populations in clusters which assume core properties that are constant with time.  These include predictions of the formation of blue stragglers \\citep[e.g.,][]{2004ApJ...605L..29M}, the evolution of the core binary fraction \\citep{2005MNRAS.358..572I}, and tidal-capture binaries \\citep[e.g.,][]{1994ApJ...423..274D}, among others. Revising the results may be as simple as scaling predicted source numbers, but may not be so simple for other quantities.  Studies of GC evolution have shown that clusters starting from very different initial conditions evolve toward a common range of values in many observable structural parameters in the binary-burning phase, including the concentration and ratio of core to half-mass radius \\citep{2003ApJ...593..772F,2006MNRAS.368..677H,2007ApJ...658.1047F}. Since most clusters may not be in the binary-burning phase after all, their observed properties are likely to be more strongly correlated with their initial conditions.  This makes modeling of clusters a bit more complicated, but on the other hand allows one to more readily deduce something about the initial properties of clusters.  Perhaps anticlimactically, our results suggest that the alternative energy sources recently proposed for supporting GC cores are not required.  These include the suggestion of IMBHs in {\\em many} Galactic GCs \\citep{2006astro.ph.12040T}, enhanced stellar mass loss from stellar evolution of physical collision products \\citep{chatterjeeposter}, mass segregation of compact remnants in young clusters \\citep{2004ApJ...608L..25M}, or evaporation of the stellar-mass black hole subsystem in young clusters \\citep{2007MNRAS.379L..40M}.  Although the picture painted in this {\\em letter} is a suggestive one, there are still several caveats and limitations to our analysis.  In the derivation of our refined $\\Gamma$ we assume that the core binary fraction and abundance of compact objects are constant over the time of integration. Neither is strictly true, although we expect they will not vary enough to significantly change the overabundance value we derive. We have used only one possible expression for the evolution of the core radius in core contraction.  While we expect the general behavior to be very similar to what we have assumed here, more work needs to be done to determine if it is universal.  Additionally, our analysis ignores the effect of Galactic tidal stripping on cluster mass, which would make a cluster appear overabundant in X-ray sources, and may be relevant for NGC 6397 \\citep{2003ApJ...591L.131P}. On the observational side, there are some uncertainties in evaluating $\\Gamma$, which is dependent on quantities that are somewhat difficult to measure for core-collapsed clusters."
    },
    "0710/0710.0375_arXiv.txt": {
        "abstract": "\\spider\\ is a long-duration, balloon-borne polarimeter designed to measure large scale Cosmic Microwave Background (CMB) polarization with very high sensitivity and control of systematics.  The instrument will map over half the sky with degree angular resolution in I, Q and U Stokes parameters, in four frequency bands from $96$ to $275$ GHz.  \\spider's ultimate goal is to detect the primordial gravity wave signal imprinted on the CMB B-mode polarization.  One of the challenges in achieving this goal is the minimization of the contamination of B-modes by systematic effects. This paper explores a number of instrument systematics and observing strategies in order to optimize B-mode sensitivity.  This is done by injecting realistic-amplitude, time-varying systematics in a set of simulated time-streams. Tests of the impact of detector noise characteristics, pointing jitter, payload pendulations, polarization angle offsets, beam systematics and receiver gain drifts are shown. \\spider's default observing strategy is to spin continuously in azimuth, with polarization modulation achieved by either a rapidly spinning half-wave plate or a rapidly spinning gondola and a slowly stepped half-wave plate.  Although the latter is more susceptible to systematics, results shown here indicate that either mode of operation can be used by \\spider. ",
        "introduction": "\\label{sec:intro}  In the past decade, a wealth of data have pointed to a ``standard model'' of the Universe, composed of $\\sim 5\\%$ ordinary matter, $\\sim 22\\%$ dark matter and $\\sim 73\\%$ dark energy in a flat geometry.  The flatness of the Universe, the near isotropy of the CMB, and the nearly-scale-invariant nature of the primordial scalar perturbations from which structure grew support the existence of an early accelerating phase dubbed ``inflation''. A necessary by-product of inflation is tensor perturbations from quantum fluctuations in gravity waves.  A detection of this Cosmological Gravity-Wave Background (CGB) would give strong evidence of an inflationary period and determine its energy scale, while a powerful upper limit would point to more radical inflationary scenarios, e.g., involving string theory, or some non-inflationary explanation of the observations.  The CGB imprints a unique signal in the curl-like, or B-mode, component of the polarization of the CMB; detection of a B-mode signal can be used to infer the presence of a CGB at the time of decoupling.  Direct detection of the gravity waves is many decades off; an advanced Big Bang Observer successor to LISA has been suggested as a way to achieve this \\citep{BBO1,BBO2}. Thus a measurement of the primordial B-modes is the only feasible near-term way to detect the CGB and have a new window to the physics of the early Universe \\citep{Bock:2006}.  A CGB with a potentially measurable amplitude is a by-product of the simplest models of single field inflation which can reproduce the scalar spectral tilt observed in current combined CMB data~\\citep{spergel2007, MacTavish:2005yk}. Examples are chaotic inflation from power law inflaton potentials~\\citep{linde:1983,linde:2005} or natural inflation from cosine inflaton potentials involving angular (axionic) degrees of freedom~\\citep{adams93}. The amplitude is usually parameterized in terms of the ratio of the tensor power spectrum to the scalar power spectrum, $r={\\cal P}_t (k_p)/{\\cal P}_s (k_p)$, evaluated at a comoving wavenumber pivot $k_p$, typically taken to be $0.002 \\ {\\rm Mpc}^{-1}$. Chaotic inflation predicts $r\\approx 0.13$ for a $\\phi^2$ potential and $r \\approx 0.26$ for a $\\phi^4$ potential, and natural inflation predicts $r \\approx 0.02-0.05$.  The potential energy $V$ driving inflation is related to $r$ by $V \\approx (10^{16} \\ {\\rm Gev})^4 r/0.1$. Low energy inflation models have low or negligible amplitudes for the CGB. To get the observed scalar slope and yet small $r$ requires special tuning of the potential. These are often more complicated, multiple field models, e.g., \\citep{linde:1993}, or string-inspired brane or moduli models~\\citep{kallosh2007}.  Given the collection of models it is difficult to predict a precise range for the expected tensor level and the prior probability for $r$ should be considered as wide open.  Recent CMB data have reached the sensitivity level required to constrain the amplitude of and possibly characterize the gradient-like, or E-mode, component of the polarization \\citep{Kovac:2002fg,hedman2002,Readhead:2004xg,Montroy:2006,page2007,quad2007}.  A significant complication of the measurements is that the E-mode amplitude is an order of magnitude lower than the total intensity. In addition, galactic foregrounds such as synchrotron and dust are expected to be significant at these amplitudes~\\citep{kogut2007}. Furthermore the polarization properties of foregrounds are largely unknown.  Constraining B-modes presents an even greater challenge as it is a near certainty that polarized foregrounds will dominate the signal.  The next generation of CMB experiments will benefit from a revolution in detector fabrication in the form of arrays of antenna-coupled bolometers~\\citep{Goldin:2002,kuo2006}. The antenna-coupled design is entirely photo-lithographically fabricated, greatly simplifying detector production.  In addition, the densely populated antennas allow a very efficient use of the focal plane area.  \\spider\\ will make use of this technological advance, in the form of 2624 polarization sensitive detectors observing in four frequency bands from $96$ to $275$ GHz.  A multi-frequency observing strategy is a necessary requirement to allow for a subtraction of the foreground signal. \\spider\\ will observe over a large fraction of the sky at degree scale resolution producing high signal-to-noise polarization maps of the foregrounds at each frequency.  Extraordinarily precise understanding of systematic effects within the telescope will be required to measure the tiny B-mode signal. This {\\it paper} presents a detailed investigation of experimental effects which may impact \\spider's measurement of B-modes. The strategy is to simulate a \\spider\\ flight time-stream injecting systematic effects in the time domain. The aim is to determine the level of B-mode contamination at subsequent stages of the analysis.  With these results one can set stringent requirements on experimental design criteria, in addition to optimizing the telescope's observing strategy.  The outline of this paper is as follows. Section~\\ref{sec:instrument} gives an overview of the instrument, flight and observing strategy. Section~\\ref{sec:method} describes the details of the simulation methodology.  Results for several systematic effects are presented in Section~\\ref{sec:results}.  Section~\\ref{sec:conclusions} concludes with a summary and discussion of the results. ",
        "conclusions": "\\label{sec:conclusions}  The results from Section~\\ref{sec:results} are summarized in terms of experimental specifications in Table~\\ref{tab:summary}.  While results are \\spider-specific the order of magnitude of various effects can be translated to other CMB polarization experiments.  The RMS B-mode signal for $r = 0.1$ is roughly 10000 $\\rm nK^2$ ($\\sim$1436 $\\rm nK^2$ at $\\ell = 8$ and $\\sim$7379 $\\rm nK^2$ at $\\ell = 80$), and scales linearly with r.  Experimental specifications are set by limiting the allowed systematic residual level to a factor of $\\sim10$ smaller than the B-mode signal for $r = 0.01$.  While the simulations were signal-only the impact of large low frequency detector noise (1/f noise) is reflected in the large scale degradation of the B-mode signal for the stepped half-wave plate mode of operation.  Rapid, continuous half-wave plate modulation mitigates this effect entirely.  Rapid, continuous gondola rotation also works but only with iterative map-making which accurately recovers the larger scale signal.  It is important to note that the effects studied in Sections~\\ref{sec:modes} and Sections~\\ref{sec:noise} (naive versus iterated maps, spinning more slowly, stepping half-wave plate versus spinning half-wave plate) will degrade the signal-to-noise achieved on the bandpowers.  These effects differ from the systematics studied in Section~\\ref{sec:point} to Section~\\ref{sec:drift} (pointing reconstruction errors, polarization angle uncertainty, uncorrected ghosting, uncorrected gain drifts) which will ultimately bias the final result.  The requirements on the biasing effects are more difficult treat than the signal-to-noise issues.  The impact of systematics on B-mode polarimeter experiments is also discussed in~\\cite{hhz:2003} and more recently ~\\cite{odea:2007}, where analytical methods are used for calculating the B-mode spectrum bias. The results are useful for setting experimental ``benchmark parameters'' at the very earliest phases of instrument design. This work goes a step further by considering the impact of systematics in the map/time-domain; A necessary step in the evolution of an experiment which aims to measure the tiny primordial, gravity wave signal.  \\begin{table*} \\centering \\begin{tabular}{|c|c|c|} \\hline\\hline \\space {\\bf Systematic} & {\\bf Experimental Spec.} & {\\bf Comments} \\\\ \\hline\\hline Receiver & &  for 110dps\\\\ 1/f knee & $< 200$ mHz & gondola spin \\\\ \\hline Receiver & &  for 36dps\\\\ 1/f knee & $< 100$ mHz & gondola spin \\\\ \\hline Pointing  & & sufficient for\\\\ Jitter & $< 10'$ & $\\ell < 50$ \\\\ \\hline Absolute&   &  \\\\ Pol. Angle Offset & $< 0.25$ deg.  &  \\\\ \\hline Relative&  & \\\\ Pol. Angle Offset & $< 1$ deg. &  \\\\ \\hline Knowledge of &   &  sufficient for\\\\ Beam Centroids & $< 1'$ & $\\ell < 30$ \\\\ \\hline Optical &  &  \\% TOD\\\\ Ghosting & $< 2\\%$ & contamination\\\\ \\hline Calibration &  &  \\\\ Drift & $< 3.0\\%$ & in phase \\\\ \\hline Calibration &  &  \\\\ Drift & $< 0.1\\%$ & out of phase \\\\ \\hline\\hline \\end{tabular} \\caption{\\rm Summary of experimental specifications based on simulation results. Realistic-amplitude, time-varying systematics are injected in the simulated time streams.  Maps are reconstructed without any attempt to correct for the systematic errors.  Experimental specifications are set by limiting the allowed systematic residual level to a factor of $\\sim10$ smaller than the B-mode signal for $r = 0.01$.  The nominal operating mode is a 36 dps gondola spin rate, with the half-wave plate stepping $22.5^{\\circ}$ once per day, with 10 iterations (sufficient to recover the residual levels of the continuously-rotating half-wave plate case) of the map-maker, a Jacobi iterative solver~\\citep{Jones:2007}. }\\label{tab:summary} \\end{table*}"
    },
    "0710/0710.5285_arXiv.txt": {
        "abstract": "Gamma-ray burst afterglow observations in the \\emph{Swift} era have a perceived lack of achromatic jet breaks compared to the \\emph{BeppoSAX} or pre-\\emph{Swift} era. Specifically, relatively few breaks, consistent with jet breaks, are observed in the X-ray light curves of these bursts. If these breaks are truly missing, it has serious consequences for the interpretation of GRB jet collimation and energy requirements, and the use of GRBs as cosmological tools. Here we address the issue of X-ray breaks that are possibly `hidden' and hence the light curves are misinterpreted as being single power laws. We do so by synthesising XRT light curves and fitting both single and broken power laws, and comparing the relative goodness of each fit via Monte Carlo analysis. Even with the well sampled light curves of the \\emph{Swift} era, these breaks may be left misidentified, hence caution is required when making definite statements on the absence of achromatic breaks. ",
        "introduction": "\\label{section:introduction}  The afterglow emission of Gamma-Ray Bursts (GRBs) is well described by the blast wave, or fireball, model \\citep{rees1992:mnras258,meszaros1998:apj499}. This model details the temporal and spectral behaviour of the emission that is created by external shocks when a collimated ultra-relativistic jet ploughs into the circumburst medium, driving a blast wave ahead of it. The level of collimation, or jet opening angle, has important implications for the energetics of the underlying physical process and the possible use of GRBs as standard candles. The signature of the collimation, according to simple analytical models, is an achromatic temporal steepening  or `jet break' at $\\sim 1$ day in an otherwise decaying, power law light curve; from the time of this break, the jet opening angle can be estimated \\citep{rhoads1997:ApJ487}.   Since the launch of the \\emph{Swift} satellite \\citep{gehrels2004:ApJ611}, this standard picture of afterglows has been called into question by the lack of observed achromatic temporal breaks, up to weeks in a few bursts (e.g., \\citealt{panaitescu2006:MNRAS369,burrows2007:astro.ph2633}). In some afterglows, a break is unobserved in both the X-ray and optical light curves, while in other bursts a break is observed in one regime but not the other (e.g., \\citealt{liang2007:arXiv0708}). In the \\emph{BeppoSAX} era, most well sampled light curves were in the optical regime, while in the \\emph{Swift} era most well sampled light curves are in the X-ray regime. Our expectations of the observable signature of a jet break are hence based on the breaks observed pre-\\emph{Swift}, predominately by optical telescopes, and the models which explained them, notably those of \\cite{rhoads1997:ApJ487,rhoads1999:ApJ525} and \\cite{sari1999:ApJ519}. It is not clear that the breaks will be identical in both the X-ray and optical regimes. In the cases of GRB\\,990510 and GRB\\,060206, both have clear breaks in the optical but only marginal breaks in the X-ray (\\citealt{kuulkers2000:ApJ538,curran2007:MNRAS381} respectively). Regardless of this, GRB\\,990510 is taken as a prototypical achromatic jet break, while the achromatic nature of the GRB\\,060206 break is only evident  when supported by the broad-band spectral indices.   In this paper we address the issue of X-ray breaks which are possibly `hidden' and hence the light curve misinterpreted as being a single power law. We do so by synthesising X-ray light curves and fitting both single and broken power laws, and comparing goodness of each fit via the F-test. In \\S\\ref{section:method} we introduce our method while in \\S\\ref{section:results} we present the results of our Monte Carlo analysis. In \\S\\ref{section:discussion} we discuss the implications of these results in the overall context of GRB observations and we summarise our findings in \\S\\ref{section:conclusion}. ",
        "conclusions": "\\label{section:conclusion}   As underlying smoothly broken power laws may be \\emph{hidden} as single power laws, in XRT light curves, we should exercise caution in ruling out breaks based solely on a comparison of the nominal fitted slopes. We hence need to be cautious in implying chromatic breaks from optical and X-ray light curves where an X-ray break is not ruled out with a high degree of certainty. Multi-wavelength temporal and spectral data are required to make a confident statement on the absence or presence of an achromatic break. There may be a bias towards detecting breaks from bursts with a constant density circumburst medium, as these are most obvious and easily detectable.  We have shown that the fitted temporal slopes of smoothly broken power laws show significant variation about central values and that, especially in the case of a smooth break, the underlying values are difficult to extract via a fit. More accurate and wavelength specific descriptions of breaks, likely involving numerical simulations of the jet dynamics, are necessary to better understand the observable signature of breaks and to clarify whether the breaks could vary somewhat between wavebands."
    },
    "0710/0710.2693_arXiv.txt": {
        "abstract": "{Recently, we have developed a method useful for mapping large-scale horizontal velocity fields in the solar photosphere. The method was developed, tuned and calibrated using the synthetic data. Now, we applied the method to the series of Michelson Doppler Imager (MDI) dopplergrams covering almost one solar cycle in order to get the information about the long-term behaviour of surface flows. We have found that our method clearly reproduces the widely accepted properties of mean flow field components, such as torsional oscillations and a pattern of meridional circulation. We also performed a periodic analysis, however due to the data series length and large gaps we did not detect any significant periods. The relation between the magnetic activity influencing the mean zonal motion is studied. We found an evidence that the emergence of compact magnetic regions locally accelerates the rotation of supergranular pattern in their vicinity and that the presence of magnetic fields generally decelerates the rotation in the equatorial region. Our results show that active regions in the equatorial region emerge exhibiting a constant velocity (faster by $60 \\pm 9$~\\mps{} than Carrington rate) suggesting that they emerge from the base of the surface radial shear at $0.95\\ R_\\odot$, disconnect from their magnetic roots, and slow down during their evolution. ",
        "introduction": "The largest scale velocity fields in the solar photosphere consist of rotation profile and a meridional flow pattern.  Basically, the differential rotation is described as an integral of the zonal component $v_\\varphi$ of the studied flow field. The integrated flow field may be obtained using spectroscopic method, using tracer-type measurements, or using helioseismic inversions. The latest case allows to measure the solar rotation not only as a function of heliographic latitude, but also as a function of depth. From the helioseismic inversion we know that throughout the convective envelope, the rotation rate decreases monotonically toward the poles by about 30~\\%. Angular velocity contours at mid-latitudes are nearly radial. Near the surface at the top of the convection zone there is a layer of a large radial shear in the angular velocity. At low and mid-latitudes there is an increase in the rotation rate immediately below the photosphere which persists down to $r \\sim 0.95~R_\\odot$. The angular velocity variation across this layer is roughly 3~\\% of the mean rotation rate and according to the helioseismic analysis of \\cite{2002SoPh..205..211C} the angular velocity $\\omega$ decreases within this layer approximately as $r^{-1}$, where $r$ is a radial coordinate. At higher latitudes, the situation is less clear. For the overview of solar differential rotation measurements see \\cite{1985SoPh..100..141S} or a more recent review by \\cite{2000SoPh..191...47B}.  The \\emph{torsional oscillations}, in which narrow bands of faster than average rotation, interpreted as zonal flows, migrate towards the solar equator during the sunspot cycle, were discovered by \\cite{1980ApJ...239L..33H}. Later research \\citep{2001ApJ...559L..67A} found that there exist two different branches of torsional oscillations. At latitudes below about 40\\,$^\\circ$, the bands propagate equatorward, but at higher latitudes they propagate poleward. The low-latitude bands are about 15\\,$^\\circ$ wide in latitude. The flows were studied in surface Doppler measurements \\citep{2001ApJ...560..466U}, and also using local helioseismology \\citep{1997ApJ...482L.207K}. The surface pattern of torsional oscillations penetrate deeply in the convection zone, possibly to its base, as suggested by \\cite{2002Sci...296..101V}. The amplitude of the angular velocity variation is about 2--5~nHz, which is roughly 1~\\% of the mean rotational rate (5--10~\\mps). The direct comparison between different techniques inferring the surface zonal flow pattern \\citep{2006SoPh..235....1H} showed that the results are sufficiently coherent.  The surface magnetic activity corresponds well with the torsional oscillation pattern -- the magnetic activity belt tends to lie on the poleward side of the faster-rotating low-latitude bands. The magnetic activity migrates towards the equator with the low-latitude bands of the torsional oscillations as the sunspot cycle progresses \\citep{2004ApJ...603..776Z}. Some studies \\citep[e.\\,g.][]{2002ApJ...575L..47B} suggest that meridional flows may diverge out from the activity belts, with the equatorward and poleward flows well correlated with the faster and slower bands of torsional oscillations. In a recent theoretical study by \\cite{2007ApJ...655..651R} it was suggested that the poleward-propagating high-latitude branch of the torsional oscillations can be explained as a response of the coupled differential rotation/meridional flow system to periodic forcing in midlatitudes of either mechanical (Lorentz force) or thermal nature. The equatorward-propagating low-latitude branch is most likely not a consequence of the mechanical forcing alone, but rather of thermal origin.  The axisymmetric flow in the meridional plane is generally known as the \\emph{meridional circulation}. The meridional circulation in the solar envelope is much weaker than the differential rotation, making it relatively difficult to measure.  Two principal methods are widely used to measure the meridional flow: feature tracking and direct Doppler measurement. There are several difficulties complicating the measurements of the meridional flow using tracers. Sunspots and filaments do not provide sufficient temporal and spatial resolution for such studies. Sunspots also cover just low latitudinal belts and do not provide any informations about the flow in higher latitudes. Doppler measurements do not suffer from the problem associated with the tracer-type measurements, however they introduce another type of noisy phenomena. It is difficult to separate the meridional flow signal from the variation of the Doppler velocity from the disc centre to the limb. Using different techniques, the parameters of the meridional flow show large discrepancies. It is generally assumed that the solar meridional flow in the close subphotospherical layers is directed poleward with one cell per hemisphere. Such flow is also produced by early global hydrodynamical simulation such as \\cite{1982ApJ...256..316G}. As reviewed by \\cite{1996ApJ...460.1027H}, the surface or near sub-surface velocities of the meridional flow are generally in the range 1--100~\\mps, the most often measured values lie within the range of 10--20~\\mps. The flow has often a complex latitudinal structure with both poleward and equatorward flows, multiple cells, and large asymmetries with respect to the equator. \\cite{2004ApJ...603..776Z} used the time-distance helioseismology to infer the properties of the meridional flow in years 1996--2002. They found the meridional flows of an order of 20~\\mps, which remained poleward during the whole period of observations. In addition to the poleward meridional flows observed at the solar minimum, extra meridional circulation cells of flows converging toward the activity belts are found in both hemispheres, which may imply plasma downdrafts in the activity belts. These converging flow cells migrate toward the solar equator together with the activity belts as the solar cycle evolves. \\cite{2002ApJ...575L..47B} measured the meridional flow (and torsional oscillations) using the time-distance helioseismology and found the residual meridional flow showing divergent flow patterns around the solar activity belts below a depth of 18~Mm.  The most complete maps of the torsional oscillations and the meridional flow available at present have been constructed on the basis of Mt.~Wilson daily magnetograms \\citep[see][]{ulrich90}. The measurements cover more than 20 years (since 1986) and the results obtained using this very homogenous material agree well with the properties described above.  The modern dynamo flux-transport models use the meridional flow and the differential rotation as the observational input. In the models by Dikpati et al. (\\citeauthor{2006ApJ...638..564D} \\citeyear{2006ApJ...638..564D} or \\citeauthor{2006ApJ...649..498D} \\citeyear{2006ApJ...649..498D}) the return meridional flow at the base of the convection zone is calculated from the continuity equation. They found the turnover time of the single meridional cell of 17--21~years. The meridional flow is assumed to be essential for the dynamo action, global magnetic field reversal and forecast of the future solar cycles.  There are known many relations of the differential rotation profile to the phase of the progressing solar cycle -- see e.\\,g. \\cite{2003SoPh..212...23J} or \\cite{2005ApJ...626..579J} -- showing for example different properties of the differential rotation profile in the odd and even solar cycles. The rotation of the sunspots in relation to their morphological type was studied e.\\,g. by \\cite{1986AA...155...87B} who found that more evolved types of sunspots (E, F, G and H type) rotate slower than less evolved types. \\cite{2004SoPh..221..225R} investigated Greenwich Photoheliographic Results for the years 1874--1976 and found a clear evidence for the deceleration of the sunspots in the photosphere with their evolution. \\cite{2002aprm.conf..427H} found that the leading part of a complex sunspot group rotate about 3~\\% faster than the following part. The dependence of the rotation of sunspot on their size and position in the bipolar region was investigated by \\cite{1994SoPh..151..213D}. They explained the observed behaviour through a subtle interplay between the forces of magnetic buoyancy and drag, coupled with the role of the Coriolis force acting on rising flux tubes. This dynamics of rising flux tubes also explains the faster rotation of smaller sunspots. In average, sunspots rotate about 5~\\% faster than the surrounding plasma.  In the theoretical study \\citep{2004SoPh..220..333B} based on 3-D numerical simulations of compressible convection under the influence of rotation and magnetic fields in spherical shells, the author stated that in the presence of magnetic field the Maxwell stresses may oppose the Reynolds stresses and therefore the angular momentum is propagated more to the poles than without the presence of magnetic fields. As a consequence the rotation profile is more differential in the periods of lowered magnetic activity and it leads to the increase of the rotation in low latitudes. This behaviour was observed in many studies, e.\\,g. \\cite{1990ApJ...357..271H}.  The subject of this work is a verification of the performance of the method described in \\cite{svanda06} (hereafter Paper~I) on the real data and the investigation of long-term properties of the flows at largest scales obtained with this method. We shall also discuss the influence of magnetic fields on the measured zonal flow in the equatorial region. This topic will be studied more in detail in one of the next papers in the series.  \\begin{figure*} \\centering \\resizebox{\\textwidth}{!}{\\includegraphics{fig01.ps}} \\caption{Mean meridional flow in time and heliographic latitude. It can be clearly seen that for almost all the processed measurements a simple model of one meridional cell per hemisphere would be sufficient. However, some local corruptions of this simple idea can be noticed on both hemispheres.} \\label{fig:meridional} \\end{figure*}   \\begin{figure*} \\centering \\resizebox{\\textwidth}{!}{\\includegraphics{fig02.ps}} \\caption{Torsional oscillations. The residua of the mean zonal flow with respect to its parabolic fit displayed in time and heliographic latitude. In the period of weak magnetic activity the pattern of belts propagating towards the equator is very clear. In the periods of stronger magnetic activity the flow field is influenced by the local motions in active regions and therefore the pattern of torsional oscillations is not clearly seen.} \\label{fig:torsional} \\end{figure*} ",
        "conclusions": "We have verified that the method developed and tested using the synthetic data (Paper~I) is suitable for application to real data obtained by the MDI on-board SoHO and maybe also to the data that will be produced by its successor Helioseismic Michelson Imager (HMI) on-board the Solar Dynamic Observatory (SDO). HMI will have a greater resolution and will cover larger time span than two months each year. We verified that the long-term evolution of the horizontal velocity fields measured using our method is in agreement with generally accepted properties.  During the periodic analysis of the equatorial area we found two suspicious periods in the real data, which are not present in the control data set containing the inclination of the solar axis towards the observer, the quantity that can bias systematically and periodically the results by a few \\mps. The periods of 1.8~year and 4.7~years need to be confirmed using a more homogenous data set.  We also found that the presence of the local magnetic field generally speeds-up the region occupied by the magnetic field. However, we cannot conclude that there exists a dependence of this behaviour for different types of sunspots. We can generally say that the more evolved types of active regions rotate slower than the young ones, however the variance of the typical rotation rate is much larger than the differences between the rates for each type. We have found that the distribution of active regions rotation is bimodal. The faster-rotating cases correspond to new and growing active regions. Their almost constant rotation speed suggests that they emerge from the base of the surface radial shear at $0.95\\ R_\\odot$. The decaying and recurrent regions rotate slower with a wider scatter in their velocities. This behaviour suggests that during the sunspot evolution, sunspots loose the connection to their magnetic roots. Both regimes alternate with a period of approximately one Carrington rotation in years 2001 and 2002, which suggests that new active regions emerge in groups and may have a linked evolution."
    },
    "0710/0710.2370_arXiv.txt": {
        "abstract": "We present an investigation of quasar colour-redshift parameter space in order to search for radio-quiet red quasars and to test the ability of a variant of the KX quasar selection method to detect quasars over a full range of colour without bias. This is achieved by combining IRIS2 imaging with the complete Fornax Cluster Spectroscopic Survey to probe parameter space unavailable to other surveys. We construct a new sample of 69 quasars with measured $b_J - K$ colours. We show that the colour distribution of these quasars is significantly different from that of the Large Bright Quasar Survey's quasars at a 99.9\\% confidence level. We find 11 of our sample of 69 quasars have signifcantly red colours ($b_J - K \\geq 3.5$) and from this, we estimate the red quasar fraction of the $K \\leq 18.4$ quasar population to be 31\\%, and robustly constrain it to be at least 22\\%. We show that the KX method variant used here is more effective than the UVX selection method, and has less colour bias than optical colour-colour selection methods. ",
        "introduction": "\\label{introduction} \\citet{1995Natur.375..469W} compared the $B - K$ colours of Parkes Half-Jansky Flat-Spectrum Sample (PHJFS) \\citep{1997MNRAS.284...85D} quasars to those of Large Bright Quasar Survey (LBQS) \\citep{1995AJ....109.1498H} quasars. The PHJFS quasars were found to have a significantly broader and redder distribution of $B - K$ colours than LBQS quasars. This disparity should not exist according to the unified model of Active Galactic Nuclei (AGN) galaxies \\citep{1993ARA&A..31..473A}. \\citet{1995Natur.375..469W} concluded that there exists a population of red quasars that at that time had not been discovered. \\citet{1995Natur.375..469W} investigated this possibility using a simple model of LBQS selection effects, and demonstrated that the LBQS was biased against the detection of red quasars. The bias was caused by the blue magnitude limit of the LBQS, because in a blue-flux limited survey, red quasars need to be intrinsically brighter than blue quasars to be detected \\citep{1995Natur.375..469W}. The bias causes fewer red quasar detections, because the luminosity function of quasars is steep at the bright end \\citep{1991MNRAS.251..482B}. Other surveys with a blue-flux limit are similarly biassed, independent of the quasar selection method. Using their simple model, \\citet{1995Natur.375..469W} estimated that the red quasars missed by previous surveys constitute as much as 80\\% of all quasars.  Red quasars are defined by a colour criterion. The best colour to use for this criterion is that made from filters separated by the largest wavelength range possible, because anomalously redder quasar colours stand out more. For the datasets used in this paper, the colour covering the largest wavelength range is $U - K$. We used the next largest colour $b_J - K$, as far more objects, $\\sim$ 3 times as many, have a measured $b_J - K$ colour compared to the number with a measured $U - K$ colour. The additional advantage of having used $b_J - K$ is that it is possible to compare the colours of the quasar sample presented here with the $B - K$ colour used in \\citet{2004ApJ...607...60G} and the $g - K$ colour of \\citet{2004AJ....128.1112H}. Throughout this paper we define a red quasar by the criterion, $b_J - K \\geq$ 3.5 \\citep{2002ApJ...564..133G}, because when describing the optical-infrared region of quasar spectra with a power-law, $S(\\lambda) \\propto \\lambda^{\\alpha}$, this colour corresponds to an effective optical-infrared spectral index of $\\alpha \\geq 1$.  The existence of red quasars, other than in radio-selected quasar samples, over most of the $b_J - K$ colour range of the PHJFS quasars, satisfying the above criterion, has already been demonstrated \\citep{2004AJ....128.1112H, 2003AJ....126.1131R, 2004ApJ...607...60G}. The \\citet{2004ApJ...607...60G} red quasar sample contains the reddest quasars found, with $B - K$ colours of $5 \\leq B - K \\leq 8$. \\citet{2004AJ....128.1112H, 2003AJ....126.1131R} found red quasars in the Sloan Digital Sky Survey (SDSS)\\citep{2000AJ....120.1579Y}, and \\citet{2004ApJ...607...60G} found them in a combination of the 2-Micron All Sky Survey (2MASS) \\citep{2006AJ....131.1163S} and an Automatic Plate Measuring CATalogue (APMCAT) of POSS/UKST sky survey plates \\citep{APMCAT}.  Despite confirming the existence of red quasars, current observations have not explained the mechanism or mechanisms responsible for their existence. Dust obscuration by dust at the quasar redshift, in either the host galaxy or the quasar accretion disk \\citep{1995Natur.375..469W}, was the first mechanism proposed for quasar reddening. \\citet{1995Natur.375..469W} proposed this mechanism as a correlation between $B - K$ colour and redshift was not observed. The result of reddening a composite optically selected quasar spectral energy distribution (SED) with dust at the quasar redshift, was found to be consistent with the dust mechanism, because the reddening resulted in the SED of a red PHJFS quasar \\citep{1995Natur.375..469W}.  As an alternative, \\citet{1998MNRAS.295..451B} suggested that the $B - K$ spread in the \\citet{1995Natur.375..469W} sample is partly or entirely caused by host galaxy starlight. Quasars with flat-radio-spectra are typically hosted by later galaxy types \\citep{1998MNRAS.295..451B}, and the old stellar population of late type galaxies is a stronger emitter in K than B, reddening quasars \\citep{1998MNRAS.295..451B}. An analysis in \\citet{1998MNRAS.301..975M} of the host galaxy contribution to the red $B - K$ colours of the PHJFS quasars in \\citet{1995Natur.375..469W} was inconclusive, but suggested host galaxy contribution was insufficient to cause the red $B - K$ colours observed. Further work in \\citet{2006MNRAS.367..717M} confirmed that host galaxy starlight can definitely redden quasars, particularly low luminosity and low redshift resolved quasars. In accordance with \\citet{2006MNRAS.367..717M}, in this paper we use the criteria that luminous unresolved quasars, which are not located at low redshift, are unlikely to be reddened by host galaxy contribution.  Alternatively, \\citet{1996Natur.379..304S} proposed that a synchrotron component of quasar emission created a K-band excess, reddening quasars. An analysis of 157 PHJFS quasars and 12 LBQS quasars found the shape of the SEDs was consistent with both line-of-sight dust and synchrotron emission mechanisms \\citep{2000PASA...17...56F}. In \\citet{2001AJ....121.2843B} it was suggested that because any synchrotron component to quasar emission is weaker in radio-quiet quasars than in radio-loud quasars, any synchrotron emission will only contribute to a small fraction of any reddening. Based on the quasar sample contained therein, \\citet{2001AJ....121.2843B} suggested that any radio-quiet quasar redder than $b_J - K \\geq$ 3.7 is too red to be the result of synchrotron emission, and is most likely the result of dust obscuration. Throughout this paper we use the criterion that any red quasar that is also radio-quiet, is unlikely to be reddened by synchrotron emission.  Finally, \\citet{1997MNRAS.288..138M} suggested that a small subset of red quasars was caused by the lensing of a quasar by a dusty galaxy. Not only does the dust in the lens redden the quasar but the dusty galaxy potentially magnifies the host galaxy starlight contribution \\citep{2002ApJ...564..133G}. This last mechanism was proposed after \\citet{1997MNRAS.288..138M} observed a similar phenomenon to \\citet{1995Natur.375..469W}; with radio-selected lenses having redder colours than optically selected lenses.  Whichever mechanism creates red quasars, they are possibly an evolutionary stage. After observing low-ionisation Broad Absorption Line (BAL) quasars, \\citet{2002AJ....123.2925L} suggested a model of quasar evolution where a quasar occupies a thick dusty torus after forming, eventually the torus dissipates leaving behind an optically blue quasar. Alternately, quasars may emerge from dusty starburst galaxies, residing in the dust stirred up by mergers \\citep{2005AAS...20717401U}. Either evolutionary sequence would naturally create a situation where quasars are reddened by dust at the quasar redshift. A third evolutionary path has been proposed where red quasars are an evolutionary link between Ultra-Luminous Infrared Galaxies (ULIRGs) and optically selected quasars \\citep{2006Ap&SS.302...17C}. Determining the fraction of red quasars and any dependence on redshift, luminosity, radio flux, etc. will indicate if red quasars are an evolutionary stage, and which evolutionary path red quasars belong to.  Depending on the mechanism, red quasars might explain the quasar contribution to the X-Ray Background (XRB). Using a theoretical model, \\citet{1994MNRAS.270L..17M} showed that a ratio of two to three times as many dust obscured to unobscured active galactic nuclei (AGN) galaxies can adequately account for the existence of the XRB. The model also accounts for the observed spectrum and the source counts in the soft and hard X-ray bands. The latest theoretical models of the XRB require a smaller ratio, lying somewhere between 0.6 to 1.5 \\citep{2007A&A...463...79G}. If it can be shown that red quasars are the result of dust obscuration, and that the red quasar fraction is within 37.5\\% to 60\\%, then red quasars are the most probable source of the quasar contribution to the XRB.  The existence of red quasars has various implications, the main one is that it can no longer be assumed that existing quasar samples are representative of the entire quasar population. The statistics of a sample reflects the properties of the sub-set of a population from which the sample is drawn, previous quasar samples, such as the LBQS, are biassed to the inclusion of blue quasars; therefore, the statistics of such samples predominantly reflects the blue quasar population. Use of existing quasar datasets requiring a sample that is representative of the entire quasar population is therefore subject to review. Direct examples of this are that the luminosity function of quasars and the change in number density with redshift need to be re-evaluated while including red quasars.   In order to learn more about red quasars, we need to spectroscopically observe quasars over their entire colour range. To achieve this, future surveys will have to select quasar candidates for spectroscopic follow-up using a technique that is not biased against either red or blue quasars. One candidate for a suitable selection technique is the KX method, developed in \\citet{2000MNRAS.312..827W}. The KX method utilises the power-law nature of quasar SEDs at long wavelengths combined with morphological classification. Stars experience a turnover in their SED in H-band while quasars follow a power-law, allowing discrimination between stars and quasars. The stellar morphology of quasars discriminates between quasars and galaxies. Observationally, this requires constructing a colour-colour plot of all objects with a stellar morphology, using an optical colour and a colour that straddles the H-band.  In this paper we explore regions of the quasar colour-redshift space unavailable in previous work, and test a KX method variant. We aim to find the boundaries of the quasar colour-redshift parameter space that is occupied. In doing so, we will improve the estimate of the red quasar fraction and possibly shed light on the mechanism that creates red quasars. In testing a KX method variant, we intend to verify if it is a viable method of selecting quasars, without bias, from the entire quasar colour-redshift parameter space. We intend to apply a KX method variant to our quasar sample, and then contrast it with established methods of selecting potential quasars for spectroscopic follow-up.  To achieve these goals we use a sub-sample of the quasars identified in the Fornax Cluster Spectroscopic Survey (FCSS) \\citep{2000A&A...355..900D}, a spectroscopic survey of all extended objects to $b_J = 19.8$ and point sources to $b_J = 21.5$. The sub-sample consists of all quasars that match to an object in Ks-band imaging to a depth of K = 18.4, obtained using the Infrared Imager and Spectrograph 2 (IRIS2) instrument on the 3.92-m Anglo-Australian Telescope (AAT). The Ks-band imaging is deeper than 2MASS allowing us to explore parts of the $b_J - K$ colour space unavailable in \\citet{2004ApJ...607...60G}, which was limited to $b_J - K \\geq$ 5 by the K = 14.5 magnitude limit of 2MASS. Furthermore, the quasar sample used covers a larger range of redshift than was examined in \\citet{2004AJ....128.1112H,2003AJ....126.1131R}, which were limited to z $\\leq$ 2.2. Because the FCSS used no selection criteria, our quasar sample is bias free, and ideal for testing the ability of the KX method to select quasars from a K magnitude limited survey. Additionally, most of the FCSS objects have a U magnitude as well as a $b_J$ magnitude, allowing the KX method to be directly contrasted with traditional methods for optically selecting quasars.  The structure of this paper is as follows. In Section \\ref{data} the datasets used in this paper are described. Section \\ref{redquasars} details the detection of red quasars in the quasar sample and an analysis of the quasar sample. This analysis involves a comparison of the $b_J - K$ colours of the sample quasars to LBQS quasars and a determination of the effect of our selection on the quasar sample constructed here. Section \\ref{redquasars} concludes with the calculation of the fraction of red quasars. The effectiveness of the KX method is assessed in Section \\ref{testKX} and following this are our conclusions in Section \\ref{conclusion}.  \\section[]{Photometric and Spectroscopic Data} \\label{data} The sample of quasars used here and their photometric information were obtained by combining three datasets. Astrometry and optical photometry were obtained from the Automatic Plate Measuring CATalogue (APMCAT) of POSS/UKST sky survey plates \\citep{APMCAT}. IRIS2 imaging was used to obtain the K-band photometry used here. The necessary spectroscopic identifications were provided by the Fornax Cluster Spectroscopic Survey (FCSS) catalogue of \\citet{2000A&A...355..900D}. The datasets were combined by independently matching the astrometry of K-band catalogue objects and FCSS objects to that of APMCAT catalogue objects. In the rest of this section, we describe the datasets in more detail.  The FCSS catalogue is a blind spectroscopic survey of Fornax objects in the APMCAT catalogue. The FCSS is an ideal source of spectroscopic identifications because the lack of target selection precludes biasing. In the FCSS, spectroscopy was carried out on all extended sources to $b_J \\leq 19.8$ and all point sources to $b_J \\leq 21.5$. Not every source was spectroscopically observed, and the fraction of observed sources is referred to as the spectroscopic completeness. The FCSS obtained a redshift and spectroscopic identification for more than 90\\% of the objects observed, referred to as the redshift completeness. In most cases where a spectroscopic identification could not be obtained, neither was a redshift because of the poor spectrum quality. Because the redshift completeness was consistently high for the entire FCSS, we have combined the redshift completeness with the spectroscopic completeness, and throughout the paper the term `completeness' refers to this combination. The completeness of the FCSS is 96\\% for $b_J < 20.5$, 82\\% for $20.5 \\leq b_J < \\leq 21$ and 36\\% for  $21 < b_J \\leq 21.5$ sources. We account for the completeness by appropriately weighting quasars during our analysis in later sections.  IRIS2 was used in December 2001 to obtain Ks-band imaging using 60s exposures. The raw data were processed using the Observatory Reduction and Acquisition Control project - Data Reduction pipeline (ORAC-DR) outlined in \\citet{1999ASPC..172...11E,1999ASPC..172..171J}. The reduced IRIS2 imaging dataset was then analysed with SExtractor \\citep{1996A&AS..117..393B} to obtain a photometric catalogue of corrected isophotal magnitudes. Corrected isophotal magnitudes were used instead of PSF or small aperture magntiudes, because PSF and small aperture magnitudes minimise any possible extended contribution to the flux of sources in the IRIS2 imaging. The Ks-band catalogue magnitude zero-point was then calibrated, by matching point sources in the Ks-band catalogue to the 2MASS point source catalogue. Calibrating also facilitated the conversion from Ks magnitudes to the K-band magnitudes used throughout this paper. After calibration, the magnitude limit of the K-band photometry was found to be 18.4 magnitudes.  In the FCSS, quasars were spectroscopically identified as objects with broad permitted lines with a measured full width half-maximum $>$ 1100 km $s^{-1}$ \\citep{2001MNRAS.324..343M}. After matching the IRIS2 imaging catalogues and the FCSS quasar identifications to the APMCAT object positions using blind matching, to a radius of 10\"; a catalogue of all the necessary photometry was created, and all quasars within both the IRIS2 imaging K magnitude limit, and the FCSS spectroscopic identifications $b_J$ magnitude limit were identified. This resulted in a sample of 69 spectroscopically identified quasars with U, $b_J$ and K magnitudes. All 69 of these quasars are APMCAT point sources; therefore, our quasar sample is limited to $K \\leq 18.4$ and $b_J \\leq 21.5$. Of these 69 quasars, 62 have an R magnitude as required by our KX method variant. The missing R magnitudes are caused by the different magnitude limits in the $b_J$ and R bands of the APMCAT catalogue, 22.5 and 21 magnitudes respectively. ",
        "conclusions": "\\label{conclusion} In this paper, we constructed a sample of 69 quasars with measured $b_J - K$ colours, using a combination of APMCAT, IRIS2 imaging and FCSS spectroscopic identifications. 11 of these quasars are red, satisfying $ b_J - K \\geq$ 3.5, and all of these red quasars are radio-quiet according to the All Sky Optical Catalogue of Radio/X-Ray Sources. In accordance with \\citet{2001AJ....121.2843B} and \\citet{2006MNRAS.367..717M}, as the 11 red quasars found here are unresolved, luminous and radio-quiet sources predominantly located at $z > 1$, this strongly indicates that the main cause of the red $b_J - K$ colours is dust at the quasar redshift. Comparing our quasar sample to LBQS quasars, we found that the two $b_J - K$ colour distributions are different at the 99.9\\% confidence level, with our sample having a significantly broader, redder distribution. Analysis of the uncertainties in the $b_J - K$ colours demonstrated that neither our red quasar detections or our comparison to the LBQS quasars is affected by them. A second analysis, on the effects of our datasets magnitude and completeness limits, revealed that they limited our measured $b_J - K$ distribution to $ b_J - K <$ 5. As red quasars are observed up to this colour limit, we concluded that the true $b_J - K$ distribution of quasars is even broader and redder than observed here. From the observed $b_J - K$ colour distribution, we robustly constrained the red quasar fraction of the $K \\leq 18.4$ quasar population to be greater than 22\\%, and from a model estimated it to be 31\\%.  Using the quasar sample constructed here, the viability of a KX method variant was tested. Using a $b_J - R$ vs $R - K$ plot, the KX method variant was capable of separating the quasar sample out from the other objects in the APMCAT catalogue. Comparing the KX method variant to the UVX and 2QZ multicolour methods, the KX method variant was found to be as effective in the number of quasars it selected; however, it was superior at selecting quasars independent of colour. For those reasons, the KX method variant used here is an ideal technique for future large surveys to use in selecting potential quasars, whether red or blue, for spectroscopic follow-up."
    },
    "0710/0710.2000_arXiv.txt": {
        "abstract": "We present a comprehensive spectroscopic imaging survey of the distribution and kinematics of atomic hydrogen (HI) in 16 nearby spiral galaxies hosting low luminosity AGN, observed with high spectral and spatial resolution (resolution: $\\sim$20\\arcsec, $\\sim$5~km~s$^{-1}$) using the NRAO Very Large Array (VLA). The sample contains a range of nuclear types, ranging from Seyfert to star-forming nuclei and was originally selected for the NUclei of GAlaxies project (NUGA) - a spectrally and spatially resolved interferometric survey of gas dynamics in nearby galaxies designed to identify the fueling mechanisms of AGN and the relation to host galaxy evolution. Here we investigate the relationship between the HI properties of these galaxies, their environment, their stellar distribution and their AGN type. The large-scale HI morphology of each galaxy is classified as ringed, spiral, or centrally concentrated; comparison of the resulting morphological classification with AGN type reveals that ring structures are significantly more common in LINER than in Seyfert host galaxies, suggesting a time evolution of the AGN activity together with the redistribution of the neutral gas. Dynamically disturbed HI disks are also more prevalent in LINER host galaxies than in Seyfert host galaxies. While several galaxies are surrounded by companions (some with associated HI emission), there is no correlation between the presence of companions and the AGN type (Seyfert/LINER). ",
        "introduction": "\\label{sec:intro} Active Galactic Nuclei (AGN) represent some of the most extreme conditions and include the most powerful individual objects to be found throughout the Universe. Nowadays, the phenomena of nuclear activity is generally understood to be the result of accretion of material onto a SuperMassive Black Hole (SMBH); very likely through the infall of gas from its host galaxy \\citep[e.g.][]{Rees}. Although SMBHs exist in most galaxies, only a small fraction of all galaxies exhibit nuclear activity \\citep{Huchra,Ho97b, Miller}. The majority of AGNs show low-ionization narrow emission-line regions (LINERs), which have luminosities lower than Seyfert galaxies and quasars. Originally, LINERs and Seyferts were considered to form a continuous distribution in the usual line-ratio classification diagrams. Recent investigations of the host properties of emission-line galaxies from the Sloan Digital Sky Survey \\citep{Kew06} confirmed that their nuclear activities are due to accretion by a SMBH. But interestingly, LINERs and Seyferts are also clearly separable in emission line ratio diagrams \\citep{Gro06}, and that these two classes have distinct host properties \\citep{Gro06}, implying that their distinction is in fact not simply an arbitrary division. An open question is therefore what mechanisms give rise to the different AGN types observed. \\par A coherent picture of AGN feeding, e.g. how to remove the angular momentum from host galaxy gas so that it can reach the central parsec, is still missing. Therefore it is expected that a hierarchy of mechanisms might combine to transport gas from large kpc scales down to the inner pc scales \\citep{Shl90, Com03, Jog04}. Additionally, trigger mechanisms driven by the galactic environment, like minor mergers and tidal forces, might play a role as well \\citep{Barnes}. A larger gas fraction for Seyferts than for normal galaxies was suggested by \\cite{Hunt99a} and a prevalence of stellar rings (RC3 classification) in active galaxies (Seyfert and LINERs) compared to non-active galaxies was suggested \\citep{Hunt99b}. NICMOS imaging of the centers of 250 nearby galaxies \\citep{Hunt04} revealed systematic morphological differences between HII/starburst (most disturbed), Seyfert (intermediate disturbed), LINER and normal galaxies (most regular). The study of the circumnuclear regions of 24 Seyfert2 galaxies using HST images \\citep{Mar99} revealed a dominance of nuclear spirals, suggesting that spiral dust lanes may be responsible for feeding gas to the central engines. \\par  Theoretical simulations have also made progress in addressing the nature of nuclear fueling with respect to different types of gravitational instabilities and their feeding efficiency: nested bars \\citep[e.g.][]{Shl89, Friedli, Mac00, Eng04}, gaseous spiral density waves \\citep[e.g.][]{Eng00, Mac02, Mac04a, Mac04b}, m = 1 perturbations \\citep[e.g.][]{Shu, Jun96, Gar00} and nuclear warps \\citep{Sch00} have all been suggested as possible transport mechanisms. \\par  In order to distinguish models for nuclear fueling, observations of neutral gas with high angular and velocity resolution are required. Therefore, the IRAM key project NUclei of GAlaxies \\citep[ NUGA; see][]{Gar03} was established - a spectroscopic imaging survey of gas in the centers of nearby low luminosity AGN. As most of the gas in galaxy nuclei is in the molecular phase, the survey used millimeter CO lines to conduct a detailed mapping of molecular gas dynamics at {\\em high-resolution} (0.5$\\arcsec$) in the central kiloparsec of AGN hosts. The CO-NUGA survey (using the IRAM Plateau de Bure Interferometer) reveals a wealth of nuclear gas distribution and kinematics which are studied in detail: (a) m=1 gravitational instabilities (one-arm spirals and lopsided disks; NGC~4826: \\cite{Gar03b}), (b) m=2 instabilities (two-arm spiral wave) expected to form in stellar bar potentials show only small amounts of molecular gas coincident with the AGN (e.g. NGC~4569: \\cite{Boo07}, NGC~4579: \\cite{Gar05} NGC~6951: Schinnerer et al., in prep.), (c) stochastic spirals that are related to non self-gravitating instabilities, and rings \\citep[NGC~7217:][]{Com04}, and (d) large scale warps that might extend into the central kiloparsec \\citep[NGC~3718:][]{Kri05}. Besides the CO studies also extended radio continuum components resembling jet emission from AGN were detected \\citep{Kri07a} and a molecular gas disk/torus of dense gas was observed in the HCN line emission \\citep[NGC~6951:][]{Kri07b}. Additionally, Garc{\\'{\\i}}a-Burillo et al. (\\citeyear{Gar05}) derived in a pilot study the gas inflow rates via measurements of the gravitational torque onto the molecular gas in the central kiloparsec for 4 galaxies. \\par  To complement these CO data and provide a more complete census of the ISM and gas flows from the outskirts of the galaxy disks to the very center, the HI-NUGA project was initiated. HI-NUGA provides observations of the HI gas distributions and kinematics for 16 galaxies (including archival data for 2 galaxies) of the NUGA sample using NRAO's Very Large Array (VLA). Determining the HI distribution and kinematics is necessary to establish whether characteristics of the nuclear dynamics, and hence the nuclear activity, depend on overall host galaxy properties. Possible dependencies might also exist between nuclear modes and large scale drivers (i.e. tidal forces exerted by companions), or the overall content/distribution of the atomic gas. Since the HI gas extends in most disk galaxies much further than the optical disk, it is more loosely bound in the outer region because of the decrease of the gravitational potential.  Hence, it provides an ideal tool to identify interactions and tidal features \\citep[e.g.,][]{Sim87, Mun95}. Furthermore, due to its dissipative nature the gas is more sensitive to dynamical disturbances, both internal ones such as non-axisymmetric potentials or external ones (e.g. tidal interactions). Thus, one could expect, for example, to find a prevalence of Seyfert activity with asymmetries in the gas distribution and kinematics.\\par  In this paper we present for 16 nearby active galaxies data obtained in the 21 cm emission line of neutral hydrogen using the VLA. The relation between HI properties, such as the large scale environment (HI/optical companions, HI disk asymmetries) and the AGN type (Seyfert, LINER, HII) are analyzed in detail. The outline of the paper is as follows: We describe the HI-NUGA sample, the observation and data reduction in \\S \\ref{sec:obs}. The results and derived HI properties are presented in \\S \\ref{sec:res}, along with a discussion of the presence of companions and tidal disturbances of the HI disks and any correlation with AGN type. A comparison of AGN types and optical (stellar) light distribution is also examined. The results are discussed in the context of correlations with the HI environment (e.g. tidal forces) and the AGN type (\\S \\ref{sec:dis}). A summary is presented in \\S \\ref{sec:sum}. ",
        "conclusions": "\\label{sec:dis}  The relationship between large scale environment and nuclear activity has been long debated in the literature. Several studies have found indications for correlations between the environment and nuclear activity \\citep{Sto01, Cha02, Mar03}. In particular it was suggested that interacting galaxies or galaxies with companions exhibit a significant excess of nuclear activity compared to isolated galaxies \\citep{Dah84, Kee85, Raf95}. On the other hand no relation was seen by other studies \\citep{Vir00, Sch01, Fue88, Mac89, Lau95}. Hence the issue of possible relations between the environment and nuclear activity appears to be still controversial. Furthermore, results from Keel et al. (\\citeyear{Kee85}) suggested that nuclear phenomena might likely be triggered by a tidally induced inflow of gas from the disk to the nuclear regions, rather than gas transfer between the interacting galaxies themselves. Keel (\\citeyear{Kee96}) also  found that Seyfert galaxies in pairs actually display smaller kinematic disturbances than non-Seyfert galaxies in pairs, which is obviously in disagreement with the hypothesis that tidal interactions are necessary for the transport of angular momentum and the fueling of the SMBH. Since most of these studies are based on optical/IR imaging, they are in principle less sensitive to distortions than our study of the atomic gas that reacts most readily to tidal disturbances. It should be pointed out that a detailed study of HI gas properties for active versus non-active galaxies (Mundell et al., in prep.) is under-way. \\par  In the context of different AGN types, the analysis of a sample of 451 active galaxies (Sy, LINER, Transition, HII, and absorption-line galaxies) from the Palomar survey \\citep{Sch01} showed no correlation between AGN-type and the percentage of galaxies with nearby companions after taking morphological differences of the host galaxies into account. We also see no evidence for a correlation between the fraction of companions and the AGN type present (Seyfert, LINER galaxies), neither from our HI study nor for cataloged optical companions listed in NED (see \\S \\ref{subsec:res_env}). Note that our sample has a limited number of LINERs (7 galaxies) and Seyfert galaxies (7) and hence the existence of possible relations can not be completely excluded. \\par   \\subsection{HI environment:  tidal forces and their correlation with disturbed HI disks} \\label{subsec:dis_env} Companion galaxies can possibly disturb the HI gas in a galactic disk via tidal forces and hence affect the fueling of the center with gas \\citep[e.g.][]{Kee85}. The sensitivity and endurance of HI to trace the strength and prevalence of tidal interactions among Seyfert galaxies is discussed in detail by Greene et al. (\\citeyear{Gre03}). In our analysis we found companions for about half of our sample. In almost all cases the produced tidal forces are smaller than the binding forces of the affected host galaxy (indicated as $\\textbf{Q}<0$, see \\S \\ref{subsec:res_env} and Tab.~\\ref{tab_sat}).   Only the system NGC5953/54 exhibits signs for very strong gravitational interaction as mentioned in \\S \\ref{subsec:res_env}. Therefore, most of the disturbances in our sample (6/7 galaxies) can not be explained simply by tidal forces presently at work. The most probable explanations for these disturbances are described in the following, listed in decreasing relevance: \\begin{itemize} \\item Interaction with a companion now further or far away. If we assume a relaxation timescale of the disturbance of 3$\\cdot 10^8$yr (typical time for one rotation of a galaxy) and a maximum fly-by velocity of $\\sim$ 500 km s$^{-1}$, the involved companion is expected to be found within a radial distance of $\\sim$ 150 kpc from the disturbed galaxy. Since 71\\% of the disturbed disks show companions in a reasonable distance (projected distance is less than 150 kpc), tidal interactions in the past seem to be primarily responsible of the disturbances identified in the HI disks. \\item Ram pressure stripping \\citep{Cay90, Vol00, Vol01}. In particular for NGC~1961 stripping by intergalactic material was suggested \\citep{Sho82}. But since none of our galaxies with disturbed HI disk lie in a massive group or cluster environment, gas stripping due to ram pressure is not very likely to explain the presence of disturbances in our sample where no nearby companion is found. Furthermore, the outer HI disk, where disturbances are usually seen, has been removed in clusters due to ram pressure stripping, indicated also by a small HI radius (e.g. NGC~4569, NGC~4579 as part of the Virgo cluster; see \\S\\ref{subsec:dis_stellar}) \\item Minor merging, whereby the companion has now fully merged and has left no optical trace. This might be the case for NGC~4736 and NGC~2782 where no companions are found in a reasonable distance for tidal interaction. \\item Large gas accretion from cosmic filaments: Asymmetries in the gas accretion may cause disturbances \\citep[for effects of gas accretion on spiral disk dynamics see][]{Bou02}. \\end{itemize}    \\subsection{HI morphology and comparison with the stellar distribution} \\label{subsec:dis_stellar}  The comparison of the radial density profiles between the HI gas distribution and the stellar distribution revealed significant deviations as expected: 1.) The extent of the HI disk is on average 1.7 times the optical radius, indicated by the Holmberg radius. Only NGC~3627, NGC~4569 and NGC~4579 show a smaller HI radius than optical radius. This can be explained by the fact that they are all members of interacting groups and/or by their rapid motion through an intracluster medium. NGC~3627 is part of the Leo Triplet group and the past encounter with NGC~3628 could explain the spatial conincidence of both the stars and the gas \\citep{Zha93}. The truncated disks of NGC~4569 (and NGC~4579) are most probably a signature of strong ram pressure stripping in the past by the intracluster medium which pervades the Virgo Cluster \\citep{Cay94}. However, most of the galaxies in our sample show a larger HI disk than their optical one. 2.) The radial density profiles exhibit for most of our galaxies a deficiency of HI gas in the inner part of the galaxy disk. This is in general explained by the phase transition from atomic to molecular gas in neutral ISM \\citep{Young}.   \\subsection{Correlations between HI gas properties and AGN activity type} \\label{subsec:dis_correl} As described in \\S \\ref{subsec:res_agn} our analysis revealed that the number of galaxies with disturbed HI disks is higher for LINER galaxies than for Seyfert galaxies. But since several mechanisms which can not easily be distinguished can cause these asymmetries (see \\S \\ref{subsec:dis_env}), it is not possible to draw any strong conclusions explaining the higher fraction for LINERs. \\par  Our study of the HI morphology revealed a significantly higher percentage of galaxies with a HI gas ring for LINER than for Seyfert galaxies. Interestingly, only the study of the Extended 12$\\mu$m Galaxy Sample \\citep{Hunt99b} indicated a prevalence of stellar rings in active galaxies (Seyfert and LINERs), where LINERs have elevated rates of inner rings, while the Seyfert host galaxies have outer ring fractions several times those in normal galaxies. However, we found no HI gas rings in Seyfert host galaxies. Note that stellar rings are not preferentially found in Seyfert, LINERs, or non-AGNs for our sample by using the optical classification from RC3 listed in NED (see classification in Tab.~\\ref{table_intro}; indicated as R or r). To summarize, HI rings are more often found in LINERs while no strong correlation with stellar rings is present. \\par  One possible explanation for an abundance of HI gas rings in LINERs may be a common evolution of the gas distribution in the disk together with the nuclear activity where both are subject to the influence of a present bar or previous one which has now dissolved. This becomes important since rings are linked observationally and phenomenologically to barred galaxy dynamics \\citep[for a general review see][]{But96}, and  hence, expected to be seen after the bar had enough time to redistribute the gas toward the end of its life-time. Thus, a time evolution of AGN types seems to be possible, where Seyfert and LINERs represent different phases of the galaxy activity cycle: Seyfert galaxies are the ones where the fueling process has just been triggered (e.g. through disturbances and/or bar dynamics) while LINERs are the ones where the triggering mechanism has already distributed the gas in a more stable new configuration (rings) that does no longer support the massive inflow of gas. Interestingly, disturbances, which are assumed to be a sign for a recent trigger of gas inflow, are preferentially found in LINERs in our sample. That would suggest that LINERS are the first stage of activity, just after the triggering: The gas is still distributed in rings, and the fueling of the AGN has just started by the redistribution of the gas. However, as most of the disturbances in our sample are probably due to tidal interactions in the past, as explained in \\S \\ref{subsec:dis_env}, the currently observed nuclear activity must not be identical to the one during the ongoing tidal interaction. \\par  A general scenario for self-regulated activity in low-luminosity AGNs was developed by \\cite{Gar05}, in which the onset of nuclear activity is explained as a recurrent phase during the typical lifetime of any galaxy. In this scenario the activity in galaxies is related to that of bar instabilities, expecting that the active phases are not necessarily coincident with the phase where the bar has its maximum strength. Since the infall of gas driven by a bar is self-destructive \\citep{Bou02}, i.e. it weakens and destroys the bar, the potential returns to axisymmetry and the gas piles up in a stable configuration (i.e. in rings at the resonances). At this stage, torques exerted by the gravitational potential are negligible and other competing mechanisms, e.g. viscous torques, must transport the gas in the center of the galaxy. The periods of Seyfert/LINER activity (each lasting $\\sim 10^8$ years) are suggested to appear during different evolutionary stages of a bar episode (typically characterized by a lifetime $\\sim 10^9$ years), depending on the competition between viscosity and gravity torques. The prevalence of HI rings in LINERS, derived  in this work, could be explained by the scenario proposed by \\cite{Gar05}, where AGN activity is linked to the evolutionary state of the bar-induced gas flow.  However, to further substantiate this link requires deriving the complete gravity torque budget based on the HI distribution (see Haan et al. in prep.).   Regarding the relative HI radius ($R_{HI}$/$R_{optical}$) we found a slightly larger HI extent for Seyfert galaxies than for LINERs, which increases even more when neglecting galaxies which lie in galaxy clusters or groups. In contrast, our study of the large environment, the HI gas content (by using the ratio $M_{HI}$/$M_{dyn}$), and the relative HI mass ($M_{HI}$/$M_{dyn}$) revealed no significant correlation with the AGN-type. This might be due to the limited number of galaxies in our sample. Hence we conclude that a detailed HI study with a larger sample may reveal more additional information on the interplay between the gaseous component and the AGN type present."
    },
    "0710/0710.0005_arXiv.txt": {
        "abstract": "{} {UVES spectra  of the very young ( $\\sim$\\,10$^{7}$ years) peculiar B-type star HR\\,6000 were analyzed in the  near-UV and visual spectral regions (3050-9460\\,\\AA{}) with the aim to extend to other spectral ranges the study  made previously in the UV using IUE spectra.} { Stellar parameters \\teff=12850\\,K, \\logg=4.10, and $\\xi$=0\\,km\\,s$^{-1}$, as determined from H$_{\\beta}$, H$_{\\gamma}$, H$_{\\delta}$ Balmer profiles and from the \\ion{Fe}{i}, \\ion{Fe}{ii} ionization equilibrium, were used to compute an individual abundances ATLAS12 model. We identified spectral peculiarities and obtained final stellar abundances by comparing  observed and computed equivalent widths and line profiles. } { The adopted model fails to reproduce the (b-y) and c color indices. The spectral analysis has revealed: the presence of emission lines for \\ion{Mn}{ii}, \\ion{Cr}{ii}, and \\ion{Fe}{ii}; isotopic anomalies for \\ion{Hg}, \\ion{Ca}; the presence of interstellar lines of \\ion{Na}{i} at $\\lambda\\lambda$ 3302.3, 3302.9, 5890, 5896\\,\\AA{}, and of \\ion{K}{i} at 7665, 7699\\,\\AA{}; the presence of a huge quantity of unidentified lines, which we presume to be mostly due to \\ion{Fe}{ii} transitions owing to the large Fe overabundance amounting to [+0.7]. The main chemical peculiarities are an extreme overabundance of Xe, followed by those of Hg, P, Y, Mn, Fe, Be, and Ti. The most underabundant element is Si, followed by  C, N, Al, S, Mg, V, Sr, Co, Cl, Sc, and Ni. The silicon underabundance [$-$2.9] is the lowest value for Si ever observed in any HgMn star. The observed lines of \\ion{He}{i} can not be reproduced by a single value of the He abundance, but they require values ranging from [$-$0.8] to [$-$1.6]. Furthermore, when the observed and computed wings of \\ion{He}{i} lines are fitted, the observed line cores are much weaker than the computed ones. From the present analysis we infer the presence of  vertical abundance stratification for He, Mn, and possibly also P. } {} ",
        "introduction": "HR\\,6000 (HD\\,144667) is one of the most remarkable chemically peculiar (CP) stars. It does not fit any of the CP subclasses, but it seems to combine abundance anomalies from a variety of Bp sub-types. It forms with the star HR\\,5999 (HD\\,144668) the common proper motion visual binary system $\\Delta$199 or Dunlop\\,199 (Bessell \\& Eggen, 1972). The angular separation between the HR\\,6000, which is the brighter more massive component, and the secondary component, the well-known Herbig Ae star HR\\,5999, is about 45\\arcsec{}.  The main interest to study the chemical composition of  HR\\,6000 comes from the generally poor understanding of the occurrence of abundance anomalies in such a young object with an estimated age of the order of 10$^{7}$ years.  In fact, the $\\Delta$\\,199 double system is located close to the center of the Lupus\\,3 molecular cloud which is populated by numerous T\\,Tauri stars. This has led to an assumption that this system has the same age as the cloud, which is estimated to be (9.1$\\pm$2.1)$\\times$10$^{6}$ years (James et al., 2006). However, after the Hipparcos mission, the membership of HR\\,5999 and HR\\,6000 to the cloud became rather questionable due to the uncertainties in the distance determination for the Lupus cloud. The distance of HR\\,5999 and HR\\,6000 measured by Hipparcos are 208$\\pm$38\\,pc and 240$\\pm$48\\,pc, respectively, while the Lupus cloud distance is 150$\\pm$10\\,pc according to Crawford (2000). Comer\\`on et al.\\ (2003) assigned a distance of about 200\\,pc to the Lupus cloud, but this determination was made by assuming a priori that $\\Delta$199 belongs to the cloud.  A ROSAT survey of Herbig Ae/Be stars presented by Zinnecker \\& Preibisch (1994) has detected strong X-ray emission coming from the direction of HR\\,6000. To explain the X-ray origin, a possible T\\,Tauri companion for HR\\,6000 was postulated by van den Ancker et al.\\ (1996) on the basis of an infrared excess in the energy distribution relative to predictions made for an effective temperature of \\teff=14000\\,K. Siebenmorgen et al.\\ (2000) fitted very short spectral regions of observed flux in the mid IR to a black body of 13000\\,K and noticed that the observed spectrum was featureless.  Neither spectrum nor radial velocity variability was found for HR\\,6000 by Andersen \\& Jaschek (1984), who studied optical spectra taken at different epochs. However, van den Ancker et al.\\ (1996) discovered long-term variations  in the u, v, and b Str\\\"omgren magnitudes with approximate amplitude of 0\\fm{}03$\\pm$0\\fm{}01 in u, 0\\fm{}02$\\pm$0\\fm{}01 in v, and 0\\fm{}01 in b. No variations in y were found.  Kurtz \\& Marang (1995) discovered  variations of 0.008 mag in V with a period near to 2\\,d and suggested that this period could be possibly caused by rotational variation in a spotted magnetic star. This would imply a pole on orientation for HR\\,6000 as they measured $v \\sin i$$\\le$5~km~s$^{-1}$.   Catanzaro et al.\\ (2004) were the first ones who provided abundances from the optical range. Previously,  optical spectra were studied only by Andersen \\& Jaschek (1984) and Andersen et al.\\ (1984) who identified the spectra from 3323\\,\\AA{} to 5317\\,\\AA{} and discussed the abundances on the basis of the line intensities. It is remarkable that the stellar parameters derived by Catanzaro et al.\\ (2004) from the Balmer profiles (\\teff=12950$\\pm$50~K, \\logg=4.05$\\pm$0.01) are 1000\\,K lower than those previously adopted and deduced from the photometry by Castelli et al.\\ (1985) (\\teff=14000\\,K, \\logg=4.0),  by Smith (1997) (\\teff=13990\\,K, \\logg=4.29) and by van den Ancker et al.\\ (1996) (\\teff=14000\\,K, \\logg=4.3).  In this paper we examine the whole optical spectrum of HR\\,6000 from 3050\\,\\AA{} to 9460\\,\\AA{} with the aim to extend the analysis performed on IUE spectra by Castelli et al.\\ (1985) to the visible region and to investigate the possible contamination of the stellar spectrum by a close T\\,Tau companion. We re-examine the parameter determination, the line identification, and the abundances. We also searched for the presence of peculiarities like emission lines and isotopic anomalies which were recently discovered in HgMn stars. Owing to the availability of spectra taken at different epochs we also investigated possible spectral variabilities.  The comparison of the observed and computed spectra as described in this paper is available at the web-address given in the footnote\\footnote{http://wwwuser.oat.ts.astro.it/castelli/hr6000/hr6000.html}. ",
        "conclusions": "The present study of HR\\,6000 has led to the following results: the measured radial velocity indicates that HR\\,6000 most likely belongs to the Lupus cloud so that its age is of the order of 10$^{7}$ years as it was found for the Lupus\\,3 cloud (James et al., 2006). Interstellar lines from the cloud can be observed in the spectrum from the ultraviolet to the visible, with displacements of the order of $-$0.05 $\\div$ $-$0.15\\,\\AA{} from the stellar lines.  UVES spectra observed with a six months time interval do not show spectrum variability for both stellar and interstellar lines, but a low variability of 1.2\\,km\\,s$^{-1}$ in the radial velocity  and marginal variabilities of the weak spectral features which are however at the level of the noise.   Different methods for the parameter determination have led to large differences in \\teff{}.  While the gravity values are found in the range between 4.1 and 4.4\\,dex with errors of the order of 0.1$\\div$0.2\\,dex, \\teff{} changes from 13900$\\pm$300\\,K, (if obtained from the UBV$\\beta$ and uvby$\\beta$ photometry), down to 12850$\\pm$50\\,K  when determined from the Balmer profiles, and can become as low as 11200$\\pm$300\\,K when derived from the  VI$_{c}$J photometry. Furthermore, as far as the Balmer profiles are concerned, the parameters which well reproduce H$_{\\beta}$, H$_{\\gamma}$, and H$_{\\delta}$ predict  too narrow H$_{\\alpha}$ wings. All these discrepancies in the parameter determinations could be accounted for by the presence of the Lupus cloud with could affect the colors in such a way to invalidate the standard reddening relations. Also a spectral contamination by a weak-emission T\\,Tauri (WTT) companion, as suggested by van den Ancker et al.\\ (1996), can not be excluded in an absolute way, although the only spectroscopic sign of its possible presence is a weak variable broad absorption at the position of \\ion{Li}{i} 6707\\,\\AA\\ and a somewhat better agreement between  observed and computed spectra longward of 6000\\,\\AA{} for a combined synthetic spectrum obtained from that  HR\\,6000  and one computed for \\teff{}=3500\\,K, \\logg=4.0, [M/H]=0.0. This companion could be either physically related with HR\\,6000 to form a close binary system or could be located in the foreground of HR\\,6000. Finally, the different parameters derived from the different determinations could be due to  the extremely peculiar nature of HR\\,6000 which may make the classical model atmospheres inadequate to reproduce all the observed stellar quantities. In fact, the abundance stratification inferred for some elements, helium in particular, would require more refined models in which the hypothesis of constant abundances throughout the atmosphere is dropped.   By using an ATLAS12 model with parameters \\teff=12850\\,K, \\logg=4.1, $\\xi$=0.0\\,km\\,s$^{-1}$ we have computed a synthetic spectrum for HR\\,6000 from 3050\\,\\AA\\ to 9460\\,\\AA\\ and have compared it with the observed spectrum. This comparison has shown that most of the elements lighter than Ca are significantly underabundant, except for Be, Na, and P. The Si underabundance ([-2.9]) is remarkable because it is even lower than that of 46 Aql ([$-$1.0]) which was claimed by Sadakane et al. (2001) to be the lowest one found in HgMn stars.  A striking peculiarity of HR\\,6000 is the lack of any overabundance for heavy elements with Z$>$40, except for Xe and Hg. The line spectrum of \\ion{Xe}{ii} is similar to that we observed in a preliminary analysis of UVES spectra of 46 Aql, another \\ion{Xe} rich star: while most of the lines lie at the laboratory wavelength, some other \\ion{Xe}{ii} lines seem to be shifted by about $-$0.1\\,\\AA\\ from the predicted position. The same behaviour can be observed also in HD\\,175640, although to a less extent, owing to its  lower \\ion{Xe} overabundance. Unfortunately, the small number of known transition probabilities and the ignorance of the \\ion{Xe}{ii} isotopic structure due to the lack of atomic data put strong limitations in the study of this element in the spectra of the CP stars.  The large overabundance  of [+0.7] for iron generates a very rich line spectrum in which numerous, still unclassified \\ion{Fe}{ii} lines are observed. A very large number of lines probably due to \\ion{Fe}{ii} remain unidentified. The similar iron overabundance of 46 Aql ([+0.65]) gives rise to an impressive  close resemblance between the spectra of the two stars.  Abundances in the ultraviolet obtained from a re-analysis of the IUE spectra studied by Castelli et al. (1985) agree with abundances derived from the UVES spectra  for all the elements, except for carbon and phosphorus. While the carbon identification in the optical spectrum is rather questionable owing to the weakness of the observed lines, the difference in the phosphorus abundance is similar to that yielded by the individual \\ion{P}{ii} and \\ion{P}{iii} lines.  There are several signs of vertical abundance stratification in HR\\,6000 for He, P, Mn, and Fe. The peculiar shape of the \\ion{He}{i} profiles is discussed for the first time in this paper. The profiles have cores too intense as compared to the wings, a fact indicating strong vertical He abundance stratification (Dworetsky, 2004; Bohlender, 2005). Furthermore, the inferred He abundance decreases with increasing wavelength ranging from an underabundance of [-0.8] at 4000\\,\\AA\\ to [-1.6] at 6000\\,\\AA.  For \\ion{Mn}{ii}, the abundance shortward of the Balmer discontinuity is larger by about 0.6\\,dex than that longward of the Balmer discontinuity; for phosphorus, the \\ion{P}{ii} abundance is larger by 0.2\\,dex than the \\ion{P}{iii} abundance, but this discrepancy lies within the mean square error of the average abundances;  for iron, the abundance from \\ion{Fe}{ii} high excitation lines is generally higher by about 0.2\\,dex than that from \\ion{Fe}{ii} low excitation lines.  Similar to other studied HgMn stars, also HR\\,6000 shows emission lines. The most numerous ones belong to \\ion{Mn}{ii}, followed by those of \\ion{Fe}{ii} and \\ion{Cr}{ii}. While the first two elements are overabundant in HR\\,6000, chromium has solar abundance. The star exhibits also isotopic anomalies for Hg and Ca. In both cases the most heavy isotope is the predominant one."
    },
    "0710/0710.3626_arXiv.txt": {
        "abstract": "Globular clusters produce orders of magnitude more millisecond pulsars per unit mass than the Galactic disk.  Since the first cluster pulsar was uncovered twenty years ago, at least 138 have been identified -- most of which are binary millisecond pulsars. Because of their origins involving stellar encounters, many of these systems are exotic objects that would never be observed in the Galactic disk.  Examples include pulsar---main sequence binaries, extremely rapid rotators (including the current record holder), and millisecond pulsars in highly eccentric orbits.  These systems are allowing new probes of the interstellar medium, the equation of state of material at supra-nuclear density, the mass distribution of neutron stars, and the dynamics of globular clusters. ",
        "introduction": "The first globular cluster (GC) pulsar was identified 20 years ago in the cluster M28 after intense efforts by an international team \\cite{lbm+87}.  Since then at least 138 GC pulsars\\footnote{For an up-to-date catalog of known GC pulsars, see Paulo Freire's website at \\url{http://www.naic.edu/~pfreire/GCpsr.html}}, the vast majority of which are millisecond pulsars (MSPs), have been found.  Finding these GC pulsars has required high-performance computing, sophisticated algorithms, state-of-the-art instrumentation, and deep observations with some of the largest radio telescopes in the world, primarily Parkes, Arecibo, and the Green Bank Telescope (GBT).  The payoff has been an extraordinarily wide variety of science.  Low-Mass X-ray Binaries (LMXBs) have been known to be orders-of-magnitude more numerous per unit mass in GCs as compared to the Galactic disk since the mid-1970s \\cite{kat75,cla75}.  This overabundance is due to the production of compact binary systems containing primordially-produced neutron stars via stellar interactions within the high-density cluster cores.  Since LMXBs are the progenitors of MSPs, this dynamics-driven production mechanism also applies to them, and it has made GCs (particularly the massive, dense, and nearby ones) lucrative targets for deep pulsar searches.  Camilo \\& Rasio \\cite{cr05} produced an excellent review of the first 100 GC pulsars in 2005.  This current review provides a significant update to Camilo \\& Rasio as it concentrates on the advances made (primarily with the GBT) within the past several years, including almost 40 additional pulsars and over 50 new timing solutions. ",
        "conclusions": ""
    },
    "0710/0710.2882_arXiv.txt": {
        "abstract": "We study the initial mass function (IMF) of one of the most massive Galactic star-forming regions NGC 3603 to answer a fundamental question in current astrophysics: is the IMF universal, or does it vary? Using our very deep, high angular resolution $JHK_{S}L'$ images obtained with NAOS-CONICA at the VLT at ESO, we have successfully revealed the stellar population down to the subsolar mass range in the core of the starburst cluster. The derived IMF of NGC 3603 is reasonably fitted by a single power law with index $\\Gamma \\sim -0.74$ within a mass range of $0.4 - 20$ \\msun, substantially flatter than the Salpeter-like IMF. A strong radial steepening of the IMF is observed mainly in the inner $r \\lesssim 30''$ field, indicating mass segregation in the cluster center. We estimate the total mass of NGC 3603 to be about $1.0 - 1.6 \\times 10^4$ \\msun. The derived core density is $\\geq 6 \\times 10^4$ \\msun pc$^{-3}$, an order of magnitude larger than e.g., the Orion Nebula Cluster. The estimate of the half-mass relaxation time for solar-mass stars is about $10 - 40$ Myr, suggesting that the intermediate- and low-mass stars have not yet been affected significantly by the dynamical relaxation in the cluster. The relaxation time for the high-mass stars can be comparable to the age of the cluster. We estimate that the stars residing outside the observed field cannot steepen the IMF significantly, indicating our IMF adequately describes the whole cluster. Analyzing thoroughly the systematic uncertainties in our IMF determination, we conclude that the power law index of the IMF of NGC 3603 is $\\Gamma = -0.74^{+0.62}_{-0.47}$. Our result thus supports the hypothesis of a potential top-heavy IMF in massive star-forming clusters and starbursts. ",
        "introduction": "\\label{introduction} One of the most interesting properties of massive star-forming regions is the stellar initial mass function (IMF). Since the pioneering work by \\citet{sal55}, which led to a standard picture of the IMF, the so-called Salpeter IMF ($dN/d\\log \\mathcal{M} \\propto \\mathcal{M}^\\Gamma$ with $\\Gamma = -1.35$, where $\\mathcal{M}$ is stellar mass), numerous efforts have been made to understand the IMF of many types of objects such as field populations, stellar associations, and open and globular clusters, covering the whole stellar mass range from massive OB stars down to substellar brown dwarfs. It is currently accepted that the IMF follows a single Salpeter-like power law in the high- to intermediate-mass range ($\\mathcal{M} \\gtrsim 1$ \\msun). It becomes flatter towards subsolar masses and peaks at a characteristic mass of several tenths of a solar mass. It then declines towards the brown dwarf mass range. Therefore, several analytical expressions have been proposed for the standard IMF, for example, a lognormal distribution \\citep{mil79,sca86}, a segmented power law distribution \\citep{sca98, kro01}, and a combination of both \\citep{cha03}. As these representations of the IMF fit reasonably well the various types of stellar populations, the idea of the \\textit{universality} of the IMF was born. A universal IMF basically suggests a universal star formation mechanism. However, it is somewhat intuitive that different physical conditions such as the density, velocity fields, chemical composition, and tidal forces in the natal molecular clouds can lead to different star formation processes and, consequently, some variability in the IMF. Indeed, there is growing evidence for the variable IMF. There are significant variations in the power law index and the characteristic masses at which the distribution shows the peak or at which the power law breaks. A typical example of an IMF that does not follow the Salpeter-like distribution is the so-called top-heavy IMF in starburst galaxies and young, massive star-forming regions.  Does the IMF really vary among stellar populations? To answer this question, we need to clarify, for at least some stellar populations, if the observed IMF variations are without any doubt true deviations from the Salpeter-like IMF or if they can simply be accommodated by the combined effects of observational, theoretical, and statistical uncertainties? Many studies have so far addressed this question. In order to address the question and to obtain new insights into the IMF of an intense starburst environment, we study the IMF of the massive star-forming HII region \\objectname{NGC 3603} based on unprecedented spacial resolution observations of its central starburst cluster obtained by NACO at the Very Large Telescope (VLT), as well as a wider field with the Infrared Spectrometer and Array Camera (ISAAC) at VLT at the European Southern Observatory (ESO). Compared to extragalactic starbursts (e.g. \\objectname{M82}), young star-forming regions in the Milky Way and in the Magellanic Clouds -- so-called nearby starburst templates -- are close enough to resolve the individual stars with current powerful adaptive optics (AO)-assisted ground-based telescopes and space-based facilities. These objects include, for example, the \\objectname{Arches} cluster, the \\objectname{Quintuplet} cluster, \\objectname{R136}, and \\objectname{NGC 3576}. More nearby but less massive star-forming regions are the \\objectname{Trapezium} cluster in the Orion Nebula Cluster (ONC), the \\objectname{Pleiades} cluster, the \\objectname{Taurus} cluster, and \\objectname{IC 348}.  The IMF of NGC 3603 has been presented in several studies, and recent works have derived somewhat flat IMFs with a power law index of $\\Gamma = -0.73$ for $1 - 30$ \\msun\\ in \\citet{eis98}, $\\Gamma = -0.9$ for $2.5 - 100$ \\msun\\ in \\citet{sun04}, and $\\Gamma = -0.91 \\pm 0.15$ for $0.4 - 20$ \\msun\\ in \\citet{sto06}. As another example of the IMF of massive Galactic star clusters, the Arches cluster has shown a slightly flat IMF with $\\Gamma = -0.6$ to $-1.1$ for intermediate- and high-mass stars \\citep{fig99,sto05,kim06}. Studies of the IMF of the Arches and other clusters are presented in \\S~\\ref{c-discussion}.  \\subsection{NGC 3603} NGC 3603 is one of the most luminous, optically visible HII regions in the Milky Way \\citep{gos69} with its global properties such as $L$(H$_{\\alpha}$) $\\sim 1.5 \\times 10^{39}$ ergs s$^{-1}$ \\citep{ken84}, and the total mass of molecular clouds $\\sim4 \\times 10^{5}$ \\msun\\ \\citep{gra88}. Its total bolometric luminosity is as large as 10$^{7}$ $L_{\\odot}$, 2 orders of magnitude larger than that of the ONC, and just an order of magnitude smaller than the other well known starburst template R136 in 30 Dor in the Large Magellanic Cloud (LMC). In particular, the starburst cluster located in the northern part of the gigantic HII complex including \\objectname{HD 97950} -- a very compact Trapezium-like system with plenty of high-mass components \\citep{wal73} -- at its center, is one of the most massive and the densest star-forming regions known in the Galaxy. The starburst cluster consists of three WNL stars, six O3 stars, and many other late O- and B-type stars in a volume of less than a cubic light year, providing most of the ionizing radiation in the giant HII region \\citep{mof83, cla86, mof94, dri95,hof95, cro98}. It thus shows remarkable similarities with R136. The distance of NGC 3603 from the Sun has been estimated in many earlier works by means of both photometric and kinematic analysis \\citep[e.g.][]{vdb78,mel82}, and the currently accepted value is about 7 $\\pm$ 1 kpc. Owing to its intrinsic properties such as its proximity, relatively low visual extinction of only $A_{V} = 4 - 5$ mag, and extreme compactness and brightness, NGC 3603 is one of the most suitable Galactic templates of starburst phenomena in distant galaxies. ",
        "conclusions": "\\label{c-summary} Our study is aiming at a fundamental question of current star formation research. Is the IMF universal, or does it vary with environment? To answer this question, we measured the IMF of NGC 3603 -- one of the most massive Galactic star-forming regions -- from our NIR observations with the AO system NACO at the VLT/ESO. In the following we summarize the main results from our study.  \\begin{enumerate} \\item \\textbf{NIR Photometry}\\\\ Our very deep, high angular resolution $JHK_{S}L'$-band images obtained by NACO show unprecedented details of the core of the starburst cluster in NGC 3603. Together with the wider field ISAAC $JHK_{S}$ images, we could successfully derive magnitudes and positions of almost 10,000 stars in the dense cluster up to $r \\sim 110''$ covering the mass range from the most massive stars down to $\\sim0.4$ \\msun. The brightest 256 stars in the NACO images of the inner $r \\leq 13''$ are listed and cross-identified with potential counterparts in previous studies.  \\item \\textbf{Age, extinction, and disk fraction}\\\\ Based on the fitting of the stellar evolution models in the CMDs and CCDs, we derived the age of the PMS stars of $0.5 - 1.0$ Myr and the upper limit for the age of the MS stars of $\\sim2.5$ Myr, suggesting a slight age spread in the cluster. The derived average foreground extinction is $A_{V} = 4.5 \\pm 0.5$ mag, and the foreground extinction increases by $\\Delta{A_{V}} \\sim 2.0$ mag towards larger radii ($r \\gtrsim 55''$). Using the $K_{S} - L'$ versus $J - H$ CCD, we derived a circumstellar disk fraction of $\\sim25 \\pm 10$\\% for stars with a mass of $\\geq 0.9$ \\msun\\ in the central cluster ($r \\leq 10''$).  \\item \\textbf{IMF and its radial variation}\\\\ Applying the field star rejection and the incompleteness correction, the \\textit{K}LF for 7514 stars simultaneously detected in the $JHK_{S}$ bands follows a power law with no obvious turnover or truncation within the detection limit of $m_{K_{S}} \\sim 17.4$ mag (based on the $J$-band 50\\% completeness of $\\sim19.4$ and the typical color of $J - K_{S} \\sim 2$). Within the mass range of $0.4 - 20$ \\msun\\, the IMF is well described by a single power law with a power law index $\\Gamma \\sim -0.74$. We found the power law index decreasing from $\\Gamma \\sim -0.3$ at $r \\leq 5''$ to $\\Gamma \\sim -0.8$ at $r \\sim 30''$. The strong steepening occurs in the inner $r \\lesssim 13''$, pointing towards mass segregation in the very center of the cluster. No significant variation of the IMF is found for larger radii ($r \\gtrsim 30''$).  \\item \\textbf{Size, mass, and dynamical status}\\\\ Fitting a King model to the radial density profile of stars with a mass of $0.5 - 2.5$ \\msun, we derived a core radius of $\\sim4''.8$ ($\\sim0.14$ pc at $d \\sim 6$ kpc). As the radial density decreases even at the limits of our field of view, we can give a firm lower limit of $r = 110''$ ($\\sim3.2$ pc) for the cluster size. We also derive an upper limit of $r = 1260''$ ($\\sim37$ pc) for the tidal radius of the cluster. The de-projected King model allowed us to extrapolate to the total mass of NGC 3603. Assuming a single power law IMF with index $\\Gamma = -0.74$ within the mass range of $0.1 - 100$ \\msun, we found a total mass of about $1.0 - 1.6 \\times 10^4$ \\msun. The half-mass radius is found to be within $25'' - 50''$ ($0.7 - 1.5$ pc). The derived core mass density of the cluster is $\\geq 6 \\times 10^4$ \\msun\\ pc$^{-3}$.  We estimate a half-mass relaxation time of approximately $10 - 40$ Myr for stars with a typical mass of 1 \\msun, an order of magnitude larger than the age of the PMS population in the cluster ($\\lesssim1$ Myr). This implies that the intermediate- and low-mass stars have not yet experienced significant dynamical relaxation. However, the relaxation time of the high-mass stars is expected to be an order of magnitude shorter and is comparable to the cluster age. We could thus not conclude if the observed mass segregation of the high-mass stars is caused by dynamical evolution or if it is primordial. Indeed it can be due to the combination of them. We compute that our images with a maximum radius of $r \\sim 110''$ cover at least $\\sim67$\\% of intermediate- and low-mass stars of NGC 3603. The stars outside the observed field cannot steepen the IMF by more than $\\Delta\\Gamma \\lesssim 0.16$. Considering also the fact that the IMF does not significantly change beyond $r \\gtrsim 30''$, we conclude that the observed IMF is representative for the whole NGC 3603 stellar cluster, irrespective of the mass segregation in the very center.  \\item \\textbf{Systematic uncertainties of the IMF}\\\\ We thoroughly analyzed the systematic errors in the IMF determination. In particular we derived the errors from uncertainties in the age, distance, foreground extinction, stellar evolution models, metallicity, individual extinction, incompleteness correction, cluster membership, unresolved binaries, and the stellar mass outside the observed field. Combining all errors we derived the power law index of $\\Gamma = -0.74^{+0.62}_{-0.47}$. Assuming a Gaussian probability distribution, we conclude that the probability that the IMF is as steep as the Salpeter IMF ($\\Gamma = -1.35$) is less than $\\sim10$\\%. \\end{enumerate}  Our result thus supports the hypothesis of a top-heavy IMF in massive star-forming clusters and starburst galaxies. One potentially common property among such clusters showing flat IMFs would be the high stellar density in the core of the cluster."
    },
    "0710/0710.0763_arXiv.txt": {
        "abstract": "Magnetic fluctuations in the solar wind are distributed according to Kolmogorov's power law $f^{-5/3}$ below the ion cyclotron frequency $f_{ci}$. Above this frequency, the observed steeper power law is usually interpreted  in two different ways:  a dissipative range of the solar wind turbulence or another turbulent cascade, the nature of which is still an open question. Using the Cluster magnetic data we show that after the spectral break the intermittency increases toward higher frequencies, indicating the presence of non-linear interactions inherent to a new inertial range and not to the dissipative range. At the same time the level of compressible fluctuations  raises. We show that the energy transfer rate and intermittency are sensitive to the level of compressibility of the magnetic fluctuations within the small scale inertial range. We conjecture that the time needed to establish this inertial range is shorter than the eddy-turnover time, and is related to dispersive effects. A simple phenomenological model, based on the compressible Hall MHD, predicts the magnetic spectrum $\\sim k^{-7/3+2\\alpha}$, which depends on the degree of plasma compression $\\alpha$. ",
        "introduction": "\\label{Intro}  Solar wind, which is highly turbulent, represents a unique opportunity  to investigate turbulence in natural plasmas via \\emph{in situ} measurements \\cite{sw2,noi}. In non-magnetized fluids, where the energy injection scale is far from the dissipation one, the intermediate scales (inertial range) are described by the universal power law Kolmogorov spectrum $k^{-s}$ with $s=5/3$. This law depends neither on the energy injection nor on the energy dissipation processes.  In the solar wind, and in the interplanetary space in general, the mean free path roughly corresponds to the Sun-Earth distance and the usual dissipation via collisions is negligible. At the same time, in a  magnetized plasma there is a number of characteristic space and temporal scales. Investigating solar wind turbulence at these scales is then a challenging topic from the point of view of basic plasma physics. Here we will focus our discussion on the ion scales, namely the ion inertial length $\\lambda_i = c/\\omega_{pi}$ and the ion cyclotron frequency $f_{ci}= eB/m_i$. At these scales the fluid-like approximation of plasma dynamics, the usual Magnetohydrodynamic (MHD) description, breaks down in favor of a more complex description of plasma.  Solar wind turbulent spectrum of magnetic field fluctuations follows a  $\\sim f^{-5/3}$ power law below the ion cyclotron frequency $f_{ci}$. For $f>f_{ci}$ the spectrum steepens significantly, but is still described by a power law $f^{-s}$, with $s \\in (2-4)$ \\citep{Leamon98,smith06}.  \\emph{In situ} solar wind measurements provide time series data, i.e. information on frequency in Fourier space. How to get any information on wave vectors? For fluctuations with velocities much smaller than the plasma bulk velocity $V$, the Taylor hypothesis is valid: the observed  variations on a time scale $\\delta t$ correspond to variations on the spatial scale $\\delta r=V \\delta t$. Therefore there is a direct correspondence between $f$ and $k$ spectra. % If the solar wind turbulence was a mixture of linear wave modes, Taylor hypothesis would be well verified only below the spectral break: the bulk velocity is superalfv\\'enic ($V>V_A$, where $V_A$ being the Alfv\\'en speed) and so the low frequency fluctuations can be considered as frozen in plasma. However, whistler waves (with $f > f_{ci}$ and phase speed $V_{\\varphi} > V_A$) do not satisfy the Taylor hypothesis assumption. In this study we assume that in the solar wind there are no whistler waves above the spectral break frequency. This last assumption is supported by results recently obtained in the Earth's magnetosheath \\citep{Mangeney2006}.   Using the Taylor hypothesis the observed solar wind spectrum below the break is usually attributed to the Kolmogorov's spectrum $\\sim k^{-5/3}$. Above the spectral break, the spectral steepening, $\\sim k^{-3}$, can be interpreted in two different ways. Some authors associate it to the dissipation range \\citep{Leamon98,Leamon99,Leamon2000,Bale2005,smith06}. Others suggest that after the spectral break another turbulent cascade takes place \\citep{Biskamp1996,Ghosh96,stawicki,Li01,Galtier06}.  In ordinary fluid flows the dissipation range is described by an exponential function \\citep{frisch}. While in the solar wind, a well-defined power law is observed after the break point and not an exponential. Note that power spectra, i.e. the second order statistics, completely describe Gaussian, or statistically independent, fluctuations. However, as is well known, fluctuations cannot be described by a Gaussian statistics in the low-frequency part of the solar wind turbulence \\citep{noi}. Deviations from Gaussianity, i.e. intermittency \\citep{frisch}, may be quantified by the flatness, the forth-order moment of fluctuations.   In this paper we investigate the nature of magnetic fluctuations in the high frequency range of the solar wind turbulence, that is usually called dissipation range. We find that in this range the flatness increase with frequency. This is similar to what is going on in the low-frequency range.  The presence of intermittency together with the well defined power law in the high-frequency part of the spectrum suggests another turbulent cascade rather than a dissipation range. This small scale cascade is observed to be much more compressible than the Kolmogorov-like inertial range, in agreement with the previous observations by the Wind spacecraft \\citep{Leamon98}. We show that the energy transfer rate and intermittency are sensitive to the level of compressibility of the turbulent fluctuations in this range. Finally, we propose a simple phenomenological model, based on the compressible Hall MHD, which allows to explain the observed range of the spectral indices in the high frequency part of the solar wind spectrum. ",
        "conclusions": "In this paper we investigate small scales turbulent fluctuations in the solar wind (i.e. at frequencies above the spectral break at the vicinity of $f_{ci}$). Taking into account the Taylor hypothesis ($\\ell=V/f$) the frequency domain above the spectral break covers $\\sim (40- 2000)$~km space scales that corresponds to $(0.3- 15)\\lambda _i$. In this range we found evidences for strong departure from Gaussian statistics and the presence of intermittency while the spectrum presents a well-defined power law. Both the presence of a power-law spectrum and the absence of global self-similarity, seems to be quite in contrast with the role of ``dissipative range\".  In usual fluid turbulence, the dissipative range  \\citep{frisch} starts with a rough exponential cutoff; in the near dissipation range the intermittency increases as far as   the Gaussian fluctuations dissipate faster than the coherent structures \\citep{near-diss}; then the fluctuations become self-similar, the singularities being smoothed by dissipation. In the solar wind turbulence we observe a completely different picture. After the spectral break in the vicinity of $f_{ci}$ the flatness increases as a power law indicating that non-linear interactions are at work to build up a new inertial range.  This small scale cascade is much more compressible than the lower frequency Alfv\\'enic cascade. This sudden change in nature of the turbulent fluctuations can happen due to a partial dissipation of magnetic fluctuations at the spectral break: the left-hand Alfv\\'enic fluctuations with $k_{\\|}\\gg k_{\\perp}$ are damped by the ion-cyclotron damping \\citep{Ghosh96}. Above the break, however, a new \\emph{magnetosonic cascade} takes place up to the electron characteristic scales. This energy cascade is seems to be dominated by the fluctuations with $k_{\\perp}\\gg k_{\\|}$ \\citep{fouad06,Mangeney2006}.  We found, as well, that the  plasma compressibility controls the statistics of the magnetic field fluctuations. Preliminary results show that the increase of the level of the plasma compressibility leads to the spectrum steepening and increase of the intermittency.  To explain this small scale compressible cascade we introduce here for the first time a simple phenomenological model based on the compressible Hall MHD: we find that the magnetic energy spectrum follows a $E(k) \\sim k^{-7/3 + 2\\alpha}$ law, depending on the degree of plasma compressibility $\\alpha$. While far from a complete model of small scale turbulence, this simple model can explain observed variations of the spectral index within the high frequency part of the solar wind turbulent spectrum \\citep{Leamon98,smith06} by different degree of plasma compressibility."
    },
    "0710/0710.5831_arXiv.txt": {
        "abstract": "{} {Microquasars are ideal natural laboratories for understanding accretion/ejection processes, studying the physics of relativistic jets, and testing gravitational phenomena. Nevertheless, these objects are difficult to find in our Galaxy. The main goal of this work is to increase the number of known systems of this kind, which should allow better testing of high-energy phenomena and more realistic statistical studies of this galactic population to be made.} {We have developed an improved search strategy based on positional cross-identification with very restrictive selection criteria to find new MQs, taking advantage of more sensitive modern X-ray data. To do this, we made combined use of the radio, infrared, and X-ray properties of the sources, using different available catalogs.} {We find 86 sources with positional coincidence in the NVSS/XMM catalogs at galactic latitudes $|$b$|$ $\\leq$ 10$^{\\circ}$. Among them, 24 are well-known objects and the remaining 62 sources are unidentified. Out of 86 sources, 31 have one or two possible infrared counterparts in the 2MASS catalog. For the fully coincident sources, whenever possible, we analyzed color-color and hardness ratio diagrams and found that at least 3 of them display high-mass X-ray binary characteristics, making them potential microquasar candidates.} {} ",
        "introduction": "Some of the most attractive objects in the Galaxy are the enigmatic microquasars (MQs). On a smaller scale, they copy the characteristics exhibited by distant quasars \\cite{mirabel99}. These systems are X-ray binaries (XRBs) containing compact objects like stellar black holes or neutron stars that accrete matter from a companion star. They are known to emit from radio to X-ray energies \\cite{mirabel94} and possibly up to TeV gamma-ray energies, as in the case of Cygnus X-1 \\cite{albert07}. Evidence that jets of accreting X-ray binaries can accelerate particles to TeV  energies has nevertheless been presented by Corbel et al. (2002). MQs combine two important aspects of relativistic astrophysics: accreting black holes or neutron stars identified by the production of hard X-rays around accreting disks and relativistic radio jets detected by means of their synchrotron emission.  These binary systems are ideal natural laboratories for understanding accretion/ejection process and other gravitational phenomena. However, they seem to be rare objects in our Galaxy. In order to enable more robust statistical studies, it is necessary to increase the number of known MQs. Of the 15 currently confirmed MQs in the galaxy, 6 belong to the high-mass X-ray binary (HMXB) class and 9 are of the low-mass X-ray binary (LMXB) kind \\cite{paredes05}.  Finding new MQs candidates is not an easy task. Considerable effort in the past has been put in to increasing their number. A few searches for radio-emitting XRB systems and therefore new MQ candidates, based on cross-identifications between radio, infrared, and X-ray catalogs, have been carried out in the past, but without enough success \\citep[e.g.][]{paredes02,ribo04}. The method of looking for such objects in the Galaxy usually includes a number of steps with very restrictive selection criteria for the sources being investigated. Mainly, a number of competing emission mechanisms and several physical parameters should be associated to the same source or system, and they basically include the properties of jet emission and XRB behavior. The detection of emission at radio wavelengths could be the signature that such  relativistic jets, mainly emitting via incoherent synchrotron emission from very high-energy electrons spiraling in magnetic fields, are present in the object. Ultraviolet and infrared emissions are characteristics displayed by the normal star companions, and X-ray radiation is most efficient at revealing accretion-powered sources, such as binary stars.  As an example, a previous list of MQ candidates obtained from cross-correlation of the radio NVSS catalog \\cite{condon98} and the ROSAT Bright Source Catalog (RBSC) sources of the Galactic plane was presented by Paredes et al. (2002). They found 35 possible MQ candidates based on a hardness ratio criterion  i.e. $HR1+\\sigma(HR1)\\geq 0.9$, where $HR1=([0.5-2.0 {\\rm ~keV}]-[0.1-0.4 {\\rm ~keV}])/([0.5-2.0 {\\rm ~keV}]+[0.1-0.4 {\\rm ~keV}])$. However, since ROSAT soft X-rays are strongly absorbed by the interstellar medium, their resulting list of MQ candidates could be fairly reduced. Given the impossibility that ROSAT will detect X-ray photons at energies above 2.4 keV, and because X-rays in MQs are expected to be highly energetic, such a criterion could not be efficient enough to select possible MQ candidates.  Here, we develop an improved search strategy that is also based on very restrictive, but improved, selection criteria aimed at finding new MQs in the Galaxy. Moreover, we take advantage of modern and more sensitive multiwavelength data. In Sect.~\\ref{search} we describe the search strategy and sample definition. The general analysis and statistical results are presented in Sect.~\\ref{results}. Finally, we summarize our main conclusions in Sect.~\\ref{summary}. ",
        "conclusions": "\\label{summary}  In this work we have presented an improved search strategy for finding new MQ candidates based on positional cross-identifications. By analyzing radio and modern more sensitive X-ray data, we found 86 positional-coincidence sources in the NVSS and XMM catalog, with galactic latitudes $|b| \\leq 10^{\\circ}$. Among them, 24 are well-known objects and the rest of the 62 sources are unidentified. Out of the 86 sources 31 have one or two infrared counterpart candidates in the 2MASS catalog. For all the sources, when possible, the hardness ratio was used to assess the low-mass or high-mass XRB properties of the objects. At least 29 unidentified objects fulfill these characteristics. Using the 2MASS counterpart of well-known MQs, we found that MQs with a high-mass nature follow a lineal behavior in the infrared color-color diagram, while LMXB MQs display a more random distribution. Based on this information, we suggest a possible region in the infrared diagram for new MQ candidates. As a result, we found 3 objects that we propose as likely galactic HMXBs and MQ candidates. A follow-up study of these three sources will be reported on a future paper (Combi et al. 2007).  Finally, it is important to mention that the whole analysis carried out here only covers a minute fraction of the expected MQ candidates. The NVSS survey does not cover all of the sky. It is complete up to $-40^{\\circ}$ in declination. In addition, the current XMM catalog only represents $\\sim$ 2\\% of the galactic plane ($|$b$|$ $\\leq$ 10$^{\\circ}$). Thus, a naive extrapolation of the 3 MQ candidates to the whole Galactic plane suggests that about 120 new candidates could be expected in the future."
    },
    "0710/0710.0139_arXiv.txt": {
        "abstract": "Recent observations of the Galactic center revealed a nuclear disk of young OB stars, in addition to many similar outlying stars with higher eccentricities and/or high inclinations relative to the disk (some of them possibly belonging to a second disk). Binaries in such nuclear disks, if they exist in non-negligible fractions, could have a major role in the evolution of the disks through binary heating of this stellar system. We suggest that interactions with/in binaries may explain some (or all) of the observed outlying young stars in the Galactic center. Such stars could have been formed in a disk, and later on kicked out from it through binary related interactions, similar to ejection of high velocity runaway OB stars in young clusters throughout the galaxy. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.0625_arXiv.txt": {
        "abstract": "In a previous paper, we modeled the oscillations of a thermally-supported (Bonnor-Ebert) sphere as non-radial, linear perturbations following a standard analysis developed for stellar pulsations. The predicted column density variations and molecular spectral line profiles are similar to those observed in the Bok globule B68 suggesting that the motions in some starless cores may be oscillating perturbations on a thermally supported equilibrium structure. However, the linear analysis is unable to address several questions, among them the stability, and  lifetime of the perturbations. In this paper we simulate the oscillations using a three-dimensional numerical hydrodynamic code.  We find that the oscillations are damped predominantly by non-linear mode-coupling, and the damping time scale is typically many oscillation periods, corresponding to a few million years, and persisting over the inferred lifetime of gobules. ",
        "introduction": "Recent observations  suggest that the internal structure of some of the small isolated molecular clouds known as starless cores or Bok globules \\citep{Bok1948} is well approximated by a balance of thermal and gravitational forces on which is superposed a pattern of non-radial oscillations \\citep{Lada2003}. Previously, the internal oscillations of these starless cores have been noted in simulations of collapse \\citep{Hennebelle2003}, magnetized clouds \\citep{Galli2005}, modeled as one-dimensional radial oscillations by means of a numerical simulation that includes radiative energy losses \\citep{KetoField2005} and as three-dimensional non-radial perturbations assuming isothermal conditions \\citep{Keto2006}. Analysis of these models and comparison of simulated molecular spectral line profiles with those observed in the dark clouds L1544  \\citep{Caselli2002a} and B68 \\citep{Lada2003} indicate that the motions in some dark clouds are consistent with large-scale hydrodynamic oscillations, i.e., sound waves. Nevertheless, a number of questions remain as yet unanswered.  Is the underlying assumption of the oscillations as perturbations valid for the amplitudes required to produce the observed velocities? Based on the observations of B68, the perturbation amplitudes in our previous three-dimensional model were $\\sim25$\\%. Does the linear model correctly predict the velocity field for such large amplitudes or would significant differences appear in the non-linear regime?   What is the lifetime of the perturbations and thus the velocities within the clouds?  For hydrodynamic oscillations to remain a viable explanation of the velocity field observed in B68 they must have lifetimes in excess of a crossing time.  Supersonic oscillations are strongly damped via shocks. However, in the subsonic regime the dominant damping mechanism is less clear.  In the one-dimensional model \\citep{KetoField2005}, the dissipation of the oscillations occurs via radiative losses with an estimated timescale of $10\\,\\Myr$.  In the non-radial perturbative model \\citep{Keto2006} damping is explicitly ignored. For mode amplitudes larger than linear perturbations, non-linear mode-mode coupling provides an additional potentially significant damping mechanism.  Are the oscillations stable or growing? Are there preferred modes? In general we expect the oscillations to damp over time, and we expect that the longest-lived modes of oscillation would be the long wavelength modes. Neither the one-dimensional non-linear analysis nor the three-dimensional linear analysis is able to adequately address this question.  In this paper we address these questions by modeling the oscillations of dark clouds with a three-dimensional, non-linear, numerical hydrodynamic simulation. We find that the dominant damping mechanism is indeed non-linear mode coupling, the longest wavelength mode lifetimes are on the order of a few $\\Myr$ and produce velocity and density fields qualitatively similar to those observed.  Section \\ref{CM} summarizes the numerical methods employed, sections \\ref{QP} and \\ref{DP} discuss the viability of quadrupole and dipole oscillation models, and the conclusions are contained in section \\ref{C}. ",
        "conclusions": "\\label{C} Despite their strongly non-linear nature, large-amplitude oscillations of pressure-supported molecular clouds reproduce many of the qualitative features of observations of starless cores.  These include the complex velocity structures and highly distorted column densities found in B68 and L1544, as well as mimicking rotation, expansion, and contraction depending upon the mode structure and viewing angle.  This may complicate efforts to discern the global dynamical state of these objects from observations of molecular line shapes alone.  However, it may be possible to distinguish oscillations from collapse by comparing the spectral line profiles and velocities of volatile molecules such as N2H+ and NH3 against those of more refractory species such as CO, HCO+, or CS \\citep{KetoField2005}.  In a simple breathing mode oscillation, the gas velocities are highest towards the outer regions of the cloud where the abundance of the carbon species is not depleted by freeze-out.  In contrast, in the gravitational collapse of unstable BE spheres, the gas velocities are highest in the center of the cloud where the emission from the non-depleting nitrogen molecules is highest.  We note that we have restricted our attention to hydrodynamic oscillations of pressure-bounded, thermally-supported spheres. This description seems appropriate for globules such as B68 that are surrounded by hotter gas and have density profiles that closely match those of Bonnor-Ebert spheres. Whether the Bonnor-Ebert description is appropriate for embedded cores remains to be determined. For example, if magnetic stresses contribute significantly to cloud support, then additional classes of oscillations exist, associated with the magnetic field \\citep[see,\\eg,][]{Hennebelle2003,Galli2005}.  The lifetimes of hydrodynamic oscillations of starless cores like B68 are sufficiently long (few $10^6$ yr) to be a significant fraction of the expected lifetime of a molecular cloud ($10^7$ yr).  The oscillation lifetimes are dominantly limited by non-linear mode-mode coupling.  Nevertheless, large amplitude oscillations persist for many periods (and thus many sound crossing times), and are therefore not expected to be particularly rare.  There are no signs of cloud fragmentation, and thus amplitudes comparable to those discussed here are insufficient to produce binaries without the inclusion of additional physical processes."
    },
    "0710/0710.4094_arXiv.txt": {
        "abstract": "A subgroup of dwarf galaxies have characteristics of a possible evolutionary transition between star-forming systems and dwarf ellipticals. These systems host significant starbursts in combination with smooth, elliptical outer envelopes and small HI content; they are low on gas and unlikely to sustain high star formation rates over significant cosmic time spans.  We explore possible origins of such starburst ``transition\" dwarfs using moderately deep optical images.  While galaxy-galaxy interactions could produce these galaxies, no optical evidence exists for tidal debris or other outer disturbances, and they also lack nearby giant neighbors which could supply recent perturbations.  Colors of the outer regions indicate that star formation ceased $>$1~Gyr in the past, a longer time span than can be reasonably associated with the current starbursts.  We consider mechanisms where the starbursts are tied either to interactions with other dwarfs or to the state of the interstellar medium, and discuss the possibility of episodic star formation events associated with gas heating and cooling in low specific angular momentum galaxies. ",
        "introduction": "The formation and evolution of dwarf galaxies remains poorly understood.  How these processes produce both rotationally supported (dIrr) as well as dynamically warm (dE/dS0s) dwarf systems remain puzzles, despite considerable effort.  An outstanding question in dwarf galaxy evolution is whether there is an evolutionary connection between the various morphological classes of dwarf galaxies (e.g. van den Bergh 1977, van Zee et al. 1998, van Zee et al. 2001, Lisker et al. 2007).  Specifically, can dwarfs of one morphological type evolve by some process into a different type?  For example, do some dIrrs evolve into dE/dS0s or did the dE/dS0s we see in the Universe originally form as that type?  These questions regarding the formation and evolution of dwarf galaxies are of particular importance in light of hierarchical cold dark matter (CDM) galaxy formation models which predict large numbers of dwarf galaxies and suggests that larger galaxies are built up from the accretion of many low mass halos (e.g. White \\& Frenk, 1991).  A related issue concerns the impact of star formation on dwarf galaxy evolution and particularly the role of episodic star formation (e.g. Searle, Sargent \\& Bagnuollo 1973, Lee et al. 2002).  How does star formation relate to morphology in low mass galaxies?  For example, where do blue compact dwarfs (BCDs) fit in the dwarf galaxy zoo as starbursting objects and are they related to either dIrrs or dE/dS0s (e.g. Sung et al. 1998, Gil de Paz \\& Madore 2005, Noeske et al. 2005)?  In particular, it has been suggested that BCD galaxies may evolve into dE/dS0 galaxies after they lose their gas, either through supernovae-driven winds during episodes of intense star formation (Marlowe et al. 1999, Tajiri \\& Kamaya 2002), or through stripping processes induced by galaxy-galaxy interactions (Drinkwater \\& Hardy 1991).  Or, are some of these objects dEs where star formation has been renewed through gas capture as suggested by Silk, Wyse, \\& Shields (1987)?  Previous studies, which have concentrated mainly on optical morphological differences between dwarf galaxies and to some extent on their neutral gas content, have not firmly established a direct evolutionary connection between dIs and dEs in galaxy groups and other low density environments (e.g. van Zee, Salzer \\& Skillman 2001, Noeske et al. 2005).  Even in the Local Group the situation is unclear.  Many Local Group dwarfs have complex star formation histories and kinematic peculiarities, and at least five of the faint dwarfs possess optical morphologies which place them in a ``transition\" or ``mixed-type\" structural category between dI and dE/dS0 (e.g. Mateo 1998, Grebel, Gallagher, \\& Harbeck 2003). While nearby dE galaxies, such as NGC~185 and NGC~205 in the local group, support star formation, the rates are extremely low and consistent with these objects' small gas supplies.  More distant galaxy samples also contain dwarfs with early-type structures that also host HI gas and often also some young stars. As in the Local Group, these transition dwarfs frequently are small, low luminosity objects containing $\\lesssim$ 10$^6$ M$_{\\odot}$ of HI and having correspondingly low star formation rates (e.g. Bouchard et al. 2005).  On the other hand, the Virgo Cluster of galaxies appears to contain a complete sequence of more luminous dwarfs with declining and increasingly concentrated star formation, extending from typical BCDs through to blue core dEs (Gallagher \\& Hunter 1989, Lisker et al. 2006,2007).  While moderate luminosity galaxies with early-type outer structures and young stellar populations in their centers also exist in less dense environments (e.g. NGC~404, del Rio et al. 2004; NGC~5102; Deharveng et al. 1997), they are sufficiently rare that it is difficult to determine if evolutionary sequences exist.   In this paper we explore possible histories for a sample of actively star forming luminous dwarf galaxies with transitional properties between BCDs and dEs residing in loose group environments.  These systems are fairly isolated and show little indication of a recent interaction or merger. They have an intriguing combination of characteristics which do not allow them to easily be categorized as either dI/BCD or dE systems, but rather show a mixture of the two.  The key properties which indicate these low luminosity galaxies {\\it may} be in the midst of an evolutionary transition from dI/BCD to dE include: \\begin{itemize} \\item They are actively forming stars with star formation rates between 0.1 and 1 M$_{\\odot}$/yr. \\item The star formation is currently centrally concentrated, with the outer regions composed of an older stellar population. \\item Although actively forming stars, the starburst is fueled by very little HI gas, with M$_{HI}/L_B \\lesssim$\\ 0.1.  This compares with M$_{HI}/L_B >$ 0.2 for typical dIs (Roberts \\& Haynes 1994) and BCDs (Pustilnik et al. 2002) and M$_{HI}/L_B <$ 0.1 for dEs.  Star formation can continue in our sample galaxies only for about another 10$^9$ years at their current rate, based on their HI content; in a few cases, including molecular gas will at most double this time (e.g.  Jackson et al. 1989, Gordon 1991). \\item Unlike many other star forming dwarf galaxies including BCDs they have high oxygen abundance ratios, with 12 + log(O/H) $>$ 8.4 as we found in Dellenbusch et al. (2007). \\item Their outer optical colors are similar to those of typical BCDs with (B-R)$_0$ of order 1 (Gil de Paz et al. 2005), but bluer than the (B-R)$_0$ $\\approx$ 1.3 - 1.4 expected for slightly fainter dEs (Conselice et al. 2003).  In this regard our sample resembles the Gil de Paz et al. \"E-type\" BCDs. \\item  These objects have smooth outer isophotes which are much more indicative of early- rather than late-type galaxy structures. \\end{itemize} Table~\\ref{properties} quantifies several of these properties.  In this paper we examine the significance of these final two characteristics.  In \\S 2 we discuss the observations and data reduction.  In \\S 3 we describe our analysis and modeling processes used to examine the data for evidence of tidal debris and fine structures.  In this section we also describe the results of this analysis.  We discuss two possible evolutionary pictures to explain the galaxies' star formation histories as a means of transition in \\S 4.  \\S 5 contains a brief summary and our conclusions. ",
        "conclusions": "In summary, we have presented results of deep R-band imaging for five starbursting transition dwarf galaxies.  Our observational results show:  \\begin{enumerate} \\item All five galaxies exhibit smooth elliptical outer isophotes. \\item We find no evidence of extended tidal debris or other indications of a recent major interaction (e.g. tails, shells or ripples). \\item Fine structure exists in the central regions that are associated with the ISM and star formation, with dust structures and HII regions being prominent in several of the galaxies. \\item Outer envelope colors are consistent with having no star formation in the outer regions for the past several Gyr.  This timescale could be longer if star formation slowly declined rather than stopping suddenly. \\end{enumerate}  We consider two possible classes of evolutionary histories which could be consistent with the observed properties and locations of transition dwarf galaxies:  an internal ISM instability and a past interaction or merger event.  The interaction or dwarf-dwarf merger scenario remains a possibility, although we find no strong evidence to support it.  After several billion years it is not clear if we should still find evidence of tidal debris in the outer regions of the galaxies.  However, long term products of dwarf interactions and dwarf-dwarf mergers in particular have not been well modeled (e.g. the effects of large dark matter contents on merger products).  The other possibility of an unstable ISM seems promising to explain the central star formation in these objects.  In this case the eventual evolution to an inactive state will be slower, possibly requiring several starburst cycles.  Observationally the model can be tested by looking for the young stellar populations in fading starburst phases in dE-like dwarfs.  A remaining question, which is particularly puzzling, is the issue of timescales.  If these objects last experienced significant star formation in their outer regions $\\gtrsim$ 1 Gyr ago, why are they currently undergoing bursts of centrally concentrated star formation? What mechanisms control the structure of the ISM such that these galaxies now support rapid star formation, and will these circumstances lead to galaxies with little or no star formation, i.e. objects resembling dEs, within the next $\\sim$ 1 Gyr?"
    },
    "0710/0710.3161_arXiv.txt": {
        "abstract": "For a density that is not too sharply peaked towards the center, the local tidal field becomes compressive in all three directions. Available gas can then collapse and form a cluster of stars in the center, including or even being dominated by a central black hole. We show that for a wide range of (deprojected) S\\'ersic profiles in a spherical potential, the tidal forces are compressive within a region which encloses most of the corresponding light of observed nuclear clusters in both late-type and early-type galaxies.  In such models, tidal forces become disruptive nearly everywhere for relatively large S\\'ersic indices $n \\ga 3.5$.  We also show that the mass of a central massive object (CMO) required to remove all radial compressive tidal forces scales linearly with the mass of the host galaxy. If CMOs formed in (progenitor) galaxies with $n \\sim 1$, we predict a mass fraction of $\\sim 0.1-0.5$\\,\\%, consistent with observations of nuclear clusters and super-massive black holes. While we find that tidal compression possibly drives the formation of CMOs in galaxies, beyond the central regions and on larger scales in clusters disruptive tidal forces might contribute to prevent gas from cooling. ",
        "introduction": "\\label{sec:intro}  It is now well-known that the masses of supermassive black holes (SMBHs) in the centres of galaxies and bulges correlate with the stellar velocity dispersion, $\\Mbh \\propto \\sigma_\\star^{\\alpha}$ with $\\alpha \\sim 4-5$ \\citep[e.g.][]{Ferrarese+01, Gebhardt+01, Tremaine+02}, as well as nearly linearly with the mass of these spheroids, $\\Mbh \\propto \\Msph^{1.12\\,\\pm\\,0.06}$ \\citep[e.g.][]{MF01, HaringRix04}. \\cite{Ferrarese+06}, \\cite{WH06} and \\cite{Rossa+06} also found that the masses of nuclear (star) clusters (NCs), which are present in many both late and early-type galaxies \\citep[see, e.g.][]{Boeker+02, Cote+06}, are similarly related to the properties of the host galaxy \\citep[see also][]{GD07}. Recently, \\cite{McLaughlin+06} proposed momentum feedback, from accretion onto SMBHs or from stellar and supernovae winds in the case of NCs, as an explanation for the observed relations.  We investigate if central massive objects (CMOs), in the form of NCs, may have formed from gas being tidally compressed in the centers of galaxies. This effect, resembling compressive shocking of globular clusters by the Galactic disk \\citep{Ostriker+72}, has been studied by \\cite{Valluri93} in the context of tidal compression of a (disk) galaxy in the core of galaxy cluster. At the scale of galaxies, \\cite{DasJog99} argue for tidal compression of molecular clouds in the centers of flat-core early-type galaxies and ultraluminous galaxies as an explanation for the presence of observed dense gas. Very recently, an independent study by \\cite{Masi07} emphasised the potential importance of compressive tidal forces.  Low luminosity early-type galaxies and late-type galaxies share an overall luminosity profile with relatively low central power-law slopes. The fact that NCs are predominantly found in such galaxies \\citep[see e.g.][]{Cote+06} may provide an interesting link between the presence of CMOs in galaxies and the properties of the host galaxy. In the present study, we investigate whether tidal forces may help explaining this link. We first derive the radial component of the tidal force associated with a (deprojected) S\\'ersic profile in \\S~\\ref{sec:tidalcompression}. We then examine in \\S~\\ref{sec:scaling} how this applies to simple models of CMO hosts, including early-type and late-type galaxies. The corresponding results are then briefly discussed in \\S~\\ref{sec:disc}, and conclusions are drawn in \\S~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  We have built simple spherical models following deprojected S\\'ersic profiles to show that compressive tidal forces are naturally present in the central region when the S\\'ersic index $n \\la 3.5$. For $n \\ga 3.5$, the radial component of the tidal forces is disruptive almost everywhere (i.e., for $r / R_e > 10^{-3}$). Observed nuclear (star) clusters in early-type and late-type galaxies have extents and/or apparent locations which are within the tidally compressed region.  If we assume that most NCs form when their host galaxies have density profiles corresponding to rather low S\\'ersic indices $n \\sim 1$, we have shown that the masses of the NCs are consistent with $M_+$, the mass above which the radial tidal force becomes disruptive due the presence of the central massive object. In this picture, the predicted $M_+$ scales linearly with the host galaxy mass (or the mass of the spheroidal component) with $M_+/\\Msph \\sim 0.1-0.5$\\,\\% for $n \\sim 1$, in agreement with what is observed for both NCs and super-massive black holes in the centers of (more luminous) galaxies.  If indeed compressive tidal forces are a key actor in the formation of CMOs, only late-type galaxies have, today, the required gas content and density profiles ($n \\sim 1$), which allow the recurrent and common formation of CMOs (in the form of NCs). This is consistent with the fact that young NCs are predominantly found in late-type spiral galaxies. Finally, while we find that tidal compression possibly drives the formation of CMOs in galaxies, beyond the central regions and on larger scales in clusters disruptive tidal forces might contribute to prevent gas from cooling.  Such a simple scenario must be tested and extended to accomodate galaxies with e.g., core-Sersic surface brightness profiles \\citep[see e.g.][]{Ferrarese+06b} as well as to allow more realistic (non-spherical, multi-component) galaxy models. Moreover, using specific (stellar) mass-to-light ratios for the NCs and virial mass estimates for the host galaxy enables a direct comparison in mass instead of luminosity. Finally, hydrodynamical simulations are needed to examine the role of compressive tidal forces in the evolution of galaxies (and cluster)."
    },
    "0710/0710.3999_arXiv.txt": {
        "abstract": "{G39$-$27/289 is a  common proper motion pair formed  by a white dwarf (WD0433$+$270) and a main-sequence  star (BD$+$26~730) that apparently has been classified  as a member of the  Hyades open cluster. Previous studies of the white dwarf component yielded a cooling time of $\\sim4$ Gyr.  Although it  has not  been pointed  out before  explicitly, this result is  6 times larger than  the age of the  Hyades cluster, giving rise to an  apparent conflict between the physics  of white dwarfs and cluster main-sequence fitting.} {We investigate whether this system belongs to the Hyades cluster and, accordingly, give a  plausible explanation to the nature  of the white dwarf member.} {We have  performed and analyzed spectroscopic  observations to better characterize  these objects,  and used  their kinematic  properties to evaluate their membership to  the Hyades.  Then, different mass-radius relations and  cooling sequences for different  core compositions (He, C/O, O/Ne  and Fe) have  been employed to  infer the mass  and cooling time of the white dwarf.} {From  kinematic and  chemical composition  considerations  we believe that  the  system  was a  former  member  of  the Hyades  cluster  and therefore has an  evolutionary link with it. However,  the evidence is not  conclusive.  With  regards  to  the nature  of  the  white  dwarf component, we find  that two core compositions --- C/O  and Fe --- are compatible with  the observed effective temperature  and radius. These compositions  yield very different  cooling times  of $\\sim$4  Gyr and $\\sim$1 Gyr, respectively.} {We distinguish two  posssible scenarios. If the pair  does not belong to  the Hyades  cluster  but only  to  the Hyades  stream, this  would indicate that  such stream contains rather old  objects and definitely not coeval  with the cluster.   This has interesting  consequences for Galactic dynamics. However,  our favoured scenario is that  of a white dwarf with a  rather exotic Fe core, having  a cooling time compatible with the  Hyades age.   This is a  tantalizing result that  would have implications  for  the thermonuclear  explosion  of  white dwarfs  and explosion theories of degenerate nuclei.} ",
        "introduction": "The determination  of reliable ages  is of obvious importance  to both astrophysics and  cosmology, but not exempt of  many complications --- see,  e.g., \\cite{von01}.   The  age of  a  star is  perhaps the  most difficult property to estimate  and, further, it nearly always depends on rather strong model assumptions (Mamajek et al.~2007). An exception is the method that uses radioactive decay to directly estimate stellar ages (Cayrel  et al.~2001;  Frebel et al.~2007),  but this  has rather restricted    applicability.     Another    emerging   technique    is asteroseismology,  which has  the potential  to provide  accurate ages (Miglio  \\& Montalb\\'an  2005), but  this will  require high-precision photometric  data from  upcoming space  missions and  also  stars with well-constrained physical properties. In the meantime, the use of open clusters and main sequence stellar evolutionary models continues to be the most widely used method to  infer the ages of stars in the Galaxy. In  spite  of  the  still-present uncertainties,  such  as  convective overshoot, chemical composition anomalies,\\ldots~the field has reached sufficient maturity to  provide ages that are reliable  to better than $\\sim$10\\% --- see, for  instance, \\cite{pau06}.  The analogous method of  using  eclipsing  binaries  (Ribas  et  al.~2000)  yields  similar (model-dependent) accuracy.  The  study of  white dwarfs  has made  very valuable  contributions to numerous  areas of  astrophysics,  and estimating  individual ages  of stars  is no  exception. The  main advantage  of white  dwarfs  is the conceptual simplicity of their evolution,  which can be described as a simple  cooling process.   Modelling  of the  cooling sequences  makes white  dwarfs powerful  stellar chronometers,  accurate to  about 25\\% (Silvestri et al.~2005).  Among the major uncertainties of this method are the pre-white dwarf evolution time and the effects of the chemical composition of the core.  \\begin{table*}[t] \\begin{center} \\caption{Photometric data and stellar parameters derived for BD$+$26~730.} \\small{ \\begin{tabular}{lccccccc} \\hline \\hline \\noalign{\\smallskip} $V$ & $J$ & $H$ & $K$  & $T_{\\rm eff}$ (K) &  [Fe/H] & $Z$ & $\\log(L/L_{\\sun})$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 8.42$\\pm$0.02 & 5.945$\\pm$0.023 & 5.400$\\pm$0.018 & 5.240$\\pm$0.023 & $4,595\\pm30$ & $0.03\\pm0.09$   & $0.021\\pm0.004$  & $-0.527\\pm0.021$ \\\\ \\noalign{\\smallskip} \\hline \\hline \\end{tabular} \\label{tab:spar}} \\end{center} \\end{table*}  In this paper we discuss a case that presents an interesting puzzle in which the age estimate stemming  from the white dwarf cooling sequence is  in apparent  conflict with  the age  estimate coming  from cluster membership. The  object is G39$-$27/289,  a common proper  motion pair formed  by a  DA white  dwarf  (WD0433$+$270) and  a K  type star  --- BD$+$26~730 (Holberg et al.~2002) --- which is a well studied variable star (V833  Tau) and also  a single-lined spectroscopic binary  with a very  low-mass companion (Tokovinin  et al.~2006).   The members  of a common proper motion pair were likely born simultaneously and with the same chemical composition (Wegner 1973; Oswalt et al.~1988). Since the components are well separated ($\\sim  2,200$ AU in this case), no mass exchange  has taken  place and  they have  evolved as  isolated stars. Thus, it  is logical  to assume that  both components  of G39$-$27/289 have  the same  age  and  the same  original  metallicity. The  K-star component, and, by extension, its common proper motion companion, were classified as Hyades members  by \\cite{per98} and therefore would have an age  ranging from about  0.6 to 0.7  Gyr.  However, this  stands in obvious conflict with the cooling time of WD0433$+$270, which has been estimated to be about 4 Gyr (Bergeron et al.~2001).  In this  work, we carry out  a detailed study of  these objects, using both  information   present  in  the  literature  and   also  our  own spectroscopic  observations,  with the  objective  of unveiling  their nature and their  possible membership to the Hyades  open cluster.  We evaluate  the   different  scenarios  and   discuss  their  respective implications to white dwarf physics and Galactic dynamics. ",
        "conclusions": "In this work  we have presented new spectroscopic  observations of the members of  the common proper  motion pair G39$-$27/289.   Using these observations and the available photometry we have better characterized these objects, which has helped us to understand their nature.  Considering the  kinematic properties of  the members of the  pair and the lithium detection in BD$+$26~730,  we favour the scenario in which the common proper motion pair is  indeed a former member of the Hyades cluster,  and thus  its members  have a  coeval age  of $\\sim$0.6--0.7 Gyr.  Having evaluated  different compositions  for  WD0433$+$270, the young age inferred from cluster membership is only compatible with the case of an Fe  core, which  would have an  associated cooling  time of $\\sim$1 Gyr.  The  agreement is not perfect, but  the modelling of the cooling  evolution  of Fe-core  white  dwarfs  could  still have  some associated uncertainties.  The existence of white  dwarfs with an Fe-rich core has important consequences  for the  models  of thermonuclear  explosion of  stellar degenerate cores.  According to the  theory of stellar  evolution, all stars that develop an Fe core  experience a collapse to a neutron star or a  black hole, regardless of  the mass loss  rate assumed. However, there is  still an  alternative possibility to  avoid this  final fate that  relies  in the  failure  of  the  thermonuclear explosion  of  a degenerate  white  dwarf  near   the  Chandrasekhar  limit  (Isern  et al.~1991).  Our  current  view  of  a thermonuclear  explosion  is  as follows. Once the thermonuclear  runaway starts in the central regions of a white  dwarf, the ignition front propagates  outwards and injects energy  at a  rate that  depends on  the velocity  at which  matter is effectively burned leading  to the expansion of the  star. At the same time, electron  captures on the incinerated  matter efficiently remove energy at a rate determined  by the density causing the contraction of the  star.  Therefore,  depending  on  the ignition  density  and  the velocity of  the burning  front, the outcome  can be different.  It is known  that  He-degenerate cores  always  explode,  that  O/Ne and  Fe degenerate  cores always  collapse and  that C/O-degenerate  cores can explode or collapse depending  on the ignition density.  The existence of Fe-rich white dwarfs would imply the possibility of an intermediate behaviour between those discussed above in which an Fe remnant is left after a mild explosion.  A detailed theory explaining the  formation of Fe-core white dwarfs is still to be developed and up to now the possibility of their existence has been suggested mainly from observational evidences. Other examples of   possible  Fe-core   white   dwarfs  have   been   found  in   the past. \\cite{pro98} singled out two objects whose radii and masses were too small to  fit the typical C/O mass-radius  relation. These authors analyzed  a sample  of  white dwarfs,  calculated  their masses  using either  orbital parameters,  gravitational redshifts  or spectroscopic fits, and  determined their radii independently from  the knowledge of effective temperatures and  distances.  \\cite{pro98} indicated that in those two cases, a core made  by Fe was the only plausible explanation that could account for their small radii.  Besides the Fe-core hypothesis,  there is an alternative scenario that we  cannot rule out.   Although we  think it  is unlikely,  the common proper  motion pair  might not  be related  in any  way to  the Hyades cluster. This  would eliminate the  age constraint and thus  relax the requirement of  an exotic  composition for the  white dwarf  core.  In this case, the age of  the objects (both WD0433$+$270 and BD$+$26~730) could  be   of  $\\sim$4  Gyr.    This  would  have   some  interesting consequences to our  current view of star streams  or moving groups in th Galaxy.  As pointed out by \\cite{che99}, the Hyades stream contains at least  three distinct age groups with  ages up to 2  Gyr. The upper limit was probably a consequence of the use of A--F stars as kinematic tracers, which would  naturally have an age cutoff  at about 2--3 Gyr. With an age of 4  Gyr, WD0433$+$270 and BD$+$26~730 would indicate the existence  of an  even  older  population in  the  Hyades stream  thus completely  ruling out any  coevality within  the members.  This would lend strong support  to the model of a resonant  origin for the Hyades stream (Famaey et al.~2007). Another implication of such an old age is the  conflict with  lithium detection,  which would  imply  much lower destruction rate than expected.  Definitive proof supporting one of  these scenarios will have to await an  unambiguous   determination  of  the  mass  of   this  object  or, alternatively, a  more conclusive study on its  evolutionary link with the Hyades cluster."
    },
    "0710/0710.1678_arXiv.txt": {
        "abstract": "Numerous cosmological hydrodynamic studies have addressed the formation of galaxies.  Here we choose to study the first stages of galaxy formation, including non-equilibrium atomic primordial gas cooling, gravity and hydrodynamics.  Using initial conditions appropriate for the concordance cosmological model of structure formation, we perform two adaptive mesh refinement simulations of $\\sim$$10^8 \\Ms$ galaxies at high redshift.  The calculations resolve the Jeans length at all times with more than 16 cells and capture over 14 orders of magnitude in length scales.  In both cases, the dense, $10^5$ solar mass, one parsec central regions are found to contract rapidly and have turbulent Mach numbers up to 4. Despite the ever decreasing Jeans length of the isothermal gas, we only find one site of fragmentation during the collapse.  However, rotational secular bar instabilities transport angular momentum outwards in the central parsec as the gas continues to collapse and lead to multiple nested unstable fragments with decreasing masses down to sub-Jupiter mass scales.  Although these numerical experiments neglect star formation and feedback, they clearly highlight the physics of turbulence in gravitationally collapsing gas.  The angular momentum segregation seen in our calculations plays an important role in theories that form supermassive black holes from gaseous collapse. ",
        "introduction": "Since the first investigations of galaxy interactions \\citep{Holmberg41} using light bulbs, the use of numerical simulations in galaxy formation has developed dramatically.  Not only gravity but also hydrodynamics and cooling are standard ingredients in the sophisticated computer models studying galaxy formation and interactions.  In hierarchical structure formation, dark matter (DM) halos merge to form larger halos while the gas infalls into these potential wells \\citep{Peebles68, White78}.  \\citeauthor{White78} provided the basis for modern galaxy formation, in which small galaxies form early and continuously merge into larger systems.  As more high redshift galaxies were observed in the following 10 years, \\citet{White91} refined the theory to address the observed characteristics in these galaxies.  In their model, the halo accumulates mass until the gas cools faster than a Hubble time, \\tH, which usually occurs when atomic hydrogen line, specifically \\lya, cooling is efficient.  This happens when the halo has \\tvir~$>$~10$^4$ K, where the cooling function sharply rises by several orders of magnitude because the number of free electrons able to excite hydrogen greatly increases at this temperature \\citep{Spitzer78}.  One can define a cooling radius, \\rcool, in which the interior material is able to cool within a Hubble time.  Once the halo reaches this first milestone, \\rcool~ increases through additional accretion and cooling. A rapid baryonic collapse ensues when \\tcool~$\\lsim$~\\tdyn~ \\citep{Rees77}.  The material accelerates towards the center, and its density quickly increases.  In the model discussed in White \\& Frenk, this collapse will halt when one of the following circumstances occurs.  First, angular momentum can prevent the gas from collapsing further, and the system becomes rotationally supported.  Afterwards, this disc fragments and star formation follows.  Alternatively, star formation does not necessarily develop in a disc component, but the energy released by stars during their main sequence and associated supernovae (SNe) terminates the collapse.  These concepts have been applied also to the earliest galaxies in the universe \\citep{Mo98, Oh02, Begelman06, Lodato06}.  Many studies \\citep[e.g.][]{Ostriker96, Haiman97b, Cen03, Somerville03, Wise05} demonstrated that OB-stars within protogalaxies at $z > 6$ can produce the majority of photons required for reionization.  These protogalaxies contain an ample gas reservoir for widespread star formation, and the accompanying radiation propagates into and ionizes the surrounding neutral intergalactic medium.  Several high redshift starburst galaxies have been observed that host ubiquitous star formation at $z > 6$ \\citep{Stanway03, Mobasher05, Bouwens06}. Additionally, supermassive black holes (SMBH) more massive than 10$^8 \\Ms$ are present at these redshifts \\citep[e.g.][]{Becker01, Fan02, Fan06}.  Finally, a reionization signature in the polarization of the cosmic microwave background (CMB) at z $\\sim$ 10 \\citep{Page07} further supports and constrains stellar and SMBH activity at high redshifts.  The distinction between SMBH formation and a starburst galaxy should depend on the initial ingredients (i.e. seed BHs, metallicity, merger histories) of the host halo, but the evolution of various initial states is debatable.  It is essential to study the hydrodynamics of high redshift halo collapses because the initial luminous object(s) that emerges will dynamically and thermally alter its surroundings. For example, as the object emits ultraviolet radiation, the nearby gas heats and thus the characteristic Jeans mass increases, which may inhibit the accretion of new gas for future star formation \\citep{Efstathiou92, Thoul96}.  The following work will attempt to clarify early galaxy formation by focusing on protogalactic (\\tvir~$>10^4$ K) halos and following their initial gaseous collapse.  \\citet[][hereafter Paper I]{Wise07a} studied the virialization of protogalactic halos and the virial generation of supersonic turbulence.  In this paper, we address the gas dynamics of the continued, turbulent collapse of a halo and study the evolution and characteristics of the central regions.  In later studies, we will introduce the effects from primordial star formation and feedback and \\hh~cooling.  The progressive introduction of new processes is essential to understand the relevance of each mechanism.  We argue that our results may be relevant for scenarios that envisage SMBH formation from gaseous collapses.  \\citet{Loeb94} and \\citet{Bromm03} conducted smoothed particle hydrodynamics (SPH) simulations that focused on the collapse of idealized, isolated protogalactic halos.  The former group concluded that a central $10^6 \\Ms$ SMBH must exist to stabilize the thin gaseous disc that forms in their calculations.  \\citeauthor{Bromm03} considered cases with and without \\hh~chemistry and a background UV radiation field.  They observed the formation of a dense object with a mass $M \\sim 10^6 \\Ms$, or $\\gsim 10\\%$ of the baryonic matter, in simulations with no or strongly suppressed \\hh~formation.  These calculations without metal cooling and stellar feedback are useful to explore the hydrodynamics of the collapse under simplified conditions. \\citet{Spaans06} analytically studied the collapse of 10$^4$ K halos with an atomic equation of state.  They find that $\\sim$0.1\\% of the baryonic mass results in a pre-galactic BH with a mass $\\sim$$10^5 \\Ms$.  \\citet{Lodato06} also found that $\\sim$5\\% of the gas mass in $M = 10^7 \\Ms$ halos at $z \\sim 10$ becomes unstable in a gaseous disc and forms a SMBH.  Recently, \\citet{Clark07} studied the effects of metal and dust cooling on the fragmentation of a collapsing protogalactic core with varying metallicities ($Z = 0, 10^{-6}, 10^{-5} Z_\\odot$) and found the gas fragmenting ten times as much in the $10^{-5} Z_\\odot$ case than the primordial case.  In addition, the fragments in the primordial case are biased toward larger masses.  A runaway gaseous collapse requires angular momentum transport so material can inflow to small scales and form a central object.  The stability of rotating gaseous clouds have been subject of much interest over the last four decades and was thoroughly detailed by the work of \\citet[][hereafter EFE]{Chandra69}.  In the 1960's and 1970's, studies utilizing virial tensor techniques \\citep[EFE;][]{Lebovitz67, Ostriker69, Ostriker73a}, variational techniques \\citep{LyndenBell67, Bardeen77}, and N-body simulations \\citep{Ostriker73b} all focused on criteria in which a stellar or gaseous system becomes secularly or dynamically unstable.  The first instability encountered is an $m = 2$ bar-like instability that is conducive for angular momentum transport in order to form a dense, central object.  \\citet{Begelman06} investigated the conditions where a gaseous disc in a pre-galactic halo would become rotationally unstable to bar formation \\citep[see][]{Christodoulou95a, Christodoulou95b}.  They adapt the ``bars within bars'' scenario \\citep{Shlosman89, Shlosman90}, which was originally formulated to drive SMBH accretion from a gaseous bar that forms within a stellar galactic bar, to the scenario of pre-galactic BH formation.  Here a cascade of bars form and transport angular momentum outwards, and the system can collapse to small scales to form a quasistar with runaway neutrino cooling, resulting in a central SMBH.  The simulations detailed below show how many central bar-like instabilities form.   In \\S2, we describe our simulations and their cosmological context. In the following section, we present our analysis of the halo collapse simulations and investigate the structural and hydrodynamical evolution, the initial halo collapse, rotational instabilities, and the importance of turbulence.  In \\S4, we discuss the relevance of angular momentum transport and rotational instabilities in early galaxy and SMBH formation.  There we also examine the applicability and limitations of our results and desired improvements for future simulations.  Finally we conclude in the last section.   \\begin{figure*} \\resizebox{\\textwidth}{!}{\\rotatebox{0}{\\includegraphics*{f1a}}} \\resizebox{\\textwidth}{!}{\\rotatebox{0}{\\includegraphics*{f1b}}} \\caption{An overview of the final state of the collapsing protogalactic gas cloud.  Slices of log gas density in cm$^{-3}$ are shown through the densest point in the halo.  The \\textit{first} and \\textit{three} rows show simulation A, and the \\textit{second} and \\textit{fourth} rows show simulation B.  The columns in the top two rows from left to right are slices with a field of view of 10 kpc, 1 kpc, 100 pc, and 1 pc.  For the bottom two rows, the fields of view are 0.01pc, 20AU, 0.2AU, and 4 R$_\\odot$.  Note that each color scale is logarithmic, spans 5 orders of magnitude, and is unique for every length scale.} \\label{fig:slices} \\end{figure*} ",
        "conclusions": "We have simulated the hydrodynamics and collapse of a protogalactic gas cloud in two cosmology AMR realizations.  Our focus on the hydrodynamics presents a basis for future studies that consider stellar and BH feedback.  In the idealized case presented, we find a central dense object forms on the order of 10$^5 \\Ms$ and $r \\lsim 5$ pc.  This central object is not rotationally supported and does not fragment in our simulations.  However our results do not dismiss disc formation in protogalaxies because rotationally supported disc formation may begin after the initial central collapse.  Disc formation may be sensitively affected by feedback from the central object.  These simulations highlight the relevance of secular bar-like instabilities in galaxy formation and turbulent collapses.  Similar bar structures are witnessed in primordial star formation simulations. As low angular momentum infalls, it gains rotational energy as it conserves angular momentum.  This induces an $m$ = 2, bar-like instability that transports angular momentum outwards, and the self-similar collapse can proceed without becoming rotationally supported and exhibits a density profile $\\rho \\propto r^{-12/5}$. This process repeats itself as material infalls to small scales that is indicative of the ``bars within bars'' scenario.  We see three and four occurrences of embedded secular instabilities in the two realizations studied here.  We also find that supersonic turbulence influences the collapse by providing a channel for the gas to preferentially segregate according to its specific angular momentum.  The low angular momentum material sinks to the center and provides the material necessary for a central collapse.  Here the possibilities of a central object include a direct collapse into a SMBH \\citep[e.g.][]{Bromm03}, a starburst \\citep[e.g.][]{Clark07}, or a combination of both \\citep[e.g.][]{Silk98}.  All of these cases are viable in the early universe, and the occurrence of these cases depends on the merger history, local abundances in the halo, and the existence of a seed BH. Moreover, star formation should occur whether a central BH exists or not.  Perhaps the frequency of these different protogalactic outcomes may be traced with either 3D numerical simulations that consider star and SMBH formation and feedback along with metal transport or Monte Carlo merger trees that trace Pop III star formation, metallicities, and BHs.  We will attempt the former approach in future studies to investigate protogalactic formation in more realistic detail."
    },
    "0710/0710.1352_arXiv.txt": {
        "abstract": "Stellar feedback, expanding {\\sc H~ii} regions, wind-blown bubbles, and supernovae are thought to be important triggering mechanisms of star formation.  Stellar associations, being hosts of significant numbers of early-type stars, are the loci where these mechanisms act. In this part of our photometric study of the star-forming region NGC~346/N~66 in the Small Magellanic Cloud, we present evidence based on previous and recent detailed studies, that it hosts at least two different events of triggered star formation and we reveal the complexity of its recent star formation history.  In our earlier studies of this region (Papers I, III) we find that besides the central part of N~66, where the bright OB stellar content of the association NGC~346 is concentrated, an arc-like nebular feature, north of the association, hosts recent star formation. This feature is characterized by a high concentration of emission-line stars and Young Stellar Objects, as well as embedded sources seen as IR-emission peaks that coincide with young compact clusters of low-mass pre-main sequence stars. All these objects indicate that the northern arc of N~66 encompasses the most current star formation event in the region. We present evidence that this star formation is the product of a different mechanism than that in the general area of the association, and that it is triggered by a wind-driven expanding {\\sc H~ii} region (or bubble) blown by a massive supernova progenitor, and possibly other bright stars, a few Myr ago. We propose a scenario according to which this mechanism triggered star formation away from the bar of N~66, while in the bar of N~66 star formation is introduced by the photo-ionizing OB stars of the association itself. ",
        "introduction": "Massive OB stars, not having an optically visible pre-main sequence (PMS) contraction phase, appear almost immediately after their birth on the main sequence (e.g. Stahler \\& Palla 2005). They are mostly grouped in stellar associations, loose concentrations of stars, which host also significant numbers of intermediate- and low-mass PMS stars (see review by Brice\\~{n}o et al. 2007). When the far-UV radiation of the bright OB stars reaches the surface of the parental molecular cloud, a photo-dissociated region (PDR) develops. Nebula LHA~115-N~66 or in short N~66 (Henize 1956), the brightest {\\sc H~ii} region in the Small Magellanic Cloud (SMC), being very rich in early-type stars, is certainly an excellent example of an extragalactic PDR. The stellar association NGC~346, located in the central part of the nebula, hosts the largest sample of spectroscopically confirmed OB stars in the SMC (Massey et al. 1989), which have been the subject of several previous investigations (Niemela et al.  1986; Walborn et al. 2000; Evans et al. 2006).  These massive stars evolve rapidly, and immediately start to ionize the cloud, blowing its material away. The degree of ionization decreases outwards, and a thin barrier develops that segregates the ionized from the atomic gas. Molecular hydrogen is not fully ionized behind this front, but partly dissociated, while CO molecules which are located a bit deeper into the cloud are more easily dissociated than H$_{2}$ by absorbing UV photons. A correlation between H$_{2}$ infrared emission and CO lines, characteristic of a PDR, has been found for the region of NGC~346/N~66 by Contursi et al. (2000) and Rubio et al. (2000). These authors define the ``bar'' of N~66 as the oblique bright emission region, extending from southeast to northwest centered on NGC~346. They suggest that {\\sl star formation has taken place as a sequential process in the bar of N~66}, resulting in several embedded sources, seen as IR-emission peaks in the 2.14~\\micron\\ H$_{2}$ line and the ISOCAM LW2 band (5 - 8 \\micron). These peaks are alphabetically numbered from ``A'' to ``I'' (Contursi et al. 2000), with the association NGC~346 itself coinciding with peak ``C''.  Recent studies reconstruct the star formation history in nearby Galactic OB associations, and provide observational evidence for sequential and triggered star formation in their vicinity. Preibisch \\& Zinnecker (2007) conclude from their study of the Scorpius-Centaurus OB association that the formation of whole OB subgroups (each consisting of several thousand stars) requires large-scale triggering mechanisms such as shocks from expanding wind and supernova driven super-bubbles surrounding older subgroups\\footnote{According to these authors, other triggering mechanisms, like radiatively driven implosion of globules, also operate, but seem to be secondary processes, forming only small stellar groups rather than whole OB subgroups with thousands of stars.}. Since the low-mass stellar members of associations remain in their PMS phase for \\lsim~30~Myr, they play a key r\\^{o}le in the understanding of star formation in the vicinity of these systems (Brice\\~{n}o et al. 2007). Consequently, the recent discovery of a plethora of low-mass PMS stars in the region of NGC~346/N~66 with photometry from the {\\sl Advanced Camera for Surveys} (ACS) onboard HST (Nota et al. 2006; Gouliermis et al. 2006, hereafter Paper~I) can contribute significantly to the clarification of the mechanisms that may act in this extraordinary star-forming region.  The subsequent investigation of these PMS stars in NGC~346/N~66 (Sabbi et al. 2007, hereafter SSN07; Hennekemper et al. 2008, hereafter Paper~III) demonstrated that indeed they are clustered in several compact concentrations, some of them coinciding with the IR-emission peaks of Contursi et al. (2000; detected also with {\\sl Spitzer}, see \\S 2.1), verifying the existence of stellar subgroups in the region, similarly to galactic associations. Not being able to resolve differences in age smaller than 1 - 2 Myr in the individual CMDs of these sub-clusters, SSN07 suggest that ``all sub-clusters appear coeval with each other''. According to these authors, this coevality is a signature of the star formation conditions predicted by the hierarchical fragmentation of a turbulent molecular cloud model (Klessen \\& Burkert 2000; Bonnell \\& Bate 2002; Bonnell et al. 2003).  However, our analysis on the clustering properties of the PMS stars in NGC~346/N~66 (Paper~III) showed that {\\sl there are significant age differences between some sub-clusters and the association itself}. Specifically, three compact PMS clusters located to the north of the bar of N~66 are found to be not older than 2.5~Myr, while the CMD of the main body of the association NGC~346 show indications of an underlying older PMS population. This clearly suggests that these northern clusters are probably formed {\\sl after} the central stellar association. As we discuss later, triggering mechanisms for star formation such as those described by models of ionization shock fronts from OB stars (e.g. Kessel-Deynet \\& Burkert 2003) or wind-driven shock waves (e.g. Vanhala \\& Cameron 1998) may explain better their formation. Indeed, while a sequential star formation mechanism has been suggested to take place {\\sl in} the bar of N~66, around the association (Rubio et al. 2000), the existence of young compact PMS clusters {\\sl away} from it to the north of the association, does not quite fit in the hypothesis that this mechanism propagated from the center of N~66 along the bar. The triggering agent of these clusters may well be located outside of the bar.  In this paper we consider the findings from earlier and recent comprehensive investigations of the region NGC~346/N~66 to attempt a clearer understanding of the mechanisms that shape the recent star formation of this outstanding extragalactic star-forming region. Our aim is to provide answers to two important questions: i) {\\sl Is the star formation away from the association NGC~346 the product of triggered fragmentation of the cloud alone?} ii) {\\sl Is the photo-dissociation by the early-type stars of NGC~346 in the center of N~66 the only triggering mechanism in the region?} In section~2 we present evidence that this region has a far more complicated recent star formation history than what was previously considered, and that the most recent star formation event could not have been triggered by the central association. In section~3 we propose a scenario for the recent star formation in the northern part of the region NGC~346/N~66 and we indicate a nearby massive supernova progenitor, located at the northeast of the bar of N~66, as the triggering agent. We also discuss the possibility that other nearby massive stars away from the association may have also contributed to this mechanism. We present supporting evidence to this scenario using analyses of massive stellar content and gas kinematics of this region. Concluding remarks are given in section~4. ",
        "conclusions": ""
    },
    "0710/0710.2398_arXiv.txt": {
        "abstract": "{% Differential rotation can be detected in single line profiles of stars rotating more rapidly than about $v\\,\\sin{i} = 10$\\,km\\,s$^{-1}$ with the Fourier transform technique. This allows to search for differential rotation in large samples to look for correlations between differential rotation and other stellar parameters. I analyze the fraction of differentially rotating stars as a function of color, rotation, and activity in a large sample of F-type stars. Color and rotation exhibit a correlation with differential rotation in the sense that more stars are rotating differentially in the cooler, less rapidly rotating stars. Effects of rotation and color, however, cannot be disentangled in the underlying sample. No trend with activity is found.  } ",
        "introduction": "Stars are born from molecular clouds carrying net angular momentum that makes every star rotate more or less rapidly. During their evolution stars are being braked more or less efficiently, but even after several Gyrs of efficient braking stars as strongly braked as the Sun still show substantial angular velocity. This rotation in general must be expected to be differential, i.e. angular velocity changes with depth and latitude. In a rotating star, it is already the temperature difference due to rotationally induced surface gravity gradients which leads to meridional flow and differential rotation. More severe effects like magnetic forces and inhomogeneities in the convective structure may lead to even more severe effects of differential rotation.  The star with the best studied rotation law is the Sun, its equatorial angular velocity is roughly 20\\,\\% higher than the polar one, which has been observed centuries ago by following the rotation of spot groups on the solar surface. Today it is believed that these spots are due to magnetic activity of which at least the cyclic part is generated by a dynamo mechanism requiring the shear that is caused by differential rotation.  Detecting differential rotation on stars other than the Sun unfortunately is not as straightforward as on our host star. Other stars cannot (yet) be spatially resolved so that we cannot just follow the latitude-dependent motion of spot groups -- on some stars such spots might even be absent. Instead, one has to apply indirect techniques. One way to detect differential rotation is the reconstruction of the stellar surface using Doppler Imaging (DI) and comparing the spot's migration at different epochs. This method was applied successfully to many objects, and a summary of the state of the art is presented by Collier Cameron in this volume. DI correlates spot configurations between different epochs that can in principle be temporally separated as long as the lifetime of the spots, i.e. on the order of month in the Sun. Hence DI allows the detection of very small angular velocity differences. This method, however, requires that for each epoch the star is observed for a full rotation period in order to reconstruct a good picture of its surface. Depending on the rotation period and the brightness of the target that requires large amounts of telescope time, which in case of faint targets have to be quite big to ensure the small exposure times necessary for a good resolution of the surface.  A different approach to detect differential rotation is to scrutinize the shape of the rotationally broadened line profiles, which yields characteristic differences between rigid and differential rotation. This method was pioneered in the seventies (Gray, 1977). The problem of the degeneracy of differential rotation, limb darkening and inclination was investigated by several authors, (see Reiners \\& Schmitt 2002a and references therein). These authors also provide a recipe how to detect differential rotation in stellar line profiles.  Since then, differential rotation was searched for in more than a thousand spectra of stars from spectral types A--G. This resulted in many detections of differential rotation. The results of this project are published in Reiners \\& Schmitt (2003a, 2003b), Reiners \\& Royer (2004), and Reiners (2006).  In this paper, I summarize the results concentrating on the dependency of differential rotation on temperature, rotation, and activity. I focus on the F-stars, i.e. I will take into account data from different samples for stars of colors $0.2 < B-V < 0.6$. ",
        "conclusions": "I have presented the results from the analysis of several hundred high resolution spectra looking for differential rotation in F-stars. The fraction of stars with signatures of differential rotation in the overall sample of 200 stars is shown as a function of $B-V$ color, projected rotation velocity $v\\,\\sin{i}$, and normalized X-ray activity $\\log{L_\\mathrm{X}/L_\\mathrm{bol}}$. The fraction of differential rotators shows a clear trend in color and rotation velocity, two parameters which unfortunately are not uncorrelated in the sample. So far the conclusion is that among the cooler, less rapidly rotating objects the fraction of differentially rotating stars is larger than among the hotter, more rapidly rotating stars. No trend is seen in the fraction of differential rotators as a function of activity.  From the current sample, it cannot be decided whether slow rotation, low temperature, or both conditions together drive the high fraction of differentially rotating stars. To decide this, observations in a well defined sample containing more slowly rotating early F-stars as well as more rapidly rotating late F-stars are required."
    },
    "0710/0710.2351_arXiv.txt": {
        "abstract": "We use a statistical sample of $\\sim$500 rich clusters taken from 72 square degrees of the Red-Sequence Cluster Survey (RCS-1) to study the evolution of $\\sim$30,000 red-sequence galaxies in clusters over the redshift range 0.35$<$z$<$0.95.  We construct red-sequence luminosity functions (RSLFs) for a well-defined, homogeneously selected, richness limited sample. The RSLF at higher redshifts shows a deficit of faint red galaxies (to $M_V\\ge$ -19.7) with their numbers increasing towards the present epoch.  This is consistent with the `down-sizing` picture in which star-formation ended at earlier times for the most massive (luminous) galaxies and more recently for less massive (fainter) galaxies.    We observe a richness dependence to the down-sizing effect in the sense that, at a given redshift, the drop-off of faint red galaxies is greater for poorer (less massive) clusters, suggesting that star-formation ended earlier for galaxies in more massive clusters.  The decrease in faint red-sequence galaxies is accompanied by an increase in faint blue galaxies, implying that the process responsible for this evolution of faint galaxies is the termination of star-formation, possibly with little or no need for merging.  At the bright end, we also see an increase in the number of blue galaxies with increasing redshift, suggesting that termination of star-formation in higher mass galaxies may also be an important formation mechanism for higher mass ellipticals.  By comparing with a low-redshift Abell Cluster sample, we find that the down-sizing trend seen within RCS-1 has continued to the local universe. ",
        "introduction": "Clusters of galaxies are ideal laboratories for studying galaxy evolution since they contain many galaxies seen at the same epoch in close proximity.  Their cores are dominated by early-type galaxies, which are the major component of the high mass end of the galaxy stellar mass function locally.  There is now a good deal of evidence that cluster early-type galaxies formed the bulk of their stars at high redshift and thereafter simply evolved passively with little or no residual star-formation.  One such line of evidence is the tight sequence they form in color--magnitude space  \\citep[e.g.,][]{visv,bow92}, `the red-sequence'.  A similar red-sequence is also seen for early-type galaxies in the field out to at least z$\\sim$1 \\citep{Bell:2004lb}. Furthermore, {\\it all} galaxies appear to be divided into two distinct populations: the passively-evolving red-sequence and the actively star-forming `blue cloud'.  Only a small amount of residual star-formation (less than $\\sim$10\\% of the galaxies' past averaged star-formation rate) is necessary to move a galaxy from the red-sequence to the blue cloud.  Therefore, early-type galaxies can provide unique insight into the history of star-formation, as traced by objects in which star-formation has already been terminated.  In the local universe, the probability of a galaxy belonging to the red-sequence or blue cloud depends on its stellar mass and its environment \\citep{Baldry:2006bq}.  It is likely that the other fundamental parameter governing the properties of a galaxy is the epoch at which it is observed.  Thus, in order to build a complete picture of galaxy evolution, we need to study the colors of galaxies as a function of mass (or luminosity), environment and redshifts.  The classical picture for the formation of galaxies proposes a single `monolithic collapse` \\citep{Eggen:1962yu}, with stars in elliptical galaxies being formed in a single burst, thereafter evolving passively \\citep{Partridge:1967nh,Sandage:1970lj}.  This very simple model predicts remarkably well many of the properties and scaling relations of elliptical galaxies.  In the current hierarchical paradigm, structure forms in a `bottom-up' sense, as galaxies and clusters are built from the merging of smaller units.  Recently, there is growing evidence that star-formation has evolved in a `top-down` sense with more massive galaxies being most actively star-forming in the past and the bulk of the star-formation activity moving toward less massive galaxies  as the universe ages.  Although this seems intuitively at odds with hierarchical models, scenarios have been proposed in which star-formation progresses in this anti-hierarchical manner \\citep{De-Lucia:2006sa}.  Whereas previous generations of semi-analytic models in this hierarchical framework suggested that the most massive early-type galaxies should be younger than less massive ones \\citep{Baugh:1996vs,Kauffmann:1998mh}, in order to reconcile with the observed down-sizing trend, the prediction is now that although the most massive early-types assembled their {\\it mass} later than lower mass early-types, the stellar mass has been built up through a series of gas-poor mergers which do not result in additional star-formation. Hence the earlier formation times of the stellar populations in more massive galaxies is recovered.  Despite numerous signs of merging in early-type galaxies \\citep[e.g.,][]{van-Dokkum:2005dl,Tran:2005xc}, it remains an open question how important mergers are in their formation and evolution.  The problem of disentangling how a galaxy assembled its mass from how it assembled its stars is a difficult one.  Several studies of field galaxies have reported this down-sizing or anti-hieararchical trend in star-formation \\citep[e.g.,][]{Bell:2004lb,Juneau:2005ft,Faber:2005yw,Bundy:2005bc,Scarlata:2007oz}.  Initial results suggested that the comoving number density of massive early-type galaxies had evolved more than could be accounted for by passive evolution alone, and that `dry merging` of massive galaxies was required \\citep{Bell:2004lb,Faber:2005yw}.  More recently, it has been suggested that most, if not all, of the evolution can be attributed to the termination of star formation and pure passive evolution; and that a significant contribution from merging is not required \\citep{Cimatti:2006vz,Scarlata:2007oz}.  In galaxy clusters, down-sizing appears to be supported by the spectroscopic ages of red-sequence galaxies as a function of mass \\citep{Nelan:2005ml}.  In distant clusters (z$\\sim$0.8), a deficit of faint red-sequence galaxies relative to local clusters has been claimed, in accordance with this picture \\citep{De-Lucia:2004xa,Tanaka:2005mk,de-Lucia:2007li}.  However, all of these high-redshift works have found a large cluster-to-cluster scatter in their samples of sizes of approximately 1-10 clusters, indicating that a large, statistical sample is crucial to such studies.  The relative contributions of passive evolution versus dry merging to explain the evolution of the number density of early-type galaxies both in the field and in clusters is still an open question.  In this paper we present results using the first statistical sample of galaxy clusters drawn from a well-defined, wide area, homogenous survey covering a large redshift range, 0.35$<$z$<$0.95.  We present the survey data in \\S2 and detail our method for constructing composite clusters in \\S3.  In \\S4 we examine the Red-Sequence Luminosity Function and use the ratio of luminous-to-faint red-sequence cluster galaxies to trace its evolution with redshift and dependence on cluster mass. In \\S5 we discuss our results and compare with other studies of early-type galaxy evolution both in clusters and the field, and in \\S6 we present our conclusions.  Throughout we assume a cosmology of $H_0=$70 km s$^{-1}$Mpc$^{-1}$ (and $h=H_0/$100 km s$^{-1}$Mpc$^{-1}$), $\\Omega_{\\rm M}=$0.3 and $\\Omega_\\Lambda=$0.7. ",
        "conclusions": "We have studied the properties of red-sequence galaxies in a well-defined statistical sample of galaxy clusters over the redshift range 0.35$<$z$\\le$0.95.  Each redshift bin of width $\\Delta$z$=$0.1 contains $\\approx$50 clusters and $\\approx$5000 red-sequence galaxies.  Our main results are:   1) The faint end of the red-sequence, as measured by the ratio of luminous-to-faint galaxies, declines with increasing redshift.  This implies that star formation has not yet ended in the faintest cluster galaxies at high redshift.  The red-sequence is built up at the faint end as star-formation proceeds to progressively less luminous (less massive) galaxies, consistent with the down-sizing scenario \\citep{Cowie:1996xw}.  2) The turnover of the faint end of the RSLF is dependent on the cluster richness (mass) in the sense that for more massive clusters, the deficit of faint red-sequence galaxies is less than that for less massive clusters.  This is an indication that star formation ended earlier for faint galaxies in richer clusters than in poorer clusters.  3) The decline in the faint end of the red-sequence toward higher redshift is accompanied by an increase in the total (i.e., including blue galaxies) cluster LF.  This suggests that the build up of faint, red galaxies may be driven largely by the termination of star-formation in low mass galaxies.  A similar increase of blue galaxies is also seen at the brighter end of the LF, suggesting that (at least some of) the build up of high mass early-type galaxies may also be attributed to the termination of star-formation.  Future work will add the $B$- and $V$-band imaging of RCS-1 fields \\citep{Hsieh:2005fq} and the accompanying photometric redshift catalog to examine the luminosity functions of RCS clusters.  The $\\sim$1000 square degree next generation survey, RCS-2 \\citep{Yee:2007qb}, will provide an order of magnitude larger sample to improve upon the statistics of the current work."
    },
    "0710/0710.0953_arXiv.txt": {
        "abstract": "{}{To investigate variability and to model the pulsational behaviour of AGB variables in the intermediate-age LMC cluster NGC 1846.}{Our own photometric monitoring has been combined with data from the MACHO archive to detect 22 variables among the cluster's AGB stars and to derive pulsation periods. According to the global parameters of the cluster we construct pulsation models taking into account the effect of the C/O ratio on the atmospheric structure.  In particular, we have used opacities appropriate for both O-rich stars and carbon stars in the pulsation calculations.} {The observed P-L-diagram of NGC 1846 can be fitted using a mass of the AGB stars of about 1.8\\,M$_{\\sun}$. We show that the period of pulsation is increased when an AGB star turns into a carbon star. Using the mass on the AGB defined by the pulsational behaviour of our sample we derive a cluster age of $1.4\\times10^{9}$ years. This is the first time the age of a cluster has been derived from the variability of its AGB stars.  The carbon stars are shown to be a mixture of fundamental and first overtone radial pulsators.}{} ",
        "introduction": "The cluster NGC 1846 belongs to an intermediate age population in the Large Magellanic Cloud (LMC). In contrast to the Milky Way the LMC harbours many clusters with an age around 1 to 5 Gyr (e.g. Girardi et al.\\,\\cite{GCBB95}). These clusters thus provide an interesting possibility to study the late stages of stellar evolution for stars around 1.5 to 2.5 M$_{\\sun}$, while the globular clusters belonging to the Milky Way trace the evolution of stars of only up to about 0.9 M$_{\\sun}$.  While there is a general agreement in the literature that NGC 1846 is indeed an intermediate age cluster, the published values on global parameters of this system show some scatter. Derived metallicities can be roughly divided into two groups, one around [Fe/H]$=$-1.5 (Dottori et al.\\,\\cite{DPB83}, Leonardi \\& Rose \\cite{LR03}) and one around [Fe/H]$=$-0.7 (Bessell et al.\\,\\cite{BWL83}, Bica et al.\\,\\cite{BDP86}, Olszewski et al.\\,\\cite{OSSH91}, Beasley et al.\\,\\cite{BHS02}). Probably the first determination of the metallicity was done by Cohen (\\cite{C82}) giving a value in between the two discussed metallicity levels, namely [Fe/H]$=$-1.1. The most recent determination of the cluster's metallicity is given by  Grocholski et al.\\,(\\cite{GCS06}), who derive a value of [Fe/H]$=-$0.49$\\pm0$ based on the Ca triplet strength of 17 individual cluster members measured with the VLT.  Age values for NGC 1846 scatter between less than 1 Gyr (Frantsman \\cite{F88}) and 4.3 Gyr (Bica et al.\\,\\cite{BDP86}). Mackey \\& Broby Nielsen (\\cite{MB07}) give cluster ages from 1.5 to $2.5 \\times 10^{9}$ years based on a comparison of the colour magnitude diagram and theoretical isochrones. The range in age thereby results from the usage of two different sets of isochrones. These ages correspond to masses on the AGB of about 1.3 to 1.8 M$_{\\odot}$. Mackey \\& Broby Nielsen (\\cite{MB07}) also report on the probable existence of two populations in NGC 1846 with similar metallicity but separated in age by about 300 Myr.  A reddening of $A_{V}$$=$0.45\\,mag has been determined by Goudfrooij et al.\\,(\\cite{GGKM06}). Grocholski et al.\\,(\\cite{GCS06}) note that NGC 1846 is probably suffering from differential reddening without giving any detailed numbers. Keller \\& Wood (\\cite{KW06}) give a somewhat lower reddening of $E(B-V)$$=$0.08, i.e. an $A_{V}$ value of 0.25\\,mag.  The first studies of the stellar content of NGC 1846 were published by Hodge (\\cite{H60}) and Hesser et al.\\,(\\cite{HHU76}). A number of luminous stars on the Asymptotic Giant Branch (AGB) were identified and published with finding charts by Lloyd-Evans (\\cite{LE80}). Throughout this paper we will use the naming given in Lloyd-Evans' publication (LE{\\it XX}) except for H39, which follows the numbering from Hodge (\\cite{H60}). Details on individual AGB stars have been published in a number of papers. Frogel et al. (\\cite{FPC80}, \\cite{FMB90}) and Aaronson \\& Mould (\\cite{AM85}) gave near infrared photometry and $m_{\\rm bol}$ values for most of Lloyd-Evans' AGB stars and a few more cluster objects, unfortunately without any finding chart. Tanab\\'{e} et al.\\,(\\cite{T98}) searched this cluster for extreme infrared stars, i.e.~stars with a high circumstellar absorption in the visual range, but found none.  No investigations on the variability of the AGB stars in NGC 1846 has been published up to now. However, as most AGB stars are pulsating (forming the group of long period variables or LPVs) it seemed very likely that some light variability would be found in these stars. In a previous paper (Lebzelter \\& Wood \\cite{LW05}) we showed that the variability analysis of LPVs in a single stellar population like the globular cluster 47 Tuc allows one to investigate various fundamental aspects of AGB stars like the evolution of the pulsation mode or mass loss. With NGC 1846 we have chosen an interesting alternative target since its AGB stars are more massive and include a number of carbon stars. ",
        "conclusions": "\\label{discus} Having matched the O-rich models to the observed values of $T_{\\rm eff}$, we ensure that the models matched the observed pulsation periods of the O-rich stars.  If not, the mass was altered, and new calculations of the giant branch track and the pulsation periods were made.  It was found that a mass of 1.8\\,M$_{\\odot}$ gave the best fit between observations and theory for the O-rich stars.  This corresponds to a cluster age of $1.4\\times10^{9}$ years using the isochrones of Girardi et al. (\\cite{GBBC00})\\footnote{Using the BaSTI isochrones (Bedin et al. \\cite{basti05}) we find an age of $1.9\\times10^{9}$ years for a 1.8\\,M$_{\\odot}$ star to reach the AGB (using the same parameters as Mackey \\& Broby Nielsen \\cite{MB07}).}. The fit is shown in Figure~\\ref{pl_fig} and the model details are given in Table~\\ref{puls_mods}. It is clear that all the O-rich stars, with the exception of LE\\,17 near $\\log P = 1.8$ and $\\log L$/L$_{\\odot} = 3.49$, are first or second overtone pulsators (the stars with long secondary periods are not considered here).  It is clear from the models that as the C/O ratio increases beyond 1.0, the periods of all modes increase rather rapidly due to the increase in radius and decrease in $T_{\\rm eff}$ caused by the molecules in the C-rich atmospheres. Furthermore, it can be seen that the C stars fall quite nicely on two sequences corresponding to pulsation in the first overtone and fundamental modes.  In the absence of these C-rich models, it would have been natural to explain the C stars as fundamental mode pulsators scattering broadly around an extension of the fundamental mode O-rich sequence.  The fit of the models to the observed C star points is not perfect.  This is not surprising given the rather elementary method used to compute the C-rich opacities, combined with the fact that we do not really know the C/O ratios in these stars.  Furthermore, as mentioned above, the stars move up and down in luminosity during a He-shell flash cycle by factors of about 2 in luminosity (Vassiliadis \\& Wood \\cite{VW93}).  The dashed lines in Figure~\\ref{pl_fig} show how the fundamental and first overtone periods would vary over a shell flash cycle for a star with C/O = 2.  It seems most likely that the two C stars with $\\log L$/L$_{\\odot} < 3.6$ are near the luminosity minimum of a shell flash cycle (where AGB stars spend about 20\\% of their time).  In general, the theoretical and observed periods for the C-stars agree reasonably well.  This indicates that the radii and $T_{\\rm eff}$ values of the models agree well with the real values for these stars and the assumption that the temperatures indicated by the near infrared colours are misleading.  It is somewhat strange that the most luminous of the C stars appear to have shorter pulsation periods than the less luminous C stars.  One way for this to occur would be to decrease the C/O ratio in the most luminous C stars by hot-bottom burning.  However, at 1.8 M$_{\\odot}$, the mass is too low for hot-bottom burning, since AGB masses greater than about 4 M$_{\\odot}$ are required for hot-bottom burning to be significant.  Another explanation may be that the nonlinear pulsation period differs from the linear pulsation period when the pulsation amplitude increases, as noted in Lebzelter \\& Wood (\\cite{LW05}). The visual light amplitudes give no indication for this as the amplitudes of the carbon stars are all very similar (see Fig.\\,\\ref{tanabe}). In the infrared, as shown in Fig.\\,\\ref{lcir}, the star with the largest amplitude in our data set is LE5 which is also the most luminous star of our sample. This may indeed indicate that the shorter periods at the tip of the AGB can be explained by nonlinear effects. However, as mentioned above, our near infrared time series are too short for a definite conclusion on this point.  A reduction in stellar mass due to significant amounts of mass loss can not explain the shorter periods of the most luminous C stars since this would {\\em increase} the pulsation period.  In general, we have not found it necessary to invoke any significant reduction in mass with luminosity up the AGB in these calculations. We note that this contrasts with the stars in 47\\,Tuc studied by Lebzelter \\& Wood (\\cite{LW05}) where a significant amount of mass loss was found on the AGB.  As noted briefly in Section 4, this is as expected.  The stars being studied here and in 47\\,Tuc appear to have relatively low mass loss rates where a Reimers' mass loss law may apply.  In such a case, the mass loss rate is proportional to $LR/M$.  For a given $L$, this functional dependence suggests that the intermediate mass ($\\sim$1.8 M$_{\\odot}$) stars in NGC\\,1846 will have lower mass loss rates than the low mass ($\\sim$0.9 M$_{\\odot}$) 47\\,Tuc stars because of the higher mass.  In addition, the giant branch temperature increases with stellar mass so the radius will be smaller at a given $L$ in NGC\\,1846. The consequence of the expected lower mass loss rate in NGC\\,1846 is that a smaller amount of mass is lost on the giant branches up to a given $L$.  The fractional change in mass in NGC\\,1846 is even smaller because of the higher mass.   An independent check of the mass loss can be obtained from mid-IR photometry looking for indications of an infrared excess due to dust emission. NGC\\,1846 was covered by the SAGE mid-IR survey (Meixner et al.\\,\\cite{sage06}). Blum et al. (\\cite{Blum07}) find [8]-[24] as an indicator for mass loss. For five of the carbon stars in our sample (LE1, LE2, LE3, LE6 and LE11) the corresponding photometry can be found in the SAGE point source data base. All of them show an [8]-[24] colour typical for low mass loss stars. We recall that Tanab\\'{e} et al. (\\cite{T98}, \\cite{T04}) did not find any dust enshrouded objects in this cluster. For two sources Tanab\\'{e} et al. (\\cite{T04}) list no identification from the literature. Their object 16 is obviously the variable LW2. Object 8, the weakest object in their list of mid infrared sources, corresponds to LW4.  Using the SAGE list of point sources we checked for bright objects at 24\\,$\\mu$m with no or a very weak optical counterpart. Indeed we found three sources fulfilling this criterion, but their membership to the cluster is not very likely. In the appendix we give a more detailed description of these three sources. Thus we can derive from our findings that the AGB stars of NGC 1846 lose their mass in a quite short time at the end of the AGB phase.  The location of LE17 in the P-L-diagram remains somehow a mystery. While the two carbon stars found at $\\log L$/L$_{\\odot}$ = 3.57 nicely agree with the expected location of a TP-AGB star close to its luminosity minimum, it is unlikely that this explanation also holds for LE17, which is the O-rich star slightly below in the P-L-diagram (Fig.\\,\\ref{pl_fig}). If it would be in the minimum of its TP cycle, it would be in a region exclusively occupied by C-rich stars during its maximum. A possibility may be that the star is located at the \"knee\" of the fundamental mode sequence at its maximum. Alternatively, the derived temperature and luminosity of this star may be wrong due to higher reddening of this object. Circumstellar reddening probably plays only a minor effect as there is no clear indication for this from mid-IR photometry (SAGE; Tanab\\'{e} et al.\\,\\cite{T04}). We further note that the star nicely fits onto the giant branch of NGC 1846 (Fig.\\,\\ref{hrd}), thus there is no indication for a significant reddening. For the same reason we may safely assume that the star is indeed a member of the cluster. Further investigations are required to understand the behaviour of this star."
    },
    "0710/0710.2484_arXiv.txt": {
        "abstract": "\\medskip \\noindent We assume the validity of the Standard Model  up to an arbitrary high-energy scale and discuss  what information on the early stages of the Universe can be extracted from a measurement of the Higgs mass. For $M_h \\simlt 130 \\gev$, the Higgs potential can develop an instability at large field values. From the absence of excessive thermal Higgs field fluctuations we derive a bound on the reheat temperature after inflation as a function of the Higgs and top masses. Then we discuss the interplay between the quantum Higgs fluctuations generated during the primordial stage of inflation and the cosmological perturbations, in the context of landscape scenarios in which the inflationary parameters scan. We show that, within the large-field models of inflation, it is highly improbable to obtain the observed cosmological perturbations in a Universe with a light Higgs. Moreover, independently of the inflationary model, the detection of primordial tensor perturbations through the $B$-mode of CMB polarization and the discovery of a light Higgs can simultaneously occur only with exponentially small probability, unless there is new physics beyond the Standard Model. ",
        "introduction": "The search for the Higgs boson and the measurement of its properties are one of the primary goals of the LHC. The mere discovery of a Higgs can be viewed as a possible indication of additional new physics not far from the electroweak scale, because of the high sensitivity to short-distance quantum corrections of the mass term associated to a fundamental scalar. However, even in the absence of any new-physics discovery at the LHC, a measurement of the Higgs mass can give us useful hints on the structure of the theory at very high energies. This is because, at large field values, the Higgs potential can develop an instability or become non-perturbative, depending on the precise value of the Higgs quartic coupling $\\lambda$ or, ultimately, on the Higgs mass $M_h$. Because of the logarithmic dependence of $\\lambda$ on the energy, such considerations~\\cite{con,Str,CEQ,thermal1,Hambye} can test the properties of the theory up to extremely small distances, which are otherwise totally unaccessible to any imaginable collider experiment.     In this paper, we want to use the same considerations for a different purpose. Rather than trying to infer new particle-physics properties at small distances, we will assume the validity of the Standard Model (SM) up to an arbitrary high-energy scale, and find what information on the early stages of the Universe can be extracted from a measurement of the Higgs mass. We will first obtain that in the Higgs mass range $114 \\gev \\simlt M_h\\simlt 130 \\gev$, where the electroweak vacuum is potentially metastable, the absence of excessive thermal Higgs field fluctuations in the early Universe imposes a bound on the reheat temperature after inflation $T_{RH}$.  Then we will discuss the interplay between the quantum Higgs fluctuations generated during the primordial stage of inflation and the cosmological perturbations which are either currently observed   in the form of  Cosmic Microwave Background (CMB) anisotropies,  or might be detected in the near future in the form of tensor (gravity waves) perturbations. The key ingredient is that all these fluctuations depend upon the Hubble rate during inflation and therefore, under certain assumptions, it is possible to relate the amount of Higgs fluctuations to observable properties of the CMB. However, excessive fluctuations of the Higgs field during inflation can pose a threat to the stability of the present vacuum, if the Higgs mass lies in the metastability window.  We will assume that the various initial inflationary patches of the Universe are characterized by different microphysical parameters. In this sense we take the point of view that the underlying theory has many vacua, realized in different patches of the Universe. This picture, usually referred to as the landscape \\cite{landscape}, has been put forward especially in the context of string theory. Under this assumption, from CMB observations and from a measurement of the Higgs mass, we can derive probabilistic conclusions on the properties of our Universe.  In particular, within the  class of large-field models of inflation, we will compute the probability to have a Universe at the end of inflation which both survived the quantum Higgs fluctuations and has the right amount of observed cosmological perturbations. Such probability is extremely (exponentially) small, if the Higgs mass is below 124~GeV (for the present central value of the top mass). Moreover, we find that the discovery of a light Higgs together with the detection of primordial tensor perturbations through the $B$-mode of CMB polarization (at a level quantitatively described in sect.~6)  would imply  that we live in a very atypical Universe, whose  probability decreases exponentially when the Higgs mass decreases. Such discovery could be interpreted as an indirect evidence for the existence of new physics beyond the Standard Model at some intermediate energy scale.  The paper is organized as follows. In sect.~2 we briefly review the Higgs mass instability window. In sect.~3 the bounds on the reheating temperature after inflation are discussed. In sect.~4 we compute the survival probability of the electroweak vacuum during inflation and in sects.~5 and~6 we relate it to the curvature and tensor perturbations, respectively. Section~7 states our conclusions. The appendix contains technical details relevant to the calculation of the Higgs mass instability window. ",
        "conclusions": "In this paper we have investigated some possible cosmological implications of the Higgs mass measurement. If the LHC discovers a Higgs with mass in the metastability window shown in fig.~\\ref{fig:window}, and does not find direct evidence for other new physics, then there is a concrete possibility that we live in a metastable state. If we assume that the pure SM is valid up to very large energy scales, then stability against field fluctuations tests properties of the early Universe. We have revisited the known considerations about thermal fluctuations, interpreting the result as an upper limit on the reheating temperature $T_{RH}$ after inflation. The bound is summarized in fig. \\ref{mhTRH}.   We have also discussed the possibility  that the inflationary vacuum fluctuations  destabilize the electroweak vacuum. If the Hubble rate is large enough during inflation and the Higgs mass is light, then the danger exists that the  classical value of the Higgs field is pushed above its instability point causing the collapse of the corresponding inflating domain. This does not necessarily pose a problem, since all the surviving domains will be characterized by the correct electroweak vacuum. However, interesting probabilistic conclusions can be reached under the assumption of a ``landscape\" scenario in which the inflationary parameters take different values in different patches of the Universe. In the context of large-field models of inflation we have argued that,  among the surviving regions, only an exponentially  tiny fraction of them  will be characterized by the correct amount of cosmological perturbations. If the Higgs mass is found below $(120-130)$~GeV, either we live in a very special, and exponentially unlikely, domain or new physics must exist below the scale $\\Lambda$.  Moreover, our considerations can be directly related to the amount of primordial gravity waves which can be measured through their imprint on the CMB. The discoveries of a light Higgs boson and of tensor modes would be mutually conflicting (at least in a probabilistic sense) if the condition in \\eq{trul} is not satisfied. Again, this evidence could be interpreted as an indication of new physics modifying the extrapolation of the Higgs potential to large field values. This result is valid, independently of the particular inflationary model considered.  Finally, a measurement of the Higgs mass below about 130~GeV provides a direct bound, shown in fig.~\\ref{conformal}, on the quantity $\\xi H^2$, where $\\xi$ is the (negative) coupling between the Higgs bilinear and the curvature, and $H$ is the maximal Hubble rate during inflation. This bound becomes particularly interesting in case of detection of primordial gravity waves, which provide a measurement of $H$. Note that the bound in fig.~\\ref{conformal} does not depend on any statistical consideration of parameter scanning."
    },
    "0710/0710.5437_arXiv.txt": {
        "abstract": "We have performed a series of three-dimensional simulations of a starburst-driven wind in an inhomogeneous interstellar medium. The introduction of an inhomogeneous disk leads to differences in the formation of a wind, most noticeably the absence of the ``blow-out'' effect seen in homogeneous models. A wind forms from a series of small bubbles that propagate into the tenuous gas between dense clouds in the disk. These bubbles merge and follow the path of least resistance out of the disk, before flowing freely into the halo. Filaments are formed from disk gas that is broken up and accelerated into the outflow. These filaments are distributed throughout a biconical structure within a more spherically distributed hot wind. The distribution of the inhomogeneous interstellar medium in the disk is important in determining the morphology of this wind, as well as the distribution of the filaments. While higher resolution simulations are required in order to ascertain the importance of mixing processes, we find that soft X-ray emission arises from gas that has been mass-loaded from clouds in the disk, as well as from bow shocks upstream of clouds, driven into the flow by the ram pressure of the wind, and the interaction between these shocks. ",
        "introduction": "In 1963, \\citeauthor{LS1963} first detected an outflow of gas along the minor axis of M82. \\citet{CC1985} proposed a model in which a galactic scale outflow could be powered by the combined kinetic energy from supernovae. Starburst galaxies, with their characteristically high star formation rates, provide the perfect environments for these winds to develop. Indeed, galactic winds are ubiquitous in starburst galaxies, having been observed in many nearby galaxies and inferred in galaxies at  high-redshifts \\citep[see][~and references therein]{VCB2005}  The best studied galactic wind is the outflow in M82, which is clearly visible in the light of H$\\alpha$, displaying a vast filamentary system extending several kpc along the minor axis of the galaxy \\citep{SB1998}. These filaments lie on the surface of a mostly hollow structure and rotate in the same direction as the disk \\citep{Greve2004}. As with other galactic winds \\citep[e.g. NGC 253:][]{SDW2003}, the wind in M82 is asymmetric, with the northern outflow more chaotic than the southern outflow. The filaments can be traced to the nuclear region and display both shell and loop-like structures \\citep{Oetal2002}. The formation of these filaments is currently not well understood, but they are thought to be either disk or halo gas that has been entrained into the outflow.  The morphology of galactic winds can vary. Outflows often display asymmetries, varying degrees of collimation and may be tilted with respect to the minor axis. While many outflows are limb-brightened \\citep[e.g. NGC 3079;][]{Vetal1994}, the optical filaments can also fill the volume rather than  remain confined to the surface of the biconical outflow \\citep{VB1997}. The host galaxy itself plays an important role in determining the morphology of a wind, with its size and structure affecting the degree of collimation \\citep{SS2000} and expansion of the outflow \\citep{SetalB2004,Getal2005,Martin2005}.  Recent Chandra observations have revealed increasing detail in the X-ray emission from galactic winds. One of the most striking results of these observations is the close spatial relationship with the H$\\alpha$ emitting gas \\citep[e.g.][]{Setal2000,Setal2002,SetalA2004,CBV2002,MKH2002,Getal2005,OWB2005b}, suggesting a close physical connection. Thus, a successful model of a galactic wind needs to explain this relationship. \\citet{Setal2002} provide a summary of several theories for the origin of the X-ray emission that could explain this correlation. These mechanisms involve shocked disk or halo gas that has been swept up into the wind, in the form of dense clouds or shells.  Over the past few decades, numerous simulations have been made of starburst-driven winds \\citep{TI1988,TB1993,Setal1994,Setal1996,DB1999,TM1998,SS2000,TSM2003}. \\citet{Setal1994} performed two-dimensional, axisymmetric simulations of a galactic wind in an isothermal ISM, with varying densities and temperatures. They concluded that the H$\\alpha$ filaments form from disk gas that has been entrained into the flow and that the X-ray emission most likely arises from shocked disk and halo gas. More recently, \\citet{SS2000} performed a series of simulations, focusing on the energetics and X-ray emission from the wind. As with \\citet{Setal1994}, their simulations were two-dimensional and axisymmetric with an isothermal ISM. They found that a large fraction of the soft X-ray emission in their model comes from shock-heated ambient gas and from the interfaces between cool dense and hot tenuous gas. While these simulations provide some insight into the origin of the X-ray and H$\\alpha$ emission, the homogeneous nature of these  models and their symmetry renders them incapable of forming significant filamentary structures, limiting their ability to constrain the emission processes.  In order to improve upon previous models and to gain a better understanding of the origin of the H$\\alpha$ filaments and X-ray emission, we have performed a series of three-dimensional simulations of a galactic wind in an inhomogeneously distributed interstellar medium (ISM). The introduction of inhomogeneity is important as the interstellar medium in a galaxy disk is highly complex in all its phases \\citep[see for example][~and references therein]{EE2001}.  Inhomogeneity is also crucial in the development of a wind, as energy from massive stars formed in dense molecular clouds in the starburst region may be radiated away before a wind could form. A wind is more likely to develop from the kinetic energy from stellar winds adjacent to the diffuse gas surrounding the clouds.  The inhomogeneous structure of the ISM is also likely to affect the distribution of filaments throughout the wind, producing asymmetric and tilted outflows. It is likely that the size and strength of the starburst itself plays an important role in determining the morphology. Many starburst galaxies, such as M82 and NGC 3079, contain circumnuclear starbursts, with their resultant outflows being strong and violent \\citep{SB1998,Vetal1994}. Other starbursts are weaker and have less prominent outflows. An example is NGC 4631, which is currently undergoing a disk-wide starburst \\citep{SetalA2004}.  \\citet{TSM2003} investigated the formation of the emission line filaments by modeling the formation of a wind from several super star clusters. They proposed that kiloparsec long filaments are formed from stationary and oblique shocks. In this paper we present a different model, which follows a similar approach to that of \\citet{SS2000}, but introduces an inhomogeneous disk. We follow the evolution of a starburst-driven wind in different ISM conditions and discuss the effect of the inhomogeneity of the disk on the morphology of the wind. We consider the morphology of the H$\\alpha$ emitting filaments separately and investigate their origin. Finally, the luminosity of the soft and hard X-ray emitting gas is calculated and we suggest an origin for the soft X-ray emission. ",
        "conclusions": "We have performed a series of three-dimensional simulations of a starburst-driven galactic wind designed to test the evolution of the wind in different ISM conditions. By conducting three-dimensional simulations we are able, for the  first time, to study the morphology and dynamics of the entire outflow. The introduction of an inhomogeneous disk enables us to study the development of asymmetries and the interaction of the wind with clouds in the disk. The results of these simulations are as follows-  \\begin{enumerate}  \\item The interstellar medium plays a pivotal role in the evolution of a galactic wind. The interaction of the wind with clouds in the disk results in asymmetries and tilted outflows on the small scale. Nevertheless, it is likely that inhomogeneities in the halo are the cause of the large-scale asymmetries in an outflow.  \\item The distribution of gas surrounding the starburst region assists in collimating the outflow. The thickness of the disk and the location of the starburst are important factors in determining the degree of collimation, with the degree of collimation increasing with the amount of gas surrounding the starburst region.  \\item The base of the outflow is well confined within a radius of 200 pc over the 2 Myr time frame of the simulation as a result of the high density of the disk gas.  \\item The H$\\alpha$ filaments form from the breakup of clouds in the starburst region, the fragments of which are then accelerated by the ram pressure of the wind. Filaments are also formed from gas that has been stripped from the sides of the starburst region. The distribution and mass of the filaments is affected by the distribution of clouds in the vicinity of the starburst region.  \\item The H$\\alpha$ filaments appear as strings of disk gas that form a filled biconical structure inside of a more spherical hot wind. The filaments are distributed throughout this structure, but do not trace the true extent to the wind defined by the hot gas.  \\item The calculated soft X-ray luminosities up until 2.0 Myr are of the order of 10$^{38}$ - 10$^{39}$ erg s$^{-1}$ and the hard X-ray luminosities of 10$^{38}$ erg s$^{-1}$. These luminosities are dependent on the volume of the wind and would be larger for a more evolved outflow.  \\item Interior to the swept-up shell of halo gas, soft X-ray emission originates in the same region as the H$\\alpha$ emitting gas. While higher resolution simulations are needed to confirm X-ray emission from mixing processes, we find 4 mechanisms that give rise to Soft X-rays: (i) The mass-loaded wind, (ii) the intermediate temperature interface between the hot wind and cool filaments, (iii) bow shocks, and (iv) interactions between bow shocks. The shell is also a major contributor to the soft X-ray emission, but has no associated H$\\alpha$ emission.  \\item The hard X-ray emission originates from gas in the starburst region.  \\end{enumerate}  The results of these simulations indicate that the host galaxy itself and the environment in which it is situated is a major determinant in the morphology of the outflow. The emission processes that contribute to the H$\\alpha$ and soft X-ray emission may vary from one galaxy to the next. Whether the H$\\alpha$ emission originates from photoionization or from shock-heating (or both) cannot be determined from these simulations. However, we do find an abundance of filamentary T $\\sim$ 10$^4$ K gas that has been accelerated into the outflow, forming a biconical shaped region that is commonly observed in starburst winds. The source of the soft X-ray emission is also likely to depend upon the environment of the host galaxy. In the case of M82 it is plausible that the interaction of the wind with the surrounding HI clouds is also a contributor to the soft X-ray emission, in addition to the processes mentioned above.  The observed spatial relationship between the H$\\alpha$ and soft X-ray emitting gas can be explained when considering emission processes interior to the wind, such as bow-shocks and the mass-loaded component of the wind. In addition, the presence of the strong X-ray emitting shell with no associated H$\\alpha$ emission is interesting. While the ultimate fate of the shell is unknown at present, this result argues for the presence of X-ray emission more extended than the filamentary H$\\alpha$ gas. \\citet{Setal2002} suggest that this emission may be detectable in more distant starburst galaxies.  In future work we shall investigate the evolution of a wind over a larger time frame and spatial extent than our current study, and look at the total energy budget of the outflow. It is also important to further test the effect of resolution on the filaments, and in particular the associated soft X-ray emission that arises through mixing of the hot wind and the cooler disk gas, and will be the subject of a subsequent paper."
    },
    "0710/0710.1434_arXiv.txt": {
        "abstract": "We discuss the possibility of accurately estimating the source number density of ultra-high-energy cosmic rays (UHECRs) using small-scale anisotropy in their arrival distribution. The arrival distribution has information on their source and source distribution. We calculate the propagation of UHE protons in a structured extragalactic magnetic field (EGMF) and simulate their arrival distribution at the Earth using our previously developed method. The source number density that can best reproduce observational results by Akeno Giant Air Shower Array is estimated at about $10^{-5}~{\\rm Mpc}^{-3}$ in a simple source model. Despite having large uncertainties of about one order of magnitude, due to small number of observed events in current status, we find that more detection of UHECRs in the Auger era can sufficiently decrease this so that the source number density can be more robustly estimated. 200 event observation above $4 \\times 10^{19}~{\\rm eV}$ in a hemisphere can discriminate between $10^{-5}$ and $10^{-6}~{\\rm Mpc}^{-3}$. Number of events to discriminate between $10^{-4}$ and $10^{-5}~{\\rm Mpc}^{-3}$ is dependent on EGMF strength. We also discuss the same in another source model in this paper. ",
        "introduction": "\\label{introduction}  The nature of the sources of ultra-high-energy cosmic rays (UHECRs) is still poorly known though many efforts in detection of these particles with UHE energies. It is one of challenging problems in astroparticle physics.  About 10 years ago, Akeno Giant Air Shower Array (AGASA) reported small-scale anisotropy of the UHECR arrival distribution within its angular resolution while it did large-scale isotropy with a harmonic analysis \\cite{takeda99,takeda01}. It found five doublets and one triplet as event clusterings. It is enough evidence that origin of UHECRs is astrophysically point-like source. However, no obvious astronomical counterparts to the observed UHECRs have been found. One of this reasons is magnetic fields in the universe.  Extragalactic magnetic field (EGMF) is poorly known theoretically and observationally. As an upper limit for its strength, $B$, and correlation length, $l_c$, $B~{l_c}^{1/2} < (1~{\\rm nG})(1~{\\rm Mpc})^{1/2}$, by Faraday rotation measurements of radio signals from distant quasars, is often adopted \\cite{kronberg94}. Based on its upper limit, UHECR propagation has been discussed in simply uniform turbulent magnetic field \\cite{yoshiguchi03,aloisio04,berezinsky06}.  It is also observationally known that clusters of galaxies have strong magnetic fields with 0.1 - a few $\\mu$G at its center \\cite{vallee04}. This is shown that there are not only uniform EGMF but also relatively strong EGMF incidental to the local structure. Recently, several groups have performed simulations of large-scale structure formation with magnetic fields \\cite{sigl03,sigl04,dolag05}. They find that magnetic field traces the local density distribution. Such magnetic field plays an important role in UHECR propagation since it is stronger than the uniform one.   In our previous study \\cite{takami06}, we constructed a structured EGMF model that reflects the local structures and discussed predicted arrival distribution of UHE protons at the Earth in consideration with their propagation processes. We also considered Galactic magnetic field (GMF) in the propagation. As a result, the number density of UHECR sources was constrained to $\\sim 10^{-4}~{\\rm Mpc}^{-3}$ under our luminosity-weighted source model (explained in section \\ref{model}) from observed arrival distribution. The EGMF strength was normalized only to $0.4~\\mu{\\rm G}$ at the center of the Virgo cluster. However, observational measurements of magnetic field in a galaxy cluster result in 0.1 - a few $\\mu$G, so that the strength of magnetic field has large uncertainty of about one order of magnitude. Thus, it is important that the propagation process and the arrival distribution are investigated in several strengths of EGMF.  In this study, we calculate propagation of UHE protons in a structured EGMF with several strengths for normalization and simulate the arrival distributions at the Earth. At first, comparing them to observational results, we constrain the source number density. We use AGASA data for this purpose because AGASA has observed the most number of events. Such constraint is dependent on UHECR source model. We discuss about two simple source models. We find that the number density has large uncertainty of about an order of magnitude due to small number of observed events at present. So, we also discuss the possibility of a decrease in the uncertainty with future observations.  In section \\ref{model}, our models of UHECR source distribution, structured EGMF, and GMF are briefly explained. In section \\ref{method}, a numerical method of UHE proton propagation, construction of the arrival distribution and statistical analysis are explained. We present our results in section \\ref{result}, and summarize in section \\ref{summary}. ",
        "conclusions": "\\label{summary}  In this study, we discuss the possibility of accurately estimating the source number density of UHECRs with small-scale anisotropy. Comparison between simulated arrival distribution and the observational results enables us to estimate the source number density. In order to construct the arrival distribution, we calculate the propagation of UHE protons in a structured EGMF with several strengths consistent with measurements of magnetic field in clusters of galaxies. The GMF is also considered. We find that the source number density of $10^{-5}~{\\rm Mpc}^{-3}$ in the normal source model and $10^{-4}~{\\rm Mpc}^{-3}$ in the luminosity-weighted source model can best reproduce the AGASA results, which are weakly dependent on strength of our structured EGMF. However, these have large uncertainty of about one order of magnitude due to the small number of observed events.  So, we discuss the possibility that future observations decrease the uncertainty. In the normal source model, Auger can distinguish $10^{-5}~{\\rm Mpc}^{-3}$ and $10^{-6}~{\\rm Mpc}^{-3}$ sufficiently by our method in the near future. If the structured EGMF is zero or very weak, $10^{-4}~{\\rm Mpc}^{-3}$ is also discriminated from the less number density in 500 event observation above $4 \\times 10^{19}~{\\rm eV}$. In stronger EGMF, more observations are requested because cosmic rays are diffused more strongly. Number of events that needed for the distinction depends on EGMF strength. In the luminosity-weighted source model, $10^{-3}~{\\rm Mpc}^{-3}$ can be distinguished from the less number density by Auger. The distinction between the less number densities is difficult due to large uncertainty which originates from different injection powers of the sources.  \\begin{figure}[t] \\begin{center} \\includegraphics[width=0.48\\linewidth]{fig8a.eps} \\hfill \\includegraphics[width=0.48\\linewidth]{fig8b.eps} \\caption{Distribution of $\\chi^2$s calculated from arrival protons above $4 \\times 10^{19}~{\\rm eV}$, in the normal source model. The GMF is not considered The strength of the structured EGMF is normalized to 0.1 $\\mu$G and that of an uniform turbulent field is 0 ({\\it upper panel}), 1 nG ({\\it middle panel}), and 10 nG ({\\it lower panel}). The numbers of events are set to be 200 ({\\it left}) and 500 events ({\\it right}) within the southern hemisphere.} \\label{fig_chi_offset} \\end{center} \\end{figure}  We exclusively adopt the AGASA results in this study. However, High Resolution Fly's Eye (HiRes) claims no significant small-scale anisotropy, contrary to AGASA \\cite{abbasi04}. This discrepancy is one of well-known problems in UHECR experiments. At present, this is not statistically significant due to the small number of observed events \\cite{yoshiguchi04}. It will be able to solved by new experiments with large aperture, such as Auger, Telescope Array \\cite{tahp}, and Extreme Universe Space Observatory \\cite{eusohp}. Hence, it is possible that these experiments do not observe sufficient small-scale anisotropy in the future. If Auger does not observe small-scale clusterings during 2007 (maybe it detects about 200 events above $4 \\times 10^{19}~{\\rm eV}$), the source number density is estimated at about $10^{-4}~{\\rm Mpc}^{-3}$ or more in the normal source model by definition of $\\chi^2$ in figure \\ref{fig_chidisL0}. It is comparable with number density of active galactic nuclei\\cite{loveday92}.   Our EGMF model within 100 Mpc has about 95\\% of volume without magnetic field. In our model, uniform magnetic field is not considered. According to an upper limit mentioned in section \\ref{introduction}, deflection angle of UHE protons with energy of $E$ during propagation of distance, $d$, is estimated as \\begin{equation} \\theta < 3^{\\circ} \\left( \\frac{E}{10^{20}~{\\rm eV}} \\right)^{-1} \\left( \\frac{d}{100~{\\rm Mpc}} \\right)^{1/2} \\left( \\frac{l_c}{1~{\\rm Mpc}} \\right)^{1/2} \\left( \\frac{B}{1~{\\rm nG}} \\right). \\end{equation} This deflection is expected to generate more isotropic arrival distribution of UHECRs. Therefore, uniform turbulent magnetic field affects determination of the source number density.  As a demonstration, we show $\\chi^2$ distribution for $B_{\\rm EG} = 0.1~\\mu{\\rm G}$, including the uniform field, in figure \\ref{fig_chi_offset}. The normal source model is adopted. The number of events is set to be 200 ({\\it left panel}), and 500 ({\\it right panel}). The strengths of the uniform turbulent field are 0 ({\\it upper panel}), 1 ({\\it middle panel}), and 10 nG ({\\it lower panel}). The distributions are shifted to lower value in stronger uniform turbulent field. The shifts are larger in more source number density. In stronger turbulent uniform magnetic field, number of events needed for the distinction between $10^{-5}$ and $10^{-6}~{\\rm Mpc}^{-3}$ is smaller while it becomes difficult to discriminate $10^{-4}$ from $10^{-5}~{\\rm Mpc}^{-3}$ due to the diffusion of cosmic rays. Strength of uniform EGMF is also important for estimating the source number density.  In this work, we adopt isotropic arrival distribution of UHECRs as a template of a future result since results of new experiments are still unpublished. Auger should detect a few times more number of UHE events than that of AGASA since it started observation about three years ago. Its result will provide us beneficial information on the nature of UHECR sources.   \\subsubsection*{Acknowledgements:} The work of H. T. is supported by Grants-in-Aid from JSPS Fellows. The work of K. S. is supported by Grants-in-Aid for Scientific Research provided by the Ministry of Education, Science and Culture of Japan through Research Grants S19104006."
    },
    "0710/0710.4482_arXiv.txt": {
        "abstract": "% We present recent results from the \\htmladdnormallinkfoot{Laboratory for Cosmological Data Mining}{http://lcdm.astro.uiuc.edu} at the National Center for Supercomputing Applications (NCSA) to provide robust classifications and photometric redshifts for objects in the terascale-class Sloan Digital Sky Survey (SDSS). Through a combination of machine learning in the form of decision trees, k-nearest neighbor, and genetic algorithms, the use of supercomputing resources at NCSA, and the cyberenvironment Data-to-Knowledge, we are able to provide improved classifications for over 100 million objects in the SDSS, improved photometric redshifts, and a full exploitation of the powerful k-nearest neighbor algorithm. This work is the first to apply the full power of these algorithms to contemporary terascale astronomical datasets, and the improvement over existing results is demonstrable. We discuss issues that we have encountered in dealing with data on the terascale, and possible solutions that can be implemented to deal with upcoming petascale datasets. ",
        "introduction": "We summarize work carried out as part of the Laboratory for Cosmological Data Mining, a partnership between the Department of Astronomy and the National Center for Supercomputing Applications (NCSA) at the University of Illinois at Urbana-Champaign (UIUC), in collaboration with the Automated Learning Group (ALG) at NCSA, and the Illinois Genetic Algorithms Laboratory at UIUC. This combination of expertise allows us to apply the full power of machine learning to contemporary terascale astronomical datasets. ",
        "conclusions": "Given the petascale datasets planned for the next decade, it is vitally important that contemporary data mining can be carried out successfully on this scale. In turn, this requires robust techniques on the terascale. We encountered numerous issues that were relevant to realizing this goal:  \\begin{itemize} \\item Because we are using tens to hundreds of parallel nodes and streaming many GB of data, D2K must be invoked via batch script, negating the advantages of its GUI interface. The resulting lack of an integrated cyberenvironment results in batch scripts that contain many tens of settings, manually set file locations and commands, making them prone to error. \\item Job submission is inflexible, subject to fixed wallclock times and numbers of nodes, unpredictable queuing times and no recourse if a job fails due a bug or hardware problem. \\item The large datasets must be stored on the Unitree mass storage system, which is occasionally subject to outages in access or significant wait times. In combination with the queuing system for batch jobs described above, this can make new scripts time-consuming to debug. \\item There is no way in which to fully explore the huge parameter space (more than $10^{15}$ combinations of settings for decision trees) of the machine learning algorithms. Genetic algorithms were used to optimize the training features, and could be used similarly to optimize the algorithm. \\item The present lack of fainter training data forces us to extrapolate in order to classify the whole SDSS. While the data and results we obtain are well-behaved, it will always be the case in astronomy that some form of extrapolation is ultimately required. This result is simply due to the fact that photometry will always be available several orders of magnitude fainter than spectroscopy, due to the physical difficulties in obtaining spectra of faint sources. Thus while our supervised learning represents a vital proof-of-concept over a whole terascale survey, ideally it should be extended with semi-supervised or unsupervised algorithms to fully explore the regions of parameter space that lie beyond the available training spectra. It is worth noting, however, that our training features, the object colors, are largely consistent beyond the limit of the spectroscopic training set. \\item The data size is such that integrating the SQL database with D2K via JDBC is impractical, and the data must be stored as flat files. As database engines become more sophisticated, however, it could in the future become possible to offload partial or entire classification rules to a database engine. Doing so, however, would require supercomputing resources for the database engine, which results in an entirely new class of problems. \\end{itemize}  In moving to the petascale, further issues include:  \\begin{itemize} \\item Conventional hardware, in the form of large clusters of multicore compute nodes, is scaleable to the petascale, however, field-programmable gate arrays (FPGAs), graphical processing units, and Cell processors may be more suited to many data mining tasks, due to their embarrassingly parallel nature. The LCDM group, in collaboration with the Innovative Systems Laboratory at NCSA, has demonstrated results on FPGAs using an SRC-6 MAPE system (Brunner, Kindratenko, \\& Myers 2007), which include running a kNN algorithm, although the implementation is not trivial because the algorithm must be rewritten. \\item For many applications on the petascale, the performance becomes I/O limited. This is quantified by Bell, Gray, \\& Szalay (2006), who apply Amdahl's law (Amdahl 1967) that one byte of memory and one bit per second of I/O are required for each instruction per second, to predict that a petaflop-scale system will require one million disks at a bandwidth of $100~{\\mathrm{MB~s^{-1}}}$ per disk. They also state that data should be stored locally (i.e., not transferred over the internet), if the task requires less than 100,000 CPU cycles per byte of data. Many contemporary scientific applications are such that local storage is favored by over an order of magnitude. \\end{itemize}"
    },
    "0710/0710.5167_arXiv.txt": {
        "abstract": "We present results from numerical relativity simulations of equal mass, non-spinning binary black hole inspirals and mergers with initial eccentricities $e\\le 0.8$ and coordinate separations $D \\ge 12\\,M$ of up to 9 orbits (18 gravitational wave cycles). We extract the mass $M_\\mathrm{f}$ and spin $a_\\mathrm{f}$ of the final black hole and find, for eccentricities $e\\lesssim 0.4$, that $a_\\mathrm{f}/M_\\mathrm{f} \\approx 0.69$ and $M_\\mathrm{f}/M_\\mathrm{adm} \\approx 0.96$ are {\\em independent} of the initial eccentricity, suggesting that the binary has circularized by the merger time. For $e \\gtsim 0.5$, the black holes plunge rather than orbit, and we obtain a maximum spin parameter $a_\\mathrm{f}/M_\\mathrm{f} \\approx 0.72$ around $e = 0.5$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.2437_arXiv.txt": {
        "abstract": "% We report the discovery of a 16--20\\,$M_\\mathrm{Jup}$ radial velocity companion around the very young ($\\sim$3\\,Myr) brown dwarf candidate Cha\\,H$\\alpha$\\,8 (M5.75--M6.5). Based on high-resolution echelle spectra of Cha\\,H$\\alpha$\\,8 taken between 2000 and 2007 with UVES at the VLT, a companion was detected through RV variability with a semi-amplitude of 1.6\\,km\\,s$^{-1}$. A Kepler fit to the data yields an orbital period of the companion of 1590~days and an eccentricity of $e$=0.49. A companion minimum mass $M_2\\sin i$ between 16 and 20\\,$M_\\mathrm{Jup}$ is derived when using model-dependent mass estimates for the primary. The mass ratio $q\\equiv M_2/M_1$ might be as small as 0.2 and, with a probability of 87\\%, it is less than 0.4. Cha\\,H$\\alpha$\\,8 harbors most certainly the lowest mass companion detected so far in a close ($\\sim$ 1\\,AU) orbit around a brown dwarf or very low-mass star. From the uncertainty in the orbit solution, it cannot completely be ruled out that the companion has a mass in the planetary regime. Its discovery is in any case an important step towards RV planet detections around BDs. Further, Cha\\,H$\\alpha$\\,8 is the fourth known spectroscopic brown dwarf or very low-mass binary system with an RV orbit solution and the second known very young one. ",
        "introduction": "Search for planetary or brown dwarf (BD) companions to BDs are of primary interest for understanding planet and BD formation. There exists no widely accepted model for the formation of BDs (e.g. Luhman et al.~2007). The frequency of BDs in multiple systems is a fundamental parameter in these models. However, it is poorly constrained for close separations: Most of the current surveys for companions to BDs are done by direct (adaptive optics or HST) imaging and are not sensitive to close binaries ($a\\la1$\\,AU and $a\\la10$\\,AU for the field and clusters, respectively), and found preferentially close to equal mass systems (e.g. Bouy et al.~2003). Spectroscopic monitoring for radial velocity (RV) variations provides a means to detect close systems. The detection of the first spectroscopic BD binary in the Pleiades, PPl\\,15 (Basri \\& Mart\\'\\i n 1999), raised hope to find many more of these systems in the following years. However, the number of confirmed close companions to BDs and very low-mass stars (VLMS, $M \\leq 0.1\\,M_{\\odot}$) is still small. To date, there are three spectroscopic BD binaries known, i.e. for which a spectroscopic orbital solution has been derived: PPl\\,15, the very young eclipsing system 2M0535-05 (Stassun, Mathieu \\& Valenti 2006), and a binary within the quadruple GJ\\,569 (Zapatero Osorio et al.~2004; Simon, Bender \\& Prato 2006). They all have a mass ratio close to unity. In particular, no RV planet of a BD/VLMS has been found yet. If BDs can harbor planets at a few AU distance is still unknown. Among the more than 200 extrasolar planets that have been detected around stars by the RV technique, 6 orbit stellar M-dwarfs showing that planets can form also around primaries of substantially lower mass than our Sun. Observations hint that basic ingredients for planet formation are present also for BDs (e.g. Apai et al.~2005). However, the only planet detection around a BD is a very wide 55\\,AU system (2M1207, Chauvin et al.~2005), which is presumably formed very differently from the Solar System and RV planets.  RV surveys for planets around such faint objects, as BD/VLMS are, require monitoring with high spectral dispersion at 8--10\\,m class telescopes. While being expensive in terms of telescope time, this is, nevertheless, extremely important for our understanding of planet and BD formation. We report here on the recent discovery of a very low-mass companion orbiting the BD candidate Cha\\,H$\\alpha$\\,8, which was detected within the course of an RV survey for (planetary and BD) companions to very young BD/VLMS in the Chamaeleon\\,I star forming region (Joergens \\& Guenther 2001; Joergens 2006; Joergens \\& M\\\"uller 2007). ",
        "conclusions": "The companion of Cha\\,H$\\alpha$\\,8 has most certainly a much smaller mass than that of any previously detected close companion of a BD/VLMS. For comparison, all other known spectroscopic BD binaries have mass ratios $>$ 0.6 and the lowest mass BD in these systems has 54\\,$M_\\mathrm{Jup}$ (Stassun et al.~2006). The discovery of the RV companion of Cha\\,H$\\alpha$\\,8 with its RV semi-amplitude of only 1.6\\,km\\,s$^{-1}$ is an important step towards RV detections of planets around BD/VLMS. In fact, from the uncertainty of the orbit solution, it cannot be excluded that the companion of Cha\\,H$\\alpha$\\,8 has a mass in the planetary regime ($<13\\,M_\\mathrm{Jup}$). Follow-up RV measurements at the next phase of periastron will clarify this.  With a semi-major axis of about 1\\,AU, the companion of Cha\\,H$\\alpha$\\,8 orbits at a much closer orbital distance than most companions detected around BDs so far. In particular, its orbit is much closer than that of recently detected very low-mass companions of BDs, like that of 2M1207 (55\\,AU; Chauvin et al.~2005) and that of CHXR\\,73 (210\\,AU; Luhman 2006; see Luhman, this volume).  The favored mechanisms for stellar binary formation, fragmentation of collapsing cloud cores or of massive circumstellar disks, seem to produce preferentially equal mass components, in particular for close separations (e.g. Bate et al.~2003). Thus, they have difficulties to explain the formation of the small mass ratio system Cha\\,H$\\alpha$\\,8. However, we know that close \\emph{stellar} binaries with small mass ratios do exist as well (e.g. q=0.2, Prato et al.~2002), and without knowing the exact mechanism by which they form, it might be also an option for Cha\\,H$\\alpha$\\,8. Considering the small mass of the companion of Cha\\,H$\\alpha$\\,8, a planet-like formation could also be possible. Giant planet formation through core accretion might be hampered for low-mass primaries, like M dwarfs, by long formation time scales (Laughlin, Bodenheimer \\& Adams 2004; Ida \\& Lin 2005), though, recent simulations hint that it can be a faster process than previously anticipated (Alibert et al.~2005). On the other hand, giant planets around M dwarfs might form by disk instability (Boss 2006a, 2006b), at least in low-mass star-forming regions, where there is no photoevaporation of the disk through nearby hot stars (e.g. Cha\\,I). The companion of Cha\\,H$\\alpha$\\,8 could have been formed through disk instability, either in situ at 1\\,AU or, alternatively, at a larger separation and subsequent inwards migration.  Cha\\,H$\\alpha$\\,8 is extremely young (3\\,Myr) and its study allows insight into the formation and early evolution at and below the substellar limit. Cha\\,H$\\alpha$\\,8 is only the 2nd known very young BD/VLM spectroscopic binary (after 2M0535-05, Stassun et al.~2006). When combined with angular distance measurements or eclipse detections, spectroscopic binaries allow valuable dynamical mass determinations. The mass is the most important input parameter for evolutionary models, which rely for $<$0.3\\,M$_\\mathrm{\\odot}$, only on the two masses determined for 2M0535-05. In order to measure absolute masses of both components of Cha\\,H$\\alpha$\\,8, it is required to resolve the spectral lines of both components. We will try this with CRIRES/VLT at IR wavelength, were the contrast ratio between primary and secondary is smaller. Having an orbital separation of the order of 13\\,milli\\,arcsec, the spatial resolution of current imaging instruments is not sufficient to directly resolve Cha\\,H$\\alpha$\\,8. However, it might be possible to detect the astrometric signal caused by the companion, e.g. with NACO/VLT. This would allow measurement of the inclination of the orbital plane and, therefore, breaking the $\\sin i$ ambiguity in the companion mass."
    },
    "0710/0710.0426_arXiv.txt": {
        "abstract": "We present the results of NICMOS imaging of two massive galaxies photometrically selected to have old stellar populations at $z\\sim2.5$.  Both galaxies are dominated by apparent disks of old stars, although one of them also has a small bulge comprising about 1/3 of the light at rest-frame 4800 \\AA.  The presence of massive disks of old stars at high redshift means that at least some massive galaxies in the early universe have formed directly from the dissipative collapse of a large mass of gas. The stars formed in disks like these may have made significant contributions to the stellar populations of  massive spheroids at the present epoch. ",
        "introduction": "Considerable observational evidence has built up over the past few years that a substantial fraction of the massive galaxies around us today were already massive at very early epochs. This evidence comes primarily from three sources: \\begin{itemize} \\item  Studies of local massive elliptical galaxies indicate that the stars in the most massive galaxies generally formed very early and over very short time intervals \\citep{pee02,tho05,nel05,ren06}.  Stars in less massive spheroids formed, on average, later and over longer time spans. \\item Massive galaxies in clusters show little evidence for significant evolution up to at least redshift $\\sim1$ \\citep[\\eg][]{deP07,sca07}. \\item Direct observations of massive galaxies at redshifts $\\gtrsim1.5$ that are dominated by already old stellar populations show that significant numbers of massive galaxies were in place at even earlier epochs \\citep*[\\eg][]{sto04,mcC04,vanD04,lab05,dad05, red06,pap06,kri06,abr07}. \\end{itemize}  Although the existence of massive galaxies at high redshifts is now well documented, there have been only a few high-resolution studies of their morphologies (e.g., \\citealt{yan03,sto04,zir07,tof07}). Morphologies are important, because they may well retain signs of the formation history of the galaxies.  This is particularly true for galaxies that show little or no recent star formation, so that we are able to observe relatively clean examples of the stellar population that formed earliest and that comprises the bulk of the mass of the galaxy. In this paper, we present deep {\\it Hubble Space Telescope} ({\\it HST}) NICMOS imaging of two galaxies with virtually pure old stellar populations at $z\\sim2.5$.  In \\S~2, we briefly recount how these galaxies were selected.  In \\S~3, we describe the observations and reduction procedures.  In \\S~4 and \\S~5, we analyze model fits to the images to determine morphologies, and in \\S~6 we discuss the implications of our conclusions.  We assume a flat cosmology with $H_0 = 73$ \\kms\\ Mpc$^{-1}$ and $\\Omega_M = 0.28$. ",
        "conclusions": "} Table~\\ref{tab1} summarizes the parameters for the two galaxies. The morphologies of both \\ea\\ and \\eb\\ appear to be dominated by disks of old stars. However, the disks are quite different in scale. \\eb, at least, also appears to have a small bulge comprising about 1/3 of the total light in the F160W filter ($\\sim4800$ \\AA, rest frame). We cannot exclude the possibility that \\ea\\ also has a weak bulge, with up to $\\sim15$\\% of the total light in the F160W filter; indeed, if the slight apparent difference in morphology between the F160W and F110W images is real, such a difference would seem to favor this possibility.  But it is the presence of massive, old disks that continues to give the strongest constraint on formation mechanisms.  Such disks also have been seen at redshifts $\\sim1.5$, where normal ellipticals with $r^{1/4}$-law profiles are also found \\citep*{iye03, cim04, yan04, fu05, sto06, mcG07}. It is difficult to imagine that these massive disks could have formed via any process other than the dissipative collapse of a large cloud of gas.  Such disks are also unlikely to have survived major merging events, although the bulge component in \\eb\\ may testify to either some level of minor merging activity or bulge building via disk instabilities.  For galaxies at $z\\sim2.5$, the evidence for a dominant old stellar population depends on the inflection in the SED shortward of the $H$ band, and establishing this inflection with optical/near-IR photometry depends on the relatively short baseline from the $H$ to the $K$ band.  Furthermore, at the present epoch, essentially all strongly disk-dominated galaxies show evidence for continued star formation.  It is therefore not too surprising that claims of passive disks at high redshift should be doubted \\citep[\\eg][]{pie05}.  However, as Fig.~\\ref{4c05sed} shows, {\\it Spitzer} IRAC data is entirely consistent with the SED of a moderately old stellar population, and no plausible SED incorporating very recent star formation combined with dust would fit the observed photometry. We have recently also obtained IRAC imaging of the field of 4C\\,23.56, and our analysis of these data shows that the IRAC photometry falls squarely on our best-fit solar-metallicity BC03 model determined from the optical/near-IR photometry alone: an instantaneous burst with an age of 2.6 Gyr and an extinction $A_V=0.16$ mag \\citep{sto07}. Using the more recent preliminary CB07 models, with their improved treatment of AGB stars, we obtain a stellar population age of 2.8 Gyr with $A_V=0$.  Again, no plausible model with significant star formation and reddening would fit these data. \\begin{deluxetable}{l c c c c} \\tablewidth{0pt} \\tablecaption{Model Parameters for \\ea\\ and \\eb} \\tablehead{ \\colhead{Galaxy} & \\colhead{Filter} & \\colhead{S\\'{e}rsic $n$} & \\colhead{$r_e$} & \\colhead{$r_e$}\\\\ & & & (\\arcsec) & (kpc) } \\startdata 4C\\,23.56\\,ER1 & F110W & $1.03\\pm0.10$ & $0\\farcs28\\pm0\\farcs02$ & $2.2\\pm0.2$ \\\\ 4C\\,23.56\\,ER1 & F160W & $1.52\\pm0.06$ & $0\\farcs24\\pm0\\farcs01$ & $1.9\\pm0.1$ \\\\ &               & 1.00\\tablenotemark{a} & $0\\farcs89\\pm0\\farcs09$ & $7.1\\pm0.8$ \\\\ \\raisebox{1.5ex}[0pt]{4C\\,05.84\\,ER1} & \\raisebox{1.5ex}[0pt]{F160W} & 4.00\\tablenotemark{a} & $0\\farcs37\\pm0.20$ & $3.0\\pm1.6$ \\\\ \\enddata \\tablenotetext{a}{The S\\'{e}rsic indices for the two model components for 4C\\,05.84\\,ER1 have been fixed at these values, which correspond to exponential and $r^{1/4}$-law profiles, respectively.} \\label{tab1} \\end{deluxetable}   Masses for these galaxies can be estimated from the model fits.  Assuming solar metallicities and a \\citet{cha03} initial mass function, we obtain a mass of $3.9\\times10^{11} M_{\\odot}$ for \\ea\\ and $3.3\\times10^{11} M_{\\odot}$ for \\eb (assuming the model at $z=2.4$ with $A_V=0.58$).  While the stellar-population age of \\eb\\ indicates that the last major star-formation episode occurred at $z\\sim3.7$, when the universe was $\\sim1.8$ Gyr old, \\ea\\ has a stellar-population age that is formally slightly greater than the age of the universe at $z=2.483$.  Clearly, the likely errors in the age determination and the usual caveats regarding the age-metallicity degeneracy mitigate any implied paradox.  Nevertheless, this massive galaxy must have formed at a very high redshift.  Models with [Fe/H] $= +0.4$ give an age of 1.9 Gyr, but with a significantly worse fit.  It therefore seems likely that galaxy formation models will have to allow for the presence of early-forming massive disks.  This means that, at least in some dense regions, it has been possible to form $\\sim3\\times10^{11}$ $M_{\\odot}$ of stars within a relatively short time via dissipative collapse and without the aid of major mergers.  While our selection criteria have ensured that the galaxies we have discussed here comprise essentially pure old stellar populations, they may well be representative of many massive galaxies at high redshift, most of which would not be in our sample if they retained even tiny amounts of residual star formation or if they had had any significant star formation within a few hundred Myr prior to the epoch at which we observe them.  \\eb\\ has a luminosity and an effective radius that are similar to those of many local galaxies.  Our best-fitting S\\'{e}rsic model has $r_e=6.3$ kpc.  For comparison, for galaxies of similar mass from the Sloan survey with S\\'{e}rsic $n<2.5$, \\citet{she03} find $r_e=7.2^{+2.9}_{-2.1}$ kpc.  This galaxy could become, with passive evolution and perhaps a few minor mergers to increase the bulge-to-disk ratio somewhat, a typical S0 galaxy at the present epoch.  On the other hand, we do not see galaxies like \\ea\\ at the present epoch.  By the prescription of \\citet{she03}, a low-S\\'{e}rsic-index galaxy with the mass of \\ea\\ would have $r_e=7.6^{+3.1}_{-2.2}$ kpc, but \\ea\\ actually has $r_e=1.9\\pm0.1$ kpc. This means that the stellar mass surface density is much higher than for local galaxies, a result that has also been found for other samples of distant red galaxies \\citep[\\eg][]{tru06,tof07}.  It would seem that the only likely path for such galaxies to evolve to objects consistent with the local population of galaxies is through dissipationless mergers.  There is recent evidence that the most massive galaxies in the local universe are likely the result of dry mergers of galaxies with stars that are already old and with very little gas \\citep[e.g.,][]{ber07}.  With the constraint that these merging components must themselves mostly be fairly massive (to avoid a large dispersion and flattening in the observed color---magnitude relation for present-day massive galaxies, \\eg\\ \\citealt*{bow98}), it seems possible that these early massive disks may well be among the sources for the old stars that today are found in the most massive elliptical galaxies."
    },
    "0710/0710.5035_arXiv.txt": {
        "abstract": "Starting from a sample of SDSS quasars appearing also in the 2MASS survey, we study the continuum properties of  $\\sim 1000$ objects observed in 8 bands, from NIR to UV. We construct the mean spectral energy distribution (SED) and compare and contrast the continua of radio loud (RLQ) and radio quiet (RQQ) objects. The SEDs of the two populations are significantly different, in the sense that RLQs are redder with power law spectral indices $\\langle\\alpha_{\\rm RLQ}\\rangle=-0.55\\pm0.04$ and $\\langle\\alpha_{\\rm RQQ}\\rangle=-0.31\\pm0.01$ in the spectral range between $10^{14.5}$ and $10^{15.35}$ Hz. This difference is discussed in terms of different extinctions, different disc temperatures, or slopes of the non-thermal component.% ",
        "introduction": "A substantial effort has been dedicated to construct the spectral energy distributions (SEDs) for sizeable samples of quasars over the whole accessible range of the electromagnetic spectrum % (see e.g.~Sanders et al.~1989; Francis et al.~1991; Elvis et al.~1994; Richards et al.~2006). % However, only a few papers focus on the comparison between the SEDs of radio loud (RLQ) and radio quiet (RQQ) quasars. Elvis et al.~1994 propose an overall spectrum from a sample of 47 quasars, divided in RLQs and RQQs, which shows that no distinction between the SEDs of the two subsamples is apparent in the range $100\\mu - 1000$\\AA. As noted by these authors, the considered sample is biased towards X-ray and optically bright quasars. Some indication of a possible difference between the NIR to optical colors of RLQs and RQQs in the 2MASS catalogue was reported by Barkhouse \\& Hall (2001). Francis, Whiting \\& Webster (2000) found that the optical--NIR continuum is significantly redder in radio selected RLQs from the PKS Half-Jansky Flat-Spectrum survey than in optically selected RQQs from the Large Bright Quasar Survey (LBQS). These and other works (e.g.~Kotilainen et al.~2007) indicate that radio loud objects are possibly redder than their radio quiet counterparts, but the samples seriously suffer of selection effects against red radio quiet quasar, because RQQs are mostly collected from optical selected samples (e.g.~the LBQS, the first selection criterion of which is ``blue color'' of candidates; see Hooper, Impey \\& Foltz 1997). White et al. suggest that redder quasars from the Sloan Digital Sky Survey (SDSS) are likely to be more radio-powerful than bluer objects, and this is an indication that the red color of radio loud objects is not completely due to the bias against red RQQs in optically selected samples, since the two classes of objects derive from the same survey.  Because of the importance of the distinction between RLQs and RQQs in the Unified Models of AGN, these results suggest and motivate a study aimed at investigating the properties of the continuum emission of quasars in the UV to NIR region, in order to compare and contrast the SEDs of radio loud and radio quiet objects.  %  In Section \\ref{secsample} we focus on the quasar sample selection criteria and on the host galaxy contribution. In section \\ref{SED} we consider the SED construction and discuss the spectral shape. % The RLQ and RQQ SEDs are then compared and contrasted. In the last section we provide a summary and a discussion of the results.     Throughout this paper, we adopt a concordant cosmology with $H_0=70$ ~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_m=0.3$ and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "\\label{sum_disc} The aim of the present paper is the study of the continuum emission in the NIR to UV region of a quasar sample, in order to compare and contrast the SEDs of the RLQ and RQQ subsamples. We selected a sample of quasars with both SLOAN and 2MASS detection, to study a spectral range from the NIR (for low-$z$ objects) to the UV (for high-$z$ QSOs). The sample consists of 887 objects, of which 113 RLQs and 774 RQQs.  For each subsample we constructed the mean SED and evaluated the average spectral index in the NIR to UV region. The slope of the underlying power-laws is significantly different: the spectral indices of the two population are statistically different at more than 99\\% confidence, both for the whole sample and dividing the sample in redshift bins. The difference is present also considering samples well matched in absolute magnitude and redshift.  If the blue bump is due to the superposition of black body emissions from an accretion disc, then the color difference between RLQs and RQQs should be interpreted in terms of different mean temperatures, in the sense that RQQs are hotter.   \\begin{figure} \\includegraphics[width=0.45\\textwidth]{Fig6} \\caption[]{Comparison of the SED of RQQs (dashed line) to the SED of RLQs (dotted line). When RLQs are corrected for an additional dust extinction ($\\Delta A_{V}=0.16$mag, solid line), no difference is apparent.} \\label{dust} \\end{figure}  We first examine the possibility that the difference in the SEDs of RLQs and RQQs is due to an enhanced dust extinction in radio loud objects, as it has been suggested e.g.~by Francis et al.~(2000). In Fig. \\ref{dust} we compare the SED of RQQs to the SED of RLQs, when RLQs are additionally corrected for dust extinction with respect to RQQs: the offset between the SEDs of the two population is minimized adopting $\\Delta A_{V}=0.16$mag. The differences are now completely negligible, supporting the hypothesis that RL objects are more subject to dust extinction than RQQs. The problem would be reconducted to understand why RLQs are more extincted than RQQs. {\\it A priori} this could be related to a difference in the inclination angle distribution. % Alternatively, it could be that the conditions of dust production are related to those justifying large radio emission. %  We then focus on the possibility that the difference of disc temperature of RLQs and RQQs is real. Since the temperature of the inner disc scales as $M_{\\rm BH}^{\\,\\,\\,-0.25}$ (e.g.~Shakura \\& Sunyaev 1973), the difference may be attributed to the fact that the black holes of radio loud quasars are supposedly more massive (e.g.~Dunlop et al.~2003; Falomo et al.~2004; Labita et al.~2007). %  The color difference may be linked to the BH spin (Stawarz, Sikora \\& Lasota 2007), as radio emission is usually ascribed to a faster spinning. However, spinning BHs are expected to have a shorter last stable orbit radius, and then a hotter disc, inconsistently with our results. %  Obviously it is possible that the non--thermal (power law) continua of RLQs and RQQs are intrinsically different. Supposing that the non--thermal component accounts for 80\\% at $10^{14.5}$Hz, while at $10^{15.35}$Hz (in the Big Blue Bump) the thermal component is dominant and accounts for 80\\%, our data are consistent with a picture where the underlying power--laws of RLQs and RQQs have $\\alpha_{\\rm RLQ}=-1.2$ and $\\alpha_{\\rm RQQ}=-1.0$ respectively, whereas the thermal bumps are indistinguishable. The problem would be reconducted to understand why the non--thermal continuum of RLQs is softer than RQQs. % It has been suggested (i.e.~Francis et al.~2000) that in radio loud samples there is a significant chance of synchrotron contamination of the rest-frame R--band nuclear luminosities, due to the presence of the relativistic jets. This effect explains the red color of QSOs in  high-frequency selected radio samples (i.e. the PKS survey, e.g.~Francis et al.~2001), that suffer from a bias towards pole-on radio sources which are relativistically boosted above the survey flux limit, but no obvious explanation can be invoked to justify the color difference in our sample."
    },
    "0710/0710.0424.txt": {
        "abstract": "A description is given of the algorithms implemented in the AstroBEAR adaptive mesh refinement code for ideal magnetohydrodynamics.  The code provides several high resolution, shock capturing schemes which are constructed to maintain conserved quantities of the flow in a finite volume sense.  Divergence free magnetic field topologies are maintained to machine precision by collating the components of the magnetic field on a cell-interface staggered grid and utilizing the constrained transport approach for integrating the induction equations.  The maintenance of magnetic field topologies on adaptive grids is achieved using prolongation and restriction operators which preserve the divergence and curl of the magnetic field across co-located grids of different resolution.  The robustness and correctness of the code is demonstrated by comparing the numerical solution of various tests with analytical solutions or previously published numerical solutions obtained by other codes. ",
        "introduction": "The development of efficient and accurate numerical algorithms for astrophysical flow has become of great interest to the astrophysical community.  Variable resolution approaches have provided an avenue for efficient simulation of hydrodynamical flow including multi-physical effects which involve substantial variation in length scale.  Adaptive Mesh Refinement (AMR) has been recognized as one of the most versatile and efficient approaches to enable the simulation of multi-scale phenomena for which fixed-resolution simulation is either impractical or impossible.  AMR discretizations employ a hierarchy of grids at different levels.  High resolution is applied only to those regions of the flow which would otherwise be subject to unacceptably large truncation error. The utility of the AMR approach is underscored by the extensive list of codes that are targeted toward astrophysical research which utilize AMR.  The list includes AstroBEAR, Enzo \\citep{enzo}, Flash \\citep{flash}, Orion \\citep{truelove,klein,crockett}, Nirvana \\citep{nirvana-amr}, Ramses \\citep{ramses}, RIEMANN \\citep{balsara-amr} and AMRVAC \\citep{amrvac} and the list of codes for which the development AMR capability is in progress including Athena \\citep{gardiner} and Pluto \\citep{pluto}.  The simulation of magnetized flow is of particular interest to astrophysical researchers owing to the utility of numerical magnetohydrodynamics (MHD) in modeling a wide range of astrophysical phenomena.  The leading line of recent research in this area has focused on the application of higher order Godunov methods to numerical MHD \\citep{rj-rs,balsara-rs}.  The conservative formulation and proper upwinding employed by these methods enable accurate simulation of strongly supersonic flow.  Because of this unique capability, such methods are often referred to ``high resolution shock capturing'' (HRSC) methods.  While HRSC methods have long been recognized as the defacto standard for the simulation of supersonic hydrodynamical phenomena, their popularity among researchers interested in magnetized flow has been slowed because standard HRSC approaches to MHD fail to maintain the solenoidality constraint on the magnetic field ($\\vec{\\nabla} \\cdot \\vec{B} = 0$).  If not corrected, local divergences in the magnetic field arising from this short coming usually grow rapidly, causing anomalous magnetic forces and unphysical plasma transport which eventually destroys the correct dynamics of the flow \\citep{brackbill}.  Early practitioners of numerical MHD therefore relied heavily on finite difference methods as employed by codes such as Zeus \\cite{zeus-mhd} which maintain the solenoidality constraint exactly despite their inferior shock capturing ability \\citep{falle}. Later works focused on improved HRSC which either eliminate the development of divergences in the magnetic field, or mitigate the effect of local divergence errors on the dynamics of the flow.  In one such approach, a projection operator is devised, usually by solving a Poisson equation, which removes numerical divergences from the grid after each time-step \\citep{balsara98,jiang,kim,zachary,rjf}.  The primary limitation of this approach is that non-trivial boundary value problems become indeterminate \\citep{rj-ct}.  In the second so called ``8-wave'' approach first explored by \\cite{powell}, alternative formulations of the MHD equations are constructed to prevent the local build up of magnetic divergence by advecting monopoles to other regions of the grid where they are of less consequence to the dynamics of the flow.  The work of \\cite{dender} augments this approach by adding source terms to the system which act to counter the effect of local divergence in the magnetic field on the dynamics of the flow.  A third approach known as constrained transport utilizes a multidimensional, divergence-preserving update procedure for the magnetic field components which are collated on a staggered grid centered on the computational volume interfaces. \\citep{evans,rj-ct,balsara-spicer,dai,londrillo,nirvana}.  This approach has been shown to provide the most accurate results in the tests of \\cite{toth} and \\cite{balsara-comp}.  The combination of AMR spatial discretizations with HRSC would seem a natural choice in order to satisfy the desire for a high accurate, computationally efficient and versatile strategy for the simulation of magnetized plasma flow.  The implementation of the aforementioned adaptations to HRSC methods for MHD in an AMR framework however, poses several challenges.  Divergence cleaning schemes utilizing a Poisson projection operator are ill-suited for AMR applications owing to difficulty in handling the projection step along internal boundaries on a patchwork of grids at different resolutions. \\cite{powell} found that the application of their 8-wave method on an AMR grid hierarchy, local divergence errors on one level of the AMR hierarchy caused local divergence of comparable magnitude on all levels. The main drawback of this method in an AMR context is that the unphysical effects of local divergences in the magnetic field are not diminished by the application of additional refinement.  \\cite{crockett}, however, have constructed an approach suitable for AMR applications which combines an approximate projection operator with the divergence advection and dampening the effects of local divergence of \\cite{powell} and \\cite{dender}.  Retaining the divergence-free property of the solution obtained through the application of the constrained transport approach on hierarchical grids requires application of a divergence-preserving prolongation operator which interpolates the magnetic field from a coarse mesh a fine mesh, a divergence-preserving restriction operator which maps the fine mesh magnetic field onto a coarser mesh and furthermore requires that the evolution of the magnetic field be consistent between collocated meshes of different resolution.  Two approaches to these challenges have emerged.  \\cite{balsara-amr} has generalized the divergence free reconstruction procedure of \\cite{balsara-spicer} to devise a prolongation operator based on a piece-wise quadratic interpolation procedure that is divergence preserving in the RIEMANN MHD code.  \\cite{li} present an adaptation of Balsara's procedure that simplifies its implementation for problems involving arbitrary refinement ratios.  \\cite{toth-roe} have devised a prolongation operator by solving an algebraic system which enforces the maintenance of the volume average curl and divergence between grids of different resolution in the AMR hierarchy.  In this paper we provide a concise description of the algorithms and tests of the AstroBEAR HRSC AMR MHD code.  AstroBEAR is comprised of several numerical solvers, integration schemes, and radiative cooling modules for astrophysical fluids.  The code's AMR capability is derived from the AMR engine of the BEARCLAW boundary embedded adaptive mesh refinement package for conservation laws.  This code utilizes the constrained transport approach to adapting HRSC methods to the MHD system of equations.  To our knowledge, AstroBEAR is the first AMR code to utilize the prolongation operator of \\cite{toth-roe} to maintain the $\\vec{\\nabla} \\cdot \\vec{B} = 0$ constraint.  By combining multi-physics capabilities relevant to simulation astrophysical plasma flow, AMR, and a wide selection of HRSC integration procedures, AstroBEAR will serve as a valuable research tool.  The authors intend that this paper will serve as a reference for future works that apply the code and provide a concise recipe for robust, reliable and accurate HRSC solution strategies for MHD on AMR grid hierarchies.  In \\S \\ref{method} we describe the several HRSC schemes and divergence preservation strategies implemented in the code.  In \\S \\ref{AMR} we provide an overview of the AMR algorithm, highlighting the stages which require special treatment of the magnetic field.  In \\S \\ref{prores} we provide a concise description of the prolongation, restriction and coarse to fine refluxing procedures required to preserve the divergence and consistency of the magnetic field across an AMR hierarchy of grids.  In \\S \\ref{tests} we comment on the results of several test and example problems with particular emphasis on the relative strengths and weaknesses of the various HRSC schemes implemented in the code.  In \\S \\ref{conclusion} we provide a synopsis and discussion of the main results of the paper. ",
        "conclusions": "\\label{conclusion} The staggered grid constrained transport schemes described in this paper enable the application of high resolution shock capturing methods to magnetized flow.  In this paper we have demonstrated that a wide cross section of high resolution shock capturing schemes for general conservation laws may be adapted for magnetized flow while preserving the divergence free constraint on the magnetic field topology exactly by conserving the surface integral of magnetic flux over each computational cell in an upwind fashion.  The use of such schemes on multi-resolution AMR grids is encumbered by the requirement that the prolongation and restriction steps preserve the divergence free topology of the magnetic fields. In this paper we have described the application of prolongation and restriction operators which maintain such topologies to machine precision.  The numerical schemes discussed here have been implemented and tested in the AstroBEAR adaptive mesh refinement code.  The code utilizes a modular design, enabling the user to choose from various methodologies to tailor the numerical integration strategy to the requirements of the application at hand.  The robustness of this approach to high resolution, shock capturing MHD on AMR grid structures, and relative advantages of the various numerical schemes implemented in the code are demonstrated in the context of several numerical example problems. The description of the numerical schemes presented in this paper provides a concise recipe for their implementation which will enable the reproduction of these outcomes by other researchers and the interpretation of future works derived from the AstroBEAR code."
    },
    "0710/0710.5173_arXiv.txt": {
        "abstract": "The results of a  world-wide coordinated observational campaign on the broad-lined Type Ic SN~2003jd are presented.  In total, 74 photometric data  points  and  26   spectra  were  collected  using  11  different telescopes.   SN  2003jd  is  one  of the  most  luminous  SN~Ic  ever observed.   A  comparison  with  other  Type  Ic  supernovae  (SNe~Ic) confirms  that  SN  2003jd  represents an  intermediate  case  between broad-line events (2002ap, 2006aj),  and highly energetic SNe (1997ef, 1998bw, 2003dh, 2003lw), with an ejected  mass of $M_{ej} = 3.0 \\pm 1$ M$_{\\odot}$ and  a kinetic energy  of $E_{k}({\\rm tot})  = 7_{-2}^{+3} \\times ~10^{51}$~erg.   SN~2003jd is similar to SN~1998bw  in terms of overall luminosity, but it is closer to SNe 2006aj and 2002ap in terms of light-curve shape and spectral evolution. The comparison with other SNe~Ic,  suggests that  the $V$-band  light  curves of  SNe~Ic can  be partially homogenized by introducing  a time stretch factor.  Finally, due to the similarity  of SN~2003jd to the SN~2006aj/XRF~060218 event, we discuss the possible connection of SN 2003jd with a GRB. ",
        "introduction": "\\label{parintroduction}  In the past decade, the  discovery of the connection of some gamma-ray bursts (GRBs) with Type Ic supernovae \\citep[SNe~Ic; see][for a review of  supernova classification]{filippenko97}  boosted  interest in  the study   of   this   type  of   SN   \\citep{galama98,stanek03,hjorth03, malesani04,pian06,soderberg06a,campana06}.  In the  current paradigm there is a  distinction between \\emph{normal} SNe~Ic   of   relatively   low   kinetic   energy   \\citep[$E_{51}   = E_K/(10^{51}~{\\rm  erg}) \\approx  1$;][]{nomoto01}, of  which SN~1994I represents  a  prototype,  and  high-energy, broad-lined  events  like SN~1998bw    \\citep[$E_{51}~>10$;][]{iwamoto98,maeda02}    which   are associated  with  GRBs.   To  understand whether  these  are  separate subclasses  or  extreme cases  of  a  continuous  distribution, it  is necessary  to   study  in  detail  intermediate  cases   such  as  the broad-lined   SN  2002ap   \\citep{galyam02,foley03},   which  had   an intermediate explosion  energy \\citep[$E_{51} =  4-10$; ][]{mazzali02} and  was not  connected with  a GRB,  or the  broad-lined,  low energy ($E_{51}   \\approx  2$)   SN   2006aj  associated   with  XRF   060218 \\citep{mazzali06a,mazzali07,maeda07}.  Here we  discuss the case  of SN~2003jd, a broad-lined  object showing clear evidence  of an asymmetric explosion  \\citep{mazzali05} but with no  confirmed GRB  connection  \\citep{gcn2434,gcn2439}.  Asymmetry  is probably a key factor in understanding the diversity of SNe~Ic.  SN 2003jd  was discovered \\citep{burket03}  on 2003 Oct. 25  (UT dates are  used throughout this  paper) with  the Katzman  Automatic Imaging Telescope  (KAIT)   during  the  Lick   Observatory  Supernova  Search \\citep{filippenko01,filippenko05}. It is located at $\\alpha~$ =  $23^h 21^m 03^s.38$  and $\\delta$  =-04$\\degr 53'  45''.5$ (equinox J2000),  which is $8.3''$  E and  $7.7''$ S  of the  centre of  the Sb spiral galaxy MCG-01-59-21 \\citep{vandenbergh05}.  \\begin{figure} \\psfig{figure=imagesn03jd_3.ps,width=8cm,height=8cm} \\caption{The field of SN 2003jd. $B$-band TNG image (14 Nov. 2003) 16.4~d after the $B$ maximum.} \\label{figseqstar} \\end{figure}  The SN was not visible on 16  Oct. in a KAIT unfiltered image (mag $<$ 19),  which sets a  tight limit  to the  explosion epoch  ($\\leq$ 13~d before  $B$  maximum). On  28  Oct., SN  2003jd  was  classified as  a peculiar Type Ic event with  very broad features analogous to those of SN 1998bw and SN 2002ap \\citep{filippenko03}.  The new interest in this kind of event prompted an intensive follow-up campaign  at  different  observing  sites,  lasting  for  about  three months. Few  more observations  in the late  nebular phase  were taken about one year later.  This  paper,  which  presents  and discusses  these  observations,  is organized as follows. In  \\S \\ref{parobs} we describe the observations and  data-reduction  techniques.   We  describe  the  photometric  and spectroscopic  data   of  SN  2003jd   in  \\S  \\ref{parphot}   and  \\S \\ref{parspe}.  In  \\S  \\ref{parintrinsic}  we  characterize  the  host galaxy,  and in  \\S \\ref{parcomparison}  we compare  SN 2003jd  with a sample  of  well-studied SNe~Ic:  1994I,  1998bw,  2002ap, 2004aw  and 2006aj. \\S \\ref{parametribolo} presents  the bolometric light curve of SN  2003jd,  computed  using   all  available  photometric  data.  The similarity   with   SN   2006aj/GRB   060218  is   discussed   in   \\S \\ref{parmissinggrb},  together  with a  discussion  of  the (lack  of) evidence  for an  associated  GRB. Finally,  in \\S  \\ref{parphyparam}, using simple  bolometric light-curve  modelling, we derive  some basic explosion parameters. ",
        "conclusions": "\\begin{tabular}{ll} \\hline Parent galaxy           &  MCG-01-59-021     \\\\ Galaxy type             &  Sb pec$^a$         \\\\ RA (2000)               &  $23^h 21^m 03^s.38$    \\\\ Dec (2000)              & $-04^\\circ53'45.5''$  \\\\ Recession velocity      [\\kms] & 5625$^b$ \\\\ Distance modulus ($H_0 = 72$) & $34.46 \\pm 0.20$ mag \\\\ $E(B-V)_{host}$         & $0.10^{+0.10}_{-0.05}$ mag$^c$\\\\ $E(B-V)_{Gal}$          & 0.044 mag$^d$ \\\\ Offset from nucleus     &  $8.3''$ E, $7.7''$ N   \\\\ \\hline \\end{tabular}\\\\ $^a$van den Bergh et al. 2005. \\\\ $^b$LEDA, corrected for Local Group infall (208 \\kms).\\\\ $^c$Calculated from the equivalent width of the \\NaID~ lines \\citep{turatto03}.\\\\ $^d$Schlegel et al. (1998). \\end{table}  The spectra of the host galaxy  taken at +788~d, long after the fading of the SN, can be  used to estimate the metallicity and star-formation rate  (SFR)   of  the  region  were   SN~2003jd  exploded.   Following \\cite{pettini04},  from  the [O\\,{\\sc  iii}]  and  [N\\,{\\sc ii}]  line strengths we derive an oxygen abundance  of $12 + {\\rm log}(O/H) = 8.4 \\pm 0.1$ dex. A slightly higher  value ($12 + {\\rm log}(O/H) = 8.7 \\pm 0.1$  dex) is  obtained using  the $R_{23}$  index (([O\\,{\\sc  iii}] + [O\\,{\\sc ii}])/H$\\beta$) of \\cite{pagel79}.  The two estimates bracket the measurement by \\cite{modjaz07};  they confirm that the metallicity of the SN~2003jd host is  below average for local galaxies and similar to that  of the SN~1998bw  host \\citep{modjaz07}.  On the  other hand, MCG-01-59-021 is a luminous galaxy,  $M_{B} = -20.3$ mag (LEDA), which is  somewhat at  odds with  the medium  to low  metallicity \\citep[see Fig.~5 of ][ and Fig.~1 of Prieto et al 2007]{modjaz07}.  The SFR at  the SN location can be estimated from  the \\Ha\\ flux, $7.1 \\pm   0.5~erg~s^{-1}~cm^{-2}$.    Following   \\cite{kennicutt98},   we estimate a SFR $\\approx  0.04$~M$_{\\odot}$ yr$^{-1}$, which is in fair agreement   with  the   value  derived   by  \\cite{modjaz07}   (SFR  = 0.07~M$_{\\odot}$ yr$^{-1}$).  This SFR  is typical for  \\HII~ regions, and similar  for instance  to that derived  for the environment  of SN 2006aj: SFR $\\approx 0.06$~M$_{\\odot}$ yr$^{-1}$ \\citep{pian06}.    We  have presented  optical photometry  and spectroscopy  of SN~2003jd spanning  from  3~d  before  $B$-band  maximum  to  $\\sim$400~d  after maximum. SN 2003jd shows the typical spectral features of a SN~Ic, but broadened as in the case of SNe 2002ap and 2006aj.  The  light curves  are similar  in shape  to those  of SNe  2002ap and 2006aj,  but the  peak  luminosity is  rather  similar to  that of  SN 1998bw.  The  comparison with  a  sample of  well-studied  SNe~Ic confirms  the heterogeneity  of this  SN  class. However,  the  application of  time stretching (as  for SNe~Ia) helps  in homogenizing this class  of SNe, reducing  the  differences among  the  light  curves and  facilitating spectral  comparisons.  Different  explosion  energies and  progenitor masses are likely  to be the main reasons  for the discrepancies.  The different viewing  angles from which  the SNe are observed  (given the evidence  for asymmetric  explosions)  may also  explain  some of  the differences, but this  effect is probably not large,  at least for the light-curve  shape  \\citep{maeda06b}.  It  may,  however,  affect  the spectra \\citep{tanaka07}.   The $E_{k}/M_{ej}$  ratio seems to  be the main factor in producing broad features.  The similarity to  SN 2006aj, the presence of  broad features, and the asymmetry shown by  the oxygen double peak in  the nebular spectra are in favour  of an asymmetric explosion  oriented away from  the line of site.   The  radio  observations  argue  that the  asymmetry  did  not extended to relativistic velocities, although uncertainties in the CSM mean that 2003jd could still be a mis-directed GRB.  Finally,  we used a  simple model  for the  bolometric light  curve to obtain the main  physical parameters of SN 2003jd.  The derived values confirm that SN 2003jd is a  broad-line SN similar to SN 2002ap and SN 2006aj,  but   with  a  larger   ejected  mass  ($M_{ej}  =   3.0  \\pm 1.0$~M$_\\odot$) and  kinetic energy ($E_{k}(tot)  = 7_{-2}^{+3} \\times E_{51}$ erg)  and producing a  large quantity of $^{56}Ni$  ($M_{Ni} = 0.36 \\pm  0.04$~M$_\\odot$).  Comparing  the physical parameters  of SN 2003jd with those of other SNe~Ic \\citep{nomoto07}, SN 2003jd appears to represent  a link between  broad-lined SNe (2002ap and  2006aj) and GRB-associated supernovae (SNe 1997ef, 1998bw, 2003dh, and 2003lw)."
    },
    "0710/0710.5345_arXiv.txt": {
        "abstract": "We extend and apply a model-independent analysis method developed earlier by Daly \\& Djorgovski to new samples of supernova standard candles, radio galaxy and cluster standard rulers, and use it to constrain physical properties of the dark energy as functions of redshift. Similar results are obtained for the radio galaxy and supernova data sets, which rely upon completely independent methods, suggesting that systematic errors are relatively small for both types of distances; distances to SZ clusters show a scatter which cannot be explained by the quoted measurement errors. The first and second derivatives of the distance are compared directly with predictions in a standard model based on General Relativity.  The good agreement indicates that General Relativity provides an accurate description of the data on look-back time scales of about ten billion years. The first and second derivatives are combined to obtain the acceleration parameter $q(z)$, assuming only the validity of the Robertson-Walker metric, independent of a theory of gravity and of the physical nature of the dark energy. The data are analyzed using a sliding window fit and using fits in independent redshift bins. The acceleration of the universe at the current epoch is indicated by the sliding window fit analysis. The effect of non-zero space curvature on $q(z)$ is explored; for a plausible range of values of $\\Omega_k$ the effect is small and causes a to shift to the redshift at which the universe transitions from deceleration to acceleration. We solve for the pressure, energy density, equation of state, and potential and kinetic energy of the dark energy as functions of redshift assuming that General Relativity is the correct theory of gravity.  Results obtained using a sliding window fit indicate that a cosmological constant in a spatially flat universe provides a good description of each of these quantities over the redshift range from zero to about one. We define a new function, the dark energy indicator, in terms of the first and second derivatives of the coordinate distance and show how this can be used to measure deviations of $w$ from $-1$ and to obtain a new and independent measure of $\\Omega_m$. ",
        "introduction": "Understanding of the physical nature of the dark energy which appears to be driving the accelerated expansion of the universe is among the most pressing and important topics in cosmology today. Studies of the expansion history of the universe allow us to constrain the physical nature of its matter and energy constituents. One way that the expansion and acceleration history of the universe can be studied is through the use of a set of coordinate distances and redshifts for some standard set of objects. Type Ia supernovae provide a modified standard candle (e.g. Phillips 1993, Hamuy et al. 1995) that allow the distance modulus, luminosity distance, and coordinate distance to each source to be determined. The recent data sets presented by Astier et al. (2006), Riess et al. (2007), Wood-Vasey et al. (2007), and Davis et al. (2007) have been analyzed by these groups and compared with numerous models by other researchers.  In a novel, largely model-independent approach to this problem, it was shown by Daly \\& Djorgovski (2003) that the first and second derivatives of the coordinate distance with respect to redshift could be obtained from the coordinate distances and combined to solve for the expansion rate $H(z)/H_0$ and acceleration rate $q(z)$ of the universe. The functions $H(y^{\\prime})$ and $q(y^{\\prime},y^{\\prime \\prime})$ are exact, that is, they are not obtained by expansions in terms of derivatives about some point. The only assumptions are that the universe is described by a Robertson-Walker metric and has zero space curvature. The results are independent of the contents of the universe and their physical properties, and even independent of whether General Relativity provides an accurate description of the universe. Here, we drop the assumption of zero space curvature; it turns out that the deceleration parameter at a redshift of zero, $q_0$, remains the same, independent of whether space curvature is zero or not.  In this paper we expand on the previous analysis done by Daly \\& Djorgovski (2003, 2004).  First, we use updated and expanded data sets, as described in section 2.1. Second, we introduce a more direct way to compare the model-independent results obtained from the data with predictions; this is done by directly comparing the first and second derivatives of the coordinate distance with respect to redshift to predicted values in various models, as described in section 2.2. Third, we analyze the data using both a sliding window fit and fits in independent redshift bins. To solve for the physical properties of the dark energy as functions of redshift, a theory of gravity must be specified. To determine the properties of the dark energy, General Relativity is taken to be the correct theory of gravity, allowing us to solve for the pressure, energy density, and equation of state of the dark energy as a function of redshift  in section 2.4.  Fourth, in section 2.4, we introduce a way to solve for the potential and kinetic energy densities of the dark energy as functions of redshift. In addition, we define a new function, the dark energy indicator, which provides a measure of deviations of $w$ from $-1$ and a new and independent measure of $\\Omega_m$. Fifth, in section 2.3, these derivatives are combined to solve for the expansion and acceleration rates of the universe as functions of redshift for both zero and non-zero space curvature; in our previous work we have not considered the effects of non-zero space curvature. The only assumption that must be made to obtain the functions $H(z)/H_0$ and $q(z)$ from the data are that the Robertson-Walker metric is valid in our universe. A discussion and conclusions follow in section 3. ",
        "conclusions": "The work presented here improves and extends our previous results. First, expanded and improved data sets are considered: three supernova samples and one radio galaxy sample. The radio galaxy data set has 11 new sources, increasing its size to 30 sources, and the supernovae data sets have increased substantially in size and quality.  In addition, SZ cluster distances and gamma-ray burst distances are considered.  The dimensionless coordinate distances (obtained directly from the data), and first and second derivatives of the distance are obtained as functions of redshift using a sliding window fit.  The good agreement obtained using supernovae and radio galaxies, two completely independent methods, with sources that cover similar redshift ranges, suggests that neither method is strongly affected by systematic effects, and that each method provides a reliable cosmological tool.  The first and second derivatives of the distance are combined to obtain the acceleration parameter $q(z)$, allowing for non-zero space curvature. It is shown that the zero redshift value of $q(z)$, $q_o$, is independent of space curvature, and can be obtained from the first and second derivatives of the coordinate distance.  Thus, $q_0$, which indicates whether the universe is accelerating at the current epoch, can be obtained directly from the supernova and radio galaxy data; our determinations of $q(z)$ only relies upon the validity of the Robertson-Walker line element, and is independent of a theory of gravity, and the contents of the universe. Each of the supernova samples, analyzed using a sliding window fit, indicate that the universe is accelerating today independent of space curvature, independent of whether General Relativity is the correct theory of gravity, and of the contents of the universe. The effect of non-zero space curvature on $q(z)$ is to shift the redshift at which the universe transitions from acceleration to deceleration, moving this to lower redshift for negative space curvature and to higher redshift for positive space. The zero redshift values of $q$ obtained using a sliding window fit. for the Davis et al. (2007) supernova sample is $q_0(192SN)= -0.48 \\pm 0.11$ and that obtained for the radio galaxy sample of Daly et al. (2007) is $q_0(30RG)=  -0.65 \\pm 0.53$ indicating that the universe is accelerating at the current epoch. The data were also binned so that only certain subsets of the data were used to solve for $y^{\\prime}$, $y^{\\prime \\prime}$, $H(z)/H_0$, and $q(z)$.  The results for $y^{\\prime}$ and $H(z)/H_0$ indicate that the standard LCDM model provides a good description of the data.  The results for $y^{\\prime \\prime}$ and $q(z)$ are consistent with the standard LCDM model, but do not independently confirm the model or the acceleration of the universe.  In addition to the evaluation of the standard cosmological parameters, in an even more direct approach, we compared $y^{\\prime}$ and $y^{\\prime \\prime}$ obtained from the fits to the data to model predictions. Comparisons of $y^{\\prime}$ and $y^{\\prime \\prime}$ with predictions based on General Relativity indicate that General Relativity provides an accurate description of the data on look-back time scales of about ten billion years, thus providing a very large scale test of General Relativity.  Another new approach is that the data were analyzed using both a sliding window fit and fits in independent redshift bins.  The fits in statistically independent redshift bins are broadly consistent with the sliding window fits, but are generally noisier (as expected).  We also explored the effects of non-zero space curvature on determinations of $H(z)$ and $q(z)$.  It is shown that the zero redshift value of $q$, obtained by applying equation (4) to $y^{\\prime}$ and $y^{\\prime \\prime}$, is independent of space curvature. This means that our method can be used to determine $q_0$, and thus the degree to which the universe is accelerating at the current epoch, with only one assumption, that the Robertson-Walker line element is valid. In addition, it is found that the effect of space curvature on the shape of $H(z)$ and $q(z)$ is small, relative to the uncertainties arising from the measurement errors.  After determining the expansion and acceleration rates of the universe as functions of redshift independent of a theory of gravity, we solve for the pressure, energy density, equation of state, and potential and kinetic energy of the dark energy as functions of redshift assuming that General Relativity is the correct theory of gravity.  We also define a new function, the dark energy indicator $s$, which provides a measure of deviations of the equation of state of the dark energy $w$ from $-1$, and provides a new and independent measure of $\\Omega_m$ if $w=-1$. The results obtained using a sliding window fit indicate that a cosmological constant in a spatially flat universe provides a good description of each of these quantities over the redshift range from zero to one. The zero redshift values of these quantities obtained with the Davis et al. (2007) supernovae sample are $P_{DE,0}/\\rho_{0c} = -0.64 \\pm 0.10$, $\\rho_{DE,0}/\\rho_{0c} = 0.67 \\pm 0.05$, $w_{DE,0} = -0.95 \\pm 0.08$, $V_{DE,0}/\\rho_{0c} = 0.65 \\pm 0.05$, $K_{DE,0}/\\rho_{0c} = 0.01 \\pm 0.03$, and $s_0 = -0.50 \\pm 0.08$. In the standard Lambda-Cold Dark Matter Model, $\\Omega_{\\Lambda} = - P_0/\\rho_{0c} = 0.64 \\pm 0.1$, obtained using the first and second derivatives of the coordinate distance, provides an independent measure of $\\Omega_{\\Lambda}$.  In addition, in this model, $w=-1$, so $s$ provides a measure of $\\Omega_m$, and the value obtained here using the first and second derivatives of the coordinate distance, is $\\Omega_m = 0.33 \\pm 0.05$. Overall, the shapes of the pressure, energy density, equation of state, and other parameters as functions are redshift are consistent with those predicted in a standard LCDM model.  There is a tantalizing hint that there may be divations from the standard model at high redshift; more observations at high redshift will be needed to investigate this further. The results obtained using fits in independent redshift bins are consistent with the standard LCDM model, but do not independently confirm the model."
    },
    "0710/0710.3755_arXiv.txt": {
        "abstract": "We argue that the Higgs boson of the Standard Model can lead to inflation and produce cosmological perturbations in accordance with observations.  An essential requirement is the non-minimal coupling of the Higgs scalar field to gravity; no new particle besides already present in the electroweak theory is required. ",
        "introduction": "The fact that our universe is almost flat, homogeneous and isotropic is often considered as a strong indication that the Standard Model (SM) of elementary particles is not complete.  Indeed, these puzzles, together with the problem of generation of (almost) scale invariant spectrum of perturbations, necessary for structure formation, are most elegantly solved by inflation \\cite{Starobinsky:1979ty,Starobinsky:1980te,Mukhanov:1981xt,Guth:1980zm,% Linde:1981mu,Albrecht:1982wi}.  The majority of present models of inflation require an introduction of an additional scalar---the ``inflaton''.  This hypothetical particle may appear in a natural or not so natural way in different extensions of the SM, involving Grand Unified Theories (GUTs), supersymmetry, string theory, extra dimensions, etc.  Inflaton properties are constrained by the observations of fluctuations of the Cosmic Microwave Background (CMB) and the matter distribution in the universe.  Though the mass and the interaction of the inflaton with matter fields are not fixed, the well known considerations prefer a heavy scalar field with a mass $\\sim \\unit[10^{13}]{GeV}$ and extremely small self-interacting quartic coupling constant $\\lambda \\sim 10^{-13}$ \\cite{Linde:1983gd}. This value of the mass is close to the GUT scale, which is often considered as an argument in favour of existence of new physics between the electroweak and Planck scales.  The aim of the present Letter is to demonstrate that the SM itself can give rise to inflation.  The spectral index and the amplitude of tensor perturbations can be predicted and be used to distinguish this possibility from other models for inflation; these parameters for the SM fall within the $1\\sigma$ confidence contours of the WMAP-3 observations \\cite{Spergel:2006hy}.  To explain our main idea, consider Lagrangian of the SM non-minimally coupled to gravity, \\begin{equation} \\label{main} L_{\\mathrm{tot}}= L_{\\mathrm{SM}} - \\frac{M^2}{2} R -\\xi H^\\dagger HR \\;, \\end{equation} where $L_{\\mathrm{SM}}$ is the SM part, $M$ is some mass parameter, $R$ is the scalar curvature, $H$ is the Higgs field, and $\\xi$ is an unknown constant to be fixed later.\\footnote{In our notations the conformal coupling is $\\xi=-1/6$.}  The third term in (\\ref{main}) is in fact required by the renormalization properties of the scalar field in a curved space-time background \\cite{Birrell:1982ix}.  If $\\xi=0$, the coupling of the Higgs field to gravity is said to be ``minimal''.  Then $M$ can be identified with Planck scale $M_P$ related to the Newton's constant as $M_P=(8\\pi G_N)^{-1/2}=\\unit[2.4\\times 10^{18}]{GeV}$.  This model has ``good'' particle physics phenomenology but gives ``bad'' inflation since the self-coupling of the Higgs field is too large and matter fluctuations are many orders of magnitude larger than those observed.  Another extreme is to put $M$ to zero and consider the ``induced'' gravity \\cite{Zee:1978wi,Smolin:1979uz,Spokoiny:1984bd,Fakir:1990iu,Salopek:1988qh}, in which the electroweak symmetry breaking generates the Planck mass \\cite{vanderBij:1993hx,CervantesCota:1995tz,Bij1995}. This happens if $\\sqrt{\\xi}\\sim1/(\\sqrt{G_N} M_W)\\sim10^{17}$, where $M_W\\sim\\unit[100]{GeV}$ is the electroweak scale.  This model may give ``good'' inflation \\cite{Spokoiny:1984bd,Fakir:1990iu,Salopek:1988qh,Kaiser:1994wj,% Kaiser:1994vs,Komatsu:1999mt} even if the scalar self-coupling is of the order of one, but most probably fails to describe particle physics experiments.  Indeed, the Higgs field in this case almost completely decouples from other fields of the SM\\footnote{This can be seen most easily by rewriting the Lagrangian (\\ref{main}), given in the Jordan frame, to the Einstein frame, see also below.} \\cite{vanderBij:1993hx,CervantesCota:1995tz,Bij1995}, which corresponds formally to the infinite Higgs mass $m_H$.  This is in conflict with the precision tests of the electroweak theory which tell that $m_H$ must be below $\\unit[285]{GeV}$ \\cite{:2005ema} or even \\unit[200]{GeV} \\cite{PDG2007} if less conservative point of view is taken.  These arguments indicate that there may exist some intermediate choice of $M$ and $\\xi$ which is ``good'' for particle physics and for inflation at the same time.  Indeed, if the parameter $\\xi$ is sufficiently small, $\\sqrt{\\xi} \\lll 10^{17}$, we are very far from the regime of induced gravity and the low energy limit of the theory (\\ref{main}) is just the SM with the usual Higgs boson.  At the same time, if $\\xi$ is sufficiently large, $\\xi \\gg 1$, the scalar field behaviour, relevant for chaotic inflation scenario \\cite{Linde:1983gd}, drastically changes, and successful inflation becomes possible.  We should note, that models of chaotic inflation with both nonzero $M$ and $\\xi$ were considered in literature \\cite{Spokoiny:1984bd,Futamase:1987ua,Salopek:1988qh,Fakir1990,Kaiser:1994vs,% Libanov1998,Komatsu:1999mt}, but in the context of either GUT or with an additional inflaton having nothing to do with the Higgs field of the Standard Model.  The Letter is organised as follows.  We start from discussion of inflation in the model, and use the slow-roll approximation to find the perturbation spectra parameters.  Then we will argue in Section \\ref{sec:radcorr} that quantum corrections are unlikely to spoil the classical analysis we used in Section \\ref{sec:cmb}.  We conclude in Section~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  In this Letter we argued that inflation can be a natural consequence of the Standard Model, rather than an indication of its weakness. The price to pay is very modest---a non-minimal coupling of the Higgs field to gravity.  An interesting consequence of this hypothesis is that the amplitude of scalar perturbations is proportional to the square of the Higgs mass (at fixed $\\xi$), revealing a non-trivial connection between electroweak symmetry breaking and the structure of the universe.  The specific prediction of the inflationary parameters (spectral index and tensor-to-scalar ratio) can distinguish it from other models (based, e.g.\\ on inflaton with quadratic potential), provided these parameters are determined with better accuracy.  The inflation mechanism we discussed has in fact a general character and can be used in many extensions of the SM. Thus, the $\\nu$MSM of \\cite{Asaka:2005an,Asaka:2005pn} (SM plus three light fermionic singlets) can explain simultaneously neutrino masses, dark matter, baryon asymmetry of the universe and inflation without introducing any additional particles (the $\\nu$MSM with the inflaton was considered in \\cite{Shaposhnikov:2006xi}). This provides an extra argument in favour of absence of a new energy scale between the electroweak and Planck scales, advocated in \\cite{Shaposhnikov:2007nj}."
    },
    "0710/0710.4519_arXiv.txt": {
        "abstract": "In the present paper we report on the difference in angular sizes between radio-loud and radio-quiet CMEs. For this purpose we compiled these two samples of events using Wind/WAVES and SOHO/LASCO observations obtained during 1996-2005. It is shown that the radio-loud CMEs are almost two times wider than the radio-quiet CMEs (considering expanding parts of CMEs). Furthermore we show that the radio-quiet CMEs have a narrow expanding bright part with a large extended diffusive structure. These results were obtained by measuring the CME widths in three different ways. ",
        "introduction": "The relation between coronal mass ejections (CMEs) and type II radio bursts has been studied for a long time, but is not fully understood (see Gopalswamy  (2006) for a recent review). Properties of the driving CMEs and the ambient medium through which the CMEs drive shocks show a large variability, which seems to contribute to the difficulties faced in understanding them (Gopalswamy {\\it et al.}, 2001). One of the major issues has been the lack of the type II radio emission in the metric (Sheeley {\\it et al.}, 1984) and decameter-hectometric (DH) wavelengths (Gopalswamy {\\it et al.}, 2001) even for CMEs with speeds exceeding 1000 km s$^{-1}$. Recently, Gopalswamy {\\it et al.} (2007) performed a systematic investigation of fast and wide (FW) CMEs that clearly lacked the metric and DH type II radio emission (``radio-quiet''CMEs) and compared them with the ones (``radio-loud'' CMEs) producing detectable radio type II. It was found that the radio-quiet CMEs can be distinguished from the radio-loud CMEs in three aspects: (1) speeds and widths, (2) a fraction of halo CMEs, and (3) solar source location of the CMEs. The radio-quiet CMEs are generally slower and narrower than the radio-loud ones. The fraction of halo CMEs is much larger for the radio-loud CMEs, which is related to the fact that the radio-quiet CMEs are narrower on the average. It is also known that halo CMEs are also faster and wider on the average (Yashiro \\emph{et al.}, 2004; Michalek, Gopalswamy, and Yashiro, 2003; Gopalswamy, 2004). When the source locations were examined, Gopalswamy (2006) and Gopalswamy \\emph{et al.} (2007) found that more than half of the radio-quiet CMEs were back-sided, while only a small fraction (25) of the radio-loud CMEs were back-sided. They attributed this result to the possibility that only a small fraction of the shock surface is visible to the observer, thereby reducing the possibility of detecting significant radio emission.  A fast but narrow CME may have a similar limitation because the CME cross-section and hence the shock surface area are expected to be smaller. One of the suggestions made in Gopalswamy \\emph{et al.} (2007) is that most of the radio-quiet CMEs may have a narrow bright part with extended diffuse structure. The purpose of this paper is to examine the evolution of the width of radio-quiet and radio-loud CMEs and compare them to confirm the smaller width of CMEs as a contributor to radio quietness.  In  Section~2 the procedure for obtaining  the samples of the radio-loud and radio-quiet CMEs is presented. In this section three different  methods for the determination of CME widths are also explained. In Section~3, we use the measured CME widths to show the spatial difference between the radio-loud and radio-quiet CME populations. ",
        "conclusions": "In this study we examined the spatial difference between the radio-loud and radio-quiet CMEs.  To get a reliable result, we determined the angular widths of the radio-loud and radio-quiet CMEs using three different methods. In all cases the radio-loud CMEs are wider than the radio-quiet CMEs. When we compare the expanding parts of CMEs (which are responsible for II type radio emission) the width difference between the events is the largest. The expanding structures of the radio-quiet CMEs are narrower ($\\approx$40$\\%$) in comparison with those of the radio-loud CMEs.   The expanding structures of CMEs are also much narrower in comparison with their total widths, especially for the radio-quite events. The ratios of the average catalog width to the average main body width are equal about 2.0 and 2.7 for the radio-loud and radio-quiet CMEs, respectively. The catalog widths for the radio-quiet CMEs are almost three times bigger in comparison with the widths of expanding structures. This means that the radio-quiet CMEs have a narrow expanding bright part with a large extended diffusive structure. It is clear that the spatial size of CME could be one of the most important factors defining the presence of type II radio emission. Our results proved the previous considerations ({\\it e.g.} Gopalswamy {\\it et al.}, 2001, 2005; Pick and Maia, 2005; Subramanian and Ebenezer, 2006). It is commonly accepted that type II radio bursts are radio signatures of coronal MHD-shock waves (Uchida, 1960; Wild, 1962). Flare-related blast waves and shock driven by CMEs have been considered as two possible pistons of metric type II bursts (see {\\it e.g.} Cliver, Webb, and Howard, 1999), while the DH and longer wavelength bursts due to CME-driven shock. CMEless type II bursts (Sheeley {\\it et al.}, 1984) and the discrepancy between the metric and IP type ( Reiner {\\it et al.}, 2001) were used to argue against the same shock causing the metric and IP type bursts. Gopalswamy (2006) demonstrated that both these discrepancies could be explained. It seems that CME-driven shock works for the entire interplanetary space but additional mechanism (blast waves) may operate for a narrow region ($\\approx 1R_{\\bigodot}$) close to the solar surface. In both cases, the width of CMEs plays an important role in generation of fast particles and radio bursts. Wider a given CME (more energetic event) wider a shock front and larger area where particles can gain energy. Additionally, larger CMEs could in bigger degree destruct  magnetic structures in corona and amplify radio emission ({\\it e.g.} Raymond {\\it et al.}, 2000; Pick {\\it et al.}, 2006). This scenario is confirmed by strong correlation between complex type III and type II radio bursts associated with CMEs (Cane, Erickson, and Prestage, 2002; Gopalswamy, 2004)."
    },
    "0710/0710.0801_arXiv.txt": {
        "abstract": "Clusters of galaxies are sites of acceleration of charged particles and sources of non-thermal radiation. We report on new constraints on the population of cosmic rays in the Intra Cluster Medium (ICM) obtained via radio observations of a fairly large sample of massive, X--ray luminous, galaxy clusters in the redshift interval 0.2--0.4. The bulk of the observed galaxy clusters does not show any hint of Mpc scale synchrotron radio emission at the cluster center (Radio Halo). We obtained solid upper limits to the diffuse radio emission and discuss their implications for the models for the origin of Radio Halos. Our measurements allow us to derive also a limit to the content of cosmic ray protons in the ICM. Assuming spectral indices of these protons $\\delta =2.1-2.4$ and $\\mu$G level magnetic fields, as from Rotation Measures, these limits are one order of magnitude deeper than present EGRET upper limits, while they are less stringent for steeper spectra and lower magnetic fields. ",
        "introduction": "Clusters of galaxies are ideal astrophysical environments for particle acceleration. Large scale shocks which form during the process of cluster formation are believed to be efficient particle accelerators (e.g. Sarazin 1999; Gabici \\& Blasi 2003; Ryu et al. 2003; Pfrommer et al. 2006). Cosmic rays (CRs) can also be injected into the ICM from ordinary galaxies and AGNs (e.g. V\\\"olk \\& Atoyan 1999) and turbulent eddies may contribute to the particle acceleration process (e.g. Brunetti \\& Lazarian 2007). CRs accelerated within the cluster volume would then be confined for cosmological times and the bulk of their energy is expected in protons since they have radiative and collisional life--times much longer than those of the electrons (e.g. Blasi et al. 2007, for a review).  While present gamma ray observations can only provide upper limits to the average energy density of CR protons in the ICM (e.g. Reimer et al. 2004), the presence of relativistic electrons in a number of clusters has been ascertained via the detection of a tenuous synchrotron radio emission: giant Radio Halos (RHs) and mini-Radio Halos, fairly symmetric sources at the cluster center, and Radio Relics, elongated sources at the cluster periphery (e.g. Feretti \\& Giovannini 2007). It is customary to classify the models for the origin of RHs in {\\it secondary electron} (e.g. Blasi \\& Colafrancesco 1999) and {\\it reacceleration models} (e.g. Brunetti et al. 2001; Petrosian 2001), depending on whether the radiating electrons are produced as secondary products of hadronic interactions or reaccelerated by turbulence from a pre-existing population of non-thermal seeds in the ICM, respectively. These models predict a different connection between radio and X-ray properties of clusters which are discussed in this Letter and compared with new observations: in Sect.2 we review the expectations of the different models, in Sect.3 we briefly present the radio observations of our cluster sample, and in Sects.4 \\& 5 we report and discuss our results. Concordance ($H_o$=70, $\\Omega_m$=0.3, $\\Omega_{\\Lambda}$=0.7) cosmology is used. ",
        "conclusions": "We have reported on constraints on the origin of RHs and on the CR content in the ICM obtained via radio observations of a fairly large sample of X--ray luminous clusters at $z = 0.2-0.4$.  \\noindent In the bulk of these clusters we do not find evidence of Mpc--scale radio emission at the level of RHs. Our conclusions become even more stringent considering radio emission on cluster--core scale, typical of smaller RHs and mini-Halos.  \\noindent We firmly confirm that RH--clusters follow a {\\it physical} correlation between synchrotron and X--ray luminosities. We find that clusters have a bimodal distribution in the $P_{1.4}$--$L_X$ plane (Fig.~4); this is in line with the expectation of the {\\it re-acceleration scenario}. On the other hand, in order to reconcile these observations with expectations from {\\it secondary} models strong dissipation of the magnetic field in the clusters with no radio emission is necessary.  \\noindent Our measurements allow us to also derive simple limits on the presence of CR protons in the ICM (Fig.~5). In the case of relatively flat spectral energy distribution of these CRs stringent upper limits can be obtained: the energy density of CRs should be $\\leq 1$\\% of the thermal energy in case of $\\geq \\mu$G field strength. This would make problematic the detection of gamma rays from $\\pi^o$--decay in clusters with GLAST. On the other hand, by assuming steeper spectral energy distributions of these CRs (or lower magnetic fields) our limits become less stringent."
    },
    "0710/0710.5938_arXiv.txt": {
        "abstract": "We present IRAC and MIPS images and photometry of a sample of previously known planetary nebulae (PNe) from the SAGE survey of the Large Magellanic Cloud (LMC) performed with the Spitzer Space Telescope.  Of the 233 known PNe in the survey field, 185 objects were detected in at least two of the IRAC bands, and 161 detected in the MIPS 24 $\\mu$m images.  Color-color and color-magnitude diagrams are presented using several combinations of IRAC, MIPS, and 2MASS magnitudes.  The location of an individual PN in the color-color diagrams is seen to depend on the relative contributions of the spectral components which include molecular hydrogen, polycyclic aromatic hydrocarbons (PAHs), infrared forbidden line emission from the ionized gas, warm dust continuum, and emission directly from the central star. The sample of LMC PNe is compared to a number of Galactic PNe and found to not significantly differ in their position in color-color space. We also explore the potential value of IR PNe luminosity functions (LFs) in the LMC. IRAC LFs appear to follow the same functional form as the well-established [\\ion{O}{3}] LFs although there are several PNe with observed IR magnitudes brighter than the cut-offs in these LFs. ",
        "introduction": "The Large Magellanic Cloud (LMC) has been important for the study of many astrophysical processes and objects because it is one of the nearest galaxies to our own, and due to its location above the Galactic plane and its favorable viewing angle \\citep[35$\\degr$;][]{vanderm01}, the system can be relatively easily surveyed and many of its global properties determined.  These properties are important in particular for the study of planetary nebulae (PNe).  The known distance to the LMC removes the relatively large uncertainty in this parameter that affects many Galactic PNe \\citep{hajian06}.  The distance of $\\sim$ 50 kpc allows individual objects to be isolated and in some cases resolved.  The effects on PNe of the lower metallicity and dust/gas mass ratio in the LMC can be explored. One can also hope to detect a large fraction of the total number of PNe in the LMC, as opposed to in the Galaxy, where confusion and extinction in the plane allow us to detect only about 10\\% of the PNe expected to exist \\citep{kwok00,frew05}.  An infrared survey of the LMC called Surveying the Agents of a Galaxy's Evolution \\citep[][SAGE]{meixner06} has recently been completed using the IRAC \\citep{fazio04} and MIPS \\citep{rieke04} instruments on the Spitzer Space Telescope \\citep{werner04}.  SAGE is an unbiased, magnitude-limited survey of a $\\sim 7\\degr \\times 7\\degr$ region centered on the LMC.  This Spitzer ``Legacy'' survey has provided a tremendous resource for the study of the stellar populations and interstellar medium (ISM) in the LMC. Some early results on the evolved stellar populations were given by \\citet{blum06}, who identified $\\sim$32,000 evolved stars brighter than the red giant tip, including oxygen-rich, carbon-rich, and ``extreme'' asymptotic giant branch (AGB) stars.  In this paper we explore the properties of a sample of known PNe as revealed by the SAGE data.  The catalog of 277 LMC PNe assembled by \\citet{leisy97} from surveys that cover an area of over 100 square degrees was used for the source of positions of the PNe.  Leisy et al. used CCD images and scanned optical plates to obtain accurate positions of the objects to better than 0\\farcs5.  They point out that the objects are in general PN candidates, with only 139 confirmed at that time with slit spectroscopy.  For simplicity we will refer to the objects in the catalog as PNe, even though this caveat still applies for many of the sources. When we began to work with the SAGE data, the Leisy et al. catalog was the largest summary list of the known PNe at the time.  During the course of this work, \\citet{reid06} published a list of PNe in the central 25 deg$^2$ of the LMC, including 169 of the previously known objects and 460 new possible, likely, or true PNe. We will present our results here for the \\citet{leisy97} catalog, and a future paper will include the new objects in the \\citet{reid06} survey. ",
        "conclusions": "We have presented images and photometry of the \\citet{leisy97} sample of PNe in the LMC.  Of the 233 known PNe in the survey field, 185 objects were detected in at least two of the IRAC bands, and 161 detected in the MIPS 24 $\\mu$m images.  Color-color and color-magnitude diagrams were presented using several combinations of IRAC, MIPS, and 2MASS magnitudes. The location of an individual PN in the color-color diagrams was seen to depend on the relative contributions of the spectral components, resulting in a wide range of colors for the objects in the sample.  A comparison to a sample of Galactic PNe shows that they do not substantially differ in their position in color-color space.  The location of PNe in the various infrared color-color and color-magnitude diagrams are in general well separated from normal stars, but overlap significantly with extragalactic sources and potential YSOs.  Any ambiguity between PNe and YSOs or galaxies can be readily resolved by the unique optical characteristics of PNe and their environs.  Therefore, an IR color-based search for new PNe in the LMC would be viable in combination with deep optical imaging and spectroscopy. The latter remains the prerequisite to confirm a candidate as a PN.  We have offered an exploration of the potential value of IR PNLFs in the LMC. IRAC LFs appear to follow the same functional form as the well-established [\\ion{O}{3}] LFs although there are several PNe with observed IR magnitudes brighter than the cut-offs in these LFs.  If these objects can be demonstrated to be true PNe and not very-low excitation variants nor symbiotic stars then their existence may confirm the long-standing suggestion that PNe with massive central stars suffer heavy internal extinction.  This extinction would eliminate optical outliers beyond the cut-off magnitude but would affect IR LF counts minimally so that all such outliers could be observed."
    },
    "0710/0710.0206_arXiv.txt": {
        "abstract": "The commonality of collisionally replenished debris around main sequence stars suggests that minor bodies are frequent around Sun--like stars. Whether or not debris disks in general are accompanied by planets is yet unknown, but debris disks with large inner cavities -- perhaps dynamically cleared -- are considered to be prime candidates for hosting large--separation massive giant planets. We present here a high--contrast VLT/NACO angular differential imaging survey for eight such cold debris disks. We investigated  the presence of  massive giant planets in the range of orbital radii  where the inner edge of the dust debris is expected. Our observations are sensitive to planets and brown dwarfs with masses  $>$3 to 7 Jupiter mass, depending on the age and distance of the target star. Our observations did not identify any planet candidates. {\\changed We compare the derived planet mass upper limits to the minimum planet mass required to dynamically clear the inner disks. While we cannot exclude that single giant planets are responsible for clearing out the inner debris disks,  our observations constrain the parameter space available for such planets.}  The non--detection of massive planets in these evacuated debris disks further reinforces the notion that the giant planet population is confined to the inner disk ($<$15~AU). ",
        "introduction": "Collisionally replenished debris dust surrounds about 10--20\\% of the main sequence Sun--like stars (e.g. \\citealt{Meyer2007}). Such widespread evidence for minor body collisions demonstrates that planetesimals orbit most stars. It is natural to ask whether or not rocky and giant planets  are also present in these systems. No convincing correlation could yet be found between close--in exoplanets and the presence of debris (e.g. \\citealt{Amaya2007}, but  see \\citealt{2005ApJ...622.1160B}). However, the presence of massive giant planets has been often invoked to account for the observed azimuthal or radial asymmetries at large radii in many debris disks (e.g. \\citealt{Greaves2005,Wilner2002}).  While theory offers several alternative mechanisms (e.g. \\citealt{Takeuchi2001,Wyatt2005}), dynamical clearing of dust parent bodies by giant planets remains a feasible and exciting theoretical possibility (e.g. \\citealt{Amaya2005,Quillen2006,Levison2007, Morbidelli2007}).  Examples for such possibly dynamically--cleared disks include two recently identified disks around the young Sun--like stars \\object{HD 105}  \\citep{Meyer2004}  and \\object{HD 107146} \\citep{Williams2004}. Both disks were found to exhibit strong excess emission at wavelengths longer than 30~$\\mu$m, while displaying no measurable excesses shortward of 20~$\\mu$m. The detailed analysis of the spectral energy distribution of  \\object{HD 105} suggests that it is consistent with a narrow dust ring ($<$4~AU) with an inner radius of $\\sim$42~AU, if the dust grains emit like black bodies \\citep{Meyer2004}. Using a similar model \\citet{Williams2004} showed that the excess emission from \\object{HD 107146} is consistent with arising from cold dust (T=51~K) emitting as a single--temperature black body. The lack of measurable infrared excess shortward of 25~\\micron ~illustrates that the inner disk regions are well cleared of dust: for HD~107146 there is at most 140$\\times$ less warm dust (T=100~K) than cold  dust (T=51~K, \\citealt{Williams2004}). The findings of the spectral energy distribution model for \\object{HD 107146} have been confirmed by direct imaging with the Hubble Space Telescope, that strengthen the case for a large featureless dust ring outside of an evacuated inner cavity \\citep{Ardila2004}.  Recent high--contrast imaging surveys have hinted on the general scarcity of giant planets at such large separations (e.g. \\citealt{Masciadri2005,Kasper2007,Biller2007,Lafreniere2007}). Quantitative statistical analysis of the non--detections demonstrates that -- at a 90\\% confidence level -- the giant planet population cannot extend beyond 30~AU if it follows a $r^{0.2}$ radial distribution, consistent with the radial velocity surveys. The statistical analysis suggests an outer cut--off  for the giant planet population at $<$15~AU \\citep{Kasper2007}. If so, dynamically cleared cold debris disks may be the ssignposts for rare large--separation giant planets, ideally suited for direct imaging studies.  In this paper we report on a VLT/NACO high--contrast imaging survey for large--separation giant planets around  HD~105, HD~107146, and six other similar disks. In the following we will review the target stars and disks, the observations, followed by a comparison of our non--detections to lower planet mass limits set by dynamical clearing simulations.  \\subsection{Targets} Our targets were selected from the sample of 328 Sun--like stars (0.7--2.2 $M_\\odot$) targeted in  the {\\em Formation and Evolution of Planetary Systems}  Spitzer Space Telescope Legacy program (FEPS, \\citealt{Meyer2006}). From this sample we identified \\ntargets southern stars, which: a) display strong infrared excess emission at long wavelengths ($\\lambda > 20 \\mu$m); b) no measurable excess emission at shorter wavelengths; and, C) are young  and close enough to permit the detection of planetary--mass objects within the inner radius of the cold debris. Table~\\ref{T:Targets} gives an overview of the key parameters of the target stars. The typical lower mass limit for the debris in the systems is $10^{-4}$ to $10^{-5}$~M$_{\\earth}$, making these disks massive analogs of our Kuiper--belt (\\citealt{Meyer2004,Kim2005}, Hillenbrand et al., in prep). {\\changed The disks of \\object{HD 105} and  \\object{HD 107146} --- included in our sample --- have  inner evacuated regions with an estimated radii of $\\sim40$~AU and $\\sim31$~AU  \\citep{Meyer2004,Williams2004}. The other six disks exhibit spectral energy distributions similar to \\object{HD 105} and \\object{HD 107146}. Based on the similarity of the excess emissions and the almost identical spectral types all eight disks are expected to have cleared--out inner disks of similar size.} The only possible exception in this sample is \\object{HD 202917}, for which the re--calibration of the IRAC fluxes after our VLT observations revealed a faint, but likely real infrared excess even at wavelengths shortward of 10\\micron, suggesting that this inner disk may harbor small, but non--negligible amounts of warm dust.  In the following we discuss briefly the results of the age determination for these sources as this has direct impact on the sensitivity of our observations to giant planets. A more detailed discussion of the ages of the whole FEPS sample will be presented in Hillenbrand et al. (in prep). We briefly summarize the upper and lower age estimates ($t_{min}$ and $t_{max}$) for each star along  with the most likely age $t_{prob}$, where available. HD~105 has already reached the main sequence ($t_{min}$=27~Myr) and its chromospheric activity suggests a $t_{max}$ of 225~Myr (Hillenbrand et al., in prep.). Very likely a member of the Tuc--Hor moving group \\citep{Mamajek2004} its $t_{prob}$ is 30~Myr \\citep{Hollenbach2005}. HD~377 is also a main sequence star ($t_{min}>25$~Myr) and the chromospheric activity suggests that $t_{max}$=220~Myr. The median of four other age indicators sets $t_{prob}$=90~Myr. For HD~107146  we adopt  the age range of 80--200~Myr. HD~202917 is a likely member of the Tuc--Hor moving group ($t_{min}=t_{prob}=30$~Myr) and its upper age limit is set by its Li--abundance, higher than that of the Pleiades ($t_{max}<100$~Myr). HD~35850 is suggested to be a $\\beta$~Pic Moving Group member ($t_{min}$=12~Myr, \\citealt{Song2003}) and its observed rotation rate sets a reliable upper age limit of $t_{max}$=100~Myr (Hillenbrand et al. in prep.; cf. \\citealt{Barnes2007}). HD~70573 is among the few stars that are known to harbor both a debris disk and a giant planet. \\citet{Setiawan2007} found an $m_2 sin i = 6.1$~\\mjup possible planet on a 1.76--AU orbit. A combination of different age indicators suggest a $t_{min}=30$~Myr for HD~70573 and a $t_{prob}=60~$Myr; \\citet{Setiawan2007} quotes $t_{max}$=125~Myr. {\\changed Based on Li--abundance and chromospheric activity, position on the color--magnitude diagram and the analysis of its space motions Mamajek et al. (in prep.)  estimates that HD~209253 has $t_{min}=200$~Myr and $t_{max}<1.6$~Gyr  with $t_{prob}$=500~Myr.} HD 25457 is a member of the AB~Dor moving group giving a very strong lower age limit ($t_{min}=50$~Myr, \\citealt{Zuckermanetal2004}). \\citet{Luhman2005} derives an age of 75--125~Myr (we adopt $t_{prob}=75$~Myr), while the upper age limit is set by the chromospheric activity ($t_{max}=170$~Myr).   \\begin{deluxetable}{lccccccc} \\tabletypesize{\\scriptsize} \\tablecaption{Target parameters. \\label{T:Targets}} \\tablewidth{0pt} \\tablehead{\\colhead{Target} &  \\colhead{R.\\,A. (J2000)} & \\colhead{Dec. (J2000)}  &  \\colhead{V--mag.$^a$}& \\colhead{Dist. [pc]$^a$} &\\colhead{Sp. Type} & \\colhead{Ages:$^b$ $t_{min}$/$t_{prob}$/$t_{max}$}  } \\startdata HD 105    & 00 05 52.6 & $-$41 45 11 & 7.51 & 40 & G0V    & 27 Myr / 30 Myr / 225 Myr\\\\ HD 377    & 00 08 25.7 & $+$06 37 01 & 7.59 & 40 & G2V    & 25~Myr / 90~Myr / 220~Myr  \\\\ HD 25457  & 04 02 36.8 & $-$00 16 08 & 5.38 & 19 & F5V    & 50 Myr / 75 Myr / 170 Myr\\\\ HD 35850  & 05 27 04.8 & $-$11 54 03 & 6.30 & 27 & F7/8V  &  12 Myr /  12 Myr / 100 Myr\\\\ HD 70573  & 08 22 50.0 & $+$01 51 34 & 8.69 & 70 & G1/2V  & 30 Myr / 60 Myr / 125 Myr \\\\ HD 107146 & 12 19 06.5 & $+$16 32 54 & 7.04 & 29 & G2V    &  80 Myr /  -- / 200 Myr \\\\ HD 202917 & 21 20 50.0 & $-$53 02 03 & 8.65 & 46 & G5V    &   30 Myr / 30 Myr / 100 Myr \\\\ HD 209253 & 22 02 33.0 & $-$32 08 02 & 6.63 & 30 & F6/7V  &  200~Myr / 500 Myr / 1.6 Gyr \\\\ \\enddata \\tablenotetext{a}{All magnitudes and distances from the Hipparcos catalog, except for the distance of HD~70573, which is a main sequence--distance. } \\tablenotetext{b}{The age estimates are discussed in the text.} \\end{deluxetable} ",
        "conclusions": "We present results from a high--contrast angular differential imaging survey of  \\ntargets cold debris disks, selected to have significantly or totally evacuated inner disks. Our observations searched for massive giant planets that may be  responsible for carving out the inner holes in the  observed cold debris disks. For most of our targets we reach typical sensitivities of  3 to 7 \\mjup between 20 to 50~AU separations, but did not identify any likely planet or brown dwarf candidates.  {\\changed By comparing the derived planet mass upper limits to lower limits derived from dynamical scattering models (typically 2-5 \\mjup between 10 and 30 AU), we limit the parameter space available for any single planet capable of efficiently clearing out the inner planetesimal disks. }  Our survey complements recent direct imaging surveys of nearby young stars indicating that massive giant planets at large separations are very rare. Cool debris disks with large inner evacuated cavities remained promising possible exceptions to this rule until now. However, the combination of  our observational upper limits and theoretical lower limits strongly suggest that massive giant planets at large separations are not present in most of these systems, reinforcing the finding that the outer cut--off for the giant planet distribution is probably at 15~AU or at even smaller semi--major axes \\citep{Kasper2007}."
    },
    "0710/0710.1941_arXiv.txt": {
        "abstract": "{Measurements at 100 TeV and above are an important goal for the next generation of high energy $\\gamma$-ray astronomy experiments to solve the still open problem of the origin of galactic cosmic rays. The most natural experimental solution to detect very low radiation fluxes is provided by the Extensive Air Shower (EAS) arrays. They benefit from a close to 90$\\%$ duty cycle and a very large field of view ($\\sim$2 sr), but the sensitivity is limited by their angular resolution and their poor cosmic ray background discrimination. Above 10 TeV the standard technique for rejecting the hadronic background consists in looking for {\\it \"muon-poor\"} showers.  In this paper we discuss the capability of a large muon detector (A$_{\\mu}$=2500 m$^2$) operated with an EAS array at very high altitude ($>$4000 m a.s.l.) to detect $\\gamma$-ray fluxes around 100 TeV. Simulation-based estimates of energy ranges and sensitivities are presented. }  \\begin{document} ",
        "introduction": "The recent TeV results of the HESS experiment suggest the existence of a population of galactic $\\gamma$-ray sources whose emission extends beyond 10 TeV in the 5 to 15$\\%$ of Crab flux range (for E $>$ 1 TeV). These sources, associated with nearby shell-type or plerionic SNRs, the most probable factories of galactic cosmic rays, can be studied detecting gamma-rays (and neutrinos) emission in the VHE/UHE energy domain. Therefore, a detector capable to perform a continuous all-sky survey at a level of about a percent of the Crab flux up to 100 TeV is needed. The search and study of {\\em \"Cosmic PeVatrons\"} and their surrounding regions is one of the main scientific issues to be addressed by the next generation of ground-based $\\gamma$-ray astronomy detectors \\cite{aharonian}.  Current experiments are not able to reach 100 TeV because their limited collection area makes the required exposures too long. Extrapolating the galactic source spectra measured by HESS, it appears that future experiments need to achieve at least 100 km$^2\\cdot$h in order to obtain meaningful measurements at 100 TeV. The most natural experimental solution that provides such a large exposure is given by the EAS arrays observing each source for $\\sim$1500 h/year. As an example, an experiment like ARGO-YBJ already results in an exposure of $\\sim$15 km$^2\\cdot$h/year, with an angular resolution ($\\sim$0.2$^{\\circ}$ at 10 TeV) near the best value attainable by a sampling array. In addition, the EAS arrays are the only ground-based detectors allowing simultaneous and continuous coverage of a significant fraction of the sky (about all that overhead). Their large field of view and high duty cycle ($>$90 $\\%$) suit to perform a $\\gamma$-ray sources population survey at VHE/UHE energies. But more important is their unique potential that allows to have an effective monitoring of the $\\gamma$-ray activity of a large number of highly variable sources like blazars and microquasars, as well as the possibility of independent detection and study of GRBs \\cite{aharonian}. In addition, the recent observations of unidentified extended sources from the Galactic plane \\cite{milagro_galpl} and in the Cygnus region \\cite{milagro_cygnus} reported by the Milagro Collaboration demonstrate the strength of EAS arrays in finding diffuse and extended sources. Therefore, the discovery science could be a feature of EAS arrays.  The limited sensitivity in detecting $\\gamma$-ray point sources, characteristic of EAS experiments, is mainly due to their poor gamma/hadron separation power, limited angular resolution and high energy threshold. Exploiting a full coverage approach at very high altitude leads to the improvement of the angular resolution to $\\sim$0.2$^{\\circ}$ and to the reduction of the energy threshold well below the TeV region. The standard technique to perform a gamma/hadron discrimination above 10 TeV with EAS arrays consists in looking for {\\it \"muon-poor\"} showers.  In this paper we discuss the capability of a large muon detector (A$_{\\mu}$=2500 m$^2$) operated with an EAS array at very high altitude ($>$4000 m a.s.l.) to detect $\\gamma$-ray fluxes up to 100 TeV. An estimation based on the simulation of energy ranges and sensitivities is reported. ",
        "conclusions": ""
    },
    "0710/0710.1310_arXiv.txt": {
        "abstract": "This Letter reports on initial Expanded Very Large Array (EVLA) observations of the 6035~MHz masers in ON~1.  The EVLA data are of good quality, lending confidence in the new receiver system.  Nineteen maser features, including six Zeeman pairs, are detected.  The overall distribution of 6035~MHz OH masers is similar to that of the 1665~MHz OH masers.  The spatial resolution is sufficient to unambiguously determine that the magnetic field is strong ($\\sim -10$~mG) at the location of the blueshifted masers in the north, consistent with Zeeman splitting detected in 13441~MHz OH masers in the same velocity range.  Left and right circularly polarized ground-state features dominate in different regions in the north of the source, which may be due to a combination of magnetic field and velocity gradients.  The combined distribution of all OH masers toward the south is suggestive of a shock structure of the sort previously seen in W3(OH). ",
        "introduction": "\\object[Onsala 1]{Onsala 1} (ON 1) is a massive star-forming region with an unusual OH maser spectrum.  The ground-state masers, which have been observed interferometrically several times, \\citep[e.g.,][]{ho83,argon00,fish05,nammahachak06,fish07}, appear in two disjoint velocity ranges: $< 6$~km\\,s$^{-1}$ and 11--17~km\\,s$^{-1}$, with no maser emission in between.  This pattern is reproduced in 6668~MHz methanol emission, which also appears near 0 and 15~km\\,s$^{-1}$ \\citep{szymczak00}.  The masers presently elude clear interpretation.  Comparison of OH maser velocities seen in projection against the ultracompact \\ion{H}{2} region \\citep[13--14~km\\,s$^{-1}$ after correcting for Zeeman splitting; e.g., from][]{fish05} with the LSR velocity of the latter derived from a hydrogen recombination line ($5.1 \\pm 2.5$~km\\,s$^{-1}$) led \\citet{zheng85} to interpret the masers as tracing infall.  However, a proper motion study of the OH masers suggests that expansion may dominate the kinematics \\citep{fish07}.  A more recent model suggests that the OH masers are associated with a molecular outflow \\citep{kumar04,nammahachak06}.  This model proposes a shocked molecular torus origin for the southern masers, but questions remain regarding the overall morphology and velocity structure of the masers in ON~1.  Few northern star-forming regions have been mapped in the 6035~MHz line of OH, in part because few radio arrays in the northern hemisphere have been capable of tuning to the frequency.  A previous three-station European VLBI Network (EVN) experiment detected seven 6035~MHz maser features in ON~1 but only observed the redshifted masers \\citep{desmurs98}.  Single-dish observations confirm that the 6030 and 6035~MHz masers in ON~1 also appear in two disjoint velocity ranges, but Zeeman pairing ambiguities have prevented a definitive measurement of the magnetic field of the blueshifted masers \\citep{baudry97,fish06}.  The upgrade of the National Radio Astronomy Observatory's (NRAO) Very Large Array (VLA) to the Expanded VLA (EVLA) presents new observational opportunities in North America \\citep{mckinnon01,ulvestad07}.  Of interest to spectral line observers is the full frequency coverage between 1 and 50~GHz that will become available.  This spring, the EVLA for the first time offered observational capabilities in the extended C-band range of 4.2 to 7.7~GHz, which includes key maser frequencies of OH and methanol. This Letter reports on initial observations of 6035~MHz OH masers with the EVLA. ",
        "conclusions": "\\subsection{Environment of the Masers in ON~1}  The distributions of the 1665 and 6035~MHz masers are the most alike of any pair of OH transitions in ON~1.  All the 1665~MHz maser regions have associated 6035~MHz masers, apart from in the center and in the extreme south, far from the exciting source where densities are likely to be low.  The brightest masers at both transitions occur in the line of maser spots just south of center in Figure~\\ref{map}.  This contrasts with the 1612, 1667, and 1720~MHz masers, which are only found in the prominent line of masers south of and on the southwestern limb of the continuum source.  A similar result was seen at VLBI resolution in W3(OH): 1665 and 6035~MHz masers appear throughout the source, 1612 and 1720~MHz masers appear only in and near the inner edge of an apparent shocked torus that is especially well traced by 6.0~GHz masers (including several of the brightest 6.0~GHz masers detected), and 1667~MHz emission is largely concentrated in areas with an apparent shock morphology \\citep{fishevn07}.  Higher-resolution observations of the 6.0~GHz masers in ON~1 would be useful to obtain a better alignment of the 6035~MHz frame with respect to the other masers and determine whether the southern masers are located preferentially along the northern edge of the southern shock front.  Some of the 6035~MHz masers may be associated with an outflow, such as the H$^{13}$CO$^{+}$ oriented northeast-southwest through ON~1 detected by \\citet{kumar04}.  \\citet{nammahachak06} note that the southern line of OH masers is oriented nearly perpendicular to this structure and propose, based additionally on the linear polarization characteristics of the ground-state masers, that the masers trace a shock in a confining molecular torus.  Similar conclusions based on maser distribution and polarization are reached for other OH maser sources hosting outflows \\citep{hutawarakorn99,hutawarakorn03,hutawarakorn05,hutawarakorn02}. This interpretation also bears a striking similarity to that proposed for W3(OH), which hosts very strong 6035~MHz masers \\citep{fishevn07}, and may therefore suggest a common mechanism for producing strong 6035~MHz maser emission.  It is probable that there is a coherent, organized structure in the northern part of ON~1.  Many LCP 1665~MHz masers, including one bright one, occur in the region, with a few weak RCP masers offset to the north and west of the LCP emission \\citep{fish05,fish07}.  Assuming a magnetic field of approximately $-10$~mG in the region, LCP and RCP 1665~MHz maser components would be separated by 6~km\\,s$^{-1}$ ($20\\, \\Delta v$ for an average line width of 0.3~km\\,s$^{-1}$) if they both appear in the region.  (It is likely that weak RCP 1665~MHz masers, seen near $-2$~km\\,s$^{-1}$ in the spectra of \\citet{clegg91}, pair with the dominant LCP masers near +4~km\\,s$^{-1}$, but no published interferometric observations of the ground-state masers have included the $-2$~km\\,s$^{-1}$ features in the observed bandpass.)  If the magnetic field and velocity gradients are aligned such that the magnetic field strength decreases (becomes less negative) in the same direction that the velocity increases, amplification of the LCP component would be favored \\citep[see][]{cook66}.  This effect is much more pronounced for ground-state OH masers than at 6035~MHz, where the Zeeman splitting is a factor of 10 smaller in velocity units, such that the LCP and RCP spectra are separated by $0.6$~km\\,s$^{-1} \\approx 2\\, \\Delta v$ per 10~mG, which is not in excess of the turbulent velocity component expected in a maser condensation \\citep{reid80}.  Hence, 6035~MHz masers appear in both polarizations with similar fluxes, while strong LCP 1665~MHz emission appears at higher velocities (and possibly weaker RCP emission at lower velocities).  The key question that remains is how the blueshifted masers in the north connect to the redshifted masers in the south.  It is tempting to associate both with the HCO$^+$ outflow, but the northern masers are significantly blueshifted compared to the HCO$^+$ velocity range of 8--16~km\\,s$^{-1}$ \\citep{kumar04}.  These authors also note a CO outflow in the region that spans a larger velocity range, but it is oriented roughly east-west.  \\citet{kumar04} interpret the outflows as coming from two different sources embedded in the ultracompact \\ion{H}{2} region.  It is possible that the blue and red OH masers are also associated with two different sources, which would complicate interpretation of their motions and morphology.  The large ($|B| > 10$~mG) magnetic fields suggest that the northern masers are near a region of higher density and therefore likely an excitation source. In W3(OH) the methanol masers in the region of highest magnetic field strength ($|B| > 10$~mG) have been modelled as undergoing conical expansion, which \\citet{moscadelli02} interpret as possibly being due to an outflow guided by the helical field from a magnetized disk.  \\subsection{Future Directions}  Excluding minor transitional issues, the performance of the EVLA antennas was as expected.  Early data from the antennas with upgraded C-band capability are of good quality and are already producing useful science \\citep[e.g.,][]{sjouwerman07}.  As the upgrade continues, more antennas with 6.0~GHz tuning capability will be added to the array, providing propotionally better sensitivity and imaging characteristics.  High spectral resolution VLBI observations of the 1665 and 6035~MHz OH masers, of the sort obtained for W3(OH) \\citep{fbs06,fishevn07}, may help answer the questions of whether there is an organized velocity structure in the north and what structure the masers are tracing.  If the segregated regions of LCP and RCP emission in the north are due to correlated magnetic and velocity field gradients, small-scale velocity gradients (i.e., on the size scale of a maser spot) of the northern masers may show similar magnitudes and position angles.  This would contrast with W3(OH), in which small-scale velocity gradients show no correlation except when masers overlap \\citep{fbs06,fishevn07}.  Any such observations of the ground-state masers should have a sufficiently wide bandwidth coverage to include the $-2$~km\\,s$^{-1}$ 1665~MHz masers in order to be able to identify magnetic field strengths (and variations thereof) throughout the northern masers. While observations of linear polarization exist for the ground-state masers \\citep{nammahachak06}, full-polarization observations at 6035~MHz would help to understand the magnetic field geometry, since linear polarization vectors are much less subject to being corrupted by external and internal Faraday rotation at the higher frequency \\citep[e.g.,][]{fishreid06}."
    },
    "0710/0710.1917_arXiv.txt": {
        "abstract": "{The standard cooling flow model has predicted a large amount of a cool gas in the clusters of galaxies. The failure of the Chandra and XXM-Newton telescopes to detect a cooling gas (below 1-2 keV) in the clusters of galaxies has suggested that some heating process must work to suppress the cooling. The most likely heating source is the heating by AGNs. There are many heating mechanisms, but we will adopt the effervescent heating model which is a result of the interaction of the bubbles inflated by AGN with the intra-cluster medium(ICM).  Using the FLASH code, we have carried out  1D- time dependent simulations to investigate the effect of the heating on the suppression of the cooling in cooling flow clusters. We have found that the effervescent heating model can not balance the radiative cooling and it is an artificial model. Furthermore, the effervescent heating  is  a function of the ICM pressure gradient but the cooling  is proportional  to the gas density square and square root of the gas temperature.} {}{}{}{} ",
        "introduction": "% According to the steady flow assumption of the standard cooling flow model, we must find a cool gas and a multi-phase medium within the cooling core in clusters of galaxies which are not observed in any wavebands  (X-ray and non X-ray). This is known as the cooling flow problem in clusters of galaxies. In others words, there is  a discrepancy between standard cooling flow model and observations (X-ray and non X-ray). This strong  discrepancy is  interpreted as either the gas is being prevented from cooling by some heating process, or  it cools without any spectroscopic signature~\\citep{fabian01} which is difficult.  But, we see that this discrepancy between the standard cooling flow model and X-ray observations indicates that either  the cooling in the center of cooling flow clusters must be suppressed by any heating mechanism, or  the steady flow assumption of the standard cooling flow model is not appropriate. Somewhere else, we have concentrated in the second point which it is found that the steady flow with cooling is impossible, i.e the cooling flow problem is due to the wrong steady flow assumption. In this work we will concentrate in the first point, heating the cooling flow.\\\\ The failure of the multi-phase model has revived the idea of a heating mechanism which can suppress the cooling. There are five main conditions for the heating \\citep[see][ for review]{gardini04}~: \\begin{enumerate} \\item  The heating must be fine tuned and distributed to get the smooth  observed temperature profile. Too much heating would result a outflow from the center region. Too little heating is not sufficient to suppress the cooling. Moreover, the heating process must be self regulated: the mass flow rate triggers the heating and the heating reduces the mass flow rate~\\citep{hans02}. That mechanism is called a heating with a feedback. \\item The heating mechanism by AGN must be sporadic because the radio activities are not observed in every cooling flow clusters. \\item  The kinetic energy injection must be subsonic and the ICM must not be shocked, as observed in general case. The shock could compress the gas producing the catastrophic cooling much faster. \\item  The heating mechanism must preserve the observed entropy profiles of the ICM,  which decrease toward the center. \\item  The heating must not destroy the metallicity profiles  which are peaked toward the centers \\citep{tamura02,grandi01,irwin01}. \\end{enumerate}  A giant elliptical or cD galaxy sits at the potential well of every cooling flow cluster~\\citep{mathews03,eilek04}. The popular heating mechanisms are a heating by AGN~\\citep{binney95,binney93,br02}, thermal conduction, cosmic rays, galaxies motions, magnetic field reconnection \\citep{soker90} and turbulent mixing \\citep{kim03}. About 71 percent of the cDs in the cooling flow clusters are radio loud compared to only 23 percent of non-cooling flow clusters ~\\citep{burns90}. This result suggests that there a relationship  between the AGN  activities and the presence of cooling flow.  In some of cooling flow  clusters, the recent X-ray observations reveal  holes in the X-rays surface brightness coincident with the radio lobes, known as X-ray cavities or bubbles~\\citep[ see][ as example]{fabian06}. The  bubbles or cavities are not a universal phenomenon, suggesting a duty cycle. {}~Radio sources are not even detected in some cooling-flow clusters.  The observed metallicity gradients make a constraint on the ability of baubles to mix the ICM~\\citep{bohringer04}. ~\\citet{bm102} have run many 1D simulations of clusters with various levels of heating. They concluded that the best fit of temperature profiles with the real clusters is without heating. The AGNs are assumed to inject buoyant bubbles into the ICM, which heat the ambient medium by doing work $PdV$ as they rise and expand. This mechanism is called the  effervescent heating model or mechanism~\\citep{begelman01}. In this work, we will concentrate on this mechanism, using time dependent hydrodynamics simulations, FLASH code~\\citep{fry02}. \\Rem{ ",
        "conclusions": "\\label{sec:sm_heat_model}                   % We have carried out time dependent simulations with cooling and heating. The general result is that the best fit to the observations is the model without heating (model A). The effervescent heating  is  a function of the ICM pressure gradient but the cooling  is proportional to the gas density square and square root of the gas temperature.~Furthermore, the conclusions are~: \\\\ \\\\ 1- From our simulations, the cooling can not be a steady flow as assumed by standard cooling flow model. The gas density must increase with the time; i.e the ICM must be compressed under the force of the inflowing gas.\\\\ \\\\ 2- The effervescent heating model profile (without the smoothness function) is more steeper than the radiative cooling profile, which makes it is very difficult to balance the cooling. \\\\ \\\\ 3-  The function $1-e^{-r_o/r}$ is not a cutoff function but it is a smoothing function in order to control the steepness of the heating function letting it ~balance the radiative cooling; i.e. it is an artificial model.  \\\\ \\\\ 4- At inner radius close to the realistic value of  the accretion radius (model D), the accretion mass rate due to flowing gas is not enough to fuel the black hole in the center of cluster. It's not reasonable to set the accretion radius of the black hole at large radius, (for example 1-1.5 kpc, as used for the three dimension simulations). \\\\ \\\\ 5- In some cooling flow clusters, the observed cooling time scale is very short, in the range $10^{8}$ yr to $10^{9}$ yr, but the cooling gas below 1-2 keV is not observed. This result is interpreted  as indication that there must be a heating mechanism to stop the cooling. As in figure~\\ref{fig:a1991_purecooling_tmscl}, the cooling time scale inside a radius of 10 kpc is shorter than $1$ Gy but there is no a cool gas. Moreover, from our simulations (model B),  we found that the heating process increases the cooling time scale. If cooling time scale is a good approximation for  the actual cooling time and the cluster is heated, then we must find that the cooling time scale is larger than observed.\\\\ \\\\"
    },
    "0710/0710.3586_arXiv.txt": {
        "abstract": "We investigate large-amplitude baryon acoustic oscillations (BAO's) in off-diagonal entries of cosmological power-spectrum covariance matrices.  These covariance-matrix BAO's describe the increased attenuation of power-spectrum BAO's caused by upward fluctuations in large-scale power.  We derive an analytic approximation to covariance-matrix entries in the BAO regime, and check the analytical predictions using $N$-body simulations.  These BAO's look much stronger than the BAO's in the power spectrum, but seem detectable only at about a one-sigma level in gigaparsec-scale galaxy surveys. In estimating cosmological parameters using matter or galaxy power spectra, including the covariance-matrix BAO's can have a several-percent effect on error-bar widths for some parameters directly related to the BAO's, such as the baryon fraction.  Also, we find that including the numerous galaxies in small haloes in a survey can reduce error bars in these cosmological parameters more than the simple reduction in shot noise might suggest. ",
        "introduction": "Humans are fortunate that there are baryons in the Universe.  Not only are we made of them, but baryons are responsible for features in the shape of cosmological power spectra and two-point correlation functions that are quite valuable to cosmologists.  These features are called baryon acoustic oscillations (BAO's), and are imprints of acoustic oscillations in the gas of the early universe \\citep{peeyu, sz, holtz, ehu, mwp}.  Their presence in galaxy clustering statistics provides a standard ruler to measure the relation between distance and redshift, and the expansion of the universe at late times \\citep{bg, se}.  They provide one of the main tools currently proposed for studying the effects of dark energy.  BAO's have been detected in modern low-redshift galaxy surveys, both in the 2dFGRS \\citep{c05} and SDSS \\citep{e05,hutsi,pshape07}.  These detections were made using the (two-point) correlation function, or its Fourier dual, the power spectrum.  Conveniently, the BAO regime is on large-enough scales that non-linear effects are mild.  Still, for precision cosmology, these mild effects must be understood, and are the topic of much recent research \\citep[e.g.][]{mill,jk,huff,sw,aea,se07,sssbao}.  Non-linear evolution of the matter power spectrum tends to dampen or smear BAO's.  For example, \\citet{esw} found that large-scale bulk flows and cluster formation produce motions that smear out the BAO peak in the correlation function, but that these motions are confined to relatively small scales of $\\sim 10 \\hMpc$ in Lagrangian space. They argued that these motions roughly preserve wiggles on the largest scales of the power spectrum, but wipe out wiggles on smaller scales. The effects we describe in this paper, using Eulerian perturbation theory, likely arise physically from the same large-scale bulk motions. The attenuation of BAO's can also be understood by considering an additive mode-coupling power spectrum, which rises on small scales as structure develops, along with a function which attenuates the linear power spectrum on non-linear scales.  In the halo model \\citep[HM, reviewed in][]{cs}, this small-scale contribution is the one-halo (1h) term, and comes from pairs of objects within single haloes.  A qualitatively similar effect occurs in renormalized perturbation theory \\citep[RPT;][]{csrpt,mrpt}, which is less empirical than the HM, and seems more accurate through translinear scales.  For example, \\citet{csrpt} show that, more physically than in the HM, the mode-coupling power spectrum in RPT goes to zero on small scales.  In this paper, we show that wiggles exist in off-diagonal terms of matter and galaxy power spectrum covariance matrices, almost entirely out-of-phase with BAO's in the power spectrum.  We interpret these BAO's in the covariance matrix as manifestations of the suppression exacted on power-spectrum BAO's by power on large scales.  Regions of the Universe with upward fluctuations in large-scale power have more-suppressed BAO's.  The body of the paper is organized as follows.  In Section \\ref{anal}, we discuss analytic predictions for the BAO's in the covariance matrix of matter.  In Section \\ref{nbody}, we test the analytic predictions against $N$-body simulations, and investigate the detectability of the wiggles in the covariance matrix.  Finally, in Section \\ref{cosmomatter}, we investigate the effect of these covariance-matrix BAO's on cosmological parameter estimation, focusing on parameters directly related to BAO's in the power spectrum.  We perform this analysis for both matter and galaxy power spectra, using the HM framework. ",
        "conclusions": "Our main points are the following:  \\begin{itemize} \\item In off-diagonal entries in power spectrum covariance matrices, BAO's exist which appear much stronger than the BAO's in the power spectrum.  These wiggles are a manifestation of the suppression which large-scale power does to BAO's in the power spectrum, and originate in the perturbation-theory trispectrum.  We give a simple analytic approximation to these wiggles in terms of the linear power spectrum and its first two derivatives, and check the analytical predictions using $N$-body simulations.  \\item These wiggles are potentially detectable in current and upcoming surveys, but because of the large noise in a covariance matrix measurement, they are only detectable at a small significance level. We estimate that a one-sigma detection could be done with a survey of size a couple of $h^{-3}{\\,}{\\rm Gpc}^3$, and that a three-sigma detection could require a survey of volume $\\sim30\\hGpcV$.  \\item The wiggles make a modest difference in cosmological parameter error bars from analysing galaxy and matter power spectra.  For example, using the true, wiggly covariance matrix in estimating the baryon fraction results in error bars several percent tighter than a no-wiggle covariance matrix.  Doing so in estimating the baryon acoustic scale has a smaller effect.  \\item At fixed volume, including galaxies in smaller-mass haloes provides tighter error bars on parameters such as the baryon fraction and the baryon acoustic scale than analysing galaxies in only large haloes.  In the context of the HM, this effect goes beyond the simple gains from analysing a sample with smaller shot noise.  We attribute this to the one-halo term of the galaxy power spectrum becoming dominant at smaller scales for small haloes than large ones. \\end{itemize}  \\begin{sloppypar} Most of the calculations in this paper made use of our package of Python code for cosmology, called {\\scshape CosmoPy}. It can be downloaded from \\url{http://www.ifa.hawaii.edu/cosmopy/}. \\end{sloppypar}"
    },
    "0710/0710.3854_arXiv.txt": {
        "abstract": "We investigate the evolution of the properties of model populations of ultraluminous X-ray sources (ULXs), consisting of a black-hole accretor in a binary with a donor star. We have computed models corresponding to three different populations of black-hole binaries, motivated by our previous studies. Two of the models invoke stellar-mass ($\\sim10\\,M_\\odot$) black-hole binaries, generated with a binary population synthesis code, and the third model utilizes intermediate-mass ($\\sim1000\\,M_\\odot$) black-hole accretors (IMBHs). For each of the three populations, we computed 30,000 binary evolution sequences using a full Henyey stellar evolution code. A scheme for calculating the optical flux from ULXs by including the reprocessed X-ray irradiation by, and the intrinsic viscous energy generation in, the accretion disk, as well as the optical flux from the donor star, is discussed. We present color-magnitude diagrams (CMDs) as ``probability images'' for the binaries as well as for the donor stars alone. ``Probability images'' in the plane of orbital period and X-ray luminosity are also computed.  We show how a population of luminous X-ray sources in a cluster of stars evolves with time. The most probable ULX system parameters correspond to high-mass donors (of initial mass $\\gtrsim 25~M_\\odot$) with effective O through late B spectral types and equivalent luminosity classes of IV or V. We also find the most probable orbital periods of these systems to lie between 1-10 days.  Estimates of the numbers of ULXs in a typical galaxy as a function of X-ray luminosity are also presented.  From these studies we conclude that if the stellar-mass black-hole binaries are allowed to have super-Eddington limited X-ray luminosities: (i) the value of the binding energy parameter for the stellar envelope of the progenitor to the black-hole accretor must be in the range of $0.01 \\lesssim \\lambda \\lesssim 0.03$ in order not to overproduce the ULXs, and (ii) the stellar-mass black-hole models still have a moderately difficult time explaining the observed ULX positions in the CMD.  Other possible explanations for the apparent overproduction of very luminous X-ray sources in the case of stellar-mass black-hole accretors are discussed. Our model CMDs are compared with six ULX counterparts that have been discussed in the literature. The observed systems seem more closely related to model systems with very high-mass donors in binaries with IMBH accretors.  We find that a significant contribution to the optical flux from the IMBH systems comes from {\\em intrinsic} accretion disk radiation whose source is viscous dissipation of gravitational potential energy.  In effect, the IMBH systems, when operating at their maximum luminosities ($10^{41}-10^{42}$ ergs s$^{-1}$), are {\\em milli-AGN}.  With regard to the IMBH scenario, while attractive from the aspect of binary evolution models, it leaves open the larger question of how the IMBHs form, and how they capture massive stellar companions into just the correct orbits. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.0585_arXiv.txt": {
        "abstract": "We describe efforts over the last six years to implement regularization methods suitable for studying one or more interacting black holes by direct N-body simulations. Three different methods have been adapted to large-N systems: (i) Time-Transformed Leapfrog, (ii) Wheel-Spoke, and (iii) Algorithmic Regularization. These methods have been tried out with some success on GRAPE-type computers. Special emphasis has also been devoted to including post-Newtonian terms, with application to moderately massive black holes in stellar clusters. Some examples of simulations leading to coalescence by gravitational radiation will be presented to illustrate the practical usefulness of such methods. ",
        "introduction": "In the study of strong gravitational interactions, utilization of the chain data structure can be very beneficial. Over many years, the original chain regularization method (Mikkola \\& Aarseth 1993) has proved to be effective in star cluster simulations containing binaries. As we shall see in the following, it is also a useful tool in connection with time transformations which do not employ the usual coordinates. By introducing one or more dominant masses these advantages become more apparent. Such problems fall naturally into three classes according to the number of massive objects and each class requires special attention. At the simplest level we have the case of one central massive body which dominates the motion of other members within a certain distance. The role of the reference body is readily seen in the case of three interacting particles which can be studied by three-body regularization (Aarseth \\& Zare 1974). This idea was extended to an arbitrary membership (Zare 1974). However, a natural application was lacking until the problem of black holes (BHs) became a challenge for simulators in recent years. The aptly named wheel-spoke regularization (Aarseth 2003a) has now been adapted to study compact subsystems containing a single massive object.  Historically speaking, a special method for a BH binary was implemented in an $N$-body code first. Here the main idea is based on a time-transformed leapfrog scheme (TTL) suitable for dealing with large mass ratios (Mikkola \\& Aarseth 2002). Remarkably, this method yields machine precision for unperturbed two-body motion. Although regularized chain coordinates are not employed directly, the accuracy is improved by using relative quantities with respect to the nearest massive body.  Alternative methods may be needed for problems involving more than two massive objects. The recent algorithmic regularization (Mikkola \\& Tanikawa 1999, Preto \\& Tremaine 1999, Mikkola \\& Merritt 2006) appears to be a promising way of studying such systems. Indeed, the masses only play a kinematical role in one formulation, suggesting that it may be applicable to systems with large mass ratios. However, it should be emphasized that for practical reasons any of the above methods are of necessity limited to relatively small particle numbers, or in other words, compact subsystems. ",
        "conclusions": ""
    },
    "0710/0710.2313_arXiv.txt": {
        "abstract": "In differentially rotating discs with no self-gravity, density waves cannot propagate around the corotation, where the wave pattern rotation speed equals the fluid rotation rate. Waves incident upon the corotation barrier may be super-reflected (commonly referred to as corotation amplifier), but the reflection can be strongly affected by wave absorptions at the corotation resonance/singularity. The sign of the absorption is related to the Rossby wave zone very near the corotation radius. We derive the explicit expressions for the complex reflection and transmission coefficients, taking into account wave absorption at the corotation resonance. We show that for generic discs, this absorption plays a much more important role than wave transmission across the corotation barrier. Depending on the sign of the gradient of the vortensity of the disc, $\\zeta=\\kappa^2/(2\\Omega\\Sigma)$ (where $\\Omega$ is the rotation rate, $\\kappa$ is the epicyclic frequency, and $\\Sigma$ is the surface density), the corotation resonance can either enhance or diminish the super-reflectivity, and this can be understood in terms of the location of the Rossby wave zone relative to the corotation radius. Our results provide the explicit conditions (in terms of disc thickness, rotation profile and vortensity gradient) for which super-reflection can be achieved. Global overstable disc modes may be possible for discs with super-reflection at the corotation barrier. ",
        "introduction": "Differentially rotating fluid discs, ubiquitous in astrophysics, are known to exhibit rich dynamics and possible instabilities (e.g. Papaloizou \\& Lin 1995; Balbus \\& Hawley 1998). While local instabilities, such as Rayleigh's centrifugal instability (for discs with specific angular momentum decreasing outwards), gravitational instability (for self-gravitational discs with too large a surface density, or more precisely, Toomre $Q\\lo 1$), and magnetorotational instability (for discs with a sub-thermal magnetic field), are well understood (at least in the linear regime), global effects and instabilities are more subtle, since they involve couplings and feedbacks of fluid at different locations (see Goldreich 1988 for an introduction/review). A well-known example is the corotation amplifier (e.g. Mark 1976; Narayan, Goldreich \\& Goodman 1987), which arises from the interaction across the corotation between waves carrying opposite signs of angular momentum. Much stronger corotation amplifications (WASER -- wave amplification by the stimulated emission of radiation, and SWING amplifiers) can be achieved for self-gravitating discs (e.g., Goldreich \\& Lynden-Bell 1965; Julian \\& Toomre 1966; Lin \\& Lau 1975; see Shu 1992 for a review). Another well-known example is the Papaloizou-Pringle instability in finite accretion tori (confined between two free surfaces), in which coupling between waves inside the corotation with those outside, combined with reflecting inner and outer boundaries, leads to violent overstable modes (Papaloizou \\& Pringle 1984; Goldreich et al.~1986). Recent works on global disc instabilities include the Rossby wave instability (for discs with a strong enough density or vortensity bump; Lovelace et al.~1999; Li et al.~2000) and the accretion-ejection instability (for magnetized discs; Tagger \\& Pellat 1999, Tagger \\& Varniere 2006).  In this paper we are interested in 2D fluid discs without self-gravity and magnetic field. For disturbances of the form $e^{im\\phi-i\\omega t}$, where $m>0$ and $\\omega$ is the wave (angular) frequency (and thus $\\omega_p=\\omega/m$ is the pattern frequency), the well-known WKB dispersion relation for density waves takes the form (e.g., Shu 1992) \\be (\\omega-m\\Omega)^2=\\tomega^2=\\kappa^2+k_r^2c^2, \\label{eq:disp} \\ee where $\\Omega$ is the disc rotation frequency, $\\tomega=\\omega-m\\Omega$ is the Doppler-shifted wave frequency, $\\kappa$ is the radial epicyclic frequency, $k_r$ is the radial wavenumber, and $c$ is the sound speed. Thus waves can propagate either inside the inner Lindblad resonance radius $r_{\\rm IL}$ (defined by $\\tomega=-\\kappa$) or outside the outer Lindblad resonance radius $r_{\\rm OL}$ (defined by $\\tomega=\\kappa$), while the region around the corotation radius $r_c$ (set by $\\tomega=0$) between $r_{\\rm IL}$ and $r_{\\rm OL}$ is evanescent. Since the wave inside $r_{\\rm IL}$ has pattern speed $\\omega_p$ smaller than the fluid rotation rate $\\Omega$, it carries negative wave action (or angular momentum), while the wave outside $r_{\\rm OL}$ carries positive wave action. As a result, a wave incident from small radii toward the corotation barrier will be super-reflected, (with the reflected wave having a larger amplitude than the incident wave amplitude) if it can excite a wave on the other side of the corotation barrier. If there exists a reflecting boundary at the inner disc radius $r_{\\rm in}$, then a global overstable mode partially trapped between $r_{\\rm in}$ and $r_{\\rm IL}$ will result (see, e.g. Narayan et al.~1987 for specific examples in the shearing sheet model, and Goodman \\& Evans,1999 and Shu et al.~ 2000 for global mode analysis of singular isothermal discs).  The simple dispersion relation (\\ref{eq:disp}), however, does not capture an important effect in the disc, i.e., corotation resonance or corotation singularity. Near corotation $|\\tomega|\\ll\\kappa$, the WKB dispersion relation for the wave is [see equation (\\ref{eq:disper}) below] \\be \\tomega={2\\Omega k_\\phi\\over k_r^2+k_\\phi^2 +\\kappa^2/c^2}\\left({d\\over dr}\\ln{\\kappa^2 \\over 2\\Omega\\Sigma}\\right)_c, \\label{eq:disp2}\\ee where $k_\\phi=m/r$ and $\\Sigma$ is the surface density, and the subscript ``c'' implies that the quantity is evaluated at $r=r_c$. The quantity \\be \\zeta\\equiv {\\kappa^2\\over 2\\Omega\\Sigma}= {(\\nabla\\times {\\bf u_0})\\cdot {\\hat z}\\over\\Sigma} \\label{eq:zeta}\\ee is the vortensity of the (unperturbed) flow (where ${\\bf u_0}$ is the flow velocity). The dispersion relation (\\ref{eq:disp2}) describes Rossby waves, analogous to those studied in geophysics (e.g. Pedlosky 1987) \\footnote{A Rossby wave propagating in the Earth's atmosphere satisfies the dispersion relation $\\tomega=(2k_\\phi/k^2R)(\\partial\\Omega_3/ \\partial\\theta)$, where $k^2=k_\\phi^2+k_\\theta^2$, $\\Omega_3=\\Omega\\cos\\theta$ is the projection of the rotation rate on the local surface normal vector and $\\theta$ is the polar angle (co-latitude).}. For $k_r^2\\gg \\kappa^2/c^2$ and $k_r^2\\gg k_\\phi^2$, we see that Rossby waves can propagate either outside the rotation radius $r_c$ (when $d\\zeta/dr>0$) or inside $r_c$ (when $d\\zeta/dr<0$). In either case, we have $k_r\\rightarrow\\infty$ as $r\\rightarrow r_c$. This infinite wavenumber signifies wave absorption (cf. Lynden-Bell \\& Kalnajs 1972 in stellar dynamical context; Goldreich \\& Tremaine 1979 in the context of wave excitation in discs by a external periodic force; see also Kato 2003, Li et. al. 2003, and Zhang \\& Lai 2006 for wave absorption at the corotation in 3D discs). At corotation, the wave pattern angular speed $\\omega/m$ matches $\\Omega$, and there can be efficient energy transfer between the wave and the background flow, analogous to Landau damping in plasma physics. Narayan et al.~(1987) treated this effect as a perturbation of the shearing sheet model, and showed that the corotational absorption can convert neutral modes in a finite shearing sheet into growing or decaying modes. Papaloizou \\& Pringle (1987) used a WKB method to examine the effect of wave absorption at corotation on the nonaxisymmetric modes in an unbound (with the outer boundary extending to infinity) cylindrical torus.  In this paper, we derive explicit expressions for the complex reflection coefficient and transmission coefficient for waves incident upon the corotation barrier. We pay particular attention to the behavior of perturbations near the corotation resonance/singularity. Our general expressions include both the effects of corotation amplifier and wave absorption at corotation (which depends on $d\\zeta/dr$). We show explicitly that depending on the sign of $d\\zeta/dr$, the corotation resonance/singularity can either enhance or diminish the super-reflectivity, and this can be understood in terms of the location of the Rossby wave zone relative to the corotation radius.  Our paper is organized as follows. After presenting the general perturbation equations (section 2), we discuss the the wave dispersion relation and propagation diagram, and derive the local solutions for the wave equation around the Lindblad resonances and corotation resonance (section 3). We then construct global WKB solution for the wave equation, and derive the wave reflection, transmission and corotational damping coefficients in section 4. An alternative derivation of the wave damping coefficient is presented in section 5. Readers not interested in technical details can skip Sections 2-5 and concentrate on Section 6, where we illustrate our results and discuss their physical interpretations. Section 6.1 contains a numerical calculation of the wave reflectivity across corotation and discusses the limitation of the WKB analysis. We discuss how global overstable modes may arise when super-reflection at the corotation is present in section 7 and conclude in section 8. ",
        "conclusions": "In this paper we have derived explicit expressions for the reflection coefficient, transmission coefficient and wave absorption coefficient when a wave is scattered by the corotation barrier in a disc. These expressions include both the effects of corotation amplifier (which exists regardless of the gradient of the vortensity $\\zeta=\\kappa^2/\\Omega\\Sigma$ of the background flow) and wave absorption at the corotation (which depends on $d\\zeta/dr$). They demonstrate clearly that the corotation wave absorption plays a dominant role in determining the reflectivity and that the sign of $d\\zeta/dr$ determines whether the corotation singularity enhances or diminishes the super-reflectivity. Our result can be understood in terms of the location of the Rossby wave zone relative to the corotation radius. We also carried out numerical calculations of the reflectivity. Our result provides the conditions (in terms of disc thickness, rotation profile and surface density profile) for which super-reflection is achieved and global overstable modes in discs are possible.  In future works we will explore global oscillation modes and their stabilities in a variety of astrophysical contexts, ranging from accreting white dwarfs to accreting black hole systems. The possible overstabilities of these modes are directly linked to the effects studied in this paper and may provide explanations for some of the quasi-periodic variabilities observed in these systems."
    },
    "0710/0710.4123_arXiv.txt": {
        "abstract": "We present new estimates of ages and metallicities, based on FORS/VLT optical (4400--5500\\,\\AA) spectroscopy, of 16 dwarf elliptical galaxies (dE's) in the Fornax Cluster and in Southern Groups. These dE's are more metal-rich and younger than previous estimates based on narrow-band photometry and low-resolution spectroscopy. For our sample we find a mean metallicity ${\\rm [Z/H]} = -0.33$\\,dex and mean age 3.5\\,Gyr, consistent with similar samples of dE's in other environments (Local Group, Virgo). Three dE's in our sample show emission lines and very young ages. This suggests that some dE's formed stars until a very recent epoch and were self-enriched by a long star formation history. Previous observations of large near-infrared ($\\sim 8500$\\,\\AA) \\ionCaT\\ absorption strengths in these dE's are in good agreement with the new metallicity estimates, solving part of the so-called Calcium puzzle. ",
        "introduction": "The observed strength of the near-infrared \\ionCaT\\ triplet absorption lines in early-type galaxies has presented astronomers with an interesting puzzle over the past couple of years. \\citet{cenarro01} defined a new CaT* index, carefully correcting for the underlying H Paschen absorption. Whereas other metallicity tracers, such as Mg$_2$, correlate with velocity dispersion $\\sigma$, it was found that CaT* \\emph{anti-correlates} with $\\sigma$ in elliptical galaxies (E's) and in bulges of spiral galaxies \\citep{saglia02, cenarro03, falcon03}. Population synthesis model predictions also show that, for sub-solar metallicities, CaT* should be sensitive to metallicity but virtually independent of age, while at super-solar metallicity, the CaT* saturates \\citep{vazdekis03}. However, taking metallicities estimated from optical spectra, \\citet{saglia02} reported that the measured CaT* values in E's are smaller by 0.5\\,{\\AA} than those predicted by population synthesis models. \\citet{michielsen03} (hereafter Paper~1) showed that the anti-correlation of CaT* with $\\sigma$ extends into the dwarf elliptical (dE) regime.  These dE's were expected to have metallicities of the order of [Z/H]$\\,\\sim -1$, and ages of the order of 10\\,Gyr \\citep{heldmould94, rakos01}. The measured CaT* values were significantly larger than those expected for such old, metal-poor stellar systems.  All of the proposed solutions to this conundrum, such as variations of the initial-mass function or the calcium yield as a function of metallicity or velocity dispersion, require considerable fine-tuning. None of them satisfactorily explains both the small CaT* in E's and the large CaT* in dE's without creating other difficulties, e.g.\\ with the FeH~$\\lambda$9916 index values observed in bright ellipticals (FeH is strong in dwarf stars but nearly absent in giants) and with stellar mass-to-light ratios \\citep{saglia02, cenarro03}. However, the stellar populations of E's and dE's are most likely not single-age, single-metallicity populations, or SSPs, as was implicity assumed in essentially all age and metallicity estimates \\citep[see e.g.][]{pasquali05}. Moreover, the stars that dominate the blue spectral range (mostly hot dwarf stars) do not necessarily have the same mean ages/metallicities as the stars producing the red light (mostly cool giants). These issues, together with systematic uncertainties inherent to population synthesis tools, potentially contribute to the CaT puzzle.  Because of their low surface brightness, accurate estimates of the ages and metallicities of dE's are still scarce. Recent studies of dE's in the Virgo cluster \\citep{g03, v04} report younger ages and higher metallicities than found in Fornax dE's. In Paper~1, the ages and metallicities of the dE's were taken from the literature, where low-resolution, modest S/N spectroscopic \\citep{heldmould94} or narrow-band photometric \\citep{rakos01} techniques were used. As a sanity check, we have now acquired high-resolution, high S/N optical spectra with FORS/VLT of all the dE's for which we presented CaT* measurements in Paper~1.  This puts us in a position where we can for the first time compare the CaT* measurements with model predictions based on robust age and metallicity estimates. ",
        "conclusions": "\\label{disc}  At least in the dE regime, the \\ionCaT\\ Triplet puzzle seems to be solved. With the new age and metallicity estimates presented in this Letter, the predicted and observed CaT* indices are in good agreement for all sample galaxies but two. The fact that CaT* values predicted using age/metallicity estimates that were derived from spectral features in the blue part of the spectrum agree with the observed CaT* values in the NIR indicates that the SSP assumption is not the cause of the CaT puzzle for dE's. Rather, the CaT puzzle in the dE regime was caused by the spuriously low metallicities and high ages, derived from lower resolution spectra using less sophisticated theoretical models, assigned to Fornax dE's. This shows that the CaT* index, and, in old stellar systems in which the PaT index is small, the CaT index as well, is indeed a good tracer of metallicity.  This also solves the apparent dichotomy between dE's on the one hand and globular clusters, Local Group dwarf spheroidals (dSph), and ultra-compact dwarfs (UCD) on the other hand. CaT line-strengths measured in individual stars of Local Group dSphs have been shown to be a very accurate tracer of metallicity \\citep{b06,to03} and [Fe/H]-values derived from CaT measurements have been used extensively to construct metallicity distributions of the stars in dSphs. Also, for UCDs \\citep{ev07} and globular clusters \\citep{saglia02} the CaT index has proved to be an excellent tracer of metallicity. Here, we have shown that in dE's as well, the CaT* index measured from integrated-light spectra can be used as a tracer of metallicity.   To summarise, we derive new age and metallicity estimates for 16 dE's in the Fornax Cluster and in Southern Groups using high S/N optical VLT/FORS1+2 spectra. We have measured the H$\\beta$, Mg$b$, Fe5270, and Fe5335 indices in the Lick/IDS system and applied the TMB03 models to them. We find that these dE's have solar [$\\alpha$/Fe] abundance ratios. A full-spectrum fit using Pegase-HR with the ELODIE.3.1 stellar library provides us with a second, independent age and metallicity estimate for these galaxies. We find both approaches to be in excellent agreement. With mean metallicity ${\\rm [Z/H]} = -0.35$\\,dex and ages younger than $\\approx 7$\\,Gyr, these dE's are more metal-rich and younger than previously thought.  Some even show strong emission lines, an indication of on-going star formation, in agreement with previous H$\\alpha$ imaging of dE's \\citep{derijcke03,michielsen04}. The ages and metallicities we derive for the dE's in the Fornax cluster and in Southern groups fall in roughly the same range as those derived by \\citet{g03} and \\citet{v04} for dE's in the Virgo cluster. This is at variance with previous estimates for Fornax dEs which yielded lower metallicities and higher ages \\citep{heldmould94, rakos01}, based on lower resolution spectra and less sophisticated theoretical models. The new age and metallicity estimates are in good agreement with the observed \\ionCaT\\ triplet absorption strengths, solving the Calcium puzzle for low-mass systems."
    },
    "0710/0710.4854_arXiv.txt": {
        "abstract": "Big Bang nucleosynthesis (BBN) is the earliest sensitive probe of the values of many fundamental particle physics parameters. We have found the leading linear dependences of primordial abundances on all relevant parameters of the standard BBN code, including binding energies and nuclear reaction rates. This enables us to set limits on possible variations of fundamental parameters. We find that ${^7}$Li is expected to be significantly more sensitive than other species to many fundamental parameters, a result which also holds for variations of coupling strengths in grand unified (GUT) models. Our work also indicates which areas of nuclear theory need further development if the values of ``constants'' are to be more accurately probed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.3190_arXiv.txt": {
        "abstract": "We set sensitive upper limits to the X-ray emission of four Type Ia supernovae (SNe~Ia) using the \\cxo.  SN~2002bo, a normal, although reddened, nearby SN~Ia, was observed 9.3 days after explosion. For an absorbed, high temperature bremsstrahlung model the flux limits are $3.2\\times 10^{-16}$ ergs cm$^{-2}$ s$^{-1}$ (0.5--2 keV band) and $4.1\\times 10^{-15}$ ergs cm$^{-2}$ s$^{-1}$ (2--10 keV band). Using conservative model assumptions and a 10 km s$^{-1}$ wind speed, we derive a mass loss rate of $\\dot{M} \\sim 2\\times 10^{-5}\\, M_\\odot$ yr$^{-1}$, which is comparable to limits set by the non-detection of H$\\alpha$ lines from other SNe~Ia.  Two other objects, SN~2002ic and SN~2005gj, observed 260 and 80 days after explosion, respectively, are the only SNe~Ia showing evidence for circumstellar interaction.  The SN~2002ic X-ray flux upper limits are $\\sim$4 times below predictions of the interaction model currently favored to explain the bright optical emission.  To resolve this discrepancy we invoke the mixing of cool dense ejecta fragments into the forward shock region, which produces increased X-ray absorption. A modest amount of mixing allows us to accommodate the \\chandra\\ upper limit. SN~2005gj is less well studied at this time. Assuming the same circumstellar environment as for SN~2002i, the X-ray flux upper limits for SN~2005gj are $\\sim$4 times below the predictions, suggesting that mixing of cool ejecta into the forward shock has also occurred here. Our reanalysis of Swift and \\chandra\\ data on SN~2005ke does not confirm a previously reported X-ray detection.  The host galaxies NGC~3190 (SN~2002bo) and NGC~1371 (SN~2005ke) each harbor a low luminosity ($L_X \\sim 3-4 \\times 10^{40}$ ergs s$^{-1}$) active nucleus in addition to wide-spread diffuse soft X-ray emission. ",
        "introduction": "Type Ia supernovae (SN~Ia) are an important subclass of supernova (SN) that are thought to arise from explosions of white dwarfs in binary systems, although the exact nature of their progenitor systems is largely a mystery.  This ignorance is not due to a lack of effort since SN~Ia are the subject of intense scrutiny, not least of all because of their importance as cosmological probes. SN~Ia have been used to measure the Hubble constant \\citep{ham95, riess96} and have even given strong evidence for a non-zero cosmological constant \\citep{riess98, perl99}.  Our failure to identify SN~Ia progenitors highlights a major gap in our understanding of stellar evolution in binary systems, and presents a stumbling-block to understanding the chemical evolution of galaxies, since Ia's are known to be efficient producers of iron.  In addition, until SN~Ia progenitor systems are clearly identified, it will remain difficult to convincingly eliminate evolutionary sources of systematic error in the observed trend of SN brightness with redshift.  The current model framework for Type Ia SNe involves carbon deflagration/detonation in a white dwarf driven close to the Chandrasekhar limit by accretion. The most likely type of progenitor system \\citep{branch95} is a C-O white dwarf accreting H/He-rich gas from a companion, either from its wind or through Roche lobe overflow. Double degenerate scenarios with coalescing pairs of C-O white dwarfs are also possible, while sub-Chandrasekhar-mass white dwarfs are less likely \\citep{branch95}.  In the non-coalescing scenarios, there will be circumstellar gas whose composition and geometry depend on the nature of the progenitor system.  If the circumstellar medium (CSM) emits radiation, or absorbs radiation from the SN, this can be used to distinguish between types of progenitor systems.  Thermal X-ray emission is expected to arise from the hot gas in the interaction region between the rapidly moving supernova ejecta and the wind from presupernova mass loss.  Detecting this X-ray emission offers a direct and potentially quite sensitive probe of the amount of CSM.  Because of its high sensitivity, broad bandwidth, unprecedented imaging capability, and ability to respond rapidly to a target of opportunity (ToO) request, the \\cxo\\ is the premier instrument to search for faint X-ray emission from transient point sources.  In the following we present the analysis and interpretation of \\chandra\\ observations of four recent SNe~Ia. Two are fairly normal SNe~Ia, while the others are peculiar cases showing strong H$\\alpha$ emission from circumstellar interaction.  The following section describes the observed targets, \\S3 presents the observations and techniques, \\S4 uses the upper limits on X-ray flux to constrain the nature of the CSM, and the final section concludes.  In an Appendix we describe the X-ray properties of the host galaxies for the two nearby systems, NGC~3190 and NGC~1371. ",
        "conclusions": "With our \\chandra\\ observation of SN~2002bo we have set the most sensitive X-ray flux upper limits at an earlier epoch than for any previous SN~Ia. The previous best case was that of SN~1992A for which a \\rosat\\ upper limit of $\\sim$$10^{-14}$ ergs cm$^{-2}$ s$^{-1}$ was set on the 0.5-2 keV band X-ray flux $\\sim$35 days after explosion \\citep{schpet93}.  Our \\chandra\\ limit in the same energy band is an order of magnitude lower and was set only 9 days after explosion.  We also set a sensitive upper limit in the important 2--10 keV band.  Converting our flux upper limits to constraints on the density of the CSM in the system depends on uncertain assumptions about the thermodynamic state of the hot plasma in the expanding ejecta and shocked wind.  As we have shown in this paper there is roughly an order of magnitude difference in the inferred wind density under the assumption of fully equilibrated electron and ion temperatures vs.\\ the case with non-equilibrated temperatures.  The latter case yields the more conservative (i.e., higher) constraint: a wind density parameter $w < 1.2 \\times 10^{15}$ g cm$^{-1}$.  In terms of a slow wind with velocity of $v_{w10}=10$ km s$^{-1}$ this corresponds to a upper limit on the mass-loss rate of $\\dot{M} < 2 \\times 10^{-5}~M_{\\odot}$ yr$^{-1}$. The much higher sensitivity of the \\chandra\\ data notwithstanding, this value is {\\em larger} by a factor of 10 or so than the mass loss upper limit derived from \\rosat\\ data by \\citet{schpet93} for SN~1992A.  This discrepancy is due to the simplicity of the model used by \\citet{schpet93} to evaluate their upper limit: they do not include absorption by residual wind material above the forward shock, assume electron-ion temperature equilibration, compare to model luminosities without applying a bolometric correction, and calculate the SN age from the time of maximum rather than from the time of explosion.  In other comparisons our \\chandra\\ upper limit is comparable to those found previously from limits on the H$\\alpha$ flux for the normal type Ia supernovae SN~1994D \\citep{cumm96} and SN~2001el \\citep{matetal05}.  In general limits on $\\dot{M}$ of $\\sim$$10^{-5}~M_{\\odot}$ yr$^{-1}$ are not stringent enough to rule out the class of symbiotic-type binaries as SN~Ia progenitors at least not for these particular cases.  The value of these results lies in our ability to calculate, using well understood physics, the expected X-ray emission from hot gas. Viewed in this light, we briefly discuss limits on the CSM obtained from radio nondetections of a number of nearby SNe~Ia \\citep{panagetal06}.  The radio results rely on semi-empirical parameterized functional forms for the time- and frequency-dependence of synchrotron emission and free-free absorption, whose relevant parameters are assumed to have values given by radio results for SNe~Ib/c. Likewise the essential parameter, i.e., the one linking the wind density parameter to the radio luminosity of the SN, cannot be calculated from theory with any accuracy, and an empirical calibration, again drawn from measurements of SNe~Ib/c, must be used. Although the sensitive radio flux limits clearly argue for low density environments around SN~Ia, the extremely low numerical limits on the mass-loss rates claimed by \\citet{panagetal06} cannot yet be considered definitive.  The other normal SN~Ia we study here is SN~2002ke.  We re-examine the claim by \\citet{immler06} of a tentative X-ray detection by Swift and find that we cannot substantiate it. We pay particular attention to the astrometric and photometric calibration of the Swift X-ray data by comparing to a \\chandra\\ observation done several months later.  We find no evidence for a significant X-ray detection of SN~2005ke by either Swift or \\chandra, to a flux limit that is several factors higher than what we obtained for SN~2002bo.  Since this limit is at a much later epoch, when the intensity of the CSM/ejecta interaction should be much reduced, we did not attempt to determine numerical limits to the wind density parameter for this SN.  We have also presented \\chandra\\ upper limits for the two examples of SNe Ia with clear evidence for circumstellar interaction: SN~2002ic and SN~2005gj.  The upper limit on the X-ray luminosity of SN~2002ic in the $0.5-6$ keV band unexpectedly reveals a serious drawback to the interaction model proposed previously (CCL): the predicted X-ray flux turns out to be larger by at least a factor of four.  We identified the major missing element of the model responsible for the controversy: macroscopic mixing of cool metal-rich ejecta fragments into the hot gas of the forward shock, which results in strong absorption of the X-rays emitted by the forward shock. Interestingly, the absorption of X-rays by mixed shocked ejecta should have the effect of decreasing the required CS density in the model to explain the late time optical luminosity. This effect together with the higher radiation efficiency of the interaction with a clumpy CS matter should result in a slightly higher expansion velocity of the shocked SN ejecta in the interaction model. These outcomes appear to be preferred by the optical spectra of SN~2002ic (see CCL).  SN~2005gj appears to belong to the SN~2002ic-like family of SN~Ia with dense CS envelopes \\citep{aldetal06,prietoetal07}. The \\chandra\\ upper limit from this object at about day 80 is a factor of four larger than the flux predicted by the interaction model of SN~2002ic recomputed for the corresponding epoch.  Although the strength of the interaction argues against it, one possible explanation is that the CS density around SN~2005gj is just somewhat lower than around SN~2002ic. A better possibility invokes mixing of the shocked ejecta, which can reduce the emergent X-ray flux. Which of these is more likely to be the correct explanation requires a better understanding of the environment of SN~2005gj than we have at present. Future studies of the optical light curve and spectra of this interesting object are strongy encouraged.  Resolving the issue of the X-ray non-detection of SN~2002ic-like objects has its dark side due to our introduction of the parameter $\\tau_{\\rm oc}$, which is essentially incalculable.  Even three-dimensional hydrodynamical simulations are unlikely to be able to recover this value in a fully self-consistent way. We, therefore, can predict in detail neither the spectrum nor the evolution of the X-ray flux from SN~2002ic-like supernovae. Two relevant remarks can be made, however. First, in the case of SN~2002ic the model without mixing predicts an increasing X-ray flux in the band below 10 keV up until day $\\sim$400 (CCL). Therefore, X-ray detection at earlier epochs is unlikely to be more favorable, as the lack of detection of SN~2005gj at roughly 80 days after explosion (versus 260 days for SN~2002ic) tends to support.  The second remark relates to the spectrum of the emergent X-ray emission: hard X-rays (i.e., with photon energies greater than $\\sim$20 keV) are not affected by the absorption that mutes the lower energy flux (see Fig.~\\ref{ic_mod}). Future sensitive X-ray observations covering a wider energy band ($1-30$ keV), therefore, should reveal a SN~2002ic-like event as a strongly absorbed hard X-ray source and thus verify the proposed mixing scenario."
    },
    "0710/0710.3159_arXiv.txt": {
        "abstract": "We present a spatially resolved comparison of the stellar-mass and total-mass surface distributions of nine early-type galaxies. The galaxies are a subset of the Sloan Lens ACS survey \\citep[or SLACS;][]{2006ApJ...638..703B}. The total-mass distributions are obtained by exploring pixelated mass models that reproduce the lensed images.  The stellar-mass distributions are derived from population-synthesis models fit to the photometry of the lensing galaxies.  Uncertainties -- mainly model degeneracies -- are also computed.  Stars can account for all the mass in the inner regions.  A Salpeter IMF actually gives too much stellar mass in the inner regions and hence appears ruled out.  Dark matter becomes significant by the half-light radius and becomes increasingly dominant at larger radii.  The stellar and dark components are closely aligned, but the actual ellipticities are not correlated.  Finally, we attempt to intuitively summarize the results by rendering the density, stellar-vs-dark ratio, and uncertainties as false-colour maps. ",
        "introduction": "In the current paradigm of galaxy formation, the building block of structure is a dark matter halo, consisting mainly of non-baryonic dark matter together with $\\sim 15\\%$ baryons in the form of gas.  Dark halos are thought to originate from the collapse of primordial density fluctuations, growing unimpeded until virial equilibrium is reached. Within these halos the baryonic component dissipates energy, collapsing further towards the center and eventually forming the visible galaxy. Subsequent mergers redistribute the matter within halos.  To work out the details of the basic picture, which goes back to \\cite{1978MNRAS.183..341W}, it is essential to determine the connection between the visible galaxies and dark halos. Strong gravitational lensing by galaxies is potentially a very useful way of doing this, since the total mass of a lensing galaxy is relatively easy to constrain.  Also, since lensing tends to be more effective for distant galaxies ($z_{\\rm lens}\\sim0.1$ to 1), it nicely complements the stellar-dynamical techniques applicable in nearby galaxies.  The difficulty with lensing galaxies (that is to say, galaxies producing multiple images of background sources) is that they are relatively rare.  Till recently, only about 80 were known.  But recently, 28 new galaxy lenses have been discovered by the Sloan Lens ACS Survey \\citep[SLACS;][]{2006ApJ...638..703B} thanks to a new survey strategy, which eventually may double the number of strong lenses. The method is to select galaxies from the Sloan survey \\citep{2000AJ....120.1579Y} that have emission lines indicating high-redshift background objects, and then to image these candidates using the Advanced Camera for Surveys on the Hubble Space Telescope.  A basic analysis of a sample of galaxy lenses is to fit simple lens models and then compute $M/L$ and its evolution with redshift.  This was first done by \\cite{1998ApJ...509..561K}.  An extension is to place lensing galaxies on the Fundamental Plane of ellipticals, using a measured or model-derived velocity dispersion \\citep{2000ApJ...543..131K,2003ApJ...587..143R,2006ApJ...640..662T}. Other work \\citep{2004ApJ...611..739T,2006ApJ...649..599K} compares lens models with the measured dispersions to constrain the mass profiles of the galaxies.  These studies have found no unexpected features or trends with redshift, and argue in favor of passive evolution.  A more detailed analysis involves modeling both the lens mass distribution and the stellar population. The star-formation history is not well-constrained by the observed fluxes and colours and must be marginalized over, but the stellar mass is fairly insensitive to model assumptions apart from the initial mass function (IMF).  The lensing mass distribution, when aggressively modeled, turns out to have much larger uncertainties than simple models assume; nevertheless, the uncertainties can be estimated and useful conclusions drawn.  \\cite{2005ApJ...623L...5F} found massive ellipticals to show a transition from no significant dark matter within $\\Re$ to dark-matter dominance by $5\\Re$, whereas lower-mass galaxies showed no significant dark halos even at $5\\Re$. The radial gradient in the dark-matter fraction agreed with the results on nearby galaxies derived from stellar dynamics \\citep{2005MNRAS.357..691N}.  In this paper we extend the detailed comparison of stellar and total mass to two dimensions, using a subsample of the SLACS lenses.  All the SLACS objects have a background galaxy that is lensed into two or four extended images. In nine of the objects, we can identify small features within the extended images.  For lenses showing point-like multiply-imaged features there is a well-developed technique for reconstructing the projected mass distribution, along with uncertainty estimates \\citep{2004AJ....127.2604S}.  Accordingly we take these nine lenses as our sample.  The stellar mass content is estimated by combining the available photometry of the lensing galaxies with population-synthesis models \\citep{2003MNRAS.344.1000B}.  \\begin{figure} \\begin{center} \\includegraphics[width=3.4in]{2DLens_f1.eps} \\caption{Stellar masses derived from the available photometry. The predictions of the Bruzual \\& Charlot (2003) models are shown for a Chabrier (2003) IMF. A galaxy with apparent magnitude $\\rm F814W=20$ is considered at two different redshifts as labelled. The shaded regions span the model predictions when assuming a wide range of $\\tau$ models (i.e., an exponentially decaying star formation history at fixed metallicity). The dark/light shaded regions correspond to a formation redshift of $z_F=5$ and $2$, respectively, and all span a range of star formation timescales ($-1<\\log \\tau ({\\rm Gyr}) < +1$) and metallicities ($-1<\\log Z/Z_\\odot <+0.3$).} \\label{fig:stellar} \\end{center} \\end{figure} ",
        "conclusions": "The potential usefulness of early-type lensing galaxies for understanding the interdependence of baryons and dark matter in galaxy formation and evolution is widely appreciated. Lensing can be used to map the total mass and starlight can be used to map the baryons.  In this paper we do this for nine galaxies from the SLACS survey by \\cite{2006ApJ...638..703B}.  From the lensing data we derived free-form pixelated models for the total mass, and from the galaxy photometry we computed stellar population-synthesis models.  In both casses we generated Monte-Carlo ensembles of models, in order to marginalize over unknowns such as lensing degeneracies and star-formation histories, thus obtaining realistic uncertainties. The technique is basically the same as in \\cite{2005ApJ...623L...5F}, but whereas the earlier work only compared radial profiles now we compare stellar and total mass in 2D.  Related work has been done on larger samples of galaxies, but is limited to fitting simple parameterized models for the lens, and does not model the stellar populations at all.  The 2D mass maps are shown in Figures \\ref{fig:mass} and \\ref{fig:cmaps}.  The former overlays contours of $\\Sigtot$ on a grayscale of $\\Sigstel$.  The latter shows the same information with uncertainties as well, all encoded in false colour: red for stellar mass, blue for dark, and pale versus coloured for uncertainty.  It is evident that (a)~these galaxies are dominated by stars in the inner regions, but mainly dark matter in the outer regions, and (b)~stellar and dark components are well-aligned, but neither has a simple elliptical shape.  One can of course still compute an ellipticity defined as a moment, and this is shown in Figure~\\ref{fig:ellip}. We see that ellipticities of the stellar and total mass are uncorrelated in magnitude but are almost perfectly aligned.  It would be interesting to see if this is true of galaxy-formation simulations.  The profiles (Figure~\\ref{fig:profile}) show that dark matter halos begin to dominate around the half-light radius, although some galaxies seem to present a stronger contribution from dark matter even inside $\\Re$ (e.g., J2300).  The present sample is too small to extract any strong correlation of the dark matter distribution with global properties such as total mass or luminosity. However, the trend found previously \\citep{2005ApJ...623L...5F}, namely that there should be more dark matter in more massive galaxies, with a scaling of roughly $\\Mtot\\propto\\Mstel^{1.2}$ (which is equivalent to the tilt of the Fundamental Plane) is compatible with the combined data of lensing galaxies (labelled CASTLES and SLACS) along with the dynamical analysis of local galaxies with the SAURON integral field unit \\citep{2006MNRAS.366.1126C}.  The top-left panel of Figure~\\ref{fig:profile} also shows that a Salpeter (1955) IMF cannot be used to estimate stellar masses as the population synthesis models predict too much stellar mass compared to the total mass obtained from our lensing studies. This result also agrees with the analysis of the SAURON sample.  We emphasize that stellar and total masses are obtained from different data through models of different physical processes.  There is no tuning to give similar results.  That the stellar and total densities come out compatible at the centers ---where the baryon content is expected to dominate the mass budget--- indicates that systematic effects are not significant."
    },
    "0710/0710.4065_arXiv.txt": {
        "abstract": "{The majority of X-ray burst sources do not display a burst rate that increases with luminosity as expected, but this is seen in the two clocked bursters XB\\th 1323-619 and GS\\th 1826-24. We present a detailed investigation of these two sources which in the case of the first source, spans 18 years. Based on measurements of the burst rate, X-ray luminosity, the $\\alpha$-parameter and the two time constants generally present in the burst decays, we demonstrate the importance of the {\\it rp} nuclear burning process. A detailed comparison with theory shows that although the burst rate in each source agrees well with the theoretical value, there is a difference of more than a factor of 5 in the burst rate at a given luminosity between the sources. We show that the main reason for this is that the two sources have substantially different emitting areas on the neutron star in non-burst emission, a factor often neglected. Variation of this area may explain the inverse relation of burst rate with luminosity in the majority of burst sources. ",
        "introduction": "% \\label{sect:intro}  X-ray bursting has been known for many years (Grindlay et al. 1976) as a phenomenon taking place in many low mass X-ray binaries (LMXB) of luminosities less than 10$^{38}$ erg s$^{-1}$ consisting of a rapid rise in intensity by a factor of $\\sim$20, followed by an exponential decay over about 50 seconds and repeating on a timescale of hours. Bursting is usually not seen in higher luminosity sources forming the Z-track class which exhibit intensity increases as flaring over much longer timescales. Thus if we divide LMXB by luminosity into Atoll or Z-track sources, X-ray bursting is generally observed in the Atoll class. Bursting appears unrelated to inclination angle since it is seen not only in Atoll sources not displaying orbital-related behaviour but also in dipping sources which are seen at high inclination. X-ray bursting has been extensively studied (Lewin et al. 1995; Strohmayer \\& Bildsten 2006) and it was realized at an early stage to consist of unstable nuclear burning. Measured values of the $\\alpha$-parameter, defined as the ratio of the energy released in steady mass accretion to the integrated energy of the burst, agreed well with values expected for unstable burning. Thus X-ray bursting was recognized as explosive burning of recently accreted material on the surface of the neutron star (Woosley \\& Taam 1976), the material building up on the neutron star over a period of hours.  In spite of this, understanding of X-ray bursting is relatively poor. On the above basis the rate of X-ray bursting would be expected to be generally stable when the luminosity of a sources is stable since then the mass accretion rate $\\dot M$ is stable, but bursting rate is rather erratic in most sources. Even worse, the rate of bursting should increase when source luminosity increases, as less time is required to accumulate the necessary mass and achieve the required plasma density and temperature on the surface of the neutron star for unstable burning, but many sources display the opposite of this (van Paradijs et al. 1988).  In two sources, which may be called the ``clocked bursters'', bursting is close to being exactly periodic and the burst rate increases with luminosity: GS\\th 1826-24 (the original clocked burster) and XB\\th 1323-619 which we have studied over an extended period as one of the $\\sim$10 dipping LMXB. In the present work, we make detailed comparisons of these two well-behaved sources with each other, and with the theory of unstable burning, to identify the reasons why the sources differ so much from each other, and to try to understand their behaviour in terms of theory. This would facilitate the understanding of the majority of burst sources that are not well-behaved.    \\begin{figure*}[!ht]                                                           % \\begin{center} \\includegraphics[width=40mm,height=140mm,angle=270]{balucinska_2007_revised_fig1} \\caption{{\\it Exosat} lightcurve of XB\\th 1323-619. Regular bursting takes place every 5.33 hr whereas X-ray dipping occurs at the orbital period of 2.93 hr.} \\label{} \\end{center} \\end{figure*} ",
        "conclusions": "We have shown that marked differences in the burst rate occur between XB\\th 1323-619 and GS\\th 1826-24. and have shown that the difference is mostly due to the very different emitting areas of non-burst emission on the neutron star, a factor that has not generally been allowed for in discussion of X-ray bursting. If we view the burst rate as a function of $\\dot m$, the accretion rate per unit emitting area, we find that the sources follow relations that agree within 40\\%, so that both can be explained to this accuracy by the same model. It is expected that the non-burst emitting area will be an important factor in the other burst sources and in future work we will test the hypothesis that in sources where the burst rate decreases with $L$, this behaviour can be explained in terms of the non-burst area increasing with $L$ so that $\\dot m$ decreases."
    },
    "0710/0710.5580_arXiv.txt": {
        "abstract": "Ram pressure stripping can remove significant amounts of gas from galaxies that orbit in clusters and massive groups, and thus has a large impact on the evolution of cluster galaxies.  In this paper, we reconstruct the present-day distribution of ram-pressure, and the ram pressure histories of cluster galaxies.  To this aim, we combine the Millennium Simulation and an associated semi-analytic model of galaxy evolution with analytic models for the gas distribution in clusters.  We find that about one quarter of galaxies in massive clusters are subject to strong ram-pressures that are likely to cause an expedient loss of all gas. Strong ram-pressures occur predominantly in the inner core of the cluster, where both the gas density and the galaxy velocity are higher.  Since their accretion onto a massive system, more than 64 per cent of galaxies that reside in a cluster today have experienced strong ram-pressures of $>10^{-11}$ dyn cm$^{-2}$ which most likely led to a substantial loss of the gas. ",
        "introduction": "In clusters, galaxies can lose some or all of their gas by ram pressure stripping (RPS) due to their motion through the intracluster medium (ICM). Both analytical estimates and hydrodynamical simulations show that RPS can remove a significant amount of gas from galaxies, and can thus explain observations such as the HI deficiency of cluster disc galaxies (see e.g.~\\citealt{Haynes_Giovanelli_1986,Solanes_etal_2001,cayatte94}), and the truncated star forming discs in the Virgo cluster (\\citealt{Warmels_1988,koopmann04b}).  RPS only affects the gaseous component of the galaxy so that a distinct feature of ram pressure stripped galaxies is that their gas discs are distorted or truncated while their stellar discs appear undisturbed. An increasing number of observations of ram-pressure stripped galaxies have become available in the last years (see for example \\citealt{Irwin_etal_1987}, \\citealt{Veilleux_etal_1999}, \\citealt{kenney04}, \\citealt{vollmer04a}, \\citealt{chung:07}, \\citealt{sakelliou:05,sun:07}).  RPS is commonly cited in early work as a possible explanation for the increased fraction of S0 galaxies in rich clusters relative to the field (\\citealt{Biermann_Tinsley_1975}, \\citealt{Butcher_Oemler_1978}). This explanation was dismissed in the original paper by \\cite{dressler_1980} on the basis of the observation that the relation between different galaxy populations and local density appears to hold independently of the cluster richness. Later studies pointed out that additional mechanisms that lead to a significant redistribution of mass and/or new star formation are required to explain the entire S0 population of galaxy clusters (see for example \\citealt{moran:07} and references therein).  The role of RPS in the chemical enrichment of the ICM has also been discussed. Observational data suggest that the ICM is enriched with metals to approximately one third of the solar value, suggesting that some fraction of the metals must have originated from cluster galaxies and since been removed from them.  Processes responsible for the supply of this enriched gas include AGN feedback (see e.g. \\citealt{roediger:07c} and references therein), galactic winds driven by supernovae explosions, and ram-pressure stripping (\\citealt{white:91,mori:00,schindler05,domainko:06}). It should be noted, however, that although RPS certainly affects the metallicity of the ICM, it may not be the dominant mechanism.  Numerical simulations by \\cite{domainko:06} indicate that RPS can account for only about 10 per cent of the ICM metal content within a radius of 1.3 Mpc.  The first analytical estimate of RPS dates back to the paper by \\citet{gunn72} who proposed that for galaxies moving face-on through the ICM the success or failure of RPS can be predicted by comparing the ram pressure with the galactic gravitational restoring force per unit area.  Later hydrodynamical simulations of RPS (\\citealt{abadi99}, \\citealt{quilis00}, \\citealt{schulz01}, \\citealt{marcolini03,acreman03}, \\citealt{roediger:06a}, \\citealt{roediger:06b}, \\citealt{roediger:07}) suggest that this analytical estimate fares fairly well as long as the galaxies are not moving close to edge-on. The ICM-ISM interaction is, however, a complex process influenced by many parameters. Different aspects have been studied by several groups. \\cite{roediger05} and others have shown that the ICM-ISM interaction is a multistage process: The most important phases are the instantaneous stripping, on a time-scale of a few 10 Myr, an intermediate phase, on a time-scale of up to a few 100 Myr, and a continuous stripping phase that, in principle, could continue until all gas is lost from the galaxy.  While numerical simulations and observations indicate that RPS has important consequences on the amount of gas in cluster galaxies, this physical process is usually not included in semi-analytic models of galaxy formation.  The effect of ram-pressure stripping has been discussed only in a couple of studies using semi-analytic models (\\citealt{okamoto:03, lanzoni:05}), and is shown to produce only mild variations in galaxy colours and star formation rates.  This happens because the stripping of the hot gas from galactic haloes (strangulation) suppresses the star formation so efficiently that the effect of ram-pressure is only marginal. We note that the studies mentioned above include ram-pressure stripping based on the analytical criterion formulated originally in Gunn \\& Gott (1972).  In recent numerical work, \\cite{roediger:07} have shown that this formulation often yields incorrect mass loss rates. Other numerical studies (e.g. \\cite{vollmer01a}) have argued that ram-pressure stripping can also temporarily enhance star formation. An updated modelling of ram-pressure stripping that takes into account these results has not been included yet in semi-analytic models. We plan to address this in future studies.  For a study of the ram pressure distribution, it is necessary to have information on the dynamics of galaxies and on the properties of the ICM. The dynamics of dark matter halos has been studied in a number of papers using numerical simulations (e.g.  \\citealt{benson:05,khochfar:05, diemand:04}). However, we know of no study on the distribution and history of ram pressures experienced by galaxies in clusters. If ram-pressure plays some role in establishing the observed morphological mix in galaxy clusters and/or the observed radial trends, it is important to quantify the distribution and history of ram pressures experienced by galaxies that reside in clusters today. This is the subject of this paper. ",
        "conclusions": "In this study we have taken advantage of the Millennium simulation by \\cite{springel:05} and of the publicly available semi-analytic model by De Lucia \\& Blaizot (2007) to reconstruct the ram-pressure distribution and ram-pressure history of galaxies that reside in clusters at the present epoch.  The gas profile in dark matter is described through two analytic models which give fairly similar results. We have not included ram-pressure stripping self-consistently in the semi-analytic model employed for our study. Instead we have simply used the available information about the orbital distribution and galaxy merging trees to estimate the importance of ram-pressure stripping on galaxies that reside in massive haloes at the present epochs.  We find that more than half of the galaxies in a $M=10^{15}M_{\\odot}$ cluster, have experienced ram pressures of $> 10^{-11}$ dyn cm$^{-2}$ after their accretion onto a massive system. This fraction is only slightly lower for $M=10^{14}M_{\\odot}$ clusters, implying that a significant fraction of galaxies in clusters at the present epoch suffered substantial gas loss due to ram pressure stripping. The fraction of galaxies that suffered significant ram-pressure after accretion increases with decreasing distance from the cluster centre, in qualitative agreement with the observed increase of early-type galaxies.  As expected, strong episodes of ram-pressure occur predominantly in the inner core of galaxy clusters, and are restricted to within the virial radius. Our results show, however, that virtually all the galaxies in clusters suffered weaker episodes of ram pressure, suggesting that this physical process might have a significant role in shaping the observed properties of the entire cluster galaxy population.  The limited number of simulation outputs does not allow us to reconstruct accurately the orbit of the cluster galaxies, and therefore their ram-pressure histories. Our result show that ram pressure fluctuates strongly so that episodes of strong ram-pressure alternate to episode of weaker ram-pressure. During these time intervals, the gaseous reservoir could be replenished and new episodes of star formation could occur. Our results indicate that ram-pressure stripping must play a significant role in the evolution of galaxies residing in massive clusters.  A more self-consistent modelling is however required in order to draw more quantitative conclusions about the importance and effects of this physical process."
    },
    "0710/0710.5307_arXiv.txt": {
        "abstract": "Several kinds of astronomical observations, interpreted in the framework of the standard Friedmann--Robertson--Walker cosmology, have indicated that our universe is dominated by a Cosmological Constant. The dimming of distant Type Ia supernovae suggests that the expansion rate is accelerating, as if driven by vacuum energy, and this has been indirectly substantiated through studies of angular anisotropies in the cosmic microwave background (CMB) and of spatial correlations in the large-scale structure (LSS) of galaxies. However there is no compelling {\\em direct} evidence yet for (the dynamical effects of) dark energy.  The precision CMB data can be equally well fitted without dark energy if the spectrum of primordial density fluctuations is not quite scale-free and if the Hubble constant is lower globally than its locally measured value. The LSS data can also be satisfactorily fitted if there is a small component of hot dark matter, as would be provided by neutrinos of mass $\\sim0.5$ eV. Although such an Einstein--de Sitter model cannot explain the SNe~Ia Hubble diagram or the position of the `baryon acoustic oscillation' peak in the autocorrelation function of galaxies, it may be possible to do so e.g. in an inhomogeneous Lemaitre--Tolman--Bondi cosmology where we are located in a void which is expanding faster than the average. Such alternatives may seem contrived but this must be weighed against our lack of any fundamental understanding of the inferred tiny energy scale of the dark energy. It may well be an artifact of an oversimplified cosmological model, rather than having physical reality. ",
        "introduction": "Following his formulation of general relativity Einstein \\cite{Einstein:1917} boldly applied the theory to the universe as a whole. The first cosmological model was {\\em static} to match the known universe, which at that time was restricted to the Milky way, and to achieve this Einstein introduced the `cosmological constant' term (for a historical perspective, see \\cite{Straumann:2002he}). Within a decade however Slipher and Hubble demonstrated that the nebulae on the sky are in fact other `island universes' like the Milky Way and that they are mainly receeding from us --- the universe is expanding. Einstein wrote to Weyl in 1933: {\\em ``If there is no quasi-static world, then away with the cosmological term''}.  This however is not a matter of choice since general coordinate invariance, which Einstein's equation is based on, permits an arbitrary constant (multiplied by the metric tensor) to be added to the lhs: \\begin{equation} R_{\\mu\\nu} - \\frac{1}{2}g_{\\mu\\nu}R + \\lambda_\\mathrm{metric} g_{\\mu\\nu} = \\frac{-T_{\\mu\\nu}}{M_\\mathrm{P}^2}. \\label{gr} \\end{equation} Here we have written Newton's constant, $G_\\mathrm{N} \\equiv 1/8\\pi\\,M_\\mathrm{P}^2$ where $M_\\mathrm{P} \\simeq 2.4\\times10^{18}$~GeV is the (reduced) Planck mass in natural units ($\\hbar = k_\\mathrm{B} = c = 1$). With the subsequent development of quantum field theory it became clear that the energy-momentum tensor on the rhs can also be freely scaled by another additive constant multiplying the metric tensor, which reflects the (Lorentz invariant) energy density of the vacuum: \\begin{equation} \\langle T_{\\mu\\nu}\\rangle_{\\rm fields} = -\\langle\\rho\\rangle_{\\rm fields}g_{\\mu\\nu}. \\end{equation} This contribution from the matter sector adds to the ``bare'' term from the background geometry, yielding an effective cosmological constant: \\begin{equation} \\Lambda = \\lambda_\\mathrm{metric} + \\frac{\\langle\\rho\\rangle_{\\rm fields}}{M_\\mathrm{P}^2}, \\label{tuning} \\end{equation} or, correspondingly, an effective vacuum energy: \\begin{equation} \\rho_\\mathrm{v} \\equiv \\Lambda M_\\mathrm{P}^2. \\label{rhov} \\end{equation}  Einstein {\\em assumed} without any observational evidence that the universe is perfectly homogeneous. We know that the universe is quite isotropic about us so this is in fact likely if we are not in a special location --- an assumption later dignified by Milne as the `Cosmological Principle'. Then using the {\\em maximally symmetric} Robertson-Walker metric to describe space-time \\begin{equation} \\mathrm{d}s^2 \\equiv g_{\\mu\\nu} \\mathrm{d}x^\\mu \\mathrm{d}x^\\nu = \\mathrm{d}t^2 - a^2(t)[ \\mathrm{d}r^2/(1 - kr^2) + r^2 \\mathrm{d}\\Omega^2], \\label{rw} \\end{equation} we obtain the Friedmann equations describing the evolution of the cosmological scale-factor $a (t)$: \\begin{equation} H^2 \\equiv \\left(\\frac{\\dot{a}}{a}\\right)^2 = \\frac{\\rho}{3M_\\mathrm{P}^2} - \\frac{\\kappa}{a^2} + \\frac{\\Lambda}{3}, \\qquad \\frac{\\ddot{a}}{a} = -\\frac{1}{6M_\\mathrm{P}^2}(\\rho + 3p) + \\frac{\\Lambda}{3}, \\label{fried} \\end{equation} where $\\kappa = 0, \\pm1$ is the 3-space curvature signature and for ordinary matter (`dust') and radiation we have used the `ideal gas' form: $T_{\\mu\\nu} = p g_{\\mu\\nu} + (p + \\rho) u_\\mu u_\\nu$, with $u_\\mu \\equiv \\mathrm{d}x_\\mu/\\mathrm{d}s$. The conservation equation $T^{\\mu\\nu}_{;\\nu} = 0$ implies $\\mathrm{d}(\\rho a^3)/\\mathrm{d}a = -3 p a^2$, so given the `equation of state parameter' $w \\equiv p/\\rho$, the evolution history can now be constructed. Since the redshift is $z \\equiv a/a_0 - 1$, for non-relativistic particles with $w \\simeq 0$, $\\rho_\\mathrm{NR} \\propto (1 + z)^{-3}$, while for relativistic particles with $w = 1/3$, $\\rho_\\mathrm{R} \\propto (1 + z)^{-4}$, but for the cosmological constant, $w = -1$ and $\\rho_\\mathrm{v} = $constant. Thus radiation was dynamically important only in the early universe (for $z \\geqsim 10^4$) and for most of the expansion history only non-relativistic matter is relevant. The Hubble equation can be rewritten with reference to the present epoch (subscript 0) as \\begin{eqnarray} H^2 &=& H_0^2 \\left[\\Omega_\\mathrm{m} (1 + z)^3 + \\Omega_\\kappa (1 + z)^2 + \\Omega_\\Lambda \\right], \\\\ \\Omega_\\mathrm{m} &\\equiv& \\frac{\\rho_{\\mathrm{m}0}}{\\rho_\\mathrm{c}}, \\quad \\Omega_\\kappa \\equiv -\\frac{\\kappa}{a_0^2 H_0^2}, \\quad \\Omega_\\Lambda \\equiv \\frac{\\Lambda}{3 H_0^2}, \\label{hub} \\end{eqnarray} where $\\rho_\\mathrm{c} \\equiv 3H_0^2M_\\mathrm{P}^2/8\\pi \\simeq (3 \\times 10^{-12} \\,\\mathrm{GeV})^4 h^2$ is the `critical density' and the present Hubble parameter is $H_0 \\equiv 100h~\\mathrm{km}~\\mathrm{s}^{-1}~\\mathrm{Mpc}^{-1}$ with $h \\simeq 0.7$, i.e. about $10^{-42}$~GeV.  This yields the sum rule \\begin{equation} \\Omega_\\mathrm{m} + \\Omega_\\kappa + \\Omega_\\Lambda = 1, \\label{sumrule} \\end{equation} so cosmological models can be usefully displayed on a `cosmic triangle' \\cite{Bahcall:1999xn}.  As emphasised in an influential review \\cite{Weinberg:1988cp}, given that the density parameters $\\Omega_\\mathrm{m}$ and $\\Omega_\\kappa$ were observationally constrained already to be not much larger than unity, the two terms in eq.(\\ref{tuning}) are required to somehow {\\em conspire} to cancel each other in order to satisfy the approximate constraint \\begin{equation} |\\Lambda| \\leqsim H_0^2, \\end{equation} thus bounding the present vacuum energy density by $\\rho_\\mathrm{v} \\leqsim 10^{-47}\\,\\mathrm{GeV}^4$ which is a factor of over $10^{120}$ below its ``natural'' value of $\\sim M_\\mathrm{P}^4$ --- the `cosmological constant problem'. Subsequently, several types of evidence have been advanced to argue that this inequality is in fact saturated with $\\Omega_\\Lambda \\simeq 0.7$ ($\\Rightarrow \\Lambda \\simeq 2H_0^2$), $\\Omega_\\mathrm{m} \\simeq 0.3$, $\\Omega_\\kappa \\simeq 0$ (see \\cite{Sahni:1999gb,Peebles:2002gy}), i.e. there is non-zero vacuum energy of {\\em just the right order of magnitude so as to be detectable today}.  In the Lagrangian of the Standard $SU(3)_\\mathrm{c} \\otimes SU(2)_\\mathrm{L} \\otimes U(1)_Y$ Model (SM) of electroweak and strong interactions, the term corresponding to the cosmological constant is one of the two `super-renormalisable' terms allowed by the gauge symmetries, the second one being the quadratic divergence in the mass of fundamental scalar fields due to radiative corrections (see \\cite{Zwirner:1995iw}). To tame the latter sufficiently in order to explain the experimental success of the SM has required the introduction of a `supersymmetry' between bosonic and fermionic fields which is `softly broken' at about the Fermi scale, $M_\\mathrm{EW}\\sim\\,G_\\mathrm{F}^{-1/2}\\simeq246$~GeV Thus the cutoff scale of the SM, viewed as an effective field theory, can be lowered from $M_\\mathrm{P}$ down to $M_\\mathrm{EW}$, at the expense of introducing over 100 new parameters in the Lagrangian, as well as requiring delicate control of the non-renormalisable operators which can generate flavour-changing neutral currents, nucleon decay etc, so as not to violate experimental bounds. This implies a {\\em minimum} contribution to the vacuum energy density from quantum fluctuations of ${\\cal O}(M_\\mathrm{EW}^4)$, i.e. halfway on a logarithmic scale down from $M_\\mathrm{P}$ to the energy scale of ${\\cal O}(M_\\mathrm{EW}^2/M_\\mathrm{P})$ corresponding to the observationally indicated vacuum energy. Thus even the introduction of supersymmetry cannot eradicate a discrepancy by a factor of at least $10^{60}$ between the natural expectation and observation.  It is likely that a satisfactory resolution of the cosmological constant problem can be achieved only in a quantum theory of gravity. Recent developments in string theory and the possibility that there exist new dimensions in Nature have generated many interesting ideas concerning possible values of $\\Lambda$ (see e.g. \\cite{Witten:2000zk,Padmanabhan:2002ji,Nobbenhuis:2004wn}). Nevertheless it is the case that there is no accepted solution to the enormous discrepancy discussed above. The problem is not new but there has always been the hope that some day we would understand why $\\Lambda$ is exactly zero, perhaps due to a new symmetry principle. However if it is in fact non-zero and dynamically important {\\em today}, the crisis is even more severe since this also raises a cosmic `coincidence' problem, viz. why is the present epoch special?~\\footnote{While $\\Lambda \\sim H_0^2$ today, we cannot have $\\Lambda \\sim H^2$ {\\em always} since this would amount to a substantial renormalisation of $M_\\mathrm{P}$ in eq.(\\ref{fried}) (taking $\\kappa=0$ as suggested by inflation); this is inconsistent with primordial nucleosynthesis which requires the ``cosmic'' Newton's Constant to be within a few \\% of its laboratory value \\cite{Cyburt:2004yc}. Thus arguably the most natural solution to the `coincidence problem' is ruled out empirically.} It has been suggested that the `dark energy' may not be a cosmological {\\em constant} but rather the slowly evolving potential energy $V (\\phi)$ of a hypothetical scalar field $\\phi$ named `quintessence' which can {\\em track} the matter energy density (see \\cite{Copeland:2006wr}). This however is equally fine-tuned since we need $V^{1/4} \\sim 10^{-12}$~GeV but $\\sqrt{\\mathrm{d}^2 V/\\mathrm{d}\\phi^2} \\sim H_0 \\sim 10^{-42}$~GeV (in order that the evolution of $\\phi$ be sufficiently slowed by the Hubble expansion), and moreover does not address the fundamental issue,\\footnote{Admittedly this criticism applies also to attempts to do away with dark energy by interpreting the data in terms of modified cosmological models, but less so since these do not invoke gravitating vacuum energy at all.}  namely why are all the other possible contributions to the vacuum energy absent?  Given the no-go theorem against any dynamical cancellation mechanism in eq.(\\ref{tuning}) in the framework of general relativity \\cite{Weinberg:1988cp}, it might appear that solving the problem will necessarily require modification of our understanding of gravity. Interesting suggestions have been made in this context e.g. the DGP model in which gravity alone propagates in a new dimension that opens up on distance scales of ${\\cal O}(H_0^{-1})$ (see \\cite{Lue:2005ya}). However since this is an unnaturally large scale for a fundamental theory, this model is clearly just as fine tuned as the cosmological constant and moreover suffers from intrinsic theoretical difficulties such as violation of unitarity due to `ghosts' (see \\cite{Koyama:2007za}).  The situation is so desperate that `anthropic' arguments have been advanced to explain why the cosmological constant is of just the right order of magnitude to allow of our existence today --- if it was much higher then galaxies would never have have formed (see \\cite{Weinberg:2000yb}). However a recent analysis \\cite{Tegmark:2005dy} shows that the most likely value from such Bayesian arguments is in fact 20--30 times bigger than the observationally suggested value. Moreover, one cannot claim to have a rigorous understanding of the prior probability distribution, lacking a theory of the cosmological constant itself. It is commonly assumed that the prior is flat in $\\Lambda$, but if it is instead flat in say $\\log(\\Lambda)$, the most likely value of $\\Lambda$ would be zero. A flat prior in $\\Lambda$ is indeed expected in quintessence-like models with a very flat potential \\cite{Weinberg:2000qm}, but such models are highly fine-tuned and have little physical basis. Whereas an uniform distribution of $\\Lambda$ does arise in the 'landscape' of the large number ($\\sim 10^{500}$) of possible vacua in string theory (see \\cite{Douglas:2006es}), there is no accepted measure for how these vacuua may come to be populated through cosmological evolution.  Given this situation, we can ask whether the observations really require dark energy or whether this is just an artifact of interpreting the data using an over-idealised description of the universe. For the FRW model we see from eq.(\\ref{fried}) that deducing $\\Lambda$ to be of ${\\cal O}(H_0^2)$ from observations should not be unexpected, this being its {\\em natural} scale in such a model universe (what seems quite unnatural is the {\\em implied} fundamental energy scale of $\\sqrt{H_0 M_\\mathrm{P}}$ since $H_0$ is so much smaller than $M_\\mathrm{P}$). From this perspective, it is easy to see why there have been recurring claims for $\\Lambda \\sim H_0^2$ --- this is effectively forced upon us by the theoretical framework in which the data is interpreted. ",
        "conclusions": "We now have a `cosmic concordance' model with $\\Omega_\\mathrm{m} \\sim 0.3, \\Omega_\\Lambda \\sim 0.7$ which is supposedly consistent with all astronomical data but has no satisfactory explanation in terms of fundamental physics. One might hope to eventually find explanations for the dark matter (and baryonic) content of the universe in the context of physics beyond the Standard Model but there appears to be little prospect of doing so for the apparently dominant component of the universe --- the dark energy which behaves as a cosmological constant. Cosmology has in the past been a data-starved science so it has been appropriate to consider the simplest possible cosmological models in the framework of general relativity. However now that we are faced with this serious confrontation between fundamental physics and cosmology, it would seem prudent to reconsider the underlying assumptions, especially that of exact homogeneity.  The observed isotropy of the CMB along with a belief in the Cosmological Principle has generally been taken to imply homogeneity but this has come to be questioned of late. There is growing interest in the inhomogeneous but isotropic Lemaitre-Tolman-Bondi model of a local void (see \\cite{Krasinski:1997}). It has been shown \\cite{Tomita:1999qn,Tomita:2000rf,Tomita:2000jj,Tomita:2001gh,Tomita:2002df} that a LTB model can in principle give an explanation of the oddities in the local Hubble flow as well as account for the SNe~Ia Hubble diagram without invoking acceleration (see also \\cite{Celerier:1999hp,Alnes:2005rw,Enqvist:2006cg,Alnes:2006uk}). Recently it has been claimed \\cite{Biswas:2006ub} that with a large local void of size $\\sim 450h^{-1}$~Mpc the angular diameter distance of the BAO peak can also be matched in this model. Such local inhomogeneity may be responsible for generating the mysteriously aligned low multipoles in the {\\sl WMAP} sky \\cite{Inoue:2006rd}. A void of this size does seem extremely unlikely, but one has recently been actually detected at $z \\sim 1$ and appears to be responsible for the anomalously `cold spot' seen by {\\sl WMAP} \\cite{Rudnick:2007kw}.  Landau famously said {\\em ``Cosmologists are often wrong, but never in doubt''}. The situation today is perhaps better captured by Pauli's enigmatic remark --- the present interpretation of the data may be {\\it ``\\ldots not even wrong''}. However we are certainly not without doubt! The crisis posed by the recent astronomical observations is not one that confronts cosmology alone; it is the spectre that haunts any attempt to unite two of the most successful creations of 20th century physics --- quantum field theory and general relativity. It would be fitting if the cosmological  constant which Einstein allegedly called his {\\it ``biggest blunder''} proves to be the catalyst for triggering a new revolution in physics in this century."
    },
    "0710/0710.0378_arXiv.txt": {
        "abstract": "Various mechanisms have been proposed to explain the inflated size of HD~209458b after it became clear that it has no companions capable of producing a stellar reflex velocity greater than around $5\\,\\ms$. Had there been such a companion, the hypothesis that it forces the eccentricity of the inflated planet thereby tidally heating it may have been readily accepted. Here we summarize a paper by the author which shows that companion planets with masses as low as a fraction of an Earth mass are capable of sustaining a non-zero eccentricity in the observed planet for at least the age of the system. While such companions produce stellar reflex velocities which are fractions of a meter per second and hence are below the stellar jitter limit, they are consistent with recent theoretical work which suggests that the planet migration process often produces low-mass companions to short-period giants. ",
        "introduction": "Of the fourteen transiting extrasolar planetary systems for which radii have been measured, at least three appear to be considerably larger than theoretical estimates suggest. It has been proposed by \\citet{b4} that undetected companions acting to excite the orbital eccentricity are responsible for these oversized planets, as they find new equilibrium radii in response to being tidally heated. In the case of HD 209458, this hypothesis has been rejected by some authors because there is no sign of such a companion at the 5 ms$^{-1}$ level, and because it is difficult to say conclusively that the eccentricity is non-zero. Transit timing analysis as well as a direct transit search have further constrained the existence of very short-period companions, especially in resonant orbits. Whether or not a companion is responsible for the large radius of HD 209458b, almost certainly some short-period systems have companions which force their eccentricities to nonzero values. This paper is a summary of \\citet{m1} which is dedicated to quantifying this effect.  The eccentricity of a short-period planet will only be excited as long as its (non-resonant) companion's eccentricity is non-zero. In fact, \\citet{m1} shows that the latter decays on a timescale which depends on the structure of the {\\it interior} planet, a timescale which is often shorter than the lifetime of the system. {\\it This includes Earth-mass planets in the habitable zones of some stars}. On the other hand, there exist parameter combinations ($Q$-value of the innermost planet, companion mass and semimajor axis) for which significant eccentricity in the short-period planet can be sustained for at least the age of the system, and these include systems with companion masses as low as a fraction of an Earth mass. ",
        "conclusions": ""
    },
    "0710/0710.2461_arXiv.txt": {
        "abstract": "We present a measure of the hard (2--8 keV) X-ray luminosity function (XLF) of Active Galactic Nuclei up to $z\\sim5$.  At high redshifts, the wide area coverage of the {\\em Chandra} Multiwavength Project is crucial to detect rare and luminous ($L_X>10^{44}$ erg s$^{-1}$) AGN. The inclusion of samples from deeper published surveys, such as the {\\em Chandra} Deep Fields, allows us to span the lower $L_X$ range of the XLF.  Our sample is selected from both the hard ($z<3$; $f_{2-8~{\\rm kev}}>6.3\\times10^{-16}$ erg cm$^{-2}$ s$^{-1}$) and soft ($z>3$; $f_{0.5-2.0~{\\rm kev}}>1.0\\times10^{-16}$ erg cm$^{-2}$ s$^{-1}$) energy band detections.  Within our optical magnitude limits ($r^{\\prime},i^{\\prime}<24$), we achieve an adequate level of completeness ($>50\\%$) regarding X-ray source identification (i.e., redshift).  We find that the luminosity function is similar to that found in previous X-ray surveys up to $z\\sim3$ with an evolution dependent upon both luminosity and redshift.  At $z>3$, there is a significant decline in the numbers of AGN with an evolution rate similar to that found by studies of optically-selected QSOs.  Based on our XLF, we assess the resolved fraction of the Cosmic X-ray Background, the cumulative mass density of Supermassive Black Holes (SMBHs), and the comparison of the mean accretion rate onto SMBHs and the star formation history of galaxies as a function of redshift.  A coevolution scenario up to $z\\sim2$ is plausible though at higher redshifts the accretion rate onto SMBHs drops more rapidly.  Finally, we highlight the need for better statistics of high redshift AGN at $z\\gtrsim3$, which is achievable with the upcoming {Chandra} surveys. ",
        "introduction": "Our present understanding of the evolution of accreting SMBHs over cosmic time comes from the measure of the luminosity function (i.e., the number undergoing a detectably luminous phase within a specific co-moving volume as a function of luminosity and redshift) of Active Galactic Nuclei (AGN).  Energy production through mass accretion onto SMBHs allows us to observationally identify these sites as the familiar AGN with Quasi-stellar objects (QSOs) the most luminous example.  Hence, the AGN luminosity function provides a key constraint to discern the underlying physical properties of the population (i.e., black hole mass and accretion rate distributions as a function of redshift) and thereby elucidate the mechanisms (i.e., galaxy mergers and/or self-regulated growth) that are instrumental in their formation and evolution.  To date, an enormous effort has been undertaken to measure the luminosity function over the wide range in luminosity spanned by AGN out to high redshift.  The bright end has been well established to $z\\sim5$ by optical surveys \\citep[e.g.,][]{ri06,cr04,wo03} which primarily select QSOs using a multi-color photometric criteria.  The most dramatic feature found is the rise and fall of the co-moving space density with peak activity at $z\\sim2.5$.  With an unprecedented sample of over 20,000 QSOs in the 2dF QSO Redshift Survey (2QZ), \\citet{cr04} convincingly show a systematic decrease in luminosity (pure luminosity evolution; PLE) from $z=2$ to the present, in agreement with past surveys \\citep[e.g.,][]{sc83,bsp88,hfc93}, which find very few bright QSOs in the local universe.  This fading of the luminous QSO population is thought to be due to a decrease in the mass accretion rate \\citep[e.g.,][]{ca00} that appears to be intimately related to the order-of-magnitude decline of the cosmic star formation rate from $z\\sim1$ to the present \\citep{bo98,fr99,me04b}.  The dropoff in the space density at $z>3$ \\citep{os82,wa94,sc95,fa01,wo03} may be indicative of either the detection of the onset of accretion onto young SMBHs or a high-redshift population that has been missed, possibly under a veil of obscuration \\citep[e.g.,][]{al05,ma05}. Excessive amounts of dust and gas may be ubiquitous in galaxies at early epochs due an increase in the merger rate \\citep{ka07} that induces high star formation rates \\citep[e.g.,][]{ch05}.  It has been evident for quite some time that optical surveys miss a significant fraction of the AGN population.  They fail to find the majority of AGN due to a steeply declining luminosity function with the low luminosity end severely affected by host-galaxy dilution.  Though current optical selection techniques do show considerable improvement \\citep{ri05,ji06}, they still fail to account for many low luminosity AGN.  Of equal significance, many AGN (e.g., Seyfert 2s, narrow line radio galaxies) are missed due to dust obscuration (causing their optical properties to differ from the type 1 QSOs) and can only be adequately selected in the low redshift universe \\citep{hao05} based on their highly ionized, narrow emission lines.  Fortunately, AGN can now be efficiently accounted for by selection techniques in other wave bands such as the X-ray as demonstrated in this work, radio \\citep[e.g.,][]{wa05} and infrared \\citep[e.g.,][]{po06}.  As further elaborated below, current models based on recent observations continue to attribute the bulk of the Cosmic X-ray Background (CXRB), the previously unresolved X-ray emission, to these various types of obscured AGN.  Over more than two decades, X-ray surveys have been improving and extending the known AGN luminosity function by including sources at low luminosity, with or without optical emission lines, and hidden by a dense obscuring medium.  The Extended Medium Sensitivity Survey \\citep{gi90} was one of the first surveys to measure the X-ray luminosity function (XLF thereafter) out to the QSO peak using a sample of just over 400 AGN detected by the $EINSTEIN$ Observatory. Since the survey only probed the more luminous AGN ($L_X>10^{44}$ erg s$^ {-1}$) above $z=0.3$ due to the bright flux limit ($f_{0.3-3.5~{\\rm keV}}>5\\times10^{-14}$ erg cm$^{-2}$ s$^{-1}$), it was quite understandable that \\citet{ma91} and \\citet{de92} found the XLF to behave similarly to the optical luminosity function (i.e., PLE) with a decreasing space density from $z\\sim2$ to the present.  The $ROSAT$ satellite with its increase in flux sensitivity ($f_{0.5-2.0~\\rm{keV}}> 10^{-15}$ erg cm$^{-2}$ s$^{-1}$) enabled AGN to be detected at lower luminosities, and out to higher redshifts ($z\\sim4.5$).  Surveys ranged from the wide area and shallow $ROSAT$ Bright Survey \\citep[f$_{lim}\\sim10^{-12}$ erg cm$^{-2}$ s$^{-1}$; 20000 deg$^{2}$;][]{sc00} to the deep and narrow Lockman Hole \\citep[f$_{lim}\\sim10^{-15}$ erg cm$^{-2}$ s$^{-1}$; 0.3 deg$^{2}$;][]{le01}.  With a compilation of 690 AGN from these fields and other available $ROSAT$ surveys \\citep{bow96,ap98,mc98,za98,ma00} that effectively fill in the parameter space of flux and area, \\citet{mi00} were able to resolve $\\approx$ 60-90\\% of the soft CXRB into point sources and generate a soft XLF that extended beyond the QSO peak.  They found that the XLF departed from a simple PLE model, now well known to describe X-ray selected samples as further elaborated below.  A $luminosity$-dependent $density$ evolution (LDDE thereafter) model was required due to the slower evolution rate of lower luminosity AGN compared to that of the bright QSOs.  The limited sky coverage at the faintest X-ray fluxes achievable with $ROSAT$ prevented an accurate measure of both the faint-end slope at $z\\gtrsim1$ and the overall XLF at high redshift since only a handful of AGN were identified above a redshift of 3.  In the current era of {\\em Chandra} and {\\em XMM-Newton}, X-ray surveys are now able to detect AGN and QSOs not only enshrouded by heavy obscuration but those at high redshift ($z>3$) with greatly improved statistics due to the superb spatial resolution and sensitivity between 0.5 to 10 keV of these observatories.  Previous X-ray missions such as $EINSTEIN$ and $ROSAT$ as described above were limited to the soft band which biases samples against absorption.  The $ASCA$ observatory \\citep[e.g.,][]{ak03} successfully found many nearby absorbed AGN but lacked the sensitivity to detect the fainter sources contributing most of the 2--8 keV CXRB.  With deep observations of the {\\em Chandra} Deep Field North \\citep[CDF-N;][]{br01}, Deep Field South \\citep[CDF-S;][]{ro02} and Lockman Hole \\citep{ha01}, a large fraction, between $\\sim70\\%$ \\citep{wo05} and $\\sim89\\%$ \\citep{mo03} of the hard (2--8 keV) CXRB has been resolved into point sources.  Many of the hard X-ray sources found so far arise in optically unremarkable bright galaxies \\citep[e.g.,][]{ba03b,to01}, which can contain heavily obscured AGN.  A more robust luminosity-dependent evolutionary scheme has emerged from recent measures of the XLF.  With the inclusion of absorbed AGN from {\\em Chandra} and {\\em XMM-Newton} surveys, lower luminosity AGN are clearly more prevalent at lower redshifts ($z<1$) than those of high luminosity that peak at $z\\sim2.5$.  This behavior is due to the flattening of the low luminosity slope at higher redshifts that has been well substantiated with hard (2--8 keV) X-ray selected surveys \\citep{cow03,ue03,fi03,ba05,si05b,laf05}. Using a highly complete soft (0.5--2.0 keV) band selected sample of over 1000 type 1 AGN, \\citet{ha05} has shown that this LDDE model accurately fits the data and shows a gradual shift of the peak in the co-moving space density to lower redshifts with declining luminosity.  This behavior may be evidence for the growth of lower mass black holes emerging in an ``anti-hierarchical'' or ``cosmic downsizing'' fashion while accreting near their Eddington limit \\citep{ma04,me04a}, or the embers of a fully matured SMBH population with sub-Eddington accretion rates.  The former scenario has been substantiated by \\citet{ba05} based on the optical luminosities of the galaxies hosting X-ray selected AGN, and \\citet{he04} using type 2 AGN from the SDSS with the [OIII] emission line luminosity as a proxy for the mass accretion rate and an estimate of the black hole mass from the $M-\\sigma$ relation.  Even though there has been much progress, there are remaining uncertainties in the current measure of the XLF.  (1) A significant number of X-ray sources in the recent surveys with {\\em Chandra} and {\\em XMM-Newton} are not identified.  \\citet{main05} find that the peak of the redshift distribution shifts to higher values ($z\\sim1.2-1.5$) when incorporating photometric redshifts for objects too faint for optical spectroscopy.  (2) \\citet{ba05} demonstrate that the XLF can be fit equally well by a PLE model at $z<1.2$.  These models only begin to substantially differ for low luminosity AGN at $z>1.5$ where AGN statistics are quite low with most being provided by the deep CDF-N and CDF-S observations.  New moderate depth surveys such as the Extended {\\em Chandra} Deep Field-South \\citep[E-CDF-S;][]{le05} and the Extended Groth Strip \\citep[EGS;][]{na05} will provide additional AGN at these luminosities and redshifts but await optical followup.  (3) How does the AGN population behave at redshifts above the peak ($z\\sim2.5$) of the optically-selected QSO population?  We have presented evidence \\citep{si04,si05b} for a similar evolution of luminous X-ray selected QSOs to those found in the optical surveys \\citep[e.g.,][]{ri06} with a decline in the co-moving space density at $z>3$ but these AGN are mainly type 1.  We don't expect the inclusion of luminous absorbed QSOs to drastically alter our measure of the XLF since they may be at most as numerous as the type 1 QSOs \\citep{gi07}.  Recent radio \\citep{wa05} and near infrared \\citep{br06} surveys are further supporting a strong decline in the co-moving density at high redshift.  In the present study, we measure the XLF of AGN in the hard X-ray (intrinsic 2--8 keV) band with an emphasis on reducing uncertainties at high redshift ($3<z<5$).  This paper is an extension to the preliminary results on the co-moving space density of AGN as reported by the ChaMP survey \\citep{si04,si05b}.  As previously described, these early epochs represent the emergence of the luminous QSOs and the rapid growth phase of young SMBHs.  To date, the limited numbers of X-ray selected AGN at $z>3$ have constrained current measures \\citep{laf05,ba05,ue03} to lower redshifts.  Motivated by \\citet{ba05}, we use the observed soft X-ray band for AGN selection above $z=3$ where we measure the rest-frame energies above 2 keV.  Due to the rarity of luminous high redshift AGN, such an endeavor requires a compilation of surveys that covers a wide enough area at sufficient depths.  As previously mentioned, a large dynamic range from the deep, narrow pencil beam to the wide, shallow surveys is required to measure a luminosity function that spans low and high luminosities at a range of redshifts. Currently, there are a handful of deep surveys with {\\em Chandra} (i.e., CDF-N, CDF-S) and {\\em XMM-Newton} (Lockman Hole) that have published catalogs with a fair sample of low luminosity ($42<log~L_X<44$) AGN out to $z\\sim5$.  To provide a significant sample of more luminous AGN ($log~L_X>44$), the {\\em Chandra} Multiwavelength Project (ChaMP) is carrying out a wider area survey of archived {\\em Chandra} fields and the CLASXS \\citep{ya04,st04} survey is imaging a contiguous area with nine {\\em Chandra} pointings.  The statistics of high redshift AGN are sure to improve with the anticipated results from the SWIRE/{\\em Chandra} (Wilkes et al., in preparation), XBootes \\citep{mu05}, E-CDF-S \\citep{le05}, EGS \\citep{na05}, XMM/COSMOS \\citep{ha07}, and the newly approved {\\em Chandra}/COSMOS surveys.  We organize the paper as follows: Section~\\ref{compile}, we describe the various surveys used in this analysis including X-ray sensitivity, sky area coverage, incompleteness as a function of not only X-ray flux but optical magnitude, and AGN selection.  Our method for measuring the luminosity function is presented in Section~\\ref{methods} and the results, including best-fit analytic models, in Section~\\ref{results} for all AGN types.  In Section~\\ref{text:cxrb}, we address the resolved fraction of the CXRB and any underrepresented source populations.  We directly compare our luminosity function to that of optically-selected samples in Section~\\ref{text:compare_opt}.  Based on our luminosity function, we derive in Section~\\ref{history} the accretion rate distribution as a function of redshift and the cumulative mass density of SMBHs. Section~\\ref{coevolution}, we compare the global mass accretion rate of SMBHs to the star formation history of galaxies out to $z\\sim5$.  We end in Section~\\ref{predict} with some predictions of the numbers of high redshift ($z>3$) AGN expected in new surveys that effectively enable us to extend the luminosity function to these high redshifts with accuracy.  Throughout this work, we assume H$_{\\circ}=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.7$, and $\\Omega_{\\rm{M}}=0.3$. ",
        "conclusions": "We have presented an extension of the hard (2--8 keV) X-ray luminosity function of AGN up to $z\\sim5$.  The ChaMP effectively covers a wide area (1.8 deg$^{2}$) at sufficient depths \\hbox{($f_{0.5-2.0~{\\rm keV}}\\sim10^{-15}$ erg cm$^{-2}$ s$^{-1}$)} to significantly improve the statistics of luminous ($L_{\\rm X}>10^{44}$ erg s$^{-1}$) AGN at high redshift.  In total, we have amassed a sample of 682 AGN with 31 at $z>3$.  The addition of lower luminosity AGN from the {\\em Chandra} Deep Fields is instrumental to characterize the faint end slope at $z\\gtrsim1$.  We have corrected for incompleteness (i.e., fraction of sources without a redshift) as both a function of X-ray flux and optical magnitude, as dictated by the limitations of optical spectroscopic followup.  Significant optical followup is still required to accurately account for the obscured population, especially at high redshifts.  We have measured the hard XLF with both a binned ($1/V_a$) and an unbinned method that fits an analytic model to our data using a maximum likelihood technique.  We found that the luminosity function is similar to that found in previous studies \\citep{ue03,laf05,ba05} up to $z=3$, with an evolution dependent upon luminosity.  At higher redshifts, there is a significant decline in the numbers of AGN with an evolution rate similar to that found with optically-selected QSO samples.  We further show that the strong evolution above the cutoff redshift may cause the LDDE model to underpredict the number of low luminosity AGN at $z>2$.  A PLE model with a faint-end slope dependent on redshift agrees better with the binned ($1/V_a$) data at $z\\sim2.5$ though it may overpredict the number of faint AGN at higher redshifts. We highlight the need to improve the statistics of both high redshift AGN at $z>4$ and lower luminosity (i.e., below the break in the luminosity function) AGN at $z>3$. Such improvements are feasible from our predictions of the numbers AGN that will be found in the latest {\\em Chandra} surveys (E-CDF-S, EGS and COSMOS)  Our new luminosity function accounts for $\\sim52\\%$ of the 2--8 keV Cosmic X-ray Background.  The integrated emission from these AGN give a $z=0$ mass density of SMBH of $1.64\\times10^5$ \\Msun~Mpc$^{-3}$, lower than other published values using X-ray selected AGN samples and the local value measured using a galaxy luminosity function and a bulge-velocity relation, possibly due to unaccounted optically-faint AGN ($r^{\\prime},i^{\\prime}>24$), and Compton-thick AGN.  Further, a comparison of the mean mass accretion rate of SMBHs to the star formation history of galaxies out to $z\\sim5$ shows the familiar co-evolution scheme up to $z\\sim2$ and a divergence at higher redshifts with perhaps star formation preceeding the formation of SMBHs."
    },
    "0710/0710.1181_arXiv.txt": {
        "abstract": "{We study the evolution and fate of solar composition super-massive stars in the mass range 60 --- 1000 \\ms. Our study is relevant for very massive objects observed in young stellar complexes as well as for super-massive stars that could potentially form through runaway stellar collisions.} {We predict the outcomes of stellar evolution by employing a mass-loss prescription that is consistent with the observed Hertzsprung-Russell Diagram location of the most massive stars.} {We compute a series of stellar models with an appropriately modified version of the Eggleton evolutionary code.} {We find that super-massive stars with initial masses up to 1000 \\ms\\ end their lives as objects less massive than $\\simeq 150$\\,\\ms. These objects are expected to collapse into black holes (with $M \\la 70 \\ms$) or explode as pair-instability supernovae.} {We argue that if ultraluminous X-ray sources (ULXs) contain intermediate-mass black holes, these are unlikely to be the result of runaway stellar collisions in the cores of young clusters.} ",
        "introduction": "\\label{sec:intro}  In this paper, we study the structure and evolution of very massive stars (VMS), defined as objects with masses of 60 up to 150\\ms, as well as super-massive stars (SMS) with masses in the range 150 - 1000 \\msun.  The interest in the upper limit of stellar masses and the evolution and fate of the most massive stars in the Universe was greatly boosted by the discovery of ultraluminous X-ray sources \\citep[ULX, see][ and references therein]{fabbiano89,colbert04,fabbiano03,fabbiano_x_sources06,soria_recipes06}. These objects are most commonly interpreted as binaries involving either sub-Eddington accretion onto an \\textit{intermediate mass black hole} (IMBH) with mass $\\sim (10^2 - 10^5)$\\,\\ms\\ or super-Eddington accretion onto a \\textit{stellar mass black hole} with mass $\\sim10$\\,\\ms.  In the latter case, beaming or support of super-Eddington luminosity by an accretion disk is required. Currently, the issue of the black hole masses in ULXs remains unresolved \\citep[see, e.g.,][]{fabbiano_x_sources06}.  It has been argued that IMBHs may also reside in the cores of some globular clusters \\citep[see, e.g.,][]{gerssen02,gerssen03,gebhardt05,Patruno:2006bw,feng_ulx_imbh06}, but see \\citet{baumgardt_m15_2003,baumgardt_g1_2003} for counter-arguments. The masses of the putative black holes in the cores of well-known globular clusters such as 47~Tuc and NGC~6397, are expected to be $\\sim (10^2 - 10^3)$\\,\\ms\\ \\citep{DeRijcke:2006tu}.  ULXs are observed in both spiral and elliptical galaxies, i.e. in environments with diverse metal abundances and star formation rates. In this paper, we focus on stars with solar initial composition ($X$=0.7,$Z$=0.02). This choice is partly motivated by the high metallicity of the starburst galaxy M82 -- with $Z\\simeq Z_{\\odot}$ \\citep{mcleod_m82_solar93,origlia_m82_chem04,mayya_m82_chem06} -- which contains one of the most promising candidate ULXs, M82~X-1, and that may host a black hole of intermediate mass, as argued by e.g. \\citet{ptak_m82_imbh99,kaaret_m82_imbh01,matsumoto_m82_ulx01,strohmayer_m82_imbh03}, but see \\citet{okajima_noimbh_m82} for arguments in favour of a stellar mass black hole in M82~X-1.  The current observational estimate of the upper cut-off of stellar masses is $\\sim 150$\\,\\ms\\ \\citep{massey_hunter98,weidner_kroupa_upper_m04,figer_upper05}. However, for solar chemical composition, it is expected that even such massive stars rarely produce black holes more massive than $\\simeq 20$\\,\\ms\\ \\citep{maeder92,fryer_kal_bh01}, due to copious  mass loss during both the hydrogen and helium burning stages. If ULXs and young globular clusters really harbour black holes with masses exceeding several 10\\,\\ms, the problem of their formation becomes a challenge for the theory of stellar evolution with mass loss.  Related intriguing problems involve the formation and evolution of the luminous stars found in the central parsec of the Galaxy, e.g., the S-stars and the IRS13 and IRS16 conglomerates \\citep{schodel_irs13,lu_irc16,maillard_13e}\\footnote{Note that a high stellar mass inferred from luminosity may be an artifact of unresolved binarity.  For instance, IRS16 SW turned out to be a binary with two almost identical components of $\\sim 50$\\,\\ms, \\citep{depoy_irs16,martins_irs16sw_bin06,peeples_irc16sw_bin06}.}, or in the R136 stellar complex of the 30 Doradus nebula in the Large Magellanic Cloud \\citep{walborn97,massey_hunter98}.  Returning to the problem of IMBH formation, there are currently two models in favour.  One involves the gradual accumulation of mass by accretion onto a seed black hole, which, while swallowing gas and stars in a stellar cluster, may grow to the intermediate mass range \\citep[see, e.g., ][ and references therein]{miller_imbh_form04}. The other scenario considers the collapse of an object which descends from an SMS formed by hierarchical runaway merger of ordinary stars in a young dense stellar cluster, during or after core collapse \\citep[see, e.g., ][]{pmmh99,pm_imbh_clust02,mbphm04,gfr04,pbhmm_clust04,fgr06}.  The motivation for the current study stems from the latter scenario. It assumes that stars are born with a wide distribution of masses ranging from the hydrogen-burning limit up to a maximum of about 150\\,\\msun.  More massive stars observed in the Galaxy may originate from stellar mergers.  Most of these may be the result of coalescence of components of binaries in common envelopes. The mass of binary merger products may be up to 300\\msun. \\citet{pmmh99,miller_hamilton02,pbhmm_clust04,gfr04} have demonstrated by means of detailed N-body simulations, in which simple stellar evolution was taken into account, that there is a range of initial conditions where stellar coagulation drives the mass of a star to $\\sim1000$\\,\\msun. The resulting star burns-up quite quickly, and the process of hierarchical merging in the cluster core terminates as soon as the first massive stars experience supernovae and collapse into black holes. The collision runaway process terminates as the mass loss from the explosions of massive stars in the cluster center drives the expansion of the cluster core. This scenario was used to explain the large black hole mass of M82 X-1 \\citep{pbhmm_clust04}.  The studies of hierarchical mergers take into account the possibilities of collisional mass loss, mass loss from stellar winds, and rejuvenation of merger products due to fresh hydrogen supply in collisions.  However, stellar evolution is treated rather crudely in these simulations, using extrapolations by several orders of magnitude for stars that typically have zero-age main-sequence (ZAMS) mass $\\leq 100$\\,\\ms. The possible formation of core-halo configurations and the existence of an upper stellar mass limit are generally ignored. The treatment of mass loss in stellar winds is particularly uncertain. In principle, the winds of very massive and super-massive stars may be so strong that the cluster core expands in such a dramatic way that it stops being dominated by collisions, in which case the hierarchical merger is terminated \\citep{pmmh99,vanbev_dyn05}.  We aim to refine evolutionary calculations for merger products, and as a first step, we study the evolution of solar composition VMS and SMS over the mass range 60-1000\\,$\\ms$. We construct chemically homogeneous models for these stars and confirm the existence of their upper mass limit ($\\simeq 1000$\\,\\ms), and study their evolution with mass loss through the core hydrogen and core helium-burning stages. This allows us to determine the mass and nature of the pre-supernova objects and to predict the fate of these objects.  The results of our evolutionary computations are presented and discussed in Sect.~\\ref{sec:evol}, a comparison with observations is given in Sect.~\\ref{sec:obs}, the fate of SMS is discussed in Sect.~\\ref{sec:fate}. Conclusions of our study follow in Sect.~\\ref{sec:concl}. In the Appendix, we discuss a number of individual stars in close proximity to the Humphreys-Davidson (HD) limit. ",
        "conclusions": "\\label{sec:concl}  In this paper we discuss the evolution of stars with initial masses in the range 60\\,\\msun\\, to 1001\\,\\msun. Our study was motivated by the results of the direct $N$-body simulations of dense star clusters by Portegies Zwart et al. (1999), in which a star grows by repeated collisions to well beyond 100\\,\\msun. In later studies, Portegies Zwart et al. (2004) proposed that the collision runaway in star clusters could explain the presence of a black hole of $\\apgt 600$\\,\\msun\\, in the star ULX M82~X-1 in cluster MGG11, which supposedly could have formed from a $\\apgt 1000$\\,\\msun\\, star. According to these models, the mass of the VMS increases over the stellar lifetime, starting as a homogeneous massive $\\sim 100$\\,\\msun\\, star that grows to $\\apgt 1000$\\,\\msun\\, within the core-hydrogen burning stage of evolution.  We have assumed that our stellar evolution calculations are representative for stars that grow in mass via the collision runaway process. This is not necessarily correct, as the hydrogen reservoir of the collision runaway product will continuously be replenished by repeated collisions, whereas in our simulations we start with a high-mass homogeneous model. Furthermore, rotation may play a relevant role in the evolution of these objects, which was also ignored in our study. Having noted this, we may nonetheless expect these massive objects to be subject to heavy mass loss, and we may provide meaningful results in terms of the fate of the objects under consideration.  We have confirmed previous results on the existence of an upper mass limit for chemically homogeneous stars, which is reached when the luminosity in the outermost layers of the stars approaches the Eddington luminosity, at which point gravity is unable to balance radiation pressure. For non-rotating solar composition stars, this limit is reached at about 1000\\,\\ms. We evolved the stars of 60\\,\\ms\\ to 1001\\,\\ms\\  to the end of the core helium-burning stage and calibrated the stellar mass-loss prescription adopted by enforcing the condition that no star should spend more than a few percent of its life above the HD-limit. We found that mass-loss rates and HR-diagram positions of the earliest O-type stars, and stars classified as transitional between O- and WR-stars, may be consistent with our computed tracks.  Based on our mass-loss prescription and stellar evolution calculations, we argue that the observed massive stars in the Galaxy and the Magellanic Clouds had birth masses of up to $\\simeq200$\\,\\ms.  For stars in the lower part of this mass-range, mass loss during the MS-stage leads to the exposure of layers moderately enriched in He, and they may be observed as late-type WN stars with hydrogen features in their spectra.  In the upper part of this range, stars become hydrogen-deficient WR stars already on the main-sequence.  Recently, \\citet{belkus07} published a study of the evolution of stars with masses up to 1000\\,\\ms.  Our study differs from that of Belkus et al. in two important aspects: (i) they provided a simple evolutionary recipe based on similarity theory, assuming that stars are homogeneous throughout the entire course of their evolution due to vigorous convection and stay in thermal equilibrium, whereas we present detailed stellar structure models; (ii) Belkus et al. used extrapolated mass-loss rates as predicted from radiation-driven wind theory, whilst we employed a mass-loss prescription that is consistent with the location of the most massive stars in the Hertzsprung-Russell diagram. In both studies, it is found that SMS are subject to dramatic mass loss that probably inhibits the formation of IMBH by the runaway stellar collision scenario.  Star clusters that experience core collapse before the most massive stars have left the main-sequence can develop a supermassive star via collision runaway. The mass of such an object is accumulated in subsequent collisions on a time scale of less than 3\\,Myr. The mass which can be grown in this time interval can be estimated using Eq.~(2) of \\citet{spz_etal06_imbh}\\footnote{These results were obtained using standard models for runaway mass accumulation without accounting for the possibility of vigorous mass loss as considered in the present paper.}: \\begin{equation} m_{\\rm r} \\simeq 0.01 m \\left( 1 + {t_{\\rm rl} \\over 100\\,{\\rm Myr}} \\right)^{-1/2}, \\end{equation} where $m_{\\rm m}$ is the mass of the object formed by runaway collision, $m$ is the system mass, and $t_{\\rm rl}$ is the system relaxation time. The average mass increase per collision is about $\\sim 20$\\,\\msun\\ \\citep{pmmh99} (for a more complete discussion, however, see \\citet{pm_imbh_clust02}). The supermassive star that accumulates $\\sim 1000$\\msun\\, has experienced some 45 collisions between the moment of gravothermal collapse of the cluster and the moment that the supermassive star dies. The mean time between collisions for this model is $\\aplt 6.0 \\times 10^4$\\,years, resulting in a mass accretion rate of $\\apgt 3\\times 10^{-4}$\\,\\msun/yr. This rate of mass accretion is lower than the rate of mass loss that we get for the most massive stars ($\\sim$ $3.8\\times 10^{-3}$\\,\\msun/yr), but since these numbers of mass-accretion and mass-loss rate are quite uncertain, it is not inconceivable that, in spite of the strong stellar winds, the objects may nonetheless experience a net gain in mass during their lifetimes. This conclusion was drawn by \\citet{suzuki07} who recently studied stellar evolution with mass loss of collisionally merged massive stars. They concluded that stellar winds would not inhibit the formation of SMS. We note, however, that the Suzuki et al. (2007) calculations were limited to masses up to $\\sim$ 100 \\ms\\ and their calculations did not consider the more important effect of mass loss for the final mass growth process, where the mass is supposed to increase beyond this value of 100 \\ms\\ by an order of magnitude.  In this study, we have found that a super-massive star is likely to shed most of its mass well before it experiences a supernova explosion.  In several cases the supernova progenitor still had $\\apgt 100$\\,\\msun, and such objects could possibly collapse to black holes of intermediate mass but with $M_{bh} << 1000$\\,\\ms. Our calculations suggest the possibility of the formation of objects that experience pair-instability supernovae, which would be interesting phenomena to observe. The recently discovered, very luminous supernova 2006gy may have been a PISN \\citep{smith_07,ofek_07,langer_07}, although this is still under debate.  Obviously there is the possible caveat that the quantitative difference between starting as a homogeneous high-mass star (as we considered in our paper), or growing to one  over the main-sequence lifetime by repeated collisions may be significant, but at this stage we cannot provide further insights about the consequences of these potential differences.  We therefore draw our conclusions on the calculations at hand, and we argue that  the majority of supermassive stars probably end up as black holes of $M \\lesssim 70$\\,\\ms, with the possible exception of some of them exploding as pair-instability supernovae.   With the approximate nature of our applied mass-loss algorithm in mind, and also taking into account the (also approximate) results of \\citet{belkus07}, we infer that the accuracy on the limit for black hole masses is $\\sim$ $\\pm 30$\\,\\ms. We therefore conclude that most supermassive stars end their lives as $70\\pm 30 $\\,\\msun\\, black holes."
    },
    "0710/0710.1148_arXiv.txt": {
        "abstract": "We present detailed geometry and kinematics of the inner outflow toward HL Tau observed using Near Infrared Integral Field Spectograph (NIFS) at the Gemini-North 8-m Observatory. We analyzed H$_2$ 2.122 $\\micron$ emission and [\\ion{Fe}{2}] 1.644 $\\micron$ line emission as well as the adjacent continuum observed at a $<$0\\farcs2 resolution. The H$_2$ emission shows (1) a bubble-like geometry to the northeast of the star, as briefly reported in the previous paper, and (2) faint emission in the southwest counterflow, which has been revealed through careful analysis. The emission on both sides of the star show an arc 1\\farcs0 away from the star, exhibiting a bipolar symmetry. Different brightness and morphologies in the northeast and southwest flows are attributed to absorption and obscuration of the latter by a flattened envelope and a circumstellar disk. The H$_2$ emission shows a remarkably different morphology from the collimated jet seen in [\\ion{Fe}{2}] emission. The positions of some features coincide with scattering continuum, indicating that these are associated with cavities in the dusty envelope. Such properties are similar to millimeter CO outflows, although the spatial scale of the H$_2$ outflow in our image ($\\sim$150 AU) is strikingly smaller than the mm outflows, which often extend over 1000--10000 AU scales. The position-velocity diagram of the H$_2$ and [\\ion{Fe}{2}] emission do not show any evidence for kinematic interaction between these flows. All results described above support the scenario that the jet is surrounded by an unseen wide-angled wind, which interacts with the ambient gas and produce the bipolar cavity and shocked H$_2$ emission. ",
        "introduction": "Jets or outflows are associated with young stellar objects (YSOs) at a variety of masses and evolutionary stages. There is growing evidence that these are powered by mass accretion, although the details of the flow launching mechanism are not yet clear. Understanding this issue is hampered by the angular resolution of the telescopes and interferometers, which is not yet high enough to resolve the central engine of the driving source. Instead, investigating flow geometry and kinematics close to the driving source provides useful information for understanding the relation between stellar mass accretion and outflows.  In particular, recent observations of near-infrared H$_2$ emission have provided new clues to understand the activities of mass ejection close to ($<$300 AU) low-mass protostars. Pioneering work includes echelle spectroscopy and Fabry-Perot imaging of several Class I protostars by Davis et. al. (2001, 2002). Davis et al. (2001) revealed the presence of two blueshifted components at high (50--150 km s$^{-1}$) and low velocities (5--20 km s$^{-1}$). Such a kinematic structure in H$_2$ emission is similar to the forbidden emission line regions close to T Tauri stars (see e.g., Hartigan et al. 1995). While the high-velocity component is clearly associated with a collimated jet, the nature of the low velocity component is less clear. This low velocity component could be either a magnetohydrodynamically driven gas (e.g., Pyo et al. 2003, 2004; Takami et al. 2004) or a component entrained by an unseen wide-angled wind (e.g., Pyo et al. 2006; Takami et al. 2006). Davis et al. (2002) revealed H$_2$ emission associated with a collimated jet and cavity walls toward some YSOs. Such emission is excited by shocks at a temperature of $\\sim$2000 $K$ (Takami et al. 2004, 2006; Beck et al. 2007).   Beck et al. (2007, hereafter Paper I) present the morphology and excitation of the near-infrared H$_2$ emission toward six T Tauri stars using NIFS (Near Infrared Integral Field Spectograph) at the Gemini North Telescope on Mauna Kea. NIFS is an image slicing IFU fed by Gemini's near-infrared adaptive optics system, Altair, to obtain integral field spectroscopy with a resolving power of $R \\sim 5000$. These observations revealed a variety of morphological distribution within 200 AU of T Tauri stars, and showed that H$_2$ molecules are presumably excited by shocks. The study we present shows a detailed analysis of the emission associated with one of the best studied YSOs, HL Tau. The YSO is in transient phase from a Class I protostar to T Tauri star (see Pyo et al. 2006, and references therein), and is known to host an extended collimated jet (Mundt et al. 1990), millimeter CO outflow (Monin et al. 1996; Cabrit et al. 1996), circumstellar disk (e.g., Wilner et al. 1996; Kitamura et al. 2002), and also a flattened gas envelope (Sargent \\& Beckwith 1990) which may be infalling toward the star (Hayashi et al. 1993). Paper I revealed the presence of a bubble-like morphology in H$_2$ emission extending toward the northeast.  In this paper we analyze the results for H$_2$ 1-0 S(1) emission (2.122 $\\micron$) together with the [\\ion{Fe}{2}] 1.644 $\\micron$ line and also adjacent continuum. Our results suggest the presence of an unseen wide-angled wind, interacting with the ambient gas and produce a bipolar cavity and shocked H$_2$ emission. Throughout the paper, we adopt the distance to the target of 140 pc (Elias 1978). ",
        "conclusions": "Mass ejection from low-mass YSOs is often observed in two different manners: (1) collimated jets often seen in optical-IR wavelengths, or millimeter wavelengths in some cases (see e.g., Bally et al. 2007 for review) (2) molecular bipolar outflows, which show a variety of morphologies in millimeter CO lines (see e.g., Arce et al. 2007 for review). The relation between these types of flows remain unclear. In particular, two major scenarios have been discussed over decades for driving of the molecular outflows. These are: (1) jet-driven scenario, for which the molecular outflow results from interaction between a collimated jet and ambient material (e.g., Raga \\& Cabrit 1993); and (2) wind-driven scenario, for which the molecular outflow results from an unseen wide-angled wind (e.g., Shu et al. 1991). Studies to date suggest that neither the jet-driven nor wind-driven models can explain a wide range of morphologies and kinematic properties observed in all outflows (see Cabrit et al. 1997; Arce et al. 2007).  The observed morphology and kinematics in H$_2$ emission from HL Tau are similar to millimeter CO outflows. It is striking that the H$_2$ flow observed in HL Tau is smaller than a typical spatial scale of the millimeter CO outflows. While the CO outflows often extend over 1000--10000 AU scales (see e.g., Lee et al. 2000, 2002), including HL Tau itself (Monin et al. 1996; Cabrit et al. 1996), the H$_2$ flow in HL Tau shows a spatial scale of only $\\sim$150 AU.  The H$_2$ emission shows a remarkably different geometry and kinematics from the [\\ion{Fe}{2}] emission associated with the collimated jet. Furthermore, there is no evidence for kinematic interaction between the H$_2$ flow and the collimated jet: i.e., there is no clear evidence of acceleration of deceleration of the H$_2$ and [Fe II] flows where these are overlapped in Figure 2. These results support the scenario that the jet is surrounded by an unseen wide-angled wind, which interacts with the ambient gas and produces the bipolar cavity and also shocked H$_2$ emission. The presence of such a wide-angled wind indeed agrees with proposed magneto-hydrodynamically driven wind models (see e.g., Brandford \\& Payne 1982; Uchida \\& Shibata 1985; Shang et al. 2006). Alternatively, Close et al. (1997) suggests that the cavities associated with HL Tau could be produced by a precessing jet. However, we emphasize that our results do not show any evidence for kinematic interaction between the [\\ion{Fe}{2}] emission associated with the collimated jet and the H$_2$ emission at the cavity walls. Furthermore, a typical cooling time scale of the near-infrared H$_2$ emission line regions is less than a year (see e.g.,Takami et al. 2004), and the gas associated with the jet would move only $<$0\\farcs2 in such a period (Mundt et al. 1990). Thus, the jet would have to be spatially located much closer to the H$_2$ emission line regions if it was responsible for the H$_2$ line excitation.  The arcs in H$_2$ emission 1''.0 away from the star suggests that the ejection activity is time-variable. Assuming a flow velocity of 10 km s$^{-1}$ inclination angle of 34$^\\circ$ from the plane of the sky (Mundt et al. 1990), we estimate the dynamical age of these arcs of $\\sim$70 yrs. Periodic monitoring observations spread over year-long timescales are necessary to measure their proposer motions and determine the accurate dynamical age. Such a time-variable wind may be a rather common activity associated with low-mass YSOs. Indeed, a wind bubble similar to the northwest H$_2$ emission is also observed in the neighboring YSO XZ Tau in optical emission lines (Krist et al. 1997, 1999; Coffey et al. 2004). Furthermore, some millimeter CO outflows show multiple shell structures extending over 1000--10000 AU scales (see e.g., Lee et al. 2002). Arce \\& Goodman (2001) show that episodic outflows well explain the presence of ``Hubble wedges'' (i.e., a jagged profile) in the position-velocity diagram, and also steep power-low slopes of mass-velocity relation ($dM(v)/dv \\propto v^{-\\gamma}$, $\\gamma > 2$) observed toward some outflows.  Time-variable mass mass ejection may be related to episodic mass accretion, observed in extreme cases, as FU Orionis (FUor) or EX Orionis (EXor) outbursts (see e.g., Hartmann \\& Kenyon 1996; Herbig 1989). Indeed, optical-IR absorption lines of FUors suggest the presence of an energetic disk wind (e.g., Calvet et al. 1993). Furthermore, spectroscopic monitoring of CO overtone absorption suggests that one of the FUors ejected an expanding shell of dense, low-temperature material within a few decades ago (Hartmann et al. 2004). Such a link between episodic mass ejection and accretion has also been proposed to explain the morphology of the collimated jet associated with YSOs (e.g., Dopita 1978; Reipurth 1989; Zinnecker et al. 1998)."
    },
    "0710/0710.3651_arXiv.txt": {
        "abstract": "In this paper we continue our study of CMB TE cross correlation as a source of information about primordial gravitational waves. In an accompanying paper, we considered the zero multipole method. In this paper we use Wiener filtering of the CMB TE data to remove the density perturbation contribution to the TE power spectrum. In principle this leaves only the contribution of PGWs. We examine two toy experiments (one ideal and one more realistic), to see how well they constrain PGWs using the TE power spectrum. We consider three tests applied to a combination of observational data and data sets generated by Monte Carlo simulations: (1) Signal-to-Noise test, (2) sign test, and (3) Wilcoxon rank sum test. We compare these tests with each other and with the zero multipole method. Finally, we compare the signal-to-noise ratio of TE correlation measurements first with corresponding signal-to-noise ratios for BB ground based measurements and later with current and future TE correlation space measurements. We found that an ideal TE correlation experiment limited only by cosmic variance can detect PGWs with a tensor-to-scalar ratio $r=0.3$ at $98\\%$ confidence level with the $S/N$ test, $93\\%$ confidence level with the sign test, and $80\\%$ confidence level for the Wilcoxon rank sum test. We also compare all results with corresponding results obtained using the zero multipole method. We demonstrate that to measure PGWs by their contribution to the TE cross correlation power spectrum in a realistic ground based experiment when real instrumental noise is taken into account, the tensor-to-scalar ratio, $r$, must be approximately four times larger. In the sense to detect PGWs, the zero multipole method is the best, next best is the S/N test, then the sign test, and the worst is the Wilcoxon rank sum test. ",
        "introduction": "Primordial gravitational waves (PGWs) (tensor) generate negative temperature-polarization (TE) correlation for low multipoles, while primordial density (scalar) perturbations generate positive TE correlation for low multipoles (see \\cite{crittenden94,baskaran06,grishchuk07,pmkpaper1}; and references in \\cite{pmkpaper1}). This signature can be to detect PGWs. The test based on this signature (the \\textit{zero multipole method}, see \\cite{pmkpaper1}) is useful as an insurance against false detection or as a monitor of imperfectly subtracted systematic effects.  In this paper, we analyze an alternative method. This method uses Wiener filtering to remove the contribution of density perturbations to the TE cross correlation power spectrum for small $\\ell$, leaving only the negative residual component of the TE power spectrum due to PGWs. Actually, this method can be treated a test of the (negative) contribution to the TE correlation power spectrum due to PGWs on large scales using uncertainties in the measurements consistent with the total TE power spectrum. We use Monte Carlo simulations to analyze the probability of detecting PGWs using this method.  By detection of PGWs we mean in this paper, the measurement of the parameter $r$, the ratio of the primordial tensor power spectrum, $P_t(k)$, to the primordial scalar power spectrum, $P_s(k)$, taken at wavenumber, $k_0$:  \\begin{equation} r = \\frac{P_t(k_0)}{P_s(k_0)} = \\frac{A_t}{A_s} \\end{equation} where $k_0 = 0.05 \\text{ Mpc$^{-1}$}$ (see \\cite{Smith2006}). For this paper we only consider the tensor-to-scalar ratio, $r$. All other parameters are assumed to be at their WMAP3 values (\\cite{Spergel2007WMAP}). The only other parameter which might affect the results is the tensor spectral index, $n_t$, however we assume, in this paper, that $n_t$ is very close to zero (\\cite{peiris03}). The problem of $n_t < 0$ will be considered in another paper.  The plan of this paper is the following. In Section \\ref{wienfilt}, we describe the method for detection of PGWs based on measurements of the TE power spectrum. In Section \\ref{simul}, we describe the numerical Monte Carlo simulations we tested. In Sections \\ref{montecarlo}, \\ref{signintro}, and \\ref{wilcoxonintro}, we introduce the statistical tests used to contstrain PGWs. In Section \\ref{comptests}, we compare the three different statistical tests used in this analysis. Section \\ref{resul} gives the results for the two toy experiments described in \\cite{pmkpaper1}. The only uncertainty in the first toy experiment is due to cosmic variance (\\ref{toy1res}). In the second toy experiment, along with cosmic variance, we take into account instrumental noise (\\ref{toy2res}). We present results of Monte Carlo simulations for WMAP (\\ref{wmapres}) and Planck (\\ref{planckres}). In Section \\ref{bbcomp}, we compare signal-to-noise ratio for BB and TE measurements. ",
        "conclusions": "The method described here is one in which we filter out the signal due to density perturbations, leaving only the contribution to the TE power spectrum due to PGWs. We then test the resulting TE power spectrum to see if it is negative. Three different statistical tests were used to see if there was a significant detection of PGWs. The $S/N$ test can give a value for $r$ using a comparison with Monte Carlo simulations, while the Wilcoxon rank sum test can only give an allowable range for $r$. The sign test will only tell us if $r \\not= 0$.  From \\cite{pmkpaper1}, we saw that we could detect $r=0.3$ to $3\\sigma$ with a measure of $\\ell_0$, the position where the TE power spectrum first changes sign, for the ideal experiment. Using the method discussed in this paper, we see that we are unable to make this significant of a detection. The best result was for the $S/N$ test which would give a $2.3\\sigma$ detection of $r=0.3$. To detect PGWs on the level of $3\\sigma$, the tensor-to-scalae ratio $r$ should be $r \\ge 0.4$. The sign test would give $2\\sigma$ detection for $r=0.3$ and a $3\\sigma$ detection for $r=0.45$. The Wilcoxon ranked sum test gives only a $1.2\\sigma$ detection for $r=0.3$ and a $3\\sigma$ detection for $r=0.7$. Similar results were gotten for the other three experiments tested. Thus in the sense of potential to detect PGWs, the zero multipole method is the best, next best is the $S/N$ test, then the sign test, and the worst is the Wilcoxon ranked sum test.  \\cite{baskaran06} present illustrative examples in which high $r$ is consistent with measured TT, EE, and TE correlations. The value of $r$ is so high in these examples that if PGWs with such $r$ really existed, current BB experiments would already detect PGWs. All models predict that the TE cross correlation power spectrum change sign only once for $\\ell < 100$. The fact WMAP cannot exclude several multipoles with $C_{\\ell}^{TE} > 0$ in between multipoles of $C_{\\ell}^{TE} < 0$ means that the TE cross correlation power spectrum either changes sign several times for $\\ell < 100$ or there is some instrumental noise which causes some anticorrelation measurements. Using instrumental noise consistent with WMAP, our Monte Carlo simulations give $\\Delta \\ell_0 \\approx 16$ and $\\ell_0 > 40$, which means that there is no evidence of PGWs in the TE correlation power spectrum."
    },
    "0710/0710.3698_arXiv.txt": {
        "abstract": "The physical processes that define the spine of the galaxy cluster X-ray luminosity -- temperature (L-T) relation are investigated using a large hydrodynamical simulation of the Universe. This simulation models the same volume and phases as the Millennium Simulation and has a linear extent of $500h^{-1}$Mpc. We demonstrate that mergers typically boost a cluster along but also slightly below the L-T relation. Due to this boost we expect that all of the very brightest clusters will be near the peak of a merger. Objects from near the top of the L-T relation tend to have assembled much of their mass earlier than an average halo of similar final mass. Conversely, objects from the bottom of the relation are often experiencing an ongoing or recent merger. ",
        "introduction": "Since the launch of the XMM-Newton and Chandra satellites \\citep{XMM,Chandra}, measurements of the X-ray emission from hot gas in clusters of galaxies have achieved unprecedented levels of accuracy and depth. However, the physical origin of the scaling relations between observable quantities, such as the luminosity of the X-ray emitting gas and its temperature, remain only partly understood.  There are currently a number of surveys \\citep{Romer, Schwope, Pierre} in progress with the potential to greatly expand our understanding of the processes that define correlations such as the luminosity-temperature (L-T) relation of clusters. For this potential to be realised we require a sound theoretical basis upon which to work. To this end, numerical hydrodynamical simulations have become indispensable tools and continue to grow in size and complexity \\citep{Pearce, Kay07, Faltenbacher} but they have to date lacked a sufficiently large dynamic range in mass. In this work we use a hydrodynamical model of a large volume that contains over a hundred galaxy clusters. For the first time we are able to study the evolutionary processes within a cosmological context as we have hundreds of well resolved objects spanning a large dynamic range rather than the more typical handful \\citep{Rowley} (hereafter R04), or idealised models \\citep{Ritchie,Poole}.  This paper is organised as follows: in the remainder of this section we summarise the work done to date on defining the physical processes that define the shape of the L-T relation. Then, in section 2, we give an account of the simulations we have undertaken, explain how our cluster sub-sample was selected and how the properties of these clusters were derived.  Section 3 details our results before we discuss their implications and conclude in section 4.  X-rays are chiefly emitted from the hot gas in clusters via thermal bremsstrahlung (for dark matter halos more massive than $10^{14} h^{-1} {\\rm M_\\odot}$ their temperature is typically above $2{\\rm keV}$). For such a homologous population \\citet{Kaiser} showed that simple scaling relations were expected between bulk properties such as the mass, temperature and luminosity.  Observational work subsequently found that the properties of X-ray clusters where indeed related but the slopes of the relations were not those derived by Kaiser.  \\citet{Kaiser} assumed that galaxy clusters were self-similar entities and that therefore only a single property, such as the mass, was required in order to describe the other bulk properties. Such a homology results in an L-T relation with a power-law slope of 2. However, as figure~\\ref{obsdata} demonstrates, X-ray observations of clusters with a median redshift of $\\sim 0.07$ found that the slope was closer to 3 \\citep{Markevitch,Arnaud,Wu99, Xue,Horner,Mulchaey, Osmond} and perhaps became even steeper on group scales \\citep{Helsdon}.  \\begin{figure*} \\begin{minipage}{150mm} \\begin{center} \\opt{bw}{ \\includegraphics[angle=0, width=\\textwidth]{figs/bw_obsdatafinalwkey} } \\opt{col}{ \\includegraphics[angle=0, width=\\textwidth]{figs/obsdatafinalwkey} } \\caption{ Compilation of low-redshift observed group and cluster X-ray luminosities within $r_{500}$ compared to their emission-weighted temperature. $r_{500}$ is a radius enclosing an overdensity of 500. The small points are the simulated groups and clusters used in this work. The data was taken from variously: \\citet{Markevitch,Arnaud,Wu99, Helsdon,Xue,Horner,Mulchaey,Osmond}.} \\label{obsdata} \\end{center} \\end{minipage} \\end{figure*}  Hydrodynamical simulations performed in the absence of cooling or any additional heat sources other than compression and shock heating have long been known to reproduce the self-similar hierarchy well \\citep{ENF,NFW}. Unfortunately they do not reproduce either the slope or the normalisation of the observations, producing clusters that are too bright for any given temperature, even at the bright end. Following this work, simulations with limited physics within a cosmological volume have been used in an attempt to reconcile the apparent discrepancy between theory and observation regarding the slope of the L-T relation \\citep{Pearce, Muanwong, Bialek,Borgani02}. These models showed that a simple cooling or preheating scheme was sufficient to match the simulated L-T relation to that observed at redshift zero.  More recently \\citet{Kay07} investigated the effects of feedback on the X-ray properties of clusters in hydrodynamical simulations, and demonstrated that their results were in good agreement with both the observed scaling relations and structural properties (e.g. entropy and temperature profiles), particularly for cool-core clusters.  \\citet{Balogh06} investigated the role that preheating, cooling and concentration of the halo profile can have on the scaling relations. They found that, for a realistic range of halo concentrations, the scatter generated was minimal in comparison with observed values. Variations in the cooling time of the gas in the centre of clusters could account for much of the scatter but is limited by the age of the universe and so could not explain the whole range. Finally, varying feedback from supernovae and AGN could explain the entire range, but required an order of magnitude difference in energy injection to cover the whole envelope.  Their result implies that it is processes in the cores of clusters that are primarily responsible for driving the scatter in the scaling relations. This confirms earlier work by \\citet{Fabian, Markevitch} and \\citet{McCarthy}.  \\citet{Kay07} identify the scatter with strong cool core clusters, and expect the scatter to be smaller at high redshift due to the diminished prevalence of such systems.  Nowadays, the general consensus is that the scatter is largely due to the strength of the X-ray core. In this work, which includes strong preheating, X-ray cores are absent. This allows us to study the shape of the relation without the additional complication of a large intrinsic scatter.  In this work we will use a sample of halos identified from the full simulation volume. With these we will show that because mergers tend to move clusters up the L-T relation they extend it beyond the point where the most massive, relaxed clusters are expected to lie. Thus many of the brightest, most luminous objects are ongoing or recent merger events which (as R04 point out) may be difficult to resolve observationally if they are close to the peak of the merger. In addition, because we have many closely spaced outputs we can track the motion of each of our clusters on the L-T plane, allowing us to define a ``mean merger'' vector. As this vector is not perfectly parallel to the L-T relation but rather falls slightly below it, a gentle roll in the relation naturally arises. ",
        "conclusions": "This work examines the physics that underlies the spine of the X-ray L-T relation. Due to our strong preheating prescription our halos do not have strong cores and as such do not reproduce the large scatter in the observed L-T relation, allowing us a clear window into the basic physics.  We intend to examine the physical origin of the observed scatter in future work \\citep{Gazzola} where a more physically motivated energy feedback prescription will be used and bright cooling cores are present. Preheating schemes such as the one used here are well known to accurately reproduce the slope and normalisation of the L-T relation as a whole \\citep{Pearce, Muanwong, Bialek,Borgani02}.  The model we have implemented also accurately reproduces the mean location of halos on the L-T plane at the present day but in a much larger volume than has typically been used previously.  In the real world bright cooling cores will further complicate matters but the processes discussed here which relate to the outer halo properties will underlie these, with the variation in core properties leading to a scatter about the relation discussed here.  By identifying mergers using the mass accretion histories of our objects and matching these episodes to the motion of each object on the L-T plane we have derived a ``mean merger'' vector in this plane. This vector lies largely parallel to the cluster L-T relation, as previously noted by RO4. At any particular time the mass function of the dark matter haloes present within a volume will be exponentially truncated at the high mass end above some characteristic mass scale. The large boost generated during a merger will produce points on the L-T plane appearing to lie above this characteristic mass, where there should be few objects. We therefore expect the majority of the brightest objects to be experiencing ongoing mergers, although they may be difficult to identify if they are close to their peak.  The mean merger vector we have derived is not exactly parallel to the L-T relation but rather lies slightly below it. This behaviour leads to all bar one of our low-scattered objects being obvious recent or ongoing merger events (figure~\\ref{lm}). We also note that at the high mass end the vast majority of our haloes lie below the mean relation shown on figure~\\ref{sample}. The fact that the mean merger vector lies slightly below the mean relation provides a natural explanation for the slight curvature evidenced in the simulated relation.  In summary, while it is straightforward to reproduce the observed slope and normalisation of the X-ray luminosity--temperature relation using a simple preheating scheme, such a scheme does not reproduce the observed scatter. As a preheating model includes the full underlying framework of the hierarchical build up of structure bulk mergers are not significant drivers of this scatter. Mergers can, however, produce objects that are brighter and hotter than would be expected from the cluster mass as merger events drive objects along the L-T relation towards the bright end. We find that a typical merger track does not exactly parallel the L-T relation but rather lies slightly below it, leading to a prevalence of recent or ongoing merger events on the low-scattered side of the relation. This process also leads to a slight curvature of the mean relation at the high mass end.  \\begin{figure} \\begin{center} \\opt{bw}{ \\includegraphics[angle=0, width=250pt]{figs/bw_meanmerg} } \\opt{col}{ \\includegraphics[angle=0, width=250pt]{figs/meanmerg} } \\caption{Relative motion on the L--T plane during each of the mergers defined in section 3.3. Each line represents a single merging event. The long dashed line indicates the mean L--T relation whereas the dotted line indicates the mean merger direction. } \\label{mergefit} \\end{center} \\end{figure}"
    },
    "0710/0710.1654_arXiv.txt": {
        "abstract": "We report an SMA interferometric identification of a bright submillimeter source, GOODS 850-5.  This source is one of the brightest 850 $\\mu$m sources in the GOODS-N but is extremely faint at all other wavelengths.  It is not detected in the GOODS \\emph{HST} ACS images and only shows a weak 2 $\\sigma$ signal at 1.4 GHz.  It is detected in the \\emph{Spitzer} IRAC bands and the MIPS 24 $\\mu$m band, however, with very low fluxes. We present evidence in the radio, submillimeter, mid-IR, near-IR, and optical that suggest GOODS 850-5 may be a $z>4$ galaxy. ",
        "introduction": "The Rayleigh-Jeans portion of the dust spectral energy distribution (SED) produces a strong negative $K$-correction and makes the observed submillimeter flux of a dusty galaxy almost invariant at $z>1$ to $z\\sim10$ \\citep{blain93}. This makes the submillimeter wavelength a potentially powerful probe to the high-redshift universe.  However, to date, all the identified submillimeter galaxies are at redshifts lower than 4, likely because of the limited resolution of current submillimeter instruments and the limited sensitivity of current radio instruments.  The Submillimeter Common-User Bolometer Array (SCUBA) on the single-dish James Clerk Maxwell Telescope resolved 20\\%--30\\% of the submillimeter extragalactic background light into point sources brighter than $\\sim2$ mJy at 850 $\\mu$m (\\citealp{barger99,eales99}).  Because of the low resolution ($\\sim15\\arcsec$) of SCUBA, identifications of the submillimeter sources have to assume the radio--FIR correlation in local galaxies \\citep[see, e.g.,][]{condon92} and rely on radio interferometry to pinpoint the location of the submillimeter emission. Optical spectroscopy of radio identified submillimeter sources shows that they are ultraluminous ($>10^{12}$ $L_{\\sun}$, corresponding to star formation rates of $10^2$--$10^3$ $M_{\\sun}$ yr$^{-1}$) sources at $z\\sim2$--3 and that they dominate the total star formation at this redshift \\citep{chapman05}. However, the positive $K$-correction of the radio synchrotron emission makes the radio wavelength insensitive to high-redshift galaxies, and radio observations can only identify 60\\%--70\\% of the blank-field submillimeter sources \\citep[e.g.,][]{barger00,ivison02}. The radio unidentified submillimeter sources are commonly thought to be at redshifts higher than the radio detection limit (typically $z\\sim3$--4) but there is no direct evidence for such a high-redshift radio-faint submillimeter population.  With recent development in submillimeter interferometry, it is now possible to locate submillimeter sources directly without relying on radio observations.  We have begun a program to target radio-faint submillimeter sources with the Submillimeter Array (SMA\\footnotemark[7]) to determine whether there is a high-redshift ($z>4$) tail in the redshift distribution of the submillimeter sources.  Here we report our first identification in this program, GOODS 850-5. GOODS 850-5 was detected by our SCUBA jiggle-map survey in the Great Observatories Origins Deep Surveys-North (GOODS-N, \\citealp{giavalisco04a}) with an 850 $\\mu$m flux of $12.9\\pm2.1$ mJy \\citep{wang04}. It was also detected in the combined jiggle and scan map of GOODS-N (GN 10, see \\citealp{pope06} and references therein). It is the second brightest submillimeter source in our jiggle-map catalog of the GOODS-N and has an IR luminosity of $\\sim2\\times10^{13}$ $L_{\\sun}$.  It does not have a $5 \\sigma$ radio counterpart in the deep Very Large Array (VLA) 1.4 GHz catalogs of \\citet{richards00} and \\citet{biggs06}. We obtained an unambiguous identification of GOODS 850-5 with the SMA and found that the counterpart to this source is remarkably faint at \\emph{all} other wavelengths.  All data in the optical, near-IR, mid-IR, submillimeter, and radio point to a source at $z>4$. This is important evidence for the existence of high-redshift submillimeter sources.  In this letter, we report the SMA observation and the multiwavelength photometry (\\S~\\ref{observation}) and the redshift constraints derived from the photometric data (\\S~\\ref{z_constraints}).  We discuss the implication for the star formation history in \\S~\\ref{discussion} and summarize in \\S~\\ref{summary}.  We adopt $H_0=71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.73$, and $\\Omega_M=0.27$.     \\begin{figure*} \\epsscale{1.17} \\plotone{f1.eps} \\caption{Multi-wavelength images of GOODS 850-5.  Each panel has a size of $24\\arcsec$.  North is up.  All four of the ACS (F435W, $b$; F606W, $v$; F775W, $i$; and F850LP, $z$) bands and all four of the IRAC channels (1, 2, 3, and 4) are included in the color pictures.  The color codes are labeled in the images. Grayscale images have inverse scales.  The SMA position is labeled with $2\\arcsec$ diameter circles.  The contours in the SMA image have levels of -2, 2, 4, and 6 $\\sigma$ where 1 $\\sigma$ is 1.4 mJy beam$^{-1}$. \\label{thumbnail}} \\end{figure*} ",
        "conclusions": "\\label{summary} We obtained accurate astrometry for the bright submillimeter source GOODS 850-5 with the SMA interferometer at 850 $\\mu$m. The counterpart is extremely faint in the optical, near-IR, mid-IR, and radio.  Both the radio and 24 $\\mu$m faintness of GOODS 850-5 suggests a high redshift of $z>4$.  The very red SED between 3.6 and 8.0 $\\mu$m and the nondetection in the optical and $K_s$ bands can only be fitted by stellar continuum at $z>4$.  This discovery provides important evidence that some of the radio undetected submillimeter sources are at high redshift and extends the redshift distribution of bright submillimeter sources to $z>4$. It suggests that a great fraction of star formation at high redshift is hidden from optical surveys."
    },
    "0710/0710.4137_arXiv.txt": {
        "abstract": "The median observed velocity width $\\v90$ of low-ionization species in damped Ly$\\alpha$ systems is close to $90\\kms$, with $\\sim 10\\%$ of all systems showing $\\v90\\gt210\\kms$ at $z=3$. We show that a relative shortage of such high-velocity neutral gas absorbers in state-of-the-art galaxy formation models is a fundamental problem, present both in grid-based and particle-based numerical simulations. Using a series of numerical simulations of varying resolution and box size to cover a wide range of halo masses, we demonstrate that energy from gravitational infall alone is insufficient to produce the velocity dispersion observed in damped Ly$\\alpha$ systems, nor does this dispersion arise from an implementation of star formation and feedback in our highest resolution ($\\sim45\\pc$) models, if we do not put any galactic winds into our models by hand. We argue that these numerical experiments highlight the need to separate dynamics of different components of the multiphase interstellar medium at $z=3$. ",
        "introduction": "Damped Ly-alpha absorbers (DLAs) provide us with a high-resolution probe of galaxy formation processes up to redshift $z\\sim5$, allowing us to study the neutral hydrogen distribution and kinematics of young galaxies and their surroundings on scales from several pc to several kpc, although detailed information about the spatial extent of the absorbing gas cannot be readily extracted from the observed velocity line profiles. DLAs contain a large fraction of high-redshift baryons potentially available for star formation (SF; Prochaska et al. 2005), however, their in-situ SF efficiency appears to be a factor of 20-30 lower than in the present-day galaxies of comparable gas surface density \\citep{wolfe.06}. Even though DLAs are presumably associated with fairly compact structures, they cover about a third of the sky, so there is a high probability that a line of sight to a remote quasar will contain such an absorber. Currently, just over a 1000 such systems have been studied, and their number will undoubtedly increase in coming years \\footnote{http://www.ucolick.org/$\\sim$xavier/SDSSDLA/}.  One of the unsolved long-standing puzzles in our understanding of DLAs is their large neutral gas velocity dispersion \\citep{prochaska.97}. The median value of $\\v90$, the velocity interval encompassing 90\\% of the optical depth, in absorption lines of low ions associated with cold neutral gas (such as SiII, NI, FeII, etc.) is close to $90\\kms$, and 10\\% of all systems have $\\v90\\gt210\\km/s$. Observations indicate that there is virtually no correlation between the column density $\\nhi$ of the absorber and its absorption line velocity width, with even lower column density systems at the DLA threshold $\\nhi=10^{20.3}\\cmsq$ demonstrating velocity widths up to $400\\kms$. On the other hand, a correlation is observed between the velocity widths and metallicities of the DLAs. In fact, the correlation is remarkably tight between the equivalent width of the \\ion{Si}{2}~1526 transition (a kinematic diagnostic because the line is saturated) and gas metallicity \\citep{prochaska....07} that shows a slope matching the local velocity/metallicity relation in dwarf galaxies \\citep{dekel.03}. These correlations may be a consequence of a relation between the mass of the galaxies and their metallicities \\citep{wolfe.98,ledoux....06}, or an indication that feedback from SF has a direct effect on the observed velocity dispersion \\citep{nulsen..98}.  To date all DLAs exhibit metal-line absorption \\citep{prochaska....03}, and one expects that they will all show significant MgII equivalent widths. The opposite is not true; the majority of MgII-selected absorbers are not DLAs. Unlike traditional low ions, MgII traces both cold and warm neutral material, as well as warm partially photoionized gas. Recently, \\citet{murphy....07} found a correlation between metallicity and the rest-frame MgII equivalent width, suggesting a link between kinematics and the metal-enrichment history of the absorber. They stressed that the absorption-line kinematics should be viewed separately from the host galaxy kinematics, in other words, the observed velocity widths need not scale directly with the mass of the DLA-hosting halos. \\citet{bouche....06} used the cross-correlation between 1806 MgII absorbers and $\\sim250,000$ luminous red galaxies from the Sloan Digital Sky Survey to find that the absorber halo mass is anti-correlated with the Mg II equivalent width, suggesting that at least part of the observed velocity dispersion in MgII regions is produced by supernova-driven winds and/or other feedback mechanisms which could drive cold and warm gas efficiently out of lower mass ($\\mhalo\\lt10^{11.5}\\msun$) galaxies. On the other hand, this anti-correlation could also be interpreted as a transition from cold to shock-heated gas in more massive $\\sim10^{12.5}\\msun$ halos accompanied by a drop in the cumulative cross-section of neutral clouds found in these halos \\citep{tinker.07}.  At least for now, observational data are insufficient to find a relation between the absorbing halo mass and the velocity width $\\v90$ of neutral absorption in DLAs. Therefore, we should consider at least three distinct physical mechanisms which could contribute especially to the high-end tail of the DLA velocity distribution:  1) If cold clouds accreting onto massive galaxies are dense enough to survive collisional heating while falling into the hot ($\\sim10^6{\\rm\\,K}$) virialized region, they could retain a sufficiently high neutral fraction with column densities above $10^{20.3}\\cmsq$ \\citep{mcdonald.99}. Note that for this mechanism to be viable, the accreting gas should be already sufficiently metal-rich, as even the most metal-poor DLAs have been found to have metallicities $Z\\gsim-2.8$, i.e. it would imply a model in which SF was ubiquitous at some time before $z=3$ in field galaxies of mass $10^8-10^{10}\\msun$. In this scenario a large fraction of metals observed in DLAs was produced far from massive ($\\gsim10^{12}\\msun$) halos where we see these metals in absorption. On the other hand, it is also possible that the low-metallicity accreting material could mix efficiently with the more metal-rich environment of the massive virialized halo.  2) Conversely, DLA kinematics could be dominated by outflow velocities of supernova-driven winds, i.e. we could see feedback directly \\citep{nulsen..98}. This hypothesis is indirectly supported by the high detected outflow velocities in Lyman-break galaxies (LBGs), the higher effective widths $W_{1526}$ of the SiII $1526\\AA$ transition in GRB-DLAs, and the anti-correlation between the absorbing halo mass and the effective line width in MgII absorbers. It would also naturally explain the metallicity-velocity correlation, since metals should be produced in the same regions which drive the winds.  3) The apparent inefficiency of in-situ SF in DLAs, coupled with their relatively high inferred cooling rates suggests a picture in which very compact star-forming regions, perhaps associated with LBGs, illuminate much more extended neutral gas structures \\citep{wolfe.06}. \\citet{maller...01} used a toy model of DLA absorption to show that a large covering factor of the cold gas in protogalactic clumps is consistent with observations. Moreover, numerical simulations confirm that massive halos at $z=3$ often consist of multiple compact galaxies embedded into larger neutral clouds. Combined outflows from stellar winds and SNe in these galaxies could drive large chunks of surrounding neutral material to larger galactic radii where they could pick up a higher velocity dispersion from the local galaxy group.  Although we now have absorption data on $\\sim$1000 DLAs, very few groups have attempted to model DLA kinematics in recent years. Both analytical and numerical models in the literature tend to rely on a velocity dispersion put in by hand, in part because of our lack of understanding of the physical processes of SF and feedback on sub-galactic scales, and in part due to numerical resolution limitations.  \\citet{mcdonald.99} used an analytical model combining the Press-Schechter formalism with a picture of individual spherically symmetric halos consisting of multiple absorbing clouds to show that the rate of energy dissipation corresponding to the velocity dispersion of these clouds necessary to produce the observed line profiles far exceeds the rate at which energy can be supplied to these clouds by gravitational collapse and mergers of halos. They noted that a large fraction of DLAs show multiple components in absorption profiles of the low-ionization lines, and that this observation can be reproduced by a model in which the missing dissipation energy comes from supernova explosions. In fact, the energy injection rate of $1.8\\times10^{50}{\\rm\\,ergs\\,yr^{-1}}/(10^{12}\\msun)$ reproduces well the fraction of multi-component DLAs and the overall absorption line velocity width distribution with the median value of $\\sim90\\kms$.  However, their model does not provide the physical mechanism for transferring the feedback energy to the turbulent motions in gas clouds, apart from specifying the source of this energy. In addition, it is a highly approximate model which assumes spherical exponential halos consisting of discrete neutral clouds moving with a given velocity dispersion. Moreover, the same feedback energy per unit mass is injected into all halos, independently of their mass, an assumption which probably puts too much velocity dispersion into small halos.  On the numerical front, \\citet{nagamine...07} used cosmological smoothed particle hydrodynamics (SPH) simulations coupled to a phenomenological galactic wind model to compute the rate of incidence of DLAs as a function of halo mass, galaxy apparent magnitude, and impact parameter, for a variable strength of hydrodynamical feedback. In their model, gas particles were driven out of dense SF regions by hand, by assigning a momentum in random directions. The wind mass loss rate was assumed to be twice the SF rate, and the wind carried a fixed fraction of the SN feedback energy, corresponding to the fixed wind velocities of $242\\kms$ (weak wind) and $484\\kms$ (strong wind). Depending on the strength of feedback, they found that it evacuated gas from predominantly low-mass galaxies and increased the cross-section and hence the rate of incidence of more massive galaxies. In their weak feedback model $\\sim10^{9.6}\\msun$ halos contributed the most to the cross-section, whereas with strong feedback the peak shifted to $\\sim10^{11}-10^{12}\\msun$ halos, depending on numerical resolution, as feedback became more and more efficient at removing gas from low-mass systems. Although, \\citet{nagamine...07} did not analyze the velocity width statistics, one would expect to see numerous line profiles with $\\v90\\gt100\\kms$ in their models, due the increased incidence rate of massive galaxies.  The purpose of this paper is twofold. First, we test a hypothesis that the observed kinematics is driven primarily by energy coming from gravitational infall in the process of hierarchical buildup of galaxies. This is essentially an extension of the idea put forward by \\citet{haehnelt..98} that the observed DLA line profiles are caused by a combination of random halo motions, rotation, and infall, coupled to the paradigm of several distinct modes of accretion onto growing proto-galaxies \\citep{dekel.06}. We use the approach developed in our earlier paper \\citep[][hereafter Paper I]{razoumov...06}, postprocessing high resolution adaptive mesh refinement (AMR) simulations of galaxy formation with high-angular resolution radiative transfer of UV ionizing photons. We improved our algorithm in several ways including a better treatment of hydrodynamical heating during the radiative transfer stage, as well as taking into account SF and feedback. We experimented with a number of SF models, including the standard four-criterion model of \\citet{cen.92} and the two-mode SF model of \\citet{sommer-larsen..03}, both of which can be tuned to give the observed SF rates in the absence of or with mild feedback. Including strong feedback in our experience tends to always suppress an ongoing SF, either unless relevant scales in the clumpy interstellar medium (ISM) are resolved, or interaction between the wind and the ambient medium has been turned off.  It is well known that without any special treatment the feedback energy is quickly lost away in cooling. Several prescriptions have been suggested to alleviate this problem; popular algorithms used in particle-based galaxy formation models are (1) turning off cooling of gas particles in the feedback regions \\citep{thacker.00,sommer-larsen..03,stinson.....06} and (2) using a subresolution model for the multiphase ISM \\citep{springel.03a,nagamine...07} usually coupled with kinematic feedback in which individual multiphase gas particles are driven out of the star-forming regions in random or specified directions.  Grid-based AMR simulations in which feedback energy is simply converted into the gas thermal energy seem to be quite effective in limiting SF, although this efficiency certainly depends on numerical and temporal resolution. In the grid-based simulations presented in Section~\\ref{amrModels} we do not suppress cooling or use any kinematic outflow model, in other words, we enforce a fairly conservative approximation to the role of SF in DLA kinematics, largely limited to conversion of some fraction of gas into stars.  In the second half of this paper (Sec.~\\ref{sphModels}) we examine DLA kinematics with TreeSPH models in which radiative cooling is suppressed in regions surrounding active SF sites. ",
        "conclusions": "We have performed a numerical study of DLA kinematics using high-resolution grid-based AMR simulations with moderate feedback, as well as particle-based SPH models with much more efficient feedback. With grid-based models our goal was to test the hypothesis that the observed velocity dispersion in DLAs could be explained by energy coming from gravitational collapse of cosmic structures. The broadest line profiles would then come from neutral clouds surviving infall into the massive virialized ($\\gsim10^{11}\\msun$) systems, as a $10^{11}\\msun$ halo at $z=3$ has a circular velocity already in excess of $100\\kms$. We used a series of models with a comoving volume size ranging from $4h^{-1}$ to $32h^{-1}\\mpc$ and the maximum physical resolution of $100\\pc$ to build a DLA population covering halos in the mass range $10^{8.5}-10^{13}\\msun$. Although our models can reproduce well the total incidence rate of DLAs at $z=3$, our median $\\v90$ velocities are of order $40-50\\kms$, well below the observed value of $90\\kms$. We can point out three possibilities to resolve this discrepancy:  1. {\\it More massive environments --} The velocity dispersion in individual clouds is a strong function of mass. Therefore, one could expect that including more massive halos through sampling of the rare density peaks could alleviate the velocity problem. On the other hand, all our 4-16 Mpc high resolution models feature similar velocity tails (Fig.~\\ref{colVel-boxSize}) even though the mass function in the $16\\mpc$ model L1 extends to nearly 50 times more massive halos than in the $4\\mpc$ model N1. Therefore it seems unlikely that the more massive ($\\mhalo\\gt8\\times10^{12}\\msun$) and consequently much rarer halos could make any substantial contribution to the overall statistics, even though each individual halo could have a velocity dispersion of several hundred $\\kms$.  Note that including only massive halos in larger simulation volumes might produce $f(N,X)$ at low $\\nhi$ which is too small \\citep{jedamzik.98} or too large, depending on the technique. In our simulations neglecting low-mass halos can actually raise $f(N,X)$ at low column densities, as all the gas which would otherwise be pulled in by low-mass halos stays in extended filaments which sometimes can produce column densities $\\nhi\\gt10^{20.3}\\cmsq$. On the other hand, our larger $32\\mpc$ models H0 and H1 do not have enough grid resolution to resolve these filaments, and hence produce small line densities $\\ell_{\\rm DLA}(X)$.  2. {\\it Resolution --} Our current simulations are already very expensive computationally as we refine adaptively everywhere in the volume, and the $128^3$ base grid models have in practice $300^3$ to $400^3$ resolution elements. If we start from $256^3$ or a larger base grid, many more halos will form drawing in a larger fraction of baryons at earlier redshifts, whereas inside individual galaxies we will start resolving clumpy interstellar gas. However, without running these simulations and sampling a large number of galaxies at higher resolution, it is impossible to predict reliably the effect on the velocity and column density distributions. In addition, higher resolution would allow us to better model propagation of supernova-driven winds in the clumpy interstellar gas.  The key question is the relative contribution of lower mass ($\\lsim10^{10.5}\\msun$) and more massive ($\\gsim10^{10.5}\\msun$) halos. In our current models these two populations contribute about equally to the total ``galactic'' DLA column density (Fig.~\\ref{incidenceMass}). It is possible that at higher resolution the relative contribution of the more massive halos might grow, as they will feature more numerous spatially separated components falling on the same line of sight, perhaps with neutral gas tidal tails and bridges, whereas isolated lower mass galaxies could become even more compact and pull the remaining cold neutral gas from extended filamentary structures decreasing its covering factor on the sky.  Unfortunately, the spatial distribution of absorbing clouds cannot be reconstructed from the observed line profiles in the velocity space, which would otherwise point to the minimum scale needed to resolve muli-component DLAs. \\citet{lopez...02} studied an almost dust free DLA at $z=2.33$ featuring 14 resolved metal line components. They found very similar Z/Fe ratios of low-ionization species in all 14 components, suggesting that these components could originate in clouds physically located in a small region of space, perhaps much smaller than the resolution limit of our current simulations.  3. {\\it Local microphysics --} The most intriguing possibility is that the observed velocity dispersion could arise as a result of feedback from SF. More efficient feedback could disrupt the highest column density absorbers ($\\nhi\\gsim10^{21.5}\\cmsq$) in the lower panel of Fig.~\\ref{scatter-colVel-boxSize} creating a population of low $\\nhi$ systems with high velocity widths. With such a disruption mechanism put in by hand with the galactic wind model, \\citet{nagamine...07} saw efficient gas removal from lower-mass ($\\lsim10^{10}\\msun$) halos and increased cross-sections in $\\gsim10^{11}\\msun$ halos. In view of our results, a model in which more efficient feedback is obtained without resorting to the unphysical assumptions of kinematic winds and/or suppressing cooling in feedback regions would seem particularly compelling. The key challenge here is to come up with a model which describes dynamics of individual components of the multiphase ISM separately from each other, either resolving them explicitly, or using a subresolution model in which different components are advected differently on a grid. In these models supernova winds would channel most of their energy into the lower density regions between the star-forming clouds, whereas the clouds would be more likely to survive the disruptive effect of the winds.      In Section~\\ref{sphModels} we used TreeSPH models to test the effect that the suppression of cooling has on neutral gas kinematics. We found the median velocity widths $\\v90$ of only 31 to $34\\kms$, even though that these models have been shown to expel gas efficiently from the star-forming regions at high redshifts and produce realistic galactic disks at $z=0$. In our view, such low velocity dispersion could be explained by a number of factors.  A. {\\it Resolution --} When numerical models do not resolve density inhomogeneities in the ISM, putting a certain amount of supernova feedback into the thermal energy of the gas and turning off cooling leads to an efficient conversion of this energy into the superwind expansion into a nearly uniform medium. However, the mass loading of the wind in a uniform medium is much larger than in a clumpy medium of the same average density, since in the latter case the wind will predominantly expand into the lower density voids between the clumps creating a high velocity outflow which will occasionally sweep chunks of neutral material off the edges of denser clouds. Clumps of gas entrained in winds are often observed in low-redshift outflows \\citet{rupke..05}. On the other hand, at low resolution winds will displace most gas in the galaxy moving it to somewhat higher galactic radii, producing ``puffy'' galaxies with a relatively low velocity dispersion. This is exactly what we see in SPH simulations at $45\\pc$ physical resolution. The expectation is that at higher resolution we should see multiphase outflows, with higher overall velocities in hot ionized gas, and numerous dense clouds giving rise to multi-component low-ionization absorption lines. Getting such galactic winds from first principles is notoriously difficult in numerical simulations. In this paper, we argue that the necessary condition for obtaining the correct DLA velocity dispersion is a separate dynamical treatment of different components of the multiphase interstellar medium at $z=3$.  In addition to the resolution effects, two other factors could have affected our SPH results. Here we just mention these effects briefly, as clearly their study goes beyond the reach of this paper.  B. {\\it Mean redshift of feedback --} A realistic population of galaxies at $z=0$ can be obtained in cosmological simulations invoking early ($z\\gsim 4-6$) episodes of star formation with energetic feedback \\citep{sommer-larsen..03}. In this scenario, the radial distribution of gas and its angular momentum arise as a cumulative result of successive episodes of star formation driving gas out of the galaxies. Other model parameters might result in similar star formation histories; not all parameters can be constrained uniquely from observations. One could speculate that a different star formation history, e.g. one featuring energetic feedback down to redshift $z\\sim3$, would produce a higher neutral hydrogen velocity dispersion at such a redshift. We note that from the spectra of $z\\sim3-4$ LBGs, outflow velocities of $300-400 \\kms$ are routinely inferred \\citep[e.g.,][]{pettini.......01, shapley...03}. Note, however, also that \\citet{laursen.07} were able to match both the magnitude and radial fall-off of Ly$\\alpha$ surface brightness of typical $z\\sim3$ galaxies, indicating that the spatial distribution of neutral hydrogen is correctly predicted by current SPH models.  C. {\\it AGN feedback --} It is possible that the core of each massive galaxy hosts a supermassive black hole which might have shown an AGN-type activity at some time in the past. Neither grid-based nor particle-based simulations presented in this paper include feedback from AGNs which could have a profound impact on gas kinematics in host galaxies.  Feedback from star formation is generally thought to play an important role in galaxy evolution. If galactic winds give rise to the observed neutral gas kinematics in DLAs, the same winds should affect the column density distribution. In Paper I and in this paper, we have attempted to look at both distributions through large-scale, ``blind-survey'' models, in which we simulate cosmological volumes using aggressive grid refinement for all galaxies. However, these models are very expensive to compute, especially as we get closer to resolving the typical scales of the ISM. Perhaps, a more efficient approach is to combine these large-scale numerical surveys with simulations of the ISM in isolated galaxies, whether or not accounting for mergers and cosmic infall. Fortunately, there is sufficient amount of observational data which can be used to constrain these models. Combining traditional QSO-DLA data with a closer look at the high-redshift star-forming regions with GRB-DLAs and with emission line studies should allow us to learn a lot more about young proto-galaxies in coming years."
    },
    "0710/0710.4301_arXiv.txt": {
        "abstract": "We develop a formalism for studying the dynamics of massive black hole binaries embedded in gravitationally-bound stellar cusps, and study the binary orbital decay by three-body interactions, the impact of stellar slingshots on the density profile of the inner cusp, and the properties of the ejected hypervelocity stars (HVSs). We find that the scattering of bound stars shrinks the binary orbit and increases its eccentricity more effectively than that of unbound ambient stars. Binaries with initial eccentricities $e\\gtrsim 0.3$ and/or unequal-mass companions ($M_2/M_1 \\lesssim 0.1$) can decay by three-body interactions to the gravitational wave emission regime in less than a Hubble time. The stellar cusp is significantly eroded, and cores as shallow as $\\rho\\propto r^{-0.7}$ may develop from a pre-existing singular isothermal density profile. A population of HVSs is ejected in the host galaxy halo, with a total mass $\\sim M_2$. We scale our results to the scattering of stars bound to Sgr A$^*$, the massive black hole in the Galactic Center, by an inspiraling companion of intermediate mass. Depending on binary mass ratio, eccentricity, and initial slope of the stellar cusp, a core of radius $\\sim 0.1$ pc typically forms in 1-10 Myr. On this timescale about 500-2500 HVSs are expelled with speeds sufficiently large to escape the gravitational potential of the Milky Way. ",
        "introduction": "In the standard paradigm of cosmic structure formation, it is expected  that many wide massive black hole binaries  (MBHBs) will form following the mergers  of two massive galaxies (e.g. Begelman, Blandford, \\& Rees 1980; Volonteri, Haardt, \\& Madau 2003; Mayer \\etal 2007). The binary will subsequently shrink due to stellar or gas dynamical processes and may ultimately coalesce by emitting a burst of gravitational waves. It was first proposed by Ebisuzaki, Makino, \\& Okumura (1991) that the heating of the surrounding stars by a decaying SMBH pair would create a low-density core out of a preexisting cuspy (e.g. $\\rho\\propto r^{-2}$) stellar profile. In a purely stellar background a `hard' binary shrinks by capturing the stars that pass close to the holes and ejecting them at much higher velocities, a super-elastic scattering process (`gravitational slingshot') that depletes the nuclear region. Observationally, there is evidence in early--type galaxies for a systematically different distribution of surface brightness profiles, with faint ellipticals showing steep power-law profiles (cusps), while bright ellipticals have much shallower stellar cores (e.g. Faber \\etal 1997; Ravindranath \\etal 2001). Dwarf ellipticals seem to elude this paradigm showing somewhat flat profiles, similarly to bright ones (Graham \\& Guzman 2003). Late--type spirals tend to show steep central cusps, as in the case of the Milky Way or Andromeda. Detailed N-body simulations have confirmed the cusp-disruption effect of a hardening MBHB (Makino \\& Ebisuzaki 1996; Quinlan \\& Hernquist 1997; Milosavljevic \\& Merritt 2001), while semi-analytic modelling in the framework of hierarchical structure formation theories has shown that the cumulative damage done to a stellar cusps by decaying black hole pairs may explain the observed correlation between the `mass deficit' (the mass needed to bring a flat inner density profile to a $r^{-2}$ cusp) and the mass of the nuclear black hole (Volonteri, Madau, \\& Haardt 2003).  This is the third paper of a series aimed at the detailed study of the interaction between MBHBs and their stellar environment. In Sesana, Haardt, \\& Madau (2006,2007a, hereafter Paper I and Paper II), we analyzed the three-body scatterings between a MBHB and background stars {\\it unbound} to the binary. The assumption of a fixed background breaks down once the binary has ejected most of the stars on intersecting orbits, and the extraction of energy and angular momentum from the binary can continue only if new stars can diffuse into low-angular momentum orbits (refilling the binary's phase-space ``loss cone''), or via gas processes. In galaxies with inner cores or shallow cusps, only a small fraction of the loss cone is confined within the sphere of influence of the binary, and the approximation of a background of unbound stars is reasonable. A similar argument holds also in the case of a galaxy with a steep cusp hosting a nearly equal-mass binary ($M_1\\sim M_2$). The radius of influence of such a pair, $r_{\\rm inf}=G(M_1+M_2)/(2\\sigma^2)$ where $\\sigma$ is the stellar velocity dispersion, is of the order of the binary hardening radius, $a_h=GM_2/4\\sigma^2$, and only few low-angular momentum stars have orbits with semi-major axis $\\lesssim r_{\\rm inf}$. This is not true for unequal-mass binaries, where $r_{\\rm inf} \\gg a_h$, and almost all interacting stars are bound to $M_1$.  In this paper we develop a formalism for studying the dynamics of MBHBs embedded in gravitationally-bound stellar cusps. The plan is as follows. In \\S~2 we describe our suite of three-body scattering experiments between the black hole pair and ambient bound stars. Our numerical results are used in \\S~3 and \\S~4 to construct a ``hybrid model'' of binary dynamics and investigate the orbital decay and shrinking of MBHBs in time-evolving stellar cusps. The properties of the ejected HVSs are discussed in \\S~5. The massive black hole (Sgr A$^*$) in the Galactic Center and the stars around it offer a unique opportunity to study stellar dynamics in the extreme environment around a relativistic potential. The scattering of stars bound to Sgr A$^*$ by an inspiraling intermediate-mass black hole (IMBH) is treated in \\S~6. Finally, we present a brief summary in \\S~7. ",
        "conclusions": "We have performed, for the first time, scattering experiments between a MBHB and stars drawn from a cusp bound to the primary hole. We have studied the dynamics of the pair and its orbital decay by three-body interactions, the impact of the gravitational slingshot on the stellar density profile, the properties of the ejected stellar population, and have scaled our results to the case of Sgr A$^*$. Our results can be quickly summarized as follows:  \\begin{enumerate}  \\item The extraction of the cusp binding energy causes the binary to shrink by a larger factor compared to the scattering of unbound stars. The effect is more noticeable in the case of small mass ratios $q$.  \\item The binary orbital eccentricity increases much more rapidly compared to the unbound case. The eccentricity growth is more pronounced in small mass-ratio binaries, and for shallower stellar cusps.  \\item The combined effects of enhanced orbital decay and eccentricity growth lead very unequal-mass binaries to the gravitational wave coalescence phase. The detailed fate of the pair depends on the absolute value of its mass. More massive binaries decay faster.  \\item The stellar cusp is eroded, and the total mass removed by strong three--body encounters is 2-to-4 times the mass of the secondary hole. While the mass deficit caused by dynamical friction in a merger event scales with the binary mass (and involves mostly distant stars), the mass of stars ejected from the inner cusp by highly energetic interactions scales with $M_2$. Ejection occurs in a ``burst\" lasting from few tenths to several thousands binary orbital periods, depending upon $q$.  \\item Scaled to the scattering of stars bound to Sgr A$^*$ by an inspiralling IMBH, our results imply the formation of a core of 0.1 pc in 1-10 Myrs, as well as the ejection of 500-2500 HVSs moving with speeds sufficient to escape the gravitational field of the Milky Way. In Sesana et al. (2007b) we have used the Brown et al. (2007) sample of unbound and bound HVSs together with numerical simulations of the propagation of HVSs in the Milky Way halo to constrain this ejection mechanisms, and shown that it appears to produce a spectrum of ejection velocities that is too flat compared to the observations. Future astrometric (as, e.g., {\\it GAIA}) and deep wide-field (as, e.g., {\\it LSST}) surveys should unambiguously identify the ejection mechanism of HVSs, and probe the Milky Way potential on scales as large as $200$ kpc (Gnedin et al. 2005; Yu \\& Madau 2007).  \\end{enumerate}"
    },
    "0710/0710.2900_arXiv.txt": {
        "abstract": "We consider the possibility of cluster membership for 13 planetary nebulae that are located in close proximity to open clusters lying in their lines of sight. The short lifetimes and low sample size of intermediate-mass planetary nebulae with respect to nearby open clusters conspire to reduce the probability of observing a true association. Not surprisingly, line of sight coincidences almost certainly exist for 7 of the 13 cases considered. Additional studies are advocated, however, for 6 planetary nebula/open cluster coincidences in which a physical association is not excluded by the available evidence, namely M 1-80/Berkeley 57, NGC 2438/NGC 2437, NGC 2452/NGC 2453, VBRC 2 \\& NGC 2899/IC 2488, and HeFa 1/NGC 6067. A number of additional potential associations between planetary nebulae and open clusters is tabulated for reference purposes. It is noteworthy that the strongest cases involve planetary nebulae lying in cluster coronae, a feature also found for short-period cluster Cepheids, which are themselves potential progenitors of planetary nebulae. ",
        "introduction": "For some time our knowledge of the intrinsic properties of the Galaxy's population of individual planetary nebulae has been restricted by large uncertainties in their derived distances. \\citet{zh95} suggests that the {\\it average} uncertainty in the distances cited to Galactic planetary nebulae is in the range 35-50\\%. Others are less optimistic. Such a large scatter may not be surprising, given that planetary nebulae exhibit various morphologies and span a large range in mass \\citep{kw05}.  In contrast, well-studied open clusters have distances and reddenings that are established to much greater precision, with distance uncertainties as small as 2.5\\% being possible \\citep{tb02}. Planetary nebulae established as members of open clusters are therefore a potential alternative means of calibrating their fundamental properties. With an inferred distance from cluster membership in conjunction with a planetary's angular diameter and expansion velocity, its true dimensions and age can be deduced. Cluster membership has the potential for a more direct calibration of the core mass-nebular He, C, and N abundance relationship expected in planetary nebulae as a result of single star evolution with asymptotic giant branch dredge-up \\citep{ka00,cm00}. Planetary nebulae confirmed as cluster members would enhance their importance as calibrators for the Shklovsky relation \\citep{of06} or other similar methods used to establish their distances \\citep{bl01}. On a cautionary note, significant improvement in such relationships may not be possible if the observed scatter is intrinsic.  Several factors conspire to reduce the probability of observing a planetary nebula associated with an open cluster. First, the effective sample of planetary nebulae includes a large number of objects that appear to populate the Galactic bulge \\citep{as71,zi75}, according to catalogue statistics \\citep{ko01} on their distribution along the Galactic plane (Figure 1), as well as their observed radial velocities \\citep{of06}. Potential calibrators lying in nearby open clusters are greatly reduced in number when that population is excluded, although many spatial coincidences still exist \\citep{zi75}. Associated open clusters with ages of less than $\\sim28\\times 10^6$ years ($\\log(\\tau)\\leq7.5$) are likely to be excluded, since stellar evolutionary models indicate that the end products of their evolved components are Type II supernovae explosions. Current knowledge of stellar evolution suggests that the immediate precursors of C/O white dwarfs were planetary nebulae central stars that did not undergo core carbon ignition.  In addition to a small sample size, the detection of an association between a planetary nebula and an open cluster is further hampered by the short lifetimes of planetary nebulae. Models indicate that their main-sequence progenitors were stars of $1-6.5\\;M_{\\sun}$ \\citep[e.g.,][]{we00}, with an upper limit of $\\sim8\\;M_{\\sun}$ being possible for production of Ne white dwarfs. The lifetime of the planetary nebula stage is very sensitive to initial progenitor mass and subsequent mass loss \\citep[e.g.,][]{sb96}, and varies significantly for main-sequence turnoff ages greater than $\\sim28\\times 10^6$ years, with estimates ranging from $10^3$ to $10^5$ years \\citep{sb96,ka00}. The most common age for nearby Galactic open clusters is $\\sim100\\times 10^6$ years ($\\log(\\tau) \\simeq 8$), according to the catalogue compilation of \\citet{di02} summarized in Figure 2. That corresponds to a main-sequence turnoff mass of $M_{TO}\\simeq 4\\;M_{\\sun}$. The lifetime of planetary nebulae associated with such progenitors is of order $10^3$ years, essentially instantaneous on the Galactic stage.  It is of interest to note that many planetary nebulae with massive central stars are found in the field, which is populated by the remnants of dissolved open clusters. Such clusters exceed the number of bound open clusters by a sizeable order \\citep{la03}, which suggests that, despite the short lifetimes of planetary nebulae with massive central stars, increasing the statistical sampling of possible spatial coincidences between planetaries and clusters may lead to successful detections. The usefulness of such surveys at extragalactic scales by \\citet{lr06} and \\citet{ma06} is therefore obvious: larger statistics dominate and planetary nebulae are readily discernable, as demonstrated by their success as standard candles \\citep{ja89}. The success of the Macquarie/AAO/Strasbourg H$\\alpha$ (MASH) survey \\citep{pa06} in detecting large numbers of additional Galactic planetary nebulae has also been extremely useful in revealing additional coincidences with Galactic clusters.  The discovery of planetary nebulae within globular clusters \\citep{ja97} raises a pertinent point that must be considered. If we consider $1-1.5\\;M_{\\sun}$ as a strict lower mass limit for the progenitors of planetary nebulae \\citep{kw05,of06}, then, for the ages assigned to globular clusters, corresponding to main sequence turnoffs of less than $1M_{\\sun}$, one must invoke binarity (mass transfer) to resolve the resulting discrepancy. That supports the scenario of \\citet{dm06}, \\citet{so06}, and \\citet{zi07}, who argue that a large fraction of observable planetaries may indeed stem from binary systems. Consequently, if a planetary nebula/open cluster association is established, we must be aware of the possibility that binarity might negate possible predictions for progenitor mass on the basis of the cluster's implied age from its main sequence turnoff.  In this paper we consider the possibility of cluster membership for a number of planetary nebulae that are located in close proximity to open clusters lying in their lines of sight. The often cited cases for planetary nebula/open cluster associations include the cluster and nebula designated as NGC 2818, as well as A9 in NGC 1912 (M38) and NGC 2438 in NGC 2437 (M46), but lesser known cases are also considered. ",
        "conclusions": "We have yet to establish a single physical association between a planetary nebula and an open cluster based on a correlation between their radial velocities, reddenings, and distances. However, further follow-up is indicated for a number of cases where the evidence is suggestive, namely M 1-80/Berkeley 57, NGC 2438/NGC 2437, NGC 2452/NGC 2453, VBRC 2 \\& NGC 2899/IC 2488, and HeFa 1/NGC 6067, six of the thirteen coincidences considered. Additional good cases may arise from closer examination of some of the other coincidences noted in Table 5, but most of the associated clusters are as yet unstudied, limiting further progress.  Almost all potential cluster planetary nebulae lie in cluster coronal regions, typically surrounding open clusters for which limited or no photometric data exist. The fact that very few Galactic open clusters have been studied to the extent that both their nuclear and coronal regions are examined \\citep{tu96a} only compounds the situation. Further progress requires not only new studies of our Galaxy's many unstudied clusters, but studies of their coronal regions as well. Spectroscopic observations of potentially-associated planetary nebulae would also be of value."
    },
    "0710/0710.2288_arXiv.txt": {
        "abstract": "{ VERITAS employs a multi-stage data acquisition chain that extends from the VME readout of custom 500 MS/s flash ADC electronics to the construction of telescope events and ultimately the compilation of information from each telescope into array level data. These systems provide access to the programming of the channel level triggers and the FADCs. They also ensure the proper synchronization of event information across the array and provide the first level of data quality monitoring. Additionally, the data acquisition includes features to handle the readout of special trigger types and to monitor channel scaler rates.  In this paper we describe the software and hardware components of the systems and the protocols used to communicate between the VME, telescope, and array levels. We also discuss the performance of the data acquisition for array operations. }  \\begin{document} ",
        "introduction": "VERITAS \\cite{icrc07:Veritas} is an array of 4 imaging Cherenkov telescopes designed to record images of gamma rays impacting the atmosphere. Photo-multiplier signals accompanying an image trigger \\cite{icrc07:VeritasL3} are processed using a 3-tiered data acquisition system that operates at the VME crate, telescope, and array levels.  The VME data acquisition guides the control and read out of the electronics channels.  At the telescope level, the Event Builder combines data from each VME crate into events.  At the array level, the Harvester collects events from each telescope and the array trigger and forms the final data product. Figure \\ref{fig:daqdiag} shows a schematic of the data flow and communication between the processes.  \\begin{figure*} \\begin{center} \\includegraphics [width=0.8\\textwidth]{icrc1166_fig1.eps} \\end{center} \\caption{Schematic of the data transfer and communication/control relationships for the data acquisition systems.  Thick black arrows indicate data transfer paths and the protocols used. Thick unfilled arrows indicate communication lines. The analog signals passed between the array trigger and the VME acquisition are included as thin arrows. The dashed line encloses processes repeated for each VME crate and the dash-dotted line encloses those repeated for each telescope. } \\label{fig:daqdiag} \\end{figure*}  \\subsection{The VME Data Acquisition}  The VME Data Acquisition (VDAQ) serves as the interface to five VME crates that participate in the digitization of the PMT signals and the channel-level triggering for each telescope. Four of these contain 500 MS/s flash ADC modules with a clock-trigger module \\cite{icrc03:VeritasFADC} and a fifth serves as an auxiliary crate housing a specialized clock-trigger module and a GPS clock (TTM637VME).  The FADC electronics and the constant fraction discriminators (CFDs)\\cite{icrc03:VeritasCFD} that produce individual channel triggers are housed on 10-channel 9U VME modules.  The required complement of FADC channels to accomodate the 499-pixel camera are distributed among 4 crates of 12 or 13 FADC modules.  Each crate is controlled by a VMIVME 7807 Intel Pentium M 1.8 GHz single board computer running Linux.  An additional Dolphin PCI mezzanine card provides a communication link using the ANSI/IEEE 1596-1992 Scaleable Coherent Interface (SCI) standard.  Each of the VME crates and the telescope Event Builder are connected as nodes of a low-latency, high-throughput network. The configuration used in VERITAS achieves transfer rates of 50 MBytes/s. Data transfers and general communication between VDAQ and the telescope Event Builder are conducted via this interface. Each crate can be accessed via GBit ethernet for starting and halting the acquisition process.  A 24-sample (48 ns) FADC trace translates into an event fragment size of 3880 bytes for a crate of 13 modules (130 channels). Event fragments are collected and buffered until the accumulated size reaches 8 MBytes ($\\sim$2200 events for FADC crates).  The buffer is then shipped to the telescope Event Builder.  The transfers from the VME crates happen asynchronously for crates containing different numbers of modules. The typical data rate for a crate is $\\sim$780 kBytes/s, which is well below the maximum rates permitted by the SCI and the Event Builder.  The chief limitation on the system throughput arises from the data transfer rate over the VME backplane.  VDAQ is an event-driven process while the telescope and array acquisition processes are buffered.  The telescope deadtime is incurred only at the VME level and is dominated by the size of the crate event fragments. For an array trigger rate of 200 Hz, the deadtime at an individual telescope is about 8.5\\%.  {\\bf Hardware Configuration.}  Upon initiation, the VDAQ process for each crate generates a physical map of the modules present by type and, if applicable, a unique board identifier.  Each crate provides this map to the Event Builder and uses it internally to procure hardware configuration parameters for each FADC channel and CFD. A variety of programmable features of the FADCs and CFDs are configured \\cite{icrc03:VeritasFADC,icrc03:VeritasCFD}. While some parameters are unique to an individual physical component and must be associated by a uniqe identifier, others are properties of the telescope and pixel to which a channel is connected.  The configuration settings and component mappings are stored in a MySQL database. This solution suits the distributed nature of the crate acquisition and accomodates the mappings required to properly configure the hardware.  An additional benefit is an accurate log of the mappings and settings used.  {\\bf Trigger Synchronization.} After initialization and configuration, the acquisition processes await signals from the array trigger or commands from the Event Builder.  The clock-trigger module in each crate generates a busy level during FADC reads when triggers cannot be accepted. To keep the five-crate system aligned, the busy levels are combined into a single telescope busy level and applied locally as a veto to incoming triggers.  This signal is sent to the array trigger \\cite{icrc07:VeritasL3} to indicate when the telescope cannot accept triggers.  {\\bf Trigger Type Handling.} Upon receipt of an array trigger, each crate is passed a serialized event mask that contains an event number and a trigger type code.  The mask is encoded into the event fragment via the clock-trigger boards to allow sychronization by the Event Builder. The trigger code is read by VDAQ and can be used to indicate special requests for FADC functions from the array trigger. Out-of-time reads of the FADC memory buffer are used to assess pedestal values periodically during observations.  Upon receipt of a pedestal trigger code from the array trigger, VDAQ invokes a dedicated hardware command in the FADCs.  Pedestal events are included in the data as normal events distinguished by their type code.   {\\bf Trigger Rate Measurements.} VDAQ accesses CFD trigger rates via the FADC modules.  Scaler counts for each channel are read every 400 events, about once every 2 seconds; this is not often enough to impact the deadtime significantly.  The CFD scalers are packed as a specially tagged event and shipped to the Event Builder as part of the regular SCI transfer.  Scaler reads are included during normal observations to provide direct diagnostics of channel-level triggering.  \\subsection{The Telescope Event Builder} Each VERITAS telescope has a dedicated Telescope Event Builder which is responsible for combining the event fragments from each of the five VME crates to produce telescope events; these are then written to local disk and sent via GBit ethernet to the Harvester system. The Telescope Event-Building system is a Dual Intel Xeon Server machine running linux and using a local RAID array. Communications with and control of the Event Builder program is achieved through use of CORBA (Common Object Request Broker Architecture, specifically OmniORB). The Event Building Software is written in C++ and is fully multithreaded, typically containing five threads whilst running: Communications, SCI Buffer Acquisition and Parsing, Event Building, Disk Writer, Network Writer.  At the start of each night, the Event Builder queries each VDAQ crate for a map of present VME modules and then dynamically configures itself. Data is buffered at each of the VDAQ machines and transferred in blocks to the Event Builder via the SCI system. The Event Builder parses these memory blocks and extracts pieces of individual events, tagged by unique event numbers, which are then stored in memory. When all of the pieces of an event have arrived, the telescope event is built and buffered. Once roughly 160 kBytes of telescope events have been accumulated, the events are then dispatched to the ``Consumers''. The Consumers are processes that receive built data buffers. They utilize a common architecture and, at run-time, any number of consumers may be registered. Typically only two are; the Disk Writer and Network Writer, but an additional Data Integrity Monitor consumer may also be used.  It is estimated that the throughput of the Telescope Event-Building system on a dual 2.4 GHz Xeon server is approximately 12 MBytes/s. In addition to receiving actual event data, the Telescope Event Builders periodically receive CFD scaler data as described in the previous section. These data are not transferred with the event data, but simply stored and made available via CORBA to any system that requests the information.   \\subsection{The Harvester}  VERITAS back-end data acquisition is the task of the Harvester -- a single eight-core machine that collects data from all telescopes in real-time in addition to a stream of meta-data from the L3 trigger.  The Harvester accomplishes four tasks:  {\\bf Storing data.}  The Harvester saves all data streams to its fast RAID in real time.  The current strategy for real-time data storage is to create a separate file for each telescope; this makes it possible to handle the separate telescope streams in parallel with minimal interaction.  {\\bf Combining data.}  In addition to saving the data, the Harvester combines it in real time into \\emph{array events}.  In this way, real time diagnostics can see a ``big picture'' of what the array is doing, rather than having to deal with telescopes individually.  {\\bf Diagnostics.}  The Harvester runs a variety of real-time diagnostics -- ranging from sanity checks to see if telescopes read out when they were supposed to, to a complete high-performance stereo analysis package that serves as the VERITAS real-time quicklook.  During a run, the observer has up-to-the-second knowledge of the performance of the system.  {\\bf Creating the final product.}  After a run completes, the Harvester immediately starts combining the data streams from the run to create a single file using the VERITAS Bank File (VBF) data format.  \\smallskip VBF groups telescope events together, such that given an event number, the user has immediate access to the corresponding events from all telescopes, in addition to meta-data from the array trigger.  Thus, the analysis does not need to correlate stereo events; this task is already accomplished by the data acquisition system.  VBF has been designed for portability, high performance reading and writing, compactness, and extensibility.  High performance access and compactness are achieved using a custom scheme for compressing FADC samples based on picking a different number of bits-per-sample depending on the dynamic range of each particular trace.  This scheme overwhelmingly outperforms gzip and bzip2 in reading and writing times while reaching similar compactness. See Table~\\ref{vbfperf} for performance measurements  \\begin{table} \\begin{center} \\begin{tabular}{|l|c|c|} \\hline VBF Variant & Space Usage & Read Time \\\\ \\hline \\hline Uncompressed                & 100\\% & 100\\% \\\\ % w/ Gzip                        & 42\\%  & 114\\% \\\\ % w/ Bzip2                       & 35\\%  & 514\\% \\\\ \\hline Compressed$^*$ & 38\\%  & 64\\% \\\\ % Comp. w/ Gzip  & 32\\%  & 93\\% \\\\ % Comp. w/ Bzip2 & 30\\%  & 471\\% \\\\ \\hline \\end{tabular} \\end{center} \\caption{Comparison of compression schemes normalized to the uncompressed VBF case. performed on a Pentium 4 with 2 GB RAM, an eight-way SCSI RAID-0, and Linux 2.4.18. For bzip2, we used a block size of 100,000 bytes. $^*$VERITAS uses custom compression alone. } \\label{vbfperf} \\end{table}   The observer interacts with the Harvester using the VERITAS array control software, as well as a suite of GUIs designed to view the results of quicklook analysis. The stereo analysis performed by the VERITAS quicklook system is capable of a very high level of performance -- both in time and in sensitivity.  A typical twenty minute run takes two minutes to analyze using quicklook. Further, the quicklook analysis results are comparable to offline analysis packages.  Thus, we have confidence that if a bright source appeared in our field of view during observations, quicklook would have as good of a chance of seeing it as a manually performed offline analysis. ",
        "conclusions": "The VERITAS data acquisition systems combine a variety of hardware and software resources to achieve efficient and reliable operation from the reading of FADC data to the building and storage of event data. Diagnostic information is available at all levels and real-time analsyis is performed to ensure data quality. The system has proven highly adaptable and meets the needs of a variety of configuration and calibration tasks. All systems operate within design parameters and leave room for the exploration of low-threshold regimes.  \\subsection*"
    },
    "0710/0710.2241_arXiv.txt": {
        "abstract": "{} {We present a uniform catalog of the images and radial profiles of the temperature, abundance, and brightness for 70 clusters of galaxies observed by {\\it XMM-Newton}.} {We use a new ``first principles'' approach to the modeling and removal of the background components; the quiescent particle background, the cosmic diffuse emission, the soft proton contamination, and the solar wind charge exchange emission. Each of the background components demonstrate significant spectral variability, several have spatial distributions that are not described by the photon vignetting function, and all except for the cosmic diffuse emission are temporally variable. Because these backgrounds strongly affect the analysis of low surface brightness objects, we provide a detailed description our methods of identification, characterization, and removal.} {We have applied these methods to a large collection of XMM-Newton observations of clusters of galaxies and present the resulting catalog. We find significant systematic differences between the {\\it Chandra} {\\rm and} {\\it XMM-Newton} {\\rm temperatures.}} {} ",
        "introduction": "Clusters of galaxies are the largest and most massive collapsed objects in the universe, and as such they are sensitive probes of the history of structure formation. While first discovered in the optical band in the 1930s \\citep[for a review see][]{b97}, in some ways the name is a misrepresentation since most of the baryons and metals are in the diffuse hot X-ray emitting intercluster medium and not in the galaxies.  Clusters are fundamentally ``X-ray objects'' as it is this energy range where this preponderance of the baryons is visible. Studies of cluster evolution can place strong constraints on all theories of large scale structure and determine precise values for many of the cosmological parameters. As opposed to galaxies, clusters probably retain all the enriched material created in them, and being essentially closed boxes they provide a record of nucleosynthesis in the universe. Thus measurement of the elemental abundances and their evolution with redshift provides fundamental data for the origin of the elements. The distribution of the elements in clusters reveals how the metals moved from stellar systems into the IGM. Clusters should be fair samples of the universe and studies of their mass and their baryon fraction should reveal the gross properties of the universe as a whole. Since most of the baryons are in the gaseous phase and clusters are dark-matter dominated, the detailed physics of cooling and star formation are much less important than in galaxies. For this reason, clusters are much more amenable to detailed simulation than galaxies or other systems in which star formation has been a dominant process.   Clusters are luminous, extended X-ray sources and are visible out to high redshifts with current technology.  The virial temperature of most groups and clusters corresponds to $T\\sim2-100\\times10^6$~K ($kT\\sim0.2-10$~keV, velocity dispersions of $180-1200$~km~s$^{-1}$), and while lower mass systems certainly exist we usually call them galaxies.  Most of the baryons in groups and clusters of galaxies lie in the hot X-ray emitting gas that is in  rough virial equilibrium with the dark matter potential well \\citep[the ratio of gas to stellar mass is $\\sim2-10:1$,][]{asf01}. This gas is enriched in heavy elements \\citep{mea78} and it thus preserves a record of the entire history of stellar evolution in these systems.  The presence of heavy elements is revealed by line emission from H and He-like transitions as well as L-shell transitions of the abundant elements.  Most clusters are too hot to have significant ($>2$~eV equivalent width) line emission from C or N, although cooler groups may have detectable emission from these elements. However, all abundant elements with $z>8$ (oxygen) have strong lines from H and He-like states in the X-ray band and their abundances can be well determined.  Clusters of galaxies were discovered as X-ray sources in the late 1960's (see \\citep[for a historical review see][]{m02} and large samples were first obtained with the {\\it Uhuru} satellite in the early 1970's \\citep{jf78}. Large samples of X-ray spectra and images were first obtained in the late 1970's with the {\\it HEAO} satellites \\citep[for an early review see][]{jf84}.  The early 1990's brought large samples of high quality images with the {\\it ROSAT} satellite and good quality spectra with {\\it ASCA} and {\\it Beppo-SAX}. In the last few years there has been an enormous increase in the capabilities of X-ray instrumentation due to the launch and operation of {\\it Chandra} and {\\it XMM-Newton}. Both {\\it Chandra} and {\\it XMM-Newton} can find and identify clusters out to $z>1.2$ and their morphologies can be clearly discerned to $z>0.8$. Their temperatures can be measured to $z\\sim1.2$ and {\\it XMM-Newton} can determine their overall chemical abundances to $z\\sim1$ with a sufficiently long exposure.  For example, a cluster at $z=1.15$ has recently had its temperature and abundance well constrained by a 1~Ms {\\it XMM-Newton} exposure \\citep{hea04}.  The temperature and abundance profiles of clusters out to redshifts of $z\\sim0.8$ can be measured and large samples of X-ray selected clusters can be derived.  {\\it Chandra} can observe correlated radio/X-ray structure out to $z>0.1$ and has discovered internal structure in clusters. The {\\it XMM-Newton} grating spectra can determine accurate abundances for the central regions of clusters in a model independent fashion for oxygen, neon, magnesium, iron, and silicon. Despite the stunning successes of the {\\it Chandra/XMM-Newton} era, clusters have not yet fulfilled their promise as a cosmological Rosetta stone; the most important tests of cluster theory require measurements of cluster properties to large radii ($R\\sim R_{virial}$) which is observationally difficult. The lack of consensus among the recent X-ray missions about, for example, temperature profiles, is a large stumbling block in the use of clusters for cosmological purposes.  \\subsection{Temperature Structure of Clusters}  As discussed in detail by \\citet{e03}, we now have a detailed understanding of the formation of the dark matter structure for clusters of galaxies. If gravity has been the only important physical effect since the formation, then the gas should be in rough hydrostatic equilibrium and its density and temperature structure should provide a detailed measurement of the dark matter distribution in the cluster. Recent theoretical work has also taken into account other processes, such as cooling, which can be important. The fundamental form of the \\citet{nfw97} dark matter potential and the assumption that the fraction of the total mass that is in gas is constant with radius results in a prediction that the cluster gas should have a declining temperature profile at a sufficiently large distance from the center (in $R/R_{viral}$ units), both from analytic \\citep{ks01} and numerical modeling \\citep{lea02}. The size of the temperature drop in the outer regions is predicted to be roughly a factor of 2 by $R/R_{viral}\\sim0.5$.  Although some observational results appear consistent with the theoretical predictions \\citep[in particular,the {\\it ASCA} results of ][]{mea98}, many others do not, and considerable controversy exists. Much of the uncertainty of the pre-{\\it Chandra}/pre-{\\it XMM-Newton} data arises from insufficient spectral and spatial resolution and the resultant difficulties in removing backgrounds, modeling the spectra, and interpreting the results. For example, the {\\it ASCA} results of \\citet{mea98} were consistent with a decline in temperature with radius, while the analysis of a similar sample of clusters by \\citet{kea99}, \\citet{wb00}, and \\citet{w00} revealed a large number of isothermal clusters.  Similar results were obtained from {\\it Beppo-SAX}, with \\citet{dgea99} finding temperature gradients and \\citet{ib00} finding isothermality. Simultaneous analysis of the higher angular resolution {\\it ROSAT} data with the {\\it ASCA} data did not resolve the issue; \\citet{fad01} finding gradients and \\citet{ibe99} isothermal profiles. The bulk of the problem with interpreting {\\it ASCA} results is the analysis of impact of the PSF on the profile \\citep{ibe99}.  {\\it XMM-Newton} and {\\it Chandra} have significantly better spectral and angular resolution than the previous generation of missions and might be expected to resolve the previous controversies. The recent {\\it Chandra} results of \\citet{vea06} show a temperature profile in good agreement with the gradients seen by \\citet{mea98} results and predicted by the standard theory. Analysis of samples of cooling flow clusters with {\\it XMM-Newton} \\citep{pea05,app05,pea07} are also mostly consistent with the \\citet{mea98} results.  However flatter, more isothermal profiles have also been found in both {\\it Chandra} and {\\it XMM-Newton} observations \\citep{asf01,kea04,app05}. Despite some early difficulties \\citep[e.g.,][]{dea06}, the {\\it Chandra} and {\\it XMM-Newton} calibrations have stabilized but agreement between the two great observatories is not assured \\citep[e.g,][]{vea06}.  The difference in the PSF between the two instruments as well as different methods of background subtraction often make direct comparison difficult. Further, an agreement between {\\it Chandra} and {\\it XMM-Newton} would not entirely resolve the problem; the smaller FOV of current instruments have led to observation of a somewhat higher redshift sample than observed by the previous generation of instruments, suggesting that part of the difference between the {\\it XMM-Newton}/{\\it Chandra} results and the {\\it ASCA}/{\\it ROSAT}/{\\it Beppo-SAX} results may be due to a real difference between clusters at lower and higher redshifts.  The measurement of the cluster mass function can provide a sensitive cosmological test but is sensitive, in turn, to the parameters that are directly measurable, and especially to the observed quantities at large radius. Recent simulations show that cluster temperature profiles decline with radius but less rapidly than is shown by previous X-ray analysis \\citep[e.g.,][]{ktea04}. Since the total mass of the cluster is quite sensitive to the measured temperature profile \\citep{rea06}, particularly at large radii, these systematic differences lead to significant uncertainties in the cosmological constraints. Thus, there is an urgent need to understand the temperature profiles of clusters at large radii and to understand the source of the systematic differences observed in the literature.  In this paper we consider a large sample of clusters observed with the {\\it XMM-Newton} observatory and derive temperature, density and abundance profiles for many of these systems out to near the virial radius. We present a new technique that should provide more accurate background subtraction at large radii, and are careful to correct for the effect of the finite {\\it XMM-Newton} PSF.  A comparison of our measurements with {\\it Chandra} measurements of the same clusters shows a simple systematic difference between the two observatories. Although we have not yet determined the source of that difference, resolution of this relatively well defined issue should significantly reduce the uncertainties in cluster cosmology.  \\subsection{Analysis of Extended Sources}  The analysis of extended sources in X-ray astronomy is typically problematic and quite often very complex.  This is particularly true for objects which subtend the entire field of view (FOV) of the observing instrument such as nearby galaxies, relatively nearby clusters of galaxies, many regions of galactic emission, and of course the cosmic diffuse background.  Even with objects smaller than the FOV, quite often the simple subtraction of a nearby ``background'' region from the same data set is inappropriate due to spectral and spatial variations in the internal background and angular variations in the cosmic background.  The use of deep ``blank sky'' observations can also be inappropriate due to the same considerations, as well as the probability that many background components are temporally varying.  Because of the temporal variation of the background and the angular variation of the cosmic background, the average of the blank-sky data, even after normalization, may not match the conditions of a specific observation of interest, and so may yield an incorrect result.  While the cores of many clusters are relatively bright in X-rays so the treatment of the background is not such a significant consideration, at the edges of clusters it is absolutely critical.  Clusters fade gently into the backgrounds at large radii, therefore improving the modeling of the backgrounds extends the reliable radial range for the determination of cluster parameters.  Critical to compensating for the various background components by filtering, subtraction, or modeling is a basic understanding of their origin and effects on the detectors.  Unfortunately this usually takes a considerable amount of time to develop, which is why useful methods for a specific observatory become available to the general community only years into the mission.  Even then, the methods will continue to evolve with greater understanding of the various background components and their temporal evolution, and the operation of the instruments.  In addition, the efforts are quite often undertaken by individuals who are not project personnel, but whose scientific interests require the improved analysis methods.  This is certainly true of the {\\it XMM-Newton} mission and observations using the EPIC instruments.  Several groups have presented methods and published scientific results based upon them \\citep{aea01,rp03,nml05,dlm04}. As opposed to these methods which derive background spectra from normalized blank-sky observations, this paper presents the details of a method based as much as possible on the specific understanding of the individual background components.  This method was used successfully in the paper identifying the solar wind charge exchange emission in the {\\it XMM-Newton} observation of the Hubble Deep Field North \\citep{sck04}.  Section~\\ref{sec:instrumentation} of this paper gives a short description of the {\\it XMM-Newton} observatory, Sect.~\\ref{sec:components} discusses the various background components and the suggested methods used to compensate for them, Sect.~\\ref{sec:example} demonstrates the data reduction method using the observation of \\object{Abell 1795}. Sect.~\\ref{sec:catalog} applies the methods to the determination of the temperature, abundance, and flux radial profiles of a catalog of 70 clusters of galaxies and presents the results, and Sect.~\\ref{sec:conclusions} discusses the conclusions.  Note that the detailed discussion of the science derived from these observations is deferred to Paper~II.  Currently the specific method and software package discussed here are only applicable to EPIC MOS data. Although the MOS and pn experience the same backgrounds, the physical difference between the two detectors (readout rates, fraction of unexposed pixels, etc.) make analysis of the pn background somewhat more difficult than that of the MOS. However, the analysis methods described here are being extended to the pn. ",
        "conclusions": "\\label{sec:conclusions}  In this paper we have outlined a robust and reliable method for analyzing extended X-ray sources observed with the {\\it XMM-Newton} EPIC MOS detectors.  The method combines screening of the data for periods of background enhancements (most notably the soft proton contamination), detailed modeling of the particle background spectrum, and the determination of other background components in the spectral fitting process (residual SP contamination, fluorescent particle background lines, and the cosmic background).  We have demonstrated our method with the bulk processing of the observation of 70 clusters of galaxies.  Comparison of the results for two separate observations of \\object{Abell 1835}, \\object{S\\'ersic 159-3}, and  \\object{Perseus} show good agreement between their fitted temperatures.  However, comparison of our results with the {\\it Chandra} results of \\citet{vea05} for the overlapping subset of clusters shows a significant discrepancy for higher temperature clusters.  The sense of this discrepancy is that the higher the fitted temperature, the greater the likelihood that {\\it Chandra} will find a higher temperature than {\\it XMM-Newton}.  The differences can be over 1~keV at $7-8$~keV.  This effect can increase the apparent temperature gradient in the outer annuli of clusters in {\\it Chandra} data.  While the detailed scientific analysis and discussion of these results are deferred to Paper~II, a few aspects are clear from plots of the entire data set.  For the combined plots, the radii of the annuli have been scaled to the $R_{500}$ value in the same manner as the individual plots (Sect.~\\ref{sec:catalog}).  Figs.~\\ref{fig:temp}, \\ref{fig:abund}, and \\ref{fig:flux} show the cumulative plots for the temperature, abundance, and flux, respectively.  Again, both the temperature and flux have also been normalized to the values in the range 5\\% -- 30\\% of $R_{500}$.  In addition, only points where the fitted values are three times the fitted uncertainty are plotted.  \\begin{figure} \\centering{\\includegraphics[angle=0,width=8.5cm]{7930-18-temp.pdf}} \\caption{Scaled temperature radial profiles for all of the analyzed clusters. \\label{fig:temp}} \\end{figure}  Inspection of Fig.~\\ref{fig:temp} shows, as seen before \\citep[e.g.][]{pea07,app05,vea06}, a wide variety of temperature profiles inside 5\\% of R(500). Most of these can be characterized by a temperature drop in the center as has long been seen in cooling-flow clusters.  However our single phase analysis may produce results slightly different than more detailed analysis. Over the range from $0.05-0.2 R_{500}$ the clusters are isothermal to better than 5\\%. Beyond $\\sim0.2 R_{500}$ a significant fraction of the clusters (Paper~II) show temperature drops, but they are not all self-similar. However a significant fraction of the clusters are relatively isothermal out to the largest radii measurable.  As noted by \\citet{app05}, many of the clusters show a self-similar surface brightness profile (Fig.~\\ref{fig:flux}).  Inside of $\\sim0.03 R_{500}$ there is significant scatter in the profile. With respect to the overall abundance, as was noted for {\\it ASCA} spectra of clusters by \\citet{fad01} and later for many {\\it XMM-Newton} and {\\it Chandra} spectra \\citep{mea07} there is, in a significant fraction of the clusters, an abundance increase in the center. However outside of the central $\\sim0.05 R_{500}$ there is little evidence for an abundance gradient and all the clusters are very close to the average value of $A=0.3$ on the \\citet{ag89} abundance scale (Fig.~\\ref{fig:abund}).  Detailed analysis of these results will appear in Paper~II.  \\begin{figure} \\centering{\\includegraphics[angle=0,width=8.5cm]{7930-19-abund.pdf}} \\caption{Abundance radial profiles for all of the analyzed clusters. \\label{fig:abund}} \\end{figure}  \\begin{figure} \\centering{\\includegraphics[angle=0,width=8.5cm]{7930-20-flux.pdf}} \\caption{Scaled flux radial profiles for all of the analyzed clusters. \\label{fig:flux}} \\end{figure}"
    },
    "0710/0710.0558_arXiv.txt": {
        "abstract": "Deep F555W and F814W {\\it Hubble Space Telescope\\/} ACS images are the basis for a study of the present day mass function (PDMF) of NGC~346, the largest active star forming region in the Small Magellanic Cloud (SMC). We find a PDMF slope of $\\Gamma=-1.43\\pm 0.18$ in the mass range $0.8-60\\, {\\rm M}_\\odot$, in excellent agreement with the Salpeter Initial Mass Function (IMF) in the solar neighborhood. Caveats on the conversion of the PDMF to the IMF are discussed.  The PDMF slope changes, as a function of the radial distance from the center of the NGC~346 star cluster, indicating a segregation of the most massive stars. This segregation is likely primordial considering the young age ($\\sim 3\\, {\\rm Myr}$) of NGC~346, and its clumpy structure which suggests that the cluster has likely not had sufficient time to relax. Comparing our results for NGC~346 with those derived for other star clusters in the SMC and the Milky Way (MW), we conclude that, while the star formation process might depend on the local cloud conditions, the IMF does not seem to be affected by general environmental effects such as galaxy type, metallicity, and dust content. ",
        "introduction": "\\label{intro}  Since the pioneering study by \\citet{salpeter55}, one of the major issues in astrophysics is whether the initial mass function (IMF), the relationship that specifies the mass distribution of a newly formed stellar population, is  universal or, alternatively, determined by environmental effects. In particular, establishing whether the IMF has been constant over the evolution of the universe, or if it varies with time (redshift) and/or metallicity has crucial consequences on the evolution, surface brightness, chemical enrichment and baryonic content of galaxies, and on the evolution of light and matter in the universe.  Due to a number of effects and observational uncertainties such as mass segregation, field star contamination, and variable extinction \\citep{elmegreen06}, the ``true\" IMF is very difficult to measure. Despite these limitations, star clusters still provide the best determination of the IMF, because they are essentially single--age, single--metallicity stellar systems. In general, all star clusters lose more than 80\\% of their stellar content in their first $\\sim 10$ Myr \\citep{kroupa02,lada03}. Therefore, in order to perform an accurate census of the initial stellar content it is necessary to focus on the youngest stellar systems. In addition, young (age $<3-5\\, {\\rm Myr}$) star clusters in nearby galaxies (e.g. the Large --LMC-- and Small Magellanic Cloud --SMC) offer to us the unique advantage to individually resolve --- because of their proximity --- statistically significant samples of stars down to the sub--solar mass regime.  In this paper, we focus our attention on the mass function (MF) of NGC 346, an extremely young \\citep[$\\sim 3\\, {\\rm Myr}$,][hereafter S07]{bouret03,sabbi07}, and active star cluster that excites the largest and brightest H{\\sc ii} region (N66) in the SMC.   NGC~346 ($\\alpha_{j2000}=00^h 59^m 05.2^s, \\delta_{j2000}=-72^{\\deg} 10'28''$) contains a major fraction of the O stars known in the entire SMC \\citep{walborn78,walborn86,niemela86,massey89}. The bright end of its stellar population ($V\\le 19.5\\, {\\rm mag}$) has been well investigated in the past 20 years. Spectral investigations of the brightest members identified several stars of spectral type O6.5 or earlier \\citep{massey89,heap06,mokiem06,evans06}. \\citet{massey89} found the  MF for  the brightest stars in NGC 346 --- down to 5 M$_{\\odot}$ --- to have a slope that is consistent with what has been found for massive stars near the Sun and in the LMC.  High resolution observations obtained with the Advanced Camera for Survey (ACS) on board of the \\emph{Hubble Space Telescope} (\\emph{HST}) recently revealed that the star formation history of the region where NGC~346 is located is quite complex. There is evidence for old episodes of star formation (SF), in the field of the SMC approximately between 3 and 5 Gyr ago. After that, the SF activity in this region appears to have been significantly diminished, with a possible moderate enhancement $\\sim 150\\, {\\rm Myr}$ ago (S07). An intermediate--age ($\\sim 4.5\\, {\\rm Gyr}$) stellar cluster (BS90), characterized by a core radius $r_c\\simeq 25\\, {\\rm arcsec}$ and a tidal radius $r_t \\simeq 130\\, {\\rm arcsec}$, is located at a projected distance of $\\sim 23\\, {\\rm pc}$ northeast from the center of NGC~346 (S07). Spectral analysis of NGC~346 members revealed the presence of stars as young as $\\sim 1\\, {\\rm Myr}$, \\citep{massey05}, but stars $\\sim 5\\, {\\rm Myr}$ old are also present \\citep{heap06,mokiem06}.  Star formation is likely still active in the region: hundreds of pre--main sequence (pre--MS) stars in the mass range $0.6 - 3\\, {\\rm M}_{\\odot}$ were discovered by \\citet{nota06} from \\emph{HST}/ACS images. \\emph{Spitzer Space Telescope} (\\emph{SST}) observations also detected a myriad of embedded young stellar objects (YSOs) scattered across the entire region \\citep{bolatto07}. By assuming a Salpeter IMF, \\citet{simon07} calculated that more than 3000 M$_\\odot$ have been formed in the last $\\sim 10^6\\, {\\rm yr}$, concluding that $\\ge 6$\\% of the current SF in the SMC is taking place in NGC~346.  This recent stellar population does not appear to be uniformly distributed within the ionized nebula, but is rather organized in many, likely coeval, subclusters (S07) that are coincident with clumps of molecular gas previously detected by \\citet{rubio00}. \\citet{simon07} noted that almost all the subclusters host one or more YSOs.  The high sensitivity of the \\emph{HST}/ACS allowed us to investigate the stellar content of NGC~346 from $\\sim 60 M_\\odot$ down to $\\sim 0.6 M_\\odot$, making NGC~346 one of the few known regions where a MF can be determined over two orders of magnitude \\citep[the other classical case being the Orion trapezium cluster in our own MW;][]{elmegreen06}. This paper gives a new derivation of the NGC 346 MF, reviews the strengths and limitations of the methods, and presents the conclusions that we can conservatively derive. The paper is organized as follows: a short description of the data reduction procedure is presented in \\S\\ref{obs}; in \\S\\ref{CMD} we present the color--magnitude diagram (CMD), and we describe the stellar populations identified. In \\S\\ref{MF} we present the MF we obtained for the NGC~346 cluster, we discuss its spatial variations, and we analyze the impact of the environment on the MF. The results of this paper are discussed in \\S\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  Our analysis of the stellar content of NGC~346 indicates that, within all the uncertainties that can affect the determination of a MF, the PDMF slope of NGC 346, at least down to the solar mass regime, is consistent with the value derived by \\citet{salpeter55} for the IMF in the solar neighborhood, further confirming that the stellar MF does not strongly depend on the metallicity, or the mass or morphological type of the parent galaxy.  We studied how the PDMF varies as a function of the distance from the center of NGC~346. We found that the PDMF is quite flat in the center, and becomes steeper in the periphery. As already found in other young star clusters \\citep{hillenbrand97,hillenbrand98,sirianni02} the increase in the MF slope with the distance from the center is due to the lack of massive stars in the periphery rather than an excess of low--mass objects there. The projected density of massive stars from the innermost annulus and the outskirts decreases by a factor of 60, whereas low-mass stars are depleted only by a factor of 6; the most massive stars are segregated in the center.  We conclude that the spatial stellar mass distribution is unlikely to be biased by age effects. SST observations indicate that there is ongoing SF throughout the complex with a lower limit on the total star formation rate of $3.2 \\times 10^{-3}\\, {\\rm M}_\\odot\\, {\\rm yr}^{-1}$ \\citep{simon07}. Furthermore \\citet{mokiem06} noted a spread of about 5 Myr in the ages of NGC~346 stars, but  did not find any correlation between the age and the position of the stars over the whole cluster.  S07 noted that stars in NGC~346 appear to be organized in several coeval sub--clusters, embedded in H{\\sc ii} gas, and coinciding with the CO clumps analyzed by \\citet{rubio00}. YSOs between Class I and Class III are found in 14 of the 16 identified sub--cluster \\citep{simon07}. The sub--clusters are also connected by filaments and arcs of gas and dust.  The complex structure of NGC~346 and its young age, compared to its dynamical time, suggest that NGC~346 is not dynamically evolved, and probably not yet relaxed. On the basis of these considerations we conclude that the observed segregation of the most massive stars is likely primordial, and reflects how the cluster formed.  S07 suggested that in NGC~346 most of the star formation was likely not triggered, but rather resulted from the turbulence driven density variations within a giant interstellar cloud complex, conditions predicted by the hierarchical fragmentation of a turbulent molecular cloud model \\citep{elmegreen00,klessen00,bonnell02,bonnell03}. According to this model, the fragmentation of the cloud is due to supersonic turbulent motions present in the gas. The turbulence induces the formation of shocks in the gas, and produces filamentary structures \\citep{bate03}. The chaotic nature of the turbulence increases locally the density in the filamentary structures. When regions of high density become self--gravitating, they start to collapse to form stars. Simulations show that star formation occurs simultaneously at several different locations in the cloud \\citep{bonnell03}, as appears to be the case for NGC~346.  In this scenario, stars forming near the cluster center would have higher accretion rates, due to the local higher gas density \\citep{larson82,bonnell01}. Simulations by \\citet{bonnell01} indicate that although low--mass stars form equally well throughout the entire cluster, the most massive stars are almost exclusively segregated to the central regions. It is interesting to note that the two most massive stars (one of them being an O3V spectral type, Figs.~\\ref{f:centro}) are $\\sim 7\\, {\\rm pc}$ off the center of the cluster (outside the core of NGC~346, but still within the half mass radius). Did these stars form where they are observed today, in a quite low stellar density region?  \\citet{vine03} studied the effect of strong stellar winds and photoionizing fluxes coming from O- and B-type stars on the gas content in a young star cluster. According to their simulations, the gas expulsion is unlikely to remove any initial mass segregation, but it can still result in the ejection of some of the most massive stars to positions outside the cluster core, although still within the half mass radius. According to \\citet{bouret03} mechanical power from O-star winds is not a dominant factor in the evolution within N66, but the molecular gas in the original cloud has been strongly photodissociated and photoionized by the O stars \\citep{rubio00}.  \\citet{niemela86} measured the radial velocity of five of the innermost and brightest stars in NGC~346, and found a velocity dispersion of $\\sim 5\\, {\\rm km\\, s}^{-1}$, indicating that, even in $\\sim 1\\, {\\rm Myr}$, they can be formed in the center of the cluster and then ejected where they are observed today. Another possibility is that these stars formed where are observed, as the result of a very efficient competitive accretion process into a potential well \\citep{bonnell03}.  Further spectroscopic analysis of the dynamics of the gas and the sub--clusters can confirm whether the giant molecular cloud which formed NGC~346 underwent under a hierarchical fragmentation. The initial supersonic turbulence originates shocks in the gas which rapidly remove kinetic energy from the gas \\citep{ostriker01}: in this case we expect to find low gas velocities around the sub--clusters. A detailed study of the sub--clusters dynamics will also reveal if NGC~346 is or is not already relaxed and virialized. We will present these results in a future paper."
    },
    "0710/0710.0591_arXiv.txt": {
        "abstract": "We report the first fully sampled maps of the distribution of interstellar CO$_2$ ices, H$_2$O ices and total hydrogen nuclei, as inferred from the 9.7 $\\mu$m silicate feature, toward the star-forming region Cepheus A East with the IRS instrument onboard the {\\it Spitzer Space Telescope}. We find that the column density distributions for these solid state features all peak at, and are distributed around, the location of HW2, the protostar believed to power one of the outflows observed in this star-forming region. A correlation between the column density distributions of CO$_2$ and water ice with that of total hydrogen indicates that the solid state features we mapped mostly arise from the same molecular clumps along the probed sight lines. We therefore derive average CO$_2$ ice and water ice abundances with respect to the total hydrogen column density of $X$(CO$_2$)$_{ice}$$\\sim$1.9$\\times$10$^{-5}$ and $X$(H$_2$O)$_{ice}$$\\sim$7.5$\\times$10$^{-5}$. Within errors, the abundances for both ices are relatively constant over the mapped region exhibiting both ice absorptions. The fraction of CO$_2$ ice with respect to H$_2$O ice is also relatively constant at a value of 22\\% over that mapped region. A clear triple-peaked structure is seen in the CO$_2$ ice profiles. Fits to those profiles using current laboratory ice analogs suggest the presence of both a low-temperature polar ice mixture and a high-temperature methanol-rich ice mixture along the probed sightlines. Our results further indicate that thermal processing of these ices occurred throughout the sampled region. ",
        "introduction": "Understanding how the structure and composition of ice mantles covering interstellar grains change with the local physical and chemical conditions is crucial to constraining the chemical evolution of protostellar envelopes, protoplanetary disks, and comets. Infrared observations toward low-to-high mass star-forming regions (e.g., Nummelin et al. 2001; Gibb et al. 2004) and quiescent molecular clouds (Whittet et al. 1998) using the {\\it Infrared Space Observatory} have shown that water (H$_2$O), carbon monoxide (CO) and carbon dioxide (CO$_2$) are common constituents of the ice mantles covering interstellar grains in dense molecular environments. While H$_2$O and CO are believed to form from grain surface reactions and gas-phase desorption, respectively (e.g., d'Hendecourt et al. 1985), the production of solid CO$_2$ remains uncertain. Mechanisms such as UV photolysis and/or grain surface chemistry are invoked. The routine detection of solid CO$_2$, with a fraction relative water ice between 9 and 37 \\% toward a number of star-forming regions (e.g., Gerakines et al. 1999; Nummelin et al. 2001), combined with the fact that CO$_2$ is easily produced via UV irradiation in the laboratory (Ehrenfreund et al. 1996), added weight to the suggestion that CO$_2$ ices are mainly produced via UV photolysis in the interstellar medium. However, the detection of CO$_2$ ice toward a number of quiescent clouds devoid of any embedded source of radiation, with abundances relative to water ice similar to those measured toward star-forming regions, demonstrated that production mechanisms such as grain surface reactions and gas-grain interaction need be considered too (Whittet et al. 1998; Whittet et al 2007).  A number of laboratory experiments studied the evolution of ice mixtures when subjected to thermal and/or UV processing (e.g., Ehrenfreund et al. 1996; 1999). Comparisons of these laboratory ice analogs with {\\it ISO} and more recent {\\it Spitzer} observations of interstellar ices allowed previous workers to better constrain the physical and chemical conditions in the clouds toward which these ices were detected (e.g., Ehrenfreund et al. 1998; 1999). In particular, the profile of the $\\nu_2$ vibrational bending mode of solid CO$_2$ (15.2 $\\mu$m) was found to be very sensitive to the composition of the ice mantle the molecules are embedded in and to the thermal history of the region under study.  Substructures that are characteristic of ice crystallization and segregation of the ice mantle constituents upon thermal processing have been routinely used to constrain the physical and chemical environments around young stellar objects (e.g., Boogert et al. 2000; Gerakines et al. 1999; Boonman et al. 2003; Gibb et al. 2004; Bergin et al. 2005). These previous studies were all based on pointed observations of individual discrete sources, and thus require observations of multiple sources to provide any spatial information. With the {\\it Spitzer} spectral mapping capability, one can now investigate the distribution of interstellar ices within a given dense cloud or star-forming region with unprecedented spatial sampling ($\\sim$3$''$). This provides a new opportunity to link the spectral properties of the ices with local physical conditions, and thus to throw new light on their origin and evolution. In this paper, we report on the spatial distribution and the evolution of the interstellar ices present toward the star-forming region Cepheus A East using this new capability offered by {\\it Spitzer}.  Cepheus A is a well-known site of star formation located at a distance of about 650 pc. The region contains a series of deeply embedded far-infrared and radio-continuum sources, one of which dominates the luminosity of the entire region (HW2 with 2.5$\\times$10$^{4}$ $L_{\\odot}$; Hughes \\& Wouterloot 1984; Lenzen et al. 1984; Evans et al. 1981). Ground- (e.g., Hartigan et al. 1996; Goetz et al. 1998) and space-based observations (e.g., Wright et al. 1996; van den Ancker et al. 2000) revealed the existence of a multipolar outflow exhibiting complex structures of shock-excited atomic and molecular gas components, indicating the presence of both dissociative (J-type) and non-dissociative (C-type) shocks resulting from successive episodes of activity (Narayanan \\& Walker 1996). The protostellar object HW2 was found to be the dominant powering source of the extremely high-velocity (EHV) outflow oriented {\\it northeast} ({\\it NE}), while the source (or sources) of the high-velocity jet (HV; oriented {\\it southeast}) is still under debate (e.g., Goetz et al. 1998; Hiriart et al. 2004; Codella et al. 2003; Mart\\'{\\i}n-Pintado et al. 2005).  Previous observations using {\\it ISO} exhibited a number of absorption bands due to the presence of ice mantles containing CO, CO$_2$ and H$_2$O  and due to silicate grains over the {\\it NE} side of this star forming region (van den Ancker et al. 2000). The CO and CO$_2$ ice abundances -- with respect to water ice and averaged over the Short-Wavelength Spectrometer (SWS) aperture -- were found to be within the range observed toward other sight lines (e.g., Whittet et al. 1996). Using new high-resolution {\\it Spitzer} data, we recently reported the first detection of gas-phase CO$_2$ emission. This emission extends over a 35$''\\times$25$''$ region associated with the EHV outflow (Sonnentrucker et al. 2006). We determined that the gaseous CO$_2$ molecules mostly result from sputtering of ice mantles due to interactions between the EHV outflow and the ambient molecular medium. We derived a CO$_2$ gas-to-ice ratio of at most 3\\% over the region showing CO$_2$ gas emission.  Our data also exhibit strong absorption from the 15 $\\mu$m CO$_2$ ice bending mode, from the 6 $\\mu$m water ice band and from the 9.7 $\\mu$m silicate feature at the position of HW2 and at numerous spatial positions preferentially located away from the EHV outflow region. In this paper, we discuss the distribution of the solid state features observed toward Cepheus A East with the {\\it Spitzer} Infrared Spectrograph (IRS). While H$_2$ $S$(0) to $S$(7), [NeII], [NeIII], [S I], [SIII] and [FeII] emissions as well as emission from C$_2$H$_2$ (Sonnentrucker et al., submitted) are also detected in our data, we focus here on the study of the absorptions from CO$_2$ ices, H$_2$O ices and silicate grains -- the latter used to estimate the total column density of hydrogen nuclei $N$(H$_{\\rm{tot}}$) -- toward Cepheus A East. Section 2 summarizes our observations and Section 3 describes our data analysis. Section 4 compares the spatial distribution of the CO$_2$ ice,  H$_2$O ice and silicate grains with that of the quiescent molecular clouds present in this region. Section 5 discusses the ice abundances as well as the variations in ice mantle composition derived from fits to representative CO$_2$ ice profiles using current laboratory ice analogs. ",
        "conclusions": "We used new {\\it Spitzer} data obtained with the IRS instrument to produce fully sampled maps of the distribution of CO$_2$ ices (15.2 $\\mu$m), H$_2$O ices (6.02 $\\mu$m) and total hydrogen nuclei, as inferred from the 9.7 $\\mu$m silicate feature, toward the star-forming region Cepheus A East. We find that all solid state features peak at, and are distributed closely around, the spatial position of the deeply embedded protostar HW2 and coinciding with the NH$_3$(1,1) molecular clump CepA-1. Otherwise, the distributions show column densities lower by about a factor 2 than those measured over the CepA-1 region.  The correlations we observe between the CO$_2$ ice, the H$_2$O ice column density distributions and that of total hydrogen imply that the solid state features predominantly arise from the same molecular clouds along the probed sight lines. We hence derived the CO$_2$ and water ice abundances with respect to the measured total hydrogen column density at each summed spatial position in the map. We find an average CO$_2$ ice abundance of (1.9 $\\pm$ 0.4) $\\times$ 10$^{-5}$ and an average H$_2$O ice abundance of (7.5 $\\pm$ 1.7) $\\times$ 10$^{-5}$ over the probed region.  A comparison of the CO$_2$ and water ice column density distributions indicates that both ices build up onto dust grains, in concert, over the probed region. The fairly good correlation between the column densities of the two ices also indicates that blending of the water ice band by unidentified features is not significant in our data.  Overall, we find that $N$(CO$_2$)$_{ice}=$ (0.22 $\\pm$ 0.03) $\\times$ $N$(H$_2$O)$_{ice}$, a fraction similar to that typically found in the intra-could medium and toward quiescent molecular clouds (17 \\%), suggesting that grain surface chemistry is the most likely production mechanism of CO$_2$ ices toward this region.  Best fits to the CO$_2$ ice bending mode using the current laboratory interstellar ice analogs database indicate that the ice mantles in Cepheus A East are composed of 2 ice mixtures, a low-temperature ($T_{lab}$=20K) polar ice mixture and a high-temperature ($T_{lab}$=119K) methanol-rich mixture. We find that while about 30\\% of the probed ice mantles is at low temperature, variations of the cold-ice fraction also seem to exist over the probed region. The presence of the high-temperature methanol-rich ice mixture is indicative of significant thermal processing of a fraction of these ice mantles over Cepheus A East history."
    },
    "0710/0710.1903_arXiv.txt": {
        "abstract": "We study the deflection of light in the background of a ``wiggly\" cosmic string, and investigate whether it is possible to detect  cosmic strings by means of weak gravitational lensing. For straight strings without small-scale structure there are no signals. In the case of strings with small-scale structure leading to a local gravitational attractive force towards the string, there is a small signal, namely a preferential elliptical distortion of the shape of background galaxies in the direction corresponding to the projection of the string onto the sky. The signal can be statistically distinguished from the signal produced by a linear distribution of black holes by employing an ellipticity axis distribution statistic. ",
        "introduction": "There has been renewed interest in cosmic strings \\cite{original} as a contributing factor to cosmological structure formation. This resurgence of interest is due firstly to the realization \\cite{Rachel} that in many supersymmetric particle physics models, cosmic strings are formed after inflation, and thus contribute to but not completely replace inflationary perturbations as the seeds for structure formation. Secondly, it has recently been realized that models with cosmic superstrings \\cite{Witten} may well be viable \\cite{CMP}. They could, for example, be generated as the remnant of brane annihilation processes in brane inflation models \\cite{Tye}, or they may play an important role in inflationary models in warped backgrounds \\cite{stringinflation}. Cosmic superstrings may also be left behind after the initial Hagedorn phase in string gas cosmology \\cite{BV}, where they would add an additional component to the spectrum of fluctuations produced by thermal string gas fluctuations \\cite{NBV}.  Unlike in the original cosmic string models of structure formation (see e.g. \\cite{CSstructure} and \\cite{CSreviews} for reviews), where it was assumed that the strings were the sole source of structure formation, in the context of the current models in which cosmic strings arise at the end of inflation, the strings contribute only a small fraction of the total power to the density fluctuations. The most stringent constraints on the fraction of the power which cosmic strings can contribute come from measurements of the angular power spectrum of coordinates cosmic microwave background (CMB) anisotropies. A large fraction $f$ of the power being due to strings is inconsistent with the observed acoustic oscillations in the angular power spectrum. The best current limits on $f$ are \\cite{CMBlimit} $f < 10^{-1}$. Recent work \\cite{ABB,Fraisse} points to the possibility that statistical analyses searching for the line discontinuities \\cite{KS} in the microwave temperature maps produced by cosmic strings might lead to even tighter limits.  Given their interest from the point of view of particle physics and superstring theory, it is of great interest to develop statistics to search for strings in observational data. There has been a substantial amount of work on identifying distinctive signals for strings in CMB temperature maps. In this paper we wish to take first steps at exploring another avenue - weak gravitational lensing. Weak gravitational lensing (see e.g. \\cite{lensing} for a general review of the applications of gravitational lensing to cosmology) is emerging as a powerful tool in observational cosmology to search for the distribution of matter in the universe. One of the advantages of weak gravitational lensing over, for example, galaxy redshift surveys, is that light deflection depends on the total mass, not just the luminous mass. As a consequence, cosmic strings (which are also dark in the sense of not emitting light) also lead to gravitational lensing. In fact, due to conical structure of the metric of space-time in the plane perpendicular to a long string \\cite{deficit}, the specific lensing pattern produced by a string can lead to interesting strong lensing patterns, e.g. double images (see \\cite{Esther,recent} for claims to have detected such events, claims which were subsequently shown to be incorrect \\cite{true}). In this short paper, we take a first step at searching for weak lensing signals from strings.  Cosmic strings \\cite{original} are one-dimensional topological defects which arise during phase transitions in the very early universe. Since they carry energy, they will lead to density fluctuations and CMB anisotropies. Causality implies that the network of strings which forms during the phase transition contains infinite strings. Once formed in the early universe, the network of strings will approach a ``scaling solution'' which is characterized by of the order one infinite string segment in each Hubble volume, and a distribution of cosmic string loops which are the remnants of the previous evolution. In particular, this implies that in any theory which admits cosmic strings, a network of strings will be present at the current time. The network will consist of a small number of strings crossing out entire Hubble volume (these are commonly called the ``infinite'' strings). The mean curvature radius $R_c$ of these strings will be comparable to the Hubble radius. Thus, for observations on smaller angular and distance scales, these strings can be approximated as straight. In addition, there is a distribution of string loops with radii $R$ between $R_c$ and a lower scale set by the strength of the gravitational radiation which the loops emit \\cite{CSstructure}: \\be G \\mu R_c \\, \\le \\, R \\, \\le \\, R_c \\, . \\ee String loops with smaller radius live for less than one Hubble expansion time. In the above, $\\mu$ is the mass per unit length of the string and $G$ is Newton's gravitational constant  The distinctive signatures of strings in observational data are a consequence of the specific form of the metric produced by a cosmic string: space perpendicular to a cosmic string is a cone with deficit angle given by \\cite{deficit} \\be \\alpha \\, = \\, 8 \\pi G \\mu \\, . \\ee  In this paper we study the specific signature of this cosmic string metric on weak gravitational lensing. We consider a cosmic string lens and investigate the shape distortions this lens induces for a screen of background galaxies as source objects.  To our knowledge, this is the first investigation of weak gravitational lensing from cosmic strings. Strong lensing signatures (lines of double images) have been studied extensity \\cite{Vilenkin2,Hogan,Gott,Paczynski,Dyer,Vachaspati,Mack} theoretically, and looked for in observational data sets \\cite{Hindmarsh,Shirasaki}.  The outline of this paper is as follows: we first derive the deflection angle induced by a cosmic string with small scale structure. We then describe a gravitational lensing simulation based on this result, and propose a statistic which could be used to search for this type of effect in future weak lensing surveys. ",
        "conclusions": "In this work we have studied the gravitational lensing by a straight cosmic string containing small-scale structure which leads to a string tension which is less than the mass per unit length of the string, and thus induces a net gravitational force towards the string which test particles feel.  Next, we studied potential weak lensing signatures of such wiggly strings. We found a shape distortion which is proportional to $G (\\mu - T)$ but depends on the location of the source relative to the line of sight between the observer and the string. Only objects which are displaced in direction of the string relative to the line of sight give rise to a shape distortion. The shape distortion increases linearly as the distance in string direction increases.  The specific distribution of the weak lensing distortion in the image plane in principle can be used to provide a signature for cosmic strings. However, in practise the signal appears to be too small to be useful. Even for optimal choices of parameters (galaxy redshifts comparable to those of galaxies in the largest current redshift surveys), distance from the line of sight comparable to the distance of the string from the observer, large amount of small scale structure and string tension close to the current observational bounds, the intrinsic ellipticity of the background galaxies needs to be very close to one (i.e. the shape needs to be very close to spherical) in order for the cosmic string signal to stand out."
    },
    "0710/0710.3906_arXiv.txt": {
        "abstract": "We model the thermal effect of young stars on their surrounding environment in order to understand clustered star formation. We take radiative heating of dust, dust-gas collisional heating, cosmic-ray heating, and molecular cooling into account. Using Dusty, a spherical continuum radiative transfer code, we model the dust temperature distribution around young stellar objects with various luminosities and surrounding gas and dust density distributions. We have created a grid of dust temperature models, based on our modeling with Dusty, which we can use to calculate the dust temperature in a field of stars with various parameters. We then determine the gas temperature assuming energy balance. Our models can be used to make large-scale simulations of clustered star formation more realistic. ",
        "introduction": "Most of the stars in our galaxy form in groups or clusters \\citep{lada}. Therefore, in order to understand the star formation history, the shape of the mass function, and the formation of massive (M $\\ga 5$ M$_{\\odot}$) stars in our galaxy, the star formation process must be studied in its most common environment -- a cluster. As stars form from their initial reservoir of gas and dust, they interact with their environment and heat the surrounding material, thus affecting future star formation. One of the first effects a protostar has on its environment is radiative heating from the accretion luminosity and, subsequently, nuclear fusion. The radiation efficiently heats the dust, which in turn heats the gas through collisions. Young stars also affect their environment via strong winds and ionization, but this only occurs when they are very massive and have evolved past the very early stages of star formation. We assume that the massive stars in our sample are very young and are accreting at very high rates ($\\dot{\\rm{M}} \\ga 1 \\ee{-5} \\msun / $yr). This high accretion rate allows the infalling mass to absorb all of the stellar UV photons \\citep{Churchwell}.  Many groups use large scale computer simulations to model clustered star formation. This is a complicated process requiring many assumptions in order to make the problem tractable. \\cite{klessen} and Martel, Evans, \\& Shapiro (2006) assume that the gas is isothermal. \\citet{bbb} go beyond this assumption by using a barotropic equation of state. However, until recently, no one has included the effect of radiatively heating the dust and gas by the stars formed in the simulation. \\citet{krumholz} have included an approximate radiative transfer method, which works well in optically thick regions. Their method assumes that the gas temperature is equal to the dust temperature throughout their simulation. This approximation is only valid at high densities when the dust and gas are collisionally coupled. Here we develop a method that explores the effect of radiative heating and the dust and gas energetics for a range of optical depths and densities.   In our method we include various heating and cooling processes to calculate the dust and gas temperature. Stars can heat dust grains more effectively than the gas because dust grains have broad-band absorption properties. Although we will not be explicitly modeling the motion or energy density of dust grains, we assume the dust and gas are well-mixed and the dust grains transfer energy to gas particles through collisions using the energy transfer rate discussed in \\citet{young}. The gas is heated by collisions with hot dust grains and cosmic rays. It can cool through CO and other molecular line emission.  In this paper, we calculate the dust and gas temperature in a field of stars. The dust temperature around a single source is calculated using a look-up table which we develop here. With this look-up table and an approximation to the flux-temperature conversion, we calculate the dust temperature in the field. Our look-up table is needed since the calculation of a single dust temperature distribution can take longer than a minute on current desktops and would take a substantial fraction of a large-scale simulation's computations. Therefore we outline our method here which can be used to decrease the time spent on the calculation of the dust temperature in future studies of clustered star formation.  With the calculated value of the dust temperature, we derive the gas temperature field for a distribution of stellar sources, as in a young stellar cluster. The effect that protostars have on heating their environment using a hydrodynamic and gravity simulation will be addressed in a future paper. In this paper, we first discuss the calculation of the dust temperature for single and multiple sources (\\S \\ref{sec:analy}), then we describe our gas temperature calculation (\\S \\ref{sec:gas}), and finally, we show some dust and gas temperature distributions when multiple sources are present (\\S \\ref{sec:multiple}). ",
        "conclusions": "We have presented a method for calculating the dust and gas temperature between stellar sources. The analytic method that we investigated for calculating the dust temperature was not accurate enough. Instead, our chosen method of calculating the dust temperature uses a simple radiative transfer code which we use to create a look-up table. Once we have derived the dust temperature, we are able to calculate the gas temperature by balancing various energy processes. We include dust-gas collisional heating, molecular cooling, and cosmic-ray heating. When we have balanced the energies, we are able to derive the gas temperature. Other methods which set the gas temperature and dust temperature equal assume the gas and dust are opaque everywhere, which is not always true.  In Figure \\ref{fig:perdiff}, we show the percentage difference between the gas and dust temperature, as well as the density, for the distribution of sources discussed in \\S \\ref{sec:three}.  These two figures show that the largest percentage difference between the dust and gas temperature occurs where the density is the lowest.  Therefore, at low densities ($n \\la 10^{5}$cm$^{-3}$), assuming equal dust and gas temperatures is not appropriate.  We plan to use the method discussed in this paper to model a region of clustered star formation with the three-dimensional hydrodynamics code discussed in \\cite{martel}. Our method of calculating the gas and dust temperature distribution in a field of young stars will enable us and others to more accurately model clustered star formation observationally and in future simulations."
    },
    "0710/0710.3517.txt": {
        "abstract": "Forty years after the discovery of rotation-powered pulsars, we still do not understand many aspects of their pulsed emission.  In the last few years there have been some fundamental developments in acceleration and emission models.  I will review both the basic physics of the models as well as the latest developments in understanding the high-energy emission of rotation-powered pulsars, with particular emphasis on the polar-cap and slot-gap models.   Special and general relativistic effects play important roles in pulsar emission, from inertial frame-dragging near the stellar surface to aberration, time-of-flight and retardation of the magnetic field near the light cylinder.  Understanding how these effects determine what we observe at different wavelengths is critical to unraveling the emission physics.  I will discuss how current and future X-ray and gamma-ray detectors can test the predictions of these models. ",
        "introduction": "\\label{sec:intro}  % Always give a unique label % and use \\ref{<label>} for cross-references % and \\cite{<label>} for bibliographic references % use \\sectionmark{} % to alter or adjust the section heading in the running head  Rotation-powered pulsars are fascinating astrophysical sources and excellent laboratories for study of fundamental physics of strong gravity, strong magnetic fields, high densities and relativity.  The major advantage we have in studying pulsars is that we know they are rotating neutron stars and that they derive their power from rotational energy loss. The challenge is then to understand how they convert this source of power into the visible radiation.  It is generally agreed that this occurs through acceleration of charged particles to extremely relativistic energies, using the rotating magnetic field as a unipolar inductor to create very high electric potentials.  Beyond this fundamental, there is a large divergence of thought on what comes next: whether the acceleration occurs in the strong field near the neutron star surface or in the outer magnetosphere near the speed of light cylinder, or even beyond the light cylinder in the wind zone.  The particle acceleration may well be occurring in all of these regions, either in the same pulsar or in pulsars of different ages.  In recent years, there has been much activity both in new detections and in theoretical study of rotation-powered pulsars.  Multibeam radio surveys at the Parkes Telescope \\cite{Manchester2001,Edwards2001} have increased the population of known radio pulsars to more than 1700. In addition, extended radio observations of supernova remnants and unidentified $\\gamma$-ray sources have discovered a number of young pulsars that are too radio-faint to be detected by surveys \\cite{Camilo2004}. Although pulsed emission at other wavelengths has been detected from only a small fraction of these, this number is growing as well.  At the present time, there are 7 pulsars with high-confidence detection of $\\gamma$-ray pulsations \\cite{Kanbach2006}, about 30 having X-ray pulsations \\cite{Kaspi2006} and 10 with optical pulsations \\cite{Mignami2004}.  This paper will review both the fundamental physics and latest theoretical developments of acceleration and radiation in polar cap models.  A complementary paper by Cheng \\cite{Cheng2006} reviews acceleration and radiation in the outer gap model. ",
        "conclusions": ""
    },
    "0710/0710.0029_arXiv.txt": {
        "abstract": "Using the OVRO, Nobeyama, and IRAM mm-arrays, we searched for ``disk''-outflow systems in three high-mass (proto)star forming regions: \\gs, \\gtt, and \\gte. These were selected from a sample of \\amm\\ cores (Codella, Testi \\& Cesaroni) associated with OH and \\wat\\ maser emission (Foster \\& Caswell) and with no or very faint continuum emission. Our imaging of molecular line (including rotational transitions of \\mcn) and 3\\,mm dust continuum emission revealed that these are compact ($\\lesssim$ 0.05 -- 0.3 pc), massive ($\\sim$ 100 -- 400 \\Msun), and hot ($\\sim$100 K) molecular cores (HMCs), that is likely sites of high-mass star formation prior to the appearance of ultracompact \\hii\\ regions. All three sources turn out to be associated with molecular outflows from \\co\\ and/or HCO$^+$ $J=$1--0 line imaging. In addition, velocity gradients of 10 -- 100 \\kms\\ pc$^{-1}$ in the innermost ($\\lesssim$0.03 -- 0.13 pc), densest regions of the \\gtt\\ and \\gte\\ HMCs are identified along directions roughly perpendicular to the axes of the corresponding outflows. All the results suggest that these cores might be rotating about the outflow axis, although the contribution of rotation to gravitational equilibrium of the HMCs appears to be negligible. Our analysis indicates that the 3 HMCs are close to virial equilibrium due to turbulent pressure support. Comparison with other similar objects where rotating toroids have been identified so far shows that in our case rotation appears to be much less prominent; this can be explained by the combined effect of unfavorable projection, large distance, and limited angular resolution with the current interferometers. ",
        "introduction": "\\label{ss:intro}   The role of disks in the formation process of low-mass stars (i.e. stars with masses $\\lesssim 1$ \\Msun) has been extensively studied in the last two decades through high angular resolution observations at various wavelengths.  Images of such disks have been obtained in the optical (e.g. Burrows et al. 1996) and at mm-wavelengths (e.g. Simon et al. 2000).  The latter have demonstrated that the majority of the disks undergo Keplerian rotation.  These findings are consistent with the fact that low-mass stars form through accretion while (partial) conservation of angular momentum during the dynamical collapse produces a flattened and rotating structure at the center of the core.\\par   What about high-mass ($M_\\ast \\gtrsim 8$ \\Msun) stars?  In this case, formation through accretion faces the problem that stars more massive than $\\sim$8~$M_\\odot$ reach the zero age main sequence still deeply embedded in their parental cores (Palla \\& Stahler 1993).  At this point radiation pressure from the newly formed early-type star can halt the infall, thus preventing further growth of the stellar mass.  Various solutions have been proposed to solve this problem: (i)~massive stars might form through merging of lower mass stars (Bonnell, Bate \\& Zinnecker, H. 1998; Bonnell \\& Bate 2002; Bally \\& Zinnecker 2005); (ii)~sufficiently large accretion rates could allow the ram pressure of the infalling material to overcome the radiation pressure form the star (Behrend \\& Maeder 2001; McKee \\& Tan 2003; Bonnell, Vine, \\& Bate 2004) (iii)~non-spherical accretion could weaken the effect of radiation pressure by allowing part of the photons to escape through evacuated regions along the outflow axis and, at the same time, enhancing the ram pressure of the accreting material by focusing it through the disk plane (Yorke \\& Sonnhalter 2002; Krumholz, McKee, \\& Klein 2005).\\par   An important test to discriminate between the different hypotheses is the presence of rotating, circumstellar disks, which would lend support to the third scenario depicted above. This can be determined by inspecting the velocity field of the innermost parts of molecular cores where young massive (proto)stars are believed to form. One possibility is to look for velocity gradients perpendicular to the direction of the larger scale outflows associated with such cores; this would suggest that one is observing rotation about the outflow axis. This technique has been adopted by us and other authors successfully inferring the existence of rotation in the gas enshrouding high-mass young stellar objects (YSOs) and leading to the discovery of circumstellar disk-like structure or ``toroids'' in a limited number of cases (see Cesaroni et al. 2007 for a review on this topic).\\par    The goal of the present study was to establish whether the presence of rotation is common in high-mass star forming regions by observing three more objects, expected to be sites of deeply embedded OB (proto)stars.  The ideal target for this type of studies are hot molecular cores (HMCs), which are believed to be the cradles of massive stars (see e.g. Kurtz et al. 2000; Cesaroni et al. 2007).  Beside other studies, the one by Codella, Testi \\& Cesaroni (1997, hereafter CTC97), has proved successful in identifying the birthplaces of young massive stars as dense \\amm\\ cores associated with OH and H$_2$O maser emission.  One of these HMCs (G\\,24.78+0.08) has been the subject of a series of articles by us (Furuya et al. 2002; Cesaroni et al. 2003; Beltr\\'an et al. 2004, 2005), which have shown that this is a unique object, characterized by the simultaneous presence of rotation, outflow, and infall towards a hypercompact \\hii\\ region ionized by an O9.5 star (Beltr\\'an et al. 2006). With this in mind, we completed the study of the sources in CTC97 by observing three more objects in the same tracers used to investigate G\\,24.78+0.08.  The selected targets for this study are G\\,16.59$-$0.05 at $d =$ 4.7 kpc, G\\,28.87$+$0.07 at 7.4 kpc, and G\\,23.01$-$0.41 at 10.7 kpc. The fact that in all three cases no or only faint free-free continuum emission has been detected, notwithstanding the large luminosities of the sources (CTC97), suggests that the embedded YSOs might be in an even earlier evolutionary phase than those in G\\,24.78+0.08 (Furuya et al. 2002; Beltr\\'an et al. 2004). ",
        "conclusions": "Using the OVRO, Nobeyama, and IRAM-PdB mm-interferometers, we carried out intensive search for rotating toroids towards the massive YSOs in \\gs, \\gte, and \\gtt\\ which exhibit no or faint free-free emission (CTC97). Our observations revealed that these objects are embedded in HMCs with masses of 95--380 \\Msun\\ and temperatures of 93--130 K, making the cores typical site of high-mass (proto)star formation. All the 3 objects harbored in the HMCs are driving powerful (\\Fco $\\simeq$ $10^{-3}-10^{-2}$ \\Msun \\kms\\ yr$^{-1}$) CO outflows. However, the nature of the outflows in \\gte\\ and \\gtt\\ is unclear; the origin of high velocity wing emission may be attributed to either single or double outflow(s). Such ambiguity made the interpretation of velocity gradients, identified through \\mcn\\ $K$-ladder line analysis, existing in the innermost densest part of the \\gte\\ and \\gtt\\ HMCs fairly difficult. The velocity gradients are almost perpendicular to their molecular outflow axes, suggesting the presence of rotating, flattened structures. However, the corresponding dynamical masses are an order of magnitude smaller than the masses derived from 3\\,mm dust continuum emission, which indicates turbulent pressure as the dominant support of the HMCs. No conclusion could be reached for the third source, \\gs, as the putative rotation axis appears to lie close to the line-of-sight, thus making the detection of the rotation velocity very difficult for projection effects. Further higher resolution imaging will allow us to establish the presence of rotation on a more solid ground."
    },
    "0710/0710.3670_arXiv.txt": {
        "abstract": "We discuss a general fourth-order theory of gravity on the brane. In general, the formulation of the junction conditions (except for Euler characteristics such as Gauss-Bonnet term) leads to the higher powers of the delta function and requires regularization. We suggest the way to avoid such a problem by imposing the metric and its first derivative to be regular at the brane, while the second derivative to have a kink, the third derivative of the metric to have a step function discontinuity, and no sooner as the fourth derivative of the metric to give the delta function contribution to the field equations. Alternatively, we discuss the reduction of the fourth-order gravity to the second-order theory by introducing an extra tensor field. We formulate the appropriate junction conditions on the brane. We prove the equivalence of both theories. In particular, we prove the equivalence of the junction conditions with different assumptions related to the continuity of the metric along the brane. ",
        "introduction": "\\label{sect1}  \\setcounter{equation}{0}  Brane universes have made great popularity during the last years \\cite{RS,brane}. However, it is remarkable that so far only the standard Einstein gravity, Gauss-Bonnet gravity \\cite{deruelle00,charmousis,davis,jim,lidsey,maeda,apostopoulos} and, in general, Euler density gravity \\cite{lovelock} on the brane have been considered in the literature \\cite{meissner01}. These can be expressed by the general action \\footnote{We use convention (-++...+) for the metric following Ref. \\cite{he}.} \\bea \\label{euler} S = \\int_M d^D x \\sqrt{-g} \\sum_n \\kappa_n I^{(n)} + S_{brane} + S_{m}~, \\eea where $I^{(n)}$ is the Euler density of the n-th order, $\\kappa_n$ is an appropriate constant of the n-th order, $M$ is a $D$-dimensional manifold, $S_{brane}$ is the brane action and $S_m$ is the matter action. The lowest order Euler densities are: the cosmological constant $I^{(0)} = 1$, the Ricci scalar $I^{(1)} = R$, and the Gauss-Bonnet density $I^{(2)} = R_{GB} = R_{abcd}R^{abcd} - 4 R_{ab}R^{ab} +R^2$ with appropriate constants $\\kappa_0 = -2\\Lambda(2\\kappa^2)^{-1} = -2\\Lambda/16 \\pi G$, $\\kappa_1 = (2\\kappa^2)^{-1}$, $\\kappa_2=\\alpha(2\\kappa^2)^{-1}$, $\\alpha=$ const. etc., $a,b,c=0,1,\\ldots, D-3, D-2, D$ \\cite{davis}.  In fact, in a general class of brane models based on an arbitrary combination of the higher-order curvature terms $f(R_{abcd}R^{abcd},R_{ab}R^{ab},R)$ the field equations are fourth-order. Because of that they are plagued by the higher power terms of the second derivative of the warp factor function $\\sigma(y) = \\mid y \\mid$. This leads to a production of the higher powers of the delta function $\\delta(y)$ which can make the field equations ambiguous. Among the general class, the models based on the Euler densities are unique in the sense that the higher powers of the second derivative of the warp factor $\\partial^2 \\sigma(y)/\\partial y^2$ exactly cancel in the field equations \\cite{meissner01}. One then is easily able to formulate appropriate junction conditions given first by Israel \\cite{israel66,visser}.  The main objective of our paper will be the study of a general {\\it fourth-order theory of gravity on the brane} \\cite{clifton} \\bea \\label{XYZ} S &=& \\chi^{-1} \\int_{M} d^{D}x \\sqrt{-g} f(X,Y,Z) + S_{brane} + S_{m}~, \\eea where $X=R$, $Y=R_{ab}R^{ab}$, $R_{abcd}R^{abcd}$ are curvature invariants, and $\\chi$ is a constant. It includes the Euler density theories with the first Euler density being just $f(X,Y,Z) = \\chi \\kappa_1 X = \\chi \\kappa_1 R $ and the second Euler density being the Gauss-Bonnet term given by $f(X,Y,Z) = \\chi \\kappa_2 (Z - 4Y +X^2)$ etc.  Up to our knowledge, the only non-Eulerian density cases were studied in Refs. \\cite{branef(R)} and \\cite{braneR2}. In Ref. \\cite{branef(R)} the fourth-order theory $f(X,Y,Z) = f(X) = f(R)$ was first reduced to the second-order theory, and then transformed into the Einstein theory. The junction conditions were then obtained, and they were obviously free from the problem of the powers of $\\delta-$function contribution. On the other hand, in Ref. \\cite{braneR2} the theories with the linear combination of the form $f(X,Y,Z) = aX^2 + bY + cZ$ ($a,b,c=$ const.) were considered and the junction conditions were obtained by the application of the appropriate Gibbons-Hawking boundary terms, again after transforming this theory to an equivalent second-order theory.  It is important to emphasize that the theories based on the functions of the Euler densities such as $f(I^{(n)})$ are fourth-order. Among them the most popular are $f(I^{(1)}) = f(R)$ - the theories of the function of the first Euler density \\cite{f(R)}. In fact, the theories which are based on the function of the second Euler density $f(I^{(2)})$ have also gained some interest recently \\cite{f(RGB)}, but they have not been studied on the brane yet.   Our paper is organized as follows. In Section II we discuss the main obstacle to formulate junction conditions for the fourth-order braneworld in a standard way which has been performed in the case of Euler densities. In Section III we make a proposal to formulate these junction conditions by imposing more regularity onto the metric tensor. Since it does not necessarily satisfy everybody's taste we present in Section IV an alternative approach. In this approach we transform our general fourth-order theory into a second-order theory by applying generalized Lagrange-multiplier approach \\cite{f(R),f(RGB),kijowski}. This method was successful in obtaining the junction conditions in $f(R)$ theory \\cite{branef(R)} and in $f = aX^2 + bY + cZ$ theory \\cite{braneR2}. In the Section V we formulate the junction conditions for the equivalent second-order theory. Finally, in Section VI we give our conclusions. ",
        "conclusions": "\\label{sum} \\setcounter{equation}{0}  In this paper we have considered the junction conditions for the fourth-order brane world given by the action which is a general function of the curvature invariants $R,R_{ab}R^{ab}$, and $R_{abcd}R^{abcd}$. We imposed the regularity conditions on the metric tensor which was taken to be of the class $ C^2$ functions of the coordinates. Explicitly, it means that the metric and its first derivative are regular at the brane, while the second derivative has a kink, the third derivative of the metric has a step function discontinuity, and the fourth derivative of the metric gives the delta function contribution to the field equations. In terms of the seminal notation given first by Israel \\cite{israel66}, these conditions describe the singular hypersurfaces of order three. The junction conditions which we obtained are quite generic and they may be applied to some special cosmological framework of interest.  As an alternative which allows less restrictive regularity conditions, we considered the reduction of the fourth-order theory to a second-order theory by applying an extra tensor field. We then formulated the junction conditions within such a theory and showed that they were equivalent to the previously obtained fourth-order theory junction conditions.  In the previously considered cases in the literature mainly the Eulerian density brane worlds were studied \\cite{deruelle00,charmousis,davis,jim,lidsey,maeda,apostopoulos,meissner01}. The only non-Eulerian density cases were investigated in Refs. \\cite{branef(R)} and \\cite{braneR2}, where the theories with at most a linear combination the curvature invariants of the form $f(X,Y,Z) = aX^2 + bY + cZ$ ($a,b,c=$ const.) were considered. In these references the junction conditions were obtained after transforming such theories into the equivalent second-order theories. In fact, we made one step further, by suggesting junction conditions for the brane world which allows a general function of curvature invariants $f(X,Y,Z)$. Also, in one of our approaches to the problem of junction conditions we do not reduce the fourth-order action into the second-order action, but suggest junction conditions to work in the fourth-order theory, though in the case of the so-called singular hypersurfaces of order three only \\cite{israel66}.  We hope that our calculations will allow us to study some cosmological applications of the general fourth-order gravity on the brane (\\ref{XYZ}) in the following papers.  \\vspace{0.3cm}"
    },
    "0710/0710.1958_arXiv.txt": {
        "abstract": "We explore the properties of dark energy from recent observational data, including the Gold Sne Ia, the baryonic acoustic oscillation peak from SDSS, the CMB shift parameter from WMAP3, the X-ray gas mass fraction in cluster and the Hubble parameter versus redshift. The $\\Lambda CDM$ model with curvature and two parameterized dark energy models are studied. For the $\\Lambda CDM$ model, we find that the flat universe is consistent with observations at the $1\\sigma$ confidence level and a closed universe is slightly favored by these data. For two parameterized dark energy models, with the prior given on the present matter density, $\\Omega_{m0}$,  with $\\Omega_{m0}=0.24$, $\\Omega_{m0}=0.28$ and $\\Omega_{m0}=0.32$, our result seems to suggest that the trend of $\\Omega_{m0}$ dependence for an evolving dark energy from a combination of the observational data sets is model-dependent. ",
        "introduction": "The present cosmic accelerating expansion has been confirmed by various observations, including the Type Ia Supernovae (Sne Ia)~\\cite{Perlmutter1999, Riess1998, Riess2004, Riess2006, Astier2006, Wood2007}, CMB \\cite{Balbi2000, de2000, Jaffe2001,Spergel2003, Spergel2006} and large scale structure (LSS) \\cite{Peacock2001, Eisenstein2005}, etc. In order to explain this observed phenomenon,  it is usually assumed that there exists, in the universe,  an exotic energy component with negative pressure, named dark energy (see  \\cite{Padmanabhan2006, Copeland2006, Sahni2006, Perivolaropoulos2006}  for recent reviews), which presumably began to dominate the evolution of the universe only recently. The simplest candidate of dark energy is the cosmological constant $\\Lambda$~\\cite{Weinberg1989, Sahni2000, peebles2003, Padmanabhan2003}. It fits the observational data very well, but at the same time, it also encounters two problems, i.e.,  the cosmological constant problem (why is the inferred value of cosmological constant so tiny ($120$ orders of magnitude lower) compared to the typical vacuum energy values predicted by particle physics?) and the coincidence problem (why is its energy density comparable to the matter density right now?). Therefore,  some dynamical scalar fields, such as quintessence~\\cite{Wetterich1988, Ratra1988, Caldwell1998}, phantom~\\cite{Cald} and quintom~\\cite{Quintom}, etc, are suggested as alternative candidates of dark energy.  One of the features of these scalar field models is that their equations of state parameter, $w$,  which embodies both gravitational and evolutionary  properties of dark energy, is evolving  with the cosmic expansion.  On the other hand, the growing number of dark energy models has prompted people to adopt a complementary approach, which assumes an arbitrary parametrization for the equation of state $w(z)$ in a model-independent way and aims to reconstruct the properties of dark energy directly from observations. Currently, there are many model independent parameterizations (see for example, ~\\cite{line1,line2,Jassal2005,Sahni2003, Sahni2004, Wetterich2004}). In general, using these parameterizations and the observational data, one can determine the present value of $w$ and whether it evolves as the universe expands, in particular, whether the phantom divide line (PDL) is crossed. In this regard,   Nesseris and Perivolaropoulos \\cite{Nesseris2006} has used the Chevallier-Polarski-Linder parametrization $ w(z) = w_0 + w_1 z/(1 + z)$~\\cite{line2} to explore the properties of dark energy with some observational data (including new Gold Sne Ia, SNLS Sne Ia, CMB, BAO, the cluster baryon gas mass fraction(CBF)  and 2dF galaxy redshift survey(2dFGRS) ) and found that the Gold data set mildly favors dynamically evolving dark energy with the crossing of the PDL while the SNLS does not,  and the combination of CMB+BAO+CBF+2dFGRS mildly favors the crossing of PDL only for low values of $\\Omega_{m0}$ ( $\\Omega_{m0}\\le 0.25$) prior considered and with  a higher prior matter density the evolving features of dark energy becomes weaker and weaker. Similar trend of  $\\Omega_{m0}$ dependence was found using the model~\\cite{Alam2006}, $w(z)=\\frac{1+z}{3}\\frac{A_1+2A_2(1+z)}{\\Omega_{DE}}-1$, with the CMB and BAO.  However, constraints from a combination of the supernovae and other observational data has not been analyzed in Ref.~\\cite{Nesseris2006}, and although that of the Sne and CMB+BAO was examined in Ref.~\\cite{Alam2006}, but the marginalization was considered only for $\\Omega_{m0}=0.28\\pm 0.03$ prior. Therefore, it remains interesting to see what happens to the conclusions reached in Refs.~\\cite{Nesseris2006,Alam2006}, when the combination of all observational data is analyzed for different $\\Omega_{m0}$ prior considered. The present paper aims to fill the gap. We discuss the constraints from the combination of different observational datasets. Besides the data sets of Sne Ia, BAO and CMB, in our analysis we add the datasets of the X-ray gas mass fraction in cluster and the Hubble parameter versus redshift. Firstly the $\\Lambda CDM$ model with curvature is discussed. Then, two parameterized  dark energy models: $ w(z) = w_0 + w_1 z/(1 + z)$ and $w(z)=\\frac{1+z}{3}\\frac{A_1+2A_2(1+z)}{\\Omega_{DE}}-1$, are studied to see if the properties of dark energy thus reconstructed are model-independent. ",
        "conclusions": "In this paper, we have reconstructed the properties of dark energy from recent observational data, including the Gold Sne Ia,  the baryonic acoustic oscillation peak from SDSS, the CMB shift parameter, the X-ray gas mass fraction in clusters and the Hubble parameter data. The $\\Lambda CDM$ model with curvature and two parameterized dark energy models are discussed. We find that a spatially flat universe is allowed by these data sets at the $68\\%$ confidence level, and a closed universe is slightly favored by the observations. For two parameterized dark energy models, we give the priors on $\\Omega_{m0}$ with $\\Omega_{m0}=0.24$, $\\Omega_{m0}=0.28$ and $\\Omega_{m0}=0.32$. For the spatially flat case, the constraints on model parameters and the evolutions of $w(z)$ and $q(z)$ are studied. The Gold + CMB +BAO give the strong constraints on model parameters. If Mod1 parametrization is used, the best fit curves in Fig.~3 show  that the combination of the data sets considered in this paper favors an evolving dark energy, a crossing of the phantom divide line in the near past, and the present value of $w$ being very likely less than $-1$. Remarkably, these conclusions are almost insensitive to the chosen value of matter density, in a sharp contrast to those obtained in~\\cite{Nesseris2006} where the Sne data are not combined with other observational ones.  However, the best fit curves in Fig.~3 indicate that the properties of dark energy reconstructed using Mod2 parametrization depend on the chosen value of matter density. For $\\Omega_{m0}=0.24$, a very mildly evolving dark energy is obtained, but with the increase of $\\Omega_{m0}$ prior considered, the evolving feature of dark energy becomes more evident. This trend is just the opposite to that found in Ref. \\cite{Alam2006} for just BAO+CMB data.  Therefore, our result seems to suggest that the trend of $\\Omega_{m0}$ dependence for an evolving dark energy is model-dependent. It should be noted, however, that at the $2\\sigma$ confidence level, the cosmological constant are allowed for both Mod1 and Mod2."
    },
    "0710/0710.0115_arXiv.txt": {
        "abstract": "Null Energy Condition (NEC) requires the equation of state (EoS) of the universe $w_u$ satisfy $w_u\\geq-1$, which implies, for instance in a universe with matter and dark energy dominating $w_u=w_m\\Omega_m+w_{de}\\Omega_{de}=w_{de}\\Omega_{de}\\geq-1$. In this paper we study constraints on the dark energy models from the requirement of the NEC. We will show that with $\\Omega_{de}\\sim0.7$, $w_{de}<-1$ at present epoch is possible. However, NEC excludes the possibility of $w_{de}<-1$ forever as happened in the Phantom model, but if $w_{de}<-1$ stays for a short period of time as predicted in the Quintom theory NEC can be satisfied. We take three examples of Quintom models of dark energy, namely the phenomenological EoS, the two-scalar-field model and the single scalar model with a modified Dirac-Born-Infeld (DBI) lagrangian to show how this happens. ",
        "introduction": "It is well known that energy conditions play an important role in classical theory of general relativity and thermodynamics\\cite{Hawking:1973uf}. In classical general relativity it is usually convenient and efficient to restrict a physical system to satisfy one or some of energy conditions for study, for example, in the proof of Hawking-Penrose singularity theorem\\cite{Penrose:1964wq,Hawking:1969sw}, the positive mass theorem\\cite{Schon:1981vd} and so on, while in thermodynamics energy conditions are the bases for obtaining entropy bounds\\cite{Bousso:1999xy,Flanagan:1999jp}. Among those energy conditions, the null energy condition is the weakest one which states that for any null vector $n^\\mu$ the stress energy tensor $T_{\\mu\\nu}$ should satisfy the relation \\begin{eqnarray} T_{\\mu\\nu} n^\\mu n^\\nu \\geq 0~. \\end{eqnarray}  In general, the violation of NEC leads to the breakdown of causality in general relativity and the violation of the second law of thermodynamics\\cite{ArkaniHamed:2007ky}. These pathologies require that the total stress tensor in a physical spacetime manifold should obey the NEC. In the framework of the standard 4-dimensional Friedmann-Robertson-Walker (FRW) cosmology the NEC implies $\\rho+p\\geq 0$, which in turn gives rise to a constraint on the equation of state of the universe (EoS) $w_u$ defined as the ratio of pressure to energy density, $w_{u}\\geq -1$. In this paper we study the constraints on the dark energy models from the requirement of $w_{u}\\geq -1$.  In the early Universe with radiation dominant the EoS of the universe $w_{u}$ is approximately equal to $\\frac{1}{3}$ and in the matter dominant period $w_{u}$ is nearly zero, so NEC is satisfied well. However when the dark energy component is not negligible we have \\begin{eqnarray}\\label{NECd} w_u=w_m\\Omega_m+w_{de}\\Omega_{de}\\geq-1~, \\end{eqnarray} where the subscripts `$m$' and `$de$' stand for matter and dark energy, respectively. With $w_m = 0$, inequality (\\ref{NECd}) becomes \\begin{eqnarray}\\label{NECde} w_{de}\\Omega_{de}\\geq-1~. \\end{eqnarray} From the inequality above, we can see that models of dark energy with $w_{de}\\geq-1$ such as the Cosmological Constant and the Quintessence satisfy the NEC, but the models with $w_{de}<-1$ predicted for instance by the Phantom theory where the kinetic term of the scalar field has a wrong sign does not. Interestingly we can see that NEC might be satisfied in models if $w_{de}<-1$ stays for a short period of time during the evolution of the universe. In this paper we will show this happens in the Quintom models of dark energy.  The paper is organized as follows: in section II we will present three examples of the Quintom models to show how the NEC is satisfied and the section III is the summary of the paper. ",
        "conclusions": "In this paper we have studied the implications of NEC in the models of dark energy. We show that NEC excludes the models with $w_{de}<-1$ forever as predicted by the Phantom dark energy, however allows the possibility of having $w_{de}<-1$ for a short period of time as it happens in the Quintom models. We have shown explicitly in this paper three examples of Quintom models where NEC is satisfied."
    },
    "0710/0710.5440_arXiv.txt": {
        "abstract": "We describe numerical tools  for the  stability analysis of extrasolar planetary systems. In particular, we consider the relative Poincar\\'e variables and symplectic integration of the equations of motion. We apply the tangent map to derive a numerically efficient algorithm of the fast indicator MEGNO (a measure of the maximal Lyapunov exponent) that helps to distinguish  chaotic and regular configurations. The results concerning the three-planet extrasolar system HD~37124 are presented and discussed. The best fit solutions found in earlier works are studied more closely. The system involves Jovian planets with similar masses. The orbits have moderate eccentricities, nevertheless the best fit solutions are found in dynamically active region of the phase space. The long term stability of the system is determined by a net of low-order two-body and three-body mean motion resonances. In particular, the three-body resonances may induce strong chaos that leads to self-destruction of the system after Myrs of apparently stable and bounded evolution. In such a case, numerically efficient dynamical maps are useful to resolve the fine structure of the phase space and to identify the sources of unstable behavior. ",
        "introduction": "{Understanding } the extrasolar planetary systems has became a major challenge for contemporary astronomy.  One of the most difficult problems in this field concerns the orbital stability of such systems.  Usually, the investigations of long-term evolution are the domain of direct, numerical integrations. The stability of extrasolar systems is often understood in terms of the Lagrange definition implying that orbits remain well bounded over an arbitrarily long time. Other definitions may be formulated as well, like the astronomical stability \\citep{Lissauer1999} requiring that the system persists over a very long, Gyr time-scale, or Hill stability \\citep{Szebehely1984} that requires the constant ordering of the planets. In our studies, we prefer a more formal and stringent approach related to the fundamental Kolmogorov-Arnold-Theorem (KAM), see \\cite{Arnold1978}. Planetary systems, involving a dominant mass of the parent star  and significantly smaller planetary masses,  are well modeled by close-to-integrable, Hamiltonian dynamical systems. It is well known, that their evolution may be quasi-periodic (with a discrete number of fundamental frequencies, forever stable),  periodic (or resonant; stable or unstable) or chaotic (with a continuous spectrum of frequencies, and unstable). In the last case, initially close phase trajectories diverge exponentially, i.e., their Maximum  Lyapunov Characteristic Exponent (MLCE, denoted also with $\\sigma$) is positive. In general, the distinction between regular and chaotic trajectories is a very difficult task that may be resolved only with numerical methods relying on efficient and accurate integrators of the equations of motion.  The main motivation of this paper is to describe numerical tools that are useful for studies of the dynamical stability and to apply them to the HD~37124 system \\citep{Vogt2005}. We recall the fundamentals of relative canonical  Poincar\\'e variables as  -- in our opinion -- one of the best frameworks for symplectic integrators. These {canonical} variables are well suited for the construction of a \\citet{LasRob} composition method that improves a classical Wisdom-Holman (W-H) algorithm \\citep{WH:91}. We supplement the integrator with a propagator of the associated symplectic tangent map that approximates the solution of variational equations \\citep{MikIn:99}. Finally, we compare two fast indicators that reveal the character of phase trajectories. The first one is a relatively simple method for resolving fundamental frequencies and spectral properties of a close-to-integrable Hamiltonian system -- a so called Spectral Number (SN), invented by \\citet{Michtchenko2001}. The second indicator belongs to the realm of the Lyapunov exponent based algorithms; we chose the numerical tool developed by \\cite{Cincotta2000,Cincotta2003} under the name of MEGNO. In this work, we refine the algorithm of MEGNO that makes explicit use of the symplectic tangent map \\citep{Gozdziewski2003b}.  As a non-trivial application of the presented numerical tools,  we consider the 3-planet system hosted by the HD~37124~star \\citep{Vogt2005}. It has been discovered by the radial velocity (RV) technique. The recent model of the RV observations of HD~37124 predicts three equal Jovian type planets with masses $\\sim 0.6$~m$_{\\idm{J}}$ in orbits with moderate eccentricities. In such a case, the application of symplectic integrators without regularization is particularly advantageous thanks to the numerical efficiency (long time-steps) and accuracy (the total energy does not have a secular error and the angular momentum integral is conserved). The number of multi-planet systems resembling the architecture of the Solar system increases\\footnote{For a recent statistics of the discoveries, see Jean Schneider's Extrasolar Planets Encyclopedia, http://exoplanets.eu.}. Hence, our approach may be useful in other cases. ",
        "conclusions": "The use of Poincar\\'e variables in the studies of  the dynamics of close-to integrable planetary systems offers many advantages. The variables are canonical and offer a simple form of a reduced Hamiltonian. The Hamiltonian can be split into a sum of two separately integrable parts: the Keplerian term and a small perturbation. As such, it can serve to construct a symplectic integrator based on any modern composition method, including the recent ones invented by \\citet{LasRob}. The tangent map computed with the same integration scheme provides an efficient way of computing the estimate of maximal Lyapunov exponent in terms of relatively recent fast indicator MEGNO. The method proves to be much more efficient than general purpose integrators (like the Bulirsh-Stoer-Gragg method). Besides, it provides the conservation of the integrals of energy and the angular momentum that is crucial for resolving the fine structure of the phase space. From the practical point of view, the symplectic algorithms are relatively simple for numerical implementation.  Using the numerical tools, we investigate the long term stability of extrasolar planetary system hosted by HD~37124. The orbital parameters in the set of our best, self-consistent Newtonian fits \\citep{Gozdziewski2006a}  are in accord with the discovery paper \\citep{Vogt2005}. Nevertheless, the observational window of the system is still narrow and the derivation of the model consistent with observations is difficult and, in fact, uncertain. The dynamical maps reveal that the relevant region of the phase space, in the neighborhood  of the mathematically best fit, is a strongly chaotic and unstable zone. The fitting algorithm (GAMP) that relies on eliminating strongly unstable fits founds solutions with a similar quality [in terms of $\\Chi$] that yields the formal solution. Moreover, they are shifted towards  larger semi-major axes and much smaller eccentricities  of the outermost planet. The orbital evolution of two outer planets is confined to a zone spanned by a number of low-order two-body and three-body MMRs. In particular, the three-body MMRs may induce very unstable behaviors that manifest themselves after many Myrs of an apparently stable and bounded evolution. To deal with such a problem, the stability of the best fits should be examined over a time-scale that is much longer than the one required when only the two-body MMRs are considered. In accord with the  dynamical maps, the  stable fits to the RV of HD~37124 should have small eccentricity of the outermost  planet~d, not larger than 0.2-0.3. Moreover, the stable configurations of the HD~37124 system are puzzling.  The best-fit mathematical three-planet model is surprisingly distant, in the phase space of initial conditions, from the zone of stable solutions consistent with the RV. It remains possible that other bodies are present in the system and the three-planet model is  not adequate to explain the RV variability, in spite that it provides apparently perfect fits. Yet, isolated initial conditions or even sets of best-fit solutions do not provide a complete answer on the system configuration. Then the fast indicator approach is essential and helpful to resolve the dynamical structure of the phase space.  The results of our experiments confirm and warn that all numerical methods should be applied with great care.  All symplectic methods are constant step integrators. In that case one should be cautious about the possibility of generating spurious resonance webs. A proper way to avoid them is to repeat computations with a different integration step in order to detect step-dependent patterns."
    },
    "0710/0710.0579_arXiv.txt": {
        "abstract": "{} {We aim to provide observational constraints on diffusion models that predict peculiar chemical abundances in the atmospheres of Am stars. We also intend to check if chemical peculiarities and slow rotation can be explained by the presence of a weak magnetic field.} {We have obtained high resolution, high signal-to-noise ratio spectra of eight previously-classified Am stars, two normal A-type stars and one Blue Straggler, considered to be members of the Praesepe cluster. For all of these stars we have determined fundamental parameters and photospheric abundances for a large number of chemical elements, with a higher precision than was ever obtained before for this cluster. For seven of these stars we also obtained spectra in circular polarization and applied the LSD technique to constrain the longitudinal magnetic field.} {No magnetic field was detected in any of the analysed stars. HD~73666, a Blue Straggler previously considered as an Ap\\,(Si) star, turns out to have the abundances of a normal A-type star. Am classification is not confirmed for HD~72942. For HD~73709 we have also calculated synthetic $\\Delta a$ photometry that is in good agreement with the observations. There is a generally good agreement between abundance predictions of diffusion models and values that we have obtained for the remaining Am stars. However, the observed Na and S abundances deviate from the predictions by 0.6 dex and $\\geq$0.25 dex respectively. Li appears to be overabundant in three stars of our sample.} {} ",
        "introduction": "Main sequence A-type stars present spectral peculiarities, usually interpreted as due to peculiar photospheric abundances and abundance distributions which are believed to be produced by the interaction of a large variety of physical processes (e.g. diffusion, magnetic field, pulsation and various kinds of mixing processes).  An interesting problem that has yet to be addressed is how these peculiarities change during main sequence evolution. The chemical composition of field A-type stars have been studied by several authors, e.g. \\citet{hillland1993}, \\citet{adelman2000}. However, it is not straightforward to use the results of these investigations to study how photospheric chemistry evolves during a star's main sequence life. First, the original composition of the cloud from which stars were born is not known and is likely somewhat different for each field star. It is therefore not possible to discriminate between evolutionary effects and differences due to original chemical composition. Secondly, it is difficult to estimate the age of field stars with the precision necessary for such evolutionary studies \\citep[for a discussion of this problem see][]{stefano2006}.  From this point of view, A-type stars belonging to open clusters are much more interesting objects. Compared to field stars, A-type stars in open clusters have three very interesting properties: \\begin{list}{-}{} \\item they were all presumably born from the interstellar gas with an approximately uniform composition; \\item they all have approximately the same age (to within a few Myr); \\item their age can be determined much more precisely than for field stars. \\end{list}  Few abundance analyses of A-type stars in open clusters have been carried out. Those that have been published have usually focused on a limited numbers of stars.  \\citet{monier1999} have determined the abundances of eleven chemical elements for a large sample of stars regularly distributed in spectral type along the main sequence in order to sample the expected masses uniformly.  All these stars were analysed in a uniform manner using spectrum synthesis. \\citet{ch2006} have performed a detailed abundance analysis for five A-type stars of the young open cluster IC~2391. \\citet{folsom2007} have performed a detailed abundance analysis for four Ap/Bp stars and one normal late B-type star of the open cluster NGC~6475.  A goal of this programme is to determine photospheric abundance patterns in A-type star members of clusters of different ages. This is crucial in order to: \\textit{i)} investigate the chemical differences between normal and peculiar stars inside the same cluster, \\textit{ii)} study the evolution with time of abundance peculiarities by studying clusters of various ages, \\textit{iii)} set constraints on the hydrodynamical processes occurring at the base of the convection zone in the non magnetic stars and \\textit{iv)} study the effects of diffusion in the presence of a magnetic field for the magnetic (Ap) stars in the cluster. The abundance analysis will be performed in an homogeneous way applying a method described in this first work.  Praesepe (NGC~2632), a nearby intermediate-age open cluster \\citep[$\\log t = 8.85\\pm 0.15$,][]{gonzalez}, is an especially interesting target because it includes a large number of A-type stars, among which are many Am stars. Furthermore, because the cluster is relatively close to the sun \\citep[d~=~180~$\\pm$~10 pc,][]{robichon1999}, many of the member A-type stars are bright enough to allow us to obtain high resolution spectra with intermediate class telescopes.  We dedicate this first paper to the Am stars of the Praesepe cluster, searching for magnetic fields in these objects and discussing the differences between \"normal\"\\footnote{We consider as \"normal\" A-type stars all the A-type stars that are classified neither as Am nor Ap} A-type stars and Am stars in the cluster.  We also compare our results with previous works and with theoretical chemical evolution models. In particular we take into account diffusion models by \\citet{richer2000}. We want to provide observational constraints to the theory of the evolution of the abundances in normal and chemically peculiar stars. Our detailed abundance analysis could provide information about the turbulence occurring in the outer stellar regions in Am stars with well determined age. In particular, our analysis can give constraints to define the depth of the zone mixed by turbulence, since it is the only parameter characterising turbulence \\citep{richer2000}. A systematic abundance analysis of normal and peculiar stars in clusters could provide information on the origin of the mixing process and show if only turbulence is needed to explain abundance anomalies, or if other hydrodynamical processes occur.  We tackle this problem using new and more precise new-generation spectrographs providing a wider wavelength coverage together with newer analysis codes and procedures (e.g. Least-Squares Deconvolution and synthetic line profile fitting instead of equivalent width measurements.)  The observed stars, the instruments employed and the target selection are described in Sect.~\\ref{observations}. The data reduction and a discussion of the continuum normalisation are provided in Sect.~\\ref{reduction,norm}. In Sect.~\\ref{LLmodels} and \\ref{abundanceanalysis} we describe the models and the procedure used to perform the abundance analysis. Our results are summarised in Sect.~\\ref{results}. Discussion and conclusions are given in Sect.~\\ref{discussion} and \\ref{conclusion} respectively. ",
        "conclusions": ""
    },
    "0710/0710.2783_arXiv.txt": {
        "abstract": "For the 2dFGRS we study the properties of voids and of fainter galaxies within voids that are defined by brighter galaxies. Our results are compared with simulated galaxy catalogues from the Millenium simulation coupled with a semianalytical galaxy formation recipe. We derive the void size distribution and discuss its dependence on the faint magnitude limit of the galaxies defining the voids. While voids among faint galaxies are typically smaller than those among bright galaxies, the ratio of the void sizes to the mean galaxy separation reaches larger values. This is well reproduced in the mock galaxy samples studied. We provide analytic fitting functions for the void size distribution. Furthermore, we study the galaxy population inside voids defined by objects with $B_J -5\\log{h}< -20$ and diameter larger than 10 \\hMpc. We find a clear bimodality of the void galaxies similar to the average comparison sample. We confirm the enhanced abundance of galaxies in the blue cloud and a depression of the number of red sequence galaxies. There is an indication of a slight blue shift of the blue cloud. Furthermore, we find that galaxies in void centers have higher specific star formation rates as measured by the $\\eta$ parameter. We determine the radial distribution of the ratio of early and late type galaxies through the voids. We find and discuss some differences between observations and the Millenium catalogues. ",
        "introduction": "The large-scale galaxy distribution is highly inhomogeneous. We observe groups, clusters and superclusters of galaxies and large voids. During last decades, much attention was paid on the analysis of bound structures as groups and clusters. Recently, new superclusters catalogues were constructed from the 2dFGRS and compared with large cosmological simulations \\citep{Einasto07a, Einasto07b}. In a complement, there are large regions in the universe without bright galaxies, so called cosmic voids. Early on very large voids over 50 \\hMpc diameter were found by \\citet{Gregory78} and \\citet{Kirshner81}. More common are voids with diameters of about 10 \\hMpc that fill most of cosmic space. The explanation of such structures is not obvious. According to the standard paradigm of cosmological structure formation, negative potential wells from primordial inhomogeneities attract all matter in bound structures. In the same way, positive potential perturbations expel matter, but observed voids are too large for completely emptying. Therefore, in addition to the dilution of matter, the galaxy formation probability should be suppressed in underdense regions, cp. e.g. \\citet{Lee98, Madsen98}. Recently, \\citet{Sheth04} and \\citet{Furlanetto06} applied these ideas within the excursion set formalism of gravitational instability. These analytical theories derived void size distributions that are peaked typically at diameters below 10 \\hMpc which seem to be smaller than observed void sizes, cp. \\citet{Mueller00}, and void sizes in CDM-simulations coupled with semianalytical galaxy formation models, \\citet{Benson03}.  Voids were routinely identified in all wide-field redshift surveys as the CfA \\citep{deLapparent86, Vogeley94}, the SSRS2 \\citep{Elad97}, the LCRS \\citep{Mueller00, Arbabi02}, the IRAS-survey \\citep{Elad00}, the 2dFGRS \\citep{Hoyle04, Croton04, Patiri06a}, the SDSS \\citep{Rojas04, Rojas05, Patiri06b}, and the DEEP2 survey with an analysis of voids up to redshift $z \\approx 1$ \\citep{Conroy05}. However, many void searches are only devoted to the identification of large voids, other void finders depend on special procedures as firstly identifying wall galaxies by an overdensity criterion and then looking for voids bounded by wall galaxies \\citep{Elad97, Hoyle04}. Furthermore, the void search depends on the galaxy sample used for defining voids, in particular on the limiting magnitude of the galaxy sample. In an influential paper, \\citet{Peebles01} derived from nearest neighbor statistics that galaxies of different brightness respect the same voids. He claimed that this contradicts the standard CDM scenario of galaxy and structure formation that seem to predict a hierarchy of galactic structures with smaller structure for fainter objects sitting in less massive dark matter halos, i.e also smaller voids for fainter objects. In a follow up theoretical study, \\citet{Mathis02} showed from high-resolution simulation that voids defined by bright galaxies are also underdense in faint galaxies, i.e. that bright and faint galaxies respect similar voids. We want to take up this question since our earlier studies of voids in LCRS and in LCDM-simulations \\citep{Mueller00, Arbabi02} showed a dependence of the void size distribution on the brightness limit of the galaxies under study and a characteristic void size scaling relation. \\citet{Benson03} confirmed this scaling relation in simulated galaxy distributions, but the quantitative parameters were different. They suspected differences in the void search algorithms as reason, but we suspect that the effective 2-dimensional nature of the LCRS is the most likely cause. But their demand for using the same void search algorithm both in data and simulations seems a prerequisite for trustworth results. More recently, Colberg et al. (2007 in preparation) compared different void search algorithms and found that most proposed algorithm find comparable locations and sizes of large voids. This is very likely not the case for the large number of small voids that fill a significant part of space. Therefore we shall present in this study an comparable analysis of voids both in simulations and in the data that explores the detailed void size distribution in dependence on the faint brightness limit of the galaxies defining the voids.  We shall use for our study the 2dFGRS \\citep{Cole05} thereby coming back to the property of the self-similarity of the void statistics. It tells that the void size distribution depends on the mean galaxy separation, in such a way that brighter galaxies define larger voids than fainter ones. Even if voids in the 2dFGRS were previously analyzed \\citep{Hoyle04, Patiri06a}, this concerned mainly large voids and not the detailed void statistics proposed by us previously. \\citet{Croton04} provided a detailed study of the void probability distribution for the 2dFGRS which is related to the void size distribution but it provides an different statistics. Essentially the void probability distribution is a weighted sum over the void size distribution \\citep{Otto86}.  The 2dFGRS is a densely sampled survey with a compact survey geometry. This is of advantage for the question of the dependence of the void sizes on the galaxy magnitudes defining voids. We shall derive phenomenological fits to the void size distribution that will be compared with simulation results. Furthermore, we shall take up the question of the faint galaxies within voids. Thereby we cut both questions, the matter content inside large underdense regions in the universe, and the change of galaxy properties. The color distribution of galaxies in the 2dFGRS employs SuperCOSMOS data \\citep{Hambly01} for the $R$-band. We will find a clear bimodality in the void galaxies so far only studied in detail for the SDSS \\citep{Rojas05, Patiri06b} but not yet for the 2dFGRS.  We shall evaluate our void analysis with model galaxy samples constructed from the Millenium simulation of \\citet{Springel05} and from semianalytical galaxy formation theory applied to the numerical merger trees \\citep{Croton06}. We analyzed specific galaxy properties within voids and found results that can be qualitatively described by the model samples. A quantitative comparison of the galaxy color distribution and the star formation efficiency hints at certain differences between observations and simulations. Tentatively we connect it with specific environmental properties of galaxy formation in underdense regions. In particular, major mergers, galaxy harassment, tidal and ram-pressure stripping will not be as effective there as in more dense regions of the universe \\citep{Avila05, Maulbetsch06}.  The outline of the paper is as follows: First we describe the galaxy extraction from the 2dFGRS, and in Section 3 we provide some details of the galaxy mock data. In Section 4 we describe our void search algorithm and in Section 5 we provide our results. Section 6 is devoted to a discussion and in the final Section we draw our conclusions. ",
        "conclusions": "In this paper we have studied properties of voids in the 2dFGRS and compared them with those predicted by semi-analytical models of galaxy formation. In particular, we were interested in the distribution of void sizes in the galaxy distribution as a function of the faint limiting magnitude of the sample. We established an almost linear dependence of the void size distribution on the mean separation $\\lambda$ of the galaxies in the sample. A similar relation was found for the magnitude dependence of the correlation length of galaxies and galaxy groups \\citep{Bahcall92, Yang05}. It is a natural consequence of the halo model of gravitational clustering and the resulting void statistics \\citep{Tinker06}.  In addition to the self-similarity, we have found that the voids among faint galaxies extend to relatively larger scales when devided by the mean void size. In the scaled radius $R/\\lambda$, the voids size distribution of faint galaxies has a longer tail at large radii than those for brighter galaxies. This indicates that the faint galaxies trace the general spatial distribution of the cosmic web of brighter galaxies.  For the 2dFGRS, we find more large voids at larger distances from the observer. This is due to the conelike structure of the survey. In our mock sample, we use exactly the same geometry of the observed samples. Thereby we found that the number density of the largest voids can be underestimated by up to 20\\%.  We confirm the previous findings by \\citet{Grogin00, Hogg04, Rojas05, Croton05, Patiri06b}, that in lower-density regions, the galaxy population is dominated by blue actively star forming galaxies. We fitted the $B_J-R$ color distribution of 2dFGRS void galaxies by double Gaussian distributions. There are significantly more void galaxies in the blue cloud than in the general field, and the red sequence is strongly suppressed. In addition, we have found the indication that the blue population of galaxies in void centers is slightly {\\it bluer} than the field population, not just more numerous. The shift is only minor and almost of the same order as the maximal quoted uncertainty by \\cite{Cole05}, so it remains unclear how robust is this effect. It is not reproduced by the semi-analytical models for galaxy formation of the Millenium simulation. Also the red sequence in the Millenium catalogue has a significantly tighter spread. This may due to a too sharp cutoff of the star formation activity in the semianalytical models which seems to be unrealistic.  The radial distribution of the ratio of early and late type galaxies inside voids coincides quantitatively for the 2dFGRS and the Millenium catalogue, but in the data, the fraction of E-galaxies stays smaller (and the fraction of S-galaxies larger) further outwards beyond the boundary of the voids. The abundance of star-forming galaxies is higher in the 2dFGRS voids than in the semianalytical model."
    },
    "0710/0710.3785_arXiv.txt": {
        "abstract": "% In order to distinguish between the various components of massive star forming regions (i.e. infalling, outflowing and rotating gas structures) within our own Galaxy, we require high angular resolution observations which are sensitive to structures on all size scales. To this end, we present observations of the molecular and ionized gas towards massive star forming regions at 230 GHz from the SMA (with zero spacing from the JCMT) and at 22 and 23 GHz from the VLA at arcsecond or better resolution.  These observations (of sources such as NGC7538, W51e2 and K3-50A) form an integral part of a multi-resolution study of the molecular and ionized gas dynamics of massive star forming regions (i.e. Klaassen \\& Wilson 2007).  Through comparison of these observations with 3D radiative transfer models, we hope to be able to distinguish between various modes of massive star formation, such as ionized or halted accretion (i.e Keto 2003 or Klaassen et al. 2006 respectively). ",
        "introduction": "At the large distances to massive star forming regions, we do not yet quite have the resolution necessary to determine whether or not they can form via a scaled up version of the processes responsible for the formation of low mass stars.  The energetics (i.e. luminosity, outflow energy and accretion rates) are orders of magnitude higher than those seen in low mass star forming regions (e.g. Arce et al. 2007), and the outward radiation and thermal pressures produced when the massive star begins ionizing its surroundings are enormous.  Yet, collimated outflows and both ionized and molecular infall signatures are present in high mass star forming regions just as they are in their lower mass counterparts.  If we assume that massive stars can form in a single accretion event similar to their low mass counterparts, how then can accretion continue beyond the formation of the H{\\sc II} region? Does accretion continue in an ionized form? Does it stop early in the star's evolution? Does it continue through a disk (not yet observed around forming O stars)?  G10.62-0.38 (hereafter G10.6) is a well studied ultracompact H{\\sc II} (UCH{\\sc II}) region (Ho et al. 1981, Keto et al. 1988, etc) through which the surrounding gas is falling in, becoming ionized, and continuing to fall in towards the central star.  The high resolution observations of Keto \\& Wood (2006) in H66$\\alpha$ taken at the VLA show an infall signature in their PV diagram, and so it appears as though accretion onto the protostar is, in this case, ionized and is continuing beyond the formation of the UCH{\\sc II} region.  G5.89-0.39 (hereafter G5.89) is the site of an UCH{\\sc II} region (the archetypal shell UHCII region from Wood \\& Churchwell 1989) and multiple outflows depending on the molecular tracer observed. This region appears to be the formation site for a number of massive stars including the O5V star which is believed to power the H{\\sc II} region (Feldt et al. 2003) and the 1.3 mm continuum source believed to be powering the SiO outflow detected by Sollins et al. (2004).  Recent maser emission studies of this region (Stark et al. 2007, Fish et al. 2005) also suggest that the two sources are independent. The star identified in the mid-infrared by Feldt et al. (2003), corresponding to `G5.89 center' in Stark et al (2007), appears to be less embedded than the Sollins et al. (2004) 1.3 mm source, suggesting it is possibly more evolved than the 1.3 mm source which drives the accelerating SiO outflow.  To date, there are no published CO maps of this region at high enough resolution to determine which of these two sources (separated by $\\sim2''$) lies at the center of the large scale decelerating outflow. However, the velocity gradient in the 1667 MHz maser emission towards G5.89 center (Stark et al. 2007) is consistent with the direction of the large scale east-west CO outflow observed by Klaassen et al. (2006) and Watson et al. (2007).    Klaassen et al. (2006) suggested that accretion onto the source powering the large scale outflow may have halted at the onset of the H{\\sc II} region. This conclusion appears consistent with the O5 star being responsible for the H{\\sc II} region, `G5.89 center' maser emission and the large scale CO outflow,  while the 1.3 mm continuum source from Sollins et al. (2004) is a younger protostar responsible for the SiO outflow.   These studies of G10.6 and G5.89 suggest that high neutral accretion rates (up to 10$^{-2}$ M$_{\\odot}$/yr, Edgar \\& Clarke 2003) and continued ionized accretion can both account for the formation of massive stars. Yet, these conclusions are based on the in depth studies of two individual regions, and from these two regions, we cannot make conclusive statements about the general nature of massive star formation.  In order to determine whether there is a preferred method, we need to observe these processes in a number of regions, and build up statistics about whether/how accretion onto a massive protostar occurs beyond the onset of dynamical expansion in the UCH{\\sc II} regions. ",
        "conclusions": ""
    },
    "0710/0710.3099_arXiv.txt": {
        "abstract": "Short--hard Gamma Ray Bursts (SGRBs) are currently thought to arise from gravitational wave driven coalescences of double neutron star systems forming either in the field or dynamically in globular clusters. For both channels we fit the peak flux distribution of BATSE SGRBs to derive the local burst formation rate and luminosity function. We then compare the resulting redshift distribution with {\\it Swift} 2-year data, showing that both formation channels are needed in order to reproduce the observations. Double neutron stars forming in globular clusters are found to dominate the distribution at $z\\lsim 0.3$, whereas the field population from primordial binaries can account for the high--$z$ SGRBs. This result is not in contradiction with the observed host galaxy type of SGRBs. ",
        "introduction": "The afterglows of several Short Gamma Ray Bursts (SGRBs) have recently  been localised on the sky with {\\it Swift} (Gehrels et al. 2004) and {\\it HETE-II} (Lamb et al. 2004), allowing for the determination of their redshift and host galaxy (Nakar 2007; Berger et al. 2007a, and references therein). For SGRBs, the data give  mean redshift $z\\sim 0.6,$ and early as well as late type galaxy hosts.  The current idea on the nature of the SGRBs is that they arise from gravitational wave  driven mergers of neutron star-neutron star and/or neutron star-black hole binaries (Belczynski et al. 2006;  Nakar 2007). Double neutron star binaries (DNSs) are observed in the field of the Milky Way and in globular clusters, so  they can form through two different channels: (i) evolution of primordial massive star binaries, formed in the galactic field (field scenario; Narayan, Paczynski \\& Piran 1992), and (ii) dynamical formation\\footnote{Formation of primordial DNSs may not be efficient in globular clusters (Ivanova et al 2008). }, through three-body interactions in globular clusters (GCs) involving the exchange of a light star companion of a neutron star with an incoming isolated neutron star (GC scenario; Grindlay, Portegies Zwart \\& Mc Millan 2006).  In this Letter, we derive the local SGRB formation rate and luminosity function by fitting the peak flux distribution of the large sample of BATSE SGRBs (Paciesas et al. 1999) exploring both formation scenarios. We compare the model results with the redshift distribution of the sample of {\\it Swift} SGRBs in the first two years of operation, with the aim of disentangling the relative importance of the two SGRB populations at different cosmic epochs. ",
        "conclusions": "We have shown that current {\\it Swift} data seem to point towards a dual nature of SGRB formation: low--$z$ SGRBs may arise from the coalescence of DNSs forming in GCs through dynamical collisions, whereas high--$z$ bursts can represent the end--result of DNS mergers born in the field from primordial massive binary stars. A way to falsify this bimodality is by considering the correlation between SGRBs and the host galaxy type (Gam-Yam et al. 2005; Zheng \\& Ramirez-Ruiz 2007; Berger et al. 2007b; Shin \\& Berger 2007).  We expect that SGRBs from the GC channel should be found preferentially in early type galaxies where the bulk of GCs resides, although we can not completely exclude a few in spirals.  On the contrary, field SGRBs can be hosted in both early and late type galaxies given the wide distribution of delay times (Section 2). Since the dynamical channel produces SGRBs at low--$z$, we expect to find a excess of SGRBs in early type galaxies below $z<0.3$.  At present it is still premature to exploit this issue in a statistically meaningful manner and we may just confine ourselves to the analysis of SGRBs in the current limited sample. Up to now, four SGRBs have been localised in early-type galaxies (GRB~050509B, 050724, 050813, and 060502B), whereas six have secure late-type hosts (GRB~051221A, 060801, 061006, 061210, and 061217, Berger 2007; {\\it HETE-II} GRB~050709 at $z=0.161$; Fox et al. 2005; Covino et al. 2006). Fig.~2 shows the fraction of SGRB hosts in early (shaded area) and late (white area) below $z=0.3$ and above, respectively. We note indeed that low--$z$ SGRBs are associated preferentially to early-type galaxies while the high--$z$ ones are located in late-type.  The large off--set of GRB~050509B and 060502B may support further the association of these SGRBs with DNSs formed dynamically in GCs.    \\begin{figure} \\begin{center} \\centerline{\\psfig{figure=host.ps,height=8cm}} \\caption{Fraction of SGRBs hosted in early--type (shaded area) and late--type (white area) galaxies below $z=0.3$ and above.} \\end{center} \\end{figure}"
    },
    "0710/0710.1113_arXiv.txt": {
        "abstract": "We investigate torsional Alfv\\'{e}n oscillations of relativistic stars with a global dipole magnetic field, via two-dimensional numerical simulations. We find that a) there exist two families of quasi-periodic oscillations (QPOs) with harmonics at integer multiples of the fundamental frequency, b) the lower-frequency QPO is related to the region of closed field lines, near the equator, while the higher-frequency QPO is generated near the magnetic axis, c) the QPOs are long-lived, d) for the chosen form of dipolar magnetic field, the frequency ratio of the lower to upper fundamental QPOs is $\\sim 0.6$, independent of the equilibrium model or of the strength of the magnetic field, and e) within a representative sample of equations of state and of various magnetar masses, the Alfv\\'{e}n QPO frequencies are given by accurate empirical relations that depend only on the compactness of the star and on the magnetic field strength. The lower and upper QPOs can be interpreted as corresponding to the edges or turning points of an Alfv\\'{e}n continuum, according to the model proposed by Levin (2007). Several of the low-frequency QPOs observed in the X-ray tail of SGR 1806-20 can readily be identified with the Alfv\\'{e}n QPOs we compute. In particular, one could identify the 18Hz and 30Hz observed frequencies with the fundamental lower and upper QPOs, correspondingly, while the observed frequencies of 92Hz and 150Hz are then integer multiples of the fundamental upper QPO frequency (three times and five times, correspondingly). With this identification, we obtain an upper limit on the strength of magnetic field of SGR 1806-20 (if is dominated by a dipolar component) between $\\sim3$ and $7\\times 10^{15}$G. Furthermore, we show that an identification of the observed frequency of 26Hz with the frequency of the fundamental torsional $\\ell=2$ oscillation of the magnetar's crust is compatible with a magnetar mass of about $1.4$ to 1.6$M_\\odot$ and an EOS that is \tvery stiff (if the magnetic field strength is near its upper limit) or moderately stiff (for lower values of the\tmagnetic field). ",
        "introduction": "\\label{sec:Intro}  The phenomenon of Soft Gamma Repeaters (SGRs) may allow us in the near future to determine fundamental properties of strongly magnetized, compact stars. Already, there exist at least two sources in which quasi-periodic oscillations (QPOs) have been observed in their X-ray tail, following the initial discovery by \\cite{Israel2005}, see \\cite{WS2006} for a recent review. The frequency of many of these oscillations is similar to what one would expect for torsional modes of the solid crust of a compact star. This observation is in support of the proposal that SGRs are magnetars (compact objects with very strong magnetic fields) \\citep{DT1992}. During an SGR event, torsional oscillations in the solid crust of the star could be excited \\citep{Duncan1998}, leading to the observed frequencies in the X-ray tail. However, not all of the observed frequencies fit the above picture. For example, the three lowest observed frequencies for SGR 1806-20 are 18, 26, 30Hz. Only one of these could be the fundamental, $\\ell=2,m=0$ torsional frequency of the crust, as the first overtone has a much higher frequency. \\cite{Levin2006} stressed the importance of crust-core coupling by a global magnetic field and of the existence of an Alfv\\'en continuum, while \\citet{GSA2006} considered model with simplified geometry, in which Alfv\\'{e}n oscillations form a discrete spectrum of normal modes, that could be associated with the observed low-frequency QPOS. In  \\cite{Levin2007}, the existence of a continuum was stressed further and it was shown that the edges or turning points of the continuum can yield long-lived QPOs. In addition, numerical simulations showed that drifting QPOs within the continuum become amplified near the frequencies of the crustal normal modes. Within this model, Levin suggested a likely identification of the 18Hz QPO in SGR 1806-20 with the lowest frequency of the MHD continuum or its first overtone. The above results were obtained in toy models with simplified geometry and Newtonian gravity.  In this Letter, we perform two-dimensional numerical simulations of linearized Alfv\\'{e}n oscillations in magnetars. Our model improves on the previously considered toy models in various ways: relativistic gravity is assumed, various realistic equations of state (EOS) are considered and a consistent dipolar magnetic field is constructed. We do not consider the presence of a solid crust, but only examine the response of the ideal magnetofluid to a chosen initial perturbation.   Spherical stars have generally two type of oscillations, {\\it spheroidal} with polar parity and {\\it toroidal} with axial parity. The observed QPOs in SGR X-ray tails may originate from toroidal oscillations, since these could be excited more easily than poloidal oscillations, because they do not involve density variations. In Newtonian theory, there have been several investigations of torsional oscillations in the crust region of neutron stars (see e.g., \\cite{Lee2007} for reference). On the other hand, only few studies have taken general relativity into account \\citep{Messios2001,Sotani2007a,Sotani2007b,SA2007,Vavoulidis2007}.  SGRs produce giant flares with peak luminosities of $10^{44}$ -- $10^{46}$ erg/s, which display a decaying tail for several hundred seconds.  Up to now, three giant flares have been detected, SGR 0526-66 in 1979, SGR 1900+14 in 1998, and SGR 1806-20 in 2004.  The timing analysis of the latter two events revealed several QPOs in the decaying tail, whose frequencies are approximately 18, 26, 30, 92, 150, 625, and 1840 Hz for SGR 1806-20, and 28, 53, 84, and 155 Hz for SGR 1900+14, see \\cite{WS2006}.  In \\cite{Sotani2007a} (hereafter Paper~I), it was suggested that some of the observational data of SGRs could agree with the crustal torsional oscillations, if, e.g., frequencies lower than 155 Hz are identified with the fundamental oscillations of different harmonic index $\\ell$, while higher frequencies are identified with overtones. However, in Paper~I and above, it will be quite challenging to identify all observed QPO frequencies with only crustal torsional oscillations.  For example, it is difficult to explain all of the frequencies of 18, 26 and 30 Hz for SGR 1806-20 with crustal models, because the actual spacing of torsional oscillations of the crust is larger than the difference between these two frequencies. Similarly, the spacing between the 625Hz and a possible 720Hz QPO in SGR 1806-20 may be too small to be explained by consecutive overtones of crustal torsional oscillations.  One can notice, however, that the frequencies of 30, 92 and 150 Hz in SGR 1806-20 are in near {\\it integer ratios}. As we will show below, the numerical results presented in this Letter are compatible with this observation, as we find two families of QPOs (corresponding to the edges or turning points of a continuum) with harmonics at near integer multiples. Furthermore, our results are compatible with the ratio of 0.6 between the 18 and 30Hz frequencies, if these are identified, as we suggest, with the edges (or turning points) of the Alfv\\'{e}n continuum. With this identification, we can set an upper limit to the dipole magnetic field of $\\sim3$ to $~7\\times 10^{15}$G. If the drifting QPOs of the continuum are amplified at the fundamental frequency of the crust, and the latter is assumed to be the observed 26Hz for SGR 1806-20, then our results  are compatible with a magnetar mass of about $1.4$ to 1.6$M_\\odot$ and an EOS that is very stiff (if the magnetic field strength is near its upper limit) or moderately stiff (for lower values of the magnetic field).  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 $(-,+,+,+)$.      \\section[]{Numerical Setup} \\label{sec:II}  The general-relativistic equilibrium stellar model is assumed to be spherically symmetric and static, i.e. a solution of the well-known TOV equations for and perfect fluid and metric described the line element \\begin{equation} ds^2 = -e^{2\\Phi(r)}dt^2 + e^{2\\Lambda(r)}dr^2 + r^2(d\\theta^2 + \\sin^2\\theta d\\phi^2). \\end{equation} We neglect the influence of the magnetic field on the structure of the star, since the magnetic field energy, ${\\cal E_M}$, is orders of magnitudes smaller than the gravitational binding energy, ${\\cal E_G}$, for magnetic field strengths considered realistic for magnetars, ${\\cal E_M}/{\\cal E_G}\\approx 10^{-4}(B/(10^{16} {\\rm G}))^2$. For simplicity, we assume that the magnetic field is a pure dipole (toroidal magnetic fields will be treated elsewhere, see \\cite{Sotani2007c}. Details on the numerical method for constructing the magnetic field, as well as representative figure of the magnetic field lines, can be found in Paper~I.  MHD oscillations of the above equilibrium model are described by the linearized equations of motion and the magnetic induction equations, presented in detail in Papers~I and II (we neglect perturbations in the spacetime metric, as these couple weakly to toroidal modes in a spherically symmetric background). The perturbative equations presented in Papers~I and II can readily be converted from an eigenvalue problem to the form of a two-dimensional time-evolution problem, by defining a displacement ${\\cal Y}(t,r,\\theta)$ due to the toroidal motion (the coefficient of shear viscosity is set to $\\mu=0$, as we neglect the presence of a solid crust in the present work). The contravariant {\\it coordinate component} of the perturbed four-velocity, $\\delta u^{\\phi}$, is then related to time derivative of ${\\cal Y}$ through \\begin{equation} \\delta u^{\\phi} = e^{-\\Phi}\\partial_t {\\cal Y}(t,r,\\theta). \\end{equation} The two-dimensional evolution equation for ${\\cal Y}(t,r,\\theta)$ is \\begin{eqnarray} {\\cal A}_{tt}\\frac{\\partial^2 {\\cal Y}}{\\partial t^2} &=& {\\cal A}_{20}\\frac{\\partial^2 {\\cal Y}}{\\partial r^2} + {\\cal A}_{11} \\frac{\\partial^2 {\\cal Y}}{\\partial r \\partial \\theta} + {\\cal A}_{02} \\frac{\\partial^2 {\\cal Y}}{\\partial \\theta^2} \\nonumber \\\\ &+& {\\cal A}_{10} \\frac{\\partial {\\cal Y}}{\\partial r} + {\\cal A}_{01} \\frac{\\partial {\\cal Y}}{\\partial \\theta} + \\varepsilon_D {\\cal D}_4 {\\cal Y}, \\label{eq2D} \\end{eqnarray} where ${\\cal A}_{tt}$, ${\\cal A}_{20}$, ${\\cal A}_{11}$, ${\\cal A}_{02}$, ${\\cal A}_{10}$, and ${\\cal A}_{01}$ are functions of $r$ and $\\theta$, given by \\begin{eqnarray} {\\cal A}_{tt} &=& \\left[\\epsilon + p + \\frac{{a_1}^2}{\\pi r^4} \\cos^2\\theta + \\frac{{{a_1}'}^2}{4\\pi r^2} e^{-2\\Lambda}\\sin^2\\theta \\right] e^{-2(\\Phi - \\Lambda)}, \\\\ {\\cal A}_{20} &=& \\frac{{a_1}^2}{\\pi r^4}\\cos^2\\theta , \\\\ {\\cal A}_{11} &=& -\\frac{a_1 {a_1}'}{\\pi r^4}\\sin\\theta \\cos\\theta, \\\\ {\\cal A}_{02} &=& \\frac{{{a_1}'}^2}{4\\pi r^4}\\sin^2\\theta, \\\\ {\\cal A}_{10} &=& \\left(\\Phi' - \\Lambda' \\right) \\frac{{a_1}^2}{\\pi r^4} \\cos^2\\theta + \\frac{a_1 {a_1}'}{2\\pi r^4} \\sin^2\\theta, \\\\ {\\cal A}_{01} &=& \\left[\\frac{a_1}{\\pi r^4}\\left(2\\pi j_1 - \\frac{a_1}{r^2}\\right)e^{2\\Lambda} + \\frac{3{{a_1}'}^2}{4\\pi r^4} \\right]\\sin\\theta\\cos\\theta, \\end{eqnarray} and where $a_1(r)$ and $j_1(r)$ are the radial components of the electromagnetic four-potential and the four-current, respectively. In the above equations, a prime denotes a partial derivative with respect to the radial coordinate.  In our numerical scheme, we employ the 2nd-order, iterative Crank-Nicholson scheme. A numerical instability, which sets in after many oscillations, was treated by adding a 4th-order Kreiss-Oliger dissipation term \\citep{KO1973}, shown as $\\varepsilon_D {\\cal D}_4 {\\cal Y}$ in Equation (\\ref{eq2D}) above. We experimented with various values of the dissipation coefficient $\\varepsilon_D$ and found the evolution to be stable for values as small as a few times $10^{-5}$. We verified that in this limit, the solution becomes independent of the strength of the numerical dissipation, as the ${\\cal D}_4$ Kreiss-Oliger dissipation operator introduces an error of higher order than the 2nd-order iterative Crank-Nicholson scheme. The numerical grid we use is equidistant, covering only the interior of the star, with (typically) 50 radial zones and 40 angular zones (we also compared our results to simulations with 100x80 points). The boundary conditions are: (a) ${\\cal Y}=0$ at $r=0$ (regularity), (b) ${\\cal Y}_{,r}=0$ at $r=R$ (vanishing traction) (c) ${\\cal Y}_{,\\theta}=0$ at $\\theta=0$ (axisymmetry) and (d) ${\\cal Y}=0$ at $\\theta=\\pi/2$ (equatorial plane symmetry of the $\\ell=2$ initial data). We obtain the same results if you use a grid that extends to $\\theta=\\pi$. We have also evolved initial data with $\\ell=3$, for which the appropriate boundary condition at $\\theta=\\pi/2$ is ${\\cal Y}_{,\\theta}=0$.     \\section[]{Alfv\\'{e}n QPOs} \\label{sec:III}  \\begin{figure} \\begin{center} \\includegraphics[width=55mm]{Fig-fft-e0001-dy1.eps} \\end{center} \\caption{FFT of the MHD oscillations at $\\theta\\sim 0$ (magnetic axis). Three lines in each figure correspond to different radial positions $r\\sim 0$, $R/2$, and $R$). A fundamental ($n=0$) QPO and several overtones (nearly integer multiples) are clearly present.} \\label{Fig:FFT} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[width=55mm]{Fig-fft-e0001-dy3.eps} \\end{center} \\caption{Same as Fig. \\ref{Fig:FFT}, but at  $\\theta\\sim \\pi/2$ (magnetic equator). A second family of QPOs is present. Arrows indicate several continuous parts, of which only the first is distinct from the others, which partially overlap.} \\label{Fig:FFT2} \\end{figure}   \\begin{figure} \\begin{center} \\begin{tabular}{cc} \\includegraphics[width=36mm]{e000002-dy-u0-unitbase.eps} & \\includegraphics[width=36mm]{e000002-dy-u1-unitbase.eps} \\\\ \\includegraphics[width=36mm]{e000002-dy-l0-unitbase.eps} & \\includegraphics[width=36mm]{e000002-dy-l1-unitbase.eps} \\\\ \\end{tabular} \\end{center} \\caption{Distribution of effective amplitude of several Alfv\\'{e}n QPOs (see text for details). The grayscale map varies from white (zero amplitude) to black (maximum amplitude).} \\label{Fig:Effective} \\end{figure}  As a fist equilibrium model, we examined a polytropic stellar model with $\\Gamma=2$, whose mass and radius are $M=1.4M_{\\odot}$ and $R=14.16$ km, respectively, with a dipole magnetic field strength of $B\\equiv B_\\mu = 4\\times 10^{15}$ G, where $B_{\\mu}$ is a typical strength. As initial data, we use the numerical eigenfunction of the $\\ell=2,m=0$ fundamental mode, for the truncated system presented in Paper~II. The grid size is $100 \\times 80$ and $\\varepsilon_D=10^{-3}$, while the total simulation time is 2s. By computing the FFT of $\\partial_t{\\cal Y}$ (which is proportional to $\\delta u^\\phi$) at various points inside the star, we obtain the following results: Examining the FFT at three different radial locations ($r\\sim 0, R/2, R$) at $\\theta \\sim 0$ (magnetic axis), Fig. \\ref{Fig:FFT}, we observe a number of narrow frequency peaks. The strength of the peaks is several orders of magnitude larger than the FFT continuum. The overtones ($n>0$) are nearly {\\it integer multiples} of the fundamental frequency ($n=0)$. The corresponding FFTs at the same radial points, but at $\\theta \\sim \\pi/2$ (magnetic equator) are shown in Fig. \\ref{Fig:FFT2}, where, in addition to the family of frequency peaks observed in Fig. \\ref{Fig:FFT}, one can also observe a second family of frequency peaks, mainly at $r\\sim R$, for which the fundamental frequency is at a ratio of 0.6 with respect to the fundamental frequency in  Fig. \\ref{Fig:FFT}. The fundamental frequency of this family, as well as the first overtone (again an integer multiple) are clearly visible, while the other overtones seem to be buried inside a continuum of other frequencies.  Using Levin's toy model for the Alfv\\'{e}n continuum in magnetars, the above numerical results can be interpreted as follows: The fundamental frequency peaks of the two families are QPOs generated by the edges or turning points of an  Alfv\\'{e}n continuum (henceforth we call these two frequencies as the fundamental lower and upper QPOs and denote them as $L_0$ and $U_0$). We denote the overtones (which are nearly integer multiples) as $L_n$ and $U_n$, where $n>0$. For the chosen form of the magnetic field, the extent of the first continuum, between the fundamental $L_0$ and $U_0$ QPOs is at lower frequencies than the continuous frequency intervals corresponding to the overtones, which all overlap partially.  Because of the Alfv\\'{e}n continuum, the phase of the oscillations at the QPO frequencies is not constant, but varies throughout the star. Since the magnetic field is axisymmetric, the axis is both a turning point and edge of the continuum. Near the magnetic axis, the background magnetic field varies slowly, which allows for the upper QPOs to remain nearly coherent for a large number of oscillations. Since $\\delta u^\\phi$ is only a coordinate component, we also obtained FFTs of $r \\sin\\theta \\partial_t {\\cal Y}$ (which is proportional to the physical velocity in a unit basis) at all grid points. In two-dimensional simulations of fluid modes in non-magnetized stars, it has been shown that the amplitude of the FFT of physical variables at every point in the numerical grid, at a chosen normal mode frequency, is correlated to the shape of the eigenfunction of a given mode \\citep{Stergioulas2004}. Similarly, we use the magnitude of the FFT of $r \\sin\\theta \\partial_t {\\cal Y} $ at every point in the grid to define an {\\it effective amplitude} for each QPO (which, of course, will be time-varying, due to the absorption of individual modes in the continuum). In Fig. \\ref{Fig:Effective}, we show the effective amplitude of the fundamental upper QPO, $U_0$ and its first overtone $U_1$ as well as for the corresponding lower QPOs $L_0$ and $L_1$, obtained after a simulation with a $50\\times 40$ grid, $\\varepsilon_D=2\\times 10^{-5}$ and a duration of 2.1s.  The effective amplitude for $U_0$ has a maximum near the magnetic axis, while there are also {\\it nodal lines} along certain magnetic field lines. For the overtone, $U_1$, there exists an additional nodal line, starting perpendicular near the {\\it magnetic axis}, which divides the region of maximum amplitude near the magnetic axis roughly in half. Similarly (not shown here), each successive overtone corresponds to an additional ``horizontal'' nodal line, dividing the region of maximum amplitude near the magnetic axis into roughly equidistant parts. This agrees well with the fact that the frequencies of the overtones are nearly integer multiples of the fundamental frequency.  The effective amplitude of the fundamental lower QPO, $L_0$, is practically limited to a region that is only somewhat larger than the region of {\\it closed magnetic field lines} inside the star. For the first overtone $L_1$, a nodal line divides the region of maximum amplitude into two parts.   In Fig. \\ref{Fig:evol1} we show the evolution of $\\partial_t {\\cal Y}$ at the location inside the star where the effective amplitude of the fundamental upper QPO, $U_0$ attains its maximum value. It is evident that the QPO is long-lived, since the amplitude of the oscillations barely diminishes with time.   \\begin{figure} \\begin{center} \\includegraphics[width=70mm]{50-40-e000002-dy-m6-x34.eps} \\end{center} \\caption{Time evolution of $\\partial_t {\\cal Y}$ at the location inside the star where the effective amplitude of the fundamental upper QPO, $U_0$ attains its maximum value. The amplitude of the oscillations barely diminish with time} \\label{Fig:evol1} \\end{figure}   \\begin{table*} \\centering \\caption{Frequencies of lower and upper Alfv\\'{e}n QPOs and their ratios, for a representative sample of equilibrium models, constructed with various EOSs and masses and for a magnetic field strength of $B=B_\\mu$ (see text for details).} \\label{Tab:eos1} \\begin{tabular}{@{}lcccccccc@{}} \\hline Model  & M/R  & $f_{L_0}$ (Hz) & $f_{U_0}$ (Hz) & ratio & $f_{L_1}$ (Hz) & $f_{U_1}$ (Hz) & ratio &  $f_{U_2}$ (Hz) \\\\ \\hline A+DH$_{14}$    & 0.218\t&  15.4  &  25.0  &  0.616  &  30.7  &  49.4  &  0.621 & 74.4\\\\ A+DH$_{16}$    & 0.264 &  11.7  &  18.3  &  0.639  &  23.5  &  35.7  &  0.658 & 54.0\\\\ WFF3+DH$_{14}$ & 0.191\t&  17.9  &  29.8  &  0.601  &  36.2  &  59.2  &  0.611 & 89.8\\\\ WFF3+DH$_{18}$ & 0.265 &  11.7  &  18.0  &  0.650  &  23.5  &  35.5  &  0.662 & 53.3\\\\ APR+DH$_{14}$  & 0.171 &  20.4  &  34.1  &  0.598  &  41.3  &  68.6  &  0.602 & 104.6\\\\ APR+DH$_{20}$  & 0.248 &  12.8  &  20.6  &  0.621  &  26.0  &  40.3  &  0.645 & 61.0\\\\ L+DH$_{14}$    & 0.141 &  23.7  &  40.8  &  0.581  &  47.5  &  81.6  &  0.582 & 123.8\\\\ L+DH$_{20}$    & 0.199 &  16.4  &  27.8  &  0.590  &  33.1  &  54.7  &  0.605 & 82.6\\\\ \\hline \\end{tabular} \\end{table*}   \\begin{figure} \\begin{center} \\includegraphics[width=55mm]{f-MoR-fit.eps} \\end{center} \\caption{Quadratic fits in terms of the compactness of the star, $M/R$, of the lower and upper fundamental Alfv\\'{e}n QPO frequencies, obtained for a representative sample of equilibrium models with various EOSs and masses. The magnetic field was set to $B=B_\\mu$. } \\label{Fig:fit} \\end{figure} ",
        "conclusions": "\\label{sec:V}   We have already verified that our main QPO frequencies agree with frequencies obtained with an independent, fully nonlinear numerical code \\citep{CerdaDuran2007}, for the same initial model. We caution, however, that we have not yet considered the crust-core interaction, different magnetic field topologies or the coupling to the exterior magnetosphere. These effects have to be taken into account and already \\cite{Sotani2007c}, find that the observed QPOs could lead to constraints on the magnetic field topology. To complete the picture, a three-dimensional numerical simulation, that includes a proper coupling of the crust to the MHD interior and to the exterior magnetosphere will be required and our current results provide a good starting point. Extensive details of our computations will be presented in \\cite{Sotani2007d}."
    },
    "0710/0710.3428_arXiv.txt": {
        "abstract": "We explore a simple model for the high luminosity of SN~2006gy involving photon diffusion of shock-deposited thermal energy. The distinguishing property of the model is that the large ``stellar'' radius of $\\sim$160 AU required to prevent adiabatic losses is not the true stellar radius, but rather, the radius of an opaque, unbound circumstellar envelope, created when $\\sim$10 M$_{\\odot}$ was ejected in the decade before the supernova in an eruption analogous to that of $\\eta$~Carinae.  The supernova light is produced primarily by diffusion of thermal energy following the passage of the blast wave through this shell.  This model differs from traditional models of supernova debris interacting with external circumstellar matter (CSM) in that here the shell is optically thick and the escape of radiation is delayed.  We show that any model attempting to account for SN~2006gy's huge luminosity with radiation emitted by ongoing CSM interaction fails for the following basic reason: the CSM density required to achieve the observed luminosity makes the same circumstellar envelope opaque ($\\tau\\ga$300), forcing a thermal diffusion solution.  In our model, the weaker CSM interaction giving rise to SN~2006gy's characteristic Type IIn spectrum and soft X-rays is not linked to the power source of the visual continuum; instead, it arises {\\it after} the blast wave breaks free of the opaque shell into the surrounding wind.  While a simple diffusion model can explain the gross properties of the early light curve of SN~2006gy, it predicts that the light curve must plummet rapidly at late-times, unless an additional power source is present. ",
        "introduction": "The recent extremely luminous supernova (SN) 2006gy has challenged our understanding of the deaths of massive stars and may represent a new class of explosions. It was discovered by Quimby (2006) and was subsequently studied in detail by Ofek et al.\\ (2007) and Smith et al.\\ (2007). It is a Type IIn event, with strong narrow lines of H in its spectrum, although the details of its light curve and spectrum are unlike most other members of this class (e.g., Filippenko 1997). Its enormous total radiated energy of $>$10$^{51}$ ergs, the late ($\\sim$70 days) peak of the light curve, and the observed constant expansion speed of $v_{exp}\\simeq$4000 km s$^{-1}$ (Smith et al.\\ 2007) place important constraints on the model.  SN~2006gy exploded in the host galaxy NGC~1260, which is a peculiar S0/Sa galaxy with ongoing star formation (see Smith et al.\\ 2007 and Ofek et al.\\ 2007).  The optical displays of all Type Ia SNe and most Type II SNe are believed to be dominated by the radioactive decay sequence $^{56}$Ni$ \\rightarrow ^{56}$Co$ \\rightarrow ^{56}$Fe.  If that is also the case for SN~2006gy, a mass M($^{56}$Ni)$\\approx$10 M$_{\\odot}$ is required to produce the observed luminosity (Smith et al.\\ 2007).  Not surprisingly, SN~2006gy has prompted varying suggestions to account for its extreme energy budget. Ofek et al.\\ (2007) proposed a SN~Ia exploding in a dense H-rich circumstellar medium (CSM) where continual interaction with an external CSM was the power source for the luminosity.  They suggested that the dense CSM was probably the result of binary star common envelope ejection, and that the extreme energy budget may need to draw upon a super-Chandrasekhar Type~Ia explosion or perhaps a massive star explosion. Smith et al.\\ (2007) argued that many different spectral properties of SN~2006gy and its CSM were incompatible with a Type~Ia explosion, but were entirely consistent with the known properties of some very massive stars.  To account for the luminosity with CSM interaction, Smith et al.\\ (2007) proposed that the star suffered a tremendous explosive but non-terminal mass-loss event in the decade preceding the SN, analogous to the 19th century eruption of $\\eta$~Carinae. Woosley et al.\\ (2007) proposed a similar model, where the $\\eta$~Car-like explosion in the decade preceding the SN could have been triggered by the pulsational pair instability (different from a genuine pair instability SN) in a star with initial mass 110 M$_{\\odot}$.  Interaction with an external CSM can, in principle, generate high luminosity as seen in other Type IIn SNe (e.g., Chugai et al.\\ 2004). However, Smith et al.\\ (2007) argued that this mechanism is difficult to reconcile with the relatively low progenitor mass-loss rate inferred from the weak X-rays detected by {\\it Chandra} and the narrow H$\\alpha$ emission in SN~2006gy's spectrum. In other words, signs of CSM interaction are seen in SN~2006gy, but that interaction appears far too weak to power the visual continuum.  Even more troubling, the CSM interaction hypothesis cannot account for why the blast wave apparently did not decelerate while it was drained of more than 10$^{51}$ ergs (Smith et al.\\ 2007). These inconsistencies, combined with the slow rise, low expansion speed, and high luminosity of SN~2006gy, led Smith et al.\\ (2007) to appeal to the alternative of a pair-instability SN (Barkat et al.\\ 1967; Bond et al.\\ 1984; Heger \\& Woosley 2002) where the core of the star is obliterated and the light is powered by radioactive decay of $\\sim$10 M$_{\\odot}$ of $^{56}$Ni. As we discuss below, however, it may be possible to account for the luminosity of SN~2006gy without placing such extreme demands on SN nucleosynthesis.  The most pressing question about SN~2006gy boils down to this: Do we {\\it need} a pair instability SN to provide a consistent explanation for its extremely high luminosity and spectral properties?  The pair-instability hypothesis requires the introduction of an exotic phenomenon; while the idea has been around for decades, it has never been observed and is only expected to have occurred in the early Universe (e.g., Heger \\& Woosley 2002). Published models of the light curves from pair instability SNe do indeed predict high luminosities, slow rise times, long durations, and slow expansion speeds (e.g., Scannapieco et al.\\ 2005) qualitatively similar to SN~2006gy. The light curve shape depends on the assumed mass of the envelope, and most studies so far have been conducted for zero-metallicity stars with no mass loss. Further work is in progress, with various assumptions about the mass of the envelope (Nomoto et al.\\ 2007; Kasen 2007; Young 2007). The progenitor of SN~2006gy is likely to have been a very massive star (Smith et al.\\ 2007), but whether it was massive enough to suffer the pair instability remains an open question.  Here we describe how a simple photon diffusion model can account for the high luminosity and long duration of SN~2006gy.  Since the explosions of core collapse SNe typically deposit $E_0 \\simeq 10^{51}$ ergs of energy in the SN envelope, an exceptional mass of $^{56}$Ni is not necessarily required for SN~2006gy if an efficient mechanism could convert the energy of the initial SN blast into emergent light.  In SNe from compact progenitors, the light from the blast wave itself normally falls short by orders of magnitude.  The trapped internal radiation will cool by virtue of adiabatic expansion, and its energy will consequently be tranferred to the kinetic energy of the expanding debris.  The fraction of the blast energy emerging as radiation will be roughly $R_0/R_{max}$, where $R_0$ is the initial radius of the progenitor and $R_{max}$ is the radius of the SN photosphere at peak light. $R_0$ ranges from $\\sim$10$^{-2}$ AU for the Wolf-Rayet progenitors of SNe Ib/c to $\\sim$1 AU for the red supergiant progenitors of SNe~II, while typically, $R_{max} \\simeq 100$ AU.  Falk \\& Arnett (1973, 1977) showed that a SN from a progenitor having a very extended envelope could produce a light curve with a long timescale and high luminosity, with very efficient conversion of blast wave energy into radiation.  Below we consider this idea as a possible explanation for the light curve of SN2006gy. While we argue that this mechanism may provide a plausible explanation for the early light curve, we caution that it {\\it does not negate} the possibility that SN~2006gy is powered by some other mechanism at late times, such as $^{56}$Co decay or continued CSM interaction. ",
        "conclusions": "The diffusion model for SN~2006gy that we present here does not require any new physical processes beyond those that account for core-collapse SNe.  The only new ingredient is the hypothesis that the SN progenitor must have ejected an opaque shell having $\\sim$10~M$_{\\odot}$ extending to $\\sim$160 AU shortly before it exploded.  According to our model, the light curve of SN~2006gy observed up to about day 170 comes entirely from energy deposited by the initial blast and does not require any source of radioactive energy or ongoing CSM interaction.  However, this internal energy source cannot last.  As the shocked shell continues to expand, the radiative diffusion time becomes shorter than the time since explosion, after which the light curve will faithfully track the energy deposition rate.  Most of the kinetic energy of the SN must reside in debris having velocity $\\ga$4000 km s$^{-1}$, and thermal energy deposition from the impact of this debris with the expanding shell must terminate in $\\la$1 yr.  After that, there are only two other obvious sources of energy deposition to account for the light curve.  The first is the decay of $^{56}$Co, which would produce an exponential tail to the light curve having luminosity $L \\approx$1.4$\\times$10$^{43}$ M($^{56}$Ni) $\\exp(t/113.6 d)$ ergs s$^{-1}$, where M($^{56}$Ni) is the mass of newly synthesized $^{56}$Ni produced by the SN.  If this is the explanation for the excess luminosity above the diffusion model after day 170, then about 8~M$_{\\odot}$ of $^{56}$Ni would be required, indicated by the dashed line in Figure 1.  On the other hand, failure to detect a continued exponential tail in the visual light curve at late times does not necessarily give an upper limit to M($^{56}$Ni) produced in SN~2006gy. The bolometric flux may shift to near-infrared wavelengths as ejecta cool, and it is quite possible that dust will form in the shocked shell, reprocessing the luminosity into far infrared radiation.  This cooling may also cause changes in the bolometric correction as the SN evolves with time, and may play a role in the tail after day $\\sim$160 seen in Figure 1.  A second source could be continued shock interaction with the transparent CSM external to the opaque shell.  We see clear signs of that interaction in the emission lines seen in the optical spectrum and in the faint X-ray emission (Smith et al.\\ 2007).  CSM interaction could produce a lower-luminosity tail to the light curve of SN~2006gy as long as the CSM density remains high."
    },
    "0710/0710.1331_arXiv.txt": {
        "abstract": "We use high resolution (2048$^2$ zones) 2D hydrodynamic simulations to study the formation of spiral substructure in the gaseous disk  of a galaxy. The obtained gaseous response  is driven by a self-consistent non-axisymmetric potential  obtained from an imposed spiral mass distribution. We highlight the importance of ultraharmonic resonances in generating these features. The temporal evolution of the system is followed with the parallel ZEUS-MP code, and we follow the steepening of perturbations induced by the spiral potential until large-scale shocks emerge. These shocks exhibit bifurcations that protrude from the gaseous arms and continue to steepen until  new shocks are formed. When the contribution from the spiral potential relative to the axisymmetric background is increased from our default value, spurs protrude from the main arms after several revolutions of the gaseous disk. Such spurs overlap on top of the aforementioned  shocks. These results support the hypothesis that a complicated gaseous response can coexist with an orderly spiral potential term, in the sense that the underlying background potential can be smooth yet drive a gaseous response that is far more spatially complex. ",
        "introduction": "Young stars, H II regions and OB associations delineate the structure that gives its name to spiral galaxies (Elmegreen 1981; Vall\\'ee 2002, 2005). This implies a correlation between the spiral structure and the process of star formation. The density wave theory attempts to explain the large-scale structure of these galaxies in terms of a wave propagating in the disk of stars. Yet this stellar wave alone cannot directly explain the narrow spiral arms as delineated by the products of star formation.  Fujimoto (1968) first proposed that the dust lanes observed on the concave side (inside corotation) of spiral arms might be the result of a Galactic shock. Large scale shocks propagating in the gaseous layer of the Galaxy could induce a sudden compression of the interstellar medium (ISM) and trigger the star formation process. Since the temporal scales involved in this process are large (of the order of the Galactic rotation period), an isothermal approximation of the ISM is a first step in modeling the gaseous response, thus rendering a highly compressible-supersonic medium in which a small amplitude disturbance, such as that provided by an underlying spiral potential, may steepen into a shock wave.  In a semi-analytical study of the gas flow in spiral density waves, Roberts (1969) found nonlinear steady-state solutions containing shocks. On the basis that these shocks were coincident with the imposed perturbing potential minima, he argued that the density jump could trigger star formation, and in this way narrow bands of young stars could delineate the spiral arms.  Shu  et al. (1973) studied the gas flow in the context of the spiral density wave by adopting a two phase model for the ISM. In their study they carried out a slightly nonlinear analysis of such flow and found that, at certain points in their solution, the amplitude  became infinite, and they called these positions ultraharmonic resonances. They argued that  additional prominent features could appear as a consequence of the ultraharmonic resonances, for instance: they found a secondary compression associated with the first one.  However, these studies were carried out  assuming steady state flow and tightly wound spirals. By removing the first constraint, Woodward (1975) showed how the convective steepening of the initial perturbation leads to large-scale shock formation (as a response to the driving potential) in a gaseous layer initially in circular orbit. He also found the effect of the first ultraharmonic resonance as a secondary peak in the density field, but numerical viscosity in his calculations inhibited the formation of secondary shocks.  In two papers, Contopoulos \\& Grosb\\o l (1986, 1988) removed the second constraint and showed that nonlinear effects are not negligible in open spirals. Using numerical calculations, they demonstrated that for open spirals the effect of the first ultraharmonic resonance (which they call the 4:1 resonance) is such that stellar orbits do not support the spiral perturbation beyond the position of this resonance and hence the length  of the stellar spiral arms is limited by this position. In an analytical study of the effects of the ultraharmonic resonances, Artymowicz \\& Lubow (1992) explained this result as a cumulative effect of higher-order resonances between  4:1 and the corotation radius. They further argued that in the vicinity of the former  the gas response is such that  it looks like a 4-arm structure with a smaller pitch angle  than that of the original two-arm pattern.  It is clear that,  due to the nature of the problem, the highly nonlinear phenomena associated with the gaseous response to a spiral density wave are best studied via numerical simulations. Results so obtained add interarm features and spiral substructure on top of  large scale shocks. Patsis et al. (1997) found that  in open spirals  a bifurcation of the main spiral arms  takes place at the 4:1 resonance position. By analyzing numerical experiments that include self-gravity in the study of the effects of the ultraharmonic resonances on gaseous disks, Chakrabarti et al. (2003, from here on CLS03) found secondary compressions, associated with the first ultraharmonic resonance, in models with a low Toomre parameter, Q, which were evolved only for a few revolutions of the disk. These compressions eventually became branches. In high-Q models, evolved for several revolutions they found the appearance of leading structures which they identified with spurs. However, spurs and branches not related to resonances have also been obtained in numerical simulations (Dobbs \\& Bonnell 2006; Wada \\& Koda 2004). By taking into account frozen magnetic fields and self-gravity in their local arm simulations, Kim \\& Ostriker (2002, 2006) showed that gaseous spurs form as a consequence of gravitational instability inside the spiral arms, a result confirmed with global 2D simulations by Shetty \\& Ostriker (2006). In those simulations spurs jut out of the spiral arms at regular intervals.  In this paper we carry out two dimensional, global, high resolution hydrodynamic simulations to further study the formation of spiral substructure in galactic gaseous disks. Our approach differs from previous work in that we employ a self-consistent model for the spiral stellar potential (in the orbital sense). This potential accounts for its own self-gravity and thus is no longer a local arm approximation to the driving term, making it  more appropriate for simulations of open galaxies.  This paper is organized as follows: in \\S 2 we describe the potential we employ, in \\S 3 we present the obtained gaseous response and the type of substructure related to it, in \\S 4 we discuss the results and compare with previous work, and in \\S 5 we present our conclusions. ",
        "conclusions": "A fully nonlinear treatment of the propagation of spiral density waves in a gaseous disk is tractable using  numerical simulations. By considering an open spiral, in the absence of self-gravity in the gaseous layer, we are able to obtain rich substructure associated with an external spiral potential. Our work differs from other published results in that we employ a self-consistent driving term that considers its own gravity and removes the local arm approximation. We have presented experiments showing the formation of a four-arm structure in response to a two-arm driving pattern. Initially the gas accumulates in the potential minima. After some time the gas responds to resonances and the main arms bifurcate. These features have a considerable azimuthal extension and continue to steepen with  time. Eventually a four-arm shock structure  emerges, in the gaseous layer that coexist with the underlying two-arm potential, made up by the stars. This result appears to agree with observational data published by Drimmel (2000) where he concludes that, using optical tracers, the spiral structure in our Galaxy is best fitted by a four-arm structure, while in the infrared a two-arm structure dominates.  If we place the first ultraharmonic resonance outside the region of influence of the spiral (in our case by changing the value of $\\Omega_p$) the main arms do not bifurcate and large-scale shocks are the only induced response in the gaseous layer. In this way we show that if the first ultraharmonic resonance is placed outside the region of influence of the perturbing term, no spiral substructure appears, thus emphasizing the connection between this resonance and the bifurcations of the gaseous arms.  If we combine the effects of the resonances with a large forcing amplitude we obtain, on top of this four-arm structure, spurs protruding from the main arms in a region between the inner Lindblad resonance and the first ultraharmonic resonance. Overlapping of nonlinear effects take place in our long-timescale global simulations, showing that an orderly spiral density wave potential can produce a gaseous response that is strongly disordered.  The inclusion of additional processes such as magnetic fields, self-gravity for the gaseous layer, and thermal processes in the ISM may lead to the appearance of additional substructure and their consequent fragmentation into bound condensations. Numerical experiments addressing these topics are needed in order to improve our understanding of the global ISM in galaxies and its relation to large-scale processes."
    },
    "0710/0710.1277_arXiv.txt": {
        "abstract": "We present new deep {\\it Chandra} observations of the Centaurus A jet, with a combined on-source exposure time of 719 ks. These data allow detailed X-ray spectral measurements to be made along the jet out to its disappearance at 4.5 kpc from the nucleus. We distinguish several regimes of high-energy particle acceleration: while the inner part of the jet is dominated by knots and has properties consistent with local particle acceleration at shocks, the particle acceleration in the outer 3.4 kpc of the jet is likely to be dominated by an unknown distributed acceleration mechanism. In addition to several compact counterjet features we detect probable extended emission from a counterjet out to 2.0 kpc from the nucleus, and argue that this implies that the diffuse acceleration process operates in the counterjet as well. A preliminary search for X-ray variability finds no jet knots with dramatic flux density variations, unlike the situation seen in M87. ",
        "introduction": "\\label{intro} Low-power, Fanaroff-Riley class I (FRI) \\citep{fr74} radio galaxies are numerically the dominant population of radio-loud active galaxies in the universe. Their dynamics and energy content are thus essential to models of AGN `feedback' in which the energy released in the process of accretion onto the central supermassive black hole is transported to large spatial scales via the interaction between the jets and the external medium. To understand the dynamics and energetics of these sources it is crucial that we should understand the particle acceleration processes by which the bulk kinetic energy of the jets is translated into the internal energy of relativistic plasma.  The discovery with {\\it Chandra} that X-ray emission is common in the inner few kpc of FRI jets \\citep{wbh01} provides a very strong argument that {\\it in situ} high-energy particle acceleration is taking place in those regions. The broad-band spectral energy distribution and X-ray spectrum of the X-ray emission imply a synchrotron origin for the X-rays \\citep{hbw01}. For magnetic fields close to the equipartition value in a typical powerful jet, the loss timescale ($\\tau = -E/\\dot E$) for the $\\gamma \\sim 10^7$ -- $10^8$ electrons emitting X-ray synchrotron emission is tens of years. Thus observations of X-ray synchrotron emission essentially tell us where particle acceleration is happening {\\it now}: in particular, resolved, diffuse X-ray emission implies a particle acceleration mechanism that must be distributed throughout the jet.  To probe the nature of the acceleration mechanisms in FRI jets we require observations that can reliably distinguish between compact and diffuse X-ray emission: this implies (at the most optimistic) a spatial resolution that is comparable to the loss spatial scale, the distance travelled by an electron before synchrotron losses remove it from the X-ray band, or roughly $c\\tau$. Even with the comparatively high angular resolution of {\\it Chandra}, this can only be achieved in the nearest FRI radio galaxy, Centaurus A. Given Cen A's distance of $\\sim 3.7$ Mpc (the average of 5 distance estimates to Cen A: \\citealt{fmst07}), {\\it Chandra}'s resolution corresponds to $\\sim 10$ pc.  {\\it Chandra} observations of Cen A show a complex, knotty and in places edge-brightened jet structure (\\citealt{kfjk00, kfjm02, hwkf03} [hereafter H03]; \\citealt{ksat06} [K06]; \\citealt{hkw06} [H06]). It has been argued (H03,K06) that at least two acceleration mechanisms are required: the dynamics and spectra of the compact knots in the inner part of the jet are consistent with a shock origin, while the diffuse emission is inconsistent both spectrally (H03) and in terms of the number distribution of knots (K06) with being the sum of many unresolved knots with the same properties as those observed, and is instead probably truly diffuse, implying a distributed acceleration mechanism such as second-order Fermi acceleration \\citep[e.g.,][]{so02} or magnetic field line reconnection \\citep[e.g.,][]{bl00}. Here we present the results of new, much deeper {\\it Chandra} observations and their consequences for the nature of particle acceleration in the Cen A jet and counterjet.  \\begin{figure*} \\epsfxsize 16.6cm \\epsfbox{f1n.eps} \\caption{False-color image of the jet and counterjet regions of Cen A using all ACIS observations. The {\\it Chandra} data from each observation have been exposure corrected at an appropriate energy and then combined, weighting according to the value of the exposure map. The images are binned in standard {\\it Chandra} pixels ($0\\farcs492$ on a side) and smoothed with a FWHM = $1\\farcs0$ Gaussian. Red shows exposure-corrected counts in the energy range 0.4--0.85 keV, green shows 0.85--1.3 keV and blue 1.3--2.5 keV. Contours [$7 \\times (1,4,16\\dots)$ mJy beam$^{-1}$] are from the 5-GHz VLA map of H06 with $6\\farcs0$ resolution.} \\label{fc} \\end{figure*}  In this analysis {\\it Chandra} data processing was done using {\\sc ciao} 3.4 and {\\sc caldb} 3.3.0.1. Spectral fitting was carried out in {\\sc xspec}. We define the spectral index $\\alpha$ such that flux $\\propto \\nu^{-\\alpha}$: the photon index $\\Gamma = 1+\\alpha$. Errors quoted and plotted throughout are $1\\sigma$ for one interesting parameter. ",
        "conclusions": "The inner part of the Cen A jet is dominated by the shock-related knots discussed by H03, although we have so far failed to see the dramatic variability seen in M87, interpreted as large changes in particle acceleration processes or in jet flow. The radio-associated compact features in the counterjet are likely also the result of small-scale shocks, presumably caused as the counterjet flow encounters compact obstacles. However, the outer 3.4 kpc of the jet -- and possibly also the newly detected extended X-ray counterjet components -- are different from the shock-dominated regions in their X-ray spectra and structure and in their radio/X-ray ratio. Our results add to the evidence that an unknown, distributed particle acceleration process operates in the jet of Cen A. Diffuse X-ray emission with a similar broad-band spectrum is seen in the jets of many other FRI sources (e.g., 3C\\,66B: \\citealt{hbw01}) and it also resembles the extended synchrotron X-ray emission seen in the jets (e.g., K06) and hotspots \\citep[e.g.,][]{hck07} of the more powerful FRIIs. The Cen A data allow us to study the spatial and spectral properties of this emission and the physical processes responsible for it in more detail than is possible in any other object. \\vspace{-10pt}"
    },
    "0710/0710.2332_arXiv.txt": {
        "abstract": "\\baselineskip=13pt A new mass table calculated by the relativistic mean field approach with the state-dependent BCS method for the pairing correlation is applied for the first time to study $r$-process nucleosynthesis. The solar $r$-process abundance is well reproduced within a waiting-point approximation approach. Using an exponential fitting procedure to find the required astrophysical conditions, the influence of mass uncertainty is investigated. $R$-process calculations using the FRDM, ETFSI-Q and HFB-13 mass tables have been used for that purpose. It is found that the nuclear physical uncertainty can significantly influence the deduced astrophysical conditions for the $r$-process site. In addition, the influence of the shell closure and shape transition have been examined in detail in the $r$-process simulations. ",
        "introduction": "It is of the utmost interest to explore the ``terra incognita'' of exotic nuclei, as evidenced by the fact that several Radioactive Ion Beam (RIB) facilities are being upgraded, under construction or planned to be constructed worldwide. Such investigations of the properties of these exotic nuclei, which may behave very differently from the nuclei around the $\\beta$-stability line, result in new discoveries such as the halo phenomenon~\\cite{THH-PRL85,Schwab94} - nucleons spread like a thin mist around the nucleus, which can significantly increase the nuclear reaction ratio. Stellar nucleosynthesis processes such as the $r$-process~\\cite{Cowan-PR91,Qian-PPNP03}, which is responsible for roughly half of the enrichment of elements heavier than iron in the universe, also require a thorough understanding of the properties of exotic nuclei. Key properties like masses for example, determine the path that the  nucleosynthesis process follow in the nuclei chart. Nevertheless, despite many experimental efforts, present knowledge of exotic nuclei still does not include much of what is required for a complete understanding of $r$-process nucleosynthesis. After the first systematic introduction to the $r$-process~\\cite{BBFH} half a century ago, $r$-process calculations for a long time could only rely on the phenomenological nuclear droplet mass formula~\\cite{Hilf} because of the lag of both experimental and theoretical development. Fortunately, in the last 15 year the theoretical study of nuclear properties has made tremendous progress and $r$-process calculations~\\cite{Kartz-APJ93,pfeiffer-ZPA97,Wanajo-APJ04} have been carried out based on the refined droplet model FRDM~\\cite{FRDM}, Hartree-Fock approach like ETFSI-Q~\\cite{ETFSIQ}, and the very recent microscopic rooted Hartree-Fock Bogliubov (HFB)~\\cite{HFB,HFB2,Samyn-PRC03}.   Despite the progress in the theoretical nuclear structure physics, mass models predictions (which by design concentrate on different nuclear structure aspects) still show a large deviation when going to very neutron-rich nuclides, even though they have achieved similar quality to describe known nuclides. This is specially troublesome since the astrophysical scenario in which an $r$-process may occur is a matter of debate and all astrophysical simulations dedicated to the nature of the stellar environment depend on the input from nuclear physics. Mass model predictions, even in models that give similar global $rms$ error still show local deviations differently.  In principle, microscopic-rooted mass models should have a more reliable extrapolation to the unknown regions, therefore these studies have received more and more interest as evidenced by the increasing number of non-relativistic HFB investigations~\\cite{HFB,HFB2,Samyn-PRC03,HFB-13}. Based on a mass-driven fitted method, the latest HFB models have achieved a similar quality ($rms\\sim$ 0.7 MeV) as the phenomenological FRDM mass model for known masses. More recently, another microscopic-rooted approach, the relativistic mean-field (RMF) theory~\\cite{Wale-AP74} has received broad attention due to its successful description of several nuclear phenomena during the past years (for recent reviews, refer to Refs.~\\cite{Bender-RMP03,Meng06}). In the framework of the RMF theory, the nucleons interact via the exchanges of mesons and photons. The corresponding large scalar and vector fields, of the order of a few hundred MeV, provide simple and efficient descriptions of several important phenomena such as the large spin-orbit splitting, the density dependence of optical potential, the observation of approximate pseudo-spin symmetry, etc.. Moreover, the RMF theory can reproduce well the isotopic shifts in the Pb region~\\cite{SLR-PLB93}, explain naturally the origin of the pseudo-spin symmetry~\\cite{AHS-PLB69,HA-69} as a relativistic symmetry~\\cite{Ginocchio-97,MSYP-PRC98,MSYA-PRC99,Chen-CPL03} and spin symmetry in the anti-nucleon spectrum~\\cite{ZMP-PRL03}.  The first RMF mass table was reported in Ref.~\\cite{Hirata-NPA97} for 2174 even-even nuclei with $8\\leq Z \\leq 120$ but without including pairing correlations. Later on, the calculation was improved by adopting a constant-gap BCS method and calculated 1200 even-even nuclei with $10 \\leq Z \\leq 98$~\\cite{Lalazissis-ADNDT99}, most of which are close to the $\\beta$-stability line. More recently, using the state-dependent BCS method with a $\\delta$-force~\\cite{Sanulescu-PRC00,Geng-PTP03}, the first systematic study of the ground state properties of over 7000 nuclei ranging from the proton drip line to the neutron drip line was performed \\cite{RMFBCS}. Comparison of this calculation with experimental data and to the predictions of other mass models will be presented in more detail in Sec II.  Considering the recent development of the microscopic mass models in both the HFB and RMF approach, it is very interesting to examine their applicability to an $r$-process calculation. The main goals of this paper is to explore to what extent the solar $r$-process abundance can be reproduced using the new RMF mass table~\\cite{RMFBCS} and by comparing with other theoretical mass models, to determine the influence of nuclear mass uncertainty in $r$-process calculations. The paper is organized as follows. In Sec.~\\ref{sec:RMF} the global agreement of the new RMF mass table with the experimental data is discussed and the RMF prediction in the very neutron-rich range is compared with the FRDM~\\cite{FRDM}, the ETFSI-Q~\\cite{ETFSIQ} and the latest HFB-13~\\cite{HFB-13} mass tables. In Sec.~\\ref{sec:$r$-process}, a short introduction to a site-independent $r$-process approach is given. In Sec.~\\ref{sec:calculation}, the new mass table is applied to reproduce the solar $r$-process abundances. In addition, the result is compared to the $r$-process abundances obtained with the predictions of the FRDM, ETFSI-Q and HFB-13 mass models. Finally the summary and conclusions are given in Sec.~\\ref{sec:conclusion}. ",
        "conclusions": "Summary}   We have applied the most recent comprehensive mass models, the non-relativistic microscopic-rooted HFB-13 and the relativistic RMF in $r$-process calculations. For the sake of comparison, we also included the widely used macro-microscopic models FRDM and ETFSI-Q. Of these models, the HFB-13 and RMF models are used for the first time in such calculations. Based on a simple $r$-process model, it is found that all mass models reproduce the main features of the solar $r$-process pattern and the position of the abundance peaks. Since $r$-process simulations have to rely on predicted nuclear physics properties of unknown regions in the nuclear chart, we have compared the predictions of different mass models. We have also made a systematic study of the influence due to the mass model uncertainty in the application of the $r$-process and thus in the required astrophysical conditions. This nuclear physical uncertainty is very important for the complete understanding of the $r$-process since the results of more modern full dynamic $r$-process calculations depend on the nuclear mass input used. It is found that the deduced astrophysical conditions like the neutron irradiation time of the $r$-process can be significantly different depending on the mass model used. Among the different models, the simulation using the RMF masses requires a longer time scale (up to a factor of 4) than those using FRDM and HFB-13 models. Furthermore, it is found that the optimal astrophysical conditions obtained using the ETFSI-Q and RMF mass models require a relatively constant weighting factor for neutron densities in the range $10^{22}$ to $10^{28}$ cm$^{-3}$, while the FRDM and HFB-13 simulations favor a large weighting factor at low densities. In addition, we have explored the possible deficiencies in different mass models, and found that the observed abundance underproduction before the abundance peaks in all the models can be a combined and complex effect of both shell structure and shape transition. An exception is the underproduction at $A\\sim 115$ in the HFB-13 model which can be attributed to incorrect $\\beta$-decay rates. Future experiments are needed to determine the strength of the shell closure towards the neutron drip line as well as the precise locations of the shape transition toward the shell-closures."
    },
    "0710/0710.2104_arXiv.txt": {
        "abstract": "In this work we show that 3-3-1 model with right-handed neutrinos has a natural weakly interacting massive particle (WIMP) dark mater candidate. It is a complex scalar with mass of order of some hundreds of GeV which carries two units of lepton number, a scalar bilepton. This makes it a very peculiar WIMP, very distinct from Supersymmetric or Extra-dimension candidates. Besides, although we have to make some reasonable assumptions concerning the several parameters in the model, no fine tunning is required in order to get the correct dark matter abundance. We also analyze the prospects for WIMP direct detection by considering recent and projected sensitivities for WIMP-nucleon elastic cross section from CDMS and XENON Collaborations. ",
        "introduction": "\\setcounter{equation}{0} \\label{sec1} The problem of matter density in the Universe seems to be one of the most intriguing and exciting subjects in modern Physics. The growing refinement achieved in cosmological data leaves no doubt about a dark component in the observed mass density, constituting roughly $22\\%$ of all energy density acording to the three year run of WMAP~\\cite{WMAP3}. This yet unknown component has to be non-baryonic, its interaction with the electroweak Standard Model (SM) particles should be negligible  and it has to be cold, i.e., non-relativistic at the time it decouples from the radiation bath, the so called  Cold Dark Matter (CDM). From the theoretical side, there are some proposals to explain the CDM in the context of Particle Physics models (see Ref.~\\cite{DMmodels,Murayama,dmoutros,Dobrescu,little} and references therein for a review of the subject). Among them there are models which present natural candidates to play this role, the weakly interacting massive particles (WIMP)'s, with mass ranging from approximately $1$~GeV to $1$~TeV. These WIMP's are nice candidates because their masses are in the GeV realm, turning them cold at decoupling, and mainly because their weakly interacting aspect not only yields a thermally averaged annihilation cross section of order of weak interactions, leading to the expected order of magnitude to CDM abundance, but also coincides with the scale of Particle Physics models to be probed at the Large Hadron Collider (LHC) that finds itself at the final stage to start its running phase~\\cite{LHC}. It also presents the possibility of being seen in direct detection experiments since its massiveness would imply an observable recoil of nuclei in elastic collisions~\\cite{DMmodels,wimpexperiments}.   The most promising scenarios where such WIMP's can be present in the particle spectrum are Supersymmetry (SUSY) and Extra Dimensions models~\\cite{DMmodels,Murayama,Dobrescu}. All these models dispose of some kind of discrete symmetry in order to stabilize their CDM candidates. Also, they have to be realized at the electroweak scale so that their new particles are potentially good candidates for CDM. Although such models may represent the greatest expectations in Particle Physics for the Physics at TeV scale to be probed by LHC, the absence of any experimental evidence to support these models allows us to work with alternatives. One of such alternative routes concerns the enlargement of the gauge symmetry group from $SU_C(3)\\otimes SU_L(2)\\otimes U_Y(1)$ to larger groups. In particular, there exists a simple extension of the SM gauge group to $SU_C(3)\\otimes SU_L(3)\\otimes U_X(1)$, the so called 3-3-1 model~\\cite{PleitezPisano}.  This class of models is interesting for several reasons, not only because they mean a different scenario, but because they possess several nice features. For example, (i) the family problem is absent in this model since it demands that there be only three families of fermions when anomalies are canceled and asymptotic freedom is considered~\\cite{PleitezPisano}; (ii) electric charge quantization is automatic~\\cite{qcharge}; (iii) right-handed neutrinos can be part of the spectrum in some versions of the model~\\cite{331valle,331RH,singer} and their tiny observed mass difference can be easily accommodated~\\cite{lightnu}; (iv) axions and majorons are a natural outcome in some versions~\\cite{axionmajoron}, providing light particles which could also contribute to the problem of dark matter origin. Besides, it is possible that a custodial symmetry exists in these models which would make them indistinguishable from SM at low energy scales~\\cite{331np}, and it would then be a strong rival to SM itself.  In this work we will concentrate on the 3-3-1 version of the model with right handed neutrinos in the spectrum (3-3-1$RH\\nu$) model~\\cite{331RH}. The reason behind this choice relies on the fact that neutrinos mass is already a mandatory property that needs to be included in all reasonable extensions of SM. Besides, the model can be implemented with just three scalar triplets instead of including a sextet as in other versions, being economical in its content. However, the most appealing motivation to deal with 3-3-1$RH\\nu$ to explain the origin of CDM is due to the possibility of having a candidate with a very distinct signature. Among its properties the model can be made lepton number conserving if some of its fields carry two units of lepton number which will be called bileptons. This peculiar property has many phenomenological implications. Namely, rare lepton decays can emerge, neutrinoless double beta decay is allowed, right handed neutrinos are going to appear as byproducts of heavy vector bileptons decays and so on. It is then automatic to ask if some of these additional bilepton fields can be a CDM candidate, once it is provided with a very specific quantum number appropriate to forbid its interaction with many of the electroweak fields. This would play a similar role as that played by the discrete symmetries in the competing models cited above. Moreover, as the sought candidate is merged in the exotic new effects just mentioned, their appearance in the coming collider experiments would represent an unquestionable evidence of our CDM candidate.  What we are going to investigate in this work is the possible realization of this scenario in the 3-3-1$RH\\nu$ model for one of the bilepton scalars. We are interested in identifying this field and characterize it as a WIMP. This can be realized if it can be shown that it is stable in the range of parameters for which the abundance is in accordance with the recent data of WMAP~\\cite{WMAP3}. It is important to stress that the natural perturbative scale for the 3-3-1 models is on the TeV scale~\\cite{alexlimit}, which is suitable for obtaining a WIMP. In view of all this, it seems that such a possibility is as welcome as any previous attempt to explain the CDM content through SUSY or  Extra dimensions, offering a completely new and distinct particle to do this job. It should be mentioned that other works exist in the literature trying to explain CDM in the context of 3-3-1 models~\\cite{DM331}, but their aim was to obtain a self-interacting dark matter to avoid excessively dense cores in the center of galaxies and clusters as well as excessive large number of halos within the local group when contrasted to observations~\\cite{spergel}, which demands a light dark matter candidate. We do not pursue this approach here. On the contrary, we want to show that the 3-3-1$RH\\nu$ model possesses a bilepton scalar that can play the role of a WIMP, which is the preferred candidate for CDM. Besides, we get this without the need to fine tune its couplings to small and unnatural values.  This work is organized as follows. In Sec.~\\ref{sec2} we present the model with its content and interactions. In Sec.~\\ref{sec3} we diagonalize the mass matrices for the particle spectrum, allowing us to characterize the new extra particles. In Sec.~\\ref{sec4} we identify our WIMP with one of the neutral scalar bileptons of the model, checking its viability as CDM candidate by computing its abundance for a range of values of the parameters, which turns out to be natural in the sense we do not need to make any fine adjustment on the parameters. Then we study the prospects of direct detection for this WIMP. We conclude with Sec.~\\ref{sec5}. ",
        "conclusions": "\\label{sec5}  The 3-3-1RH$\\nu$ model admits a couple of bilepton particles in its spectrum, raising the possibility of having a  CDM candidate, since bileptons carry two units of lepton number. This is so because such a very specific quantum number is appropriate to forbid its interaction with many of the electroweak fields, since they are allowed to decay only on other bileptons, which have to be heavier than SM particles. Considering this scenario we obtained the particle mass spectrum of scalars in 3-3-1RH$\\nu$ model and, by assuming some conditions over the parameter space, we have shown that the lightest bilepton in the model turns out to be a scalar, a combination of two scalar interaction eigenstates that we called $\\phi$. For the region where the values of the parameters guarantee the stability of this scalar, we computed the $\\phi$ abundance and obtained stringent constraints for the parameters in order to have agreement with WMAP results for CDM abundance. We have found that $\\phi$ can have mass ranging from about $600$~GeV to some Few TeV, characterizing it as a heavy WIMP. It is opportune to say that, although we have restricted our parameter space due  to lack of knowledge on several couplings in the model, we had no need to unnaturally adjust them to very small values as generally happens in several models, including Supersymmetry. In fact we assumed that these couplings are close to one, and checked that $\\phi$ is an excellent candidate to represent a WIMP and explain the presently observed CDM abundance.  We also studied the possibility of observing this WIMP in direct detection experiments. For this we  have computed the elastic scattering $\\phi$-nucleon cross section and contrasted our results with present and future experiments. We have seen that $\\phi$ is still far from the range of detection for current and near future CDMS and XENON sensitivities, at least for the short parameter space considered in this work. However, even this limited scenario can be at reach for projected  Phases B and C of Super-CDMS~\\cite{SuperCDMS}. Besides this, it would be interesting to pursue the production of $\\phi$ at collider experiments, mainly at LHC, and also extend our search including a larger region of the parameter space considering additional phenomenological constraints on 3-3-1RH$\\nu$ model from Collider physics and Cosmology, as well as include prospects for $\\phi$ indirect detection too, a gap we wish to fill soon.  Finally, we would like to stress that our proposed WIMP is not only feasible but a reasonable alternative in the sense that the Particle Physics model we are dealing with is only a small extension of the SM gauge group, whose scale is about to be assessed at LHC. There are several features that distinguishes the 3-3-1RH$\\nu$ model from other extensions, like SUSY and Extra dimensions models. Namely, we have not only bilepton scalars in the spectrum but vector bosons and quark bileptons, all of them acquiring mass at hundreds of GeV. Certainly their signal at detectors are worth to be studied. Besides, new phenomena are predicted in this model~\\cite{331valle,lightnu,axionmajoron,SCPV}, including neutrinoless double beta decay, rare decays, new sources of CP violation and so on. Presence of such signals  would reinforce our expectation concerning a bilepton WIMP to explain CDM in the Universe. \\\\ \\\\           \\noindent {\\bf Acknowledgments:}\\\\ The authors acknowledge the support of the Conselho Nacional de Desenvolvimento Cient\\'{\\i}fico e Tecnol\\'ogico (CNPq)."
    },
    "0710/0710.1727_arXiv.txt": {
        "abstract": "Stellar kinematics show no evidence of hidden mass concentrations at the centre of M83. We show the clearest evidence yet of an age gradient along the starburst arc and interpret the arc to have formed from orbital motion away from  a starforming region in the dust lane. ",
        "introduction": "\\label{sec:intro} The nucleus of M83 is offset from the bulge centre and surrounded by a semicircular starburst arc (Fig. 1). \\citet[][hereafter H01]{Harris01} dated the star clusters in the arc with WFPC2 photometry and found evidence of an age gradients along it. However, the reddening vector parallelled the tracks in the two-colour diagram and clusters may be overlooked or poorly sampled in the visible because of extinction. \\citet[hereafter T00]{Thatte00} proposed the existence of a second mass concentration after NIR long-slit kinematics revealed an additional peak in the velocity dispersion, 2.7\\hbox{$^{\\prime\\prime}$}\\ SW of the nucleus. Further studies with IFUs linked velocity gradients in gas kinematics to hidden mass concentrations at different locations: \\citet{Mast06} report a gradient in \\textrm{H}{\\large $\\alpha$}\\ at 3\\hbox{$.\\!\\!^{\\prime\\prime}$}9$\\pm$0\\hbox{$.\\!\\!^{\\prime\\prime}$}5 W of the nucleus while Diaz et al. (\\citeyear[hereafter D06a and D06b]{D06a,D06b}) report a gradient in \\textrm{Pa}{$\\beta$}\\ 7\\hbox{$^{\\prime\\prime}$}\\ WNW of the nucleus. However, gas kinematics are known to suffer non-gravitational effects \\citep{KR95}.  We have analysed new VLT ISAAC K band long-slit spectroscopy together with archival HST data (\\textrm{Pa}{\\large $\\alpha$}\\ from NICMOS and \\textrm{H}{\\large $\\alpha$}\\ from WFPC2). Fig. 1 illustrates slit positions on the HST data. The combined data gives equivalent width (EW) measurements of \\textrm{H}{\\large $\\alpha$}, \\textrm{Pa}{\\large $\\alpha$}\\ and CO (2.3\\hbox{$\\umu$m}), as well as stellar and gas kinematics.  \\begin{figure} \\centering \\begin{minipage}[t]{0.7\\textwidth} \\centering \\resizebox{0.9\\textwidth}{!}{\\includegraphics{HOUGHTON_fig1.ps} } \\end{minipage}% \\begin{minipage}[c]{0.3\\textwidth} \\vspace{-1.5cm} \\caption[]{\\textbf{(a)} A BVR image of M83 (logarithmic intensity) using data from \\citet{Larsen99}, with the footprint of the homogenised HST data overlaid. Axes are scaled in degrees. \\textbf{(b)} F222N NICMOS image of the central 20\\hbox{$^{\\prime\\prime}$}$\\times$20\\hbox{$^{\\prime\\prime}$} with ISAAC slit positions overlaid. Positions of the putative hidden mass concentrations are also shown as purple (T00) and red (D06a,D06b) triangles. Flux is given in ergs s$^{-1}$ cm$^{-2}$ and axes are scaled in arcseconds. } \\end{minipage}  \\label{fig:m83} \\end{figure}  \\begin{figure} \\centering \\begin{minipage}[r]{0.6\\textwidth} \\centering \\vspace{-2.5cm} \\includegraphics[width=0.7\\textwidth]{HOUGHTON_fig2.ps} \\end{minipage}% \\begin{minipage}[c]{0.4\\textwidth} \\caption[]{An image of the centre of M83 with ages overplotted; note that the circumnuclear arc extends into and along the dust lane; the position, scaling and orientation is the same as Fig. 1b; NICMOS F222M is red, WFPC2 F814W is green and WFPC2 F300W is blue.} \\end{minipage}  \\label{fig:sb99} \\end{figure} ",
        "conclusions": ""
    },
    "0710/0710.1382_arXiv.txt": {
        "abstract": "Planned space-based ultra-high-energy cosmic-ray detectors (TUS, JEM-EUSO and S-EUSO) are best suited for searches of global anisotropies in the distribution of arrival directions of cosmic-ray particles because they will be able to observe the full sky with a single instrument. We calculate quantitatively the strength of anisotropies associated with two models of the origin of the highest-energy particles: the extragalactic model (sources follow the distribution of galaxies in the Universe) and the superheavy dark-matter model (sources follow the distribution of dark matter in the Galactic halo). Based on the expected exposure of the experiments, we estimate the optimal strategy for efficient search of these effects. ",
        "introduction": "\\label{sec:intro} The next step in the studies ultra-high-energy (UHE) cosmic rays (CRs) is related to the use of space-based detectors observing the fluorescent light induced by air showers in the terrestrial atmosphere with large exposures. It is expected that these instruments will help us to shed light on the origin of the highest-energy cosmic particles which remains unknown up to now.  One of the important signatures of particular UHECR models is the global anisotropy of arrival directions of the highest-energy events. For instance, models where the origin of UHECRs is attributed to the acceleration in astrophysical objects (so-called ``bottom-up'' scenarios) naturally predict that the distribution of arrival directions follows the distribution of these cosmic accelerators. In the most common scenario with extragalactic protons and/or nuclei, the patterns of the distribution of galaxies in the nearby Universe should be seen~\\cite{EGmogila} on the UHECR skymap because of the limited propagation length of these particles due to the GZK effect~\\cite{GZK} or nuclear photodisintegration. Strong suppression of the cosmic-ray flux is predicted at highest energies in these models. On the other hand, the ``top-down'' models~\\cite{TD} (see Refs.~\\cite{TDrev} for reviews and references) with the distribution of sources in the Galactic halo following that of the dark matter (which is the case for the superheavy dark-matter (SHDM) particles and some of topological defects) predict~\\cite{DubTin} the Galactic center-anticenter asymmetry due to the non-central position of the Sun in the Galaxy (see Refs.~\\cite{DMmogila,ABK1,ABK2} for extensive discussions). Models of this kind predict continuation of the cosmic-ray spectrum and gamma-ray dominance beyond $10^{20}$~eV.  Currently, neither the spectrum nor anisotropy observations can definitely favour one of the two scenarios at the highest energies, $E \\sim 10^{20}$~eV. Indeed, the AGASA experiment claims~\\cite{AGASAspectrum} the super-GZK continuation of the spectrum while the HiRes collaboration reports~\\cite{HiRes-cutoff} the observation of the Greisen--Zatsepin--Kuzmin (GZK) cutoff (data of other experiments, including recent results of the Pierre Auger Observatory (PAO)~\\cite{PAO-SDspectrum,Engel-CIC}, are not yet conclusive: though the unsuppressed continuation of the spectrum  is excluded by the Auger data, the cutoff in the spectrum is not clearly seen). The limits on the gamma-ray fraction in the primary cosmic-ray flux (currently the most restrictive ones arise from the AGASA and Yakutsk muon data at $E>10^{20}$~eV~\\cite{gamma}, from the Yakutsk muon data at $E>4 \\cdot 10^{19}$~eV~\\cite{Yak-gamma} and from the Pierre Auger Observatory data on the shower geometry at $E>2 \\cdot 10^{19}$~eV and $E>10^{19}$~eV~\\cite{PAO-gamma}) disfavour the SHDM scenario and even exclude it for particular values of the dark-matter parameters (see e.g.\\ Ref.~\\cite{Risse} for a more detailed discussion of some of these limits). Current experiments do not report any significant deviations from the global isotropy at the highest energies\\footnote{ After this paper was submitted, PAO reported a significant deviation from isotropy~\\cite{PAOagn}. The definite interpretation of this effect awaits further data (see e.g.\\ Ref.~\\cite{Comment} for discussion). }. This is however not conclusive both due to the low statistics and due to a limited field of view of any terrestrial-based installation.  The steeply falling spectrum of cosmic rays makes it very difficult to obtain a reliable measure of the global anisotropy in {\\em any} combination of terrestrial experiments. Indeed, the relative systematic difference in the energy estimation between two installations located in different parts of the Earth (and thus observing different parts of the sky) can hardly be made smaller than some $\\sim 15\\%$. Such a relative error would give $\\sim 30\\%$ higher integral flux seen by one of the experiments with respect to the other at the same reconstructed energy. Thus, possible observations of global anisotropy can be attributed both to a physical effect and to unknown systematics in the energy determination. Moreover, a seemingly isotropic distribution of the arrival directions over the sky might represent a physically anisotropic one masked by the systematic effects.  On the other hand, the planned space-based experiments, e.g.\\ TUS~\\cite{TUS}, JEM-EUSO~\\cite{EUSO} or S-EUSO~\\cite{S-EUSO}, will provide a unique opportunity to observe full sky with a single detector. While not being free from systematic uncertainties in the energy determination, an experiment of this kind would not introduce strong direction-dependent systematics and thus would be able to perform the studies of the global anisotropy at high confidence.  One may expect that in future space-based experiments with their huge exposures, particular sources of UHECR will be determined on event-by-event basis for the case of the astrophysical scenario. This is not an easy task, however: limited angular resolution together with large numbers of events would lead to identification problems similar to thoose of the gamma-ray astronomy\\footnote{See e.g.\\ Refs.~\\cite{EGRET-statistics} for discussions of statistical methods of analysis of photon-by-photon EGRET data. Currently, the gamma-ray sky at energies $(\\sim 0.1 \\div 1)$~GeV contains 101 identified source, 170 unidentified ones and a strong non-uniform diffuse background of unidentified origin.} but strongly enhanced due to magnetic deflections of the charged cosmic-ray particles. The searches for global patterns in the distribution of UHECR arrival directions will thus be important in any case.  A robust method for the study of any global asymmetry in the arrival directions is the harmonic analysis (see e.g.\\ Ref.~\\cite{Sommers}). It works perfectly if the predicted effect may be clearly seen in the first few harmonics but requires large statistics to reveal/exclude more complicated patterns. Here, we determine the optimal strategy for the searches of global anisotropies even with the low-resolution experiments (TUS) and for short exposure times.  The strategy is to fix two regions of the sky (not necessary covering full $4\\pi$) which are expected to provide the strongest contrast  in over/underdensity of events with respect to the null hypothesis of the isotropic distribution. The shape and the size of these regions, as well as the energy range of the events, are determined {\\it a priory} in order to balance the strength of the expected anisotropy (increasing for smaller regions and higher energies) and its statistical significance. The aim of this paper is to simulate optimal regions and energies for TUS, JEM-EUSO and S-EUSO, suitable to test the scenarios of extragalactic sources and of decays of superheavy dark matter concentrated in the Galactic halo. To this end, we perform new and improved (with respect to previous works) simulations of the distribution of arrival directions expected in both cases.  The paper is organised as follows. In Sec.~\\ref{sec:EG} and Sec.~\\ref{sec:DM}, we discuss our simulations for extragalactic sources and for SHDM decays, respectively. Sec.~\\ref{sec:experiments} gives specific predictions for three coming spaceborn experiments, TUS, JEM-EUSO and S-EUSO. Sec.~\\ref{sec:concl} contains our brief conclusions. Some technical details are collected in \\ref{app:density-function}. ",
        "conclusions": "\\label{sec:concl} In this work, we made quantitative predictions for the global anisotropy of the UHECR arrival directions expected in two distinct scenarios (``top-down'' and ``bottom-up'') of the origin of the highest-energy cosmic rays.  Several refinements and improvements resulted in considerable changes in the predictions as compared to previous studies. In particular, the patterns in the distribution of arrival directions of cosmic rays from astrophysical sources are more pronounced than expected before. The superheavy dark matter scenarios consistent with current data predict less pronounced anisotropy.  We developed optimal observables for distinction (exclusion) of these two scenarios. The actual required observational time depends strongly on currently uncertain duty cycle, fluorescent yield and the actual spectrum. In any case, some of the scenarios will be tested even with the limited statistics of the first space-based UHECR detector, TUS, before (or at the time of) the planned launch of JEM/EUSO. However, only JEM-EUSO and/or S-EUSO will be able to detect the SHDM component predicted by the ``minimal'' fit of the AGASA data consistent at 95\\% CL with the spectrum and with photon limits.  To observe, at the 95\\% CL, the patterns correlated with cosmic large-scale structure, one needs $\\sim 30$ events in the full-sky sample with energies $E>7 \\cdot 10^{19}$~eV. If these patterns show up at high confidence with lower statistics, this might indicate either significant underestimation of the particle energies or a problem in theoretical understanding of the origin and/or propagation of ultra-high-energy cosmic particles. One may try to use these patterns as a rough but independent tool for the energy calibration, quite important for space-based detectors.   \\ack The authors are indebted to D.~Gorbunov, G.~Rubtsov and D.~Semikoz for discussions. We acknowledge the use of online tools~\\cite{XSC,LEDA,GalacticExtinction}. This work was supported in part by the grants RFBR 07-02-00820 (OK and ST) and NS-7293.2006.2 (government contract 02.445.11.7370, OK and ST) and by the Russian Science Support Foundation fellowship (ST). Simulations of the cosmic-ray propagation have been performed at the computer cluster of the Theoretical Division of INR RAS.  \\appendix"
    },
    "0710/0710.0384_arXiv.txt": {
        "abstract": "Galactic winds from starbursts and Active Galactic Nuclei (AGN) are thought to play an important role in driving galaxies along the starburst-AGN sequence. Here, we assess the impact of these winds on the CO emission from galaxy mergers, and, in particular, search for signatures of starburst and AGN-feedback driven winds in the simulated CO morphologies and emission line profiles. We do so by combining a 3D non-LTE molecular line radiative transfer code with smoothed particle hydrodynamics (SPH) simulations of galaxy mergers that include prescriptions for star formation, black hole growth, a multiphase interstellar medium (ISM), and the winds associated with star formation and black hole growth. Our main results are: (1) Galactic winds can drive outflows of masses $\\sim$10$^8$-10$^9$\\msun which may be imaged via CO emission line mapping. (2) AGN feedback-driven winds are able to drive detectable CO outflows for longer periods of time than starburst-driven winds owing to the greater amount of energy imparted to the ISM by AGN feedback compared to star formation. (3) Galactic winds can control the spatial extent of the CO emission in post-merger galaxies, and may serve as a physical motivation for the sub-kiloparsec scale CO emission radii observed in local advanced mergers. (4) Secondary emission peaks at velocities greater than the circular velocity are seen in the CO emission lines in all models, regardless of the associated wind model. In models with winds, however, these high velocity peaks are seen to preferentially correspond to outflowing gas entrained in winds, which is not the case in the model without winds. The high velocity peaks seen in models without winds are typically confined to velocity offsets (from the systemic) $\\la$1.7 times the circular velocity, whereas the models with AGN feedback-driven winds can drive high velocity peaks to $\\sim$2.5 times the circular velocity. ",
        "introduction": "Observed relationships in galaxies between central black hole mass and stellar mass (e.g. Magorrian et al. 1998), velocity dispersion (i.e. the $M_BH$-$\\sigma$ relation; e.g. Gebhardt et al. 2000; Ferrarese \\& Merritt 2000), or galaxy structural properties (e.g. the black hole fundamental plane: Hopkins et al. 2007a,b) indicate a co-eval nature in supermassive black hole growth and star formation in galaxies.  These results have prompted a number of investigations in recent years to quantify this apparent self-regulation in star formation and black hole growth in galaxies, and their relationship to the formation and evolutionary history of the host system (e.g. Kauffmann \\& Haehnelt, 2000; Hopkins et al. 2006a,b; 2007c,d).   Over the last two decades, observations of local galaxies have painted a compelling picture in which galaxy mergers provide the link between massive starbursts and central black hole growth and activity.  ULIRGs (Ultraluminous Infrared Galaxies), for example, are a class of starburst galaxies with elevated infrared luminosities (\\lir $\\geq$ 10$^{12}$\\lsun) which typically show signs of interactions (e.g. Downes \\& Solomon, 1998; Sanders et al. 1988a; Scoville et al. 2000). While the intense infrared luminosity in these sources certainly owes in large part to the merger induced starbursts (e.g. Sanders et al. 1988a,b; Sanders \\& Mirabel, 1996), in some cases a contribution from a buried active galactic nucleus (AGN) may be non-negligible.  For example, many ULIRGs show spectral energy distributions (SEDs), infrared color ratios (e.g. F[25/60$\\mu$m]), polycyclic aromatic hydrocarbon emission (PAH) deficits, and emission line fluxes (e.g. [Ne V] at 14.3$\\mu$m) consistent with central AGN activity (e.g. Armus et al. 2004, 2006; Farrah et al. 2003; de Grijp et al. 1985).  Indeed, 35-50\\% of ULIRGs with luminosity above \\lir of 10$^{12.3} L_{\\sun}$ show optical and NIR spectra consistent with AGN activity (Kim, Veilleux \\& Sanders, 2002; Tran et al. 2001; Veilleux, Kim \\& Sanders, 1998). These results suggest that these merging systems are simultaneously undergoing a massive star formation and central black hole growth phase.  Observed similarities such as these between starburst galaxies, ULIRGs and quasars prompted Soifer et al. (1987) and Sanders et al. (1988a,b) to propose an empirically derived evolutionary sequence which connects galactic starbursts to quasars through galaxy mergers. In this picture, the fueling of the central black hole and AGN growth phase is realized during the major merger when gaseous inflows trigger nuclear starbursts as well as central black hole growth. The dusty galaxy transitions from a cold, starburst dominated ULIRG, to a warm, AGN-dominated ULIRG (where 'cold' and 'warm' refer to F[25/60$\\mu$m] ratios), and, as supernovae and stellar winds clear the obscuring gas and dust, to an optical quasar. In this sense, galaxy mergers provide a unique laboratory for studying the possible co-evolution of starbursts, black hole growth and activity, and spheroid formation.  Recent models have provided a theoretical foundation and further evidence for a merger-driven starburst-AGN connection in galaxies. Specifically, simulations by Springel, Di Matteo \\& Hernquist (2005a) have shown that galaxy mergers can fuel large-scale gaseous inflows (e.g. Barnes \\& Hernquist, 1991, 1996) which trigger nuclear starbursts (Mihos \\& Hernquist, 1996; Springel et al. 2005a) as well as promote central black hole growth (Di Matteo, Springel \\& Hernquist, 2005). Subsequent winds associated with the growth of central black holes can lift the veil of obscuring gas and dust, and along several sightlines produce a quasar with comparable lifetimes, luminosity functions, and observed $B$-band and X-ray properties to those observed (Cox et al., 2006b; Hopkins et al. 2005a-d; 2006a-d). The merger remnants quickly redden owing to gas depletion and the impact of feedback from star formation and black hole growth (e.g. Springel et al. 2005b) and resemble elliptical galaxies in their kinematic and structural properties.  Indeed, the population of stars formed during the starbursts accounts for the central ``excess light'' seen in ongoing mergers (e.g. Rothberg \\& Joseph 2004, 2006) and provides a detailed explanation for the luminosity profiles of old ellipticals (e.g.  Kormendy et al. 2007; see Mihos \\& Hernquist 1994a; Hopkins et al. 2007e,f,g).  A consensus picture has thus been borne out from these observations and simulations over the last two decades in which so called 'feedback' processes associated with winds from star formation and black holes act to self-regulate the growth of both the stellar and black hole masses in galaxies (e.g. Fabian, 1999; Silk \\& Rees, 1998). Indeed, the effects of starburst and AGN feedback-driven winds have been observed in local galaxies (e.g. Heckman et al. 2000; Martin, 2005; Rupke, Veilleux \\& Sanders, 2005a-c; Rupke \\& Veilleux, 2005; Tremonti, Moustakas \\& Diamond-Stanic, 2007), as well as those at high-\\z \\ (e.g. Narayanan et al. 2004; Pettini et al. 2002; Shapley et al. 2003) by way of absorption line outflows. However, the direct effect of feedback processes (especially from the highly efficient central AGN) on the emission properties of galaxies is not yet well characterized (see Veilleux, Cecil \\& Bland-Hawthorn, 2005 for an extensive review and associated references). In this sense, it is important to relate theories of galaxy formation and evolution which incorporate physically motivated models of feedback processes to observational signatures of winds across the electromagnetic spectrum.  In particular, observations of molecular gas in galaxies have proven valuable in characterizing the physics related to nuclear star formation and central AGN fueling in galaxy mergers as the molecular gas serves as fuel both for star formation as well as the central black hole(s).  For example, interferometric observations of molecular gas in ULIRGs show that most typically harbor $\\sim$10$^{10}$\\msun of molecular gas within the central 1.5 kpc (Scoville et al. 1986; Bryant \\& Scoville, 1999). Moreover, high resolution maps of dense molecular gas at submillimeter wavelengths in ULIRGs have revealed kinematic structures of double nuclei in mergers (Scoville, Yun \\& Bryant, 1997), bar-driven inflows (Sakamoto et al. 2004), and the density structure of gas fueling the central AGN (Iono et al. 2004).  With the increased spatial resolution and sensitivity afforded by the latest generation of (sub)mm-wave interferometers (e.g. the SMA, CARMA, PdBI), a number of recent observations have been able to pioneer investigations as to the effects of galactic winds on molecular gas emission in galaxies (e.g. Iono et al. 2007; Sakamoto, Ho \\& Peck, 2006; Walter, Wei\\ss \\ \\& Scoville, 2002). Detections such as these are expected to become more routine in upcoming years as the ALMA interferometer becomes available. An important complement to these current and forthcoming observations of molecular line emission from starburst galaxies and AGN are physical models which directly relate CO emission properties to galactic scale winds.  In this context, it is our aim to investigate the role that galactic winds can play on CO emission properties from starburst galaxies and AGN via numerical simulations. In particular, we focus on specific signatures of winds imprinted on CO morphologies and emission line profiles. In this work, we present self-consistent radiative transfer calculations for the emission properties of CO molecular gas in gas-rich galaxy mergers which account for the winds associated with both star formation and black hole growth.  In \\S~\\ref{section:numericalmethods}, we describe the hydrodynamic simulations, and radiative transfer methodology. In \\S~\\ref{section:spiral} we provide an example of our methods by applying our radiative transfer calculations to a model of a star-forming disk galaxy.  In \\S~\\ref{section:morphology}, we discuss the effect of winds on observed CO morphologies. We explore the response of emission line profiles to winds in \\S~\\ref{section:lineprofiles}, present a broader discussion of these results with respect to observations in \\S~\\ref{section:observations}, and summarize in \\S~\\ref{section:conclusions}. Throughout this paper, we assume a $\\Lambda$CDM cosmology with $h$=0.7, $\\Omega_\\Lambda$=0.7, $\\Omega_{\\rm M}$=0.3. ",
        "conclusions": "\\label{section:conclusions}   We have combined 3D non-LTE radiative transfer calculations with SPH simulations of galaxy mergers in order to investigate the effects of galactic-scale winds on the molecular line emission in starburst galaxies and AGN. We find that galactic winds are a natural result of merger-induced star formation and black hole growth. These winds can entrain molecular gas of order $\\sim$10$^8$-10$^9$ \\msun which imprints generic signatures in both the CO morphology as well as unresolved emission line profiles. The specifics of the morphological and emission line indicators of molecular outflows depend on physical parameter choices within the galaxy merger models. In particular, the energy source (i.e. BH accretion or star formation), as well as the merger orientation can vary the strength, direction, and duration of molecular outflows in emission contour maps, as well as the velocity separation of high velocity peaks in multi-component emission lines. Many of the results including the halflight radius as a function of excitation, or the velocity offsets of high velocity peaks in the emission lines vary enough between physical parameter choices such that these models may be used to constrain physical origins for observed CO morphologies and line profiles. In detail, we find the following:     \\begin{enumerate}  \\item Molecular outflows entrained in AGN feedback and starburst-driven winds can be directly imageable via CO emission line mapping. These types of outflows have been recorded in at least some local systems (e.g. Iono et al. 2007). These outflows will be detectable at cosmological distances given the predicted spatial resolution of ALMA.   \\item Molecular outflows entrained in winds driven by AGN feedback are typically longer lived than those entrained in starburst driven winds. This owes to the relative strength of AGN feedback-driven winds versus starburst driven winds.  \\item Winds from AGN feedback in coplanar merger models are typically more powerful than mergers which occur at more random orientations. This results in outflows in these models being visible for the majority of the $\\sim$200 Myr 'active period' studied here.  \\item The spatial extent of CO emission can be controlled by the presence of galactic winds.  The emission is seen to be stratified with transition in all models such that the CO (J=1-0) halflight radius is larger than the radius from higher lying transitions (e.g. CO J=7-6), though the degree of stratification depends on the inclusion of winds.  The relatively compact nature of observed CO emission in local mergers (R$_{1/2}$ typically confined to the central kpc; Scoville \\& Bryant, 1999) may be a consequence of AGN feedback or starburst-driven winds.  \\item In all models, high velocity peaks (peaks at velocities greater than the circular velocity) can exist superposed on the post-merger galaxy's broad CO emission line. In models without winds, these peaks owe to random kinematics of molecular gas. In models with winds, these peaks are seen to originate primarily from gas entrained in outflow, at least during the period of peak black hole accretion/star formation.  \\item For the models studied here (Table~\\ref{table:ICs}), high velocity peaks driven by random kinematics do not typically appear at velocity offsets (from systemic) greater than $\\sim$1.7 times the circular velocity of the post-merger galaxy. In contrast, peaks entrained in AGN feedback-powered winds can be driven to velocities near 2.5 times the circular velocity. The centroid velocities of the simulated lines are typically (1$\\sigma$) within $\\sim$20 \\kms of the true systemic velocity, and can generally be used as a reliable substitute for the systemic velocity in these models. Thus the above results hold true when measuring the velocity offsets of high velocity peaks with respect to line centroids as well.      \\end{enumerate}"
    },
    "0710/0710.2518_arXiv.txt": {
        "abstract": "{ {} {Using the high-level Balmer lines and continuum, we trace the density structure of two magnetospheric disk segments of the prototypical Bp star $\\sigma$\\,Orionis\\,E (B2p) as these segments occult portions of the star during the rotational cycle.  } {High-resolution spectra of the Balmer lines $\\ge$H9 and Balmer edge were obtained on seven nights in January-February 2007 at an average sampling of 0.01 cycles. We measured equivalent width variations due to the star occultations by two disk segments 0.4 cycles apart and constructed differential spectra of the migrations of the corresponding absorptions across the Balmer line profiles. We first estimated the rotational and magnetic obliquity angles. We then simulated the observed Balmer jump variation using the model atmosphere codes \\mbox{\\it synspec/circus} and evaluated the disk geometry and gas thermodynamics. } {We find that the two occultations are caused by two disk segments.  The first of these transits quickly, indicating that the segment resides in a range of distances, perhaps 2.5$-$6R$_{*}$, from the star.  The second consists of a more slowly moving segment situated closer to the surface and causing two semi-resolved absorbing maxima. During its transit this segment brushes across the star's ``lower\" limb. Judging from the line visibility up to H23-H24 during the occultations, both disk segments have mean densities near 10$^{12}$ cm$^{-3}$ and are opaque in the lines and continuum. They have semiheights less than ${\\frac 12}$R$_{*}$, and their temperatures are near 10\\,500\\,K and 12\\,000\\,K, respectively. } {In all, the disks of Bp stars have a much more complicated geometry than has been anticipated, as evidenced by their (sometimes) non-coplanarity, de-centerness, and from star to star, differences in disk height. } ",
        "introduction": "Over the last several years the interest in the class of magnetic Bp stars has spurred the accumulation of datasets at many wavelengths over the stars' rotational/magnetic cycles. These datasets support the view that the circumstellar environment is the site of a balance of physical processes that channels the wind toward the magnetic plane. If the magnetic field is strong, that is if the circumstellar magnetic energy density exceeds the wind energy density, the star's wind is channeled into the magnetic plane where a shock is formed and the cooled shock settles to form a stable disk. The magnetic axis is typically inclined with respect to the rotational axis. Then, as different portions of the disk wobble in front of the star during the rotational cycle, they produce alternate absorption and emission components in the H$\\alpha$, UV resonance lines, and low-excitation metallic lines  (e.g., Bohlender et al. 1987, Shore 1987, Bolton 1994). To complicate the picture, it appears that slow infall of disk matter is responsible for producing inhomogeneous metal-poor patches or a belt along the magnetic equator (e.g., Khokhlova et al. 2000, Groote 2003). Also, although there remains a theoretical difficulty with the fractionation hypothesis for helium (Krti\\v{c}ka et al. 2006), enough neutral helium atoms in fact seem likely to separate from the polar wind, return to the surface, and accumulate there as helium caps.  Recent theoretical refinements in the picture (e.g., Preuss et al. 2004, Townsend \\& Owocki 2005; hereafter TO05) suggest that the wind will settle onto low equipotential surfaces determined by the balance of radiative, gravitational, magnetic, and centrifugal forces. These surfaces reside primarily, though not exclusively, in the disk plane. Because the disk is likely to have a nonaxisymmetric density distribution, we will refer to the accumulations near the plane responsible for time-variable spectral features as disk segments. TO05 have presented a ``rigidly rotating magnetospheric\" {\\it ab initio} model to explain physical properties of these segments for the case of $\\sigma$\\,Ori\\,E and it can be applied to the cases of other Bp stars with arbitrary magnetic inclinations and viewed from any rotational inclinations. TO05 find that the disk plane is not precisely perpendicular to the magnetic polar axis for intermediate magnetic obliquities, including those they obtained for $\\sigma$\\,Ori\\,E.  The spectroscopic anomalies of  $\\sigma$\\,Ori\\,E (HD\\,37479, B2\\,Vp; Lesh 1968) were first noticed by Berger (1956) and discussed further by Osmer and Peterson (1972). This star is the prototype and most studied of the helium-strong subclass of Bp stars. Walborn (1974) discovered H$\\alpha$ emissions in this star's spectrum that varied on a timescale of hours. This variability was subsequently established to be associated with the rotation of a frozen-in disk, as described above. Hesser et al. (1977) determined a period of 1.19081 days based on several nights of photometric observations distributed over a few years. Subsequently, various studies (Landstreet \\& Borra 1978, Bohlender et al. 1987) determined that the magnetic variations (suggestive of an approximately dipolar field with polar field strength B$_p \\approx$ 10\\,kG) are consistent with the photometric period. Various periods over the range 1.19081$-$1.19084 days have since been determined by several authors from analyses of variations of photospheric and circumstellar H$\\alpha$ line emissions (Hesser, Moreno, \\& Ugarte 1977, Reiners et al. 2000, hereafter R00; Townsend, Owocki, \\& Groote  2005; hereafter TOG05).  Groote \\& Hunger (1976, 1977 hereafter GH76 and GH77) discussed the variations of high-level Balmer lines based on only photographic spectra. The GH76 data indicated that two disk occultations  of the star, about 0.4 cycles apart, cause high-level Balmer line absorptions visible at these times. At the same times, the intervening disk material absorbs flux just shortward of the Balmer edge. Although the analysis of line absorptions do not generally allow us to infer the extent and total volume of such a corotating structure, these absorptions do yield considerable information about its three dimensional geometrical character. Smith \\& Groote (2001, hereafter SG01) combined strength and shape information of many ultraviolet metallic lines obtained by the {\\it International Ultraviolet Explorer} around the rotation cycle to exploit this potential. Their study provides some of the first estimates of the physical parameters of the disk gas, including disk temperature, column density, and areal coverage of the star.  The limitations of this study were the low signal-to-noise ratio and the paucity of phase sampling of the archival database. What has been needed is a study that includes complete rotational phase coverage  of absorption lines sensitive to electron density.  To respond to this need, we have obtained complete phase coverage of the hydrogen lines from H9 (H$\\zeta$) to the Balmer edge over the entire rotation period of $\\sigma$\\,Ori\\,E. Such datasets have the potential to better elucidate areal coverage and density structure of the absorbing disk segments than has previously been possible. ",
        "conclusions": "We have used high-level Balmer absorption data finely sampled around the rotation cycle to show that the distribution of plasma in the magnetosphere of $\\sigma$\\,Ori\\,E is more complicated than previously believed. Most especially, the secondary occultation is comprised of two semi-resolved events.  In addition, there seem to be at least two brief appearances of absorbing matter well beyond the disk plane. The disk plane is dominated by two segments which occult portions of the star some 0.4 cycles apart. We have no direct information from absorption profiles alone on the distribution of matter at other azimuths in the plane and so look forward to future analyses based on the H$\\alpha$ and/or H$\\beta$ emissions over the cycle.  The occulting disk segments are sections of a ring which occult the star as the frozen disk corotates around the star. The segment causing the primary occultation lies further from the star (between limits of $\\approx$2.5R$_{*}$ and 6R$_{*}$) than the opposing segment. There are at least three arguments that point to the disk segments lying at different orbital radii. First, gas in the segment causing the primary occultation has a lower temperatures (10\\,500\\,K, compared to 12\\,000\\,K). This can be inferred first from the lack of an enhanced He\\,I feature and also from the absence of weak absorptions of metallic-line blends referred to in $\\S$4.3. Second, the acceleration of spectral migrating components across the line profile is higher for this segment. Third, our occultation analysis program indicates that one cannot reconcile the radii of the two segments even if one attributes the broad secondary absorption maximum as being due to viewing two occultations of a single warped disk segment.   Perpendicular to the magnetic disk, the semiheights of the two segments are some 0.3$-$0.45R$_{*}$. It is important to stress that the disk of $\\sigma$\\,Ori\\,E has a smaller semiheight than disks for other disk-harboring Bp stars that we have analyzed with similar modeling techniques (e.g., SG01, Smith et al. 2006). It is tempting to speculate that a disk should be more compressed toward the plane due to the pressure provided by a large magnetic field (e.g., Babel \\& Montmerle 1997). However, this cannot be the full story because the most relevant quantity, the ratio ``$\\eta$\" of magnetic to wind energy densities is actually larger for 36\\,Lyn than for $\\sigma$\\,Ori\\,E and yet the disk of the former has a somewhat larger height (Smith et al. 2006). Might other considerations, such as the larger radial extent of the disk (e.g. in the case of 36\\,Lyn), have a bearing on this question?  In addition to these dimensional properties, we find that  a phase discrepancy of +0.018 cycles exists between the phase at which the central absorption transits the center of the hydrogen line and the phase of EW maximum. This agrees with a similar phase difference found by GH77 and calls into definition of the star's period when determined from different observation modes. The secondary occultation is distinguished by the passage of two weak lobes, the phases of which are consistent with the semi-resolved structure of the light curves in the literature. The interpretation of this structure is not clear, but it may be related to a disk ``warping\" or a corrugation of the disk plane predicted in the TO05 models.  When we view the disk edge-on through the magnetic plane, the Balmer line absorptions are opaque, and thus the absorptions of the high-level members of this series increase.  For H9 the optical depth is about 100, while it is $\\approx 10$ for H20, and $\\approx 2$ at the Balmer edge. As the disk ingresses and egresses there is an absence of an absorption tail in the equivalent width curves. This confirms the theoretical expectation (e.g., TO05) that the density scale height is very small compared to the stellar radius.  This implies in turn that our homogeneous slab models of the disk and the total mass estimate of about 1$\\times$10$^{-10}$ M$_{\\sun}$ are only rough approximations.  How would hypothetical data of indefinitely high quality (signal-to-noise, spectral resolution, and cadence) improve on our analysis of the disk of $\\sigma$\\,Ori E? We can think of the following applications of enhanced quality data:  One improvement would be in the mapping of vertical density distribution and indeed the separation of the two density maxima associated with the secondary occultation. This would result from tracing details of the absorption subfeatures as they migrate across the line profiles. An improvement in our understanding of the vertical stratification should indirectly improve out understanding of whether the disk is warped and the inner edge of the secondary occultation-causing segment extends almost to the star or not. It will also clarify the cause of the phase delays that both we and GH76 have found among various hydrogen lines and with respect to the continuum. A resolution of this issue would allow an improved determination of ephemerides from flux curves derived from different diagnostics, thereby leading to a definitive determination of the period and perhaps its first derivative.  Second, the changes in the highest-level lines detected during the ingress and egress of the occultations can be used in principle to further map the vertical density distribution of the disk segments because of the different optical depth regimes in which the various lines are formed.  Third, highly accurate difference profiles should allow one to search for asymmetric absorption features in the EW curves that could indicate the rate at which disk particles leak from the disk either through the inner or outer edge.  Fourth, increased data quality improvements (assuming a stable spectrograph and atmospheric transmission) can lead to the observation of subtle changes in disk structure. We suggest that these might occur on timescales of one to a few months or longer.  Finding such changes could allow one ultimately to correlate them with the expected ``break out\" of disk matter through the outer edge (ud-Doula, Townsend, \\& Owocki 2006), possibly associated with X-ray flares from this star, (Groote \\& Schmitt 2004, Sanz-Forcada, Franciosini, \\& Pallavicini 2004)."
    },
    "0710/0710.2668_arXiv.txt": {
        "abstract": "We report the detection of 10 new X-ray filaments using the data from the {\\sl Chandra} X-ray satellite for the inner $6^{\\prime}$ ($\\sim 15$ parsec) around the Galactic center (GC). All these X-ray filaments are characterized by non-thermal energy spectra, and most of them have point-like features at their heads that point inward. Fitted with the simple absorbed power-law model, the measured X-ray flux from an individual filament in the $2-10$ keV band is $\\sim 2.8\\times10^{-14}$ to $10^{-13}$ ergs cm$^{-2}$ s$^{-1}$ and the absorption-corrected X-ray luminosity is $\\sim 10^{32}-10^{33}$ ergs s$^{-1}$ at a presumed distance of 8 kpc to the GC. We speculate the origin(s) of these filaments by morphologies and by comparing their X-ray images with the corresponding radio and infrared images. On the basis of combined information available, we suspect that these X-ray filaments might be pulsar wind nebulae (PWNe) associated with pulsars of age $10^3 \\sim 3\\times 10^5$ yr. The fact that most of the filament tails point outward may further suggest a high velocity wind blowing away form the GC. ",
        "introduction": "The Galactic center (GC) is the only place where we can observe parsec details of various interaction in and around the Galactic nucleus. Advances in this research frontier rely primarily on observations at radio, infrared and X-ray wavelengths, because the optical band suffers seriously from a considerable extinction with an $Av\\sim 30$ (e.g., Becklin, Matthews, Neugebauer \\& Willner 1978). One of the most important discoveries in the GC region is perhaps the presence of many structured non-thermal radio filaments (NTFs) (e.g., Yusef-Zadeh, Morris, \\& Chance 1984; Morris \\& Serabyn 1996; LaRosa, Kassim, Lazio, \\& Hyman 2000). While these non-thermal radio filaments have been intensively studied, their origins and implications on the underlying physical processes around the GC region remain largely unclear. In the current analysis of X-ray filaments around the GC region, these X-ray emitting particles are usually expected to be fairly close to their acceleration zone and evolve very rapidly in time. Thus, the X-ray study of the same region would be essential to probe the origin of these energetic particles. The {\\sl Chandra} Galactic center Survey (CGS), with its unprecedented high spatial resolution of $\\sim 0.5''$ and moderate spectroscopy capability, has already revealed remarkable X-ray structures (including thousands of X-ray bright point sources and some filaments, as well as clumps of diffuse emission) within the central $200$ pc of our Galaxy (e.g., Wang, Lu, \\& Lang 2002a). In this paper, we mainly concentrate on the nature of those thread-like linear structures or filaments as observed in X-ray bands.  Up to this point within $15^\\prime$ ($\\sim 37$ pc at 8 kpc) from \\hbox{Sgr~A$^*$} where a $\\sim 4\\times 10^6M_{\\odot}$ black hole resides inside a compact region of less than $\\sim 1$AU (e.g., Shen et al. 2005), 5 X-ray filaments have been studied in details (see also Table 1). For example, G359.95-0.04, a comet-like filament at $\\sim 0.32$ pc north of \\hbox{Sgr~A$^*$}, is thought to be a ram-pressure confined pulsar wind nebula (PWN) (Wang et al. 2006). Another prominent filament, G359.89-0.08 (SgrA-E), at $\\sim 7$ pc southeast of \\hbox{Sgr~A$^*$}, was first noticed by Sakano et al. (2003) and an interpretation of a possible PWN origin was discussed in details by Lu, Wang \\& Lang (2003). An alternative picture of G359.89-0.08 as a source of synchrotron emission from relativistic particles accelerated by a shock wave of W28 SNR was suggested recently by Yusef-Zadeh et al. (2005); in that same paper, they also detected a new X-ray filament G359.90-0.06, which coincides spatially with a radio filament at $\\sim 5.8$ pc southwest of \\hbox{Sgr~A$^*$}. They explored the mechanism of an inverse Compton scattering (ICS) for the X-ray emission of G359.90-0.06. Another 2 filaments were found in more extensive regions. G0.13-0.11 in the Radio Arc region was first reported by Wang, Lu \\& Lang (2003) and was suspected to be also a PWN. Of particular interest is the X-ray filament G359.54+0.18; it associates with only one strand of the obviously bifurcated radio threads (e.g., Yusef-Zadeh et al. 2005; Lu et al. 2003). A common feature shared in the X-ray energy spectra of these filaments is that they all appear to be non-thermal. It should be noted however that any thermal component to these sources would likely be completely absorbed and unobservable, given the high foreground column density.  In these earlier investigations, pulsar wind nebulae (PWNs) and supernova remnants (SNRs) seem to offer natural explanations for the appearance of such X-ray filaments. Indeed, it is believed that a considerable number of supernovae should have happened in the GC region (e.g., Figer et al. 1999, 2004; Wang et al. 2006). One would naturally expect to find some of their end-products such as pulsars and SNRs in the GC region. However, no radio pulsars have yet been found within $\\sim 1^\\circ$ of the GC (Wang et al. 2006). This might be caused by difficulties in radio observations (Cordes \\& Lazio 1997; Johnston et al. 2006). Seeking observational clues in X-ray bands might shed new light to the search for radio pulsars embedded in the GC region.   Another tempting idea is to use these X-ray filaments as potential tracers for the magnetic field and gas dynamics around the GC region, since the magnetic fields should have played a significant role in producing such non-thermal spectra and the thread-like shapes of these filaments and the gas motion is usually coupled with the magnetic fields. (e.g., Chevalier 1992; Boldyrev \\& Yusef-Zadeh 2006). Magnetic fields exist on all scales of the Galaxy as well as in other spiral galaxies and generally trace out spiral patterns on large scales (e.g., Beck \\& Hoernes 1996; Fan \\& Lou 1996; Zweibel \\& Heiles 1997; Wielebinski 2005; Ferriere 2001, 2007), and great progress has been made in measuring them and inferring their influence by various means. For example, the observed high-energy cosmic ray anisotropy at a few times $10^{-3}$ (e.g., Amenomori et al. 2006) might be physically related to large-scale structures of Galactic magnetic fields and inhomogeneous cosmic ray source distribution. Using diffuse synchrotron radio emissions at 74 and 330 MHz frequencies produced by relativistic cosmic-ray electrons and the magnetic field around the Galactic center, LaRosa et al. (2005) inferred a weak magnetic field of order of $\\sim 10 \\mu$G on size scales $\\gsim 125''$ based on a minimum-energy analysis. This is about 2 orders of magnitude lower than $\\sim 1$ mG usually estimated for the GC region. Very recently, Cowin \\& Morris (2007) argued that the assumption of LaRosa et al. (2005) that the magnetic field and cosmic rays are in a minimum-energy state across this region is unlikely to be valid. According to their model estimates, the mean magnetic field is at least 100 microgauss on a scale of several hundred parsecs and peaks at approximately 500 microgauss at the center of the diffuse nonthermal source (DNS). This is an important issue to be settled for the GC magnetic environment. Most of the GC radio NTFs are found to be perpendicular to the Galactic plane, implying a local poloidal magnetic fields of about milli-gauss strengths (e.g., Yusef-Zadeh \\& Morris 1987). However, recent observations revealed that the GC magnetic field may be more complex than a simple globally ordered dipolar field (e.g., LaRosa et al. 2004; Nord et al. 2004). It might be possible for X-ray NTFs to also provide clues of the configuration of the local magnetic field as well as the interaction between it and the ambient gas flow.  We report morphological and spectral properties of another 10 newly discovered X-ray filaments within a region of $6^\\prime .5 \\times 6^\\prime .5$ surrounding the \\hbox{Sgr~A$^*$} (roughly corresponding to a projected sky area of $\\sim 15$ pc by $\\sim 15$ pc at a presumed distance of 8 kpc to the GC). Their plausible physical origins are discussed in section 4. All error bars in the X-ray spectrum parameter measurements are at the 90\\% confidence level, and we express the fitted parameters in the format of $y\\ (y_l,\\ y_u)$, where $y$ is the best fit value while $y_l$ and $y_u$ are the lower and upper limits of the 90\\% confidence interval, respectively. For all the images in the paper, north is up and east is to the left. We shall adopt a distance of 8 kpc from the solar system to the GC throughout this paper. We note that during the review process of this manuscript, Muno et al. (2007) submitted a paper also on X-ray filaments around the Galactic center. ",
        "conclusions": "\\subsection{\\label{subsec:origins}Properties and possible origins of GC X-ray filaments }  While X-ray photon numbers may not be high enough to constrain the exact shapes of the energy spectra, all spectra appear featureless except for filament F3 showing weak iron line features (see Section 3.3) and can be well fitted with power-law models. Fitting of some of these energy spectra with thermal emission models is acceptable statistically; nevertheless, this always gives quite high temperatures, i.e., $>10$ keV. We therefore incline to the view that X-ray emissions from these filaments are non-thermal in nature.  Table 2 sums up the inferred parameters for these 15 non-thermal X-ray filaments in the inner $15^{\\prime}$ around the GC. In addition to the 10 filaments studied in this paper, we also include the other 5 filaments, namely G359.89-0.08, G359.90-0.06, G359.95-0.04, G359.983-0.046, and G0.13-0.11, reported and studied earlier in the literature. Most of their hydrogen column density $N_H$ are of the order of $\\sim 10^{22}-10^{23}~{\\rm cm}^{-2}$, consistent with other $N_H$ estimates around the GC region. Most of their photon indices $\\Gamma$ fall within the range of $\\sim 1-2.5$. The chance of these X-ray filaments being background extragalactic sources is very small according to the spectral and morphological properties. For instance, it would be very difficult to explain the linear filamentary morphology using the hypothesis of extragalactic origins. Therefore, these X-ray filaments are most likely unique objects around the GC region.  As already discussed in Section 1, there were suggestions that these X-ray filaments may be ram-pressure confined PWNe (Wang et al. 2003, 2006), or synchrotron emissions from MHD shocks associated SNRs or emissions resulting from inverse Compton scattering (Yusef-Zadeh et al. 2005; Figer et al. 1999). The non-thermal X-ray emission mechanisms may be either synchrotron emission or inverse-Compton scattering. Magnetohydrodynamic (MHD) relativistic pulsar winds (Michel 1969; Goldreich \\& Julian 1970; Kennel \\& Coroniti 1984a, b; Lou 1996, 1998) and MHD shock interactions of magnetized outflows (e.g., Yu \\& Lou 2005; Yu et al. 2006; Lou \\& Wang 2006, 2007) with the interstellar medium (ISM) in SNRs and PWNe could provide high-energy electrons needed in these two radiation mechanisms (e.g., Sakano et al. 1993; Lu et al. 2003; Wang et al. 2006). If these X-ray filaments are SNRs, their elongations would probably represent MHD shock fronts and therefore, one would not expect to see a tendency of spectral softening along a filament. The two bright filaments F3 and F10 both show evidence for such a softening tendency. The fact that most of these filaments have point-like sources at the heads also againsts the SNR origin. On the other hand, as nonthermal X-ray emission is only detected in several SNRs within the entire Galaxy, it would be highly unlikely that there are so many nonthermal SNRs around the GC region. For this reason, we would argue that most of these X-ray filaments are not SNRs. Observed properties of these X-ray filaments may be more consistent with those of PWNe. Typical features of a PWN are: non-thermal X-ray spectrum, with photon index of $1.1-2.4$ and a X-ray luminosity $L_x$ range from $4\\times 10^{32}$ to $2\\times 10^{37}~{\\rm erg}~{\\rm s}^{-1}$ in $0.2-10$ keV band (e.g., Gaensler \\& Slane 2006; Kaspi, Roberts \\& Harding, 2006). The $\\Gamma$ and $L_x$ of these 10 filaments are consistent with the values of a PWN. The existence of point-like X-ray sources as indicated by the image study also tends to favor a PWN scenario.   We may estimate the ages of the putative pulsars with the X-ray luminosities of these X-ray filaments. Li et al. (2007) studied statistically the nonthermal X-ray emission from young rotation powered pulsars and PWNe. They noted that there exists a correlation between the pulsar age $\\tau$ and the $2-10$ keV PWN luminosity $L_{x,pwn}$, which can be expressed as $L_{x,pwn}=10^{41.7}\\tau^{-2.0\\pm 0.3}$. The X-ray luminosities of these 10 filaments are in the range of 0.2-2.2$\\times 10^{33}$ erg s$^{-1}$. Using this empirical formula, ages of these putative pulsars are possibly between $\\sim 10^3$ to $3\\times10^5$ yr. However, given the dispersion about the above empirical relationship (Li et al. 2007), the estimate may be uncertain probably by a factor of 10.  Since the ages of pulsars with bright PWNe are usually younger than a few tens of thousand years, one may doubt if a pulsar at the age of several $10^5$ years can produce a detectable X-ray nebula. However, the PWN of a relatively old pulsar can be enhanced in surface brightness and thus become detectable if the pulsar wind materials are confined to one direction. PSR B0355+54 is $\\sim 5.6\\times10^{5}$ yr old. It converts $\\sim 1$\\% of its spin-down luminosity to the cometary-like X-ray nebula (e.g., Tepedelenlio\\v{g}lu \\& \\\"{O}gelman 2007). The old pulsar PSR B1929+10 ($\\tau\\sim3\\times10^6$ yr) also converts $2.1\\times10^{-4}$ of its spin-down luminosity $\\sim 3.9\\times10^{33}$ erg s$^{-1}$ into the emission of the cometary nebula (e.g., Becker et al. 2006). The X-ray filaments identified in the GC region are similar to these two systems and thus probably powered by pulsars.  Now we discuss whether the number of X-ray filaments, if identified with PWNe, would be consistent with the estimated star formation rate in the GC region. According to Figer et al. (2004), the star formation rate at the GC is about $10^{-7} M_{\\odot}\\hbox{ yr}^{-1}\\hbox{ pc}^{-3}$ which is some 250 times higher that the mean star formation rate in the Galaxy. In the field of view of our Fig. 1, we take a radius of about $7$ pc and estimate the star formation rate to be $1.4\\times 10^{-3}M_{\\odot}\\hbox{ yr}^{-1}$. If the mean mass of a star is 10$M_{\\odot}$, the frequency of supernova explosions would be $1.4\\times 10^{-4}$ per year, leading to about 40 pulsars in the field of Fig. 1 younger than $\\sim 3\\times 10^5$ yr as estimated above. This number is roughly consistent with the 15 candidate PWNe identified in the field.  \\subsection{\\label{subsec: F10_radio}G0.029-0.06 (F10) and its radio counterpart}  Filament F10 bears certain unique features to be noted here. First, it has the longest linear structure with the entire image slightly bent towards the northeast, more or less like an arc. Second, it is the farthest away from \\hbox{Sgr~A$^*$} and thus has much less contamination from the strong diffuse X-ray emission of Sgr A. Third, there is an obvious 20cm radio NTF coincident spatially with X-ray filament F10.  The spectral indices for different regions along filament F10 show evidence of spectral steepening from the ``head\" to ``tail\" (see Table 1). When $N_H$ is fixed at the best fit value $\\sim 7\\times 10^{23}~{\\rm cm}^{-2}$, the $\\Gamma$ values for the ``head\", ``middle\", and ``tail\" regions are 1.1(1.0, 1.3), 1.5(1.4, 1.6), and 1.8(1.7, 1.9), respectively. This might suggest an energetic particle flow direction from the southeast (head) to the northwest (tail). A pulsar moving through the magnetized interstellar medium seems to give a plausible explanation of this scenario. Indeed, the morphology of F10 does imply a point source in the ``head\" region. The corresponding point spread function (PSF) at G0.029-0.06 is an ellipse with a size of $\\sim 2\\arcsec\\times 4\\arcsec$. For an updated X-ray versus spin-down luminosity correlation of rotation powered pulsars, a modified empirical relation is given by equation (3) of Possenti et al. (2002), namely, $\\log L_{x,(2-10)}=1.34~\\log\\dot{E}-15.34$ where $L_{x,(2-10)}$ is the X-ray luminosity in $2-10$keV energy band; using this empirical relation, we would have a $\\dot E\\sim 10^{36}~{\\rm ergs}~{\\rm s}^{-1}$. Since PSRs J1747-2958 and B1929+10 convert about 2.5\\%  and 2.1$\\times10^{-4}$ of their spin-down powers to their cometary X-ray nebulae (e.g., Gaensler et al. 2004; Becker et al. 2006), the ratio $L_x$/$\\dot{E}$ of F10 ($\\sim 10^{-3}$) indicates that the above estimate for $\\dot{E}$ is reasonable.  The arc-like X-ray morphology of F10 and its coincidence with a radio NTF might be a good indicator of its interaction with the interstellar magnetic field environment of the GC region (Lang et al. 1999; Wang et al. 2002b). Similar to the discussion about G0.13-0.11 by Wang et al. (2002b), we may estimate the magnetic field strength $B$ in the current context. First, the lifetime $\\tau$ of synchrotron X-ray emitting particles is given by $\\tau\\sim (1.3 {\\rm~yrs})\\epsilon^{-0.5} B_{\\rm mG}^{-1.5}$, where $\\epsilon$ is the X-ray photon energy in unit of keV (a value of 4 keV is adopted here) and $B_{\\rm mG}$ is the magnetic field strength in the filament volume in units of mG. The simulations of Bucciantini et al. (2005) show that the average flow speed in the tail is about 0.8-0.9 $c$. For a sustained X-ray linear structure, we estimate by requiring $\\tau \\geq L_{obs}/(0.85 c)$. Adopting a characteristic angular length $L_{obs}$ of $47\\arcsec$ ($\\sim 2 {\\rm pc}$), we thus infer a magnetic field strength $B\\sim 0.3$mG, similar to those in the bright radio NTFs (e.g, Yusef-Zadeh \\& Morris 1987; Lang et al. 1999).  We try to outline a few plausible scenarios in the present context and discuss relevant aspects qualitatively. Magnetized neutron stars move with peculiar speeds in the range of a few tens of kilometers per second (a mean space velocities of $\\sim 300-400\\hbox{ km s}^{-1}$ for young pulsars; Hobbs et al. 2005; Faucher-Gigu\\`ere \\& Kaspi 2006) to well over one thousand kilometers per second ($\\sim 1600\\hbox{ km s}^{-1}$) and the surrounding ISM is generally magnetized. Generally speaking, a typical peculiar velocity of a neutron star is supersonic and super-Alfv\\'enic in a magnetized ISM. Neutron stars or pulsars have different ranges of surface magnetic field strengths: $10^9-10^{10}$G for millisecond pulsars in binaries, $10^{11}-10^{12}$G for a wide range of pulsars, and $10^{14}-10^{15}$G inferred for several magnetars. Several situations may happen. (1) If a pulsar does not involve an active pulsar wind, its peculiar motion through the surrounding magnetized ISM would sustain an MHD bow shock by its magnetosphere as well as a magnetotail. The faster the pulsar moves, the more linear the system would appear. This is basically like a bullet moving through an air supersonically and generating a Mach cone or wake. Relativistic electrons can be produced at the MHD bow shock and synchrotron emissions can be generated and sustained at the same time. (2) In a binary system, the fast wind (say, with a speed higher than 1000 km s$^{-1}$) from a companion star can blow towards a spinning magnetized pulsar in orbital motion. Here, the situations of a companion fast wind blowing across a magnetized pulsar and a pulsar moving through the ISM with a high speed are more or less equivalent. Again, an MHD bow shock and a magnetotail can form in association with the pulsar system. The stronger the companion wind and the faster the pulsar moves, the more linear the pulsar system would appear. Relativistic electron and/or positrons can be generated and sustained to power synchrotron emissions in the bow shock draped around the pulsar magnetosphere. For such a system, one might be able to detect the presence of the companion by various independent means. (3) For a pulsar emitting an active pulsar wind and with misaligned magnetic and spin axes, spiral forward and reverse shock pairs can be generated in the relativistic pulsar wind as a result of inhomogeneous wind and eventually the pulsar wind is stopped by the ISM through a MHD termination shock (e.g., Lou 1993, 1996, 1998). (4) Case (3) can also happen for a pulsar moving with a high peculiar velocity through a magnetized ISM. (5) Case (3) can also happen for a pulsar in binary orbital motion with the companion blows a powerful wind with a speed higher than 1000 km s$^{-1}$. In both cases of (2) and (5), the center of mass of the binary system may also move with a high speed through the ISM. One can further speculate several possible combinations along this line of reasoning (e.g., Chevalier 2000; Toropina et al. 2001; Romanova, Chulsky \\& Lovelace 2005).  There are two clumps (referred to as the east and west clumps hereafter) of diffuse X-ray emission surrounding filament F10. Although the west clump is also elongated, it is not called a filament because it contains many substructures. To see if these clumps are physically related to F10, we extract their energy spectra separately (see the middle and bottom panels of Fig. 6). The fitted parameters are listed in Table 3. The much higher absorbing column densities of the two clumps indicate strongly that they are located farther away from us than F10 is, while these two clumps themselves are almost at the same distance (see Table 3). Moreover, their X-ray emissions are very likely powered by the same mechanism, as hinted by the characters of their spectra, which can be fitted well with an absorbed power-law model plus a 6.4 keV emission line. The total emission comes mostly from the photon energy $4-7$ keV band, with a strong 6.4 keV neutral Fe K line. In contrast, the Fe line feature is not present in the spectrum of F10. A possible explanation for this non-thermal, apparently broadened iron line emission at 6.4 keV is the collisions of low-energy cosmic-ray electrons with irons in molecular clouds (e.g., Valinia et al. 2000) or by the radiative illumination from the GC massive black hole that was suggested to be very bright in the past (e.g., Koyama et al. 1989). In conclusion, F10 and the surrounding clumps do not seem to interact directly.  \\subsection{\\label{subsec:B-fields}X-ray NTFs as tracers of the small-scale magnetic fields and gas dynamics}  The oritentation of the X-ray filaments provides an opportunity to probe the physics conditions of the GC region. As discussed in section 4.1, the cometary shapes of the X-ray filaments imply that the pulsar wind particles are confined to one direction by the ambient materials.  Mechanism shaping the filamentary structure could be ordered magnetic fields (e.g., Yusef-Zadeh \\& Morris 1987; Lang, Morris, \\& Echevarria 1999) and/or high relative velocity between the pulsars and the surrounding gas (e.g., Wang et al. 1993; Shore \\& LaRosa 1999). The magnetic field could be the product of the gas motion, and the magnetic field could also control the motion of the gas (e.g., Heyvaerts et al. 1988; Chevalier 1992). Radio NTFs suggests that the magnetic field is poloidal at large-scale (e.g., Yusef-Zadeh \\& Morris 1987; Lang, Morris, \\& Echevarria 1999) with some more complex smaller structures around the GC region (e.g., Nord et al. 2004). By looking at the X-ray images shown by Figures 1, 2 and 3, there seems to be a tendency that the PWN tails point away from the GC, indicating that the pulsar wind particles are blown outward.  This might imply the presence of a Galactic wind of hot plasma blowing away from the center, given the high star formation rate (and so plenty of hot gas) in this region. Pulsars may have typical peculiar velocities of $\\sim 400\\hbox{ km s}^{-1}$ (e.g., Hobbs et al. 2005; Faucher-Gigu\\`ere \\& Kaspi 2006) and we would expect them to move in random directions. The tendency for the structures of ten PWNs to orient away from the center seems to suggest that the Galactic wind has a speed comparable to or greater than $\\sim 400\\hbox{ km s}^{-1}$.  In the above scenario, the particle flow direction of the X-ray filament F3 should be from the southwest (closer to Sgr A*) to the northeast. This suggests that the pulsar, the origin site of the particles, is actually in the ``tail'' region defined in Fig 2. Then the evident spectral steepening from the southwest to the northeast (see section 3) can be naturally expained. Therefore, the spectral evolution along F3 also supports the existence of a radial high velocity wind in the GC region.  While the X-ray filaments may not completely overlap with their radio NTFs, their overall orientations are similar. This is supported by the four X-ray NTFs (including filament F10 in our analysis) that have radio counterparts: G359.54+0.18 overlays exactly on the northern part of the two radio filaments (e.g., Yusef et al. 2005; Lu et al. 2003); G359.89-0.08 and its radio counterpart SgrA-E overlaps partly and extends in the same direction, with a centroid offset of $\\sim 10 \\arcsec$ (e.g., Yusef et al. 2005; Lu et al. 2003); G359.90-0.06 (SgrA-F) (e.g., Yusef et al. 2005) and G0.029-0.06 (F10) (see our subsection 3.10) also show similar spatial property. Generally speaking, the X-ray NTFs tend to be shorter than the radio NTFs. This centroid offset and smaller extent of X-ray filaments could be both explained by the much shorter synchrotron cooling lifetime in X-ray than in radio (e.g., Ginzburg \\& Syrovatskii 1965)."
    },
    "0710/0710.1946_arXiv.txt": {
        "abstract": "We report on the first wide-field, very long baseline interferometry (VLBI) survey at 90 cm.  The survey area consists of two overlapping 28 deg$^{2}$ fields centred on the quasar J0226$+$3421 and the gravitational lens B0218$+$357. A total of 618 sources were targeted in these fields, based on identifications from Westerbork Northern Sky Survey (WENSS) data. Of these sources, 272 had flux densities that, if unresolved, would fall above the sensitivity limit of the VLBI observations. A total of 27 sources were detected as far as $2\\arcdeg$ from the phase centre. The results of the survey suggest that at least $10\\%$ of moderately faint (S$\\sim100$ mJy) sources found at 90 cm contain compact components smaller than $\\sim0.1$ to $0.3$ arcsec and stronger than $10\\%$ of their total flux densities. A $\\sim90$ mJy source was detected in the VLBI data that was not seen in the WENSS and NRAO VLA Sky Survey (NVSS) data and may be a transient or highly variable source that has been serendipitously detected. This survey is the first systematic (and non-biased), deep, high-resolution survey of the low-frequency radio sky. It is also the widest field of view VLBI survey with a single pointing to date, exceeding the total survey area of previous higher frequency surveys by two orders of magnitude. These initial results suggest that new low frequency telescopes, such as LOFAR, should detect many compact radio sources and that plans to extend these arrays to baselines of several thousand kilometres are warranted. ",
        "introduction": "\\label{sec:introduction}  The general properties of the 90 cm sky are not very well known and even less is known at VLBI resolution. Previous snapshot surveys at these wavelengths have only targeted the brightest sources and were plagued by poor sensitivity, radio interference and limited coherence times. Furthermore, the field of view that could be imaged was typically limited by the poor spectral and temporal resolution of early generation hardware correlators and the available data storage and computing performance at the time. As a result, although several hundred 90 cm VLBI observations have been made over the past two decades, images of only a few tens of sources have been published e.g. \\citet{alt95}; \\citet{laz98}; \\citet{chu99}; \\citet{cai02}. With such a small sample it is difficult to quantify the total population and nature of these sources. In particular, the sub-arcsecond and sub-Jansky population of 90 cm sources is largely unexplored.  Recent improvements to the EVN hardware correlator at JIVE \\citep{van04}, have enabled significantly finer temporal and spectral resolution. Combined with vast improvements in storage and computing facilities, it is now possible to image fields as wide as, or even wider than, the FWHM of the primary beam of the observing instrument. To complement the hardware improvements, new approaches to calibration and imaging have been developed to better utilise the available data and processing platforms. For example, \\cite{gar05} performed a deep VLBI survey at 20 cm of a $36\\arcmin$ wide field by using a central bright source as an in-beam calibrator. The approach was ideal for survey work as it permitted the imaging of many potential target sources simultaneously by taking advantage of the full sensitivity of the observation across the entire field of view. We have applied a similar technique at 90 cm by piggybacking on an existing VLBI observation of the gravitational lens B0218$+$357 and the nearby quasar J0226$+$3421, with the aim of surveying a 28 deg$^{2}$ field around each of the sources. The results provide an important indication of what may be seen by future low-frequency instruments such as the Low Frequency Array (LOFAR), European LOFAR (E-LOFAR) and the Square Kilometre Array (SKA).  In this paper, we present the results of a 90 cm wide-field VLBI survey that covers two partially overlapping regions of 28 deg$^{2}$ each, surveying 618 radio source targets at angular resolutions ranging between 30 and 300 mas. For sources located at a redshift of $z=1$, the linear resolution corresponding to 30 mas is 230 pc. A \\emph{WMAP} cosmology \\citep{spe03} with a flat Universe, $H_{0}=72$ km s$^{-1}$ Mpc$^{-1}$ and $\\Omega_{m}=0.29$ is assumed throughout this paper. ",
        "conclusions": "Our survey results indicate that at least $10\\%$ of moderately faint (S$\\sim100$ mJy) sources found at 90 cm contain compact components smaller than $\\sim0.1$ to $0.3$ arcsec and stronger than $10\\%$ of their total flux densities. This is a strict lower limit as the sensitivity of our observation was limited by the primary beam at the edge of the survey fields. None of the surveyed sources that were even slightly resolved by WENSS were detected. Similarly, none of the WENSS sources that were below the sensitivity limits of the VLBI observation were detected either, suggesting that none of these sources had significantly increased in brightness since the WENSS observations were carried out.  The apparent lack of sources varying above our detection threshold must at least in part be due to resolution effects.  As 90\\% of the WENSS sources above the VLBI detection threshold are not detected they must be at least partially resolved at the VLBI resolution and the compact component of the radio emission would only be a fraction of the WENSS flux density.  For the compact component of these sources to vary enough to be detected with VLBI, they must increase in strength by factors of perhaps at least a few (the reciprocal of the ratio of compact flux to WENSS flux) to be detectable with VLBI; the compact component of the flux needs to increase above the VLBI sensitivity limit.  Resolution effects are masking variability in these sources.  As discussed below, the detection of one apparent highly variable source in the VLBI data is rather remarkable.  The interpretation of our detection statistics is complicated, in that the survey has a non-uniform sensitivity over both fields, due to the primary beam response of the VLBA antennas and the fact that we are imaging objects well beyond the half-power points of the primary beam.  In addition, due to time and bandwidth smearing effects, as one images objects further from the phase centre, data on the long baselines is discarded, since the smearing effects make imaging difficult.  A consequence of this is that the angular resolution is also non-uniform across the surveyed fields, with low resolution far from the phase centre.  Not only is the flux limit variable across the field, the brightness temperature sensitivity also varies.  It is possible to estimate the detection statistics of our survey for a uniform flux density and brightness temperature limit by considering sources not too far from the phase centre and for a flux sensitivity between the extremes at the phase centre and field edge.  For example, if a sensitivity limit of 30 mJy beam$^{-1}$ is considered (achieved in the $0.25 - 0.5$ degree annulus of field 1 and in the $0.5 - 1$ degree annulus of field 2, and exceeded in the lower radius annuli in each field), 11 out of 55 possible sources are detected, a detection rate of 20\\%, higher than the strict lower limit of 10\\% estimated above for all sources at all annuli.  \\citet{gar05} performed a similar survey of the NOAO Bootes field at 1.4 GHz, using the NRAO VLBA and 100 m Green Bank Telescope. The survey covered a total of 0.28 deg$^{2}$, one hundredth of the area covered by our survey, and detected a total of 9 sources. The survey achieved sensitivities of $0.074-1.2$ mJy beam$^{-1}$ that enabled the detection of both weak and extended sources, whereas our 90 cm observations detected mainly compact sources or slightly resolved bright sources. Nonetheless, we can estimate the number of detections in this region that could be achieved using the 90 cm survey techniques described in this paper. The 0.28 deg$^2$ NOAO Bootes field contains a total of 13 WENSS sources, 6 of which have integrated flux densities $>30$ mJy. Based on our detection rate of 20\\% for such sources we would expect to detect one WENSS source at 90 cm. Assuming a median spectral index of -0.77, only two of the \\citet{gar05} sources have integrated flux densities above our 30 mJy beam$^{-1}$ limit at 90 cm, however, one of these is extended and would have a VLBI peak flux density that falls below our limit. Thus the observations of \\citet{gar05} are consistent with our 90 cm VLBI results for sources with a peak flux density above 30 mJy beam$^{-1}$.  Estimates of the percentage of sources detected with VLBI gives an estimate of the relative contribution of AGN (that contain compact radio emission and are detectable with VLBI) and starburst galaxies (which contain low brightness temperature radio emission not detectable with VLBI).  Analysis of the ratio of starburst galaxies to AGN as a function of redshift (at high redshifts) can help to determine the initial sources of ionising radiation early in the Universe.  As very little redshift data for our surveyed sources are available, such an analysis is not currently possible with this dataset.  In practice, VLBI data at an additional frequency is also required, to confirm that the compact radio emission attributed to AGN has plausible spectral indices.  The distribution of morphologies in the detected survey sources are typical of AGN.  10/27 sources are unresolved point sources, consistent with core-dominated AGN.  A further 8/27 are clearly resolved into double component sources, consistent with being core-jet AGN or double-lobed radio galaxies.  7/27 sources have complex or extended structures, not obviously clear double components.  Again, these sources may be core-jet AGN or radio galaxies.  The remaining 2/27 sources are the gravitational lens and the quasar at the phase centres of the two fields.  The serendipitous detection of a likely highly variable, very compact source near the target WENSS source B0223.8$+$3533 is intriguing. The total area imaged by this survey represents $\\sim0.5\\%$ of the area within the $0\\arcdeg-2\\arcdeg$ annulus and is equivalent to $\\sim2.2\\%$ of the FWHM of the VLBA primary beam. While it is difficult to place any limits on the real population of variable sources based on this one observation, it does highlight the importance of imaging wide-fields completely, in order to improve our understanding of such sources.  \\subsection{Future Prospects}  The observations presented here demonstrate that extremely wide-field surveys can now be piggybacked on current and future VLBI observations at 90 cm. While this survey has mainly concentrated on detecting and imaging sources already detected by other surveys, we find tantalising evidence of a transient or highly variable source. We were fortunate to have found one that appeared in close proximity to one of our target sources but this may not always be the case. This provides a motive to take on a more ambitious survey of the entire field. Such a survey is not beyond the reach of current technology, it would require at most $\\sim45$ times more processing compared to the project presented here, in order to image the entire primary beam of the VLBA using a similar faceted approach. While this is not the most efficient means of detecting transients, it will help progress the development of algorithms and techniques needed for next generation, survey-class instruments that operate at wavelengths or sensitivities not matched by current instruments.  The observations presented in this paper were limited by the spectral and temporal resolution of the EVN correlator at the time of the observation. To minimise the effects of bandwidth and time-averaging smearing it was necessary to compromise resolution and image noise. Future technical developments in the capabilities of correlators will allow wide-field, global VLBI studies to be conducted without such restrictions. In particular, software correlators can provide extremely high temporal and spectral resolution, limited only by the time it takes to process the data \\citep{del07}. They also allow for some pre-processing to be applied during the correlation process to, for example, mitigate the effects of radio interference or to correlate against multiple phase centres simultaneously.  \\subsection{Implications for LOFAR and SKA}  The results of these observations provide important information on the nature and incidence of compact, low-frequency radio sources, with consequences for next generation, low-frequency instruments such as LOFAR and the SKA. LOFAR is currently being deployed across The Netherlands but remote stations are already under construction in neighbouring countries, in particular Germany. Other countries (e.g. UK, France, Sweden, Italy and Poland) are also expected to join this European expansion of LOFAR (E-LOFAR), extending the longest baseline from a few hundred, to a few thousand km. This development will provide LOFAR with sub-arcsecond resolution at its highest observing frequency (the $120-240$ MHz high-band). One concern associated with extending LOFAR to much longer baselines is whether enough cosmic sources will remain unresolved - this characteristic is required in order to ensure there are enough calibrator sources in the sky in order to calibrate the instrument across its full, very wide, field-of-view.  The observations presented here suggest that at least one tenth of all radio sources (at the several tens of mJy level) are likely to exhibit compact VLBI radio structure in the LOFAR high-band. In all likelyhood, an even larger fraction of the E-LOFAR source population will therefore be bright and compact enough to form a grid of calibrator sources across the sky. From our results, we estimate the spatial number density of relatively bright (S$>10$ mJy) and compact (LAS$<200$ mas) sources at 240 MHz to be $\\sim3$ deg$^{-2}$. The aggregate total of these compact sources within a beam should serve as a good calibrator for E-LOFAR and enable most of the low-frequency radio sky to be imaged with excellent sub-arcsecond resolution and high dynamic range. Extrapolation to LOFAR's low-band (10$-$80 MHz) is probably very dangerous, but there is every reason to believe that a large number of these sources will remain compact.  In order to assess the relative numbers of starburst galaxies and AGN as a function of redshift, obviously large redshift surveys need to take place for these radio continuum objects.  Such a survey could be conducted using the redshifted HI signal from these galaxies, using the SKA."
    },
    "0710/0710.1207_arXiv.txt": {
        "abstract": "The total mass of a distant star cluster is often derived from the virial theorem, using line-of-sight velocity dispersion measurements and half-light radii, under the implicit assumption that all stars are single (although it is {\\em known} that most stars form part of binary systems). The components of binary stars exhibit orbital motion, which increases the measured velocity dispersion, resulting in a dynamical mass overestimation. In these proceedings we quantify the effect of neglecting the binary population on the derivation of the dynamical mass of a star cluster. We find that the presence of binaries plays an important role for clusters with total mass $M_{\\rm cl} \\leqslant 10^5~{\\rm M}_\\odot$; the dynamical mass can be significantly overestimated (by a factor of two or more). For the more massive clusters, with $M_{\\rm cl} \\geqslant 10^5~{\\rm M}_\\odot$, binaries do not affect the dynamical mass estimation significantly, provided that the cluster is significantly compact (half-mass radius $\\leqslant 5$~pc). ",
        "introduction": "Young star clusters, with typical masses of $M_{\\rm cl}=10^{3-6}~{\\rm M}_\\odot$, indicate recent or ongoing violent star formation, and are often triggered by mergers and close encounters between galaxies. Only a fraction of these young massive star clusters evolve into old globular clusters, while the majority ($60-90\\%$) will dissolve into the field star population within about 30\\,Myr (e.g., de Grijs \\& Parmentier 2007). In order to understand the formation and fate of these clusters, it is important to study these in detail, and obtain good estimates of the mass, stellar content, dynamics, and binary population. The dynamical mass for a cluster in virial equilibrium, consisting of single, equal-mass stars, is given by: \\begin{equation} M_{\\rm dyn} = \\eta \\, \\frac{ R_{\\rm hm} \\sigma_{\\rm los}^2 }{G} \\end{equation} (Spitzer 1987), where $R_{\\rm hm}$ is the (projected) half-mass radius, $\\sigma_{\\rm los}$ the measured line-of-sight velocity dispersion, and $\\eta \\approx 9.75$. For unresolved clusters, $\\sigma_{\\rm los}$ is usually derived from spectral-line analysis, neglecting the presence of binaries. However, observations have shown that the majority of stars form in binary or multiple systems (e.g., Duquennoy \\& Mayor 1991; Kouwenhoven et al. 2005, 2007; Kobulnicky et al. 2007). When binaries are present, $\\sigma_{\\rm los}$ does not only include the motion of the binaries (i.e., their centre-of-mass) in the cluster potential, but additionally the velocity component of the orbital motion. This results in an overestimation of the velocity dispersion, and hence of $M_{\\rm dyn}$. ",
        "conclusions": ""
    },
    "0710/0710.3202_arXiv.txt": {
        "abstract": "We present spectroscopic observations of the quiescent black hole binary A0620-00 with the the 6.5-m Magellan Clay telescope at Las Campanas Observatory. We measure absorption-line radial velocities of the secondary and make the most precise determination to date ($K_{2}=435.4\\pm0.5$ km s$^{-1}$). By fitting the rotational broadening of the secondary, we refine the mass ratio to $q=0.060\\pm0.004$; these results, combined with the orbital period, imply a minimum mass for the compact object of $3.10\\pm0.04$ M$_{\\sun}.$ Although quiescence implies little accretion activity, we find that the disc contributes $56\\pm7$ per cent of the light in B and V, and is subject to significant flickering. Doppler maps of the Balmer lines reveal bright emission from the gas stream-disc impact point and unusual crescent-shaped features. We also find that the disc centre of symmetry does not coincide with the predicted black hole velocity. By comparison with SPH simulations, we identify this source with an eccentric disc. With high S/N, we pursue modulation tomography of H$\\alpha$ and find that the aforementioned bright regions are strongly modulated at the orbital period. We interpret this modulation in the context of disc precession, and discuss cases for the accretion disc evolution. ",
        "introduction": "A0620-00 (V616 Mon) is the prototype Soft X-ray Transient, a class of low-mass binary stars which exhibit infrequent but intense X-ray bursts (Gelino, Harrison, and Orosz, 2001). In 1975 it became the brightest X-ray nova ever detected, at approximately 50 Crab \\citep{Elvis75}, and it was the first nova to be identified with a black hole primary (McClintock \\& Remillard, 1986; hereafter MR86). MR86 measured an orbital period of 7.75 hr and a radial velocity semiamplitude for the K-type secondary of 457 km s$^{-1}$, leading to a mass function $f(M)=3.18$ M$_{\\sun};$ estimates of $K_{2}$ and $f(M)$ have decreased slightly since then (i.e. 433 km s$^{-1}$ and 3.09 M$_{\\sun}$) (Marsh, Robinson, \\& Wood 1994, hereafter MRW94). Given this minimum mass, it is likely that A0620-00 is a black hole.  A substantial amount of work has gone into the analysis of A0620-00 in the last twenty years, with particular emphasis on ellipsoidal variations in the light curve and the contamination of the K-star flux by light from the accretion disc. As yet, no real consensus has been reached, mostly due to the complexity of the lightcurves. While ellipsoidal variations are obvious, they are highly asymmetric (Leibowitz, Hemar, \\& Orio 1998); the origin of the asymmetry is undetermined.  Modelling this lightcurve, \\citet{Gelino01} determined in inclination of 41$\\pm 3\\degr$, invoking starspots to explain the asymmetries. Shahbaz, Naylor, and Charles (1994) found a 90 per cent confidence interval of $i=$30--45$\\degr$ given the mass ratio of A0620-00, modelling their asymmetries with the bright spot where the accretion stream hits the disc.  Lightcurve modelling is also complicated by the variability of the disc itself. To quantify ellipsoidal variations, most authors assume the disc to be constant, and justify the claim by noting that A0620-00 is quiescent. They do not mention that estimates of the disc contamination range from $<$3 per cent \\citep{Gelino01} to $\\la50$ per cent (MR86). The contribution from this disc is not only unclear, but apparently not constant. More than half of A0620-00's 58-year burst cycle has passed, and it is important to note that quiescent does not mean inactive. We will argue that the variability of the accretion disc cannot be neglected. In order to determine definitively the mass of the compact object, it is very important to understand the structure and variation of the accretion disc. MRW94 made enormous progress towards this goal. In 2004, Shahbaz et al. (hereafter S04) noticed signatures of an eccentric disc not seen in previous Doppler maps, but lacked the phase coverage to verify their hypothesis.  Therefore, as follow-up to the work of MRW94 and S04, and as part of a Doppler imaging survey of black hole and neutron star binaries, we undertook phase-resolved optical spectroscopy of A0620-00. In $\\S 2$ we describe our observational methods and data reduction. In $\\S 3$ we measure the radial velocity of the secondary star, the system mass ratio, and attempt to quantify flickering. In $\\S 4$ we present Doppler images of the accretion disc at several wavelengths, investigate evidence for an eccentric disc, and report results of modulation tomography of the H$\\alpha$ line. We discuss conclusions from the variability of the disc and our Doppler maps in $\\S 5.$ ",
        "conclusions": "We have presented spectroscopic analysis of the black hole binary A0620-00. We measure an absorption-line radial velocity $K_{2}=435.4\\pm0.5$ km s$^{-1}$. With two measurements of the rotational broadening of the secondary, we find a mass ratio of $q=0.060\\pm0.004$ and a minimum mass of $3.10\\pm0.04$ M$_{\\sun}$ for the primary object. With the most likely inclination of 41$\\degr$ from \\citet{Gelino01}, measured in $J,~H,$ and $K$, the black hole has a mass of 11.1 M$_{\\sun}.$ The strong infrared flickering discussed earlier, in conjunction with unexplained smooth variability in the lightcurve and uncertainty in the disc spectrum itself, makes it difficult to estimate the true uncertainty in the inclination. The range of reported inclinations, 31$\\degr$ to 70.5$\\degr$(Gelino et al. 2001 and references therein), results in a black hole mass between 3.7 M$_{\\sun}$ and 22.7 M$_{\\sun}$. Until the nature and variability of the light from the disc is revealed in full, this conservative error estimate must be sufficient.  We also find that the secondary contributes $44\\pm7$ per cent of the light near 5500 \\AA. As this means that the disc contributes a significant fraction of the light, especially in emission line regions, it becomes important to assess the variability of the disc, particularly if the inclination is to be determined by lightcurve modelling. As noted, it is common to assume that the disc is a constant source of light. While it may be valid when $f$ is large, three observational points cast doubt on this assumption: \\begin{enumerate} \\item S04 performed a detailed study of flares from A0620-00, which in their observations have amplitudes nearing 20 per cent of the source flux. \\item Measurements of the fraction $f$ of light contributed by the secondary have not been consistent. MR86 found $40\\pm10$ per cent at 5100 \\AA, MRW94 found $94\\pm3$ per cent at H$\\alpha$, and \\citet{Gelino01} found $f\\ga97$ per cent in $J,~H,$ and $K$, assuming a \\textit{constant} diluting source of light. \\item Observations in all four Spitzer bands, taken approximately two weeks before our observations on the Clay telescope, show strong flickering which is highly correlated with simultaneous R-band lightcurves from the 1.2m telescope on Mt. Hopkins. \\end{enumerate} While McClintock, Horne, and Remillard (1995) rightly point out that the absorption line strength of the template will affect the observed dilution fraction, unless it can be shown that the measured fraction is strongly correlated with template line strength, a physical origin for this variation cannot be ruled out. Future studies could assess the true dependence of $f$ on template star, as well as the long-term variability of $f$, by observing both BS 753 (MWR94's template) and HD 7142. Our measurements of the H$\\alpha$ equivalent widths, when compared to those of MRW94, suggest a physically real origin for the variation, because equivalent widths are independent of the template star and the instrument. So there appears to be evidence for a disc which is brighter relative to the secondary than it used to be.  Indeed, the modulations of the equivalent width even point to a physically non-uniform flickering, because the disc emission lines vary relative to the non-stellar continuum. We can tentatively identify the flickering with the crescents, neither of which was present in 1994, so our conclusion seems viable. The line emission and the continuum may have different radial emissivity dependencies, which could result in slower modulations. We also detect several bright regions in the disc: one near the gas stream impact point, and two crescent-shaped regions on opposite sides of the disc. The reality of these features, as well as the non-uniform flickering, are confirmed by modulation tomography of the H$\\alpha$ disc line, which reveals variation near the bright crescents.  First noticed by S04, the crescent-like features may indicate an eccentric disc, which is predicted for systems like A0620-00 with small mass ratios and large discs. It seems that we have observed an eccentric disc, but let us consider the evidence. From our observations, the following are clear: \\begin{enumerate} \\item The disc is bright and variable. The brightness is evident in the increased dilution of the secondary spectrum, and the variability is clear from a number of phenomena. First, the trailed H$\\alpha$ line shows clear evidence of flickering events. Second, the subtracted equivalent width of the same line is variable beyond explanation by noise alone. Third, the phase-resolved light fraction cannot be reproduced by ellipsoidal variability on top of a constant source of light.  \\item The disc extends to the 3:1 tidal resonance. This is a simple point, clear from the Doppler and modulation maps.  \\item The disc is not centered on the radial velocity of the black hole.  \\item The disc is not radially symmetric, but characterized by bright crescent-shaped regions.  \\item The crescent regions are modulated at the orbital period. \\end{enumerate}  Each of these points alone would be insufficient evidence to conclude that the accretion disc is eccentric and precessing. But with the exception of a direct image of the elliptical disc, we can present a complete and coherent argument that this is the case. The disc has grown to tidal resonance, where the enhanced disc viscosity results in bright and variable rims of extra dissipation. We then observe crescents of extra dissipation at relatively low velocities, as expected. Given the viscous effects, it is predicted that the disc will receive a gravitational torque from the secondary, and begin to precess. The asymmetries introduced here shift the velocity center of disc emission away from the black hole, and we find that the disc is not centered on the black hole in velocity space. Furthermore, given the beat period between the precession and orbital motion, the regions of viscous dissipation should be modulated at roughly the orbital period. Modulation tomography reveals this to be the case.  In retrospect, knowing that portions of the disc are modulated on the orbital period, we look closer at the fraction of light contributed by the secondary star, and see that it is not well fit by ellipsoidal modulations for a system at the inclination of A0620. But if another component of the system was variable on the orbital period, as we have observed the disc to be, then there is no need for concern. The physical picture, a precessing elliptical disc torqued by the secondary star, predicts and produces all the phenomena we have discussed in our data, which are of high quality.  To put it another way, the eccentric disc hypothesis is nicely self-consistent. It explains why and how the accretion disc has changed, allows the disc to be large enough for the growth of eccentric modes, and predicts the phenomena that we observe. It is unfortunately not possible at this point to make an estimate of the disc eccentricity. \\citet{Smith07} have shed a great deal of light on the evolution of disc eccentricity and energy dissipation with 3D SPH simulations. They find that systems with $q$ between 0.08 and 0.24 develop low-mass eccentric discs withsuperhumps; for $q=0.0526,$ the disc exhibits a short-lived superhump and decaying eccentricity. All mass ratios show enhanced dissipation in the disc from the thermal-tidal instability, even without the eccentric modes.  Since we have not observed a superhump, we cannot place A0620-00 in either category. If it falls in the more extreme group, the disc eccentricity is likely zero (reached after about 300 orbital periods) \\citep{Smith07}. In that case, the steady state is a massive disc. If the steady state is very long-lived, and the disc continues to grow, this could explain the enormous intensity of novae like A0620-00. If it fits among the less extreme mass ratios, the disc eccentricity is around 0.1--0.2, and a superhump should be observable with better photometry and a longer baseline \\citep{Smith07}. A0620-00 may also be at a transition between those cases, and its evolution might be somewhat more erratic, as suggested by the SMARTS data discussed earlier. For example, it may toggle between states of quiescence, superhumps, and variability (like what we have observed here). It might, then, be erroneous to interpret this recent increase in brightness as the build towards outburst.  While we have strong evidence that the accretion disc around the black hole has grown out to the tidal distortion radius, evolved into an eccentric disc, and started to precess, further study is required to verify our conclusion. Data from SMARTS, FLWO, and Spitzer will further quantify flickering, and may reveal a superhump, or some new period consistent with our results, and future programs of tomography will track the evolution of the accretion disc. In anticipation of the impending outburst, and in light of progress in simulations, we suggest that this well-studied system not be disregarded or ignored, for it affords us the opportunity to watch the evolution of an accretion disc from quiescence to outburst, and the chance to test models for disc instabilities in X-ray novae."
    },
    "0710/0710.3805_arXiv.txt": {
        "abstract": "Multi-GeV and TeVs gamma sources are currently observed by their Cherenkov flashes on Telescopes (as Magic, Hess and Veritas), looking vertically up into sky. These detectors while pointing horizontally should reveal also the fluorescence flare tails of nearby down-going airshowers. Such airshowers, born at higher (tens km) altitudes, are growing and extending up to lowest atmospheres (EeVs) or up to higher (few km) quotas (PeVs). These fluorescence signals extend the Cherenkov telescopes to a much higher Cosmic Ray Spectroscopy. Viceversa, as it has been foreseen \\cite{Fargion2005} and only recently observed, the opposite takes place. Fluorescence Telescopes made for UHECR detection (as AUGER ones) may be blazed by inclined Cherenkov lights: less energetic, but frequent (PeVs) CR are expected to be often detected. Nearly dozens blazing Cherenkov at EeV should be already found each year in AUGER, possibly in hybrid mode (FD-SD, Fluorescence and/or Surface Detector). Many more CR events (tens or hundred of thousands) at PeVs energies should be blaze Cherenkov lights each year on the AUGER Fluorescence Telescopes. Their UV filter may partially hide their signals and they cannot, unfortunately, be seen yet in any hybrid mode. At these comparable energy the rarest UHE resonant antineutrino $\\bar{\\nu}_e+e$ interactions in air at $\\frac{{M_{W}}^2}{2m_e}=6.3$ PeV  energy, offer enhanced $W^-$ Neutrino Astronomy showering at air horizon, at $\\sim90^\\circ$, while crossing deep atmosphere column depth or Earth (Ande) boundaries. However, AUGER FD are facing opposite way. An additional decay channel rises also (after resonant neutrino skimming Earth) via their secondary $\\tau$ exit in air, by decay in flight via amplified showering: $\\bar{\\nu}_e+e\\rightarrow W^-\\rightarrow\\bar{\\nu}_{\\tau}+\\tau$. Moreover, expected horizontal UHE GZK neutrinos $\\nu_{\\tau}\\,\\bar{\\nu}_{\\tau}$ at EeVs energy, powered by guaranteed cosmogenic GZK \\cite{Greisen:1966jv, za66}, $\\nu_{\\mu}\\,\\bar{\\nu}_{\\mu}$ flavor conversions (in cosmic distances), are also producing penetrating  UHE EeV lepton taus that could sample, better and deeper than PeVs ones, the Earth skin. Such almost horizontal and up going tau showers, originated by UHE astronomical neutrino, may shower and flash by Fluorescence and/or Cherenkov diffused lights at Auger Sky in a few years (nearly three). Viceversa,  at Hess, MAGIC and VERITAS Horizons, at tens or a hundred kilometer distances, the same up going $\\tau\\,\\bar{\\tau}$ airshowers  might rise via fluorescence. On axis they might blaze (rarely) as a Cherenkov flashes below the horizons, possibly correlated to BL Lac or GRB activity. Also UHE ($1-0.1$ EeV) GZK $\\tau$ showering, can be observed upward once reflected onto clouds. The geomagnetic splitting may tag the energy as well as the inclined shower footprint as seen in a recent peculiar event in  AUGER. Additional stereoscopic detection may define the event origination distance and its consequent primary composition, extending our understanding on UHECR composition. ",
        "introduction": "\\begin{figure} \\centering \\includegraphics[width=7.5cm]{Fargion_Ricap1.eps} \\caption{The different geometry for a Magic-like telescope searching at horizon airshower tails. Nearby PeVs CR can be observed by fluorescence at few km altitude while a more powerful airshower (EeV), might be observed at lower altitudes because of the larger slant depth on far edges. Cherenkov airshowers might blaze in the telescope if in axis. Rarest up-going Tau airshower may escape the Earth at far horizon, leading to up-going flashes. Guaranteed nearby CR (PeVs-EeVs) airshower Cherenkov lights, may also be  observed by reflection from nearby hills or mountains (as in Magic or Veritas), or from the sea \\cite{Fargion07}}\\label{fig1} \\end{figure} \\label{intro} Cosmic Rays is a mature, century old  Science. It still hides its secrets beyond the amazing homogeneity and isotropy in all energy range. Only photons, our best  neutral courier in Astronomy, offered up to now a view of the Universe from lowest (radio) to highest (TeVs) energies. Because of the TeVs-PeVs-EeVs photon self-interaction with relic IR (Infrared), CBBR (Cosmic Black Body Radiations) and Radio backgrounds, UHE (Ultra High Energy) photons are bounded in nearby Universe. Therefore neutral neutrinos may offer a far insight of CR sources because of their weak interaction  for we could also reveal the inner core of their mysterious accelerators.  After a decade, the exciting hopes of a discover by AGASA and by Hires of (an almost undeflected) UHECR traces, a long waited new Particle Astronomy have been frustrated by AUGER results. No UHECR clusterings toward BL Lacs or AGN seems to arise up now. Even if GZK cut-off (the CBBR opacity to nucleons above few $10^{19}$ eV flux \\cite{Greisen:1966jv, za66}) appeared to be finally confirmed by Hires \\cite{Hires06} and AUGER \\cite{Yamamoto2007}: no Galactic or nearby Universe map within GZK cut off volume seems to be correlated with these UHECR events. If homogeneity and isotropy will survive AUGER, Z-Burst model \\cite{Fargion-Mele-Salis99}, linking far cosmic UHE ZeV $\\nu$ sources scattering on  $\\bar{\\nu}$  relic ones, remains the unique natural option. Otherwise, our Near Universe as Virgo and our Super- Galactic group and/or plane, must rise soon. Moreover an unexpected heavy composition of extreme UHECR makes more urgent an independent UHECR spectroscopy. As well as the discover of UHE $\\nu\\,\\bar{\\nu}$ GZK, rare but guaranteed GZK \\cite{Greisen:1966jv, za66} secondary neutrino traces. To this project and to its solution we address in present paper. \\begin{figure}[] \\centering \\includegraphics[width=7.5cm]{Fargion_Ricap2.eps} \\caption{The rare inclined UHECR event seen in axis from above. To this picture derived by an AUGER presentation, we overlapped a drawing of the airshower components, their bending and the consequent geomagnetic $\\vec{B}_\\oplus$ splitting. Note the $\\vec{B}_\\oplus$ vector pointing to North and upward respect the ground. Note also the consequent vertical and lateral charge bundle separations. Each charge-bundle follow its bend trajectory that generate its own Cherenkov beam. The comparability between $\\vec{B}_{\\oplus\\perp}$ and $\\vec{B}_{\\oplus\\parallel}$ field vector modules, is the cause of a similar angular (lateral-vertical) deflection. Their projection on the ground is not symmetric at all. The dashed-ellipse on the right side marks our forecast Cherenkov spot made by $e^+$ split shower component, undetected (out of a Cherenkov  reflection on the ground hardly recorded) by Los Leones FD station which instead detected the main Florescence flare. The top-left side ellipse, marks the probable Cherenkov spot born by negative electrons blazing  on Coihueco telescope.}\\label{fig2a} \\includegraphics[width=7.5cm]{Fargion_Ricap3.eps} \\caption{The same Auger event seen from another angulation \\cite{Auger07}. The inclined UHECR is reaching the ground on the center, where the muon pairs are clustering into a twin overlapped ellipses. The arrival angle is approximated at $80^\\circ$ zenith angle \\cite{Auger07}. Our consequent estimated primary energy is above few EeV (possibly around ten EeV). The  Cherenkov blaze time on Coihueco $\\tau'$, being slightly off axis ($\\sim5^\\circ$), must appear much shorter (one or two order of magnitude) than the Fluorescence duration signal $\\simeq\\tau$ reaching Los Leones, because of both the geometric and the relativistic shrinkage, $\\tau'=\\tau_o\\cdot(1-\\beta\\cos(\\theta))$. The FD, because of lack of angular resolution in Auger Telescope, might be unable to reveal the electron pair splitting. The Cherenkov bending angle extends a few degrees: at Coihueco the airshower beam  is roughly at $5^\\circ$ above horizons, while its shower beginning reaches $7^\\circ-9^\\circ$ angle. Therefore the Cherenkov splitting shape might be marginally detectable by Coihueco Telescope, by a two-three pixels separation along its inclined polarized axis.}\\label{fig2b} \\end{figure}  \\subsection{Fluorescence  flares within Cherenkov Telescope}  We foresee that Cherenkov telescopes while pointing at horizons ($\\geq80^\\circ$ zenith angles) may observe a truncate image (a cylinder like) of downward fluorescence airshower, lightening (See Fig.\\ref{fig1}). These flaring views may appear often at few tens km distance, toward $80^\\circ$ angle for tens $PeVs$ (or a hundred km at $\\simeq85^\\circ$ for EeVs energies), as often as once a night. Nearby hills or reflecting sea may disturb the detection. We foresee that such a discover must occur soon, amplifying MAGIC, VERITAS and HESS high energy CR yields.  \\subsection{Cherenkov blazing photons on Fluorescence Telescopes}  The opposite also take place: Cherenkov photons may hit Fluorescence Telescopes, even if  most Fluorescence  detectors are masked by UV filter. Indeed the blazing Cherenkov lights are collimated into a narrow cone; therefore they are more rarer than Fluorescence isotropic signals. We may estimate that few hundreds of EeV airshowers in a year might hit with Cherenkov the AUGER FD. Probably only a few dozens are well revealed with the FD and by muons on SD, as well as Cherenkov lights, as it has been foreseen \\cite{Fargion2005}. A very peculiar and pedagogical event has been shared on line by AUGER collaboration \\cite{Auger07}. We shall analyze that event in the next sections. ",
        "conclusions": "Cherenkov and Fluorescence Telescope may enlarge their view and role tracing both lights. The  wider CR range will  allow a better spectroscopy at PeV-EeV (knee-ankle) regions and a more detailed anatomy of UHECR composition. MAGIC, HESS and VERITAS may soon trace the Fluorescence lights of downward UHECR airshowers. MAGIC and VERITAS must reveal the Cherenkov reflections also on nearby Mountains. The recent inclined UHECR event in Auger \\cite{Auger07} clearly foreseen in \\cite{Fargion2005} discussed in this paper offer the first footprint for such rich information derived by muons bundles, electron pairs showering and splitting into polarized Cherenkov and Fluorescence traces. Within this novel spectroscopy a hidden Neutrino Astronomy wait to be finally unveiled \\cite{Fargion2007}."
    },
    "0710/0710.0421_arXiv.txt": {
        "abstract": "{In this note we present global string solutions which are a generalization of the usual field theory global vortices when the kinetic term is DBI. Such vortices can result from the spontaneous symmetry breaking in the potential felt by a D$3$-brane. In a previous paper (hep-th/$0706.0485$), the DBI instanton solution was constructed which develops a \"wrinkle\" for stringy heights of the potential. A similar effect is also seen for the DBI vortex solution. The wrinkle develops for stringy heights of the potential. One recovers the usual field theory global string for substringy potentials. As an example of the symmetry breaking, we consider a mobile D$3$-brane on the warped deformed conifold. Symmetry breaking can occur if the structure of the vacuum manifold of the potential for the D$3$-brane changes as it moves through the throat region. }    \\begin{document} ",
        "introduction": "String theory embeddings of inflationary scenarios typically involve D-branes which are either static or undergoing some motion in the transverse directions. The world volume theory of the D-brane is described by the DBI action. Since the DBI action is a higher dimensional generalization of the Lorentz invariant relativistic action for a point particle, relativistic effects are to be expected for the D-brane dynamics under appropriate relativistic conditions. Such examples of DBI relativistic effect are  studied in \\cite{Brown:2007ce,Brown:2007vh,Silverstein:2003hf, Chen:2004gc}. In \\cite{Silverstein:2003hf}, the speed limit arising due to the relativistic nature of the DBI action is responsible for inflation. In \\cite{Chen:2004gc} a variant of this scenario was proposed, again based on the DBI action. Further, the DBI effect has been shown to prevent the occurence of slow roll eternal inflation in DBI inflation scenarios \\cite{Chen:2006hs}. These inflationary scenarios employ mobile branes and, therefore, when the D$3$-brane speed becomes relativistic see DBI effects. However, the nonlinear nature of the DBI action can lead to new behavior in certain phenomena even when the D-brane is static. In \\cite{Brown:2007ce} the tunneling of a D$3$-brane from a metastable to a true vacuum was investigated along the lines of the usual quantum field theory analysis given in \\cite{Coleman:1977py}. For substringy barriers the tunneling picture for the DBI action matches with the usual QFT picture given in \\cite{Coleman:1977py}. But once the barrier between the metastable and the true vacuum assumes stringy heights, tunneling is much more enhanced than what one would have expected based on the usual QFT intuition. This is surprising as the usual QFT intuition has taught us to expect an exponential suppression in the tunneling probability with an increase in the height of the barrier. This counterintuitive result is easily understood, however, based on a similar treatement of a point charged particle tunneling through an electrostatic barrier \\cite{Brown:2007vh}. Once the height of the barrier is large enough, there is a significant Schwinger pair production of point particle- antiparticle and this Schwinger effect can enhance tunneling. Similar effect for the D$3$-brane enhances the tunneling rate. The D-brane, therefore, need not be moving at relativistic speeds to see interesting DBI effects. In this paper we investigate global vortex solutions in the world volume of a D$3$-brane described by the DBI action. Similar to the findings in \\cite{Brown:2007ce}, these vortex solutions display a departure from the usual field theory behavior once the height of the potential of the complex scalar field becomes stringy.  The DBI vortex solutions that we will construct should not be confused with the vortex solutions of the tachyon field formed due to tachyon condensation after brane-antibrane annihilation (constructed, for example, in \\cite{Jones:2002si} ) which correspond to codimension two branes. The DBI vortices we construct are generalizations of the well known field theory global vortices when the kinetic term is DBI. Indeed, for small gradients (which will correspond to small heights of the potential) the DBI vortex resembles the usual field theory vortex. Further, these vortices can form even when the separation between a brane and an antibrane is much greater than the critical distance when tachyon condensation initiates. For the tachyon vortices to form, the brane-antibrane separation should be string length. ",
        "conclusions": "We have considered the DBI action for a complex field $\\psi$ with a potential term $V(\\psi)$ and constructed the vortex solutions which are generalizations of the usual global vortex of the complex scalar field theory. For substringy heights of the potential the vortex resembles the usual field theory global vortex. For stringy potentials the vortex develops a wrinkle analogous to the instanton wrinkle found in \\cite{Brown:2007ce} \\footnote{Vortex solutions for the DBI action have been studied by various authors including \\cite{Callan:1997kz, Gibbons:1997xz, Hashimoto:1997px,Hashimoto:2003pu}. These vortex solutions (BIons) can develop a throat and have double valued $\\phi$.}. We have neglected the world volume gauge fields. Including the world volume gauge fields might lead to the DBI analogues of field theory local strings. These vortices should not be confused with the vortex solutions on a tachyon profile that are produced upon tachyon condensation initiated by the brane-antibrane instability. Unlike tachyon vortices, the DBI vortices can form even when the brane-antibrane separation is large in string length units.  A similar effect was seen in the context of DBI instanton in \\cite{Brown:2007ce} where the instanton interpolating between the metastable and the true vacuum of the potential $V(\\phi)$ was studied in the thin wall limit. When the height of the barrier $V_0 < 1$ one gets the usual QFT instanton. However when $V_0 > 1$ the instanton develops a wrinkle due to the double valuedness (and turning around) of the instanton configuration. For the vortex solution the wrinkle appears at $V_0 \\approx 0.8$. This difference in the critical $V_0$ is due to the friction term present in the vortex equation of motion (Eq.(\\ref{dbi})). The wrinkled instanton in \\cite{Brown:2007ce} was constructed in the thin-wall limit where the friction term is set to zero.  The appearance of the wrinkle can be understood by following the analogy with the DBI instanton of \\cite{Brown:2007ce}. Assuming that the vortex configuration will always have a large enough value of the radial coordinate $r$ (i.e. apriori assuming the formation of a wrinkle), the second and the third terms in Eq.(\\ref{dbics}) (i.e. $ \\frac{\\gamma}{r} \\frac{d\\chi}{dr}$ and $\\gamma \\frac{n^2\\chi}{r^2} $) can be neglected. The equation of motion then resembles the thin wall equation of motion for the DBI instanton in \\cite{Brown:2007ce} and following the reasoning for the instanton, a wrinkle must develop. This is, of course, a heuristic way to see the appearance of a wrinkle in the vortex, which will fail to give the exact value of the radial coordinate $r$ or height of the potential $V_0$ where the wrinkle appears.  The asymptotics of the DBI vortex at radial infinity will remain the same as that for the usual field theory vortex of Sec.(\\ref{qft}). This is because at radial infinity the gradient term vanishes and the DBI kinetic term does not play any role. The DBI effect only occurs near the origin (i.e. near the core of the vortex) where the gradient term is big and sources the DBI term. Since the DBI global vortex asymptotics at radial infinity are the same as the usual field theory global vortex asymptotics at infinity, there will still be a logarithmic divergence in the total energy of the field configuration. However, as explained before, inspite of the double-valuedness of the DBI vortex field configuration, the Hamiltonian density will remain finite everywhere.  In Sec. (\\ref{moduli}) we considered the possibility of the cosmological formation of these vortices in a brane inflation scenario. Brane inflation consists of two steps. First, before the inflation can begin, a D$3$-antibrane migrates down the throat and settles at the tip. In the presence of an appropriate moduli space, vortices can form on the antibrane world volume when it settles at the tip. However, in the next stage a D$3$-brane is attracted to the antibrane and this leads to inflation. The inflationary stage washes away the vortices. Further, the different domains on the antibrane which have the $\\psi$ field at different points of the $S^1$ become exponentially large during inflation. Consequently, the D$3$-brane effectively falls towards a single point on the $S^1$. No DBI vortices form at the end of inflation.  Our approach in this note has been phenomenological. Motivated by the results for the D$3$-brane moduli space on a warped deformed conifold in \\cite{DeWolfe:2007hd}, we consider a potential $V(\\psi)$ along the $S^3$ angular coordinates for the D$3$-brane and examine the global vortex solution. Whether or not a wrinkle will form depends on the height of the potential. If $V(\\psi)$ never attains stringy heights, then there is no possibility of a wrinkle. At this point we simply note that the simplest DBI inflation scenario \\cite{Silverstein:2003hf} considers a DBI action of the form Eq.(\\ref{eom}) (with gravity added). The inflaton potential $V(\\phi)$ is generated by radiative or bulk effects whose value must be large when measured in terms of the warp factor/local string units $f(\\phi)$. In particular, power law inflation occurs at late times (i.e. when $t \\to \\infty$) as the mobile D$3$-brane is nearing the tip. The warp factor and the potential have late time dependence $f(\\phi) \\to t^4$ and $V(\\phi) \\to 1/t^2$ which leads to $f(\\phi)V(\\phi) \\to t^2$ in the late time DBI regime. In such a setting it appears that one can generate large potentials needed for interesting DBI effects. For the appearance of a wrinkle we require a potential whose value is order one in local string units. Another issue while considering the appearance of the wrinkle is the trustability of the DBI description. The DBI description is valid whenever the extrinsic curvature of the solution $\\psi(r)$ is low. In the absence of any warp factor this extrinsic curvature is given by \\cite{Brown:2007ce, Sarangi:2007jb} \\baray K(\\psi) = \\frac{1}{\\sqrt{\\alpha'}} \\frac{\\partial V(\\psi)}{\\partial \\psi}, \\earay i.e. as long as the slope of the potential is small in string units, the DBI action has small higher derivative corrections.   It would be interesting to see if there is some version of brane inflation on a warped deformed conifold which can lead to the formation of these DBI vortices. Further, it would be interesting to find the local string version of these DBI global vortices as local strings can be realistic cosmic string candidates. If such local strings still develop a wrinkle due to the DBI effect, this wrinkle would be the result of the stringy DBI action. Such cosmic strings would, in principle, differ from their field theory cousins and perhaps lead to novel phenomenological consequences.  Further one could construct DBI defects, that are extensions of the DBI vortices we have studied, on the world volume of multiple D$3$-branes. The theory would then be a non-abelian one. One could then look , for example, for the DBI extension of the t' Hooft-Polyakov monopole. It is tempting to speculate that just as what we saw for the DBI vortex solution, there will be no DBI effects present at radial infinity of the monopole solution. This is because the field gradient vanishes at large distances from the core of the monopole. The DBI effect would occur for large field gradients near the core of the monopole and lead to the formation of a wrinkle near the core for potentials with stringy heights."
    },
    "0710/0710.2338_arXiv.txt": {
        "abstract": "We present the first large-scale effort of creating composite spectra of high-redshift type Ia supernovae (SNe~Ia) and comparing them to low-redshift counterparts.  Through the ESSENCE project, we have obtained 107 spectra of 88 high-redshift SNe~Ia with excellent light-curve information.  In addition, we have obtained 397 spectra of low-redshift SNe through a multiple-decade effort at Lick and Keck Observatories, and we have used 45 ultraviolet spectra obtained by \\hstiue.  The low-redshift spectra act as a control sample when comparing to the ESSENCE spectra.  In all instances, the ESSENCE and Lick composite spectra appear very similar.  The addition of galaxy light to the Lick composite spectra allows a nearly perfect match of the overall spectral-energy distribution with the ESSENCE composite spectra, indicating that the high-redshift SNe are more contaminated with host-galaxy light than their low-redshift counterparts.  This is caused by observing objects at all redshifts with similar slit widths, which corresponds to different projected distances.  After correcting for the galaxy-light contamination, subtle differences in the spectra remain.  We have estimated the systematic errors when using current spectral templates for K-corrections to be \\about 0.02 mag.  The variance in the composite spectra give an estimate of the intrinsic variance in low-redshift maximum-light SN spectra of \\about 3\\% in the optical and growing toward the ultraviolet.  The difference between the maximum-light low and high-redshift spectra constrain SN evolution between our samples to be $< 10$\\% in the rest-frame optical. ",
        "introduction": "\\label{s:intro}  Type Ia supernovae (SNe~Ia) are the most precise known distance indicators at cosmological redshifts.  The meticulous measurement of several hundred SNe~Ia at both low and high redshifts has shown that the expansion of the Universe is currently accelerating \\citep{Riess98:lambda, Riess07, Perlmutter99, Astier06, Wood-Vasey07}; see \\citet{Filippenko05} for a recent review. The underlying assumption behind that work is that high-redshift SNe~Ia have the same peak luminosity as low-redshift SNe~Ia (after corrections based on light-curve shape; e.g., \\citealt{Phillips93}).  The luminosity of a given SN and its light-curve shape are determined by initial conditions of the white dwarf progenitor star (e.g., mass at explosion, C/O abundance, and metallicity), and the properties of the explosion (e.g., deflagration/detonation transition, the amount of unburned material, and the density at the ignition point).  The progenitor properties are set by the initial conditions at the formation of the progenitor system, presumably having properties similar to the global galactic properties at that time.  Since low-redshift SN~Ia progenitor systems likely form, on average, in significantly different environments than high-redshift SN~Ia progenitors, one may assume that some amount of evolution is inevitable \\citep[for a discussion of different causes and effects of SN evolution, see][]{Leibundgut01}.  Theoretical studies of SN~Ia evolution have focused on the composition, particularly metallicity, of the progenitor system as the primary potential difference between the two samples.  There have been two major studies with conflicting results.  For their study, \\citet{Hoflich98} changed the progenitor metallicity and modeled the explosion, including a full nuclear-reaction network.  \\citet{Lentz00} changed the metallicity of the results of W7 models \\citep{Nomoto84} and input those parameters into their PHOENIX code \\citep{Hauschildt96} to produce synthetic spectra.  The main difference between these methods is the definition of ``metallicity.'' \\citet{Hoflich98} uses the term to mean the metallicity of the progenitor star, while \\citet{Lentz00} uses it to mean the metallicity of the ejecta.  The differing definitions of metallicity yield different initial conditions, which resulted in contradictory results from these studies.  \\citet{Hoflich98} suggest that with increasing metallicity, the ultraviolet (UV) continuum of the SN increases, while \\citet{Lentz00} suggest that it decreases. Ultimately, the differences are the result of differing density structures \\citep{Lentz00, Dominguez01}.  Although the method of \\citet{Lentz00} seems less physical than that of \\citet{Hoflich98} (simply scaling the metallicity of the ejecta by solar abundances does not take into account, for example, that the Fe-group elements are mainly produced in the SN explosion), they provide model spectra for varying metallicities, which may elucidate differences between low and high-redshift SN spectra.  Further predictions for lower metallicity include faster rise times \\citep{Hoflich98}, faster light-curve decline \\citep{Hoflich98}, lower $^{54}$Fe production \\citep{Hoflich98}, smaller blueshifting of \\ion{Si}{2} $\\lambda 6355$ \\citep{Lentz00}, decreasing $B-V$ color \\citep{Dominguez01, Podsiadlowski06}, and changing luminosity \\citep{Hoflich98, Dominguez01, Podsiadlowski06, Timmes03}. \\citet{Ropke04} suggest that the C/O ratio of the progenitor does not significantly affect peak luminosity.  Observationally, a lack of evolution has been supported by investigating various SN quantities such as rise time \\citep{Riess99:risetime}, line velocities \\citep{Blondin06, Garavini07}, multi-epoch temporal evolution \\citep{Foley05}, line strengths \\citep{Garavini07}, and line-strength ratios \\citep{Altavilla06}. There have also been studies comparing the spectra of individual high-redshift SNe~Ia to low-redshift SNe~Ia \\citep{Riess98:lambda, Coil00, Hook05, Matheson05, Balland07}, all of which have concluded that there is no clear difference in spectral properties between the two samples.  \\citet{Bronder07} recently presented measurements of line strengths that suggest a difference between low and high-redshift SNe~Ia in one of three features measured.  They find that the difference is highly dependent on the galaxy contamination at high redshift and might be affected by their small low-redshift SN sample.  Consequently, they note that the difference is interesting but not significant.  Despite the consistencies in spectral properties, \\citet{Howell07} note a slight shift in the mean photometric properties of SNe~Ia with redshift.  They explain this evolution as a change in the ratio of progenitors from the ``prompt'' and ``delayed'' channels \\citep{Scannapieco05}, corresponding to young and old progenitor systems at the time of explosion, respectively.  In particular, the light-curve shape parameter ``stretch'' \\citep{Goldhaber01} increases with redshift.  \\hst\\, observations of ESSENCE objects suggested that the sample may have a large proportion of objects with slow-declining (large stretch) light curves, but this is probably the result of a selection bias\\citep{Krisciunas05}.  Since stretch (and other luminosity light-curve parameters) is correlated with spectral properties, one might expect the spectra of high-redshift SNe~Ia, on average, to differ from those of low-redshift SNe~Ia.  Since all galactic environments at redshift $0 < z < 1.5$ are also present in the local Universe, SN~Ia evolution does not necessarily mean that there are not local analogs.  For instance, if the distribution of observables is on average different at high redshift, as long as for each high-redshift SN there is a similar low-redshift counterpart, the peak brightness could, in principle, be correctly translated into an accurate distance.  Within the local sample, there is no indication of a correlation between host-galaxy metallicity and light-curve shape \\citep{Gallagher05}.  In the process of classifying and finding the redshifts for SNe from the ESSENCE (Equation of State: SupErNovae trace Cosmic Expansion) survey \\citep{Miknaitis07, Wood-Vasey07}, we have obtained 107 spectra which have accurate light-curve parameters such as $\\Delta$ (a light-curve width parameter), time of maximum light, and visual extinction \\citep{Matheson05, Foley08a}.  Most spectra in this sample have low signal-to-noise ratios (S/N) compared to spectra of low-redshift SNe.  This makes impractical a detailed analysis of each object individually to test for outliers.  However, by combining the spectra to make composite spectra, we are able to study the mean spectral properties of the samples.  In Section~\\ref{s:sample} we discuss our low and high-redshift SN~Ia spectral samples.  We describe our methods of creating composite spectra in Section~\\ref{s:method}. In Section~\\ref{s:results} we present the composite spectra and compare the two samples, while in Section~\\ref{s:discussion} we discuss the implications of these results.  We present our conclusions in Section~\\ref{s:conclusions}.  Throughout this paper we assume the standard cosmological model with $(h, \\Omega_{m}, \\Omega_{\\Lambda}) = (0.7, 0.3, 0.7)$. ",
        "conclusions": "\\label{s:conclusions}  By combining many low-S/N, high-redshift SN~Ia spectra, we are able to construct the first series of composite SN~Ia spectra based on the parameters of redshift, phase, $\\Delta$, and $A_{V}$.  In addition, we constructed similar composite spectra from a high-quality sample of low-redshift SN~Ia spectra obtained over the last two decades. Comparison of the composite spectra has shown that once we account for galaxy-light contamination, the two samples are remarkably similar. There are several minor deviations between low and high-redshift samples.  These deviations fall into three categories:  related to metallicity, related to $^{56}$Ni production, and unknown.  The UV excess in the ESSENCE premaximum spectrum is indicative of a different metallicity for the low and high-redshift SNe.  Depending on the model, a UV excess is the result of higher or lower metallicity \\citep{Hoflich98, Lentz00}.  The stronger, more blueshifted \\ion{Si}{2} $\\lambda 6355$ line in the ESSENCE premaximum spectrum indicates a higher metallicity \\citep{Lentz00}.  The most significant difference between our samples is the varying strength of the \\ion{Fe}{3} $\\lambda 5129$ line.  The evolution of the \\ion{Fe}{3} line may indicate that high-redshift SNe~Ia have lower temperatures than low-redshift SNe~Ia.  This, in turn, suggests that SNe~Ia should be less luminous at high redshift.  The weak \\ion{Fe}{3} line may indicate lower $^{54}$Fe production, which could be the result of lower metallicity.  Alternatively, lower metallicity may cause less backwarming from UV photons, decreasing the temperature and the \\ion{Fe}{3}/\\ion{Fe}{2} ratio.  It is unclear if this difference is from evolution in all SNe~Ia, evolution just in the low-$\\Delta$ SNe~Ia, changing demographics, or a selection effect. Low-$\\Delta$ SNe~Ia tend to come from the short-delay progenitor channel.  Because of the short delay, the progenitors of these SNe are more biased by galactic environment, and thus galactic evolution, than the progenitors of long-delay channel SNe~Ia.  It is therefore not surprising that the difference in the \\ion{Fe}{3} line is more obvious in the low-$\\Delta$ objects.   It is possible that the different strengths of the \\about3000~\\AA\\ \\ion{Fe}{2} feature between low and high redshift is an artifact of the construction of the composite spectra.   An analysis of the individual spectra of both samples indicates that the samples are not significantly different; however, the small number of UV objects hampers this study.  It is difficult to definitively detect metallicity differences between the Lick and ESSENCE samples.  First, we have observed three differences between the samples which harbinger a difference in metallicity: a UV excess, a stronger and more blueshifted \\ion{Si}{2} $\\lambda 6355$ line, and a weaker \\ion{Fe}{3} $\\lambda 5129$ line. The UV excess is an ambiguous indicator since the models disagree if it indicates lower or higher metallicity.  The \\ion{Si}{2} line in the premaximum spectrum suggests higher metallicity, and the \\ion{Fe}{3} line suggests lower metallicity.  We have also shown that the previously published low-redshift template spectra have multiple drawbacks when comparing to high-redshift composite spectra.  We therefore caution against using these templates for studies of SN~Ia evolution.  Furthermore, deriving K-corrections from any low-redshift template is difficult; however, the systematic errors are likely to be relatively small.  We see that the intrinsic variation of low-redshift SN spectra is \\about3\\% in the optical.  The spectra vary more in the near-UV and UV as suggested by photometry \\citep{Jha06}.  We are able to put the first constraints of SN~Ia evolution to $\\lesssim 10$\\%.  The results of this study are very suggestive, but require further investigation.  In order to improve our understanding of SN~Ia evolution, we propose three future studies related to this work. First, the theoretical models of the effects of metallicity on SN~Ia spectra should be expanded.  With the current ambiguity amongst models, we cannot determine the direction of the trend in metallicity.  Second, we should gather many more high-redshift spectra to disentangle the redshift-$\\Delta$ ambiguity.  Finally, further UV observations of nearby SNe~Ia are desperately needed.  The only current instrument available for the task is the {\\it Swift} \\uvot. However, previous attempts at obtaining SN~Ia UV spectra have been disappointing \\citep{Brown05}.  We suggest an intense campaign spending several hours per spectrum (similar to \\iue) with the \\uvot. In the near future, we may once again have the ability to obtain high-quality UV spectra with \\hst\\, using \\stis\\ or COS. If the upcoming \\hst\\, servicing mission is successful, we strongly suggest a massive campaign to observe local SNe~Ia in the UV.  Since \\jwst\\, does not have the capabilities to observe the UV, this may be our last opportunity for many years."
    },
    "0710/0710.2879_arXiv.txt": {
        "abstract": "% The physical ingredients to describe the epoch of cosmological recombination are amazingly simple and well-understood. This fact allows us to take into account a very large variety of processes, still finding potentially measurable consequences. In this contribution we highlight some of the detailed physics that were recently studied in connection with cosmological hydrogen recombination. The impact of these considerations is two-fold: (i) the associated release of photons during this epoch leads to interesting and {\\it unique deviations} of the Cosmic Microwave Background (CMB) energy spectrum {\\it from a perfect blackbody}, which, in particular at decimeter wavelength, may become observable in the near future. Observing these distortions, in principle would provide an additional way to determine some of the key parameters of the Universe (e.g. the specific entropy, the CMB monopole temperature and the pre-stellar abundance of helium), {\\it not suffering} from limitations set by {\\it cosmic variance}. Also it permits us to confront our detailed understanding of the recombination process with {\\it direct observational evidence}. In this contribution we illustrate how the theoretical {\\it spectral template} for the cosmological recombination spectrum may be utilized for this purpose. (ii) with the advent of high precision CMB data, e.g. as will be available using the {\\sc Planck} Surveyor or {\\sc Cmbpol}, a very accurate theoretical understanding of the {\\it ionization history} of the Universe becomes necessary for the interpretation of the CMB temperature and polarization anisotropies. Here we show that the uncertainty in the ionization history due to several processes that until now are not taken in to account in the standard recombination code {\\sc Recfast} exceed the level of $0.1\\%$ to $0.5\\%$ for each of them. However, it is indeed surprising how {\\it inert} the cosmological recombination history is even at percent-level accuracy. ",
        "introduction": "\\label{RS:sec:Intro}  \\subsection{What is so rich and beautiful about cosmological recombination?} \\label{RS:sec:Intro1} Within the cosmological concordance model the physical environment during the epoch of cosmological recombination (redshifts $500 \\lesssim z\\lesssim 2000$ for hydrogen, $1600 \\lesssim z\\lesssim 3500$ for \\ion{He}{II}$\\rightarrow$\\ion{He}{I} and $5000 \\lesssim z\\lesssim 8000$ for \\ion{He}{III}$\\rightarrow$\\ion{He}{II} recombination) is extremely simple: the Universe is homogeneous and isotropic, globally neutral and is expanding at a rate that can be computed knowing a small set of cosmological parameters. The baryonic matter component is dominated by hydrogen ($\\sim 76\\%$) and helium ($\\sim 24\\%$), with negligibly small traces of other light elements, such as deuterium and lithium, and it is continuously exposed to a bath of isotropic blackbody radiation, which contains roughly $1.6\\times 10^9$ photons per baryon. These initially simple and very unique settings in principle allows us to predict the {\\it ionization history} of the Universe and the {\\it cosmological recombination spectrum} (see Sect.~\\ref{RS:sec:spectrum}) with extremely high accuracy, where the limitations are mainly set by our understanding of the {\\it atomic processes} and associated transition rates. It is this simplicity that offers us the possibility to enter a {\\it rich} field of physical processes and to challenge our understanding of atomic physics, radiative transfer and cosmology, eventually leading to a {\\it beautiful} variety of potentially observable effects.  \\subsection{What is so special about cosmological recombination?} \\label{RS:sec:Intro2} The main reason for this simplicity is the {\\it extremely large specific entropy} and the {\\it slow expansion} of our Universe. Due to the huge number of CMB photons, the free electrons are tightly coupled to the radiation field due to tiny energy exchange during {\\it Compton scattering} off thermal electrons until rather low redshifts, such that during recombination the thermodynamic temperature of electrons is equal to the CMB blackbody temperature with very high precision. In addition, the very fast {\\it Coulomb interaction} and atom-ion collisions allows to maintain full thermodynamic equilibrium among the electrons, ions and neutral atoms down to $z\\sim 150$ \\citep{RS_Zeldovich68}. Also, processes in the baryonic sector cannot severely affect any of the radiation properties, down to redshift where the first stars and galaxies appear, and the atomic rates are largely dominated by radiative processes, including {\\it stimulated recombination}, {\\it induced emission} and absorption of photons. On the other hand, the slow expansion of the Universe allows us to consider the evolution of the atomic species along a sequence of {\\it quasi-stationary} stages, where the populations of the levels are nearly in full equilibrium with the radiation field, but only subsequently and very slowly drop out of equilibrium, finally leading to {\\it recombination} and the {\\it release of additional photons} in uncompensated bound-bound and free-bound transitions.   \\subsection{Historical overview.} \\label{RS:sec:history} It was realized at the end of the 60's \\citep{RS_Zeldovich68, RS_Peebles68}, that during the epoch of cosmological hydrogen recombination (typical redshifts $800\\lesssim z \\lesssim 1600$) any direct recombination of electrons to the ground state of hydrogen is immediately followed by the ionization of a neighboring neutral atom due to re-absorption of the newly released Lyman-continuum photon. In addition, because of the enormous difference in the $2{\\rm p}\\leftrightarrow 1{\\rm s}$ dipole transition rate and the Hubble expansion {rate}, photons emitted close to the center of the Lyman-$\\alpha$ line scatter $\\sim 10^8-10^9$ times before they can finally escape further interaction with the medium and thereby permit a successful settling of electrons in the 1s-level. It is due to these very peculiar circumstances that the $2{\\rm s}\\leftrightarrow 1{\\rm s}$-two-photon decay process, being $\\sim 10^8$ orders of magnitude slower than the Lyman-$\\alpha$ resonance transition, is able to substantially control the dynamics of cosmological hydrogen recombination \\citep{RS_Zeldovich68, RS_Peebles68}, allowing about 57\\% of all hydrogen atoms in the Universe to recombine at redshift $z\\lesssim 1400$ through this channel \\citep{RS_Chluba2006b}.  Shortly afterwards \\citep{RS_Sunyaev1970, RS_Peebles1970}, the importance of the ionization history as one of the key ingredients for the theoretical predictions of the Cosmic Microwave Background (CMB) temperature and polarization anisotropies became clear, and today these tiny directional variations of the CMB temperature ($\\Delta T/T_0\\sim 10^{-5}$) around the mean value $T_0=2.725\\pm 0.001\\,$K \\citep{RS_Fixsen2002} have been observed for the whole sky using the {\\sc Cobe}        % and {\\sc Wmap} % satellites, beyond doubt with great success. The high quality data coming from balloon-borne and ground-based CMB experiments ({\\sc Boomerang, Maxima, Archeops, Cbi, Dasi} and {\\sc Vsa} etc.) today certainly provides one of the mayor pillars for the {\\it cosmological concordance model} \\citep{RS_Bennett2003}, and planned CMB mission like the {\\sc Planck} Surveyor will help to further establish the {\\it era of precision cosmology}.   \\begin{figure} \\centering \\includegraphics[width=0.8\\columnwidth]{./RS.plot.1.eps} \\caption{Ionization history of the Universe and the origin of CMB signals. The observed angular fluctuations in the CMB temperature are created close to the maximum of the Thomson visibility function around $z\\sim 1089$, whereas the direct information carried by the photons in the cosmological hydrogen recombination spectrum is from earlier times.} \\label{RS:fig:plot1} \\end{figure} In September 1966, one of the authors (RS) was explaining during a seminar at the Shternberg Institute in Moscow how according to the Saha formula this recombination should occur. After the talk his friend (UV astronomer) {\\it Vladimir Kurt} asked him: {\\it 'but where are all the redshifted Lyman-$\\alpha$ photons that were released during recombination?'} Indeed this was a great question, which was then addressed in detail by \\citet{RS_Zeldovich68}, leading to an understanding of the role of the 2s-two-photon decay, the delay of recombination as compared to the Saha-solution, the spectral distortions of the CMB due to two-photon continuum and Lyman-$\\alpha$ emission, the frozen remnant of ionized atoms, and the radiation and matter temperature equality until $z\\sim 150$.   All recombined electrons in hydrogen lead to the release of $\\sim 13.6\\,$eV in form of photons, but due to the large specific entropy of the Universe this will only add some fraction of $\\Delta \\rho_\\gamma/\\rho_\\gamma\\sim 10^{-9}-10^{-8}$ to the total energy density of the CMB spectrum, and hence the corresponding distortions are expected to be very small. However, all the photons connected with the Lyman-$\\alpha$ transition and the 2s-two-photon continuum today appear in the Wien part of the CMB spectrum, where the number of photons in the CMB blackbody is dropping exponentially, and, as realized earlier \\citep{RS_Zeldovich68, RS_Peebles68}, these distortions are significant (see Sect.~\\ref{RS:sec:spectrum}).  In 1975, {\\it Victor Dubrovich} pointed out that the transitions among highly excited levels in hydrogen are producing additional photons, which after redshifting are reaching us in the cm-spectral band. This band is actually accessible from the ground. Later these early estimates were significantly refined by several groups (e.g. see \\citet{RS_Kholu2005} and \\citet{RS_Jose2006} for references), with the most recent calculation performed by \\citet{RS_Chluba2006b}, also including the previously neglected free-bound component, and showing in detail that the relative distortions are becoming more significant in the decimeter Rayleigh-Jeans part of the CMB blackbody spectrum (Fig.~\\ref{RS:fig:DI_results}). These kind of precise computations are becoming feasible today, because (i) our knowledge of atomic data (in particular for neutral helium) has significantly improved, and (ii) it is now possible to handle large systems of strongly coupled differential equations using modern computers.  The most interesting aspect of this radiation is that it has a very {\\it peculiar} but {\\it well-defined, quasi-periodic spectral dependence}, where the photons are coming from redshifts $z\\sim 1300-1400$, i.e. {\\it before} the time of the formation of the CMB angular fluctuation close to the maximum of the Thomson visibility function (see Fig.~\\ref{RS:fig:plot1}). Therefore, measuring these distortions of the CMB spectrum would provide a way to confront our understanding of the recombination epoch with {\\it direct experimental evidence}, and in principle may open another independent way to determine some of the key parameters of the Universe, in particular the value of the CMB monopole temperature, $T_0$, the number density of baryons, $\\propto \\Omega_{\\rm b}h^2$, or alternatively the specific entropy, and the primordial helium abundance (e.g.  see \\citet{RS_Chluba2007d} and references therein). ",
        "conclusions": "\\label{RS:sec:conclusion} It took several decades until measurements of the CMB temperature fluctuations became a reality. After {\\sc Cobe} the progress in experimental technology has accelerated by orders of magnitude. Today CMB scientists are even able to measure $E$-mode polarization, and the future will likely allow to access the $B$-mode component of the CMB in addition. Similarly, one may hope that the development of new technologies will render the consequences of the discussed physical processes observable. Therefore, also the photons emerging during the epoch of cosmological (hydrogen) recombination could open another way to refine our understanding of the Universe.  As we illustrated in this contribution, by observing the CMB spectral distortions from the epoch of cosmological recombination we can in principle directly measure cosmological parameters like the value of the CMB monopole temperature, the specific entropy, and the pre-stellar helium abundance, {\\it not suffering} from limitations set by {\\it cosmic variance}. Furthermore, we could directly test our detailed understanding of the recombination process using {\\it observational data}. It is also remarkable that the discussed CMB signal is coming from redshifts $z\\sim 1300-1400$, i.e. before bulk of the CMB angular fluctuations were actually formed. To achieve this task, {\\it no absolute measurement} is necessary, but one only has to look for a modulated signal at the $\\sim\\mu$K level, with typical amplitude of $\\sim 30\\,$nK and $\\Delta \\nu/\\nu\\sim 0.1$, where this signal can be predicted with high accuracy, yielding a {\\it spectral template} for the full cosmological recombination spectrum, also including the contributions from helium."
    },
    "0710/0710.0798_arXiv.txt": {
        "abstract": "We demonstrate that low resolution Ca\\,{\\small II} triplet (CaT) spectroscopic estimates of the overall metallicity ([Fe/H]) of individual Red Giant Branch (RGB) stars in two nearby dwarf spheroidal galaxies (dSphs) agree to $\\pm$0.1-0.2 dex with detailed high resolution spectroscopic determinations for the same stars over the range $-2.5 <$ [Fe/H] $< -0.5$. For this study we used a sample of 129 stars observed in low and high resolution mode with VLT/FLAMES in the Sculptor and Fornax dSphs. We also present the data reduction steps we used in our low resolution analysis and show that the typical accuracy of our velocity and CaT [Fe/H] measurement is $\\sim$2 \\kms and 0.1 dex respectively. We conclude that CaT-[Fe/H] relations calibrated on globular clusters can be applied with confidence to RGB stars in composite stellar populations over the range $-2.5 <$ [Fe/H] $< -0.5$. ",
        "introduction": "An important aspect for a full understanding of galactic evolution is the metallicity distribution function of the stellar population with time.  Carrying out detailed abundance analyses with high resolution (HR) spectroscopy to trace the patterns that allow one to distinguish between the different galactic chemical enrichment processes is time consuming for large samples of individual stars in a galaxy. This is partly due to the observing time required, but also because of the complex data reduction and analysis necessary.  Fortunately, there is an empirically developed, simply calibrated method available which can make an efficient estimate of metallicity ([Fe/H]) for individual Red Giant Branch (RGB) stars using the strength of the Ca\\,{\\small II} triplet (CaT) lines at 8498, 8542, 8662 \\AA. This method was pioneered for use on individual stars by \\citet{arm1991}. It has the advantage that the lines are broad enough that they can be accurately measured with moderate spectral resolution \\citep[e.g,][]{cole2004}.  The CaT method is routinely used to estimate [Fe/H] for nearby resolved stellar systems and also provides an accurate radial velocity estimate.  Both measurements are facilitated by the strength of the CaT lines and by the generally red colours of the target stars. However, the CaT-derived abundances are empirically defined, with a poorly understood physical basis. Therefore it is important to check the results against HR spectroscopic (\\ie direct) measurements of [Fe/H] and other elements.  The ``classical'' CaT calibration is based on the use of globular cluster stars, all which are drawn from a single age and metallicity stellar population.  The CaT equivalent widths are directly compared to HR spectroscopic measurements of [Fe/H] over a range of metallicity, and this comparison is used to define the relation between CaT equivalent width (EW) and [Fe/H] for all observations taken with the same set up.  This approach has been extensively tested for a large sample of globular clusters \\citep[and references therein]{rutledge1997a, rutledge1997b}. However, globular clusters typically exhibit a constant [Ca/Fe] for a large range of [Fe/H]. This leads to uncertainty in the effect of varying [Ca/Fe] ratios such as is seen in the more complex stellar populations found in galaxies. Furthermore, stars in dwarf galaxies invariably cover a significant range of ages as well as metallicities.  This mismatch in the properties of calibrators and targets has led to suggestions that the CaT method may not be a very accurate indicator of [Fe/H] for more complex stellar populations, especially in those cases where [Ca/Fe] varies significantly \\citep[e.g,][]{pont2004}.  In this paper we investigate the validity of the CaT method for complex stellar populations. We compare large samples of [Fe/H] measurements coming from VLT/FLAMES made using both the CaT method and direct HR spectroscopic measurements {\\it for the same stars} in two nearby dwarf spheroidal galaxies (dSphs), Sculptor and Fornax, over a range of [Fe/H] and [Ca/H]. This is the first time such a detailed comparison has been made for stars outside globular clusters.  We also investigate the theoretically predicted behaviour of the CaT method for a range of stellar atmospheric parameters using a grid of model atmosphere spectra from \\citet{munari2005}.  The paper is organised as follows.  In Section~\\ref{sect:datareduction} we describe the data reduction steps we use within the DART (Dwarf galaxy Abundances and Radial velocities Team) collaboration to estimate EW and velocities from observations in the CaT region, as the accuracy with which this can be done clearly has important implications for the reliability of our conclusions for these galaxies.  We also discuss the verification of the overall calibration and accuracy of the velocity and EW measurements by comparison of results from independent CaT observations and by comparing with theoretical expectations based on signal-to-noise, resolution and line profile properties. In Section~\\ref{sec:calibration_gcs} we derive the standard CaT-[Fe/H] globular cluster calibration for low resolution (LR) VLT/FLAMES data. In Section~\\ref{sec:calibration_hr} we compare the derived [Fe/H] from the CaT to the HR [Fe/H] for the Sculptor and Fornax dSphs. Finally, in Section~\\ref{sec:uncert} we discuss the uncertainties that come from using Ca\\,{\\small II} lines to derive an [Fe/H] abundance for stellar populations where the $\\alpha$-abundance varies and use a comparison with stellar model atmospheres to further investigate age, metallicity and $\\alpha$-abundance effects. ",
        "conclusions": "We described the data reduction steps we use within the DART collaboration to estimate velocities and CaT metallicities for LR data for RGB stars in dSphs. We showed that we obtain accurate velocities and [Fe/H] measurements, with internal error in velocity $\\sim$ 2 \\kms and in [Fe/H] $\\sim$ 0.1 dex at a S/N per \\AA\\, of 20.  We used 4 Galactic globular clusters observed with VLT/FLAMES in the CaT region to test the performance of several CaT $W'$-[Fe/H] relations existing in the literature. We also derived the best calibration from these globular cluster data. The relation here derived is consistent at the 1-$\\sigma$ level with the calibration derived in \\citet{tolstoy2001}, which we used in \\citet{tolstoy2004}, \\citet{battaglia2006} and \\citet{helmi2006}.  We used a sample of 93 and 36 RGB stars in the Sculptor and Fornax dSphs, respectively, overlapping between our LR and HR VLT/FLAMES observations, to test the globular clusters CaT calibration. This is the first time that the CaT calibration is tested on field stars in galaxies. We find a good agreement between the metallicities derived with these two methods. However, a systematic trend is present with [Fe/H], such that using the globular cluster calibration derived in this work the [Fe/H] measurement from CaT is overestimated of $\\sim$0.1 dex at [Fe/H]$_{\\rm HR} < -2.2$, whilst at [Fe/H]$_{\\rm HR} > -1.2$ it is underestimated of $\\sim$0.1-0.2 dex. No clear systematic trend is instead derived from our data for [Fe/H]$_{\\rm HR} > -0.8$. In order to understand this systematic effect, we explored the possible contribution of Ca abundance to the calibration, and showed that there are indications that it might well affect the CaT $W'$, although much less than [Fe/H]. From our dataset we also show that, contrary to previous claims, it is not advisable to use the CaT $W'$ as a linear indicator of [Ca/H].  Finally we investigated the effect of varying stellar atmosphere parameters on the CaT method by analysing a large sample of model atmosphere spectra \\citep{munari2005}.   We again demonstrated that the CaT method is (surprisingly) robust to the usual combination of age, metallicities and [$\\alpha$/Fe] variations seen in nearby dSphs.  From our analysis we see that even for large differences in [Ca/Fe] between calibrating globular clusters and our sample of dSphs ($\\sim$ 0.5 dex at [Fe/H]$\\sim -0.7$) and large difference in ages (at [Fe/H] $\\sim -0.7$ Fornax stars are $\\sim$ 10 Gyr younger than globular clusters stars) the error in estimating [Fe/H] using globular clusters as calibrators is just 0.1-0.2 dex. The Fornax dSph is likely to represent the most extreme case as it has had one of the most extended star formation histories among the dSphs in the Local Group.  We conclude that CaT-[Fe/H] relations calibrated on globular clusters can be applied with confidence to RGB stars in composite stellar populations such as galaxies, at least in the [Fe/H] range probed by the above analyses, $-2.5 <$ [Fe/H] $<-0.5$. Hence, the CaT method provides a good indicator of the overall metallicity of resolved stars.  This has implications for the efficiency with which we can obtain metallicity distribution functions of nearby resolved galaxies. The HR data collected in this paper required more than 6 nights of VLT observing time for $\\sim$ 150 spectra, whereas $\\sim 120$ CaT spectra were obtained in one hour VLT observing time."
    },
    "0710/0710.1082_arXiv.txt": {
        "abstract": "We report a new determination of the faint end of the galaxy luminosity function in the nearby clusters Virgo and Abell 2199 using data from SDSS and the Hectospec multifiber spectrograph on the MMT. The luminosity function of A2199 is consistent with a single Schechter function to $M_r=-15.6$ + 5 log $h_{70}$ with a faint-end slope of $\\alpha=-1.13\\pm0.07$ (statistical).  The LF in Virgo extends to $M_r\\approx-13.5\\approx M^*+8$ and has a slope of $\\alpha=-1.28\\pm0.06$ (statistical).  The red sequence of cluster members is prominent in both clusters, and almost no cluster galaxies are redder than this sequence.  A large fraction of photometric red-sequence galaxies lie behind the cluster.  We compare our results to previous estimates and find poor agreement with estimates based on statistical background subtraction but good agreement with estimates based on photometric membership classifications (e.g., colors, morphology, surface brightness).  We conclude that spectroscopic data are critical for estimating the faint end of the luminosity function in clusters.  The faint-end slope we find is consistent with values found for field galaxies, weakening any argument for environmental evolution in the relative abundance of dwarf galaxies.  However, dwarf galaxies in clusters are significantly redder than field galaxies of similar luminosity or mass, indicating that star formation processes in dwarfs do depend on environment. ",
        "introduction": "The luminosity function of galaxies is fundamental to understanding galaxy formation and evolution.  The luminosity function differs dramatically from the expected mass function of dark matter halos, indicating that baryonic physics is very important for understanding galaxies.  In particular, a well-determined luminosity function enables accurate modeling linking the masses of dark matter haloes to galaxy luminosities \\citep[e.g.,][and references therein]{vale06,yang07}.  These empirical models provide a powerful test of any model of galaxy formation and evolution.  Early studies of the luminosity function used the large galaxy density in clusters as a tool for measuring the shape of the luminosity function \\citep[e.g.,][]{sandage85}.  The obvious drawback of this method is that the luminosity function in dense environments may differ from that in more typical galaxy environments \\citep[][]{binggeli88,driver94,depropris95}.  Environmental trends in the luminosity function may reflect differences in galaxy formation in different environments \\citep[][]{tully02,benson03}. For instance, tidal stripping or ``threshing'' of larger galaxies may produce dwarf galaxies \\citep{bekki01}, or dwarf galaxies may be formed in tidal tails of intractions among giant galaxies \\citep{barnes92}.  Alternatively, the denser environments of protoclusters may have shielded low-mass galaxies from the ultraviolet radiation responsible for reionization \\citep{tully02,benson03}.  Many studies suggest an environmental influence on the LF; others provide no such evidence.  The main difficulty in resolving this important issue is the challenge of determining cluter membership for faint galaxies where background galaxy counts are large.  Because few deep spectroscopic surveys of clusters extend into the dwarf galaxy regime \\citep[$M_r\\gtrsim$-18; for exceptions, see][and references therein]{mobasher03,christlein03,mahdavi05}, cluster membership is usually determined via statistical subtraction of background galaxies \\citep[e.g.,][and references therein]{popesso06b,jenkins07,milne07,adami07,yamanoi07,barkhouse07}.  Because galaxy number counts increase much more steeply than cluster member counts (even for very steep faint-end slopes), small systematic uncertainties in background subtraction can produce large uncertainties in the abundance of faint cluster galaxies.   Here, we use MMT/Hectospec spectroscopy and data from the Sloan Digital Sky Survey \\citep[SDSS,][]{sdss} to estimate the luminosity function (LF) in the clusters Abell 2199 and Virgo.  These data enable very deep sampling of the luminosity function.  In particular, we report an estimate of the faint-end slope of the luminosity function with much smaller systematic uncertainties than most previous investigations.  We demonstrate that photometric properties of galaxies such as color and surface brightness correlate well with cluster membership (in agreement with many previous studies).  Very few galaxies redder than the red sequence are cluster members.   We discuss the photometric and spectroscopic data in $\\S 2$.  We present the luminosity functions in $\\S 3$.  We compare our results to previous studies and discuss possible systematic effects and uncertainties in $\\S 4$.  We conclude in $\\S 5$.  An Appendix details the construction of our catalog of confirmed and probable Virgo cluster members.     We assume cosmological parameters of $\\Omega_m$=0.3, $\\Omega_\\Lambda$=0.7, $H_0$=70~$h_{70} \\kms \\mbox{Mpc}^{-1}$.  The spatial scale at the distance of A2199 is 1\\arcs=0.61~$\\kpc$ and 1\\arcs=0.080~$\\kpc$ at the distance of Virgo. ",
        "conclusions": "Determining the faint end of the luminosity function in clusters has remained an unresolved problem for many years.  We report a new estimate of the luminosity function in A2199 from MMT/Hectospec spectroscopy, and of the Virgo cluster from SDSS DR6 data.  Both LFs extends to fainter absolute magnitudes than most previous determinations of the LF in massive clusters.  The LF closely follows a Schechter function to $M_r\\approx-15\\approx M^*+6$ in A2199 and to $M_r\\approx-13\\approx M^*+8$ in Virgo.  There are no large populations of high surface brightness galaxies or galaxies redder than the red sequence that contribute significantly to the LF in either cluster.   In A2199, we find that essentially no galaxies redward of the photometric red sequence are cluster members.  However, the red sequence itself contains a significant number of background galaxies.  We find no evidence of an upturn in the A2199 LF at faint magnitudes as claimed by some recent studies \\citep[][]{popesso06b,jenkins07,milne07,adami07,yamanoi07,barkhouse07}, although the range of absolute magnitudes precludes a conclusive result.  A simple extrapolation of the A2199 LF using the membership fractions at the spectroscopic limit provides an approximate upper limit of the LF at fainter magnitudes; this extrapolation suggests that the faint end slope is probably comparable to that of Virgo.  In the Virgo cluster, we find an LF consistent with a moderate faint-end slope ($\\alpha=-1.28\\pm0.06$).  The Virgo LF extends much fainter than the A2199 LF, and we conclusively demonstrate that the Virgo LF is inconsistent with the steeper LFs found by \\citet{popesso06b}.  The discrepancy between the LFs may be due to systematic uncertainties in statistical background subtraction, and we discuss some possibilities.  Other estimates of the LF using deep spectroscopy find slopes similar to ours \\citep{mobasher03,christlein03,mahdavi05}.  Recent estimates of the LF in Fornax using surface brightness as a membership classification \\citep{hilker03} or surface brightness fluctuations to estimate distances \\citep{mieske07} indicate a shallow LF similar to the Virgo SDSS LF.  Similarly, a clever application of gravitational lensing in A1689 by \\citet{medezinski07} suggests a LF consistent with those we find in Virgo and A2199.  Low surface brightness galaxies remain problematic for determining cluster LFs.  Their low surface brightnesses prohibit reliable redshift estimates using SDSS DR6 spectroscopy.  Perhaps the most significant impact of the SDSS spectra is to demonstrate conclusively that higher surface brightness galaxies are virtually all in the background of Virgo.  A simple division in apparent magnitude versus surface brightness is a surprisingly powerful membership classification (Figure \\ref{virgolsb}, $\\S 2.3$).  Careful inspection of SDSS galaxies failing the spectroscopic target selection criteria reveals many low surface brightness galaxies that are likely Virgo members.  It is somewhat ironic that almost all $r<17.5$ galaxies within 1 Mpc of M87 without well-measured redshifts are the LSB galaxies most likely to be Virgo members. We list these galaxies in the Appendix as an aid for future efforts to obtain spectroscopy of these galaxies.  The spectroscopically determined LFs we find for the A2199 and Virgo clusters are similar to the field LF, an important result for models of galaxy formation.  Future studies of the LF in more clusters and at larger clustrocentric radii will constrain cluster-to-cluster variations in the LF and any radial dependence.  The fraction of galaxies belonging to a cluster decreases dramatically with both increasing magnitude and increasing projected radius.  Even at the relatively small radii we probe here, we show that statistical background subtraction is problematic due to the high precision required.  We therefore recommend that future investigations of the LF in clusters avoid statistical background subtraction and instead identify member galaxies via spectroscopy or photometric information such as colors or surface brightness fluctuations."
    },
    "0710/0710.3614_arXiv.txt": {
        "abstract": "Ultraviolet, optical and near infrared images of the nearby ($D \\approx 5.5$ Mpc) SBm galaxy NGC 1311, obtained with the {\\it Hubble Space Telescope}, reveal a small population of 13 candidate star clusters.  We identify candidate star clusters based on a combination of their luminosity, extent and spectral energy distribution.  The masses of the cluster candidates range from $\\sim$10$^3$ up to $\\sim$10$^5$ $M_{\\odot}$, and show a strong positive trend of larger mass with increasing with cluster age.  Such a trend follows from the fading and dissolution of old, low-mass clusters, and the lack of any young super star-clusters of the sort often formed in strong starbursts.  The cluster age distribution is consistent with a bursting mode of cluster formation, with active episodes of age $\\sim$10 Myr, $\\sim$100 Myr and $\\ga$1 Gyr.  The ranges of age and mass we probe are consistent with those of the star clusters found in quiescent Local Group dwarf galaxies. ",
        "introduction": "Star clusters are powerful tools for probing the star-formation history and chemical enrichment of galaxies (e.g., \\citeauthor{h61}\\citeyear{h61}; \\citeauthor{whit99}\\citeyear{whit99}; \\citeauthor{dk2}\\citeyear{dk2}; \\citeauthor{grijs}\\citeyear{grijs}).  In nearby systems, star clusters provide us with spatially resolved examples of simple stellar populations (SSPs). Observations of individual cluster stars are, therefore, crucial for understanding the effect of metallicity on stellar evolution (e.g., \\citeauthor{carme}\\citeyear{carme} and references therein).  In more distant systems, star clusters provide us with examples of unresolved SSPs.  Analysis of statistical samples of such clusters offers us a wealth of information on the history of star formation in their parent galaxies (e.g., \\citeauthor{grijs}\\citeyear{grijs} and references therein).  Understanding star and cluster formation in a range of environments is essential for understanding galaxy evolution because star cluster formation traces the strongest episodes of galactic star formation (e.g., \\citeauthor{dGNA}\\citeyear{dGNA}).  There has been a great deal of work on star cluster populations and formation histories in nearby luminous spiral galaxies over the last decade (e.g., \\citeauthor{whit99}\\citeyear{whit99}; \\citeauthor{grijs}\\citeyear{grijs}; \\citeauthor{dk2}\\citeyear{dk2}; \\citeauthor{lar04}\\citeyear{lar04}).  By comparison, the star cluster formation history in nearby, low-luminosity galaxies has received relatively little attention (but see, e.g., \\citeauthor{bhe}\\citeyear{bhe}; \\citeauthor{a04}\\citeyear{a04}; \\citeauthor{seth04}\\citeyear{seth04}; \\citeauthor{dGA06}\\citeyear{dGA06}). Studies of nearby galaxies are important because we can resolve abundant small-scale detail in them.  Moreover, the low metallicities of low-luminosity nearby galaxies offer us a close-up view of the star-formation process that may better resemble that in the high-redshift, early Universe.  Studies of detailed stellar formation histories are still restricted to fairly nearby galaxies, and are most powerful when they combine high angular resolution with broad wavelength coverage.  The number of very nearby galaxies with such data available is still quite small.  NGC 1311 is a very nearby ($D \\approx 5.5$ Mpc), but little-studied late-type (SBm) galaxy. \\citeauthor{tully}(\\citeyear{tully}) identify it as a member of the 14$+$14 Association, a loose group dominated by the luminous spiral NGC 1313.  Table 1 summarizes the basic properties of the system.  NGC 1311 was a target in several {\\it Hubble Space Telescope} (HST) snapshot surveys (GO programs 9124 and 9824; \\citeauthor{wind02}\\citeyear{wind02}, \\citeauthor{vio}\\citeyear{vio}, and \\S 2 below).  As a result, a set of broad-band images spanning a wide wavelength interval (0.3---1.6$\\mu$m) at sub-arcsecond resolution now exists for this galaxy.  As NGC 1311 is quite nearby, the bright star clusters and luminous individual stars are detected as discrete sources.  We can thus probe the spatially resolved star-formation history of NGC 1311 by studying the broad-band spectral energy distributions of the star clusters, the individual stars, and the unresolved light.  This paper is concerned with the star clusters of NGC 1311.  We shall address the individual stars and unresolved light in future publications.  In \\S 2 we describe the HST observational data.  We present the observed properties of the candidate star clusters in \\S 3, and our analysis of these observations in \\S 4.  We summarize our conclusions, and discuss issues for further research in \\S 5. ",
        "conclusions": "NGC 1311 is a nearby, but little studied, low-luminosity late-type spiral.  It has optical and \\hi\\ properties (see Table 1) typical for such galaxies (\\citeauthor{m98}\\citeyear{m98}; \\citeauthor{lvz}\\citeyear{lvz}).  It is a member of a loose association of galaxies (\\citeauthor{tully}\\citeyear{tully}), but displays no obvious sign of any recent interaction.  We have used HST WFPC2 and NICMOS images of NGC 1311, that span the near-UV through the near-IR, to identify a small population of 13 candidate star clusters.  Their masses increase systematically with cluster age, as would follow from the fading and dissolution of old, low-mass clusters, and the lack of any young super star clusters associated with strong starbursts.  Half the cluster candidates are significantly fainter than the turnover of the globular cluster luminosity function.  We are thus probing the range of luminosities typical of the faint star clusters found in Local Group dwarf galaxies, and the open cluster population of the Galactic disk.  Analysis of the photometry of the candidate star clusters suggests that NGC 1311 has had three cluster-forming episodes in its history, occuring $\\sim$10 Myr, $\\sim$100 Myr, and $\\ga$1 Gyr ago.  This is consistent with observational work on other nearby low-luminosity galaxies indicating a bursting mode of star-formation.  The recent star formation, as traced by the NUV continuum (see Fig.~1a), is concentrated at the east and west ends of the central bulge-like concentration.  This is reminiscent of stochastic star formation models (\\citeauthor{ssg}\\citeyear{ssg}) as well as the observed properties of other low-luminosity star-forming galaxies (e.g., Sextans A; \\citeauthor{rdp}\\citeyear{rdp}).  NGC 1311 is an excellent example of a nearby, low-luminosity, star-forming, gas-rich galaxy that is evolving in relative isolation.  Understanding the star- and cluster-formation history and chemical evolution of such galaxies is an essential part of unraveling the problem of galaxy evolution.  Our next step is a study of the resolved stellar populations of NGC 1311 with our combined WFPC2/NICMOS HST data.  The large wavelength range of our data allows us to sample both the very recent star formation (dominating the UV light) and the ancient stellar populations (from the red/near-IR light) that appear ubiquitous even in very late-type galaxies (e.g., \\citeauthor{baa}\\citeyear{baa}; \\citeauthor{vio}\\citeyear{vio}).  This should give us a clearer picture of the star formation history in this system, and a fuller understanding of the process of star formation in low-luminosity late-type galaxies in general."
    },
    "0710/0710.4202_arXiv.txt": {
        "abstract": "We have carried out sub-mm $^{12}$CO($J=3-2$) observations of 6 giant molecular clouds (GMCs) in the Large Magellanic Cloud (LMC) with the ASTE 10m sub-mm telescope at a spatial resolution of 5 pc and very high sensitivity. We have identified 32 molecular clumps in the GMCs and revealed significant details of the warm and dense molecular gas with $n$(H$_2$) $\\sim$ 10$^{3-5}$ cm$^{-3}$ and $T_{\\mathrm{kin}}$ $\\sim$ 60 K. These data are combined with $^{12}$CO($J=1-0$) and $^{13}$CO($J=1-0$) results and compared with LVG calculations. The results indicate that clumps we detected are distributed continuously from cool ($\\sim$ 10 -- 30 K) to warm ($\\sim$ higher than 30 -- 200 K), and warm clumps are distributed from less dense ($\\sim$ 10$^3$ cm$^{-3}$) to dense ($\\sim$ 10$^{3.5 - 5}$ cm$^{-3}$).% We found that the ratio of $^{12}$CO($J=3-2$) to $^{12}$CO($J=1-0$) emission is sensitive to and is well correlated with the local \\Halpha flux. We interpret that differences of clump propeties represent an evolutionary sequence of GMCs in terms of density increase leading to star formation.Type I and II GMCs (starless GMCs and GMCs with HII regions only, respectively) are at the young phase of star formation where density does not yet become high enough to show active star formation and Type III GMCs (GMCs with HII regions and young star clusters) represents the later phase where the average density is increased and the GMCs are forming massive stars. The high kinetic temperature correlated with \\Halpha flux suggests that FUV heating is dominant in the molecular gas of the LMC. ",
        "introduction": "It is of a fundamental importance in astronomy to understand the evolution of galaxies. Since a major constituent of galaxies is stars, the formation of stars is a fundamental process in galactic evolution. The properties of stars characterize the basic contents of galaxies and their time evolution. We understand from studies of the Milky Way galaxy that giant molecular clouds (GMCs), whose mass ranges from 10$^5$ to 10$^7$ \\Msun, are the principal sites of star formation and that it perhaps holds the true in other galaxies as well. We also recognize that the GMC properties (e.g., $L_{\\mathrm{CO}}$ -- line width relation, index of mass spectrum) are similar among five galaxies in the Local Group according to the spatially resolved studies \\citep{Blitz2006}. This supports the idea that studies of GMCs will be useful in understanding the fundamentals of galactic evolution through the formation and evolution of GMCs and star formation therein.  Observational studies of GMCs have been most effectively made by the mm interstellar carbon monoxide emission line at 2.6 mm which allows us to probe molecular gas whose density is greater than $\\sim$ 100 cm$^{-3}$. We note that the most abundant species, molecular hydrogen, does not have appropriate line emissions in the mm and sub-mm region due to its zero permanent electric dipole moment and large separation between the lowest energy levels, which are not excited significantly in the typical physical conditions of molecular clouds.  Recent advances in sub-mm observations have allowed us to determine physical parameters of molecular clouds over much larger ranges than in the mm region by comparing line intensities between different transitions. These sub-mm studies were initiated by the SEST 15m telescope in Chile followed by instruments in Hawaii, Mauna Kea and in the Swiss Alps at an altitude range from 3700 m to 4200 m, including the CSO 10m, JCMT 15m, and KOSMA 3m telescopes, and the AST/RO 1.6m telescope in Antarctica. Subsequently, in the 2000's, the developements of new instruments at an altitude of $\\sim$ 5000 m in Atacama in northern Chile resulted in a superior capability because of the high altitude and dry characteristics of the site. The instruments installed in Atacama include the ASTE 10m, APEX 12m and NANTEN2 4m telescopes. All these instruments are beginning to take new molecular data with significantly better quality than before in terms of noise level as well as angular resolution. It is also noteworthy that the current frequency coverage extends as high as the 800GHz band and even the THz region.  Among nearby galaxies we can observe at reasonably high spatial resolutions, the Large and Small Magellanic Clouds offer us a unique opportunity to achieve the highest resolutions due to their unrivaled closeness, 50 -- 60 kpc. In particular, the Large Magellanic Cloud (LMC) is actively forming stars in clusters and is an ideal laboratory for us to study star formation, particularly massive star formation in star cluster. In the LMC, the metallicity is a factor of $\\sim$ 3 lower than in the Solar neighborhood \\citep{Dufour1982, Dufour1984, Rolleston2002}. Also, the visual extinctions are lower and the FUV field is stronger in the LMC than in the Milky Way \\citep{Israel1986}, characterizing the initial conditions of star formation.  The first spatially resolved complete sample of GMCs in a single galaxy has been obtained towards the whole LMC with the NANTEN 4m telescope in 2.6 mm CO emission at 40 pc resolution \\citep{Fukui1999, Fukui2001, Fukui2006b, Mizuno2001}. These studies revealed the three types of GMCs in terms of star formation activities; Type I is starless, Type II is with HII regions only, and Type III is associated with active star formation indicated by huge HII regions and young star clusters, where the stars identified are only O stars due to the sensitivity limitation. It also revealed that the lifetime of a GMC is as short as $\\sim$ 30 Myrs \\citep{Fukui2006a, Kawamura2006, Kawamura2007}. These previous studies naturally place the LMC as one of the prime targets for sub-mm studies to derive the physical parameters of GMCs.  Another aspect which deserves our attention is that very young, rich stellar clusters are forming in the LMC. These are so called populous clusters which are very rare in the Milky Way and resemble globulars formed in the primeval Milky Way. The open clusters forming in the Milky Way are small in the number of stars and loose in spatial distribution. Along with the low metallicity of the LMC, it is an interesting possibility to use molecular data to investigate the formation mechanism of super star clusters at the molecular cloud stage.   In the past, there have been some studies that used the higher transitions ($J=2-1$, $J=3-2$, $J=4-3$, $J=7-6$) of CO spectra of the molecular clouds in the LMC \\citep[e.g.,][]{Sorai2001,Johansson1998,Heikkila1999,Bolatto2005,Israel2003,Kim2004,Kim2006}. These studies suggest the molecular gas may be warmer and/or denser than in the Milky Way.  \\citet{Johansson1998} used the SEST 15m telescope to observe the central part of the 30 Doradus nebula (rms $\\sim$ 0.2 K in 0.5 km s$^{-1}$ velocity resolution for $J=1-0$ and rms $\\sim$ 1.0 K in 0.5 km s$^{-1}$ velocity resolution for $J=3-2$), and the southern HII regions N 158C, N 159, and N 160 with a few prominent CO clouds in the $J=2-1$ and $J=3-2$ transitions of CO. They find that the kinetic temperatures are 10 -- 80 K and the highest temperature is towards 30 Dor. The smallest beam size and grid spacing are 15$\\arcsec$ and 11$\\arcsec$  respectively in the $J=3-2$ emission. \\citet{Heikkila1999} used SEST to observe the $J=3-2$ transition of CO in N 159 and 30 Doradus, as well as other rarer molecular species. This study aimed at obtaining chemical abundance, while it also provides more information on cloud temperature etc, from CO($J=3-2$) data. The kinetic temperatures they derived are 50 K in 30 Dor-10, 15 K in 30 Dor-27, and 20 -- 25 K in N 159W and N 160. \\citet{Bolatto2005} employed the AST/RO to observe the $^{12}$CO($J=4-3$) transition at 461 GHz with a 109$\\arcsec$ beam. They observed 9 regions in the LMC at 6$\\arcmin$$\\times$6$\\arcmin$  field all with HII regions and derived kinetic temperatures from a comparison between the CO($J=4-3$) and ($J=1-0$) transitions. N 48, N 55A, N 79, N 83A, N 113, N 159W, N 167, N 214C, and LIRL 648 are included. They derive temperatures of 100 -- 300 K and note a trend that higher temperatures occur in moderate density regions, 100 -- 1000 cm$^{-3}$, and the lower temperatures in much denser regions 10$^{4-5}$ cm$^{-3}$. These studies were preceded by a suggestion that significant amounts of warm molecular gas may exist in the LMC \\citep{Israel2003}. \\citet{Kim2004} also made similar observations towards an HII region, N 44, and suggest very dense gas of $\\sim$ 10$^5$ cm$^{-3}$. Most recently, \\citet{Kim2006} derived $T_{\\mathrm{kin}}$ = 100 K, $n$ $\\sim$ 10$^{4.3}$ cm$^{-3}$ for 30 Dor from the intensity ratios of $^{12}$CO($J=7-6$) to $^{12}$CO($J=4-3$) and $^{12}$CO($J=1-0$) to $^{13}$CO($J=1-0$).  In the present study, we aim to obtain sub-mm molecular data at better S/N ratios than in the previous studies to make estimates of temperatures and densities over a large sample in the LMC. We will combine the $^{12}$CO($J=3-2$) data obtained with the ASTE telescope and CO($J=1-0$) data obtained with the SEST and Mopra telescopes. In order to make reasonable comparisons between the two transitions, $J=3-2$ and $J=1-0$, we shall smooth the ASTE results (22$\\arcsec$  beam) to the same resolution as the SEST data (45$\\arcsec$) and use LVG calculations to estimate density and temperature. We shall also employ the $^{13}$CO($J=1-0$) data where available to place constraints on the physical parameters.  This paper is organized as follows: Section 2 describes the observations. Section 3 and 4 show the results and data analysis, respectively. In section 5, we discuss the physical properties of clumps and evolutional sequence of GMCs. Finally, we present a summary in section 6. ",
        "conclusions": "\\subsection{Dense and compact clumps as candidates for proto-cluster condensations}  We have carried out sub-mm $^{12}$CO($J=3-2$) observations of GMCs in the LMC which are most extensive and highly sensitive compared to the previous studies. Six GMCs were selected based on the NANTEN CO survey of the LMC, including 3 Type III GMCs actively forming O stars in addition to 3 Type I/II GMCs which are quiet in O-star formation or cluster formation, although the formation of low to intermediate mass star is not excluded. The spatial resolution of $\\sim$ 5 pc and the high sensitivity allowed us to identify 32 molecular clumps in these GMCs and to reveal significant details of the warm and dense molecular gas with $n$(H$_2$) $\\sim$ 10$^{3-5}$ cm$^{-3}$ and $T_{\\mathrm{kin}}$ $\\sim$ 10 -- 200 K.  The typical mass of the molecular clumps is large, in the range of 5$\\times$10$^3$ -- 2$\\times$10$^5$ M$_{\\odot}$ with radii of 1 -- 12 pc. Of all of our objects, N 159 No.1 or -W shows the strongest concentration of mass of $\\sim$ 7$\\times$10$^4$ M$_{\\odot}$ within a radius of $\\sim$ 5 pc. The masses seem to be larger than those of typical Milky Way GMCs such as those in the eta Car region \\citep[e.g.,][]{Yonekura2005}% , although the propeties of these galactic GMCs are based on optically thin C$^{18}$O data. We suggest that these are good candidates for the precursors of rich super clusters like R136 in 30 Dor which includes more than 10$^4$ stars in a small volume with a radius of $\\sim$ 1 pc. It is of particular interest to look for even denser gas towards them in higher excitation transitions of the sub-mm region.   \\subsection{Density and temperature of the clumps and implications}  The results of our LVG analysis indicate that clumps are distributed from cool to warm in temperature and from less dense to dense in density. These differences of clump properties in density and temperature show good correspondence with the GMC Types based on the star formation activity, as well as with the \\Halpha emission of ionized gas associated with each clump. Clumps in Type III GMCs are warm ($T_{\\mathrm{kin}}$ $\\sim$ 30 -- 200 K) and are either dense ($n$(H$_2$) $\\sim$ 10$^{3.5-5}$ cm$^{-3}$) or less dense ($n$(H$_2$) $\\sim$ 10$^3$cm$^{-3}$). Clumps in Type II GMCs are either warm ($T_{\\mathrm{kin}}$ $\\sim$ 30 -- 200 K) or cool ($T_{\\mathrm{kin}}$ $\\sim$ 10 -- 30 K) and less dense ($n$(H$_2$) $\\sim$ 10$^3$ cm$^{-3}$). Clumps in Type I GMC are cool ($T_{\\mathrm{kin}}$ $\\sim$ 10 -- 30 K) and less dense ($n$(H$_2$) $\\sim$ 10$^3$cm$^{-3}$). The physical parameters of clumps are generally correlated with the star formation activity of GMCs and can perhaps be interpreted in terms of evolutionary effects.  Our interpretation is that defferences of clump density and temperature represent an evolutionary sequence of GMCs in terms of density increase leading to star formation; Type I/II GMCs are at a young phase of star formation where density has not yet reached high enough values to cause active massive star formation, and Type III GMCs represent the later phase where the average density is higher, including both high and low density sub-types. The high density clumps in Type III GMCs show high $R_{3-2/1-0,clump}$ of 1.0 -- 1.5 and are associated with strong \\Halpha flux while the low density clumps in Type III GMCs show low $R_{3-2/1-0,clump}$ of 0.5 -- 1.0 with weak \\Halpha flux.  We suggest two alternative ideas to explain % the density difference of the clumps % in Type III GMCs; one is that density is being enhanced by shock compression driven by HII regions and the other is that gravitational condensation of each clump plays a role in the density increase. The former may be difficult because the shock front may occupy a small volume which is likely missed with the present 5 pc beam. It seems thus favorable that the latter scenario is working mainly to enhance density.  The present study, which resolved the smaller clumps in GMCs at 5 -- 10 pc scales, indicates that the clumps may have physical properties affected by local properties such as the \\Halpha distribution. It should be interesting to investigate the variations among these internal clumps and their relation to star formation.   \\subsection{FUV Heating of the molecular gas in the LMC}  The present findings that the $R_{3-2/1-0,clump}$ is well correlated with \\Halpha flux suggests that the heating of molecular gas by far-ultraviolet (FUV) photons may be effective in the LMC where the dust opacity is lower and the FUV intensity is higher than in the Milky Way. The molecular gas in the Milky Way is mainly heated by cosmic ray protons of $\\sim$ 100 MeV as discussed by a number of authors, although the surface layers of molecular clouds with small visual extinctions at $A$v $\\sim$ a few mag or less may be dominated by the FUV heating \\citep[e.g.,][]{Kaufman1999}. Some authors have made detailed calculations of gas heating and cooling under the effects of FUV radiation fileds \\citep{Kaufman1999}. We shall try to present a picture that can be applied to the present results below.  First, the gas temperature is determined through the balance between the cooling and heating. According to Table 4 in \\citet{Goldsmith1978}, the total cooling rate is 6.8$\\times$10$^{-27}T^{2.2}$ ergs cm$^{-3}$ s$^{-1}$ for $X$(CO)/(d$v$/d$r$) = 4$\\times$10$^{-5}$ and $n$(H$_2$) = 10$^3$ cm$^{-3}$. In 30 Dor region, since $X$(CO)/(d$v$/d$r$) = 3$\\times$10$^{-6}$ / 0.8 = 3.75$\\times$10$^{-6}$ and this value is 10 times lower than the value used in \\citet{Goldsmith1978}, $n$(H$_2$) = 10$^4$ cm$^{-3}$ can be read 10$^3$ cm$^{-3}$, then the cooling rate is estimated as 1.7$\\times$10$^{-22}$ ergs cm$^{-3}$ s$^{-1}$ for $T$ = 100 K. This value is a factor 2 -- 3 smaller than that of Galactic clouds with $n$(H$_2$) = 10$^4$ cm$^{-3}$ and $T$ = 50 K.  We shall assume that the heating by cosmic ray electrons is not important in the LMC. This assumption is not directly confirmed, but it is consistent with the low non-thermal fraction of the LMC's radio continuun emission \\citep[e.g.,][]{Hughes2006} and studies that suggest a significant fraction of cosmic ray electrons are able to escape from low luminosity galaxies \\citep[e.g.,][]{Bell2003, Skillman1988}  The FUV flux ($G_0$) is estimated as 3500 for 30 Doradus \\citep{Bolatto1999, Poglitsch1995, Werner1978, Israel1979}, and 300 for N 159 \\citep{Bolatto1999, Israel1996, Israel1979}. In Orion, it is estimated as 25 \\citep{Bolatto1999, Stacey1993}. The FUV flux in the LMC is larger than that in the Milky Way.  PDR models are calculated by \\citet{Kaufman1999} which incorporate the chemical and physical processes that form and destruct atoms or molecules, as well as ionization effects. Figuer 1 of \\citet{Kaufman1999} shows the kinetic temperature for a molecular gas layer with density of $n$ (cm$^{-3}$) under FUV flux of $G_0$ at the surface. PDR surface temperatures are estimated as listed in Table \\ref{tbl07}. These indicate that temperature becomes as high as 100 -- 300 K on the PDR surface under the conditions of the clumps in Type III GMCs in the LMC. These temperatures are basically consistent with the temperatures of the warm clumps in the present sample.  Generally speaking, at a scale of $\\sim$ 10 pc, $T_{\\mathrm{kin}}$ $\\sim$ 100 K seems to be higher than the kinetic temperatures typical in Milky Way GMCs, where the Milky Way values are usually derived from the $^{12}$CO($J=1-0$) emission only \\citep[e.g., eta Car $T_{\\mathrm{kin}}$ $\\sim$ 50 K by][]{Yonekura2005}. This suggests that the heating of molecular clouds may be stronger in the LMC than in the Milky Way and the molecular temperature may be higher. If this is correct, the lower metallicity, resulting in lower extinction, is the basic cause for the higher temperature in addition to the stronger FUV field in the LMC. We shall note in the end that this higher temperature in the molecular gas possibly leads to an increase of the Jeans mass of molecular clumps, which may favor the formation of rich super clusters in the LMC. This is consistent with the higher mass of the molecular clumps which may represent precursors of the clusters.  The present work has undertaken to sample 6 GMCs (7 regions) to have a uniform determination of the density and temperature in the LMC. The number of GMCs is still limited to 6 among $\\sim$ 300 detected with NANTEN. We should make more efforts to collect appropriate data sets in the sub-mm wavelengths to improve our understanding of the cloud properties. NANTEN2, ASTE and others will certainly be powerful in achieving this goal."
    },
    "0710/0710.3600_arXiv.txt": {
        "abstract": "{Two separate statistical tests are described and developed in order to test un-binned data sets for adherence to the power-law form. The first test employs the TP-statistic, a function defined to deviate from zero when the sample deviates from the power-law form, regardless of the value of the power index. The second test employs a likelihood ratio test to reject a power-law background in favor of a model signal distribution with a cut-off. }    \\begin{document} ",
        "introduction": "The question of whether the cosmic ray energy spectrum exhibits a cut-off at the very highest energies is of central interest to the cosmic ray (CR) physics\\cite{refG,refZK}. The flux of CR's at these energies is very small - about $3 /$km$^{2}$ steradian century - and, therefore, statistical analysis techniques which clearly quantify ones knowledge of flux suppression are useful. In this note we apply the statistics first developed for binned CR data sets in \\cite{refHague} to an un-binned analysis. We also introduce a new test based on a likelihood ratio test and show that both statistics can quantify our knowledge of a flux suppression.  We first establish the mathematical foundations of the analysis. The CR flux follows a power-law for over 10 orders of magnitude. The fundamental probability distribution function (p.d.f.) governing the power-law assumption (normalized such that $\\langle \\rangle_{X} \\equiv \\int_{x_{min}}^{\\infty} f_{{\\scriptscriptstyle X}}(x; x_{min}, \\tg)dx=1$) is \\begin{equation}\\label{equpwlpdf} f_{{\\scriptscriptstyle X}}(x; x_{min}, \\tg) = A\\,x^{-\\tg}, \\end{equation} where $A=(\\tg-1)x^{\\tg-1}_{min}$ and the parameter $\\tg$ is referred to as the {\\it spectral index}.  The $n^{th}$ raw moment of this distribution diverges\\cite{refNewm} for $n \\geq 2$ with $\\tg \\leq 3$. Alternatively, the expected value of $\\ln(x/x_{min})$ is better behaved and offers a crucial result of this analysis. Analytically we find, \\begin{equation}\\label{equnu} \\nu_{n}        \\equiv \\left\\langle \\ln^n \\left( \\frac{x}{x_{min}} \\right) \\right\\rangle_{X} = \\frac{n!}{(\\tg-1)^{n}}. \\, \\end{equation} For a given sample we use, \\begin{equation} \\label{equnuhat} {\\hat \\nu}_{n}(X_{(j)}) \\equiv \\frac{1}{N-(j-1)} \\sum_{i=j}^N \\ln^{n} \\left( \\frac{X_{(i)}}{X_{(j)}} \\right). \\, \\end{equation} In eq.\\ref{equnuhat} we denote the sorted (from least to greatest) data set as $\\left\\{X_{(1)},X_{(2)},\\ldots,X_{(N)}\\right\\}$. To apply these statistics to an un-binned data set we calculate ${\\hat \\nu}_{n}(X_{(j)})$ for each minimum $X_{(j)}$.  We also study a toy p.d.f. which is designed to mimic a power-law up to a certain energy but then exhibit a sharp ``Fermi-Dirac like'' cut-off above that energy\\cite{refHague}. We follow the parameterization used in \\cite{refPAOicrc318}, \\begin{equation}\\label{equFDpdf} f_{{\\scriptscriptstyle FD}}(x; x_{c}, w_{c}, \\tg) = \\frac{B\\,x^{-\\tg}}{1 + \\exp \\left( \\frac{\\log x - \\log x_{c}}{ w_{c}} \\right) }, \\end{equation} where $B$ is chosen such that $f_{{\\scriptscriptstyle FD}}$ is normalized over the interval $[x_{min},\\infty)$, i.e. $\\langle \\rangle_{{\\scriptscriptstyle FD}} = 1$. ",
        "conclusions": "We began this note by verifying that the log-binned spectral index estimator has more bias and a larger error than the un-binned (maximum likelihood) estimator. We then detailed two un-binned statistical tests sensitive to flux suppression. We show that both tests show high sensitivity for rejecting the power-law hypothesis in favor of a toy flux suppression model and depend only weakly on the true spectral index. Applying these tests to $3500$ events drawn from a toy cut-off distribution (see eq. \\ref{equFDpdf}) we can reject the power-law model in favor of the cut-off model at a confidence level $\\sim 4$ standard deviations."
    },
    "0710/0710.4485_arXiv.txt": {
        "abstract": "The dissolution process of star clusters is rather intricate for theory. We investigate it in the context of chaotic dynamics.  We use the simple Plummer model for the gravitational field of a star cluster and treat the tidal field of the Galaxy within the tidal approximation. That is, a linear approximation of tidal forces from the Galaxy based on epicyclic theory in a rotating reference frame. The Poincar\\'e surfaces of section reveal the effect of a Coriolis asymmetry. The system is non-hyperbolic which has important consequences for the dynamics. We calculated the basins of escape with respect to the Lagrangian points $L_1$ and $L_2$. The longest escape times have been measured for initial conditions in the vicinity of the fractal basin boundaries. Furthermore, we computed the chaotic saddle for the system and its stable and unstable manifolds. The chaotic saddle is a fractal structure in phase space which has the form of a Cantor set and introduces chaos into the system. ",
        "introduction": "The dissolution process of star clusters is an old problem in stellar dynamics. Once a star cluster has formed somewhere in a galaxy, it tends to lose mass due to dynamical interactions until it has completely dissolved. It turned out that the physics behind the dissolution of star clusters is intricate and fascinating for theory. If a star cluster of finite mass were isolated, in virial equilibrium (i.e. $\\langle v_e^2 \\rangle = 12\\sigma_{1D}^2$) and the velocity distribution given by a Maxwellian $f_M(X) = \\left(4/\\sqrt{\\pi}\\right) X^2 \\exp(-X^2)$, where  $X=v/(\\sqrt{2}\\sigma_{\\rm 1D})$ and $v, v_e$ and $\\sigma_{1D}$ are the velocity, the escape velocity and the velocity dispersion, respectively, the fraction of stars which are faster than the rms escape speed were given by $\\chi_e = \\int_{\\sqrt{6}}^\\infty f_M(X) dX = 2\\sqrt{6/\\pi} \\exp(-6) + {\\rm erfc}\\sqrt{6} \\simeq 0.00738316$. This simple analytical result was published by Ambartsumian (1938) and two years later, independantly by Spitzer (1940) who named this effect ``evaporation''. The relevant process which brings stars above the escape speed and lets them evaporate, is two-body relaxation. The time scale of relaxation, which determines the rate of dynamical evolution of a star cluster, yields thus an upper limit to the lifetime of any star cluster. However, since $\\chi_e$ is so small (and relaxation time relatively long), the evaporation time is much longer than a Hubble time for typical globular clusters. Following a suggestion of Chandrasekhar (1942),  King (1959) studied the effect of ``potential escapers''. These are stars which have been scattered above the escape energy but which have not yet left the cluster. These may be scattered back to negative energies within a crossing time and remain bound.  H\\'enon (1960, 1969) stressed the importance of few close encounters between stars for the rate of mass loss of star clusters.  The Fokker-Planck approximation, which is widely used to study the dynamical evolution of star clusters, neglects strong encounters by construction. Nevertheless, close encounters could still be interpreted statistically as a certain discontinuous Markov process (Tscharnuter 1971). However, direct $N$-body models seem to be ideally suited to study this phenomenon in more detail. Spitzer \\& Shapiro (1972) estimate that `` close encounters may produce effects perhaps as great as 10 percent of the ``dominant'' distant encounters'' and ignore them.  Nature provides an environment for star clusters in which the escape rate is typically strongly enhanced as compared with the slow evaporation rate of isolated star clusters: The tidal field of a galaxy induces saddle-like troughs in the walls of the potential well of a star cluster (cf. Figure 1). It therefore lowers the energy threshold in star clusters above which stars can escape from zero to a negative value (Wielen 1972, 1974). Moreover, if we consider a star cluster in the tidal field of a galaxy as a dynamical system, the tidal field can change the system's dynamics in a dramatical way as compared with an isolated system.  In general, the escape process from star clusters in a tidal field proceeds in two stages: (1) Scattering of stars into the ``escaping phase space'' by two-body encounters and (2) leakage through openings in the equipotential sufaces around saddle points of the potential. The ``escaping phase space'' is defined as the subset of phase space, from which escape is possible.  It is well understood that the time scale for a star to complete stage (1) scales with relaxation time. On the other hand, the time scale for a star to complete stage (2) depends mainly on its energy (but also on its location in phase space as we will see). When we neglect the effect of two-body relaxation for the consideration of stage (2), the motion  of a single star in the star cluster is determined between times $t_1$ and $t_2$ only by the smooth gravitational potential in which the star moves. The potential itself is generated by the other stars in the star cluster disregarding their ``grainyness'' and by the superposed galactic gravitational field, which is due to the matter distribution of the galaxy. Within this framework, we will study stage (2) of the escape process in this paper. In this connection, the work of Fukushige \\& Heggie (2000) is of major interest. Their main result is an expression for the time scale of escape for a star in stage (2) which has just completed stage (1). The dependance of the escape process on two time scales which scale differently with the particle number $N$ imposes a severe scaling problem for $N$-body simulations. The scaling problem is of relevance since the it is on today's general-purpose hardware architectures not yet simply feasible to simulate the evolution of globular clusters with realistic particle numbers of a few hundred thousands or even millions of stars by means of direct $N$-body simulations. The result of Fukushige \\& Heggie has been applied in Baumgardt (2001) to solve the important scaling problem for the dissolution time of star clusters in the special case of circular cluster orbits. The obtained scaling law $t_{dis}\\propto t_{rh}^{3/4}$, where $t_{dis}$ and $t_{rh}$ are the dissolution and half-mass relaxation times, respectively has been verified, e.g. in Spurzem et al. (2005).  The problem of escape has also a long history in the context of the theories of dynamical systems and chaos. It is well-known for a long time, that certain Hamiltonian systems allow for escape of particles towards infinity. Such ``open'' Hamiltonian systems have been studied by Rod (1973), Churchill et al. (1975), Contopoulos (1990), Contopoulos \\& Kaufmann (1992), Siopis et al. (1997), Navarro \\& Henrard (2001) and Schneider, T\\'el \\& Neufeld (2002). The related chaotic scattering process, in which a particle approaches a dynamical system from infinity, interacts with the system and leaves it, escaping to infinity, was investigated by many authors, as Eckhardt \\& Jung (1986), Jung (1987), Jung \\& Scholz (1987), Eckhardt (1987), Jung \\& Pott (1989), Bleher, Ott \\& Grebogi (1989) and Jung \\& Ziemniak (1992). Chaotic scattering in the restricted three-body problem has been studied by Benet et al. (1997, 1999). Typically, the infinity acts as an attractor for an escaping particle, which may escape through different exits in the equipotential surfaces. Thus it is possible to obtain basins of escape (or ``exit'' basins), similar to basins of attraction in dissipative systems or the well-known Newton-Raphson fractals. Special types of basins of attraction (i.e. ``riddled'' or ``intermingled'' basins)  have been explored by Ott et al. (1993) and Sommerer \\& Ott (1993, 1996). Basins of escape have been studied by Bleher et al. (1988), and they are discussed in Contopoulos (2002). Reasearch on escape from the paradigmatic H\\'enon-Heiles system has been done by de Moura \\& Letelier (1999), Aguirre, Vallejo \\& Sanju\\'an (2001), Aguirre \\& Sanju\\'an (2003), Aguirre, Vallejo \\& Sanju\\'an (2003), Aguirre (2004) and Seoane Seoane Sep\\'ulveda (2007). These papers served as the basis of this work. Relatively early, it was recognized, that the key to the understanding of the the chaotic scattering process is a fractal structure in phase space which has the form of a Cantor set (Cantor 1884) and is called the chaotic saddle. Its skeleton consists of unstable periodic orbits (of any period) which are dense on the chaotic saddle (e.g. Lai 1997) and introduce chaos into the system (e.g. Contopoulos 2002). The properties of chaotic saddles have been investigated by different authors, as Hunt (1996), Lai et al. (1993), Lai (1997) or Motter \\& Lai (2001). Both hyperbolic and non-hyperbolic chaotic saddles occur in dynamical systems. In the first case, there are no Kolmogorov-Arnold-Moser (KAM) tori, which means that all periodic orbits are unstable. In the second case, there are both KAM tori and chaotic sets in the phase space (J. C. Vallejo, priv. comm. and e.g. Lai et al. 1993). We note that all of the above references on the chaotic dynamics are exemplary rather than exhaustive since there exists a vast amount of literature on these topics.  The aim of this paper is to allude to the importance of this last-mentioned branch of research for the field of stellar dynamics of star clusters. We will study the escape process from star clusters within the framework of chaotic dynamics. In section 2, we introduce the tidal approximation, i.e. approximate equations of motion for stellar orbits in a star cluster which is embedded in the tidal field of a galaxy. In section 3, we describe our (very simple) model of the gravitational potential. In section 4, we discuss Poincar\\'e surfaces of section, which show the effect of a Coriolis asymmetry. Furthermore, we discuss the basins of escape in Section 5 and the chaotic saddle and its stable and unstable invariant manifolds in Section 6. Section 7 contains the discussion and conclusions. ",
        "conclusions": "We studied the chaotic dynamics within a star cluster which is embedded in the tidal field of a galaxy. We calculated within the framework of the tidal approximation Poincar\\'e surfaces of section, the basins of escape, the chaotic saddle and its stable and unstable manifolds as well. The system is non-hyperbolic which has important consequences for the dynamics, i.e. there are orbits which do not escape if relaxation is neglected. These are mainly the retrograde orbits as has been shown earlier by Fukushige \\& Heggie (2000) and, more recently, in the $N$-body study by Ernst et al. (2007). The escape times are longest for initial conditions near the fractal basin boundaries. The decay law is a power law for those stars which escape from the regions without sensitive dependance on the initial conditions in Figure \\ref{fig:basins} (i.e. with short escape times, as can be seen in Figure \\ref{fig:times}). On the other hand, the decay law is exponential for orbits which escape from the regions with sensitive dependance on the initial conditions (i.e. with long escape times). The effect of relaxation (i.e. a diffusion in the Jacobian and the third integral among different stellar orbits) on the chaotic dynamics which we investigated in this work may be a very interesting topic for future research."
    },
    "0710/0710.1096_arXiv.txt": {
        "abstract": "We report first results of a multifrequency campaign from radio to hard X-ray energies of the prominent \\gray\\ blazar 3C~279, which was organised around an INTEGRAL ToO observation in January 2006, and triggered on its optical state. The variable blazar was observed at an intermediate optical state, and a well-covered multifrequency spectrum from radio to hard X-ray energies could be derived. The SED shows the typical two-hump shape, the signature  of non-thermal synchrotron and inverse-Compton (IC) emission from a relativistic jet. By the significant exposure times of INTEGRAL and Chandra, the IC spectrum (0.3 - 100 keV) was most accurately measured, showing -- for the first time -- a possible bending. A comparison of this 2006 SED to the one observed in 2003, also centered on an INTEGRAL observation, during an optical low-state, reveals the surprising fact that -- despite a significant change at the high-energy synchrotron emission (near-IR/optical/UV) -- the rest of the SED remains unchanged. In particular, the low-energy IC emission (X- and hard X-ray energies) remains the same as in 2003, proving that the two emission components do not vary simultaneously, and provides strong constraints on the modelling of the overall emission of 3C~279. ",
        "introduction": "The discovery by the experiments aboard the Compton Gamma-Ray Observatory (CGRO) that blazars can radiate a large -- sometimes even the major -- fraction of their luminosity at \\gray\\ energies marked a milestone in our knowledge on blazars. During the CGRO mission about 90 blazars were detected by the different CGRO experiments at \\gray\\ energies, the majority by the EGRET experiment at energies above $\\sim$100~MeV \\citep{Hartman99}. 3C~279, an optically violently variable (OVV) quasar located at a redshift of 0.538, is one of the most prominent representatives of these sources. The source shows rapid variability in all wavelength bands, polarized emission in radio and optical, superluminal motion, and a compact radio core with a flat radio spectrum. These properties put the quasar 3C~279 into the blazar sub-class of AGN. According to the unified model of Active Galactic Nuclei (AGN), blazars are sources which expel jets close to our line-of-sight.  3C~279 was already detected by INTEGRAL in June 2003 \\citep{Collmar04}. Those high-energy observations were supplemented in X-rays by a short (5 ksec) Chandra  pointing, and by ground based monitoring from radio to optical bands. Since the blazar was found in 2003 at the faintest optical brightness (optical R-band: $\\sim$17~mag) of the last 10 years, roughly 5 mag fainter than the maximum, and about 2.5 to 3 mag fainter than average, a simultaneous spectral energy distribution (SED) of an exceptional optical low-state could be compiled \\citep{Collmar04}. In order to measure emission changes as function of optical brigthness, we proposed for an INTEGRAL ToO observation during an optical high state (criterion: optical R-band brigther than 14.5~mag). In January 2006, the trigger citerion was met, and the INTEGRAL observations together with supplementing multifrequency observations were carried out. This campaign resulted in a well covered SED from radio to hard X-ray energies.   In this paper we present first results of this 2006 multiwavelength campaign on 3C~279. Because of the page restrictions, we concentrate on presenting the main observational results, focussing on the observed SED and its comparison to the one of 2003. A more detailed presentation, including a discussion on the scientific implications of the new results, will be given in a later paper (Collmar et al., in prep.; including also all participating individual WEBT (Whole Earth Blazar Telescope) collaborators in the author list). In addition, the results on variability analyses and time correlations of the different wavelength bands will be given elsewhere (B\\\"ottcher et al., in prep.).  \\begin{figure}[th] \\centering \\epsfig{figure=fig_1a.eps,width=8.0cm,clip=} \\epsfig{figure=fig_1b.eps,width=8.1cm,clip=} \\caption{ Top: The ISGRI image shows a 6.4\\sig\\ detection of 3C~279 in the 20-60~keV band. In addition, the Seyfert galaxy NGC 4593 is even more clearly detected. \\newline Bottom: The ISGRI hard X-ray spectrum between 20 and 100 keV is shown, together with the best-fit power-law shape. \\label{fig:1} } \\end{figure} ",
        "conclusions": ""
    },
    "0710/0710.2789_arXiv.txt": {
        "abstract": "{In December 2004, the soft gamma-ray repeater \\src\\ emitted the most powerful giant flare ever observed. This probably involved a large-scale rearrangement of the magnetosphere leading to observable variations in the properties of its X-ray emission. Here we present the results of the first \\suz\\ observation of \\src, together with almost simultaneous observations with \\xmm\\ and \\int. The source seems to have reached a state characterized by a flux close to the pre-flare level and by a relatively soft spectrum. Despite this, \\src\\ remained quite active also after the giant flare, allowing us to study several short bursts observed by \\suz\\ in the 1--100 keV range. We discuss the broad-band spectral properties of \\src, covering both persistent and bursting emission, in the context  of the magnetar model, and consider its recent theoretical developments. ",
        "introduction": "The four known Soft Gamma-ray Repeaters (SGRs) were discovered as transient sources of high-energy photons; they emit sporadic and short ($\\sim$0.1 s) bursts of (relatively) soft gamma-rays with luminosity $L\\sim10^{40}$--$10^{41}$ \\lum during periods of activity, that are often broken by long intervals of quiescence. Three ``giant'' flares with luminosity \\mbox{$\\gtrsim$$10^{43}$ \\lum} have also been observed to date, each one from a different SGR: on March 5, 1979 from SGR\\,0526--66 in the Large Magellanic Cloud \\citep{mazets79}, on August 27, 1998 from SGR\\,1900+14 \\citep{hurley99}, and on December 27, 2004 from \\src\\ \\citep{hurley05}. Persistent emission with $L\\sim10^{35}$ \\lum\\ is also observed from SGRs in  the soft X--ray range (\\mbox{$<$10 keV}) and, for \\src\\ and SGR\\,1900+14, also in the hard X-ray range \\citep{mgm05,gotz06}. In three cases, periodic pulsations at a few seconds have been detected. The bursts, the giant flares, the quiescent X-ray counterparts, and the pulsations have been interpreted in the framework of the magnetar model \\citep[see][and references therein]{tlk02}. Magnetars are highly magnetized neutron stars with field strengths of \\mbox{$10^{14}$--$10^{15}$ G}, larger than those of the majority of radio pulsars. The ultimate source of energy for the bursts and the quiescent emission is believed to be the ultra-strong magnetic field.\\\\ \\indent \\src\\ was discovered in 1979 \\citep{laros86,laros87} and its persistent X-ray counterpart was observed for the first time with the \\asca\\ satellite in 1993 \\citep{murakami94}. A \\xte\\ observation led to the discovery of pulsations in the persistent emission with period $P\\simeq7.47$ s and period derivative $\\dot{P}\\simeq2.6\\times10^{-3}$ s  yr$^{-1}$ \\citep{kouveliotou98}. Under the assumption of pure magnetic dipole braking, these values imply a surface magnetic field strength of \\mbox{$8\\times10^{14}$ G}, strongly supporting the magnetar model. Both the burst rate and the X-ray persistent emission of \\src\\ started increasing during 2003 and throughout 2004 \\citep{mte05,tiengo05,met07,woods07}, culminating with the giant flare of December 27, 2004, during which $\\sim$$10^{47}$ erg were released\\footnote{Assuming isotropic luminosity and for a distance $d=15$ kpc \\citep{corbel97,mcclure05}. } \\citep{hurley05,mereghetti05,terasawa05}. This giant flare was exceptionally intense, $\\sim$100 times more energetic than those from  SGR\\,0526--66 and SGR\\,1900+14. Observations with \\xte\\ unveiled, for the first time in an isolated neutron star, rapid quasi-periodic oscillations in the pulsating tail of the flare, likely related to global seismic oscillations on the neutron star surface \\citep{ibs05}. The flare produced a hard X-ray (\\mbox{$>$80 keV}) afterglow lasting a few hours \\citep{mereghetti05,frederiks07} and a radio afterglow that faded in a few days \\citep{cameron05}. The small positional uncertainty of the radio observations permitted to identify the likely IR counterpart of the SGR \\citep{kosugi05,israel05}. The fluxes observed in the IR and gamma energy bands show a variability correlated with that observed in the 2--10 keV energy range \\citep{met07}.\\\\ \\indent After the giant flare, the persistent X-ray flux of \\src\\ started to decrease from its outburst level, and its X-ray spectrum to soften \\citep{rea05,rea05_atel,met07,tiengo05,woods07}. A Similar flux decrease have been observed from its radio afterglow \\citep{gaensler05,taylor05} and its newly discovered IR counterpart \\citep{israel05,rea05_atel,met07}.\\\\ \\indent Here we present the results of the first \\suz\\ observation of \\src, covering both persistent and bursting emission in the 1--100 keV energy band. We also report on the analysis of a simultaneous observation performed with \\xmm\\ and the  latest outcomes of the monitoring of \\src\\ with \\int, comparing them with what is seen in the same energy ranges with \\suz. ",
        "conclusions": "The \\suz, \\xmm, and \\int\\ observations reported here represent a complementary data set that allows us to study the spectral properties of \\src\\ in the broad \\mbox{1--100 keV} energy range. Although the  \\suz/HXD does not have imaging capabilities, we know thanks to  \\int\\ that no other bright hard X-ray sources are present in its field of view.  The uncertainties in the instrumental background (currently at the $\\sim$5\\% level) and in the modelling of the Galactic Ridge emission are a more relevant concern. Future improvements in the knowledge of these components may eventually allow us to obtain more robust conclusions. Thanks to its imaging capabilities, background issues do not affect the \\int\\ observations. However, the IBIS/ISGRI data required long integration times, with discontinuous observations spanning several months. Thus they provide information on the average properties only. The possible presence of long term variability was in fact one of the main motivations to perform the \\suz\\ and \\xmm\\ observations simultaneously. With all these caveats in mind we proceed now to discuss the broad band spectrum of \\src.\\\\ \\indent The \\xmm\\ and \\int\\ spectra are consistent with an extrapolation of the power-law plus blackbody model measured in the 2--10 keV band. Between 12 and 30 keV the \\suz/HXD sensitivity is better than that of \\int, allowing to detect \\src\\ during a single 50 ks long observation. With respect to the \\xmm\\ and \\int\\ joint fit, the HXD data show an ``excess'' (see Fig.\\,\\ref{broad}) that cannot be completely ascribed to calibration uncertainties between the various instruments.\\\\ \\indent Given the lack of a direct measure of the Galactic Ridge emission around \\src, we cannot exclude that this excess is due to an underestimation of such contribution to the background. If instead the excess is a real feature of the spectrum of \\src, its broadband spectrum could be empirically modeled adopting a power-law with the photon-index changing from $\\sim$1 to $\\sim$2 at $\\sim$16 keV, and a blackbody component with $k_BT\\sim0.8$ keV. This would agree with the results reported by \\citet{gotz06}, who point out that the hard tails of the SGRs are softer than the power-law components measured below \\mbox{10 keV}.\\\\ \\indent The presence of a down-break in the 10--20 keV spectrum of \\src\\ would have remarkable physical implications. The soft X-ray emission from magnetar candidates (SGRs and AXPs) is usually interpreted within the twisted magnetosphere model as due to resonant cyclotron up-scattering of soft photons from the star surface onto charges flowing into the magnetosphere \\citep{tlk02}. Detailed calculations of resonant compton scattering (RCS) spectra have been recently performed \\citep{lyutikov06,fernandez07} and successfully applied to fit AXP spectra \\citep{rzl07}. Quite interestingly, some of the model spectra presented by \\citet{fernandez07} exhibit a downward break in the tens of keV range. Their overall shape is quite reminiscent of the \\suz\\ XIS/HXD-PIN spectrum of \\src, and, as noted by \\citet{fernandez07}, they may also play a role in the interpretation of the broadband X-ray spectrum of SGR\\,1900+14. In particular, when assuming a (non-thermal) top-hat or a broadband velocity distribution for the magnetospheric charges, multiple peaks can appear in the spectrum (see their Fig.\\,6 and Fig.\\,11). The downturn possibly present in our data may then be due the presence of a second ``hump'' (in addition to the main thermal one) in the range \\mbox{10--20~keV}. Nobili, Turolla, \\& Zane (in preparation) assuming a 1-D thermal electron distribution superimposed to a (constant) bulk velocity, found also double humped spectra. In this case the second (and only) hump occurs when resonant scattering is efficient enough to fill the Wien peak at the temperature of the comptonising particles. A spectral break at $\\sim$15~keV would translate then in a temperature of $\\sim$5 keV for the magnetospheric electrons. If a more refined treatment of background subtraction confirms the spectral break in the X-ray data of \\src\\, this would provide important diagnostics for the physical parameters of the model.\\\\ \\indent In 2003 we started a long-term monitoring program to study the time evolution of the spectral properties of \\src\\ using the \\xmm\\ X-ray satellite. The December 27, 2004 giant flare was a fortunate occurrence that allowed us to observe how the source properties evolved in the two years leading up to the flare and how they changed after this dramatic event \\citep[see][for details]{mte05,met07,rea05_atel,tiengo05}. The \\xmm\\ data showed a doubling of the flux in the September--October 2004 followed by a gradual recovery to the ``historical'' level after the giant flare. A direct comparison of the \\xmm/pn count rates measured in the different observations shows that before the giant flare the flux of \\src\\ in the \\mbox{1--10 keV} band was monotonically increasing, while the three observations after the flare, and preceding the one reported here, followed a steady decreasing trend. The September 2006 observation breaks this long term decay, having a count rate higher by 5\\% with respect to the last \\xmm\\ observation performed 5 months before. This slight (but statistically significant) re-brightening might indicate either a temporary oscillation around an equilibrium flux level or the start of a  new monotonic flux increase, similar to the one that preceded the December 2004 giant flare. This phenomenon was interpreted as due to the building up of a magnetospheric twisting, that determined also the hardening of the X-ray spectrum, an increase of the spin-down and a more intense bursting activity \\citep{mte05}. The relatively high burst rate observed during the \\xmm\\ and \\suz\\ observations of September 2006 (see Section\\,\\ref{burstingemission}) is therefore another indication of a possible increase of the magnetospheric twisting in \\src, but, before a new \\xmm\\ observation will be performed, only the monitoring of the frequency and intensity of \\src\\ bursts can tell us if the evolution is erratic or follows a stable trend. The recent report of a bright burst from \\src\\ \\citep{golenetskii07,perotti07} seems actually to favour the second hypothesis.\\\\ \\indent Two of the bursts detected during the \\suz\\ observation were bright enough to allow spectral analysis. In both cases, the broadband spectrum (\\mbox{2--100 keV}) revealed the presence of two components: a soft component which is well reproduced by a blackbody with $k_BT \\sim 2$--4~keV, and a harder one whose spectral shape is not firmly established and can be equally well fit with a power-law, a hot bremsstrahlung or a second blackbody (See Table\\,\\ref{bfits} and Fig.\\,\\ref{burstfits}). In absence of robust theoretical predictions, we can not exclude that a two component model simply reflects our ignorance of the correct spectral shape, and has therefore a purely phenomenological significance. However, it is worth noticing that, from their recent analysis of a sample of 50 bursts detected from \\src\\ with \\hete\\ from 2001 to 2005, \\citet{nakagawa07} have suggested the presence of a time delay between the 30--100~keV and the 2--10~keV emission. Although such a delay can be attributed to an intrinsic, rapid spectral softening, an alternative, and simpler, interpretation invokes the presence of two separate emitting regions.\\\\ \\indent Let us consider a scenario in which the two components are physically distinct and let us consider the hard component first. In the magnetar scenario, short bursts are usually ascribed to either reconnection phenomena in the external magnetosphere \\citep[eventually modulated by a tearing instability, see][]{lyutikov03} or movements of the footprints of the external magnetic field, produced by crustal deformations or fractures driven by the stress exerted by the internal field helicity \\citep{thompson95,thompson01,tlk02}. Both kind of processes lead to the generation and launch of an Alfv\\`en wave, which produces and accelerates a particle cascade, and ultimately is detected as a burst. The emerging spectrum is expected to be synchrotron dominated, unless the Alfv\\`en wave is temporarily trapped in a fireball and thermalized. Therefore, both the \\mbox{BB + PL} and BB + BB spectral fits are consistent with a scenario in which such an Alfv\\`en wave is responsible for the hard component. We notice that, although a fireball formation is not required to explain short bursts (and therefore usually not invoked in such cases), to our knowledge there is no a priori reasons why a small fireball can not be created and evaporated in a sub-second time scale, giving rise to a thermal spectrum. A point in favour of this interpretation is that, in the BB + BB fit, the temperature of the hot blackbody is remarkably close to the minimal temperature above which a fireball thermalizes, $k_BT \\approx 11\\,(R_{\\rm{NS}}/10\\,\\rm{km})^{-1/5}$~keV, according to \\citet{thompson01}, eq. [71] \\citep[see also][for similar findings based on longer duration burst]{olive04}. The third model of the hard component which is compatible with our data is a bremsstrahlung emission at $\\sim$100~keV. Quite recently, \\citet{thompson05} and \\citet{beloborodov07} discussed the electrodynamics of the magnetar coronae and the production mechanisms for soft gamma-rays. In particular, their model predicts the existence of a thin transition layer between the corona and the thermal photosphere, where Langmuir turbulence can be excited by a downward beam of current-carrying charges. As a result, the transition layer can be heated up to a typical temperature of $\\sim$100~keV, and emit,  approximatively, an optically thin bremsstrahlung at a single temperature. Although \\cite{thompson05} model was originally developed in connection with the persistent hard emission of magnetars, the predicted bremsstrahlung temperatures are remarkably close to those we detected during the two bursts, suggesting that a similar mechanism may instead be activated during periods of activity.\\\\ \\indent Our results about the spectral modelling of the soft X-ray component are more robust, inasmuch as all our spectral fits require the presence of a cold blackbody with \\mbox{$k_BT \\sim 2$--4~keV}. This is in agreement with similar findings by \\cite{olive04}, \\cite{nakagawa07}, and \\citet{feroci04}. This component is usually interpreted as due to emission from a fraction of the star surface (which can be as large as the whole star in the case of our BB + BB fit) heated by returning currents. Alternatively, it has been suggested that the soft component may originate up in the magnetosphere ($\\leq$700~km), presumably due to a delayed emission process \\citep{nakagawa07}. Here we only notice that, although the spectra of our two events are compatible with emission from the star surface, the radius of the cold blackbody as measured during other short bursts can reach values much higher than 50--100~km \\citep[][similar findings have been found in the case of SGR\\,1900+14 bursts measured with \\emph{Swift}, Israel et al. (in preparation)]{nakagawa07}. One possible explanation is that part of the flare energy is intercepted and reprocessed in a larger region and re-emitted at a lower temperature. In such scenario, the radius of the reprocessing region can then vary depending on the fraction of material that is intercepted, and is not bounded by the value of the star radius. \\citet{thompson01} considered the equilibrium state of a pair corona sufficiently extended that the local value of the magnetic field is $B\\ll B_{\\rm{QED}}$, so that photon splitting can be ignored (if the magnetic field scales as a dipole, this occurs above $\\sim$3 star radii, for a polar surface value of $\\sim$$10^{15}$~G). They found that a stable balance between heating and diffusive cooling requires a continuous source of ordinary photons that can be provided, for instance, by external illumination. If the corona intercepts part of the flare beam (which in their treatment is assumed to originate in a trapped fireball, although in general this is not necessarily required), equilibrium is possible below a critical luminosity given by \\begin{displaymath} L<L_{\\rm{max}} = 1.5 \\times 10^{42}\\, \\tau_{EO}^{-1} \\left ( \\frac {k_BT_e} {20\\ {\\rm keV} } \\right )^4 \\left ( \\frac {R }{10\\ {\\rm km}} \\right )^2 \\, {\\rm erg\\ s^{-1}} \\, , \\end{displaymath} where $T_e$ is the pair temperature in the corona, $\\tau_{EO} \\sim 1$ is the scattering depth for ordinary to extraordinary mode conversion and $R$ the radius of the emitting part of the corona (see equations [84] and [89] in \\cite{thompson01} and note that there is a typo in their equation [89]: it should contain $R^2$ instead of $R^{-2}$). Even for an emission region as small as 5--20~km (as inferred by our best fit of the low temperature blackbody) and a temperature of $k_BT \\sim 2$--4~keV, this is well above the luminosity emitted during the two bursts detected by \\suz. Therefore, simply on the basis of energetics, relatively large emitting regions for the cold blackbody are compatible with the temporary formation of a pair corona, sustained by a fraction of the flare energy. When the heating rate ceases, the pair atmosphere contracts and quickly evaporates. In order to derive firmer conclusions a more detailed analysis is needed, mainly in assessing the possibility that the intercepted beam is thermalized and re-emitted as a blackbody. This study is beyond the purpose of this paper, and will be presented elsewhere (Israel et al. 2007, in preparation)."
    },
    "0710/0710.0090_arXiv.txt": {
        "abstract": "An $\\sim$800~arcmin$^2$ mosaic image of the W3 star forming complex obtained with the {\\it Chandra X-ray Observatory} gives a valuable new view of the spatial structure of its young stellar populations. The {\\it Chandra} image reveals $\\sim 1300$ faint X-ray sources, most of which are PMS stars in the cloud.  Some, but not all, of the high-mass stars producing hypercompact and ultracompact H{\\sc II} (UCHII) regions are also seen, as reported in a previous study.  The {\\it Chandra} images reveal three dramatically different embedded stellar populations.  The W3~Main cluster extends over 7~pc with $\\sim 900$ X-ray stars in a nearly-spherical distribution centered on the well-studied UCHII regions and high-mass protostars. The cluster surrounding the prototypical UCHII region W3(OH) shows a much smaller ($\\leq 0.6$~pc), asymmetrical, and clumpy distribution of $\\sim 50$ PMS stars. The massive star ionizing the W3~North H{\\sc II} region is completely isolated without any accompanying PMS stars.  In W3~Main, the inferred ages of the widely distributed PMS stars are significantly older than the inferred ages of the central OB stars illuminating the UCHIIs.  We suggest that different formation mechanisms are necessary to explain the diversity of the W3 stellar populations: cluster-wide gravitational collapse with delayed OB star formation in W3~Main, collect-and-collapse triggering by shock fronts in W3(OH), and a runaway O star or isolated massive star formation in W3~North. ",
        "introduction": "}  W3 (Westerhout 3) is perhaps the most active region of current star formation in the nearby Galaxy.  Extending 30~pc along the edge of a $M \\simeq 5 \\times 10^4$~M$_\\odot$ giant molecular cloud (GMC), the star-forming complex has dozens of embedded young massive stars producing a variety of pre-stellar condensations, hot molecular cores, hypercompact to small H{\\sc II} regions, maser clusters, and molecular outflows \\citep[e.g.,][]{Lada78, Reid80, Dreher81, Tieftrunk97, Chen06}. Its infrared sources have an integrated luminosity of several~$\\times~ 10^5$~L$_\\odot$.  Situated just east of W3 are the older IC~1795 and IC~1805 clusters, the latter lying within the enormous W4 superbubble/chimney structure blown by generations of massive stars.  The W4--IC~1795--W3 complex is widely considered to be an examplar of sequential triggered star formation  \\citep{Lada78, Oey05}. Recent SCUBA observations of the W3 GMC find a higher percentage of the gas mass gathered into dense molecular clumps  at the eastern edge compared to the undisturbed parts of the W3 GMC, supporting this triggering scenario \\citep{Moore07}. A detailed description of the W3 and W4 complexes and a thorough review of the literature are given by \\citet{Megeath07}.  The richest site of massive star formation in W3 is the W3~Main cluster of embedded OB stars, dominated by the very young and luminous IRS~4 and IRS~5 sources.  IRS~5 lies at the center of a 0.1~pc concentration of massive stars resembling a nascent counterpart of the Orion Trapezium \\citep{Megeath05}.  W3(OH) to the southeast and W3~North to the north have massive stars but appear less active than W3~Main.  The distance to the complex is accurately measured from maser kinematics to be 2.0~kpc \\citep{Xu06, Hachisuka06}.  Despite the intense study of W3 at radio, millimeter, and infrared wavelengths, little is known about its low mass stellar population. For example, a $JHK$ near-infrared (NIR) survey of $5\\arcmin \\times 5\\arcmin$ in W3~Main reveals $\\sim 40$ sources with $K$-band excesses indicative of Class~I--II pre-main sequence (PMS) stars with disks \\citep{Ojha04}.  Hundreds of other stars are detected, but infrared photometry cannot discriminate disk-free Class~III PMS stars from the strongly contaminating population of unrelated Galactic field stars (mostly red giants).  A new mid-infrared (MIR) photometric survey of W3~Main, IC~1795, and W3(OH) with the {\\it Spitzer Space Telescope} helps to identify cluster members \\citep{Ruch07}, but it suffers confusion from three effects: foreground and background Galactic field stars \\citep{Benjamin05}, bright diffuse emission produced by heated dust around the H{\\sc II} regions \\citep{Povich07}, and extragalactic objects with MIR excesses \\citep{Harvey07}.  X-ray surveys of young stellar clusters (YSCs) with the {\\it Chandra X-ray Observatory} are surprisingly efficient at detecting low mass PMS populations, even at distances around 2~kpc and at obscurations $10<A_V<150$ mag typical for W3 stars.  PMS X-ray emission arises primarily from violent magnetic reconnection events, similar to solar flares but far more powerful, and is largely independent of circumstellar disks or accretion \\citep[see reviews by][]{Feigelson99, Feigelson07}. Luminous and spectrally hard X-ray flares are present throughout the PMS phases of Class~I--II--III at levels $10^2-10^3$ above that seen in old disk populations \\citep{Preibisch05a}, so relatively few Galactic disk interlopers appear in X-ray samples.  These field star X-ray sources and extragalactic contaminants are easily removed \\citep{Getman06}.  Due to a poorly-understood statistical association between X-ray luminosity and PMS stellar mass \\citep{Preibisch05b,Guedel07}, a flux-limited X-ray observation of a young stellar cluster will be roughly complete down to a corresponding mass limit.  Taking together these properties of X-ray studies, we find that X-ray surveys at sufficiently high spatial resolution and sensitivity provide uniquely rich, largely disk-unbiased, mass-limited, and nearly uncontaminated samples of PMS stars in both embedded and unobscured YSCs. These samples complement MIR surveys obtained with the {\\it Spitzer Space Telescope}, which generally extend down to lower masses (including the brown dwarf regime) but cannot readily discriminate disk-free PMS stars from field stars. {\\it Spitzer} thus detects more disky Class~0-I-II systems while {\\it Chandra} effectively samples Class~III systems in addition to many Class I and II stars.  The X-ray samples are useful for various astrophysical purposes such as probing the stellar Initial Mass Function, protoplanetary disk evolution, and magnetic activity.  In an early {\\it Chandra} study, two $\\sim$20~ks exposures of a $\\sim 300$~arcmin$^2$ field in W3~Main revealed 236 X-ray sources \\citep{Hofner02}.  Several are associated with massive stars ionizing H{\\sc II} regions but most do not have counterparts in $JHK$ images. We report here an extension of those efforts with a {\\it Chandra} mosaic of 7 exposures totaling $\\sim 230$~ks over $\\sim 800$~arcmin$^2$, spanning much of the W3 star forming complex (Figure~\\ref{fig:ACIS_mosaic}).  A preliminary discussion of this mosaic, a {\\it Chandra} Large Project, is given by \\citet{Townsley06}.  Over 1300 X-ray sources are seen; a full listing and study of their properties will be presented in a separate paper.  For each source, {\\it Chandra} observations provide a sub-arcsecond position, line-of-sight absorption, and rough mass estimate in addition to magnetic activity characteristics.  We discuss here insights into the global structure and origins of the W3 stellar populations derived from the new {\\it Chandra} data. The brief presentation of the observations in \\S\\ref{sec:obs} will be expanded in a forthcoming paper with complete source lists similar to our group's recent studies of the Cep~OB3b \\citep{Getman06}, Pismis~24 \\citep{Wang07a}, M~17 \\citep{Broos07}, RCW~49 \\citep{Tsujimoto07}, and Rosette star forming region \\citep{Wang07b, Wang07c} YSCs. The three well-studied star forming regions in W3 are described and contrasted in \\S\\ref{sec:Xray}, and explanations for their origin are considered in \\S\\ref{sec:diversity}.  Section~\\ref{sec:W3Main.origin} considers in more detail the implications of the W3~Main results for astrophysical models of star cluster formation. ",
        "conclusions": "This paper introduces a new high-resolution X-ray mosaic of the W3 star forming complex, a Large Project of the {\\it Chandra X-ray Observatory}. A rich population of $\\sim 1300$ young stars is imaged and the three well-known regions of high-mass star formation are shown to have very different populations of low mass stars:  W3~Main is a large, rich, nearly spherical cluster; W3(OH) lies in an elongated group of sparse stellar clumps; and W3~North is an isolated O star without low-mass companions. Suggestions of these differences were inferred from earlier infrared studies, but they are more apparent here because the X-ray selection has the advantage of low contamination by the Galactic field population or diffuse interstellar emission, high penetration into molecular environments, and little bias towards stars with massive protoplanetary disks.  We emerge from this study with an improved view of star formation in the region. The W3(OH) structures are consistent with collect-and-collapse triggering process caused by by shocks from the older IC~1795 cluster,  as previously suggested. The W3~Main cluster, however, does not show the elongated and patchy structure of a recently triggered star cluster and appears to have formed in an earlier episode.  Its PMS population strongly resembles those seen in other {\\em Chandra} studies of massive star-forming regions such as those ionizing the Orion, M~17 and Rosette Nebulae. A major difference is that the individual H{\\sc II} regions in these other clusters have already merged into a large blister and dispersed their natal clouds.  In contrast, the W3~Main OB stars are very recently formed with small individual H{\\sc II} regions still embedded in a dense, clumpy molecular medium.  Star formation in W3 has proceeded in a prolonged fashion, and apparently with a time-dependent Initial Mass Function.  The OB stars exciting the hypercompact and ultracompact HII regions at the center of W3~Main formed more recently than the hundreds of X-ray emitting PMS stars distributed over several parsecs. W3~Main thus becomes a critical testbed for theories of rich cluster formation."
    },
    "0710/0710.5026_arXiv.txt": {
        "abstract": "Detection and study of gravitational waves from astrophysical sources is a major goal of current astrophysics. Ground-based laser-interferometer systems such as LIGO and VIRGO are sensitive to gravitational waves with frequencies of order 100 Hz, whereas space-based systems such as LISA are sensitive in the millihertz regime. Precise timing observations of a sample of millisecond pulsars widely distributed on the sky have the potential to detect gravitational waves at nanohertz frequencies. Potential sources of such waves include binary super-massive black holes in the cores of galaxies, relic radiation from the inflationary era and oscillations of cosmic strings. The Parkes Pulsar Timing Array (PPTA) is an implementation of such a system in which 20 millisecond pulsars have been observed using the Parkes radio telescope at three frequencies at intervals of two -- three weeks for more than two years. Analysis of these data has been used to limit the gravitational wave background in our Galaxy and to constrain some models for its generation. The data have also been used to investigate fluctuations in the interstellar and Solar-wind electron density and have the potential to investigate the stability of terrestrial time standards and the accuracy of solar-system ephemerides. ",
        "introduction": "The existence of gravitational radiation is a key prediction of relativistic theories of gravity. These waves propagate at the speed of light and are generated by the acceleration of massive bodies. Astrophysical sources of gravitational waves (GW) include relic radiation from the inflation era \\citep{tur97,gri05}, radiation from reconnection and oscillations of cosmic strings \\citep{dv05}, supernovae and formation of compact stars and black holes \\citep{bsr+05}, binary super-massive black holes in the cores of galaxies \\citep{jb03,wl03a,eins04}, coalescence of double-neutron-star binary systems \\citep{kkl+04a} and short-period X-ray binaries in our Galaxy \\citep{nyp04}. Some of these sources may be individually detectable, others combine to form an essentially isotropic and stochastic background of GW which permeates all of space.  Observations of orbital decay of double-neutron-star binary systems have provided irrefutable evidence that gravitational radiation exists and that its power is accurately described by Einstein's general theory of relativity \\citep{wt05,ksm+06}. However the signal expected at the Earth from any realistic source is exceedingly weak, with typical strain amplitudes of order $10^{-22}$ at frequencies of order 1 Hz. Despite considerable efforts over more than 40 years, up to now there has been no confirmed direct detection of GW. Initial efforts used the massive bar detectors pioneered by Joseph Weber \\citep{web69} but more recent detectors with higher sensitivity are based on laser interferometer systems, for example, the ground-based systems LIGO \\citep{aad+92} and VIRGO \\citep{gb02} and the proposed space interferometer LISA \\citep{dan00}.  The ground-based interferometers are sensitive to GW with frequencies in the range 10 -- 500 Hz, whereas LISA is sensitive to frequencies in the range 0.1 -- 100 mHz. Initial LIGO is now operating and has set limits on various sources \\citep[e.g.,][]{aaa+04}; higher sensitivity will be achieved with Advanced LIGO which is due for completion in 2011. The launch date for LISA is rather uncertain but is unlikely to be before 2017.  Pulsars, especially millisecond pulsars (MSPs) are incredibly precise clocks making possible many interesting applications. Of most interest to us here is the use of MSPs as GW detectors. GW passing over pulsars and over the Earth will modulate the received pulsar period; the net effect is the difference in the modulation at the two ends of the path \\citep{det79}. Pulsar timing experiments measure variations of pulse phase relative to model predictions. They are therefore most sensitive to long-period GW with periods comparable to the data span, typically several years, which corresponds to frequencies in the nanoHertz regime. Even for these long periods, the expected timing residuals are very small. Simulations using {\\sc tempo2} \\citep{hem06,hjl+07} show that a binary system consisting of two $10^9\\;\\Msun$ black holes with a 4-yr orbital period in a galaxy at redshift 0.5 will produce a timing residual of amplitude just 1.5 ns. Although such a signal would be very difficult to detect with current technology, expected levels of the stochastic GW background from binary super-massive black holes in galaxies are considerably higher with millions of galaxies throughout the Universe contributing. Predicted levels of the GW background from other sources such as the inflation era and cosmic strings, while much more uncertain, are at comparable levels and are potentially detectable.  Other sources of ``noise'' exist in pulsar timing data and if we wish to detect GW using pulsar timing we have to be able to separate these different effects. With timing observations of just one or even a few pulsars, upper limits may be set but a positive detection is not possible. However, a large sample of pulsars widely distributed on the sky --- a pulsar timing array --- can in principle {\\em detect} GW with frequencies in the nanoHertz range. ",
        "conclusions": "Pulsar timing arrays exploit the remarkable stability of MSP periods to enable investigation of a range of phenomena. Direct detection of gravitational waves from astrophysical sources is a major goal of current astrophysics and pulsar timing arrays have the potential to achieve this goal. They are sensitive to GW at frequencies of a few nanoHertz, complementing ground-based and space-based laser interferometer systems which are sensitive at much higher frequencies. PTA systems also have the potential to establish a ``pulsar timescale'' which is more stable than the best terrestrial timescales over intervals of several years or more and to detect errors or omissions in models of Solar-system dynamics, for example, the existence of currently unknown trans-Neptunian objects.  The Parkes Pulsar Timing Array (PPTA) project is using the Parkes 64-m radio telescope to time a sample of 20 millisecond pulsars at three frequencies every 2 -- 3 weeks. Observations commenced in early 2005, so we now have over two years of timing data. Sub-microsecond timing residuals have been achieved on about half the sample but we still need to improve timing precisions by a factor of a few in order to have a realistic chance of detecting the stochastic GW background. New instrumentation and other improvements will help us to achieve that goal. We are also actively seeking international collaborations with other timing array projects to increase the sky coverage and the density of observations. Already, limits on the GW background are starting to limit some inflation-era and cosmic string models. There seems little doubt that the proposed Square Kilometer Array radio telescope will be able to not only detect GW but to also study the properties of the GW sources in some detail, opening up a new era in astrophysics.   \\begin{theacknowledgments} The Parkes Pulsar Timing Array project has a great team of people based at the collaborating institutions. They contribute in different ways to helping us achieve our goals and I thank them all for their efforts. The Parkes radio telescope is part of the Australia Telescope which is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO. \\end{theacknowledgments}"
    },
    "0710/0710.1480_arXiv.txt": {
        "abstract": "{} {By reevaluating a 13-month stretch of Ulysses SWICS H pickup ion measurements near 5~AU close to the ecliptic right after the previous solar minimum, this paper presents a determination of the neutral interstellar H density at the solar wind termination shock and implications for the density and ionization degree of hydrogen in the LIC.} {The density of neutral interstellar hydrogen at the termination shock was determined from the local pickup ion production rate as obtained close to the cut-off in the distribution function at aphelion of Ulysses. As shown in an analytical treatment for the upwind axis and through kinetic modeling of the pickup ion production rate at the observer location, with variations in the ionization rate, radiation pressure, and the modeling of the particle behavior, this analysis turns out to be very robust against uncertainties in these parameters and the modeling. } {Analysis using current heliospheric parameters yields the H density at the termination shock equal to $0.087\\pm0.022$~cm$^{-3}$, including observational and modeling uncertainties. } {} ",
        "introduction": "Neutral interstellar gas of the local interstellar cloud (LIC) penetrates into the inner heliosphere as a neutral wind due to the relative motion between the Sun and the LIC. Apparently, the Sun is found near the boundary of a warm, relatively dilute cloud of interstellar gas, possibly with a significant gradient in the ionization fraction of H and He \\citep[e.g.][]{cheng_bruhweiler:90a, wolff_etal:99, slavin_frisch:02} within a very structured surrounding \\citep[e.g. reviews by][]{cox_reynolds:87a, frisch:95a}. In a companion paper within this special section, Frisch and Slavin (2008) lay out how the physical parameters and composition of the LIC at the location of the Sun, as derived from in-situ observations and from absorption line measurements, constrain the ionization state and radiation environment of the LIC. In situ observations of the two main constituents of the LIC, H and He, have been obtained with increasing accuracy, starting with the analysis of backscattered solar Lyman-$\\alpha$ intensity sky maps \\citep{bertaux_blamont:71, thomas_krassa:71} for H as well as with rocket-borne \\citep{paresce_etal:74a} and satellite-borne \\citep{weller_meier:74} observations of interstellar He using the solar He I 58.4 nm line. The optical diagnostics was followed by discovery of pickup ions for He \\citep{mobius_etal:85a} and for H \\citep{gloeckler_etal:92} and finally by direct neutral He observations \\citep{witte_etal:93}.  Such in-situ diagnostics, even at 1~AU, is made possible by the neutral gas flow deep into the inner heliosphere. Through the interplay between this wind, the ionization of the neutrals upon their approach to the Sun, and the Sun's gravitational field (distinctly modified by radiation pressure for H) a characteristic flow pattern and density structure is formed, with a cavity close to the Sun and gravitational focusing on the downwind side (for all species except H). The basic understanding of the related heliosphere -- LIC interaction has been summarized in early reviews by \\citet{axford:72, fahr:74, holzer:77, thomas:78}. While He provides us with almost completely unbiased information about the physical parameters of the LIC since it enters the heliosphere unimpeded, the abundance of H and O, along with other species, is significantly depleted, their speed decreased, and their temperature increased through charge exchange in the heliospheric interface \\citep{fahr:91, rucinski_etal:93a, izmodenov_etal:99a, mueller_etal:00, izmodenov_etal:04a}. A consolidation of the physical parameters of interstellar He, including the flow velocity vector relative to the Sun, as determined from neutral gas, pickup ion, and UV backscattering observations, was achieved through the effort of an ISSI Team \\citep[][and references therein]{mobius_etal:04a}, thus leading to a benchmark for the physical parameters of the LIC. This paper is part of a follow-up effort within an ISSI Team to also consolidate the determination of the LIC H density. The determination of the H density in the LIC proper not only involves a measurement inside the heliosphere, but is also dependent on the filtration of H in the heliospheric boundary. Therefore, consolidating the observational results concentrates on the determination of the H density at the termination shock, which still requires taking into account of the dynamics of the flow into the inner heliosphere as well as ionization effects. This paper deals with the determination of the H density from the pickup ion observations made with Ulysses SWICS. The effort to obtain the density from mass-loading of the solar wind by H pickup ions and its resulting slowdown at large distances from the Sun is described in the paper by \\citet[this volume]{richardson_etal:08a}, while \\citet[this volume]{pryor_etal:08a} discuss a determination of the H density based on the reduction of the modulation of the UV backscatter signal with distance from the Sun. To illustrate the state of the model-dependent H density value in the LIC \\citet[this volume]{mueller_etal:08a} compare different global models of the heliosphere and their results for the distances of the key boundary structures and the filtration factor, which also connects the inner heliosphere observations to the ionization state of the LIC, discussed by \\citet[this volume]{slavin_frisch:08a}.  In their previous work \\citet{gloeckler_geiss:01b} used pickup ion fluxes as observed at 5~AU with Ulysses SWICS, the charge exchange rates from SWOOPS, and a Vasyliunas \\& Siscoe distribution function to deduce the local neutral H density; they then used a hot interstellar gas model with the ionization rate significantly modified by electron ionization to deduce the density at the termination shock. In the present paper we use the same data set, but follow a complementary approach.  After discussing previous derivations of the local neutral gas density and its extrapolation to the termination shock at the beginning of section 2 we present an alternative approach. We make use of the fact that the ionization rate in the pickup ion production appears both as production rate of PUI and as loss rate of the parent neutral gas population and any variations balance close to the aphelion of Ulysses. As a consequence, the PUI production rate is almost exactly proportional to the H density at the termination shock. In the same section we illustrate this behavior in a simplified analytical model that applies to the upwind region.  In section 3 we simulate the local PUI production rate, starting with the density in the interstellar medium, compare it with the observations, and confirm the robustness of this approach by varying the parameters. We start with Monte-Carlo simulations of the flow through the heliospheric interface for two different LIC parameter sets, which result in two different H densities at the nose of the termination shock. In a second step, we hand these results over to a 3D time dependent test-particle code to calculate the H densities and H$^+$ PUI production rates at Ulysses during the observation interval, while accurately taking into account losses and radiation pressure along the trajectories of interstellar gas. We find the density at the termination shock and the LIC parameters that fit the observations best by interpolating between the two initial models. We study the response of the resulting PUI production rates to variations in the ionization rate, radiation pressure, and details in the modeling. In section 4, we present the results and show that our method is very robust against uncertainties of these parameters, and, in fact, against details of the simulations. \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[width=8cm]{aanda07a2_gr3.eps}}  \\caption{The interface correction: ratio of results of the Moscow Monte Carlo model for the geometry of Ulysses H PUI observations to the results of a sum of classical hot models (two populations), evaluated for identical parameters of the gas at the nose of the termination shock and identical ionization rate and radiation pressure as used in the MC model. The dotted horizontal line marks the heliospheric correction value equal to 1, the shaded area corresponds to the range of Ulysses heliocentric distances during the observations.} \\label{aa} \\end{figure} \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[width=8cm]{aanda07a3_gr1.eps}}  \\caption{Change of Ulysses position during the observations. The horizontal axis shows the heliocentric distance, the left-hand vertical scale heliolatitude, and the right-hand vertical scale ecliptic latitude. The times are indicated at the plot. Ecliptic longitude varied from $153.6\\degr$ at the beginning of observations interval to $157.5\\degr$ at the end, which corresponds to a change from $78.4\\degr$ to $81.8\\degr$ in heliolongitude. } \\label{obs1} \\end{figure} \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[width=8cm]{puiProdProf.eps}}  \\caption{Absolute production rates of H PUI as a function of heliocentric distances, obtained from the observed PUI spectrum during the interval discussed in the paper. Shown are the rates after normalizing the spectrum to the production rate at 1~AU and then multiplying by $1/r^2$. } \\label{obs2} \\end{figure} ",
        "conclusions": "We have used the accumulation of the H$^+$ pickup ion production rate from SWICS/Ulysses over a $\\sim 13$~month period in 1997 -- 1998 at the Ulysses passage through the solar equator plane to infer the interstellar H density at the termination shock. By extensive simulations we demonstrated that the H$^+$ PUI production rate in this location of the heliosphere is only weakly dependent on the values of solar radiation pressure and neutral H ionization rate, but sensitively depend on the density at the termination shock. We have found that the H density at the termination shock inferred from these pickup ion production rates is very robust against any variations in the ionization rate, radiation pressure, and the actual modeling approaches for the density distribution in the inner heliosphere.  In the present analysis we have, for the first time, included explicitly both the observational uncertainty and the modeling uncertainties. While in our approach the modeling uncertainties are minimized to a few percent, a larger uncertainty is incurred for the observation because absolute flux values are used, which results in an uncertainty of the obtained termination shock density equal to $\\pm 25$\\%. In the previous approach by \\citet{gloeckler_geiss:01b} the observational uncertainty was minimized by making use of the ratio of the pickup ion and solar wind flux, but the $\\sim 10$\\% uncertainty quoted in \\citet{izmodenov_etal:03a} does not include a range of values for the ionization rate and radiation pressure and the uncertainty of the geometric factor of the instrument. But, as demonstrated in the previous sections, the density at the termination shock scales linearly with the observed pickup ion production rate, which is directly related to the pickup ion flux and/or distribution function. Hence, any observational uncertainties will transfer linearly into the resulting densities. Since \\citet{izmodenov_etal:03a} started from the local neutral gas density at Ulysses, any uncertainty in the ionization rate will appear approximately linearly in the extrapolated density at the termination shock.  The determination of the PUI production rate at Ulysses near the aphelion, on which our derivation of $n_{\\mathrm{H,TS}}$ is based, is not entirely model-free. Although in the present determination of the PUI production rate at Ulysses a simple hot model was used for forward-modeling of the PUI distribution function, our method is robust against simplifications inherent to that kind of modeling because it uses a quantity (i.e., the production rate of PUI at Ulysses) which weakly depends on details of such modeling.  The determination of the H density at the termination shock by \\citet{gloeckler_etal:08a}, equal to $0.080\\pm0.008$~cm$^{-3}$, is free of the uncertainty of the geometric factor of the instrument, but is subject to a combination of uncertainties in the determination of the He density and that of the He abundance relative to H from the Voyager LECP observations, including uncertainties in the ratios of the production rates of these species. Nevertheless, after including of all uncertainties, all three approaches (i.e. the present one and those from \\citet{gloeckler_geiss:01b} and \\citet{gloeckler_etal:08a}) should be read with a similar uncertainty band. The density value presented here agrees very well with the new determination by \\citet{gloeckler_etal:08a}. Although our value still agrees with the previous determination of the H density from SWICS pickup ion observations by \\citet{gloeckler_geiss:01b} within their mutual uncertainty bands, the combination of the two new results suggest a somewhat lower density than 0.1~cm$^{-3}$.  Our results also agree comfortably with the TS density values found from the analysis of the heliospheric Lyman-$\\alpha$ glow \\citep[this volume]{pryor_etal:08a} and from the solar wind slowdown \\citep[this volume]{richardson_etal:08a} within the uncertainty bands. It should be noted here that the coupling between the neutral and ionized component of the interstellar medium between the bow shock and the heliopause appears to be somewhat stronger than suggested previously. This is a consequence of an updated relation for the energy dependence of the charge exchange cross-section between protons and H atoms \\citep{lindsay_stebbings:05a}.  The parameters of the interstellar gas in front of the heliospheric bow shock, as assumed in simulation (1), also seem to be robust, as shown by \\citet[this volume]{mueller_etal:08a}, who discussed the present status of the modeling of heliospheric interface and showed that differences in the filtration rate returned by different models of the heliosphere evaluated with identical initial parameters are about 15\\% and this result can be adopted as the uncertainty of the H density in the CHISM.  \\begin{appendix}"
    },
    "0710/0710.1952_arXiv.txt": {
        "abstract": "We present a Markov Chain Monte Carlo global analysis of neutrino parameters using both cosmological and experimental data. Results are presented for the combination of all presently available data from oscillation experiments, cosmology, and neutrinoless double beta decay. In addition we explicitly study the interplay between cosmological, tritium decay and neutrinoless double beta decay data in determining the neutrino mass parameters. We furthermore discuss how the inference of non-neutrino cosmological parameters can benefit from future neutrino mass experiments such as the KATRIN tritium decay experiment or neutrinoless double beta decay experiments. ",
        "introduction": "The question of neutrino mass is one of the most profound in modern particle physics. Most plausible models of neutrino mass solve the puzzle of why neutrino masses are so small by introducing a new scale at high energy, and precision studies of neutrino physics therefore hold the potential to investigate physics at scales beyond those reachable in current accelerator experiments. They also make the study of the possible Majorana nature of neutrinos possible (see \\cite{Mohapatra:2004ht,Mohapatra:2004vr,Mohapatra:2006gs} for a thorough discussion of this). While the neutrino mass differences have now been measured at about 10\\% precision by oscillation experiments (see e.g.\\ \\cite{Maltoni:2004ei,fogli}) the absolute mass scale remains unknown and inaccessible to oscillation experiments.  There are, however, several possible paths to measuring the absolute neutrino mass. The kinematical effect of neutrino mass can be probed either via its effect on the beta decay spectrum or via its effect on cosmological structure formation. If neutrinos are Majorana particles a different possibility is to search for neutrinoless double beta decay because the transition probability for this process is proportional to the neutrino mass squared.  In the past year there have been several papers discussing how to unify the data analysis for the various approaches \\cite{fogli,Host:2007wh}. This is a non-trivial issue, given that completely different physics is involved and that the three probes are actually sensitive to three distinct observables.  Here we present a new Markov Chain Monte Carlo global analysis of neutrino parameters using both cosmological and experimental data. The analysis software is based on the CosmoMC Markov Chain Monte Carlo (MCMC) package for cosmological parameter estimation \\cite{Lewis:2002ah,cosmomc}, appropriately modified to incorporate all parameters related to neutrino physics. This approach uses Bayesian inference instead of the frequentist method commonly used in particle physics. The approach is somewhat similar to the MCMC technique developed in \\cite{trotta} to constrain MSSM parameters. However, a key difference is that here we keep the full cosmological parameter estimation which allows for a closer study of the interplay between neutrino data and cosmological parameter estimation.  In Section II we describe the methodology used and in Section III we present the main results for various different assumptions about present and future data, as well as different parameter spaces. Finally we present our conclusions in Section IV. ",
        "conclusions": "A detailed neutrino parameter estimation study has been carried out using the Markov Chain Monte Carlo technique with the goal of unifying the various techniques for measuring the absolute neutrino mass scale. The MCMC technique is extremely powerful in this regard and allows for a very fast scanning many-dimensional likelihood spaces. In the concrete example here we have used 8 parameters describing the properties of light, active Majorana neutrinos, and 6 further parameters which specify the cosmology.  We find that for present data the combination of cosmological data with the upper limit on $m_{\\beta \\beta}$ from Heidelberg-Moscow slightly improves the existing cosmological bound on the sum of neutrino masses.  More interestingly we have studied the interplay between various future constraints from cosmology, tritium decay and neutrinoless double beta decay. If all probes come up with a negative result the addition of data sets does not yield any radically new information. However, we have also studied an example in which the upcoming KATRIN and GERDA experiments are both assumed to provide tentative evidence for neutrino mass. In this case the combination of all three types of data allows for a much stronger constraint on neutrino properties than otherwise allowed.  Finally we have also studied how experimental data from tritium decay or neutrinoless double beta decay can help in cosmological parameter estimation, particularly concerning the dark energy equation of state.  It should be noted that in the present analysis only presently available cosmological data has been used. In the same time frame as KATRIN and GERDA new cosmological data will become available and is likely to improve the cosmological neutrino mass bound significantly (see \\cite{Xia:2007gz,Gratton:2007tb,Hannestad:2007cp,Perotto:2006rj,Takada:2005si,% Lesgourgues:2005yv,Wang:2005vr,Song:2003gg,Hannestad:2002cn} for a non-exhaustive list). In the somewhat longer term cosmological constraints can be potentially be pushed below 0.1 eV sensitivity to $\\sum m_\\nu$. At the same time neutrinoless double beta decay experiments will have equally improved sensitivity and it will very likely be possible to determine the absolute neutrino mass as well as the nature of the mass hierarchy.  In conclusion, the combination of cosmological data with experimental neutrino data in a global analysis will be extremely useful in the future, when more precise experimental data becomes available."
    },
    "0710/0710.3166_arXiv.txt": {
        "abstract": "Despite intense scrutiny, the progenitor system(s) that gives rise to Type Ia supernovae remains unknown.  The favored theory invokes a carbon-oxygen white dwarf accreting hydrogen-rich material from a close companion until a thermonuclear runaway ensues that incinerates the white dwarf.  However, simulations resulting from this single-degenerate, binary channel demand the presence of low-velocity H$\\alpha$ emission in spectra taken during the late nebular phase, since a portion of the companion's envelope becomes entrained in the ejecta.  This hydrogen has never been detected, but has only rarely been sought. Here we present results from a campaign to obtain deep, nebular-phase spectroscopy of nearby Type Ia supernovae, and include multi-epoch observations of two events: SN~2005am (slightly subluminous) and SN~2005cf (normally bright).  No H$\\alpha$ emission is detected in the spectra of either object. An upper limit of $0.01\\ M_\\odot$ of solar abundance material in the ejecta is established from the models of \\citet{Mattila05} which, when coupled with the mass-stripping simulations of \\citet{Marietta00} and \\citet{Meng07}, effectively rules out progenitor systems for these supernovae with secondaries close enough to the white dwarf to be experiencing Roche lobe overflow at the time of explosion.  Alternative explanations for the absence of H$\\alpha$ emission, along with suggestions for future investigations necessary to confidently exclude them as possibilities, are critically evaluated. ",
        "introduction": "\\label{sec:1}  Ever since the currently favored single-degenerate, binary channel was proposed as the progenitor system for Type Ia supernovae \\citep{Whelan73}, all models of the impact of the exploded white dwarf (WD) on the secondary star have indicated that significant amounts of solar-abundance material, stripped from the secondary's envelope, become entrained in the ejecta \\citep{Wheeler75,Fryxell81,Taam84,Chugai86,Livne92,Marietta00,Meng07}. Observational evidence for this material, however, still eludes us \\citep{Mattila05}, and serves as one reminder among many that we still lack direct observational proof for the single-degenerate scenario \\citep{Branch95,Livio01}.  The most detailed theoretical investigation of the expected amount and distribution of stripped material within a young Type Ia supernova (SN~Ia) remnant is that of \\citet{Marietta00}, who studied the problem with two-dimensional numerical simulations.  Four basic progenitor systems were investigated, including three with secondaries (a main-sequence, subgiant, or red giant star) close enough for mass-transfer to occur through Roche lobe overflow (RLOF), and one containing a secondary (a red giant) donating material through a strong stellar wind --- the symbiotic case.  All secondaries were given masses of $\\sim 1\\ M_\\odot$ at the time of the explosion and were placed either just within (the RLOF cases), or just beyond (the symbiotic case), the limiting distance from the WD within which RLOF can occur (i.e., $a/R = 3$, where $a$ is the orbital separation in units of the secondary star's radius, $R$; see \\citealt{Eggleton83}).  Three additional systems, in which a main-sequence secondary was placed too far away to experience RLOF (thus rendering it an unlikely progenitor system) were also included to establish the scaling between orbital separation and amount of stripped material.  The numerical results of the \\citet{Marietta00} study confirmed predictions from earlier analytic work \\citep[e.g.,][]{Wheeler75,Chugai86} that substantial material is indeed stripped, with the amount ranging from a minimum of $0.15\\ M_\\odot$ for a close ($a/R = 3$) main-sequence secondary, up to nearly the entire envelope ($\\sim 0.5\\ M_\\odot$) for a similarly placed red giant. Increasing the orbital separation beyond the RLOF limit in the \\citet{Marietta00} simulations resulted in a dramatic decrease in the amount of stripped material: For $a/R = 12$, a main-sequence secondary loses only $0.0018\\ M_\\odot$.  However, since such systems lack an efficient mechanism for mass transfer, they are not considered to be viable SN~Ia progenitors.  A red giant secondary placed at a similarly distant location, however, could potentially donate mass through a strong stellar wind, making the symbiotic case a possibility if large orbital separations are required \\citep[e.g.,][]{Munari92}.  As discussed by \\citet{Meng07}, one shortcoming of the \\citet{Marietta00} study is their use of standard solar-model stars for the companion, rather than companions whose structures have been appropriately modified due to having evolved in a binary system \\citep[e.g.,][]{Eggleton73}.  To investigate the effect this simplification might have had on the results, \\citet{Meng07} use an analytic model to estimate the amount of mass expected to be stripped from a variety of evolved secondaries.  (Their analytic approach was first tested using the unevolved secondaries used by \\citealt{Marietta00}, and was demonstrated to approximate the results obtained numerically.)  The result is that the quantity of material expected to be stripped from evolved secondaries is considerably lower than that predicted for standard solar-model companions. In fact, \\citet{Meng07} find that the minimum value for systems experiencing RLOF is diminished from $0.15\\ M_\\odot$ to only $0.035\\ M_\\odot$.  The reduction arises primarily from the pre-explosion mass-loss producing a more compact companion star whose material is more difficult to strip than it is in the unevolved case.  \\citet{Meng07} stress, however, that their new values are really lower bounds on the amount of stripped material, since their analytic approach does not consider the thermal energy imparted by the ejecta to the companion envelope, which likely serves to heat and vaporize a portion of it and thereby increase the amount of stripped material \\citep[e.g.,][]{Fryxell81,Mattila05}.  Thus, $0.035\\ M_\\odot$ serves as a conservative lower bound on the expected amount of stripped material resulting from their models.  The typical velocity of the stripped material is found in all studies to be far slower than the $\\sim 10,000$~\\kms\\ velocity that characterizes the bulk of the iron-rich ejecta \\citep{Chugai86,Marietta00,Meng07}.  This has the effect of placing it almost entirely in the central region of the supernova remnant, with the majority of it predicted to be moving with a velocity of under $1,000$~\\kms\\ \\citep{Marietta00}.  The stripped material is largely confined to the downstream region behind the companion star, where it contaminates a solid angle that ranges from $66^\\circ$ for the main-sequence companion to $115^\\circ$ for the red-giant companion \\citep{Marietta00}.  The low velocity of the stripped material renders it undetectable when the faster-moving, iron-rich ejecta are optically thick.  However, detailed radiative transfer calculations performed by \\cite{Mattila05} predict that it should become visible at late times, when the outer ejecta have thinned out and become transparent enough to reveal the slower-moving gas in the central regions.  The most prominent expected spectral signature of the companion star's stripped material is narrow H$\\alpha$ emission in nebular spectra \\citep{Mattila05}, taken more than $\\sim 250$ days after maximum light.  The H$\\alpha$ emission should be present within $\\pm 1,000$~\\kms\\ {\\rm of\\ } $\\lambda_0 = 6563$ \\AA\\ but, due to the expected asymmetry in the distribution of the solar-abundance material, could present an H$\\alpha$ profile ranging from a very narrow spike to a broader emission line in the observed spectrum. Obtaining spectra with high enough resolution to resolve fairly narrow lines is thus a useful component of a targeted search for this H$\\alpha$.  To date, only a few nebular SN~Ia spectra have been obtained and H$\\alpha$ has never been detected, although the majority of the spectra lack the spectral resolution and signal-to-noise ratio to place interesting constraints on the companion.  The best limits, by far, come from the recent study by \\cite{Mattila05}, which constrains hydrogen-rich material in SN~2001el to be $\\lesssi 0.03\\ M_\\odot$ from a low-resolution ($\\sim 700$~\\kms) spectrum obtained 398 days past maximum light.  In an effort to expand and improve on earlier work, both in terms of the number of objects studied as well as the resolution, sensitivity, and temporal coverage of the spectra, we have initiated a program to obtain deep, moderate resolution ($\\lesssim 150$~\\kms, or $\\sim 3$~\\AA, at H$\\alpha$), late-time spectra of SNe~Ia at multiple epochs using the Keck and Gemini telescopes.  The first phase of this project has garnered data on two objects: SN~2005am, a slightly subluminous \\citep{Li06} event and SN~2005cf, an SN~Ia of normal brightness \\citep{Pastorello07}.  We present and analyze our observations in \\S~\\ref{sec:2}, discuss the results in \\S~\\ref{sec:3}, and conclude in \\S~\\ref{sec:4}. ",
        "conclusions": "\\label{sec:4}  We obtained five deep, moderate-resolution, nebular-phase spectra of two SNe~Ia (SN~2005am and SN~2005cf) in order to search for narrow H$\\alpha$ emission that would betray the existence of material stripped from the envelope of a mass-donating stellar companion to the exploding WD.  No such emission is detected in either object at any epoch.  From the models of \\citet{Mattila05}, we establish upper limits of $0.01\\ M_\\odot$ of solar abundance material in the inner ejecta of both objects, which are the tightest constraints yet established by such studies.  Our non-detections of H$\\alpha$, coupled with the mass-stripping results of \\citet{Marietta00} and \\citet{Meng07}, rule out all hydrogen-donating companions close enough to the WD to have been experiencing RLOF at the time of explosion for these events.  Additional theoretical work is needed in several areas to buttress this conclusion, including most critically verification of the transparency of the outer, more rapidly moving ejecta that could potentially block H$\\alpha$ photons from escaping from the inner region. Bearing this caveat in mind, we propose that symbiotics are, at this time, the most likely progenitor class that remains consistent with these data.  Definitive proof of the identity of the progenitor system(s) that gives rise to SNe~Ia remains elusive, and it must be admitted that our conclusion, which is based on the {\\it lack} of a detection, is not as satisfying as one based {\\it on} a detection.  Should future modeling efforts prove unable to ``hide the hydrogen'' for even widely separated binaries, then the continued viability of the single-degenerate, hydrogen-donating progenitor will require that H$\\alpha$, no matter how weak, must ultimately be detected."
    },
    "0710/0710.3399_arXiv.txt": {
        "abstract": "There is currently no explanation of why the corona has the temperature and density it has.  We present a model which explains how the dynamics of magnetic reconnection regulates the conditions in the corona.  A bifurcation in magnetic reconnection at a critical state enforces an upper bound on the coronal temperature for a given density.  We present observational evidence from 107 flares in 37 sun-like stars that stellar coronae are near this critical state.  The model may be important to self-organized criticality models of the solar corona. ",
        "introduction": "Dynamics in the solar corona takes on a wide range of forms.  On one hand, the corona is the setting for the most violent eruptions in the solar system: solar flares and coronal mass ejections \\citep{Aschwanden01}.  On the other, coronal heating makes the corona almost a thousand times hotter than the photosphere, even in the quiet sun \\citep{Klimchuk06}.  \\citet{Parker83,Parker88} unified these two phenomena by proposing that micro- and nano-flares, less energetic cousins of eruptive flares, heat the corona.  This model gained credence from studies showing that solar flares exhibit power law statistics \\citep{Lin84,Dennis85,Crosby93,Feldman97,Wheatland00, Nita02,Paczuski05} over a wide range of scales for many quantities. [See \\citet{Charbonneau01} for a review.]  In addition, stellar flares have similar light curves to solar flares \\citep{Gershberg05} and also exhibit power law statistics \\citep{Collura88,Shakhovskaya89, Audard00}, suggesting that the physics of the solar corona is generic to sun-like stars.  Coronal dynamics remains an active research area \\citep{Hudson91, Georgoulis98,Shibata02,Hughes03}.  Details of the eruption process including how magnetic energy is stored, how eruptions onset, and how the stored energy is converted to other forms are still open questions.  In addition, while micro- and nano-flares are believed to be a major contributor to coronal heating, the authors know of no theory which explains why the coronal temperature and density have the values they have, as opposed to larger or smaller values.  In this paper, we propose that the condition of the corona is regulated by magnetic reconnection \\citep{Cassak06b}, a dynamical process which converts magnetic energy into kinetic energy and heat and energizes particles.  Magnetic energy is stored during collisional (slow) reconnection, which has been shown to drive the coronal plasma toward lower collisionality \\citep{Cassak06}.  If the plasma becomes marginally collisionless, a bifurcation in the underlying dynamics of reconnection occurs \\citep{Cassak07c}.  This bifurcation, which occurs when two length scales $\\delta_{SP}$ and $\\rho_{i}$ (to be defined below) are comparable, catastrophically initiates fast (Hall) reconnection, releasing the stored energy in the form of an eruption. The condition of marginal collisionality, therefore, sets an upper bound on how hot the coronal plasma can be for a given density.  The continual driving toward lower collisionality of the pre-flare corona by slow reconnection enforces the self-organization of the corona to a state of marginal collisionality where $\\delta_{SP} \\sim \\rho_{i}$. We present finer details of this process below.  Then, we perform the first observational test of this model using a large sample of data from stellar flares on sun-like stars.  We find that $\\delta_{SP}$ and $\\rho_{i}$ are comparable for every event in the sample, indicating that stellar coronae do self-organize into a marginally collisional state. ",
        "conclusions": "\\label{sec-disc}  The data analyzed in this paper pertain to flares in sun-like stars, but the underlying dynamics of reconnection is general.  Our model applies equally well to micro- and nano-flares in the quiet corona. Using values for the quiet sun of $T \\sim$ 1 MK, $n \\sim 10^{9}$ ${\\rm cm}^{-3}$, $B \\sim $ 5 G, and $L \\sim 10^{10}$ cm, we find $\\delta_{SP} \\sim 770$ cm and $d_{i} \\sim 720$ cm, in agreement with the model.  The present result may have important implications for self-organized criticality (SOC) models of the solar corona.  SOC occurs in driven, dissipative systems when the system is driven to a critical state where it undergoes a major reconfiguration \\citep{Bak87}.  SOC leads to power law statistics, which encouraged \\citet{Lu91} to propose the corona undergoes SOC.  Subsequent studies of SOC in the corona exist \\citep{Lu93,Vlahos95,Longcope00,Isliker01}, but a firm physical foundation of the mechanism for self-driving and the physical condition setting the critical state is often traded for the ease of performing cellular automaton simulations [see \\citet{Charbonneau01} for a review].  The present result provides a physical mechanism for self-driving (embedded Sweet-Parker reconnection) and the critical state (marginal collisionality), which may provide an avenue for developing quantitative predictions of SOC to compare with coronal observations.  An alternate mechanism \\citep{Uzdensky06,Uzdensky07,Uzdensky07b} for heating the solar corona uses a change in density to achieve self-regulation.  After an eruption, chromospheric evaporation increases the coronal density, decreasing the ion gyroradius [eq.~(\\ref{didef})] and making subsequent eruptions more difficult. The extent to which Uzdensky's and our mechanisms regulate coronal heating is an open question.  The present model assumes that Sweet-Parker scaling is appropriate for thin current sheets of large extent.  Long current sheets are known to fragment due to secondary instabilities, but the effect of this on the reconnection rate is unknown.  Verification of the present model would entail testing whether Sweet-Parker reconnection in extended current sheets remains much slower than Hall reconnection.  [See \\citet{Uzdensky07b} for further discussion of this point as well as other future research directions.]  The authors thank J.~F.~Drake, A.~Klimas, E.~Ott, S.~Owocki, P.~So and D.~Uzdensky for helpful conversations.  This work was supported in part by the Delaware Space Grant.  \\clearpage"
    },
    "0710/0710.5789_arXiv.txt": {
        "abstract": "Recent stellar evolutionary calculations of low-metallicity massive fast-rotating main-sequence stars yield iron cores at collapse endowed with high angular momentum. It is thought that high angular momentum and black hole formation are critical ingredients of the collapsar model of long-soft $\\gamma$-ray bursts (GRBs). Here, we present 2D multi-group, flux-limited-diffusion MHD simulations of the collapse, bounce, and immediate post-bounce phases of a 35-\\msun\\, collapsar-candidate model of Woosley \\& Heger. We find that, provided the magneto-rotational instability (MRI) operates in the differentially-rotating surface layers of the millisecond-period neutron star, a magnetically-driven explosion ensues during the proto-neutron star phase, in the form of a baryon-loaded non-relativistic jet, and that a black hole, central to the collapsar model, does not form. Paradoxically, and although much uncertainty surrounds stellar mass loss, angular momentum transport, magnetic fields, and the MRI, current models of chemically homogeneous evolution at low metallicity yield massive stars with iron cores that may have {\\it too much} angular momentum to avoid a magnetically-driven, hypernova-like, explosion in the immediate post-bounce phase. We surmise that fast rotation in the iron core may inhibit, rather than enable, collapsar formation, which requires a large angular momentum not in the core but {\\it above} it. Variations in the angular momentum distribution of massive stars at core collapse might explain both the diversity of Type Ic supernovae/hypernovae and their possible association with a GRB. A corollary might be that, rather than the progenitor mass, the angular momentum distribution, through its effect on magnetic field amplification, distinguishes these outcomes. ",
        "introduction": "There is mounting observational evidence for the association between long-soft $\\gamma$-ray Bursts (GRBs) and broad-lined Type Ic supernovae (SNe; see Woosley \\& Bloom 2006 for a review). Such hydrogen-deficient (and, perhaps, also helium-deficient) progenitors are compact and, if fast rotating in their % core at collapse, fulfill critical requirements for the formation of a collapsar (Woosley 1993). The engine that converts energy from long-term accretion of disk material onto the black-hole (BH) may power a relativistic jet in the excavated polar regions. The jet breaks out of the progenitor surface while equatorial accretion continues. Depending on the BH mass and the angular momentum budget in the collapsing envelope, this ``engine'' may operate for seconds, i.e. as long as typical long-soft GRBs. Accompanying this beamed relativistic polar jet might be a disk wind, fueled by neutrinos or MHD processes, which would explode the Wolf-Rayet envelope. This explosion and the radioactive $^{56}$Ni material produced might lead to very energetic, broad-lined, Type Ic SN of the hypernova variety (Iwamoto et al. 1998; MacFadyen \\& Woosley 1999; Hjorth et al. 2003; Stanek et al. 2003).  State-of-the-art radiation-hydrodynamics simulations including a sophisticated equation of state (EOS) and detailed neutrino transport (Buras et al. 2003; Burrows et al. 2006,2007a; Kitaura et al. 2006; Marek \\& Janka 2007; Mezzacappa et al. 2007) suggest that while the neutrino mechanism of supernova explosions may work for the lower-mass massive progenitors, it may not for the more massive progenitors, characterized by an ever higher post-bounce accretion rate onto the proto-neutron star (PNS). Burrows et al. (2006,2007a) have suggested that an acoustic mechanism will work for all slowly rotating progenitors that do not explode by other means within the first second after bounce. However, massive star cores endowed with a large angular momentum at the time of collapse should experience the magneto-rotational instability (MRI; Balbus \\& Hawley 1991; Akiyama et al. 2003; Pessah et al. 2006; Shibata et al. 2006; Etienne et al. 2006), with the potential to exponentially amplify weak initial fields on a rotation timescale. The saturation values of such fields are ultimately set by the free-energy of differential rotation available in the surface layers of the PNS (Ott et al. 2006), and can be large, i.e., on the order of 10$^{15}$\\,G at a radius of a few tens of kilometers. The corresponding magnetic stresses at the neutron star surface lead systematically to powerful jet-like explosions $\\sim$100\\,ms after bounce (see, e.g., Ardeljan et al. 2005; Yamada \\& Sawai 2004; Kotake et al. 2004; Sawai et al. 2005; Moiseenko et al. 2006; Obergaulinger et al. 2006; Burrows et al. 2007b, hereafter B07; Dessart et al. 2007).  In this letter, we investigate, in the context of the collapsar model, the potential implications of this magnetic explosion mechanism. Our study focuses on the immediate post-bounce phase, whose importance was emphasized by Wheeler et al. (2000,2002). This is in contrast to previous work which explored only the phase subsequent to BH formation (MacFadyen \\& Woosley 1999; Aloy et al. 2001; Zhang et al. 2003; Proga 2005). Indeed, two terms sometimes used in the collapsar context are ``failed SN'' (MacFadyen \\& Woosley 1999) and ``prompt BH formation'' (MacFadyen et al. 2001). Our analysis supports the idea that the conditions for the collapsar model, as stated so far, are also suitable for a magnetically-driven explosion in the immediate post-core-bounce PNS phase, and that BH formation may be so delayed for a range of putative progenitor models that it does not in fact occur\\footnote{In the present context, BH formation is never prompt, since it takes a finite time, on the order of seconds, for the PNS to accumulate the critical mass at which it experiences the gravitational instability. This is in contrast with super-massive stars, such as the progenitors of pair-instability SNe, which may form an apparent horizon during collapse and thus ``directly'' transition to a BH (Liu et al. 2007).}. In \\S2, we present radiation MHD simulations with the code VULCAN/2D (Livne et al. 2004,2007) of a collapsar-candidate model that support this thesis. In \\S3, we discuss the implications of our results for stellar evolutionary models that might lead to collapsars and/or hypernovae.  \\begin{deluxetable}{lcccccccc} \\tablewidth{0pt} \\tabletypesize{\\scriptsize} \\tablecaption{Properties of our two MHD-VULCAN/2D simulations of the 35OC collapsar model of WH06.\\label{tab_model}} \\tablehead{ \\colhead{}& \\colhead{$t_{\\rm end}$}& \\colhead{$t_0$}& \\colhead{M$_{10}$}& \\colhead{P$_{10}$}& \\colhead{E$_{\\rm expl}$}& \\colhead{\\edot$_{\\rm gas}$}& \\colhead{\\edot$_{\\rm \\vec{E} \\times \\vec{B}}$}& \\colhead{$v_{\\rm max}$}\\\\ \\colhead{}& \\colhead{ms}& \\colhead{ms}& \\colhead{\\mo}& \\colhead{ms}& \\colhead{B}& \\multicolumn{2}{c}{B\\,s$^{-1}$}& \\colhead{km\\,s$^{-1}$} } \\startdata M0 & 369 & ...  & 2.1 &  4 & 0.03 & 0.5 & 0.25 & 43,000 \\\\ M1 & 666 & 349  & 1.7 & 12 & 3.31 & 9.4 & 3.0  & 58,000 \\\\ \\enddata \\tablecomments{ $t_{\\rm end}$ gives the time at the end of each simulation, while $t_0$ is the time when the rate of polar mass ejection first overcomes equatorial mass accretion. All quoted quantities in the table correspond to the final time in each simulation, while times are given with respect to core bounce. M$_{10}$ (P$_{10}$) corresponds to the total baryonic mass (average rotation period) inside the 10$^{10}$\\,g\\,cm$^{-3}$ isodensity contour. \\edot$_{\\rm gas}$ (\\edot$_{\\rm \\vec{E} \\times \\vec{B}}$) is the Bernoulli (Poynting) power in the ejecta, obtained by integrating the corresponding flux over a shell with a radius of 500\\,km. [See text for additional information.] } \\end{deluxetable} ",
        "conclusions": "The potential for exponential growth on a rotational timescale of initial seed magnetic fields by the MRI (Shibata et al. 2006; Etienne et al. 2006), fueled by the free energy of core rotation, makes the initial angular momentum budget of the progenitor star the key parameter in determining the outcome during the immediate post-bounce phase (B07). A magnetically-driven, baryon-loaded, and non-relativistic, explosion is obtained for WH06's 35OC collapsar candidate model, evolved at low metallicity from a 35\\,\\msun\\, fast-rotating main-sequence star. The explosion occurs $\\sim$200\\,ms after bounce and reaches $\\sim$3\\,B $\\sim$400\\,ms later. After an initial accretion phase, the steadily decreasing PNS mass reaches only $\\sim$1.7\\,\\msun\\, at the end of the simulation, and, thus, the quasi-steady explosion we observe suggests that BH formation is unlikely to occur. Moreover, baryon contamination prevents the ejecta from becoming relativistic. Note that the production of a GRB in the collapsar context is contingent on the gravitational collapse of the PNS to a BH.  The recent stellar evolutionary calculations of Yoon \\& Langer (2005), WH06, and Meynet \\& Maeder (2007) of fast-rotating main-sequence objects at low metallicity systematically predict such fast-rotating cores at collapse. Starting from similar conditions for a 35-40\\,\\msun\\, star, but using different mass-loss ``recipes,'' they obtain very similar rotational profiles in the inner core. % Allowing for anisotropic mass loss (Meynet \\& Maeder 2007), a model of C. Georgy (2007, priv. comm.) suggests an even larger (by a factor of two) specific angular momentum in the inner 3\\,\\msun\\, at the end of silicon core burning. % Despite the agreement between these different evolutionary computations, the magnetically-driven explosion and the ``failed'' BH formation described here are conditional on the uncertain treatment of mass loss, angular momentum transport, and magnetic processes (Spruit 2002) during the pre-collapse evolution.  At very low metallicities, radiatively-driven winds of massive stars are inhibited by the lack of metals (Kudritzki 2002; Vink et al. 2001; Vink \\& de Koter 2005), whose optically-thick lines intercept radiation momentum (Castor et al. 1975). Recent revisions downward of mass-loss rates due to clumping (Owocki et al. 1988; Bouret et al. 2005; Fullerton et al. 2006) suggest, however, the potential importance of episodic outbursts, akin to the 1843 giant eruption of Eta Carina (Smith \\& Owocki 2006). The metallicity dependence of such phenomena is entirely unknown, mostly  because the fundamental cause of the outburst remains a mystery. While line driving seems excluded, continuum driving of a porous medium at super-Eddington luminosities has been proposed by Owocki et al. (2004) as an alternative. Finally, mass loss in fast-rotating, and sometimes critically-rotating (Townsend et al. 2004), envelopes is complicated by the effects associated with centrifugal support, surface oblateness, and gravity-darkening (Cranmer \\& Owocki 1995; Owocki et al. 1996), so that the mass-loss ``recipes'' used in stellar evolutionary models are not always substantiated by observational and theoretical evidence. At present, and in light of our simulations, it appears that chemically-homogeneous evolution of fast-rotating main-sequence massive stars at low metallicity systematically yields iron cores at collapse that may have {\\it too much} angular momentum, a property that prevents the formation of a collapsar. Uncertainties in the modeling of the pre-collapse evolution may result, however, in slower-rotating iron cores\\footnote{Note that the rotational energy $E_{\\rm rot}$ is a stiff function of angular velocity $w$, i.e., $E_{\\rm rot} \\propto w^2$.} and, thus, might inhibit an early magnetically-driven explosion in favor of black hole, and perhaps collapsar, formation.  \\begin{figure} \\includegraphics[width=8cm]{f2.ps} \\caption{Colormap of the entropy at 666\\,ms after bounce for model M1, overplotted with white iso-density contours (every decade downward from 10$^{10}$\\,g\\,cm$^{-3}$) and velocity vectors (length saturated to 15\\% of the width of the display and corresponding to a velocity of 30,000\\,km\\,s$^{-1}$.) } \\label{fig_still} \\end{figure}  We conclude that variations in the angular momentum distribution of pre-collapse massive stars may lead to different post-bounce scenarios. Non- or slowly-rotating progenitors may explode with weak/moderate energy ($\\sles$1\\,B) through a neutrino or an acoustic mechanism $\\sles$1\\,s after bounce, or may collapse to a BH. Objects with large angular momentum in the envelope, but little in the core, may proceed through the PNS phase, transition to a BH and form a collapsar with a GRB signature. Owing to the modest magnetic-field amplification above the PNS, a weak precursor polar jet may be launched, soon overtaken by a baryon-free, collimated relativistic jet. At the same time, the progenitor envelope  is exploded by a disk wind, resulting in a hypernova-like SN with a large luminosity (large $^{56}$Ni mass). Finally, and this is what we conclude here, objects with large angular momentum in the core may not transition to a BH. Instead, and fueled by core-rotation energy, a magnetically-driven baryon-loaded non-relativistic jet is obtained without any GRB signature. The explosion has the potential of reaching energies of a few B to 10\\,B, and for viewers along the poles of looking like a Type Ic hypernova-like SN with broad lines. For a viewer at lower latitudes, the delayed and less energetic explosion nearer the equator may look more like a standard Type Ic SN (H\\\"{o}flich et al. 1999). This volume-restricted jet-like explosion is dimmer, as the amount of processed $^{56}$Ni may be significantly less than the $\\sim$0.5\\,\\msun\\, obtained in the collapsar context (MacFadyen \\& Woosley 1999). Hence, magnetic processes during the post-bounce phase of fast-rotating iron cores offer a potential alternative to collapsar formation and long-soft GRBs by producing non-relativistic non-Poynting-flux-dominated baryon-loaded hypernova-like explosions without any GRB signature. Importantly, while our study narrows the range over which the collapsar model may exist, it also offers additional routes to explain the existence of GRB/SN-hypernova events like SN 1998bw (Woosley et al. 1999), and hypernova events like SN 2002ap without a GRB signature (Mazzali et al. 2002).  More generally, magnetic effects should naturally arise in the context of gravitational collapse and fast rotation. The resulting angular momentum of newly-formed BHs and magnetars, for example, would be reduced, perhaps considerably, by any prior magnetically-driven explosion, and, thus, may decrease the power of subsequent mass ejections from compact objects (see, e.g., Thompson et al 2004)."
    },
    "0710/0710.5740_arXiv.txt": {
        "abstract": "{} {The properties of the very high energy (VHE; E$>$100 GeV) $\\gamma$-ray emission from the high-frequency peaked BL\\,Lac PG\\,1553+113 are investigated.  An attempt is made to measure the currently unknown redshift of this object.} {VHE Observations of PG\\,1553+113 were made with the High Energy Stereoscopic System (HESS) in 2005 and 2006. H+K (1.45$-$2.45$\\mu$m) spectroscopy of PG\\,1553+113 was performed in March 2006 with SINFONI, an integral field spectrometer of the ESO Very Large Telescope (VLT) in Chile.} {A VHE signal, $\\sim$10 standard deviations, is detected by HESS during the 2 years of observations (24.8 hours live time). The integral flux above 300 GeV is $(4.6\\pm0.6_{\\rm stat}\\pm0.9_{\\rm syst}) \\times 10^{-12}$ cm$^{-2}$\\,s$^{-1}$, corresponding to $\\sim$3.4\\% of the flux from the Crab Nebula above the same threshold.  The time-averaged energy spectrum is measured from 225 GeV to $\\sim$1.3 TeV, and is characterized by a very soft power law (photon index of $\\Gamma = 4.5\\pm0.3_{\\rm stat}\\pm0.1_{\\rm syst}$). No evidence for any flux or spectral variations is found on any sampled time scale within the VHE data. The redshift of PG\\,1553+113 could not be determined. Indeed, even though the measured SINFONI spectrum is the most sensitive ever reported for this object at near infrared wavelengths, and the sensitivity is comparable to the best spectroscopy at other wavelengths, no absorption or emission lines were found in the H+K spectrum presented here.} {} ",
        "introduction": "Evidence for VHE ($>$100 GeV) $\\gamma$-ray emission from the active galactic nucleus (AGN) PG\\,1553+113 was first reported by the HESS collaboration (\\cite{HESS_discovery}) based on observations made in 2005. This evidence was later confirmed (\\cite{MAGIC_1553}) with MAGIC observations in 2005 and 2006. Similar to essentially all AGN detected at VHE energies, PG\\,1553+113 is classified as a high-frequency peaked BL\\,Lac \\cite{classification} and is therefore believed to possess the double-humped broad-band spectral energy distribution (SED) typical of blazars. The low-energy (i.e. from the radio to the X-ray regime) portion of the SED of PG\\,1553+113 is well-studied, including several simultaneous multi-wavelength observation campaigns (see, e.g., \\cite{Example_1553_MWL}). However, the only data in the high-energy hump are from HESS and MAGIC.  The measured VHE spectra are unusually soft (photon index $\\Gamma$=4.0$\\pm$0.6 and $\\Gamma$=4.2$\\pm$0.3, respectively) but the errors are large, clearly requiring improved measurements before detailed interpretation of the complete SED is possible.  Further complicating any SED interpretation is the absorption of VHE photons (\\cite{EBL_effect3,EBL_effect2}) by pair-production on the Extragalactic Background Light (EBL). This absorption, which is energy dependent and increases strongly with redshift, distorts the VHE energy spectra observed from distant objects. For a given redshift and a given EBL model, the effects of the latter on the observed spectrum can be reasonably accounted for during SED modeling. Unfortunately, the redshift of PG\\,1553+113 is currently unknown\\footnote{\\cite{redshift_question} demonstrate the catalog redshift of $z=0.36$ is incorrect.}. To date no emission or absorption lines have been measured from PG\\,1553+113 despite more than ten observation campaigns with optical instruments, including EMMI at the NTT and FORS2 with the 8-meter VLT telescopes (\\cite{Carangelo_03,no_lines}).  Lower limits of $z>0.09$ (\\cite{no_lines}) and $z>0.3$ (\\cite{Carangelo_03}) were determined from the lack of detected absorption/emission lines, implying that the effect of the EBL is large in the observed VHE data. The absence of absorption and emission lines suggests that the non-thermal component of the emission from PG\\,1553+113 is largely dominant over that of the host galaxy. This is consistent with the fact that, although some hints for a host galaxy have been suspected, no clear detection has yet been found, even in Hubble Space Telescope (HST) images of PG\\,1553+113 taken during the HST survey of 110 BL Lac objects \\cite{Hubble_image}.  Interestingly,  approximately 80\\% (88/110) of the BL\\,Lacs initially surveyed by HST have known redshifts, of which all 39 with $z<0.25$ and 21 of the 28 with $0.25<z<0.6$ have their hosts resolved, suggesting that the redshift of PG\\,1553+113 is indeed large. Recently, using a re-analysis of the HST snapshot survey of BL\\,Lacs, \\cite{z_extreme} claimed the dispersion of the absolute magnitude of BL\\,Lac hosts is sufficiently small that the measurement of the host-galaxy brightness allows a reliable estimate of their redshift. With the assumption that there is no strong evolution, these authors set a possible lower limit of $z>0.78$ for PG\\,1553+113. However, adding new STIS-HST data of $z>0.6$ BL\\,Lacs to the snapshot survey, \\cite{HST_new} found on the contrary that host galaxies of BL\\,Lacs evolve strongly, making any photometric redshift determination questionable. A conservative upper limit ($z<0.74$) was determined (\\cite{HESS_discovery}) from the photon spectrum measured by HESS.  The same limit was similarly determined from the MAGIC spectral measurement (\\cite{MAGIC_1553}). Using both the MAGIC and HESS data, a stronger upper limit of $z<0.42$ was later reported \\cite{Mazin_limit} based on the assumption that there is no break in the intrinsic VHE spectrum of the object.  The range of allowed redshift is clearly large enough that the significant effects of EBL absorption cannot be reliably removed from the observed VHE spectrum of PG\\,1553+113. This correspondingly makes modeling the high-energy portion of its SED unreliable. A clear detection of the object's redshift would dramatically improve the understanding of PG\\,1553+113.  Further, if it were found to be distant, its VHE spectrum could potentially provide strong constraints on the poorly-measured EBL (assuming a reasonable intrinsic spectrum) and contribute to establishing the VHE $\\gamma$-ray horizon.  Results from 17.2 hours of new HESS observations of PG\\,1553+113 in 2006 are reported here.  In addition, a re-analysis of the previously published 2005 HESS data (7.6 hours), with an improved calibration of the absolute energy scale of the detector, is presented. HESS and the Suzaku X-ray satellite observed the blazar simultaneously in July 2006. The HESS results from this epoch are also discussed. This will enable, for the first time, future modeling of an SED determined from simultaneous observations at VHE and lower energies. Finally, results of a March 2006 VLT SINFONI spectroscopy campaign to determine the redshift of PG\\,1553+113 are also reported. ",
        "conclusions": "With a data set that is $\\sim$3 times larger than previously published (\\cite{HESS_discovery}), the HESS signal from PG\\,1553+113 is now highly significant ($\\sim$10$\\sigma$). Thus, the evidence for VHE emission previously reported is clearly verified.  However, the flux measured in 2005 is now $\\sim$3 times higher than first reported due to an improved calibration of the absolute energy scale of HESS. The statistical error on the VHE photon index is now reduced from $\\sim$0.6 to $\\sim$0.3.  Nevertheless, the error of 0.34 is still rather large, primarily due to the extreme softness of the observed spectrum ($\\Gamma=4.46$). The total HESS exposure on PG\\,1553+113 is $\\sim$25 hours.  Barring a flaring episode, not yet seen in two years of observations, a considerably larger total exposure ($\\sim$100 hours) would be required to significantly improve the spectral measurement.  This large exposure is unlikely to be quickly achieved.  However, the VHE flux from other AGN is known to vary dramatically and even a factor of a few would reduce the observation requirement considerably. Should such a VHE flare occur, not only will the error on the measured VHE spectrum be smaller, but the measured photon index may also be harder (see, e.g., \\cite{VHE_hardening}).  Both effects would dramatically improve the redshift constraints and correspondingly the accuracy of the source modeling. Therefore, the VHE flux from PG\\,1553$+$113 will continue to be monitored by HESS. In addition, the soft VHE spectrum makes it an ideal target for the lower-threshold HESS Phase-II \\cite{HESSII} which should make its first observations in 2009."
    },
    "0710/0710.2160_arXiv.txt": {
        "abstract": "We have searched for pulsed radio emission from magnetar 4U 0142+61 at the frequency of 111 MHz. No pulsed signal was detected from this source. Upper limits for mean flux density are 0.9~-~9~mJy depending on assumed duty cycle (.05~-~.5) of the pulsar. ",
        "introduction": "4U~0142+61 is an anomalous X-ray pulsar (AXP) with 8.7 seconds period. At period derivative $\\dot P = 2 \\times 10^{-12}$ magnetic field on a surface of a neutron star equals to $1.3 \\times 10^{+14}$~G - typical value for magnetars. The pulsar can be seen in hard X-ray (\\cite{den06}, \\cite{kui06}), optical (\\cite{ker02}), and infrared (\\cite{hul04}, \\cite {wan06}) bands. We searched for pulsed radio emission of this AXP at the frequency of 111~MHz. ",
        "conclusions": "The search of pulsed radio emission from anomalous X-ray pulsar (magnetar) 4U~0142+61 at the frequency of 111~MHz give no positive results.  Upper limits for mean flux density are 0.9~-~9~mJy depending on assumed duty cycle (.05 - .5) of the pulsar."
    },
    "0710/0710.2210_arXiv.txt": {
        "abstract": "We study the evolution of simple cosmic string loop solutions in an inflationary universe. We show, for the particular case of circular loops, that periodic solutions do exist in a de Sitter universe, below a critical loop radius $R_c H=1/2$. On the other hand, larger loops freeze in comoving coordinates, and we explicitly show that they can survive more $e$-foldings of inflation than point-like objects. We discuss the implications of these findings for the survival of realistic cosmic string loops during inflation, and for the general characteristics of post-inflationary cosmic string networks. We also consider the analogous solutions for domain walls, in which case the critical radius is $R_c H=2/3$. ",
        "introduction": "Introduction} Topological defects are unavoidably formed at cosmological phase transitions \\cite{Kibble,Book}. Studying their physical properties, evolution and cosmological consequences is therefore mandatory for a proper understanding of the early universe. The last three decades have seen dramatic progress in this task (see \\cite{Book} for a review), but significant knowledge gaps still exist. The aim of the present paper is to eliminate one of these gaps.  Defect-forming phase transitions often occur near or at the end of inflation. Moreover, it is possible that various stages of inflation occur, with defects being formed in between them \\cite{openinf}. It is therefore important to understand the effects of inflation on the defects, as well as to quantify their ability to survive any inflationary periods that occur after they form. The effects of inflation on the internal (microscopic) structure of defects have been studied in \\cite{basu}, which shows that those with thickness $\\delta H>1/\\sqrt{2}$ get smeared by expansion, while those with smaller thickness survive. Effects on macroscopic (cosmological) scales are well known for the case of long string networks \\cite{quantitative,extending}, but this is not so for the loop populations, despite the fact that they are known to contain a total amount of energy which is comparable to (if not greater than) that in the long strings \\cite{quantitative,protyloops}.  Here we study the effects of inflation on specific loop solutions, and consider both their microscopic and macroscopic evolution. We also discuss the consequences of our findings for cosmological scenarios involving string networks. Incidentally, we note that loops in a flat anisotropic universe were studied in \\cite{anisot}, where it was shown that the anisotropy of the background has an effect on the loop motion; the reasons for this will become clearer in what follows. We will also present a brief analysis of the analogous solutions for domain walls. We shall assume that the source of inflation is a perfect fluid with equation of state $p=w\\rho$ (with $ w < -1/3$) and the scale factor behaves as $a \\propto t^{2/3(1+w)}$; if $w = -1$ then $a \\propto \\exp(Ht)$ with $H$ being constant. ",
        "conclusions": ""
    },
    "0710/0710.5430_arXiv.txt": {
        "abstract": "\\footnotesize\\ This paper\\footnote{Published in 1999 in the book {\\it Solar Polarization}, edited by K.N. Nagendra \\& J.O. Stenflo. Kluwer Academic Publishers, 1999. (Astrophysics and Space Science Library ; Vol. 243), p. 143-156} addresses the problem of scattering line polarization and the Hanle effect in one-dimensional (1D), two-dimensional (2D) and three-dimensional (3D) media for the case of a two-level model atom without lower-level polarization and assuming complete frequency redistribution. The theoretical framework chosen for its formulation is the QED theory of Landi Degl'Innocenti (1983), which specifies the excitation state of the atoms in terms of the irreducible tensor components of the atomic density matrix. The self-consistent values of these density-matrix elements is to be determined by solving jointly the kinetic and radiative transfer equations for the Stokes parameters. We show how to achieve this by generalizing to Non-LTE polarization transfer the Jacobi-based ALI method of Olson {\\it et al.} (1986) and the iterative schemes based on Gauss-Seidel iteration of Trujillo Bueno and Fabiani Bendicho (1995). These methods essentially maintain the simplicity of the $\\Lambda-$iteration method, but their convergence rate is extremely high. Finally, some 1D and 2D model calculations are presented that illustrate the effect of horizontal atmospheric inhomogeneities on magnetic and non-magnetic resonance line polarization signals. ",
        "introduction": "The scattering line polarization, and its modification due to a weak magnetic field ---such that the Zeeman splitting is negligible compared with the line width (the so called Hanle effect; Hanle, 1924)---, sensitively depends on the {\\it anisotropy} of the radiation field and on the magnetic field vector geometry (Landi Degl'Innocenti, 1985; Stenflo, 1994). The solar atmospheric plasma is spatially inhomogeneous, with vertical and horizontal variations, not only in temperature, macroscopic velocity and density, but also in the orientation and intensity of the magnetic field (see S\\'anchez Almeida, 1999 for new insights in this respect). It is thus clear that, in order to fully exploit the Hanle effect as a diagnostic tool for weak magnetic fields, we also need to address the problem of resonance line polarization and the Hanle effect in 2D and 3D media where the radiation field's anisotropy is different from that corresponding to the currently-assumed 1D atmospheric models.  To this end, this contribution begins presenting a formulation of resonance line polarization and the Hanle effect that we consider as the most suitable one for practical RT applications. It is based on the density-matrix theory for the generation and transfer of polarized radiation (see Landi Degl'Innocenti, 1983; 1984; 1985). In this paper we consider the standard case of a two-level model atom neglecting atomic polarization in its lower level (i.e. it is assumed that the lower-level Zeeman sublevels are equally populated and that there are no coherences among them). The quantities whose {\\it self-consistent} values are to be determined are the irreducible tensor components of the density matrix ($\\rho^K_Q$), which depend only on the spatial coordinates. The statistical equilibrium (SE) and RT equations to be solved are valid independently of whether we assume 1D, 2D or 3D geometries.  A summary of previous work done in the subject of the numerical solution of Non-LTE polarization transfer problems can be found in Trujillo Bueno and Manso Sainz (1999). In this respect, we should mention the recent work of Nagendra {\\it et al.} (1998; see also their contribution in these proceedings) where the Hanle effect in 1D is considered using a different theoretical approach. For information concerning the numerical solution of more general polarization transfer problems formulated with the density-matrix theory see Trujillo Bueno (1999).  The outline of this paper is as follows. Section 2 presents the basic equations for the case of a two level atom without lower-level atomic polarization and assuming complete frequency redistribution. Section 3 shows how the very efficient iterative methods of solution investigated by Trujillo Bueno and Fabiani Bendicho (1995) (Jacobi, Gauss-Seidel and successive over-relaxation) can be suitably generalized to the problem of Non-LTE polarized radiative transfer in the Hanle effect regime. Finally, in Sect. 4 we present the results of some illustrative 2D Hanle-effect calculations for triplet lines and discuss the ensuing horizontal transfer effects.% ",
        "conclusions": "We have developed a Hanle effect code that allows the numerical simulation of resonance line polarization signals in the presence of weak magnetic fields in 1D, 2D and 3D media. The governing equations have been formulated working within the framework of the density matrix polarization transfer theory of Landi Degl'Innocenti (1983, 1985). These SE and RT equations are the same independently of whether we are considering 1D, 2D or 3D atmospheric models. The six $\\rho^K_Q$-unknowns of the problem are neither frequency nor angle dependent, since they only vary with the spatial position.  Three different iterative schemes that were originally developed for RT applications in the unpolarized case have been generalized to solve this set of equations: Jacobi (ALI), Gauss-Seidel and SOR (Trujillo Bueno \\& Fabiani Bendicho, 1995, Trujillo Bueno \\& Manso Sainz 1999). This kind of iterative methods does not make use of any matrix inversion, and essentially maintain the $\\Lambda$-iteration simplicity. The only difference between the 1D, 2D and 3D versions of our Hanle effect code lies in the formal solution routine that calculates the radiation field tensors from the current values of the density-matrix elements. To this end, in 2D we use the formal solver developed by Auer, Fabiani Bendicho and Trujillo Bueno (1994) and in 3D we use the one presented at this workshop by Fabiani Bendicho and Trujillo Bueno (1999).  We have also shown some Hanle effect results for 1D and 2D media. These calculations illustrate how weak magnetic fields and horizontal radiative transfer effects compete to modify the scattering line polarization signals expected from plane-parallel 1D atmospheres. Thus, further careful investigations must be done in order to separate both effects, when diagnosing weak solar magnetic fields via the Hanle effect. This type of future studies should be done thinking in the interpretation of {\\it low} spatial resolution scattering line polarization observations. Our numerical approach is very efficient and suitable to investigate scattering polarization signals for a variety of atmospheric models having any desired temperature, density and magnetic field vector variations. Another useful research that can be done with our Hanle effect codes concerns the simulation of polarization signals emerging from realistic MHD and semi-empirical 2D and 3D models."
    },
    "0710/0710.5882.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory \\abstract % context heading (optional) % {} leave it empty if necessary {The young $\\sigma$~Orionis cluster is an indispensable basis for understanding the formation and evolution of stars, brown dwarfs and planetary-mass objects. Our knowledge of its stellar population is, however, incomplete.} % aims heading (mandatory) {I present the Mayrit catalogue, that comprises most of the stars and high-mass brown dwarfs of the cluster.} % methods heading (mandatory) {The basis of this work is an optical-near infrared correlation between the 2MASS and DENIS catalogues in a circular area of radius 30\\,arcmin centred on the OB-type binary $\\sigma$~Ori~AB. The analysis is supported on a bibliographic search of confirmed cluster members with features of youth and on additional X-ray, mid-infrared and astrometric data.} % results heading (mandatory) {I list 241 $\\sigma$~Orionis stars and brown dwarfs with known features of youth, {97} candidate cluster members (40 are new) and {115} back- and foreground sources in the survey area. The {338} cluster members and member candidates constitute the Mayrit catalogue.} % conclusions heading (optional), leave it empty if necessary {This catalogue is a suitable input for studying the spatial ditribution, multiplicity, properties and frequency of discs and the complete mass function of $\\sigma$~Orionis.} ",
        "introduction": "The $\\sigma$ Orionis cluster in the \\object{Ori~OB~1b} Association is getting as important for the study of the formation, evolution and characterisation of stars and substellar objects as other famous clusters and star-forming regions, like the \\object{Hyades}, the \\object{Pleiades}, the \\object{Orion Nebula Cluster} or the \\object{Taurus-Auriga Complex}. The $\\sigma$~Orionis cluster is young (3$\\pm$2\\,Ma), nearby ( $\\sim$385\\,pc) and relatively free of extinction (Lee 1968; Brown et~al. 1994; Oliveira et~al. 2002; Zapatero Osorio et~al. 2002a; Sherry et~al. 2004; B\\'ejar et~al. 2004b; Caballero 2007d). Firstly identified by Garrison (1967) and Lyng\\aa~(1981), $\\sigma$~Orionis was rediscovered by Wolk (1996) and Walter et~al. (1997). They reported a clustering of young low-mass stars, many of them positionally coincident with X-ray sources, surrounding the Trapezium-like, multiple stellar system $\\sigma$~Ori in the vicinity of the \\object{Horsehead Nebula} (see Caballero 2007b for a description of the multiple system $\\sigma$~Ori and its surroundings). Previously, the area had been investigated during wide searches in the Orion complex with prism-objective and Schmidt plates, detecting a wealth of emission stars (e.g. Haro \\& Moreno 1953; Wiramihardja et~al. 1989), but the cluster had not been treated as an independent entity within the complex. Complete compilations of the determinations in the literature of the age, heliocentric distance, frequency of discs and mass function of the $\\sigma$~Orionis cluster are in Caballero (2007a). A chapter of the Handbook of Star Forming Regions, edited by B.~Reipurth, will be exclusively devoted to $\\sigma$~Orionis (F.~M.~Walter et~al., in~press).  After the seminal work by B\\'ejar et~al. (1999), who found for the first time a rich population of young brown dwarfs in $\\sigma$~Orionis, the cluster has turned into an excellent laboratory for the study of, e.g.:  \\begin{itemize} \\item the search for free-floating planetary-mass objects (with masses below the deuterium-burning limit) and the study of the substellar mass function down to a few Jupiter masses (Zapatero Osorio et~al. 2000; B\\'ejar et~al. 2001; Gonz\\'alez-Garc\\'{\\i}a et~al. 2006; Caballero et~al. 2007); \\item the frequency and the properties of $\\sim$3\\,Ma-old discs at different mass intervals (Jayawardhana et~al. 2003; Oliveira et~al. 2004, 2006; Hern\\'andez et~al. 2007; Caballero et~al. 2007; Zapatero Osorio et~al. 2007a); \\item the masses of OB-type stars in resolved binary systems (Heintz 1997; Mason et~al. 1998; Caballero 2007d); \\item the X-ray emission of young stars and brown dwarfs (Mokler \\& Stelzer 2002; Sanz-Forcada et~al. 2004; Franciosini et~al. 2006; Caballero 2007b); \\item the characterisitics of jets and Herbig-Haro objects (Reipurth et~al. 1998; L\\'opez-Mart\\'{\\i}n et~al. 2001; Andrews et~al. 2004); and \\item the photometric variability of low-mass stars and brown dwarfs (Bailer-Jones \\& Mundt 2001; Caballero et~al. 2004; Scholz \\& Eisl\\\"offel~2004). \\end{itemize}  Many interesting star-like objects have been discovered in the cluster, from the helium-rich, B2.0Vp-type magnetic star $\\sigma$~Ori~E (Greenstein \\& Wallerstein 1958), through the Class~I object candidate IRAS~05358--0238 (Oliveira \\& van Loon 2004), to the hypothetical proplyd $\\sigma$~Ori~IRS1 (van Loon \\& Oliveira 2003; Caballero 2005, 2007b). The most interesting objects in the cluster are, however, below the hydrogen-burning mass limit. Some of these substellar objects are the $\\sim$T6-type object \\object{S\\,Ori~70} (which may be the least massive body directly detected out of the Solar System, $\\sim$3\\,$M_{\\rm Jup}$ -- Zapatero Osorio et~al. 2002c, 2007b; Burgasser et~al. 2004), the T Tauri substellar analog S\\,Ori~J053825.4--024241 (which is the most variable brown dwarf yet found; Caballero et~al. 2006a), the two strong H$\\alpha$ emitters at the planetary boundary \\object{S\\,Ori~55} and \\object{S\\,Ori~71} (with masses of only 10--20\\,$M_{\\rm Jup}$ and equivalent widths of the H$\\alpha$ line of up to --700\\,\\AA; Zapatero Osorio et~al. 2002b; Barrado y Navascu\\'es et~al. 2002a) and the brown dwarf-exoplanet system candidate \\object{SE~70} + \\object{S\\,Ori~68} (which could be the widest planetary system known so far; Caballero et~al. 2006b). The number of known substellar objects in $\\sigma$~Orionis is comparable to those of other rich, more massive, younger star-forming regions like Chamaeleon, Ophiuchus or the Orion Nebula Cluster. However, $\\sigma$~Orionis is by far the region with the largest amount of brown dwarfs with membership confirmation ($>$ 30; Caballero et~al. 2007) and planetary-mass object candidates (29; Zapatero Osorio et~al. 2000; Gonz\\'alez-Garc\\'{\\i}a et~al. 2006; Caballero 2007b; Caballero et~al. 2007). Many of the latter bodies have measured L spectral types and/or flux excess longwards of 5\\,$\\mu$m (Zapatero Osorio et~al.~(2007a).  Important efforts have been recently carried out to characterise the $\\sigma$~Orionis cluster in general and to investigate the connection between its stellar and substellar populations in particular (e.g. B\\'ejar et~al. 2004a; Sherry et~al. 2004; Kenyon et~al. 2005; Burningham et~al. 2005; Caballero 2005, 2007a, 2007b, 2007c). The works by Kenyon et~al. (2005), who investigated membership, binarity and accretion among very low-mass stars and brown dwarfs surrounding $\\sigma$~Ori, and, especially, Sherry et~al. (2004) stand out. The latter authors presented an ambitious study, estimating the number of cluster members in the mass range 0.2\\,$M_\\odot$ $\\lesssim M \\lesssim$ 1.0\\,$M_\\odot$ and the radius, age and total mass of the cluster. All these works are, however, incomplete or biased in some way: no comprehensive, homogenous, multi-band study, fully covering the whole stellar mass interval (from $\\sim$20\\,$M_\\odot$ down to the substellar boundary) and the cluster area without gaps (from the very centre to the border) exists so far in $\\sigma$~Orionis.  I extend the study of the brightest stars of the cluster shown in Caballero (2007a) down to the hydrogen-burning mass limit and beyond. I start out from a correlation between the 2MASS and DENIS catalogues within the environment of the Virual Observatory\\footnote{See {http://www.ivoa.net}.} and a bibliographic search of confirmed cluster members, and support them with spectroscopic and photometric data from the X-ray region to 120\\,$\\mu$m when available. The outcome of this work is the Mayrit catalogue, that comprises the majority of the stars and high-mass brown dwarfs of $\\sigma$~Orionis. ",
        "conclusions": "The $\\sim$3\\,Ma-old $\\sigma$~Orionis cluster is a new cornerstone for observational and theoretical studies with the aim to understand the general processes of collapse and fragmentation of a molecular cloud, formation of stars and substellar objects and evolution of circumstellar discs. I~present a comprehensive, rather complete catalogue of cluster members that can be used for further studies in $\\sigma$~Orionis (e.g. spatial distribution, multiplicity, mass function and frequency and characterisation of~discs). This investigation covers the whole stellar and part of the brown dwarf domain of the cluster from $\\sim$18 to $\\sim$0.03\\,$M_\\odot$.  I~have performed an $IK_{\\rm s}$ survey in a 30\\,arcmin-radius region centred on the O9.5V+B0.5V binary $\\sigma$~Ori~AB using the Aladin tool and optical and near-infrared data from the DENIS and 2MASS catalogues. The photometric data have been complemented with information from the literature regarding the membership of known sources. I~have compiled a list of 241 cluster stars and brown dwarfs with known features of youth (e.g. early spectral types, Li~{\\sc i} in absorption, near and mid infrared excess attributed to surrounding discs, strong X-ray and H$\\alpha$ emissions), 85 fore- and background stars with astrometric and spectroscopic data (nine are new), {18} galaxies with extended FWHMs and 12 probable reddened sources in a nebulosity northeastern of the survey area, associated to the Horsehead Nebula. From the $I$ vs. $I-K_{\\rm s}$ diagram, I have selected {97} additional photometric cluster member candidates without reliable membership information. This makes a list of {338} $\\sigma$~Orionis members and member candidates, the Mayrit catalogue, of which more than 70\\,\\% display features of extreme youth. %Global contamination and incompleteness are at the order of magnitude of %only 10\\,\\%. I~tabulate precise coordinates, $IJHK_{\\rm s}$ and suplementary information of all the cluster members, member candidates and non-members.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0710/0710.3558.txt": {
        "abstract": "We investigate the dynamics of magnetic fields in spiral galaxies by performing 3D Magnetohydrodynamics (MHD) simulations of galactic discs subject to a spiral potential using cold gas, warm gas and a two phase mixture of both. Recent hydrodynamic simulations have demonstrated the formation of inter-arm spurs as well as spiral arm molecular clouds provided the ISM model includes a cold HI phase. We find that the main effect of adding a magnetic field to these calculations is to inhibit the formation of structure in the disc. However, provided a cold phase is included, spurs and spiral arm clumps are still present if $\\beta \\gtrsim 0.1$ in the cold gas. A caveat to the two phase calculations though is that by assuming a uniform initial distribution,  $\\beta \\gtrsim 10$ in the warm gas, emphasizing that models with more consistent initial conditions and thermodynamics are required. Our simulations with only warm gas do not show such structure, irrespective of the magnetic field strength.  Furthermore, we find that the introduction of a cold HI phase naturally produces the observed degree of disorder in the magnetic field, which is again absent from simulations using only warm gas. Whilst the global magnetic field follows the large scale gas flow, the magnetic field also contains a substantial random component that is produced by the velocity dispersion induced in the cold gas during the passage through a spiral shock. Without any cold gas, the magnetic field in the warm phase remains relatively well ordered apart from becoming compressed in the spiral shocks. Our results provide a natural explanation for the observed high proportions of disordered magnetic field in spiral galaxies and we thus predict that the relative strengths of the random and ordered components of the magnetic field observed in spiral galaxies will depend on the dynamics of spiral shocks. ",
        "introduction": "Observations of the magnetic field in galaxies indicate that the magnetic pressure is comparable in magnitude to the thermal pressure and interstellar turbulence \\citep{Heiles2005}. Consequently magnetic fields are expected to play an important part in the dynamics of the ISM, including spiral shocks and the formation of molecular clouds.  The response of a gaseous disc to a spiral potential has been well examined both theoretically and numerically \\citep*{Roberts1969,Shu1972,Woodward1976,Gittins2004,Wada2004,DBP2006}. The gas flow evolves into a quasi-stationary solution which contains spiral shocks, providing the potential is of sufficient strength. \\citet{Roberts1970} found that the strength of the shock decreases with increasing magnetic field strength, whilst the magnetic field strength is amplified in the spiral shock, the latter agreeing with observations at the time for galactic magnetic fields.  There are numerous hydrodynamic simulations of gas discs subject to a spiral potential (e.g. \\citealt{Kim2002,Chak2003,Slyz2003,Wada2004,Gittins2004,Dobbs2006}) which discuss the formation of substructure and location of the shock. Whilst 2D simulations of warm gas show the formation of small-scale spurs perpendicular to the spiral arms \\citep{Gittins2004,Wada2004}, these features appear to be suppressed in 3D simulations \\citep{Kim2006}. On the other hand, \\citet{Dobbs2006} do find more extensive substructure and spiral arm clouds \\citep*{DBP2006,Dobbs2007} in 3D simulations, but only in models where the gas temperature is $\\leq 1000$ K. There are comparatively fewer MHD calculations, largely due to the relatively recent addition of magnetic fields into numerical codes, and the limitations of numerical resolution. \\citet{Shetty2006} investigate the formation of spurs using 2D MHD grid-based calculations of the warm ISM with self gravity, concluding that self gravity is necessary for spurs to evolve. \\citet{Gomez2002,Gomez2004} also perform 3D calculations with warm gas, but whilst both groups compare MHD with non-MHD results, neither investigate the effects of varying the field strength or ISM temperature.  Recent observations now provide much more detailed measurements of the magnetic field strength and direction in the diffuse (warm) ISM (see reviews by \\citealt{Beck1996} and \\citealt{Beck2007}). Spiral magnetic arms appear to occur in all disc galaxies, regardless of the presence of optical spiral arms \\citep{Beck2005}. Typically the field in the inter-arm regions is several $\\mu$G, but may be 10 or more $\\mu$G in the spiral arms, and the total field contains random and ordered (regular) components of comparable strength. The general consensus currently favours dynamo theory to explain the origin of these relatively strong magnetic fields \\citep*{Parker1971,P1971,Balsara2004,Beck2005}, whereby the magnetic field is generated by turbulence in the ISM, and can produce or enhance magnetic spiral arm arms \\citep*{Panesar1992,Rohde1998}. However in grand design galaxies where the gas experiences strong spiral shocks, such as M51, the magnetic spiral arms tend to be aligned with the dust lanes \\citep{Nein1992} suggesting that the magnetic field is strongly related to the dynamics of the spiral shock.  Magnetic fields have been long associated with molecular cloud formation, originally through the Parker instability \\citep{Parker1966}, and more recently via the magneto-rotational \\citep*{Kim2003,Piontek2005} instability. The Parker instability is expected to produce sinusoidal motions in the $z$ direction, with the magnetic field channeling gas into dense concentrations in the plane of the disc. Molecular clouds formed in this way were originally believed to be magnetically supported, with lifetimes of order $10^8$ years \\citep*{Zweibel1975,Shu1987}. However the Parker instability is found to be relatively weak, and singularly insufficient to induce molecular cloud formation \\citep*{Elmegreen1982,OtherKim2001,KOS2002}. The magneto-rotational instability (MRI) has been proposed as a formation mechanism away from spiral arms \\citep{Kim2003}. However the MRI generally takes a few orbits to initialise, particularly if the field is toroidal (e.g. 10's of orbits in simulations by \\citealt{Nish2006}), and is therefore of less importance where spiral shocks occur. Magnetic fields may be more relevant to GMC (giant molecular cloud) formation in conjunction with gravitational instabilities \\citep{ElmegreenMJI1987,Kim2001,KOS2002} [also known as the magneto-Jeans-instability, MJI] or cooling \\citep{Kosinski2007}.  In this paper we describe 3D MHD calculations of a galactic disc subject to a spiral perturbation. These are the first fully Lagrangian MHD calculations to model this problem, using the Smoothed Particle Magneto-hydrodynamics (SPMHD) code \\citep{Price2004,Price2005,PB2007}. We compare the structure of the disc for a range of magnetic field strengths, assuming a single phase medium of cold or warm gas, or a two-phase medium. We also describe the strength and morphology of the magnetic field, and relate these to observations. In this paper we focus on the effect of the spiral shock inducing structure in the gas rather than Parker or MRI instabilities, and we leave a discussion of results including self-gravity to a future paper. ",
        "conclusions": "We have performed simulations of galactic discs subject to a spiral potential with a range of field strengths. The main results we have discussed are 1) the reduction in structure across the disc as the magnetic field strength increases, and 2) the possibility of spiral shocks inducing an irregular magnetic field in the ISM.  As the strength of the magnetic field increases, the strength of the spiral shocks and therefore density of the spiral arms are reduced. This is as expected from the analysis of \\citet{Roberts1970}. Consequently the formation of spurs is increasingly suppressed for higher magnetic field strength. The spiral arms themselves are more continuous and clumps along the spiral arms are less dense. This supports previous 2D results \\citep{Shetty2006,Tanaka2005} which find that both Kelvin-Helmholtz and gravitational  instabilities perpendicular to the magnetic field are reduced by stronger magnetic fields (although we do not relate the structure in our simulations to these instabilities, our overall conclusions agree). We find that the addition of the magnetic field is similar in effect to an increase in thermal pressure, in that both provide a pressure which oppose the formation of structure, and smooth out the gas.  Nonetheless, we still find significant inter-arm structure with the presence of a magnetic field, unlike the results (those which are non-self gravitating) of \\citet{Shetty2006}. This structure is present in the cold component of the ISM, which has generally not been included in previous simulations. Inhomogeneities present in the initial random distribution of gas become amplified by spiral shocks. With warm gas or strong magnetic fields, these inhomogeneities are smoothed out by the pressure. For our calculations with cold gas, we find that magnetic fields only prevent substructure when the ratio of gas to magnetic pressure ($\\beta$) is $\\lesssim 0.1$. For $\\beta \\gtrsim 1$, spurs perpendicular to the arm still form in the cold gas. The additional pressure provided by warm gas in the two-phase results increases the longevity of structure in the inter-arm regions, even when the magnetic field dominates for the cold gas.  The main difference between ours and previous results is that we use cold gas, although we also use a weaker magnetic field than \\citet{Shetty2006}. We find structure is present for values of $\\beta$ in the cold gas similar to observations, but the volume averaged magnetic field strengths in our two phase models are typically $< 1 \\mu$G, lower than observations. This discrepancy arises because we start with an initially uniform distribution of cold gas. To obtain a similar $\\beta$ for both the cold and warm gas in our two phase models, we would need the cold gas to be 100 times denser than the warm gas, i.e. this just indicates that structure is \\emph{already} present in the cold gas. If we used higher field strengths, we would find, like the results of \\citet{Shetty2006}, very little structure emerges. In this case, $\\beta$ for the warm gas would be closer to observations, but the magnetic pressure would dominate the thermal pressure by a factor of 100 or so in the cold gas, inconsistent with observations. Thus we either need to distribute the cold gas initially in clumps, or ideally include a much more detailed thermal treatment of the ISM, in order to obtain magnetic field strengths \\emph{and} $\\beta$ consistent with observations.  We find that whilst the ISM appears highly structured in observations, we would expect a higher degree of structure with relatively weak magnetic fields. For instance, the two phase model where $\\beta_{cold}=4$ retains much more structure typical of grand design galaxies compared to the case where $\\beta_{cold}=0.4$. Current observations suggest that magnetic pressure exceeds thermal pressure \\citep{HT2005}. We however note that $\\beta$ exhibits a range of values in our simulations, and $\\beta$ tends to be take smaller values in the spiral arms and dense gas (for a given temperature), which are more likely to correspond with observations. Again a more complete treatment of the ISM in future work, and further observations of the CNM will allow a better comparison.  In our simulations, the magnetic field is compressed by the spiral shocks, as expected from a straightforward analysis of MHD shocks \\citep{Roberts1970,Priest1982}. The relative increase in the magnetic field strength is greater where the shock is stronger. However the most intriguing result from our simulations is the possibility that spiral shocks generate an irregular magnetic field. This process has not been identified in previous simulations, which we attribute to the fact that they have not included a cold phase. Galaxies are known to contain a random component of the field, but it is usually supposed that this is due to supernovae and/or feedback from stars. We therefore postulate that spiral shocks are important in generating disorder in the magnetic field, whilst simultaneously inducing a velocity dispersion and density structure in the gas (\\citealt*{Bonnell2006,KKO2006}; \\citealt{DBP2006}). The degree of order in the disc, similar to the presence of arm/inter-arm structure is related to the strength of the shock. When the gas is cold and the magnetic field is weak, the shock is much stronger, and a higher velocity dispersion is induced in the gas. Consequently the magnetic field lines follow the gas and the field becomes tangled. In simulations with warm gas, we find the magnetic field is almost entirely regular, although with a stronger shock, the field may become more irregular. However in our two-phase results, the velocities and irregularities in the magnetic field of the cold gas induces comparatively more disorder in the warm gas than the single phase calculations. Although there are no measurements of regular versus disordered components of the magnetic field in cold HI, we would predict from our results that the field in the cold gas will be more disordered than that of warm HI."
    },
    "0710/0710.1554_arXiv.txt": {
        "abstract": "{}{We have investigated a sample of 28 well-known spectroscopically-identified magnetic Ap/Bp stars, with weak, poorly-determined or previously undetected magnetic fields. The aim of this study is to explore the weak part of the magnetic field distribution of Ap/Bp stars.} {Using the MuSiCoS and NARVAL spectropolarimeters at T\\'elescope Bernard Lyot (Observatoire du Pic du Midi, France) and the cross-correlation technique Least Squares Deconvolution (LSD), we have obtained 282 LSD Stokes $V$ signatures of our 28 sample stars, in order to detect the magnetic field and to infer its longitudinal component with high precision (median $\\sigma=40$~G). } {For the 28 studied stars, we have obtained 27 detections of Stokes $V$ Zeeman signatures from the MuSiCoS observations. Detection of the Stokes $V$ signature of the $28^{\\rm th}$ star (HD~32650) was obtained during science demonstration time of the new NARVAL spectropolarimeter at Pic du Midi.  This result shows clearly that when observed with sufficient precision, all firmly classified Ap/Bp stars show detectable surface magnetic fields. Furthermore, all detected magnetic fields correspond to longitudinal fields which are significantly greater than some tens of G. To better characterise the surface magnetic field intensities and geometries of the sample, we have phased the longitudinal field measurements of each star using new and previously-published rotational periods, and modeled them to infer the dipolar field intensity and the magnetic obliquity. The distribution of derived dipole strengths for these stars exhibits a plateau at about 1 kG, falling off to larger and smaller field strengths. Remarkably, in this sample of stars selected for their presumably weak magnetic fields, we find only 2 stars for which the derived dipole strength is weaker than 300~G. We interpret this ``magnetic threshold'' as a critical value necessary for the stability of large-scale magnetic fields, and develop a simple quantitative model that is able to approximately reproduce the observed threshold characteristics. This scenario leads to a natural explanation of the small fraction of intermediate-mass magnetic stars. It may also explain the near-absence of magnetic fields in more massive B and O-type stars.}{}{} ",
        "introduction": "The magnetic chemically peculiar Ap/Bp stars are the non-degenerate stars for which the strongest magnetic fields have been measured (Landstreet 1992). Although the fields are thought to be fossil remnants of flux swept up during star formation or produced via dynamo action on the pre-main sequence, their origin is not understood in any real detail (e.g. Moss 2001). Furthermore, the role of the magnetic field in the diffusion processes which are responsible for their chemical peculiarity has been studied in only a schematic fashion.  Although more than one thousand main sequence {A-type stars} have been catalogued as magnetic Ap/Bp stars (Renson et al. 1991) following the scheme of Preston (1974), direct measurements of the magnetic field have been obtained for only a few hundred of them (Romanyuk 2000, Bychkov et al. 2003). Examination of the published measurements shows that the majority of the reported values are rather large. For example, $55\\%$ of the 210 stars of the catalogue of Romanyuk (2000) with published magnetic field measurements have a maximum unsigned line-of-sight (longitudinal) magnetic field $B_\\ell$ larger than 1 kG. On the other hand, according to Bohlender \\& Landstreet (1990), the median root-mean-square (rms) longitudinal magnetic field of Ap stars (based on a small magnitude-limited sample observed by Borra \\& Landstreet 1980) is only about 300~G (the largest rms field they report is only 710 G). This implies that most Ap stars have relatively weak ($\\ltsim 1$~kG) magnetic fields, and that the available observations are strongly biased toward stars with the strongest and most easily-measured fields. One consequence of this bias is that the weak-field part of the magnetic field distribution of Ap stars is poorly studied. It is not known if it increases monotonically toward arbitrarily small field strength, or if it is truncated at a minimum magnetic field strength (as proposed by Glagolevskij \\& Chountonov 2002).  In order to improve our knowledge of the weak-field part of the magnetic field distribution of Ap stars, we have undertaken a study of a sample of 28 well-known spectroscopically-identified Ap/Bp stars, with very weak, poorly-determined or previously-undetected magnetic fields. We describe our survey in Sect. 2 and report our observational and modeling results in Sect. 3 and Sect. 4. We discuss the implications of our results, suggesting one possible interpretation involving the stability of large scale magnetic fields, in Sect. 5 and give our conclusions in Sect. 6.  \\section {The weak-field Ap stars survey}  \\subsection{The selected sample}  Our sample is composed of spectroscopically-identified Ap/Bp stars belonging to the HD catalogue. Twelve stars were selected based on the observations of Borra \\& Landstreet (1980) and Bohlender et al. (1993), identifying stars for which no significant detection of the magnetic field was obtained. Thirteen additional targets are stars for which only old photographic measurements of the magnetic field were available, typically by Babcock (1958), and for which measurements have poor precision and do not provide a significant detection of the magnetic field; these stars were generally selected using the catalogues of Romanyuk (2000) or Bychkov et al. (2003). Finally, 3 stars of our sample were selected which had not been observed for magnetic field before this work. The observational properties of the 28 stars are presented in Table 1. Section 3.2 gives more details on each star and on the obtained results.  Because all of these stars are relatively bright, most have been known for decades and have been well studied. All appear in the Hipparcos catalogue (Perryman et al.,  1997), and all but 6 have $\\sigma_\\pi/\\pi<0.2$. The majority have photometrically-determined rotational periods and published values of $v\\sin i$. Many have been studied using high-resolution spectroscopy and Doppler Imaging. Therefore the classification of this sample as {\\em bona fide} Ap/Bp stars is generally quite firm. As a consequence of this careful selection (and as will be described later in the paper), {\\em no} non-Ap/Bp star has been mistakenly included in the sample.  \\subsection{Observations and reduction}  Stokes $V$ and Stokes $I$ spectra of the 28 sample stars were obtained during 10 observing runs, from July 2001 to June 2006.  We used the MuSiCoS spectropolarimeter attached the Bernard Lyot telescope (TBL) at Observatoire du Pic du Midi. The MuSiCoS spectropolarimeter is composed of a cross-dispersed echelle spectrograph (Baudrand \\& B\\\" ohm 1992) and a dedicated polarimeter module (Donati et al. 1999). The spectrograph is a table-top instrument, fed by a double optical fibre directly from the Cassegrain-mounted polarimeter. In one single exposure, this apparatus allows the acquisition of a stellar spectrum in a given polarisation (Stokes $V$ in this case) throughout the spectral range 450 to 660 nm with a resolving power of about 35000. Spectra in both orthogonal polarisations are recorded simultaneously by the CCD detector. A complete Stokes $V$ exposure consists of a sequence of four subexposures, between which the quarter-wave plate is rotated by 90$^{\\rm o}$. This has the effect of exchanging the beams in the whole instrument, and in particular switching the positions of the two orthogonally polarised spectra on the CCD, thereby reducing spurious polarisation signatures.  The echelle polarisation spectra were reduced using the ESpRIT package (Donati et al. 1997). The observation and reduction procedures are more thoroughly described by Shorlin et al. (2002).  The correct operation of the MuSiCoS instrument, and in particular the absence of spurious magnetic field detections, is supported by other data obtained during these same observing runs, including studies of non-magnetic A-type stars (e.g. Shorlin et al. 2002), magnetic A, B and O-type stars (e.g., Ryabchikova et al. 2005a,  Donati et al. 2001, Wade et al. 2006a) and magnetic late-type stars (e.g. Petit et al. 2005).  MuSiCoS has recently been decomissioned, and has been replaced with NARVAL (Auri\\`ere 2003), the new-generation spectropolarimeter which is a copy of the ESPaDOnS instrument in operation at the Canada-France-Hawaii Telescope (Donati 2004, 2007: in preparation). The main improvements of NARVAL in polarisation mode with respect to MuSiCoS are a spectral resolution of about 65000, spectral response between 370 nm and 1000 nm, and an overall sensitivity increased by a factor of about 30.  \\begin{table*}[t] \\caption{Observational properties of the weak-field Ap star sample. Columns give ID and HD number, visual magnitude, spectral classification, effective temperature, luminosity and radius (with associated $1\\sigma$ error bars), adopted LSD mask temperature, number of observations obtained and detection level (d=definite detection; m=marginal detection), maximum observed unsigned longitudinal field in G, $1\\sigma$ error in G, and peak longitudinal field detection significance {\\it z } = $ |B^{max}_{l}|$ / $\\sigma$. $B_{\\rm d}^{\\rm min, 3.3}$ is the minimum dipole field (at $2\\sigma$) inferred from the maximum measured longitudinal field and Eq. (7).} \\begin{tabular}{lcrr|ccc|cccccc} \\\\ \\hline ID & HD      & $m_V$ & Spec & $T_{\\rm eff}$ & $\\log L$ & $R$ & Mask &   \\# & Det.& $|B_{\\ell}|^{max}\\pm \\sigma$ & {\\it z } & $B_{\\rm d}^{\\rm min, 3.3}$\\\\ &      &  & Type &  (K) & ($L_\\odot$) & ($R_\\odot$) &  (kK) &  & level & (G)& &(G)\\\\ \\hline \\\\ HN And& 8441     &  6.7 &A2p  &$ 9060\\pm 300$ & $ 1.90\\pm\t0.16$&$3.6\t\\pm\t0.9 $   &  9  & 8& 8d       & $157\\pm 18 $& 8.7 &  399    \\\\ 43 Cas& 10221    &  5.5 &A0sp &$10660\\pm 350$ & $ 2.11\\pm\t0.09$&$3.3\t\\pm\t0.7 $   &  11  & 10& 8d     & $148\\pm 34 $& 4.3 &  264    \\\\ $\\iota$ Cas &15089 & 4.5 &A5p &$ 8360\\pm 275$ & $ 1.38\\pm\t0.05$&$2.3\t\\pm\t0.4 $   & 9  & 12& 12d      & $486\\pm 23 $& 20.8&  1452   \\\\ &15144 & 5.8 &A6Vsp           &$ 8480\\pm 280$ & $ 1.21\\pm\t0.07$&$1.9\t\\pm\t0.3 $   & 9  & 6& 6d        & $631\\pm 15 $& 42.0&  1983   \\\\ 21 Per& 18296  & 5.0 &B9p     &$ 9360\\pm 310$ & $ 2.08\\pm\t0.11$&$4.2\t\\pm\t1.0 $   & 10  & 2& 2d       & $213\\pm 20 $& 10.6&  571    \\\\ 9 Tau & 22374 & 6.7 &A2p      &$8390\\pm  275$ & $ 1.48\\pm\t0.13$&$2.6\t\\pm\t0.7 $   &  9  & 2& 2d       & $523\\pm 24 $& 21.7&  1568   \\\\ 56 Tau &27309    &  5.3 &A0sp &$12730\\pm 420$ & $ 2.06\\pm\t0.08$&$2.2\t\\pm\t0.4 $   & 12  & 12& 12d     & $804\\pm 50 $& 16.1&  2323   \\\\ 11 Ori& 32549    &  4.7 &A0sp &$10220\\pm 335$ & $ 2.35\\pm\t0.12$&$4.7\t\\pm\t1.2 $   & 11  & 11 & 1d3m &  $186 \\pm 39$& 4.7 &  356    \\\\ &32650 & 5.4 &B9sp            &$11920\\pm 390$ & $ 2.11\\pm\t0.07$&$2.7\t\\pm\t0.5 $   & 12  & 18 & 2m3d&   $91  \\pm 18$& 5.0 &  237     \\\\ &37687 & 7.0 &B8              &$9450\\pm  310$ & $ 2.18\\pm\t0.25$&$4.6\t\\pm\t1.9 $  & 10  & 2 & 2d       & $766\\pm 119$& 6.4 &  1742   \\\\ 137 Tau& 39317 & 5.5 &B9spe   &$10130\\pm 330$ & $ 2.19\\pm\t0.14$&$4.0\t\\pm\t1.1 $   & 11  & 8 & 3d1m  &  $216 \\pm 59$& 3.6 &  323    \\\\ &40711 & 8.5 &A0              &$8070\\pm  265$ & $ 1.94\\pm\t0.61$&$4.8\t\\pm\t5.3 $   &  9  & 3 & 1d1m  &  $528 \\pm 38$& 13.8&  1492   \\\\ &43819 & 6.2 &B9IIIsp         &$10880\\pm 355$ & $ 2.15\\pm\t0.20$&$3.3\t\\pm\t1.2 $   &  11  & 8 & 8d     & $628\\pm 25 $& 25.1&  1907   \\\\ 15 Cnc& 68351    &  5.6 &B9sp &$10290\\pm 340$ & $ 2.65\\pm\t0.21$&$6.6\t\\pm\t2.4 $   &  10  & 16& 1d4m &  $325 \\pm 47$& 6.9 &  762    \\\\ 3 Hya &72968    &  5.7 &A1spe &$9840\\pm  320$ & $ 1.55\\pm\t0.08$&$2.0\t\\pm\t0.4 $   & 10  & 13& 13d     & $427\\pm 16 $& 27.4&  1304   \\\\ 45 Leo& 90569    &  6.0 &A0sp &$10250\\pm 335$ & $ 1.78\\pm\t0.10$&$2.5\t\\pm\t0.5 $   & 11  & 10& 10d     & $541\\pm 23 $& 23.5&  1634   \\\\ &94427    &  7.3 &A5          &$7250\\pm  240$ & $ 1.05\\pm\t0.11$&$2.1\t\\pm\t0.5 $   & 8  & 8& 8d        & $356\\pm 41 $& 8.6 &  904    \\\\ EP Uma &96707 & 6.0 &F0sp     &$7780\\pm  255$ & $ 1.54\\pm\t0.08$&$3.2\t\\pm\t0.6 $   & 8  & 21& 7d7m   &  $69  \\pm 33$& 2.3 &  128     \\\\ 65 Uma& 103498 & 6.9 &A1spe   &$9220\\pm  300$ & $ 2.06\\pm\t0.20$&$4.2\t\\pm\t1.5 $   & 9  & 14& 12d1m  &  $169 \\pm 19$& 8.9 &  432    \\\\ 21 Com& 108945   &  5.4&A2pvar&$8870\\pm  290$ & $ 1.72\\pm\t0.09$&$3.1\t\\pm\t0.6 $   & 9  & 13& 12d1m  &  $234 \\pm 54$& 4.3 &  416    \\\\ $\\omega$ Her &  148112&4.6&B9p&$9330\\pm  305$ & $ 1.86\\pm\t0.08$&$3.2\t\\pm\t0.6 $   & 10  & 12& 11d     & $204\\pm 21 $& 9.7 &  535    \\\\ 45 Her &151525   &  5.2 &B9p  &$9380\\pm  310$ & $ 2.18\\pm\t0.13$&$4.7\t\\pm\t1.2 $   & 11  & 14 & 2d3m &  $146 \\pm 38$& 3.8 &  231    \\\\ &171586 & 6.4&A2pvar          &$8760\\pm  290$ & $ 1.37\\pm\t0.10$&$2.1\t\\pm\t0.5 $   & 10  & 5 & 5d      & $375\\pm 56 $& 6.6 &  868    \\\\ &171782 & 7.8 &A0p            &$9660\\pm  315$ & $ 1.76\\pm\t0.30$&$2.7\t\\pm\t1.3 $  & 10  & 6 & 2d1m   &  $333 \\pm 78$& 4.2 &  584    \\\\ 19 Lyr &179527   &  5.9 &B9sp &$10370\\pm 340$ & $ 2.63\\pm\t0.16$&$6.4\t\\pm\t1.9 $   &   11  & 11& 8d2m&  $156 \\pm 46$& 3.4 &  211    \\\\ 4 Cyg& 183056   &  5.1 &B9sp  &$11710\\pm 385$ & $ 2.69\\pm\t0.11$&$5.3\t\\pm\t1.2 $   & 12  & 13& 13d     & $290\\pm 42 $& 6.9 &  680    \\\\ &204411   &  5.3 &A6pe        &$8750\\pm  290$ & $ 1.97\\pm\t0.07$&$4.2\t\\pm\t0.8 $   &   9  & 12& 12d    & $88 \\pm 14 $& 6.0 &  198    \\\\ $\\kappa$ Psc &220825  &4.9&A0p&$9450\\pm  310$ & $ 1.40\\pm\t0.05$&$1.9\t\\pm\t0.3 $   &   10  & 12& 12d   & $312\\pm 25 $& 12.8&  865    \\\\ \\hline  \\end{tabular} \\end{table*}   \\subsection{Physical properties of the sample}  For each of the sample stars, we have determined effective temperature $T_{\\rm eff}$, luminosity $L$ and radius $R$, to allow us to identify the appropriate line mask for Least-Squares Deconvolution (Sect. 2.5) and for determination of the rotational axis inclination for the dipole magnetic field model (Sect. 2.7).  Effective temperatures of stars of our sample were derived using Geneva and Str\\\"omgren photometry (obtained from the General Catalogue of Photometric Data (GCPD); Mermilliod et al. 1997) using the calibrations of Hauck \\& North (1982) and Moon \\& Dworetsky (1985). Effective temperatures reported in Table 1 are the average of the two estimates when both were available, or the single one which could be derived when only one photometric set was available. We have assumed for $T_{\\rm eff}$ an uncertainty of the order of 3\\% for propagation of uncertainties in all calculations using the effective temperature.  Luminosity was inferred using the GCPD-reported visual magnitude, the Hipparcos parallax and the bolometric correction relations of Balona (1994). Radius was then inferred directly from the luminosity and temperature via the Stefan-Boltzmann equation for a uniform spherical star.  The inferred values of $T_{\\rm eff}$, $\\log L/L_\\odot$ and $R/R_\\odot$ are reported in Table 1. For the 20 stars which we have in common with the study of Kochukhov \\&  Bagnulo (2006), these values are all in good agreement. As pointed out by Landstreet et al. (2007), the assumed uncertainties of  Kochukhov \\& Bagnulo (2006), which are comparable to our own, are probably somewhat underestimated. However, as our fundamental parameters are not to be used for detailed evolutionary studies, we consider them to be sufficient for this study.  \\subsection{Least-Squares Deconvolution and magnetic field detection}  The primary aim of our study is to detect line circular polarisation (a ``Stokes $V$ Zeeman signature'') which is characteristic of the longitudinal Zeeman effect produced by the presence of a magnetic field in the stellar photosphere. For this we used the Least-Squares Deconvolution (LSD) procedure, first used by Donati et al. (1997) to study the magnetic fields of active late-type stars and by Wade et al. (2000 a,b) for Ap stars. This method enables the ``averaging'' of several hundred (and possibly several thousand in some stars) lines and thus to obtain Stokes $I$ and Stokes $V$ profiles with greatly improved S/N.  LSD provides a single quantitative criterion for the detection of Stokes $V$ Zeeman signatures: we perform a statistical test in which the reduced $\\chi^2$ statistic is computed for the Stokes $V$ profile, both inside and outside the spectral line (Donati et al. 1997). The statistics are then converted into detection probabilities, which are assessed to determine if we have a definite detection (dd, false alarm probability smaller than $10^{-5}$), a marginal detection (md, false alarm probability greater than $10^{-5}$ and smaller than $10^{-3}$), or no detection at all (nd). A diagnostic null spectrum (called $N$ in the following) is also obtained using the same subexposures obtained for Stokes $V$, but by pair processing those corresponding to identical azimuths of the quarter-wave plate. By checking that a signal is detected only in $V$ and not in $N$, and that any detected signature is located within the line profile velocity interval, we can distinguish between real magnetic signatures and (infrequent) spurious signatures. In addition, using the full resolved Stokes $V$ profile enables the detection of the magnetic field, even if the integrated line-of-sight component is very weak, or even null.  \\subsection {LSD masks}  LSD is a cross-correlation method which requires comparison of our observed spectra with synthetic line masks (Donati et al. 1997, Shorlin et al. 2002). To obtain the most realistic masks, we used spectral line lists from the Vienna Atomic Line Database (VALD; Piskunov et al 1995; Ryabchikova et al. 1997; Kupka et al. 1999). To take into account the chemical peculiarities of Ap/Bp stars, we employed an abundance table in which the abundance of metals (Al, Si, S, Ti, V, Mn, Fe, Co, Ni, Zn, Sr, Y, Zr, Ba, La, Ce, Pr, Nd, Eu, Gd, Dy) is 10x solar, except for Cr which was increased to 100x solar (e.g. Shorlin et al. 2002). Masks were then compiled for effective temperature ranging from 7000-13000~K, with $\\log g=4.0$, a microturbulence of 2~km/s, and including all metal lines with a central depth greater than 10\\% of the continuum. We also compiled a series of masks assuming solar abundances. We have found in other studies that the magnetic field measurements are not very sensitive to the mask temperature within a couple of thousand K (Wade et al., in preparation). We therefore computed masks spaced every 1000~K.  Least-Squares Deconvolution was performed for several temperatures and in some cases the most significant magnetic field detection was obtained for a temperature somewhat hotter than that given by the photometric data. This occured several times for the hottest sample stars - this can be seen in Table 1. In the case of the cool Ap star EP UMa, a solar abundance mask gave a better result (better detection of magnetic field and smaller error bars on $B_\\ell$) than Ap abundances as described above. These discrepancies probably result from differences between the true chemical peculiarities of individual stars and those assumed in the line masks. For the discrepant hot stars, we computed the longitudinal magnetic field for a mask temperature between that corresponding to the best detection and the derived effective temperature. For EP UMa, we used solar abundance masks in our analysis. The number of lines used in the LSD ranged from 1500 to 3000, and is anticorrelated with the temperature.  \\subsection{Longitudinal magnetic field}  The longitudinal magnetic field was inferred from each of the Stokes $I$ and $V$ profile sets, using the first-order moment method. According to this method, the longitudinal field $B_\\ell$ (in G) is calculated from the Stokes $I$ and $V$ profiles in velocity units as:  \\begin{equation} B_{\\ell} = -2.14 \\times 10^{11} \\frac{\\int_{}^{}vV(v)dv}{\\lambda gc\\int_{}^{}[I_c - I(v)]dv}, \\end{equation}  \\noindent (Rees \\& Semel 1979; Donati et al. 1997; Wade et al. 2000b) where $\\lambda$, in nm, is the mean wavelength of the LSD profile, $c$ is the velocity of light (in the same units as $v$), and $g$ is the mean value of the Land\\'e factors of all lines used to construct the LSD profile. Integration ranges used for evaluation of Eq. (1) were computed automatically, beginning and ending 15~km/s before/after the location in the line wings at which the residual flux was equal to 85\\% of the continuum flux. The accuracy of this technique for determining high-precision longitudinal field measurements has been clearly demonstrated by Wade et al. (2000b), Donati et al. (2001) and Shorlin et al. (2002).  The resultant longitudinal magnetic field measurements, which are reported in Table~3, are remarkably precise (this Table is only available on line). The 282 measurements, with a median $1\\sigma$ uncertainty of 40 G, represent the largest compilation of high-precision stellar magnetic field measurements ever published.   \\subsection{Modeling the longitudinal field variation}  To characterise the dipole components of the magnetic fields of our sample stars, we use the oblique rotator model (ORM, Stibbs 1950) as formulated by Preston (1967). This model provides a good first approximation of the large scale magnetic field of Ap stars (e.g. Landstreet 1988). Because of the weakness of the longitudinal magnetic field observed for the stars of our sample, we do not expect to be able to detect departures from a global dipolar configuration.  To begin, we added to our high-precision data set additional good-quality published magnetic field measurements collected by Bychkov et al. (2003). Details of these collected measurements were kindly provided by Dr. Victor Bychkov. Then, we searched the literature for rotational periods for each of our sample stars. For many stars, published rotational periods were available which provided an acceptable folded phase variation $B_\\ell(\\phi)$ of the magnetic measurements. However, for some stars, the published rotational period or periods did not provide an acceptable folded magnetic field variation, and for others, no published period was available. For these latter stars, we used a modified Lomb-Scargle technique to attempt to infer the rotational period, both directly from the longitudinal field measurements, as well as from the variations of the LSD Stokes $I$ and $V$ profiles. The period searches of LSD profiles were performed by treating each pixel in the Stokes $I$ and $V$ profiles as an independent timeseries (similar to the technique described by Adelman et al., 2002). Individual periodograms were subsequently weighted according to their amplitude of variation and averaged to characterise variability of the whole LSD profile. Acceptable periods were identified by establishing the 99\\% confidence threshold, and candidate periods were evaluated by phasing the LSD profiles and longitudinal field measurements.  Results of the period searches for individual stars are provided in their appropriate subsections in Sect. 3. Ultimately, acceptable rotational periods were obtained for 24 stars, and these periods are reported in Table 2.  We also searched the literature for published values of the projected rotational velocity ($v\\sin i$) of each star, which we compared to the value measured from the LSD Stokes $I$ profile by fitting rotationally-broadened synthetic profiles. Sometimes significant discrepancies were found between our values of $v\\sin i$ and those reported in the literature. These discrepacies are discussed in Sect. 3, and the adopted rotational velocities (generally those obtained from the LSD profiles) are shown in Table 2.  Each phased longitudinal field variation $B_\\ell(\\phi)$ was then fit using a 1$^{\\rm st}$ order sine function:   \\begin{equation} B_\\ell(\\phi) = B_0 + B_1\\sin {2\\pi(\\phi+\\phi_0)}. \\end{equation}   The phased and fit longitudinal field variations are shown in Fig. 1. The reduced $\\chi^2$ of this fit ($\\chi^2_2$), along with those of linear fits through $B_\\ell=0$ (the ``null field'' model, $\\chi^2_0$ ) and through the weighted mean of the measurements (the ``constant field'' model, $\\chi^2_1$),  of each star, are reported in Table 2.  $\\chi^2_{\\rm lim}$,  the $2\\sigma$ upper limit for admissible models, computed according to Press et al. (1992), is reported in Table 2 as well . A comparison of these reduced $\\chi^2$ values with each other allows us to evaluate the significance of the detection of the longitudinal magnetic field and its variability. Although variability of the longitudinal field cannot be established for a few stars ($\\chi^2_1<\\chi^2_{\\rm lim}$), the only star for which the longitudinal field is not detected with more than 2$\\sigma$ confidence is HD 96707 (for which $\\chi^2_0<\\chi^2_{\\rm lim}$).  \\begin{table*}[ht] \\caption{Results of the magnetic field modeling. The contents of the columns are described in sect. 2.7. The uncertainties associated with the derived dipole parameters $i, \\beta$ and $B_{\\rm d}$ correspond to 2$\\sigma$.} \\begin{tabular}{lcc|cccc|cccccc} \\\\ \\hline star     &   Period   &  $v\\sin i$& $\\chi^2_0$ & $\\chi^2_1$ &$\\chi^2_2$ &$\\chi^2_{\\rm lim}$& $i$ &$\\beta$& $B_{\\rm d}$& $B_{\\rm d}^{\\rm min}$& $B_{\\rm d}^{\\rm max}$\\\\ & (d)             & (km/s)           &            &            &           &                  & ($\\degr$) & ($\\degr$) & ($10^3$ G)       & ($10^3$ G)                  & ($10^3$ G)\\\\ \\hline 8441  & 69.2         & 2  &    13.71 &     6.69 &     1.75 &3.35 &$  49\\pm33$ & $ 73_{- 61}^{+ 17}$ &      0.683 &      0.415 &     2.931 \\\\ 10221 & 3.15459      & 24 &     5.90 &     1.81 &     1.56 &2.71 &$  27\\pm11$ & $ 42_{- 41}^{+ 38}$ &      0.375 &      0.195 &     1.202 \\\\ 15089 & 1.74033      & 48 &    57.83 &    56.37 &     1.83 &2.14 &$  45\\pm11$ & $ 80_{- 12}^{+  7}$ &     2.031 &     1.560 &     2.999 \\\\ 15144  & 2.99787     & 13 &  1305.83 &     4.57 &     0.04 &2.71 &$  24\\pm 8$ & $  9_{-  3}^{+  6}$ &     2.100 &     2.007 &     2.281 \\\\ 27309 & 1.568884     & 57 &   157.42 &     2.43 &     1.74 &2.63 &$  53\\pm17$ & $  5_{-  5}^{+ 11}$ &     3.673 &     2.325 &     8.022 \\\\ 32549 & 4.6393       & 47 &     7.19 &     8.56 &     1.75 &2.75 &$  65\\pm27$ & $ 77_{- 74}^{+ 11}$ &      0.546 &      0.312 &    28.176 \\\\ 32650 & 2.7347       & 30 &     4.48 &     1.92 &     1.19 &1.72 &$  37\\pm12$ & $ 45_{- 30}^{+ 31}$ &      0.229 &      0.153 &      0.477 \\\\ 39317 & 2.6541       & 45 &     4.97 &     1.84 &     2.24 &3.18 &$  36\\pm12$ & $ 20_{- 20}^{+ 69}$ &      0.560 &      0.113 &     2.252 \\\\ 43819 & 15.02        & 10 &   135.66 &    59.32 &     3.45 &4.18 &$  63\\pm66$ & $ 42_{- 42}^{+ 47}$ &     2.626 &     2.488 &    78.367 \\\\ 68351 & 4.16         & 33 &     7.85 &     2.17 &     1.38 &2.00 &$  28\\pm37$ & $ 46_{- 41}^{+ 36}$ &      0.649 &      0.437 &    71.486 \\\\ 72968 & 5.6525       & 16 &   360.16 &     4.51 &     1.22 &2.03 &$  61\\pm18$ & $  5_{-  4}^{+  7}$ &     2.388 &     1.451 &     6.702 \\\\ 90569 & 1.04404      & 13 &   170.38 &    36.08 &     1.93 &3.07 &$   9\\pm 4$ & $ 81_{-  7}^{+  5}$ &     5.157 &     2.946 &    11.284 \\\\ 94427  & 1.9625      &  8 &    37.58 &    46.61 &     2.12 &3.72 &$   8\\pm 4$ & $ 89_{-  4}^{+  1}$ &     8.957 &     3.806 &    25.519 \\\\ 96707  & 3.515       & 37 &     0.82 &     0.91 &     0.77 &1.21 &$  53\\pm16$ & $ 90_{- 90}^{+  0}$ &      0.100 &        0.0 &      0.492 \\\\ 103498 & 15.830      & 13 &    25.43 &    29.71 &     6.55 &7.44 &$  75\\pm68$ & $ 80_{- 11}^{+ 10}$ &      0.600 &      0.572 &     6.751 \\\\ 108945 & 2.01011     & 65 &     5.31 &     6.00 &     2.10 &2.83 &$  57\\pm18$ & $ 85_{- 61}^{+  3}$ &      0.735 &      0.333 &     1.509 \\\\ 148112 & 3.04296     & 44.5 &  40.67 &     0.93 &     0.95 &1.73 &$  56\\pm16$ & $  3_{-  3}^{+ 11}$ &     1.042 &      0.579 &     2.370 \\\\ 151525 & 4.1164      & 35 &     2.74 &     2.86 &     0.97 &1.70 &$  37\\pm19$ & $ 78_{- 43}^{+ 11}$ &      0.545 &      0.208 &     1.927 \\\\ 171586 & 2.1308      & 37 &    10.86 &     8.27 &     2.59 &6.60 &$  48\\pm19$ & $ 46_{- 40}^{+ 40}$ &     1.422 &      0.716 &     4.413 \\\\ 171782 & 4.4674      & 24 &     6.88 &     1.12 &     1.79 &3.79 &$  51\\pm51$ & $  5_{-  5}^{+ 85}$ &     1.651 &      0.213 & 22257.276 \\\\ 179527 & 7.098       & 33 &     6.57 &     7.70 &     0.30 &1.30 &$  74\\pm32$ & $ 81_{- 34}^{+  8}$ &      0.522 &      0.409 &     1.233 \\\\ 183056 & 2.9919      & 26 &    25.02 &    28.74 &     1.59 &2.32 &$  74\\pm 8$ & $ 49_{- 32}^{+ 37}$ &     1.558 &     1.172 &     3.938 \\\\ 204411 & 4.8456      & 5.4&    20.75 &     5.43 &     0.57 &1.46 &$   7\\pm 5$ & $ 81_{- 12}^{+  7}$ &      0.968 &      0.416 &     4.509 \\\\ 220825 & 1.42539     & 38 &    78.54 &    90.05 &     2.29 &3.18 &$  35\\pm54$ & $ 83_{- 80}^{+  7}$ &     1.957 &     1.141 &    21.045 \\\\ \\\\ \\hline \\end{tabular} \\end{table*}   For a tilted, centred magnetic dipole, the surface polar field strength $B_{\\rm d}$ is derived from the variation of the longitudinal magnetic field $B_\\ell$ with rotational phase $\\phi$ using Preston's (1967) well-known relation:  \\begin{equation} B_d = B_\\ell^{\\rm max} \\biggl({{15 + u}\\over{20(3-u)}} (\\cos\\beta\\cos i + \\sin\\beta\\sin i)\\biggr)^{-1}, \\end{equation}   \\noindent where $B_\\ell^{\\rm max}=|B_0|+B_1$ and $u$ denotes the limb darkening parameter (equal to approximately $u=0.5$ for our sample).  The rotational axis inclination and obliquity angles $i$ and $\\beta$ are related by  \\begin{equation} \\tan\\beta={{1-r}\\over{1+r}} \\cot i, \\end{equation}  \\noindent where $r=(|B_0|-B_1)/(|B_0|+B_1)$.  We have determined the inclination $i$ for each of our stars assuming rigid rotation, and computing:  \\begin{equation} \\sin i = {{P_{\\rm rot} v\\sin i}\\over {50.6 R}}, \\end{equation}  \\noindent where $P_{\\rm rot}$ is the adopted stellar rotational period in days, $v\\sin i$ is the adopted projected rotational velocity in km/s, and $R$ is the computed stellar radius in solar units. The magnetic obliquity $\\beta$ was then inferred from Eq. (4) and the polar strength of the dipole $B_{\\rm d}$ from Eq. (3). Uncertainties associated with all parameters (at 2$\\sigma$, including a max$(2~{\\rm km/s}, 10\\%)$ uncertainty on $v\\sin i$) were propagated through the calculations of $i$, $\\beta$ and $B_{\\rm d}$. The resultant dipole magnetic field models are reported in Table 2.  An important and interesting consequence of Eq. (3), independent of all parameters except limb-darkening, is:  \\begin{equation} B_{\\rm d} \\geq {{20(3-u)}\\over{15+u}} B_\\ell^{\\rm max}, \\end{equation}  \\noindent which, for typical limb-darkening $u$, yields:  \\begin{equation} B_{\\rm d} \\gtrsim 3.3 B_\\ell^{\\rm max}. \\end{equation}  Therefore, exclusive of the model geometry, a lower limit on the surface dipole component of the magnetic field is obtained directly from the maximum measured value of the longitudinal field. The maximum measured longitudinal field is reported in column 11 of Table 1, and the inferred lower limit ($2\\sigma$) of $B_{\\rm d}$,  $B_{\\rm d}^{\\rm min, 3.3}$ is reported in column 13 of Table 1. It is clear from these data that $B_{\\rm d}$ for most of our sample is larger than a few hundred G at $2\\sigma$. ",
        "conclusions": "We have investigated a sample of 28 well-known spectroscopically-identified magnetic Ap/Bp stars, obtaining 282 new Stokes $V$ Zeeman signatures and longitudinal magnetic field measurements using the MuSiCoS spectropolarimeter. Magnetic field is detected in all sample stars, and the inferred longitudinal fields are significantly greater than some tens of G. To characterise the surface magnetic field intensities of the sample, we modeled the longitudinal field data to infer the intensity of the dipolar field component.  The distribution of derived dipole strengths for these stars exhibits a plateau at about 1 kG, falling off to larger and smaller field strengths. Remarkably, in this sample of stars selected for their presumably weak magnetic fields, we find only 2 stars for which the derived dipole strength is weaker than 300~G. Interpreting this ``magnetic threshold'' as a critical value necessary for the stability of large-scale magnetic fields leads to a natural explanation of the small fraction of intermediate-mass magnetic stars. It may also explain the near-absence of magnetic fields in more massive B and O-type stars."
    },
    "0710/0710.3883_arXiv.txt": {
        "abstract": " ",
        "introduction": "Current observational cosmology allows us to test fundamental physics with a continuously improving precision. To evaluate and understand these data, theoretical models about the evolution of our universe are crucial. One promising paradigm receiving growing experimental support is cosmological inflation \\cite{original_inflation}. Inflation postulates a period of exponential expansion of the universe driven by scalar fields slowly rolling in an almost flat potential. This enormous growth stretches quantum fluctuations present in the early universe to currently observable astrophysical scales. The imprints of such a process can be found, for example, in the cosmic microwave background and the large scale structure of the universe \\cite{WMAP,Tegmark, Sanchez}.  In recent years much effort has focused on the realization of inflation within string theory \\cite{Cosmo_rev,Hertzberg:2007ke}. In string theory there are potentially many scalar fields which could drive inflation and hence different opportunities to model inflation. Specific scenarios include realizations of K\\\"ahler moduli inflation~\\cite{Banks:1995dp,Conlon:2005jm}, racetrack models \\cite{BlancoPillado:2006he} and the most intensively studied possibility of D-brane inflation \\cite{Dvali:1998pa,Cosmo_rev}. However, as of today, it remains challenging to establish explicit scenarios in a controlled compactification without employing extreme fine-tuning to obtain a sufficiently flat potential \\cite{BDKMcAS}. Having found a realization of inflation reproducing the current cosmological observables, it is important to establish which models can incorporate possible future observations. For example, as argued in refs.~\\cite{BMcA}, many string scenarios do not allow for a high ratio $r$ of gravitational waves produced in the early universe. The current experimental bound on gravitational waves is $r<0.3$ \\cite{Tegmark}, but future experiments, including Planck, BICEP and Spider \\cite{Efstathiou:2006ak}, might allow the observation of $r$ with a precision down to $r >0.01$. It is thus desirable to study string embeddings of inflation which can incorporate an $r$ observable in these experiments. Recent attempts to do that can be found, for example, in refs.~\\cite{Olsson:2007he,Krause:2007jr,Becker:2007ui,Kobayashi:2007hm, KSS}.   A possible scenario able to incorporate primordial gravitational waves was suggested by Dimopoulos, Kachru, McGreevy and Wacker \\cite{DKMW}.  The authors argue for an embedding of multi-field inflation with a large number $N$ of axion fields. Such models use an assistance effect studied in refs.~\\cite{LMS,KO} ensuring that the small fraction $1/N$ controls the flatness of the potential. Indeed, generic compactifications of type II string theory on a six-dimensional manifold can admit  $10^4$, or more, axions from the NS-NS B-field and the R-R form fields. An appropriate subset of $N$ such axions were proposed to  drive inflation in ref.~\\cite{DKMW} and the authors termed these scenarios $N$-flation. For a sufficiently large $N$ the scenarios can be interpreted as a realization of natural inflation \\cite{Freese:1990rb,Kim:2004rp}. If the inflatons can produce the desired amount of e-foldings already in the quadratic regime of the potential the models admit chaotic inflation \\cite{Linde:1983gd}. This implies that $N$-flation can yield a possibly observable signature of gravitational waves with $r<0.14$ and thus distinguishes it from most other string realizations of inflation.   The aim of this paper is to study a specific realization of $N$-flation in type IIB string theory. The inflating axions correspond to the zero-modes of the R-R forms in the compactification to four space-time dimensions. In compactifications preserving $\\cN=2$ supersymmetry the R-R axions sit  in the same supermultiplets as the K\\\"ahler structure moduli parametrizing the volumes of two-dimensional cycles in the internal space. Manifolds with a large number of  non-trivial two-cycles are thus candidate backgrounds for $N$-flation. As will be shown, the density perturbations and slow roll parameters depend on the volume of the compact space and it has to be ensured that they do not become large with increasing $N$. We will thus argue that explicit examples always involve compact manifolds which admit many very small or vanishing cycles.  In the presence of small cycles stringy effects become important and significantly alter the structure of the four-dimensional effective theory. In order to analyze these contributions we will concentrate on axions arising from the R-R two-form evaluated on  vanishing two-cycles of the compact geometry. The standard example of a vanishing two-cycle is the resolved conifold \\cite{Candelas:1989js}. In this case a conical singularity is resolved by a two-sphere supported by a geometric volume or an NS-NS B-field. If this $S^2$ becomes smaller than the square string length, world-sheets will start  to wrap and significantly contribute to the metric of the R-R axions. Fortunately, in $\\cN=2$ compactifications these corrections can be computed for the conifold and many other Calabi-Yau geometries \\cite{Hori:2003ic,Hosono:1993qy} allowing the evaluation of the metric of the axions. We will illustrate this general fact on a toy model with $N$ conifold singularities. For such examples it can be shown that the kinetic terms of the axions can be close to the Planck scale which is crucial to obtain inflation \\cite{BDFG}.  In addition to the fundamental strings also D-branes can become relevant in geometries with vanishing cycles \\cite{Strominger:1995cz}. In particular, D1 instantons can wrap the small cycles and correct the metric of the R-R axions. Such contributions appear with the exponential of the D1 instanton action which depends on the R-R two-form axions themselves. In contrast to the string world-sheet action the D-instanton action also contains a factor of the inverse string coupling. This implies that for small string coupling and finite volume or B-field of the vanishing cycles the D1 instanton corrections are subleading in the axion metric. However, correction due to D-branes will be the leading contributions in the scalar potential and can induce the desired potential for the R-R axions.   To study the scalar potential for the axions and non-axionic moduli we will focus on $\\cN=1$ orientifold compactifications of type IIB string theory. Such compactifications have been studied intensively in the last years \\cite{review_flux}. It was shown that $\\cN=1$ potentials can be induced by background fluxes, D-brane instantons or gaugino condensates on space-time filling D-branes. The desired axion potentials can arise through non-trivial superpotentials from either of the three sources \\cite{review_flux}. We will briefly discuss their properties in various orientifold compactifications: D1 superpotentials in type I~\\cite{Witten2}, D1 dependences through the determinants in the D3 instanton superpotentials \\cite{Witten3,Ganor,TG}, and gaugino condensates on space-time filling D5 branes \\cite{review_flux,Vafa:2000wi,Cachazo:2001jy,Dijkgraaf:2002dh}. Remarkably, explicit computations of the non-perturbative superpotentials can often be performed in a dual flux picture where geometry dictates the form of the corrections \\cite{Vafa:2000wi,Cachazo:2001jy,Dijkgraaf:2002dh,Heckman:2007ub,ABK}.   In the final part of this work we will study the effective theory of orientifold compactifications with O3 and O7 planes in more detail.  Using earlier results \\cite{GL1,GL2} we argue that the $\\cN=2$ world-sheet corrections are inherited by the $\\cN=1$ orientifold theory.  In particular, they correct the $\\cN=1$ K\\\"ahler potential and complex coordinates in a calculable way. This general fact can be applied to a simplistic compact toy model with $2N$ conifolds pairwise identified under the orientifold projection. Including a flux and D-instanton superpotential, we make first steps in establishing an effective theory with a large number of light axions and all other moduli stabilized. An explicit numerical evaluation indicates the presence of a non-supersymmetric axion valley \\cite{Kallosh}. This ensures the desired mass hierarchy between the axions and their non-axionic partners.  The paper is organized as follows. In section \\ref{rev_N} we review the $N$-flation scenario of \\cite{DKMW} and discuss some of its cosmological implications. We recall that the kinetic terms of the axions, set by the axion decay constants, have to be large in order to obtain inflation. A discussion of axion decay constants in Type IIB string theory is presented in section~\\ref{axion_decay}. The four-dimensional effective Lagrangian and the general form of the axion decay constants are studied in section \\ref{four-dim_Lagr} supplemented with appendix \\ref{gen_axion_const}. In section \\ref{large_vol} we argue that in compactifications with all cycles larger than string length, the axion decay constants typically become very small with an increasing number of axions. This implies that only compactification manifolds with small or vanishing cycles are candidate backgrounds to obtain $N$-flation. The quantum corrected axion decay constants for the resolved conifold and geometries with $N$ resolved singularities are discussed in sections \\ref{res_cone} and \\ref{multiax}. Non-perturbative D-brane effects can induce a scalar potential through a non-vanishing superpotential as discussed in section \\ref{axion_potential}. In section \\ref{D1_corrections} it is shown that such a superpotential can arise from D1 instanton corrections, while section~\\ref{D5_gaugino} discusses superpotentials originating from gaugino condensates on D5 branes. In the final part of this work, section \\ref{N=1_ori}, the embedding of $N$-flation into a concrete $\\cN=1$ orientifold compactification is addressed. The general form of the $\\cN=1$ data including the inherited $\\cN=2$ perturbative and non-perturbative string world-sheet corrections is presented in section~\\ref{four-dim_N=1} and appendix~\\ref{N=1_recall}. In section~\\ref{N_conifolds}, a toy model with $N$ conifold pairs is used to illustrate that the string world-sheet corrections in the K\\\"ahler potential and the presence of a non-perturbative superpotential can ensure an effective theory with light axions. This indicates the possibility of $N$-flation in string theory. ",
        "conclusions": "In this paper we discussed the possibility of a string theoretical embedding of cosmological inflation driven by a large number of axions. Such scenarios use the fact that in the dimensional reduction of string theory  to four space-time dimensions a vast number of scalar fields arise in the low energy theory. In particular, we considered type IIB string theory on a compact Calabi-Yau manifold with many non-trivial two-cycles. To each of these two-cycles an axion from the R-R two-form and an axion from the R-R four-form can be associated. The effective theory for the axionic fields will sensitively depend on the values of the geometric moduli, i.e.~the volume of the two-cycles, as well as the NS-NS B-field. We have argued that a possible realization of axion inflation might only exist in special corners in the landscape of vacua. In these regimes various stringy effects become relevant and have to be included. In this work, we have made first steps one region in the $\\cN=1$ landscape, where axion inflation might be realized.  In recent years, $\\cN=1$ type IIB compactifications with all volumes stabilized at scales much larger than string scale have been investigated intensively \\cite{review_flux}. For vacua in this regime of the field space, stringy corrections, such as wrapping stings and D-branes, play a sub-leading role in the derivation of the kinetic terms of the axions. Non-perturbative D-brane effects can induce a potential exponentially suppressed by the large volume of the cycles times the inverse of a small string coupling. We have argued that such scenarios are not suitable for axion inflation. They naturally admit very small axion decay constants which become even smaller if the number of axions increases. Moreover, due to the strong exponential suppression, the scale of inflation is typically too low for semi-realistic scenarios. This conclusion lead us to the consideration of manifolds with small or vanishing cycles.  In compactifications with cycles smaller than string scale the caveats of the large volume compactifications can be avoided. To obtain such geometries, we considered resolutions of singularities supported by a volume and a NS-NS B-field. The standard example of such a blown-up singularity is the resolved conifold. In the vicinity of small cycles, new corrections will become relevant and alter the effective theory. We have discussed a subset of such correction in $\\cN=2$ compactifications of type IIB string theory on a Calabi-Yau manifold. Including the leading singular non-perturbative string world-sheet contributions, we have argued that the axion decay constants can take values close to the Planck scale also for scenarios with many axions. In $\\cN=1$ compactifications further perturbative and non-perturbative corrections can become relevant and might alter the structure of the effective four-dimensional theory. However, in case these do not cancel the $\\cN=2$ effects, large axion decay constants will remain accessible also in these less supersymmetric scenarios.  In addition to being close to the Planck scale, $N$-flation also requires the axion decay constants to be independent of the axions themselves. Corrections depending on the R-R axions are believed to only arise from non-perturbative D-brane effects. If the small cycles remain to be of finite size and we work at sufficiently small string coupling, D-brane corrections are subleading in the axion decay constants, since the instanton action is larger by a factor of the inverse string coupling. This will also be the case for vacua in an $\\cN=1$ supergravity theory. In a consistent analysis, the vacuum values of the $\\cN=1$ moduli fields have to be inside an appropriate field range to ensure that D-brane instantons are subleading in the K\\\"ahler potential encoding the kinetic terms of the scalar fields.  One of the remaining complications in realizing $N$-flation is to ensure that there indeed exists an effective theory for axions with appropriate masses during inflation. These masses have to be lower than the masses of their non-axionic partner and other bulk moduli fields, but still sufficiently large to match the observed density perturbations during inflation. Again, this forces us to work away from the large volume regime, where at least some of the moduli masses are suppressed by large instanton actions. A careful investigation of the effective potential for the moduli is necessary to study the vacua of the theory. In this work we briefly discussed effective $\\cN=1$ potentials arising from D-instantons and gaugino condensates on space-time filling branes. We pointed out that for potentials induced by large rank gaugino condensates on D5 branes, the quadratic region of the axion potential can stretch over the entire accessible field range. In practice, D1 corrections arising through the pre-factors of the D3-instantons are of particular interest. If these are the leading axion-dependent contributions to the $\\cN=1$ superpotential, a mass hierarchy can be ensured due to the suppression by both the D1 and D3 instanton action.  To investigate the properties of the effective theory more explicitly, we discussed the embedding of axion inflation into $\\cN=1$ Calabi-Yau orientifold compactifications with O3 and O7 planes. We showed that the $\\cN=1$ characteristic data remain calculable even in the case that the non-perturbative $\\cN=2$ string world-sheet corrections are included. Utilizing a flux superpotential together with a superpotential from non-perturbative D-brane effects a potential for all bulk moduli fields is generated. We illustrated that a theory of light axions could exist, if the axion dependence is suppressed in the superpotential and the non-axionic partners of the axions in the $\\cN=1$ chiral multiplet is stabilized due to the string world-sheet corrections to the K\\\"ahler potential. In an optimistic scenario, one can hope that such an effective theory of a large number of relatively light axions from the R-R forms will survive also further perturbative corrections.  Even though an explicit embedding of $N$-flation into string theory still remains to be constructed, the scenarios outlined and studied in this work might provide a promising route to achieve this goal. Likely, such an embedding will not solve intrinsic issues related to the fine-tuning of initial conditions in chaotic and natural inflation with many inflatons. However, it might provide a way to accommodate possible future observations of primordial gravitational waves in a string theoretic model."
    },
    "0710/0710.5188_arXiv.txt": {
        "abstract": "By-now photons are the unique universal messengers. Cosmological sources like far-away galaxies or quasars are well-known light-emitters. Here we demonstrate that the nonlinear electrodynamics  (NLED)  description of photon propagation  through the weak background intergalactic magnetic fields modifies in a fundamental way the cosmological redshift that a direct computation within a specific cosmological model can abscribe to a distant source. Independently of the class of NLED Lagrangian, the effective redshift turns out to be $1 + \\tilde{z} = (1 + z)~\\Delta$, where $\\Delta \\equiv (1 + \\Phi_e)/(1 + \\Phi_o)$, with $\\Phi \\equiv {8}/{3} ({L_{FF}}/{L_F}) B^2$, being $L_F = {dL}/{dF}$, $L_{FF} = {d^2L }/{dF^2}$, the field $F\\equiv F_{\\alpha \\beta} F^{\\alpha \\beta}$, and $B$ the magnetic field strength. Thus the effective redshift is always much higher then the standard redshift, but recovers such limit when the NLED correction $\\Delta(\\Phi_e, \\Phi_o) \\longrightarrow 1$. This result may provide a physical foundation for the current observation-inspired interpretation that the universe undergoes an accelerate expansion. However, under the situation analyzed here, for any NLED the actual (spatial) position of the light-emitting far-away source remains untouched. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.2761_arXiv.txt": {
        "abstract": "{(Ultra)Luminous Infrared Galaxies [(U)LIRGs] are much more numerous at cosmological distances than locally, and are likely the precursors of elliptical galaxies. Studies of the physical structure and kinematics of representative samples of these galaxies at low redshift are needed in order to understand the interrelated physical processes involved. Integral field spectroscopy (IFS) offers the possibility of performing  such a detailed analysis.} {Our goal is to carry out IFS of 42 southern systems which are part of a representative sample of about 70 low redshift (z $\\le$ 0.26) (U)LIRGs, covering different morphologies from spirals to mergers over the entire infrared luminosity range (L$_{IR}$=10$^{11}-10^{12.6}$L$_{\\odot}$).} {The present study is based on optical IFS obtained with the VIMOS instrument on the VLT.} {The first results of the survey are presented here with the study of two galaxies representative of the two major morphological types observed in (U)LIRGs, interacting pairs and morphologically- regular, weakly-interacting spirals, respectively. We have found that IRAS F06076$-$2139 consists of two low-intermediate mass (0.15 and 0.4 m$_*$) galaxies with relative velocities of $\\sim$  550 km s$^{-1}$ and, therefore, it is unlikely that they will ever merge. The VIMOS IFS has also discovered the presence of a rapidly expanding and rotating ring of gas in the southern galaxy (G$_s$). This ring is interpreted as the result of a nearly central head-on passage of an intruder about 140 million years ago. The mass, location and relative velocity of the northern galaxy (G$_n$) rules out this galaxy as the intruder. IRASF 12115$-$4656 is a spiral for which we have found a mass of 1.2 m$_*$. The ionized gas shows all the kinematic characteristics of a massive (8.7 $\\times$ 10$^{10}$M$_{\\odot}$), fast rotating disk. The neutral gas traced by the NaI doublet shows distinct features not completely compatible with pure rotation. The neutral and ionized gas components are spatially and kinematically decoupled. The analysis presented here illustrates the full potential of IFS in two important aspects: (i) the study of the kinematics and ionization structure of complex interacting/merging systems, and (ii) the study of the kinematics of the different gas phases, neutral (cool) and ionized (warm), traced by the NaD and H$\\alpha$ lines, respectively.} {} ",
        "introduction": "Most of the observational efforts devoted in recent years to the study of galaxy formation and evolution have been focused on obtaining and analyzing large surveys of different characteristics such as, for instance,   2MASS (Skrutskie et al. 2006), SDSS (Adelman-McCarthy et al. 2006), 2DF (Colless et al. 2001), UDF (Beckwith et al.  2006),  GOODS (Giavalisco et al. 2004),  COSMOS (Capak et al. 2007),  GEMS (Rix et al. 2004),  etc.  Thanks to these studies observational constrains on integrated properties such as  the galaxy luminosity functions, galaxy (stellar) masses,  sizes,  spectral energy distributions,  etc., as well as their evolution with redshift,  have drastically improved (see, for instance, Rix et al. 2006 and references therein).  However, the observational limits imposed by the HST and 10 meter class telescopes have been reached and a substantial new improvement using this methodology will probably have to wait until a new generation of telescopes (JWST, ELT, etc) comes into operation.  In addition to those integrated properties, a comprehensive picture of galaxy formation should explain the internal structure of galaxies (i.e., internal kinematics, internal stellar population gradients, dust distribution, ionization structure,  nuclear properties and effects, interaction with the IGM, etc).  Such detailed studies are more demanding in terms of 'data quality per object'  requiring high S/N two dimensional spectroscopic information with high angular and spectral resolution. Thanks to the advent and popularization of Integral Field Spectroscopic instruments it is now possible to obtain high quality two-dimensional spectroscopy (i.e. 3D data) over a wide wavelength range ($>$ 1000 $\\AA$) in one spectral setting. Although most such studies have been based on local samples of galaxies (e.g.,  Mediavilla et al. 1997; Peletier et al. 1999, 2007; del Burgo et al. 2001; de Zeeuw et al. 2002 and references therein), recent work also includes high-z studies (e.g., Smith et al. 2004; Forster-Schreiber et al. 2006; Flores et al. 2006; Puech et al. 2007; Law et al. 2007). However, limitations in sensitivity and angular resolution make the studies of the internal structure of high-z galaxy populations via IFS extremely difficult, and the observation of comprehensive samples (i.e., not just individuals at the tip of the luminosity function) will require larger aperture telescopes, and higher angular resolution than those currently available.  In the context of the high-z universe, the study of representative samples of some low-z galaxies populations are of particular importance. That is the case of the luminous and ultraluminous infrared galaxies  (LIRGs,  L$_{IR}$=L[8-1000$\\mu$m]=10$^{11}$-10$^{12}$L$_{\\odot}$, and   ULIRGs, L$_{IR}$$>$10$^{12}-10^{13}$L$_{\\odot}$, respectively). These objects are thought to be the local counterparts of cosmologically important galaxy populations at z $>$ 1 such as the submillimiter galaxies (SMG; e.g., Frayer et al. 2004), and the majority of the infrared sources found in recent {\\it Spitzer} cosmological surveys  (e.g., P\\'erez-Gonz\\'alez et al. 2005).  ULIRGs and LIRGs  have large amounts of gas and dust, and are undergoing an intense star formation in their (circum)nuclear regions (e.g. Scoville et al. 2000, Alonso-Herrero et al. 2006, and references therein). This starburst activity is believed to be their major energy reservoir, although AGN activity may also be present (Genzel et al. 1998).   In many cases (especially for ULIRGs) these objects show clear signs of an on-going merging process producing moderate-mass (0.1$-$1 m$_*$) ellipticals as the end product of the merger (Dasyra et al. 2006, Genzel et al. 2001).  Although we have a qualitative picture, details on how the physical and dynamical processes involved are interrelated and how they determine the internal structure of these cosmologically important objects are far from being well understood.  We have already started a program aimed at studying the internal structure and kinematics of a representative sample of low-redshift LIRGs and ULIRGs using optical and near-IR Integral Field Spectroscopy. These objects are natural laboratories that allow to investigate in detail the internal structure of these galaxy populations in a high S/N, high spatial resolution regime, capturing their complex two dimensional structure. We are mainly analyzing rest-frame optical and near-IR spectral diagnostic features on linear scales of about 1 kpc, that will appear in the near- and mid-IR in the high-z galaxy populations to be studied with future instruments such as the NIRSpec and MIRI instruments on board the JWST. This program initially has been focused on the most luminous objects, the ULIRGs  (e.g., Colina, Arribas \\& Monreal-Ibero, 2005 and references therein; Garc\\'{\\i}a-Mar\\'\\i n 2007) and now is being extended towards lower luminosities, i.e. the LIRG luminosity range.  The present paper is the first of a series of studies based on a representative sample of LIRGs observed with VLT-VIMOS. With $\\sim$ 2000 spectra per object, this study involves a large analysis effort. In this first paper we describe the sample, the experimental set-up used for all the observations with VIMOS, as well as the reduction procedures and analysis methods. The analysis of two of the galaxies representative of the sample is shown to illustrate and emphasize the power of the survey in several areas, which after completion will include, among others, the following: (i) two dimensional kinematics of different phases (neutral and ionized) of the interstellar gas as a function of luminosity and morphology, (ii) two dimensional  ionization structure, and the interplay between ionization and kinematics (i.e., the role of shocks), and (iii) star formation in the (circum)nuclear and external regions (candidate Tidal Dwarf Galaxies and/or precursors of globular clusters)  in interacting/merging systems.  Throughout the paper we will consider H$_{0}$=70 km~s~$^{-1}$Mpc$^{-1}$, $\\Omega_{M}$= 0.7, $\\Omega_{\\Lambda}$= 0.3. ",
        "conclusions": "This paper presents the first results obtained from the integral field spectroscopic survey of local luminous and ultraluminous infrared galaxies being carried out with the VIMOS instrument on the VLT.  This sample includes 42 local  (U)LIRGs systems (65 galaxies) and it is part of a larger representative sample of about 70 systems (z $<$ 0.26) also observed from the northern hemisphere, which covers the different morphologies from spirals to mergers and spans the entire luminosity range L$_{IR}$=10$^{11}-10^{12.6}$L$_{\\odot}$.  The two galaxies analysed here (IRAS F06076$-$2139 and IRAS F12115$-$4656) have been selected as prototypes of the most common morphological types among (U)LIRGs, i.e.  interacting pairs and regular/weakly-interacting spirals, while also covering a wide range in IR luminosity.  A) The VIMOS data of IRAS F06076$-$2139 clearly illustrate the potential of IFS in disentangling the complex ionizing and kinematic structure detected in pairs or mergers. The main conclusions for this system are:  a1) We have found that this system is formed by two independent low-intermediate mass galaxies, the G$_n$ (northern)  and G$_s$(southern), of 0.4 m$_*$  and 0.15 m$_*$, respectively, and  with a relative projected velocity of about 550 km s$^{-1}$. Therefore, it is unlikely that these galaxies will ever merge. G$_n$ is characterized by an AGN-like ionization, and large velocity dispersions ($\\geq$ 100 km s$^{-1}$), while G$_s$ has HII characteristics and lower velocity dispersions  (by at least a factor of two)  than  G$_n$.  a2) The morphology and velocity field of the ionized gas in G$_s$ is consistent with the presence of an expanding, rotating ring of about 20 kpc (deprojected major axis) in size. This expanding ring is interpreted as the result of a head-on collision with an intruder that took place some 1.4 $\\times$ 10$^{8}$ years ago. The characteristics of G$_n$ rule out this galaxy  as the intruder.  a3) A chain of H$\\alpha$ emitting regions kinematically associated with G$_s$ at distances of about 12 kpc have been detected. These regions are  identified as young, tidally-induced star-forming regions (some may be candidate tidal dwarf galaxies) formed as a consequence of a past interaction.   B) The VIMOS data of IRAS F12115$-$4656 show the potential of IFS in the kinematic study of the different phases of the interstellar gas: the neutral gas traced by the NaI absorption doublet, and the ionized gas traced by the H$\\alpha$ or [NII] emission lines. The main conclusions for this system are:  b1) The neutral and ionized gas morphology show a different structure indicating that they have a different spatial distribution. The neutral gas closely follows the stellar distribution while the ionized gas traces mostly the location of HII regions in the circumnuclear and external regions of the galaxy.  b2) The ionized gas shows typical kinematic characteristics of massive rotating disks,  with a large amplitude (peak-to-peak of 520 km s$^{-1}$), a maximum of the velocity dispersion at the kinematic/light centre, and a major kinematic axis coincident with the stellar photometric axis. A dynamical mass of 8.7 $\\times$ 10$^{10}$ M$_{\\odot}$ is measured inside a region of 6 kpc radius. The neutral gas, although consistent  on a large scale with a rotating system, shows clear departures from pure rotation and it is dynamically hotter than the ionized gas, being more affected by random motions.  b3) Despite being dominated by rotation the ionized gas shows large-scale (10 kpc),  low-amplitude (20-30 km s$^{-1}$) non-rotational motion, which spatially correlates  with regions of increased velocity dispersion.  The neutral gas also exhibits an even clearer spatial correlation between departures from rotation and high velocity dispersion values."
    },
    "0710/0710.5836_arXiv.txt": {
        "abstract": "{ {\\it Aims.} Tidally locked rotation is a frequently applied assumption that helps to measure masses of invisible compact companions in close binaries. The calculations of synchronization times are affected by large uncertainties in particular for stars with radiative envelopes calling for observational constraints. We aim at verifying tidally locked rotation for the binary PG~0101$+$039, a subdwarf B star + white dwarf binary from its tiny (0.025\\%) light variations measured with the MOST satellite (Randall et al. \\cite{randall}).\\\\ {\\it Methods.} Binary parameters were derived from the mass function, apparent rotation and surface gravity of PG~0101$+$039 assuming a canonical mass of 0.47\\,M$_{\\rm \\odot}$ and tidally locked rotation. The light curve was then synthesised and was found to match the observed amplitude well.\\\\ {\\it Results.} We verified that the light variations are due to ellipsoidal deformation and that tidal synchronization is established for PG~0101$+$039. We conclude that this assumption should hold for all sdB binaries with orbital periods of less than half a day. Hence the masses can be derived from systems too faint to measure tiny light variations. ",
        "introduction": "}  The masses of compact objects like white dwarfs, neutron stars and black holes are fundamental to astrophysics, but very difficult to measure. Close binary systems consisting of a visible primary and an invisible compact object are very useful to this end as the companion mass can be constrained from the radial velocity  and the light curve of the primary if the primary mass and the orbital inclination are known. The latter can be measured in systems that are eclipsing or show tidally locked rotation. If the mass of the primary can be estimated, the companion mass can then be derived. The assumption of tidally locked rotation has often been used to determine the masses of neutron stars and black holes in known X-ray binaries (see Charles \\& Coe \\cite{charles} for a review). The same technique has recently been applied to KPD~1930$+$2752, a short period binary consisting of a subluminous B stars and a white dwarf (Geier et al. \\cite{geier1})\\footnote{The companion is so massive that the system mass may exceed the Chandrasekhar mass, making the system a viable Supernova Ia progenitor candidate in the double degenerate scenario.}.  sdB stars are core helium burning stars with very thin hydrogen envelopes and masses around $0.5M_{\\rm \\odot}$ (Heber et al. \\cite{heber1}). A large fraction of the sdB stars are members of short period binaries (Maxted et. al \\cite{maxted2}; Napiwotzki et al. \\cite{napiwotzki6}). For these systems common envelope ejection is the most probable formation channel (Han et al. \\cite{han2}). The companions of sdBs in these systems are predominantly white dwarfs implying that the system has undergone two phases of common envelope ejection.      In order to establish bound rotation for an sdB star the synchronization time scale has to be smaller than its evolutionary life time ($t_{\\rm EHB}\\approx10^{8}\\,{\\rm yrs}$). Two theoretical concepts to compute synchronization times have been developed by Zahn (\\cite{zahn}) and Tassoul \\& Tassoul (\\cite{tassoul}), respectively. While the mechanism proposed by Zahn is not efficient enough to fully fit the observed levels of synchronization, the more efficient one of Tassoul \\& Tassoul is matter of a controversy, given its free parameter dependence (see Claret et al. \\cite{claret1}, \\cite{claret2} and references therein). Especially in the case of hot stars with radiative envelopes, where tidal forces are less effective in synchronizing the stars, the predictions of the two theoretical models at hand can differ by orders of magnitude. All prior studies were undertaken to match observations of hot main sequence stars. Hot subdwarf stars have similar temperatures as B-type main sequence stars, but are much smaller and the internal structure of these helium core burning objects is different. In addition, the fraction of sdBs residing in close binary systems is among the highest known of all types of stars. Observations of hot subdwarfs could provide a new benchmark to study the yet unresolved problem of tidal dissipation in radiative stellar envelopes.  Independent observational constraints are needed to prove or disprove synchronized rotation in hot subdwarf stars. Ellipsoidal variations can be used to verify synchronization of the stellar surface because the light variations would then have to occur at exactly half the orbital period. Two sdB + white dwarf binaries are known to show ellipsoidal variations at half of the orbital period (KPD\\,0422+5421 Orosz \\& Wade \\cite{orosz}; KPD\\,1930+2752, Geier et al. \\cite{geier1}). However, both systems have short orbital periods of about $0.1\\,{\\rm d}$ and theory predicts synchronization times much smaller the evolutionary time scale. As the synchronization time strongly increases with increasing period, we expect an upper limit to the period to exist at which the assumption of tidally locked rotation breaks down. To this end it would be of utmost importance to find ellipsoidal variations in an sdB binary of longer period and to provide a stringent test for the theory of synchronization.  Recently a suitable object has been found. PG~0101$+$039, an sdB+WD binary (P=0.567 d, Maxted et al. \\cite{maxted2}) was discovered to show very weak luminosity variation at half the orbital period in a light curve in a 16.9 day long, almost uninterrupted light curve obtained with the MOST satellite (Randall et al. \\cite{randall}).    In order to verify that we indeed see ellipsoidal variations, we have to show that the observed light curve can be consistently modelled. Beforehand, we have to derive the complete set of system parameters. As the spectrum is single lined, the analysis of the radial velocity curve yields the mass function only. Complementing it with an estimate of the sdB mass and with measurements of the sdB's projected rotational velocity as well as its gravity allows to solve for all binary parameters and compute the light curve. ",
        "conclusions": "Tidally locked rotation in close binary systems has been assumed to measure masses of invisible compact companions, in particular in X-ray binaries. The synchronization time is very difficult to calculate for stars with radiative envelopes and plagued with large uncertainties. Therefore, observational constraints are of utmost importance. Ellipsoidal variations can be used to verify the assumption at least for the surface layers. We applied this technique to the sdB/WD binary PG~0101$+$039, for which a very weak luminosity variation at half the orbital period has been discovered in a light curve in a 16.9 day long, almost uninterrupted light curve obtained with the MOST satellite.  From spectroscopy we measured the mass function, apparent rotation and surface gravity of PG~0101$+$039. Stellar evolution models suggest that the sdB mass is close to 0.47\\,M$_{\\rm \\odot}$. Assuming tidally locked rotation, this information is sufficient to solve for all parameters of the binary system. The companion mass is found to be $M_{\\rm WD}=0.72 \\pm 0.10\\,M_{\\rm \\odot}$, typical for a white dwarf. The light curve was then synthesised and was found to match the observed amplitude well. However, a problem with the phasing of the light curve to the radial velocity curve became apparent. Due to a six year difference between the MOST photometry and published radial velocities, the phase errors were far too large for any conclusion to be drawn. Therefore we added 16 radial velocities from three observatories. The statistical error of the period decreased sufficiently to enable proper phasing of the photometry. The synthesised light curve was found to be offset by 0.1 cycles from the observed one indicating that our systematic error estimate may be overly optimistic. Alternative explanations like supersynchronous rotation of the sdB that may cause the observed phase shift seem to be unlikely because a deviation of $10\\%$ from equilibrium would require fast rotation of the sdB. In this case the inclination would be very low and the companion mass would rise dramatically.  A simultaneous measurement of the radial velocity curve and a high precision light curve would be necessary to solve this problem since the theoretical understanding of angular momentum transfer in hot stars with radiative envelopes is still very limited. In conclusion, we found strong indication that the surface rotation of the sdB star PG~0101$+$039 is tidally locked to its orbit.   The synchronization times for any given type of primary depend strongly on the orbital period (Zahn \\cite{zahn}, Tassoul \\& Tassoul \\cite{tassoul}). Hence, other sdB stars in close binaries should also be synchronized if their orbital period is less than that of PG~0101$+$039 (P\\,=\\,$0.567\\,{\\rm d}$). Hence we conclude that tidally locked surface rotation is established in sdB binaries with orbital periods of less than half a day. Hence the assumption of tidally locked rotation can be safely applied to such systems, even if they are too faint to measure such extremely small light variations as observed here."
    },
    "0710/0710.1718_arXiv.txt": {
        "abstract": "We derive accretion rate functions (ARFs) and kinetic luminosity functions (KLF) for jet-launching supermassive black holes. The accretion rate as well as the kinetic power of an active galaxy is estimated from the radio emission of the jet. For compact low-power jets, we use the core radio emission while the jet power of high-power radio-loud quasars is estimated using the extended low-frequency emission to avoid beaming effects.  We find that at low luminosities the ARF derived from the radio emission is in agreement with the measured bolometric luminosity function (BLF) of AGN, i.e., all low-luminosity AGN launch strong jets.  We present a simple model, inspired by the analogy between X-ray binaries and AGN, that can reproduce both the measured ARF of jet-emitting sources as well as the BLF. The model suggests that the break in power law slope of the BLF is due to the inefficient accretion of strongly sub-Eddington sources. As our accretion measure is based on the jet power it also allows us to calculate the KLF and therefore the total kinetic power injected by jets into the ambient medium. We compare this with the kinetic power output from SNRs and XRBs, and determine its cosmological evolution. ",
        "introduction": "\\label{s:intro} There is increasing support for the idea that feedback from active galactic nuclei (AGN) plays an important role for galaxy formation \\citep[e.g.,][]{CowieBinney1977,BinneyTabor1995,SilkRees1998,ChurazovKaiser2001,diMatteoSpringelHernquist2005}. This feedback is also invoked to explain the M-$\\sigma$ relation of supermassive black holes \\citep[e.g.,][]{HaehneltNatarajanRees1998,King2003,RobertsonHernquistCox2006,FabianCelottiErlund2006}. It is not yet clear whether kinetic or radiative feedback is dominant and how efficient each is at a given accretion rate. In this paper we will exploit the analogy between X-ray binaries (XRBs) and AGN to obtain information about the efficiency of the different feedback processes and calculate the total power available for feedback at a given redshift.  The central engines of XRBs and AGN seem to be similar \\citep[e.g.,][]{MirabelRodriguez1998,Meier2001} and recently relations have been found which scale spectral and variability properties from one class to the other \\citep{MerloniHeinzdiMatteo2003,FalckeKoerdingMarkoff2004,KoerdingJesterFender2006,McHardyKoerding2006}. In the nearby universe ($z\\la 0.2$), there are very few high-luminosity quasars like those that exist at high redshifts. These bright high redshift quasars are likely to have a strong effect on the evolution of the galaxy luminosity function. However, there are several XRBs that go through transient phases of very high accretion rates. These ``very-high state'' objects may be better templates for the bright quasars at a redshift of two than any nearby AGN. Thus, we will use our knowledge of XRBs and their states to obtain information about the kinetic and radiative properties of AGN.  For XRBs one can observe a full outburst cycle, in which the accretion rate increases from very low ($10^{-7} \\leq \\frac{\\dot{M}}{\\dot{M}_{Edd}} \\leq 10^{-5}$) to near the Eddington limit and then decays again. At low accretion rates, the source is generally found in the \\emph{hard state}, which is characterized by a hard power law in the X-ray spectrum. The hard X-ray emission is usually accompanied by radio emission associated with a compact jet \\citep{Fender2001} which sometimes can be directly imaged \\citep[e.g.,]{StirlingSpencer2001}.  The accretion flow in the hard state is likely to be radiatively inefficient \\citep*[e.g.,][]{EsinMcClintockNarayan1997,KoerdingFenderMigliari2006}. The source can stay in the hard state to fairly high luminosities ($\\sim 30 \\%$ Eddington) until it changes its state. During the state transition, it first enters the \\emph{hard intermediate state} (IMS), which is characterized by a hard spectral component and band-limited noise in the power spectrum. It is usually accompanied by an increasingly unstable jet \\citep{FenderBelloniGallo2004}.  After the hard IMS it enters the soft IMS, which is dominated by a soft spectral component and power-law noise in the power spectrum. Near this transition one often observes a bright radio flare; however, after the flare the jet seems to be quenched. After the soft IMS, the source is often found in the \\emph{soft state} where the X-ray spectrum is dominated by a soft multi-colour black body component. Here the jet is typically quenched by a factor $\\ga 50$ in GHz radio luminosity \\citep{FenderCorbelTzioumis1999,CorbelFenderTzioumis2000}.  In the soft state it is generally assumed that the source is efficiently accreting, i.e., it has a standard geometrically thin, optically thick accretion disk \\citep{ShakuraSunyaev1973}. On the further way back down to low accretion rates, the source stays in the soft state until it reaches a critical accretion rate around $\\sim 2\\%$ of the Eddington rate \\citep{Maccarone2003}, where the reverse state transition begins. This proceeds again via the soft IMS and the hard IMS to the hard state. The transition from the hard to the soft state typically occurs at a higher luminosity than the transition from the soft to the hard state. Due to this hysteresis effect there is no one-to-one correspondence of accretion rates to accretion states.  For a  detailed discussion of states and their exact definitions see \\citet{Nowak1995,BelloniHomanCasella2005,HomanBelloni2005}, but see also \\citet{McClintockRemillard2003} for slightly different definitions (mainly concerning the intermediate states).  Both in X-ray binaries and AGN, the unbeamed radio luminosity can be used to estimate the accretion rate. \\citet{KoerdingFenderMigliari2006} do this for the \\emph{core} radio emission, while \\citet{WillottRawlingsBlundell1999} present a correlation between \\emph{extended} radio luminosity and tracers of the accretion rate of luminous double-lobed radio sources. Here, we use these methods to translate radio luminosity functions into accretion rate functions (ARF) of jet-emitting sources.  In order to turn a \\emph{radiative} bolometric luminosity function (BLF) into an ARF as well, it is necessary to know the radiative efficiency of the accretion flow. There have been several suggestions that inefficient accretion is visible both in XRBs \\citep[e.g.,][]{Ichimaru1977,EsinMcClintockNarayan1997} and AGN (e.g., \\citealt{ReesPhinneyBegelman1982,ChiabergeGilliMacchetto2006}), and a number of authors have attempted to link both source classes and establish a detailed correspondence (see, e.g., \\citealt*{Meier2001,LivioPringleKing2003,MaccaroneGalloFender2003,MerloniHeinzdiMatteo2003,FalckeKoerdingMarkoff2004,Jester2005}; \\citealt*{KoerdingFalckeCorbel2005}; \\citealt{KoerdingFenderMigliari2006}). These sources are not only likely to be inefficiently accreting, but the total energy output of the sources may be dominated by the jet power \\citep[e.g.,][]{FenderGalloJonker2003,KoerdingFenderMigliari2006}.  To compare the ARF with the BLF we will explore the effect of inefficient accretion flows for low-luminosity objects that are included in the bolometric luminosity function of \\citep{HopkinsRichardsHernquist2007}.  Comparing our radio-derived ARF obtained in this way with this BLF, we can also determine a radio-loud fraction that has an intuitive physical meaning: the ratio of volume densities of radio-loud and radio-quiet sources at a given total accretion rate.  As the jet power is employed in our method for estimating the accretion rates from radio data, we can also directly estimate the jets' kinetic power from the radio luminosities \\citep[for an approach estimating jet powers from modeling blazar SEDs, see][]{MaraschiTavecchio2003}. This allows us to construct kinetic luminosity functions for jet emitting sources in the local universe as well as for higher redshifts. Such kinetic luminosity functions have already been constructed by \\citet{HeinzMerloniSchwab2007}. However, we use a slightly different methodology and are able to include high-luminosity sources (i.e., FR-II radio galaxies). By integrating the kinetic luminosity functions over all luminosities we then estimate the total power available for kinetic and radiative feedback at a given redshift.  We also consider constraints on the Eddington ratio distribution of accreting black holes which reproduce the observed luminosity functions of \\emph{all} classes of AGN using only the black-hole mass function and prescriptions for the radiative and kinetic outputs of accreting black holes as function of accretion rate \\citep[compare][]{VolonteriSalvaterraHaardt2006,MarulliBranchiniMoscardini2007}. ",
        "conclusions": "For both AGN and X-ray binaries, both the accretion rate and the jet power can be estimated from either the core radio luminosity \\citep{KoerdingFenderMigliari2006} or the extended low-frequency radio luminosity \\citep{WillottRawlingsBlundell1999}. Using these accretion rate estimates, we construct accretion rate functions (ARFs) of jet-emitting sources. The luminosity function (LF) of low-luminosity sources -- which likely have slow jets -- is roughly in agreement with the bolometric luminosity function (BLF) of AGN (Fig.~\\ref{FigLumZ0}) as determined by \\citet{HopkinsRichardsHernquist2007}. The ARF of RLQ-like jet sources is roughly 1\\,dex below the BLF (Fig.~\\ref{Figfrac}).  This provides a measurement of the radio-loud fraction with a direct physical meaning: the ratio of volume densities of radio-loud and radio-quiet sources at a given accretion rate. However, the exact value is unfortunately strongly dependent on the normalizations used to derive accretion rates.  We have developed a simple model based on the universality of accretion physics in XRBs and AGN that reproduces both our ARF as well as the BLF. At low luminosities, all sources are radiatively inefficient and the luminosity depends quadratically on the accretion rate (eq.~\\ref{eqbol}). Sources are radiatively efficient only at high luminosities ($\\ga 3\\%$ Eddington). In our model, we assumed that the distribution of accretion rates does not depend on the black-hole mass and can be written as a simple broken power law for the Eddington-scaled accretion rate (eq.~\\ref{eqmdotdist}). With this model we can reproduce the measured ARF as well as the local BLF (Fig.~\\ref{fiARFQLF}). Thus, this model can solve the problem that there seem to be too few LLAGN compared to the number of weak quasars. In this simple model the break in the luminosity function is due to the inefficient accretion at low accretion rates \\citep[compare][]{Jester2005}.  Using the corresponding jet power measures, we have calculated kinetic luminosity functions and their cosmological evolution. Our findings support the idea that the majority of the kinetic power is created by lower-luminosity AGN with their likely slow jets. The highly luminous radio-loud quasars do not contribute significantly due to their low number density. The total power injected into the ambient medium by jets from AGN is comparable or even exceeds the effect of supernovae (the situation may and will be different for individual objects). Only at high redshifts, ``quasar-mode'' feedback mechanisms provide a comparable amount of energy for heating the interstellar/intergalactic gas. Finally, we discuss the effects of the different feedback mechanisms and find that a roughly constant fraction (5-10$\\%$) of the accretion power is available for feedback, independent of the nature of the AGN."
    },
    "0710/0710.4417_arXiv.txt": {
        "abstract": "{We present the results of a deep (1.1 Ms) observation of the Coma cluster of galaxies in the 18-30 keV band with the ISGRI imager on board the \\emph{INTEGRAL} satellite. We show that the source extension  in the North-East to South-West (SW) direction ($\\sim 17'$)  significantly exceeds the size of the point spread function of ISGRI, and that the centroid of the image of the source in the 18-30~keV band is displaced in the SW direction compared to the centroid in the 1-10~keV band. To test the nature of the SW extension we fit the data assuming different models of source morphology. The best fit is achieved with a diffuse source of elliptical shape, although an acceptable fit can be achieved assuming an additional point source SW of the cluster core. In the case of an elliptical source, the direction  of extension of the source coincides with the direction toward the subcluster falling onto the Coma cluster. If the SW excess is due to the presence of a point source with a hard spectrum, we show that there is no obvious X-ray counterpart for this additional source, and that the closest X-ray source is the quasar EXO 1256+281, which is located $6.1'$ from the centroid of the excess. Finally, we show that the hard X-ray emission coincides with the 1.4 GHz radio emission, which suggests that the hard X-ray emission comes from the same population of electrons that is responsible for radio haloes through synchrotron emission.}  \\begin{document} ",
        "introduction": "In the hierarchical scenario of structure formation, clusters of galaxies are the latest and biggest structures to form. Hence, we expect some of them to be still forming, and experiencing major merging events with smaller clusters. This is the case of the Coma cluster, that is currently merging with the NGC 4839 group.  In such events, the merging of the ICM of the two clusters creates shock fronts, in which theory predicts that an important population of particles would be accelerated to high energies \\citep{sarazin}. This phenomenon should then produce a reheating of the gas, and create a higher temperature plasma that would radiate more strongly in hard X-rays. Alternatively, interaction of the population of mildly relativistic electrons that produce the halos of galaxy clusters via synchrotron radiation \\citep{feretti} with the Cosmic Microwave Background would then produce hard X-ray emission through inverse Compton processes, and thus add a power-law tail to the spectrum in the hard X-ray domain. Detection of this hard X-ray excess would help to learn more about the cosmic ray population detected by radio observations. Furthermore, characterization of the morphology of the hard X-ray emission would bring a possible identification of acceleration sites.  Recent reports of detection of a hard X-ray excess by \\emph{Beppo-SAX} \\citep{fusco} and \\emph{RXTE} \\citep{rephaeli} in the Coma cluster seem to confirm the existence of a high energy tail of the spectrum of merging clusters, and thus prove the existence of particle acceleration sites in these clusters. However, these detections are rather weak and controversial \\citep{rossetti}, and since the hard X-ray instruments on both \\emph{Beppo-SAX} and \\emph{RXTE} are non-imaging, contamination by very hard point sources inside the cluster could not be excluded (e.g. by the central galaxy NGC 4874, NGC 4889 or the QSO EXO 1256+281). Besides, it was not possible to have any information on the morphology of the hard X-ray emission. ",
        "conclusions": "Thanks to the imaging capabilities of ISGRI, we were able for the first time to resolve spatially the Coma cluster in the hard X-ray domain. We showed that the hard X-ray emission is displaced compared to the purely thermal emission in soft X-rays. Investigating this displacement, we found that the contribution of point sources to the observed spectrum is likely negligible. Our analysis shows that a hotter plasma can explain this excess. However, the correlation between the radio and hard X-ray morphology strongly suggests that the SW excess is due to IC scattering of mildly relativistic electrons.  Together with the spectral results obtained by \\emph{Beppo-SAX} \\citep{fusco2}, it is now becoming clear that the presence of an additional X-ray spectral component is required by the data. Moreover, the strong correlation between radio and hard X-ray morphology clarifies the nature of this excess, and seems to confirm the existence of a particle acceleration site in the Coma cluster."
    },
    "0710/0710.1697_arXiv.txt": {
        "abstract": "{A large fraction of otherwise similar asymptotic giant branch stars (AGB) do not show OH maser emission. As shown recently, a restricted lifetime may give a natural explanation as to why only part of any sample emits maser emission at a given epoch.} % {We wish to probe the lifetime of 1612 MHz OH masers in circumstellar shells of AGB stars.} {We reobserved a sample of OH/IR stars discovered more than 28 years ago to determine the number of stars that may have since lost their masers.} {We redetected all 114 OH masers. The minimum lifetime inferred is 2800 years (1$\\sigma$). This maser lifetime applies to AGB stars with strong mass loss leading to very red infrared colors. The velocities and mean flux density levels have not changed since their discovery. As the minimum lifetime is of the same order as the wind crossing time, strong variations in the mass-loss process affecting the excitation conditions on timescales of $\\approx$3000 years or less are unlikely.} {} ",
        "introduction": "Molecular OH maser emission at 1612 MHz is frequently observed from stars approaching the tip of their evolutionary track on the asymptotic giant branch (AGB). These stars are in general long-period variables with pulsation periods $>$1~year.  The pulsation is accompanied by heavy mass loss, which forms a circumstellar envelope (CSE) of gas and dust, and if the mass-loss rate surpasses \\mdot\\ $\\approx$10$^{-6}$ \\Myr\\ the dust envelope becomes opaque for visible light. Classically these `invisible' stars were called OH/IR stars, but nowadays this term is often used in general reference to OH emitting AGB stars selected in the infrared (see review by Habing \\cite{Habing96}).  Interestingly, a large fraction ($\\approx$40\\%) of any sample of IRAS sources, selected to match the infrared colors of known OH/IR stars, does not exhibit a detectable OH 1612 MHz maser (Lewis \\cite{Lewis92}).  These infrared sources were baptized by Lewis as `OH/IR star color mimics'.  Some of them do exhibit mainline OH (at 1665 or 1667 MHz) and/or 22 GHz \\water\\ masers (Lewis \\& Engels \\cite{Lewis95}), corroborating their cousinhood with OH/IR stars. OH/IR stars are therefore only part of the population of evolved and obscured AGB stars with oxygen-rich chemistry.  The reasons for the absence of 1612 MHz OH masers in a large fraction of oxygen-rich AGB stars are not known.  The maser photons are emitted by a transition between hyperfine levels of the groundstate of OH, which are inverted with the help of pump photons at $\\lambda$=35 and 53\\micron, emitted from dust in the CSEs. A requirement for excitation is sufficient OH column density, which might be low in mimics due to the destruction of OH by the interstellar UV field. This effect cannot account for mimics in general, as mimics are also present at higher galactic latitudes, where the influence of the UV field is low. In addition, the presence of a hot white dwarf companion leading to photodissociation of OH is not generally able to suppress the OH maser (Howe \\& Rawlings \\cite{Howe94}). A further requirement is velocity coherence over large distances to allow amplification. Turbulence may disrupt this coherence in the case of mimics.  An alternative explanation is the assumption that the OH maser is only present temporarily on the AGB and therefore these stars may change between an OH/IR status and that of a mimic. This explanation has been triggered by recent observations showing the fading of the OH maser in \\object{IRAS~18455+0448} over a time range of a decade (Lewis et al. \\cite{Lewis01}) and the absence of masers in four stars during a revisit of 328 OH/IR stars 12 years after their first detection (Lewis \\cite{Lewis02}).  The high rate of `dead' OH/IR stars among a sub-sample of 112 OH/IR stars with relatively blue colors led Lewis to conclude that the mean 1612 MHz emission life is in the range 100--400 years.  To test this lifetime we reobserved another sample of N$>$ 100 OH/IR stars, which was drawn from the first surveys for OH maser emission prior to 1980. With a difference of almost 30 years between the two observations several masers were expected to have disappeared. ",
        "conclusions": "We find that the lifetime of OH masers in classical OH/IR stars is $>2800$ years, in contrast to the lifetime implied by the observed disappearance of several masers in AGB stars with bluer envelopes. The lower limit of the lifetime is of the order of the wind crossing time in the CSEs, implying that, in general, no drastic variations in the structure of the CSEs happen on such timescales, or shorter ones. If non-variable OH/IR stars are in transition to the post-AGB stage, a sudden decline of the mass-loss rates at the end of AGB evolution is ruled out.  The previous observed disappearance of OH masers in stars with bluer envelopes is probably not associated with major changes in the mass-loss process. The susceptibility of masers to smaller variances in their environment may lead to OH masers as transient phenomena on the AGB.  This gives a natural explanation for the detection of OH masers in only a part of any AGB star sample with otherwise similar properties."
    },
    "0710/0710.1142_arXiv.txt": {
        "abstract": "{} % {We present new millimetre 43 GHz observations of a sample of radio-bright Planetary Nebulae. Such observations were carried out to have a good determination of the high-frequency radio spectra of the sample in order to evaluate, together with far-IR measurements (IRAS), the fluxes emitted by the selected source in the millimetre and  sub-millimetre band. This spectral range,  even very important to constraint the physics of circumstellar environment, is still far to be completely exploited.} {To estimate the millimetre and sub-millimetre fluxes, we extrapolated and summed together the ionized gas  (free-free radio emission) and dust (thermal emission) contributions at this frequency range. By comparison of the derived flux densities to the foreseen sensitivity we investigate the possible detection of such source for all the channels of the forthcoming ESA's PLANCK mission.} {We conclude that almost 80\\% of our sample will be detected by PLANCK, with the higher detection rate in the higher frequency channels, where there is a good combination of brighter intrinsic flux from the sources and reduced extended Galactic foregrounds contamination despite a worst instrumental sensitivity. From the new 43 GHz, combined with single-dish 5 GHz observations from the literature, we derive radio spectral indexes, which are consistent with optically thin free-free nebula. This result indicates that the high frequency radio spectrum of our sample sources is dominated by thermal free-free and other emission, if present, are negligible. } {} % ",
        "introduction": "The PLANCK ESA mission will provide us with  nearly full sky maps over a wide range of frequencies, from 30 to 900 GHz. Therefore, even if designed  for cosmological studies, the mission  will have profund impact on fundamental Physics and Galactic and extragalactic astrophysics. Planck will be sensitive to the millimetre emission from dusty envelopes of stars and we expect the dusty circumstellar envelopes, that characterize the latest stages of stellar evolution, to be source of relevant foreground contamination. When a low or intermediate mass star is approaching the end of its evolution, it goes through a period of heavy mass-loss known as Asymptotic Giant Branch (AGB) phase. The ejected envelopes are partially condensed in dust grains and completely obscure the central star. Immediately after the AGB evolutionary phase, the mass-loss stops and the stars may become optically visible as the dusty shells disperse (Proto Planetary Nebula, PPN phase). Eventually, once it reaches a temperature of 2 -- 3 $\\times 10 ^4$ K, the central star starts to ionize the AGB shell and a Planetary Nebula will form. Dusty envelopes re-radiate the absorbed stellar light showing a clear signature in the far-infrared spectrum, i.e. an IR excess with peculiar IRAS colours. In addition to that, PNe show  a radio continuum due to free-free emission from the fraction of the CSE ionized by the central star.         \\begin{table*} \\caption{The selected sample} \\label{tab_posizioni} \\begin{center}   \\vspace{0.3cm}  \\begin{tabular}{cccc|cccc} \\hline \\hline IAU Name    & Other Name   & R.A.\\ (J2000) &Dec. \\ (J2000)    &IAU Name & Other Name  & R.A.\\ (J2000)  &   Dec. \\ (J2000)  \\\\ PN G&& [\\ h\\ m\\ s] &[$^{\\circ}\\ ^{\\prime}\\ ^{\\prime\\prime}$]  &PN G&&[\\ h\\ m\\ s] &[$^{\\circ}\\ ^{\\prime}\\ ^{\\prime\\prime}$] \\\\ \\hline $000.3+12.2$ &IC 4634         & 17 01 33.6    & $-21$ 49 33.1 &$093.4+05.4$ &NGC 7008         & 21 00 32.7      & $+54$ 32 39.4   \\\\ $002.4+05.8$ &NGC 6369        & 17 29 20.5    & $-23$ 45 35.0 &$093.5+01.4$ &PN M 1-78        & 21 20 44.8      & $+51$ 53 27.5   \\\\ $003.1+02.9$ &PN Hb 4         & 17 41 52.8    & $-24$ 42 09.3 &$096.4+29.9$ &NGC 6543         & 17 58 33.4      & $+66$ 37 58.8   \\\\ $006.7-02.2$ &PN M 1-41       & 18 09 30.6    & $-24$ 12 28.7 &$097.5+03.1$ &PN A66 77        & 21 32 10.2      & $+55$ 52 43.2   \\\\ $007.2+01.8$ &PN HB 6         & 17 55 07.0    & $-21$ 44 41.0 &$106.5-17.6$ &NGC 7662         & 23 25 53.9      & $+42$ 32 04.7   \\\\ $008.0+03.9$ &NGC 6445        & 17 49 15.0    & $-20$ 00 33.7 &$107.8+02.3$ &NGC 7354         & 22 40 19.9      & $+61$ 17 08.0   \\\\ $008.3-01.1$ &PN M 1-40       & 18 08 26.0    & $-22$ 16 53.4 &$120.0+09.8$ &NGC 40           & 00 13 01.0      & $+72$ 31 19.6   \\\\ $009.4-05.0$ &NGC 6629        & 18 25 42.5    & $-23$ 12 11.3 &$130.9-10.5$ &NGC 650-51       & 01 42 19.7      & $+51$ 34 31.7   \\\\ $009.6+14.8$ &NGC 6309        & 17 14 04.3    & $-12$ 54 37.2 &$138.8+02.8$ &IC 289           & 03 10 19.3      & $+61$ 19 00.4   \\\\ $010.1+00.7$ &NGC 6537        & 18 05 13.1    & $-19$ 50 34.4 &$144.5+06.5$ &NGC 1501         & 04 06 59.3      & $+60$ 55 14.7   \\\\ $010.8-01.8$ &NGC 6578        & 18 16 16.5    & $-20$ 27 03.4 &$165.5-15.2$ &NGC 1514         & 04 09 16.9      & $+30$ 46 32.0   \\\\ $011.7-00.6$ &NGC 6567        & 18 13 45.2    & $-19$ 04 35.6 &$166.1+10.4$ &IC 2149          & 05 56 23.9      & $+46$ 06 17.4   \\\\ $020.9-01.1$ &PN M 1-51       & 18 33 29.0    & $-11$ 07 26.3 &$173.7+02.7$ &PP 40            & 05 40 52.7      & $+35$ 42 18.6   \\\\ $025.8-17.9$ &NGC 6818        & 19 43 57.8    & $-14$ 09 11.8 &$194.2+02.5$ &J 900            & 06 25 57.3      & $+17$ 47 27.6   \\\\ $027.7+00.7$ &PN M 2-45       & 18 39 21.9    & $-04$ 19 52.6 &$197.8+17.3$ &NGC 2392         & 07 29 10.8      & $+20$ 54 41.6   \\\\ $033.8-02.6$ &NGC 6741        & 19 02 37.0    & $-00$ 26 57.2 &$206.4-40.5$ &NGC 1535         & 04 14 15.8      & $-12$ 44 22.3   \\\\ $034.6+11.8$ &NGC 6572        & 18 12 06.3    & $+06$ 51 12.4 &$215.2-24.2$ &IC 418           & 05 27 28.2      & $-12$ 41 50.2   \\\\ $035.1-00.7$ &PN Ap 2-1       & 18 58 10.5    & $+01$ 36 57.5 &$221.3-12.3$ &IC 2165          & 06 21 42.8      & $-12$ 59 13.9   \\\\ $037.7-34.5$ &NGC 7009        & 21 04 10.8    & $-11$ 21 48.5 &$234.8+02.4$ &NGC 2440         & 07 41 55.4      & $-18$ 12 30.5   \\\\ $039.8+02.1$ &PN K 3-17       & 18 56 18.2    & $+07$ 07 26.2 &$254.6+00.2$ &NGC 2579         & 08 20 54.1      & $-36$ 13 00.0   \\\\ $041.8-02.9$ &NGC 6781        & 19 18 28.1    & $+06$ 32 20.0 &$258.1-00.3$ &Hen 2-9          & 08 28 28.0      & $-39$ 23 39.4   \\\\ $043.1+37.7$ &NGC 6210        & 16 44 29.5    & $+23$ 47 59.9 &$259.1+00.9$ &Hen 2-11         & 08 37 08.1      & $-39$ 25 04.9   \\\\ $045.7-04.5$ &NGC 6804        & 19 31 35.1    & $+09$ 13 30.2 &$261.0+32.0$ &NGC 3242         & 10 24 46.1      & $-18$ 38 32.3   \\\\ $050.1+03.3$ &PN M 1-67       & 19 11 31.1    & $+16$ 51 32.0 &$294.1+43.6$ &NGC 4361         & 12 24 30.8      & $-18$ 47 04.0   \\\\ $054.1-12.1$ &NGC 6891        & 20 15 08.9    & $+12$ 42 15.4 &$342.1+10.8$ &NGC 6072         & 16 12 58.4      & $-36$ 13 46.6   \\\\ $063.1+13.9$ &NGC 6720        & 18 53 35.1    & $+33$ 01 45.1 &$349.5+01.0$ &NGC 6302         & 17 13 44.5      & $-37$ 06 11.6   \\\\ $064.7+05.0$ &BD+30 3639      & 19 34 45.2    & $+30$ 30 59.2 &$352.6+00.1$ &PN H 1-12        & 17 26 24.3      & $-35$ 01 41.8   \\\\ $082.1+07.0$ &NGC 6884        & 20 10 23.7    & $+46$ 27 40.0 &$352.8-00.2$ &PN H 1-13        & 17 28 27.7      & $-35$ 07 30.4   \\\\ $083.5+12.7$ &NGC 6826        & 19 44 48.2    & $+50$ 31 31.3 &$358.5+02.6$ &PN HDW 8         & 17 31 47.3      & $-28$ 42 03.5   \\\\ $086.5-08.8$ &PN Hu 1-2       & 21 33 08.2    & $+39$ 38 08.3 &$358.5+05.4$ &PN M 3-39        & 17 21 11.5      & $-27$ 11 37.0   \\\\ $089.0+00.3$ &NGC 7026        & 21 06 18.7    & $+47$ 51 07.5 &$359.3-00.9$ &Pn HB 5          & 17 47 56.3      & $-29$ 59 40.6   \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}   From a feasibility study, Umana et al. (\\cite{Umana06}) concluded that a sizable ($\\approx 300$) sample of AGB and post-AGB stars would be detected during the mission and derived estimates for the expected flux densities at various Planck channels. However, the simulations carried out by Umana et al. (\\cite{Umana06}), on a sub-sample of PNe rely only on NVSS fluxes, obtained at 1.4 GHz, extrapolated to the Planck frequencies. This leads to underestimate the contribution due to free-free as, at this frequency, PNe are often optically thick (Siodmiank \\& Tylenda \\cite{st01},  Luo et al. \\cite{luo_etal05}).  PNe  are among the brightest Galactic radio sources. Some of them could also reach a flux density of Jy level. More than 800 have been  detected at least at one frequency  and for 200 morphological and spectral information (between 1.4 and 22 GHz) were obtained. Most of the higher frequency (22 GHz) data were obtained with interferometers (i.e. VLA) while very little single-dish, high frequency measurements are available.   In this paper we present new 43 GHz, single-dish observations by using the 32 m INAF-IRA Radiotelescope at Noto of a sample of PNe, which are potential foregrounds for PLANCK. The main goal of this project is to obtain reliable estimates of flux density expected at PLANCK channels, by building  and modeling their spectral energy distribution (SED). This in turn will contribute to the compilation of the PLANCK pre-launch catalogue. We stress here that 43 GHz is one of the observing channels of the forthcoming PLANCK mission. Therefore, at this frequency band, we would obtain a direct measurement of the expected flux and not an extrapolation.  As added value, the present work provides the first sizeable dataset of 43 GHz measurements of PNe, that constitute strong  constraints to the observed SEDs in the very important spectral region where free-free emission and thermal dust emission may overlap.  While interferometric high frequency observations provide us with detailed morphological information they quite often fail to entirely recover the extended emission. This eventually leads to underestimate the total radio flux density. This problem is  overcame by single-dish observations. Assesting the SEDs in the radio-millimetre spectral range is a crucial step for the study of the physics of dusty envelopes around PN. A correct evaluation of the free-free contribution, up to millimetre range, when combined with information provide by far-IR observations, would allow to determine the presence of an excess due to the presence of a cold dust component/s (Gee et al. \\cite{gee_etal84}; Hoare et al. \\cite{hoare_etal92}; Kemper et al. \\cite{kemper_etal02}) or of alternative emission mechanisms (Casassus et al. \\cite{cas_etal07}). Since PNe and their progenitors are believed to be among the major sources of recycled interstellar matter,  determining the properties of the dust ejected in the ISM is  very important to the study of the Galaxy evolution in general.     \\begin{table*} \\begin{center} \\caption{Results} \\label{tab_misure} \\begin{tabular}{crrclr|crrclr} \\hline \\hline IAU Name     & $S_\\mathrm{43 ~GHz}$ & $\\sigma_\\mathrm{43 ~GHz}$  &$\\theta_\\mathrm{1.4 ~GHz}$ & Ref.$^{\\ast}$ &$S_\\mathrm{c~43 ~GHz}$ &IAU Name  &$S_\\mathrm{43 ~GHz}$ & $\\sigma_\\mathrm{43 ~GHz}$  &$\\theta_\\mathrm{1.4 ~GHz}$ & Ref.$^{\\ast}$ &$S_\\mathrm{c~43 ~GHz}$ \\\\ PN G         & [mJy]   & [mJy]  & [arcsec]  &  & [mJy] &PN G          & [mJy]   & [mJy]  & [arcsec]  &  & [mJy]    \\\\ \\hline $000.3+12.2$ & $<$ 180    & 60  & 10.0  & C H     &      &$093.4+05.4$  & $<$ 240    & 80   & 49.2    & A E   &      \\\\ $002.4+05.8$ & 1330       & 110 & 20.7  & C H     &1530  &$093.5+01.4$  & 560        & 35   & 14.6    & A E   &600   \\\\ $003.1+02.9$ & 100        & 20  & 11.6  & C H     &105   &$096.4+29.9$  & 400        & 70   & 11.9    & A E   &420   \\\\ $006.7-02.2$ & 300        & 90  & 21.0  &         &345   &$097.5+03.1$  & 230        & 25   & 33.4    & A E   &320   \\\\ $007.2+01.8$ & 310        & 50  & 10.6  & C H     &320   &$106.5-17.6$  & 610        & 40   & 12.0    & A E   &640   \\\\ $008.0+03.9$ & 180        & 20  & 34.6  & C G     &250   &$107.8+02.3$  & 280        & 25   & 16.4    & A E   &305   \\\\ $008.3-01.1$ & 90         & 20  & 9.9   & C H     &95    &$120.0+09.8$  & 420        & 70   & 27.6    & A E   &530   \\\\ $009.4-05.0$ & 130        & 30  & 12.8  & C H     &135   &$130.9-10.5$  & 140        & 40   & 59.3    & A E   &310   \\\\ $009.6+14.8$ & 270        & 60  & 14.3  & C H     &290   &$138.8+02.8$  & $<$ 150    & 50   & 23.1    & A E   &      \\\\ $010.1+00.7$ & 400        & 60  & 11.8  & C F H   &420   &$144.5+06.5$  & 215        & 30   & 34.9    & A E   &305   \\\\ $010.8-01.8$ & 130        & 20  & 11.9  & C H     &135   &$165.5-15.2$  & 220        & 30   & 94.4    & A E   &890   \\\\ $011.7-00.6$ & $<$ 150    & 50  & 8.7   & C H     &      &$166.1+10.4$  & 100        & 30   & 9.1     & A E   &105   \\\\ $020.9-01.1$ & 420        & 90  & 13.5  & C       &445   &$173.7+02.7$  & 260        & 20   & 11.5    & E     &270   \\\\ $025.8-17.9$ & 270        & 30  & 14.7  & C H     &290   &$194.2+02.5$  & 90         & 20   & 0.0     & A E H &90    \\\\ $027.7+00.7$ & 110        & 20  & 0.0   & B G     &110   &$197.8+17.3$  & 240        & 50   & 22.2    & A E H &280   \\\\ $033.8-02.6$ & 290        & 50  & 14.7  & E H     &310   &$206.4-40.5$  & 160        & 20   & 20.6    & C H   &180   \\\\ $034.6+11.8$ & 1220       & 100 & 11.9  & A E H   &1280  &$215.2-24.2$  & 1100       & 100  & 0.0     & C H   &1100  \\\\ $035.1-00.7$ & 170        & 20  & 13.0  & A E     &180   &$221.3-12.3$  & 350        & 50   & 11.0    & C H   &365   \\\\ $037.7-34.5$ & 375        & 40  & 14.9  & C H     &400   &$234.8+02.4$  & 350        & 20   & 16.5    & C H   &380   \\\\ $039.8+02.1$ & 240        & 20  & 0.0   & A E     &240   &$254.6+00.2$  & 1770       & 290  & 41.2    & F     &2800  \\\\ $041.8-02.9$ & 230        & 20  & 78.4  & A E H   &715   &$258.1-00.3$  & 180        & 40   & 4.5     & D H   &181   \\\\ $043.1+37.7$ & 230        & 30  & 10.8  & A E H   &240   &$259.1+00.9$  & 270        & 60   & 42.2    & D     &435   \\\\ $045.7-04.5$ & 140        & 20  & 25.4  & A B E H &170   &$261.0+32.0$  & 290        & 30   & 19.5    & C H   &330   \\\\ $050.1+03.3$ & 140        & 30  & 41.4  & A E H   &220   &$294.1+43.6$  & 130        & 20   & 47.3    & C H   &230   \\\\ $054.1-12.1$ & 130        & 40  & 8.9   & A E H   &135   &$342.1+10.8$  & $<$ 210    & 70   & 32.9    & H     &      \\\\ $063.1+13.9$ & 255        & 75  & 48.2  & A E     &460   &$349.5+01.0$  & 2150       & 220  & 14.8    & D F H &2310  \\\\ $064.7+05.0$ & 565        & 20  & 10.5  & A E     &585   &$352.6+00.1$  & 815        & 125  & 10.3    & F     &845   \\\\ $082.1+07.0$ & 250        & 50  & 11.0  & A E     &260   &$352.8-00.2$  & 420        & 50   & 14.8    & F     &450   \\\\ $083.5+12.7$ & 320        & 40  & 15.9  & A E     &350   &$358.5+02.6$  & $<$ 120    & 40   & 21.1    & C G   &      \\\\ $086.5-08.8$ & $<$ 120    & 40  & 10.4  & A E     &      &$358.5+05.4$  & 380        & 40   & 0.0     & C G   &380   \\\\ $089.0+00.3$ & 220        & 30  & 11.9  & A E     &230   &$359.3-00.9$  & 290        & 70   & 0.0     & F H   &290   \\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{\\ast}$] References for 5 GHz measurements: A) Gregory et al. (\\cite{g_etal96}); B) Griffith et al. (\\cite{g_etal95}); C) Griffith et al. (\\cite{g_etal94}); D) Wright et al. (\\cite{w_etal94}); E) Becker et al. (\\cite{b_etal91}); F) Haynes et al. (\\cite{h_etal79}); G) Milne (\\cite{m79}); H) Milne (\\cite{ma75}) \\end{list} \\end{center} \\end{table*} ",
        "conclusions": "We have presented new 7 mm (43 GHz) observations of a sample of radio-bright PNe, carried out with the INAF-IRA Noto Radiotelescope. Such observations have been used to derive the high-frequency free-free contribution, due to the ionized fraction of the circumstellar envelope. So far, the majority of high frequency radio observations of PNe have been conducted with interferometers, that reveal details of source radio morphology but could, in principle, resolve out some of the extended emission. Our single dish measurements provide an extended database of millimetre observations of PNe, to be used when is necessary to know the overall emission, such as when building a SED. We used our measurements, that trace the free-free contribution, together with IRAS measurements, which trace the thermal dust emission, to build up the observed SED, from radio to far-IR and by extrapolation, to estimate the expected fluxes in the spectral range between radio and sub-mm, where observational data are missing.  When comparing the expected millimetre-sub-millimetre fluxes with the total foreseen sensitivity of the forthcoming ESA mission PLANCK we estimate that a consistent number of our targets will be easily detected by PLANCK, mostly at higher frequency channels.  Therefore, even if designed for cosmological study, PLANCK could also contribute to the PNe science. Results from such kinds of observations, once modelled with appropriate code (i.e. CLOUDY, DUSTY), would provide important constraints on the chemical composition and structure of the CSEs. It would  point out the presence of multi-shells, related to multiple mass-loss event suffered from the central object during its previous evolution (AGB), or a contribution of different emission processes, besides free-free and thermal from dust, as recently claimed by Cassassus et al.~(\\cite{cas_etal07}). Moreover, PLANCK results can be considered as pathfinder for other future instrumentations, with the same frequency coverage, such as ALMA, as they will help in planning more focused experiments aimed to investigate the morphological details of the sources."
    },
    "0710/0710.4883_arXiv.txt": {
        "abstract": "Knowledge of the stellar parameters for the parent stars of transiting exoplanets is pre-requisite for establishing the planet properties themselves, and often relies on stellar evolution models. GJ~436, which is orbited by a transiting Neptune-mass object, presents a difficult case because it is an M dwarf. Stellar models in this mass regime are not as reliable as for higher mass stars, and tend to underestimate the radius. Here we use constraints from published transit light curve solutions for GJ~436 along with other spectroscopic quantities to show how the models can still be used to infer the mass and radius accurately, and at the same time allow the radius discrepancy to be estimated. Similar systems should be found during the upcoming \\emph{Kepler\\/} mission, and could provide in this way valuable constraints to stellar evolution models in the lower main sequence. The stellar mass and radius of GJ~436 are $M_{\\star} = 0.452_{-0.012}^{+0.014}$~M$_{\\sun}$ and $R_{\\star} = 0.464_{-0.011}^{+0.009}$~R$_{\\sun}$, and the radius is 10\\% larger than predicted by the standard models, in agreement with previous results from well studied double-lined eclipsing binaries.  We obtain an improved planet mass and radius of $M_p = 23.17 \\pm 0.79$~M$_{\\earth}$ and $R_p = 4.22_{-0.10}^{+0.09}$~R$_{\\earth}$, a density of $\\rho_p = 1.69_{-0.12}^{+0.14}$ g~cm$^{-3}$, and an orbital semimajor axis of $a = 0.02872 \\pm 0.00027$ AU. ",
        "introduction": "\\label{sec:introduction}  The applications of stellar evolution theory to astrophysics are so widespread, and its validity so often taken for granted, that it is easy to forget that it took decades to develop, and significant effort to validate by comparison with careful measurements, a process that still continues. It is usually only when those theoretical predictions fail that the classical discipline of stellar evolution ``makes the headlines'', and even then it draws the attention of relatively few. One such instance has occurred for low mass main-sequence stars. Over the last 10 years or so it has become clear that our understanding of the structure and evolution of these objects is still incomplete.  Discrepancies between theory and observation in the radii of stars under 1~M$_{\\sun}$, first mentioned by \\cite{Hoxie:73}, \\cite{Lacy:77}, and others, are now well documented for several low-mass eclipsing binaries \\citep[see, e.g.,][]{Popper:97, Clausen:99, Torres:02, Ribas:03, Lopez-Morales:05}. Differences in the effective temperatures have been observed as well. The direction of these disagreements is such that model radii are underestimated by roughly 10\\%, while effective temperatures are overestimated.  In recent years stellar evolution has had important applications in the field of transiting extrasolar planets. This is because the planetary parameters of interest (mass $M_p$, radius $R_p$) depend rather directly on those of the star ($M_{\\star}$, $R_{\\star}$), and in most cases models provide the only means of determining the latter. The subject of this paper is GJ~436, a late-type star found by \\cite{Butler:04} to be orbited by a Neptune-mass planet with a period of 2.644 days. This object was later discovered by \\cite{Gillon:07a} to undergo transits, enabling its size to be determined ($\\sim$4~R$_{\\earth}$).  As the only M dwarf among the 22 currently known transiting planet host stars, GJ~436 (M2.5V) presents a special challenge for establishing the stellar parameters, because of the disagreements noted above. Little mention seems to have been made of this, and for the most part past studies have relied instead on empirical mass-luminosity ($M\\!-\\!L$) relations to set the mass of GJ~436. Radius estimates have often rested on the assumption of numerical equality between $M_{\\star}$ and $R_{\\star}$ for M stars. Despite being the closest transiting planet system (only 10 pc away), it is rather surprising that the mass of the star is only known to about 10\\% \\citep{Maness:07, Gillon:07b}. This is currently limiting the precision of the planetary mass, and some of that uncertainty translates also to the radius. These two properties are critical for studying the structure of the object.  Given the importance of GJ~436 as the parent star of the only Neptune-mass transiting exoplanet found so far, and hence the closest analog to our Earth with a mass and radius determination, one of the motivations of this paper is to improve the precision of the stellar and planetary parameters by making use of additional observational constraints not used before. Specifically, we incorporate the information on the stellar \\emph{density} directly available from the transit light curve \\citep{Sozzetti:07}, which provides a strong handle on the size of the star.  The nature of the discrepancies between evolutionary models and observations for low-mass stars has been examined recently from both the observational and theoretical points of view by \\cite{Lopez-Morales:07} and \\cite{Chabrier:07}. Further progress depends on gathering more evidence to supplement the few available highly accurate mass and radius measurements based on double-lined eclipsing binaries.  Thus, a second motivation for this work, despite the fact that GJ~436 is not a double-lined eclipsing binary, is to present a way of using all observational constraints simultaneously to show that the star presents the same radius anomaly found for the other systems, or more generally, to test the models.  Because similar constraints may become available in the future for other M dwarfs given the keen interest in finding smaller and smaller transiting planets, we anticipate that the indirect technique described here may yield valuable information on this problem and eventually help improve our understanding of low-mass stars. ",
        "conclusions": "With the host star properties known, the planet parameters we infer are given in the bottom section of Table~\\ref{tab:properties}.  The required orbital period and velocity semi-amplitude $K_{\\star}$ are adopted from \\cite{Maness:07}, along with the eccentricity from \\cite{Demory:07}, who obtain a similar value for $K_{\\star}$. The improved precision of these derived parameters is a reflection of the better stellar parameters. The slightly larger planet radius than in previous studies confirms with even greater statistical significance the conclusions of earlier authors regarding the presence of a hydrogen/helium envelope \\citep{Gillon:07a, Deming:07, Gillon:07b}, and agrees very well with the models by \\cite{Fortney:07} for a 10\\% fraction of those elements.  In this paper we have shown that GJ~436 provides a valuable test of stellar evolution theory near the bottom of the main sequence, made possible by the fact that it has a transiting planet.  Traditional studies in the area of low-mass stars have made the comparison with models by measuring the mass and radius directly for double-lined eclipsing systems containing M dwarfs. In a few other cases angular diameters have been measured interferometrically for single stars, and the mass has been inferred from empirical $M\\!-\\!L$ relations \\citep[e.g.,][]{Lane:01, Segransan:03}. More recently, a variety of constraints and assumptions have been used to infer the mass and radius of the late-type secondaries in F+M systems observed as part of transiting planet surveys \\citep[e.g.,][]{Bouchy:05, Pont:05, Beatty:07}. Though perhaps not as compelling as having actual model-independent mass and radius measurements, the approach in the present work is able to make use of available information for GJ~436 and compare the models directly with the observational constraints \\emph{without} requiring a direct measurement of the mass and radius. The discrepancy in $R_{\\star}$ is derived by parameterizing it in terms of a single adjustment factor to the model radii ($\\beta$), assuming the luminosity from theory is accurate, as other observations seem to indicate.  NASA's upcoming \\emph{Kepler\\/} mission, currently slated to launch in early 2009, will emphasize the search for transiting Earth-size planets.  These should be easier to detect around late-type stars. Therefore, we anticipate that many systems similar to GJ~436 could be found and become a significant source of information on radii for low-mass stars, since they will have all the observational constraints needed (including trigonometric parallaxes) to test models of stellar evolution in the way we have done here."
    },
    "0710/0710.0372_arXiv.txt": {
        "abstract": "We present a sample of \\total\\ galaxy nuclei from 12 nearby ($z<4500$\\kms) Hickson Compact Groups (HCGs) with a complete suite of 1--24\\micron\\ 2MASS+\\spitzer\\ nuclear photometry.  For all objects in the sample, blue emission from stellar photospheres dominates in the near-infrared through the 3.6\\micron\\ IRAC band.  Twenty-five of \\total\\ (54\\%) galaxy nuclei show red, mid-infrared continua characteristic of hot dust powered by ongoing star formation and/or accretion onto a central black hole. We introduce \\alphairac, the spectral index of a power-law fit to the 4.5--8.0\\micron\\ IRAC data, and demonstrate that it cleanly separates the mid-infrared active and non-active HCG nuclei.  This parameter is more powerful for identifying low to moderate-luminosity mid-infrared activity than other measures which include data at rest-frame $\\lambda<3.6$\\micron\\ that may be dominated by stellar photospheric emission.  While the HCG galaxies clearly have a bimodal distribution in this parameter space, a comparison sample from the \\spitzer\\ Nearby Galaxy Survey (SINGS) matched in $J$-band total galaxy luminosity is continuously distributed.  A second diagnostic, the fraction of 24\\micron\\ emission in excess of that expected from quiescent galaxies, \\fracd, reveals an additional three nuclei to be active at 24\\micron.  Comparing these two mid-infrared diagnostics of nuclear activity to optical spectroscopic identifications from the literature reveals some discrepancies, and we discuss the challenges of distinguishing the source of ionizing radiation in these and other lower luminosity systems.  We find a significant correlation between the fraction of mid-infrared active galaxies and the total \\HI\\ mass in a group, and investigate possible interpretations of these results in light of galaxy evolution in the highly interactive system of a compact group environment. ",
        "introduction": "\\label{sec:intro}  The first galaxies and their environments differed substantially from those locally, often involving multiple interactions as seen in the \\hst\\ Ultra Deep Field \\citep[e.g.,][]{malhotra05}. Compared to all other nearby environments, present-day compact galaxy groups most closely reproduce the interaction environment of the early Universe ($z\\sim4$) when galaxies assembled through hierarchical formation \\citep[e.g.,][]{baron}, and galaxy groups combined to form proto-clusters (in dense regions; e.g., \\altcite{rudick+06}) or massive ellipticals (in the field; \\altcite{white}).  Because of their high space densities (with comparable surface densities to the centers of rich galaxy clusters; e.g., \\altcite{rubin+91}) and low velocity dispersions ($\\sigma_{\\rm v}\\sim10^2$~\\kms), compact groups of galaxies are ideal environments for studying the mechanisms of interaction-induced star formation and nuclear activity.  From optical spectroscopic surveys, Hickson Compact Groups (HCGs) are known to host a population of galaxies with emission-line nuclear spectra characteristic of star-formation and/or active galactic nuclei (AGNs).  Based on the optical spectroscopic survey of \\citet{coziol98a,coziol98b}, the AGN fraction in HCGs is found to be $\\sim40\\%$, perhaps consistent with the $43\\%$ nuclear activity level found for nearby $M_V<-19$ galaxies (with greater detection sensitivity; \\altcite{HoEtal97a}) and significantly higher than the $\\sim1\\%$ AGN fraction identified optically in cluster galaxies (with $M_{V}<-21$; \\altcite{dressler85}).  Further, HCG AGNs (including low-luminosity AGNs, hereafter LLAGNs; HCGs host no known Seyfert~1-luminosity AGNs) are preferentially found in optically luminous, early-type galaxies with little or no ongoing star formation in the cores of evolved groups.  Similarly, many galaxies in clusters host LLAGNs in the local Universe \\citep[e.g.,][]{martini07}, in particular brightest cluster galaxies \\citep{best07}.  Typically, the most active star-forming galaxies are late-type spirals in the outskirts of groups. \\citet{coziol98b} interpreted relative distributions of star-forming and AGN-hosting galaxies as indicating a clear evolutionary scenario whereby group cores are mature collapsed systems in which the high galaxy densities led to increased gravitational interactions and hence more rapid exhaustion of gas reservoirs in the past through star formation.  (The radiatively weak AGNs in the evolved group cores require only a small amount of gas for fueling.)  A similar study by \\citet{shimada00} found that the fraction of emission-line HCG galaxies was comparable to the field, and conversely concluded that the HCG environment does {\\it not} trigger either star formation or AGN activity.  The initial expectation that the interactions evident in the compact group environment would naturally lead to markedly enhanced levels of star formation compared with the field has not been satisfied, and a coherent understanding of the history of gas and cold dust in compact groups has proven elusive.  An analysis of the CO content in HCG galaxies found them to be similar to those in the field, in loose groups, and in other environments; a notable exception is the $\\sim20\\%$ of HCG spirals that are CO deficient.  This suggests less, not more, star formation in HCGs compared to other environments \\citep{verdes98}. At the same time, the detection of a few HCG elliptical and S0 galaxies in both CO and the far-infrared (unlike typical galaxies of these types) suggests that tidal interactions are influencing galaxy evolution to some extent. While the far-infrared power is similar to comparison samples, the ratios of 25 to 100\\micron\\ \\iras\\ fluxes implies a greater number of intense, nuclear starbursts in HCG galaxies \\citep{verdes98}.  Clearly, a robust and consistent understanding of the impact of the compact group environment on galaxy properties has not yet emerged. One possible difficulty to date is the predominant use of optical emission-line studies to identify activity --- both star-forming and accretion-dominated.  Ground-based studies of bright galaxies can easily obscure low-contrast emission lines through their dilution by a strong stellar contribution \\citep[e.g.,][]{HoEtal97c}; intrinsic absorption can also mask spectroscopic signatures.  The clear discrepancy between the $\\sim1\\%$ AGN fraction in galaxy clusters from optical spectroscopic surveys \\citep[e.g.,][]{dressler85} compared to the larger fraction ($\\sim5\\%$ for luminous galaxies) revealed by X-ray observations \\citep[e.g.,][]{martini06} is illustrative of this problem. As highlighted by \\citet{martini07}, mismatched sample selection and detection techniques can create apparent (and false) discrepancies in AGN fractions between environments. In this paper, we take an alternate approach, focusing on the mid-infrared spectral energy distributions (SEDs) of individual galaxy nuclei to clearly identify the thermal, hot dust continua that signify neutral gas heated by ionizing photons from either young stars or an AGN.  The clear discrepancy between a blue, quiescent galaxy SED where the mid-infrared is dominated by the Raleigh-Jeans tail of stellar photospheres and the red, mid-infrared SED of warm to hot dust emission offers promise for reducing the ambiguity of previous compact group studies.  \\citet[][hereafter J07]{kelsey07} have presented the first results from a Cycle~1 \\spitzer\\ IRAC (3.6--8.0\\micron) and MIPS (24\\micron) imaging survey of \\total\\ galaxies in 12 nearby HCGs.  In brief, this work revealed trends between the evolutionary states of compact groups (determined from their dynamical and \\HI\\ masses), and their mid-infrared colors and luminosities.  Galaxies in relatively gas-rich groups tend to have colors most indicative of star formation and AGN activity, and galaxies in gas-poor groups predominantly exhibit a narrow range of mid-infrared colors that are consistent with the light from quiescent stellar populations.  The galaxies in this sample of 12 compact groups also occupy infrared color space in a distinctly different way than the population of galaxies in the {\\it Spitzer} First Look Survey (FLS; e.g., \\altcite{lacy04}), notably exhibiting a ``gap'' in their color distribution not found in the FLS sample.  All of these results suggest that the environment of a compact group is intimately connected to the mid-infrared activity of the member galaxies.  We present additional analysis of the nuclei of these galaxies, focusing in particular on developing diagnostics for identifying quantitative measures of mid-infrared activity with the ultimate goal of exploring the nature of mid-infrared emission, i.e., quiescent, star-forming, and/or accretion-powered.  We also explore the connections between mid-infrared nuclear properties and other features of individual galaxies including morphology, optical spectral type, and the host group's evolutionary stage.  Because of their proximity ($v<4500$\\kms), we can spatially extract nuclear photometry for the HCG galaxies to reduce contamination from the extended galaxy that is inevitable in surveys of more distant galaxies.  Throughout we assume a $\\Lambda$-CDM cosmology with $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$, and $H_0=70$\\kms\\,Mpc$^{-1}$ \\citep[e.g.,][]{sperg07}. ",
        "conclusions": "Of the complete set of \\total\\ HCG galaxies presented in the {\\em Spitzer} imaging survey of J07, we have identified 25 with red (\\alphairac$<0.0$) mid-infrared continua.  All eight known, spectroscopically identified star-forming galaxies (7a, 16c, 19c, 22c, 31ace, and 31b) are within this group. An additional three galaxies (16b, 61a, and 62b --- all optically identified AGNs) are apparently quiescent in the IRAC bands but show evidence for a 24\\micron-excess above a stellar continuum, bringing the total number of 24\\micron-active HCG galaxies to 61\\% (as all of the mid-infrared active galaxies are also 24\\micron-active).  This is higher than the fraction of emission-line galaxies (including AGNs and star-forming galaxies) identified in optical spectroscopic surveys of HCGs \\citep{coziol98a,coziol98b,shimada00}, and highlights the utility of multiwavelength studies to capture all activity.  Though the mid-infrared analysis presented here has the advantage of being unambiguous in terms of identifying mid-infrared activity -- the distribution of HCG nuclei is clearly bimodal in terms of both \\alphairac\\ and \\fracd\\ -- distinguishing between accretion and star formation as the dominant source of mid-infrared power remains challenging.  The evident trend in Figure~\\ref{fig:alphalum24} between \\alphairac\\ and \\ltwofour\\ seen in mid-infrared active galaxy nuclei could be interpreted as an increasing contribution from star formation at higher \\ltwofour\\ and redder (more negative) values of \\alphairac. Establishing this would be extremely powerful for interpreting photometric data of fainter objects in wide field surveys. However, the significant correlation between \\frachd\\ and \\leight, which might similarly be read to imply a stronger AGN component (because of the greater hot dust contribution at 4.5\\micron) at larger values of \\leight, is not generally consistent with the known nuclear identifications from optical spectroscopy.  This implies that simple measures of the steepness of the mid-infrared continuum are not diagnostic -- underscoring the difficulty of using mid-infrared photometry alone to categorize objects.  Even with high signal-to-noise mid-infrared spectroscopy, the underlying source of mid-infrared power can remain ambiguous in some objects \\citep[e.g.,][]{weedman05}. Furthermore, the strong stellar contamination at 3.6\\micron\\ in all of these galaxies means that color-color diagnostics that use the ratio of 3.6 to 4.5\\micron\\ fluxes to isolate AGNs \\citep[e.g.,][]{stern05} likely miss a large fraction of them, particularly at lower infrared luminosities. Selection criteria with a wider dynamic range in color will be less sensitive to this effect \\citep[e.g.,][]{lacy04}.  However, as noted by J07, none of our sample would fulfill the AGN mid-infrared color-color selection criteria of \\citet{stern05} or \\citet{lacy04}, which were designed to target objects where the AGN dominates the mid-infrared SED.  Our entire sample would also fail the power-law AGN selection criterion of \\citet{donley07} that required a monotonic flux increase from 3.6 to 8.0\\micron.  A few optically identified LLAGNs -- 22a, 42a, and 62a -- show no evidence for any excess mid-infrared (including 24\\micron) emission, indicating that warm to hot dust emission does not always accompany accretion onto a black hole. However, without a reservoir of cold material, an optically detectable AGN may be short-lived.  In any case, such a system is unlikely to contribute significantly to the luminous energy budget of its host galaxy.  Given their weak emission lines (which would be strongly diluted by host galaxy light at larger distances; \\altcite{moran02}), these three galaxies may be identified with the optically dull but X-ray detected AGNs found in clusters \\citep{martini06}.  The evolved nature of their host groups and the X-ray detections of 42a and 62a in particular are consistent with identifying these systems as mini-clusters.  In contrast, the less evolved compact groups may be less likely to host X-ray bright, optically quiet galaxies, as found in the loose group survey of \\citet{shen+07}.  At present, the lack of a complete suite of optical spectroscopy and high quality X-ray observations of our galaxy sample complicates definitively fitting the compact groups in context with the AGN surveys in the loose group and cluster environments done to date.  The striking difference between the distribution of \\alphairac\\ values for the HCG and SINGS galaxies supports the role of environment in affecting HCG galaxies --- a conclusion consistent with the marked difference in the mid-infrared color-color distribution of HCG and FLS galaxies noted by J07. In particular, the strongly bimodal nature of HCG nuclear (and galaxy) \\alphairac\\ values as well as the correlation between group \\HI\\ content and 24\\micron\\ activity suggests that star formation in a compact group galaxy, likely induced by interactions, occurs in a burst that then exhausts its reservoir of cold gas, as suggested previously from far-infrared \\citep{verdes98} and optical spectroscopic \\citep{coziol98b} studies. The end result is a group dominated by mid-infrared quiescent galaxies.  If local compact galaxy groups are analogous to the building blocks of clusters in the early Universe, this implies that the exhaustion of cold gas through enhanced star-formation occurs as a result of interactions (though not necessarily mergers) prior to cluster infall.  Any AGN activity in those galaxies at that point would proceed in a low-luminosity, X-ray bright mode \\citep{shen07}. However, Tully-Fisher \\citep{Mendes2003} and fundamental plane studies \\citep{delaRosa01} of HCG galaxies reveal that they are consistent with control galaxy populations.  Therefore, there is no strong evidence for a significant effect of the compact group environment on the current states of galaxy morphology in such a significant way that would alter the end states of galaxy evolution.  From a stellar population study of HCG ellipticals, \\citet{delaRosa07} concluded that these galaxies showed evidence for truncated star formation in comparison to a (very small) sample of field ellipticals. While this may be the case, the mechanism for star formation truncation is unlikely to be extensive feedback from AGN activity. None of the optically known AGNs in any HCG has sufficient luminosity to clear a galaxy of its interstellar medium.  The sole possible exception to this observation is 62a, where X-ray cavities in the intragroup medium suggest that a relativistic jet has ejected a sigificant amount of kinetic energy from a radiatively quiet AGN \\citep[e.g.,][]{morita+06,gu+07}.  However, this ongoing process likely postdates the epoch of star formation in this evolved group."
    },
    "0710/0710.0002_arXiv.txt": {
        "abstract": "The spectral shape of the unresolved emission from different classes of gamma-ray emitters can be used to disentangle the contributions from these populations to the extragalactic gamma-ray background (EGRB). We present a calculation of the unabsorbed {\\em spectral shape} of the unresolved blazar contribution to the EGRB starting from the spectral index distribution (SID) of resolved EGRET blazars derived through a maximum-likelihood analysis accounting for measurement errors. In addition, we explicitly calculate the {\\em uncertainty} in this theoretically predicted spectral shape, which enters through the spectral index distribution parameters.  We find that: (a) the unresolved blazar emission spectrum is only mildly convex, and thus, even if blazars are shown by GLAST to be a dominant contribution to the ERGB at lower energies, they may be insufficient to explain the EGRB at higher energies; (b) the theoretically predicted unresolved spectral shape involves significant uncertainties due to the limited constraints provided by EGRET data on the SID parameters, which are comparable to the statistical uncertainties of the observed EGRET EGRB at high energies; (c) the increased number statistics which will be provided by GLAST will be sufficient to reduce this uncertainty by at least a factor of three. ",
        "introduction": "The isotropic, and presumably, extragalactic gamma-ray background emission (EGRB) detected by the Energetic Gamma-ray Experiment Telescope (EGRET) aboard the {\\it Compton Gamma-ray Observatory} (Sreekumar et al.\\ 1998) is one of the most important observational constraints on known or theorized populations of faint, unresolved gamma-ray emitters. With the imminent launch of the {\\it Gamma-ray Large Area Space Telescope} (GLAST), which is expected to represent an  unprecedented leap in observational capabilities in GeV energies, the timing is especially opportune to consider the information content of the diffuse background and methods for maximizing the scientific return from its study.  One of the primary challenges in using the EGRB to constrain properties of extragalactic gamma-ray emitters and exotic physics is disentangling the convolved contributions of guaranteed participating populations. Estimates of the levels of the collective unresolved emission even from established classes of extragalactic sources (such as blazars and normal galaxies) involve significant uncertainties and are at the order-of-magnitude level at best (e.g., Padovani et al.\\ 1993; Stecker \\& Salamon 1996; Kazanas \\& Perlman 1997; Mukherjee \\& Chiang 1999; M\\\"ucke \\& Pohl 2000; Dermer 2006; Lichti et al.\\ 1978; Pavlidou \\& Fields 2002).  A very promising approach for the study of the EGRB and its components is through the use of spectral shape information. Let us consider the optimal case where the expected spectral shapes of the unresolved emission from known classes of gamma-ray sources can be confidently predicted. In this case, a series of conclusions can be drawn regarding the potential contributions of these classes to the EGRB even without detailed calculations of the {\\em magnitudes} of their collective emission. For example, in comparing the spectral shape of the spectrum due to a particular class with that of the EGRB, one can identify whether this class could, in principle, comprise most of the EGRB or require the existence of contributions from other classes (Stecker \\& Salamon 1996a,b; Strong et al. 2004; Pavlidou \\& Fields 2002)\\footnote{Note that the method of spectral comparison can only be used to {\\em reject} a population from being the sole source of the EGRB; spectral consistency does not constitute in itself proof of the importance of a class of objects as an EGRB contributor, since the overall normalization of the emission may, in fact, be low depending on the gamma-ray luminosity function of the population.}. Potentially identifiable spectral features could be predicted and searched for in GLAST data (e.g., Pavlidou \\& Fields 2002). Finally, spectral information could be used to ultimately disentangle different components and contributions (as in e.g. de Boer et al.\\ 2004 for the case of the diffuse emission from the Milky Way). An additional attractive feature of such calculations is that the associated uncertainties are largely independent of those entering the calculations of the overall unresolved emission flux.  As blazars are the most populated class of gamma-ray emitters, unresolved blazars are guaranteed to contribute significantly, if not dominantly, to the EGRB. Thus, it is especially important to understand the expected spectral shape of their collective unresolved emission and the uncertainties involved in its calculation. Individual blazars have been measured to have power-law spectra in the EGRET range, $F_E \\propto E^{-\\alpha}$, with spectral index $\\alpha$ ranging approximately from 1.5 to 3. The unresolved emission from a collection of power-law emitters with variable spectral indices\\footnote{Statistically significant spread in the observed and intrinsic spectral index of blazars has also been confirmed in other energy bands (see e.g.\\ Shen et al.\\ 2006 for the case of X-ray emission).} has, invariably, a convex spectral shape (Brecher \\& Burbidge 1972; Stecker \\& Salamon 1996; Pohl et al.\\ 1997; Pavlidou et al.\\ 2007). The exact shape of the unresolved spectrum depends on the spectral index distribution (SID) of gamma-ray loud blazars. Recognizing that this is the case, Stecker \\& Salamon (1996a) explicitly reconstructed a spectrum from the observed SID of blazars in the 2nd EGRET catalog (Thompson et al.\\ 1995), deriving a spectrum that was indeed significantly convex. However, measurement errors in individual spectral indices smear the SID and exaggerate the curvature of the spectrum (Pohl et al.\\ 1997).  Recently, in Venters \\& Pavlidou 2007 (hereafter VP07),  we have applied a maximum-likelihood analysis to recover the intrinsic spectral index distribution (ISID) of gamma-ray loud blazars from EGRET observations. We found that (1) the maximum-likelihood ISID is appreciably narrower than the observed SID, so the best-guess spectrum is likely to have only a mild curvature; (2) BL Lacs and flat spectrum radio quasars (FSRQs) are likely to be spectrally distinct populations with spectrally distinct contributions to the EGRB; (3) there is no evidence for a systematic shift of spectral variability with flaring implying that although variability may be important in the level of the contribution from blazars, ignoring variability effects in spectral shape studies is likely to be a good approximation.  Here, we use the ISIDs derived in VP07 to calculate the spectral shape of the blazar contribution to the EGRB. We examine the sensitivity of the shape to the exact values of the ISID parameters and report on the range of possible shapes given our uncertainties in the determination of these parameters. We also investigate how the spectral shapes of the BL Lac and FSRQ contributions may differ.  Finally, we predict how our understanding of the spectral shape of the unresolved blazar emission will improve after GLAST observations become available. ",
        "conclusions": "\\label{Disc}  We have calculated the expected spectral shape of the unresolved gamma-ray emission from blazars under the assumptions that the ISIDs of blazars do not evolve with redshift and are independent of blazar luminosity and flaring state. We have also explicitly calculated the $1\\sigma$ uncertainty in the spectral shape entering through the limited constraints on the blazar ISIDs derived from EGRET data. Finally, we have predicted by how much these uncertainties will be reduced if GLAST observations are used to determine the blazar ISIDs.  The unresolved emission spectral shape can be used as an indicator of the potential importance of a given population's contribution to the EGRB, and  it constrains the maximal contribution at high energies relative to that at low energies.  If the curvature tentatively seen in the observed EGRB is real, then a population with little such curvature in its unresolved spectrum (such as the FSRQs) will not be the dominant contributor to the EGRB at high energies even if it is dominant at low energies. The unresolved BL Lac spectrum does seem to be more convex than that of the FSRQs, but the level of uncertainty in the spectral shape is high in this case because only a few BL Lacs were detected by EGRET. However, GLAST observations will dramatically improve our understanding of the blazar ISIDs and the associated unresolved spectral shapes allowing us to use spectral shape information to calculate the minimal additional contribution required from other classes of sources to explain the observed EGRB spectrum.  It should be stressed that here we have only calculated {\\em unabsorbed} spectra. At energies higher than 10 ${\\rm \\, GeV}$, gamma-ray absorption through interactions (pair production) with the extragalactic background light becomes important (e.g., Salamon \\& Stecker 1998), and the spectral shape of any contribution to the EGRB will be accordingly changed. We will return to this effect in a future publication."
    },
    "0710/0710.2225_arXiv.txt": {
        "abstract": "We have determined the mineralogical composition of dust in the Broad Absorption Line (BAL) quasar \\pgbal\\ using mid-infrared spectroscopy obtained with the {\\em Spitzer Space Telescope}. From spectral fitting of the solid state features, we find evidence for Mg-rich amorphous silicates with olivine stoichiometry, as well as the first detection of corundum (Al$_2$O$_3$) and periclase (MgO) in quasars. This mixed composition provides the first direct evidence for a clumpy density structure of the grain forming region.  The silicates in total encompass ($56.5 \\pm 1.4$)\\% of the identified dust mass, while corundum takes up ($38 \\pm 3$) wt.\\%. Depending on the choice of continuum, a range of mass fractions is observed for periclase ranging from ($2.7 \\pm 1.7$)\\% in the most conservative case to ($9 \\pm 2$)\\% in a less constrained continuum.  In addition, we identify a feature at 11.2 $\\mu$m as the crystalline silicate forsterite, with only a minor contribution from polycyclic aromatic hydrocarbons. The ($5 \\pm 3$)\\% crystalline silicate fraction requires high temperatures such as those found in the immediate quasar environment in order to counteract rapid destruction from cosmic rays. ",
        "introduction": "\\label{sec:intro}  In the local Universe, Asymptotic Giant Branch (AGB) stars are believed to dominate dust production \\citep[e.g.,][]{D_03_dust}.  AGB stars are a late phase of stellar evolution for stars with initial masses of $M$=1--8\\msun, which develop outflows where the low temperatures and high densities become optimal for dust condensation. However, a significant amount of dust originating from AGB stars does not build up within $\\sim1$~Gyr of the birth of the first generation of low-mass stars \\citep{ME_03_earlydust}.  Nevertheless, large quantities of dust are clearly present at these early times: quasar host galaxies at $z\\sim6$ show evidence for 10$^{8-9}$~\\msun\\ of dust heated by star formation in their host galaxies as seen in submm and far-infrared emission \\citep{PriddeyEtal2003,BCB_06_dust}. Furthermore, extinction curves for a $z=6.2$ quasar and GRB~050904 ($z=6.29$) are not consistent with those locally \\citep{maio04,stratta07}.  Another mechanism for dust production is therefore required, and \\citet{elvis-dust} explored the possibility that quasar winds might provide environments suitable for efficient dust formation.  They determined that the temperatures and pressures in these regions could reach the values found in cool, dust-producing stars.  The quasar winds are predicted to produce dust masses up to $10^7$ \\msun\\ \\citep{elvis-dust}, implying that the quasar wind may in some cases only account for part of the dust production, and supernovae have been suggested as an alternative source of dust in high-z galaxies \\citep{SEB_06_SNdust}.  Quasar outflows are most obviously manifested in Broad Absorption Line (BAL) quasars.  This population, approximately 20\\% of optically selected type~1 quasar samples \\citep[e.g.,][]{HewFol2003,reich+03a}, is notable for broad, blueshifted absorption evident in common ultraviolet resonance transitions such as \\CIV, Ly$\\alpha$, and \\OVI. These P~Cygni-type features arise because the observer is looking through an outflowing wind. \\balqs\\ are thus a natural population to consider when investigating the grain properties of dust in the quasar environment.  In this letter, we present a detailed analysis of the mid-infrared spectrum of the luminous \\balq, \\pgbal, in order to determine its dust composition.  \\pgbal, an {\\em IRAS} source, is mid-infrared bright, and one of the most luminous low-redshift ($z=0.466$) Bright Quasar Survey objects with $M_{\\rm V}=-26.9$ (\\hnot$=70$\\hunits, $\\Omega_{\\rm M} = 0.3$, and $\\Omega_{\\Lambda} = 0.7$ are assumed throughout).  \\hst\\ spectra revealed broad, shallow \\CIV\\ absorption, and it has been well-studied in the UV and X-ray \\citep[e.g.,][]{gall+04}. ",
        "conclusions": "For the first time, the composition of the dust in a quasar has been determined, albeit limited to the components with resonances in the infrared. While the origin of the dust could lie in stellar ejecta, the dust properties are consistent with their formation in the quasar wind itself.  The dust in \\pgbal\\ clearly bears similarity to dust in other astrophysical environments, i.e.~in the properties of the amorphous silicates, but the presence of large amounts of MgO and Al$_2$O$_3$ sets it apart from the composition of interstellar dust in the local universe.  Assuming the dust is formed in the quasar wind, the co-existence of the highly refractory corundum, the crystalline silicates and and the less refractory MgO and amorphous silicates indicates that the wind of \\pgbal\\ has an inhomogeneous temperature and density structure."
    },
    "0710/0710.0827_arXiv.txt": {
        "abstract": "{Multiple systems are the product of protostellar core fragmentation. Studying their statistical properties in young stellar populations therefore probes the physical processes at play during star formation.} {Our project endeavors to obtain a robust view of multiplicity among embedded Class\\,I and Flat Spectrum protostars in a wide array of nearby molecular clouds to disentangle ``universal'' from cloud-dependent processes.} {We have used near-infrared adaptive optics observations at the VLT through the $H$, $K_s$ and $L'$ filters to search for tight companions to 45 Class\\,I and Flat Spectrum protostars located in 4 different molecular clouds (Taurus-Auriga, Ophiuchus, Serpens and L1641 in Orion). We complemented these observations with published high-resolution surveys of 13 additional objects in Taurus and Ophiuchus.} {We found multiplicity rates of 32$\\pm$6\\% and 47$\\pm$8\\% over the 45--1400\\,AU and 14--1400\\,AU separation ranges, respectively. These rates are in excellent agreement with those previously found among T\\,Tauri stars in Taurus and Ophiuchus, and represent an excess of a factor $\\sim$1.7 over the multiplicity rate of solar-type field stars. We found no non-hierarchical triple systems, nor any quadruple or higher-order systems. No significant cloud-to-cloud difference has been found, except for the fact that all companions to low-mass Orion protostars are found within 100\\,AU of their primaries whereas companions found in other clouds span the whole range probed here.} {Based on this survey, we conclude that core fragmentation always yields a high initial multiplicity rate, even in giant molecular clouds such as the Orion cloud or in clustered stellar populations as in Serpens, in contrast with predictions of numerical simulations. The lower multiplicity rate observed in clustered Class\\,II and Class\\,III populations can be accounted for by a universal set of properties for young systems and subsequent ejections through close encounters with unrelated cluster members.} ",
        "introduction": "The prevalence of binary and multiple systems among stellar populations in our Galaxy is generally understood as a consequence of the natural tendency of prestellar cores for fragmentation during or immediately after their free-fall collapse (see Tohline 2002 for a detailed review). Numerical simulations have long predicted that this fragmentation, hence the frequency and properties of multiple systems, is strongly dependent on the initial conditions reigning in the core (e.g., Bonnell et al. 1992; Durisen \\& Sterzik 1994; Boss 2002; Goodwin et al. 2004b). Millimeter observations over the last two decades have provided a detailed view of the initial conditions of star formation. For instance, the basic properties of individual prestellar cores, outer radius and central density, differ significantly from one molecular cloud to another (Motte \\& Andr\\'e 2001). It is also likely that the temperature in prestellar cores vary from cloud to cloud depending on the strength of the interstellar radiation field (e.g., Stamatellos et al. 2007). Finally, there is now good evidence that protostellar collapse is generally more violent in a cluster-forming environment than in isolated dense cores, e.g., induced by strong external disturbances as opposed to spontaneous or self-initiated (Andr\\'e et al. 2003; Belloche et al. 2002, 2006). Considering these differences, one may therefore expect to observe substantial differences in the properties of multiple systems in independent stellar populations.  Early high-angular resolution surveys of Myr-old T\\,Tauri stars in nearby T associations, such as Taurus and Ophiuchus (Ghez et al. 1993; Leinert et al. 1993; Reipurth \\& Zinnecker 1993) revealed a significant excess over the multiplicity rate among field stars (Duquennoy \\& Mayor 1991, hereafter DM91; Fischer \\& Marcy 1992). It was then rapidly discovered that clustered populations of equally young stars, such as the Orion Nebula Cluster (ONC), do not possess such a high multiplicity rate and rather resemble field stars from this point of view (Padgett et al. 1997; Petr et al. 1998; Duch\\^ene et al. 1999). While this could be evidence that supports the environment-dependent fragmentation of prestellar cores, alternative explanations cannot be excluded. In particular, dynamical interactions with unrelated cloud members can disrupt most wide companions in less than 1\\,Myr in the densest clusters (Kroupa 1995). In other words, T\\,Tauri multiple systems have had sufficient time to evolve since their formation, so that their observed properties may not be considered as directly representative of core fragmentation. To circumvent this problem, observations of less evolved, embedded, young stellar objects (YSOs) are required in order to determine the pristine properties of multiple systems.  Radio interferometric studies of embedded Class\\,0 and Class\\,I sources have hinted at a high overall multiplicity rate, comparable to that observed among T\\,Tauri stars in non-clustered populations (Looney et al. 2000; Reipurth et al. 2002; 2004). Non uniform and limited samples, as well as the fact that not all YSOs emits strongly at centimeter wavelengths, prevent conclusive comparisons at this point, however. In parallel to this effort, Haisch et al. (2002, 2004) and Duch\\^ene et al. (2004, hereafter D04) conducted direct near-infrared imaging surveys of Class\\,I protostars. These surveys also found a high multiplicity rate, consistent with that of somewhat more evolved T\\,Tauri stars in the same clouds. D04 found marginal evidence that the multiplicity rate of these sources decreases on a timescale of $\\sim10^5$\\,yr, possibly a result of the internal or external dynamical ejection of wide companions. No strong evidence was found for a variation of the multiplicity rate of Class\\,I sources from one cloud to another, in part because of small sample sizes and of the limited range of projected separations probed by these surveys. Whether fragmentation does indeed depend on environmental conditions remains to date a theoretical/numerical predictions that has not yet been confirmed observationally. In particular, the outcome of core fragmentation in a very rich molecular cloud such as Orion remains unknown.  To extend the analysis of D04, we have undertaken a high-angular resolution survey of embedded Class\\,I protostars sampling the Serpens and Orion molecular clouds in addition to Taurus and Ophiuchus. This paper is organized as follows: we present our sample and observations in Section\\,\\ref{sect:obs} and the results in Section\\,\\ref{sect:results}. In Section\\,\\ref{sect:discus}, we analyze these results in view of other multiplicity surveys and of predictions of star formation theories. We summarize our findings in Section\\,\\ref{sect:concl}. ",
        "conclusions": "\\label{sect:concl}  We have used diffraction-limited imaging from 1.6 to 3.7\\,$\\mu$m with the 8m-VLT adaptive optics system to search for tight companions around 45 embedded Class\\,I and FS protostars in the Taurus, Ophiuchus, Serpens and L1641 (Orion\\,A) molecular clouds. We complement our analysis with published high-resolution surveys of similar objects in Taurus and Ophiuchus to build a sample of 58 Class\\,I and FS targets. We derive an average multiplicity rate of 32$\\pm$6\\% over the separation 45--1400\\,AU. In the Taurus and Ophiuchus clouds, the closest clouds in our sample, we further derive a multiplicity rate of 47$\\pm$8\\% over the separation range 14--1400\\,AU. These rates are a factor of $\\sim1.7$ higher than those derived for nearby solar-type field stars, extending the multiplicity excess found among several populations of T\\,Tauri star to even younger ages. Most importantly, we find that the embedded protostars in L1641 show a multiplicity rate similar to that in other clouds, indicating for the first time that a high multiplicity rate is achieved after core fragmentation in all types of nearby molecular clouds, including giant molecular clouds such as Orion\\,A, which also hosts the dense ONC cluster. We also find a high multiplicity rate in Serpens, the densest cluster in our survey.  Our results support the view that core fragmentation results in a multiplicity rate for wide companions that does not depend on the initial conditions reigning in the cores, as opposed to predictions of most numerical simulations. Rather, our findings support a scenario in which all YSO populations start with a similar set of multiplicity properties and only evolve as a consequence of disruptive system-system interactions prior to dilution of the clusters in the field. Follow-up multiplicity surveys of the Class\\,II and Class\\,III populations of the Serpens and L1641 clouds would provide key empirical verifications of this scenario. We found 6 triple systems, all of them hierarchical, and higher-order systems are very rare among Class\\,I/FS sources within the separation ranges studied here. It is unlikely that many companions will be ejected as a result of internal decay of unstable systems past the Class\\,I phase. Finally, we also find a possible trend for the embedded Orion multiple systems to have different orbital properties than those in other clouds, namely systematically tighter projected separations, a hint that environmental conditions may impact on the properties of protobinaries."
    },
    "0710/0710.5891_arXiv.txt": {
        "abstract": "{ Spectral differential imaging is an increasingly used technique for ground-based direct imaging searches for brown dwarf and planetary mass companions to stars. The technique takes advantage of absorption features that exist in these cool objects, but not in stars, and is normally implemented through simultaneous narrow-band imagers in 2 to 4 adjacent channels. However, by instead using an integral field unit, different spectral features could be used depending on the actual spectrum, potentially leading to greater flexibility and stronger detection limits. In this paper, we present the results of a test of spectral differential imaging using the SINFONI integral field unit at the VLT to study the nearby active star L449-1. No convincing companion candidates are found. We find that the method provides a $3 \\sigma$ contrast limit of 7.5 mag at 0.35\", which is about 1.5 mag lower than for NACO-SDI at the same telescope, using the same integration time. We discuss the reasons for this, and the implications. In addition, we use the SINFONI data to constrain the spectral type in the NIR for L 449-1, and find a result between M3.0 and M4.0, in close agreement with a previous classification in the visual range. ",
        "introduction": "In recent years, developments in techniques and instrumentation have led to a strongly increased capacity for high-contrast imaging of substellar companions at small angular separations to stars, from the ground. In particular, the use of high-order adaptive optics (AO) in combination with various differential imaging techniques provides a sensitivity for companions that are $10^4$-$10^5$ times fainter than the primary at a separation of 0.5\"-1\" at near-infrared wavelengths (see Janson et al. 2007 and Biller et al. 2007). According to theoretical models (see e.g. Baraffe et al. 2003), this corresponds to objects of a few times the mass of Jupiter around young stars.  A particularly efficient technique for such observations is simultaneous spectral differential imaging (SDI, see e.g. Rosenthal et al. 1996 and Racine et al. 1999), in which the flux from a star is observed in two narrow, adjacent wavelength bands simultaneously. The bands are chosen such that one is inside and one outside of an absorption feature which arises uniquely in cool atmospheres. One example of such a feature is the methane band at 1.6 $\\mu$m, which only occurs in dwarfs of spectral type T or later, whereas a stellar spectrum is flat in this range. By subtracting one image from the other, the stellar point spread function (PSF) can be largely subtracted out, since it is approximately flat around 1.6 $\\mu$m. Along with the main PSF, most of the random PSF substructure (speckle noise) is subtracted out as well, since the images are simultaneous and hence the atmospheric speckle pattern is the same in both images.  NACO-SDI at the VLT is an excellent instrument for the purpose of SDI observation (see Lenzen et al. 2004). NACO-SDI images the same field in three different narrow bands around the 1.6 $\\mu$m methane feature simultaneously. It is designed to minimize non-common path aberrations, and has been shown to be capable of achieving a contrast of more than 13 mag between star and companion at 1\" angular separation at the 3$\\sigma$ level (see Janson et al. 2007), without the use of a coronagraph. One limitation of the NACO-SDI instrument is that it can only attain information about the methane feature in question. It has been suggested (Berton et al. 2006) that if $N > 1$ different absorption features are used at once for SDI purposes, the achievable contrast would increase by a factor $N^{1/2}$ for the same integration time, simply due to the more efficient use of photons. Hence, in principle, an instrument with the capacity to do multi-feature SDI (MSDI) would be preferable to NACO-SDI, if the instruments perform equally well in all other respects (in terms of differential aberrations, etc.).  SINFONI is an integral field spectroscopy instrument with AO capability at the VLT (see Eisenhauer et al. 2003 and Bonnet et al. 2004). It uses an image slicer to divide the FOV into pseudo-slits that are individually dispersed on a grating, hence retaining both the spatial and spectral information of the incoming light. The output data can be used to construct a data cube that stores spatial information along two axes, and spectral information along the third. In this way, SINFONI can be used as an SDI or MSDI instrument, since the cube contains simultaneous narrow-band images over a large range of wavelengths.  We have acquired SINFONI data to test its capacity for SDI. In the following sections, we describe the observations and data reduction, and compare the results with NACO-SDI. We also compare the results to a spectral deconvolution (SD) scheme applied to another set of SINFONI data by Thatte et al. (2007). We dicuss the advantages and disadvantages of using SINFONI for high-contrast imaging and characterization of exoplanet and brown dwarf companions to stars. Finally, we analyze the non-differential collapsed images, and constrain the spectral type of L 449-1 from H- and K-band spectroscopy. ",
        "conclusions": "We have investigated the potential of the SDI technique with SINFONI for detecting cool substellar companions to stars. It was discovered that the performance is considerably worse than for NACO-SDI under very similar circumstances. In addition, the error is partly static, leading to a poor increase of performance with increasing integration time. The quasi-static nature of the noise, along with the fact that it is constant over most of the image space, implies that non-common path aberrations and errors from the slicing and reconstruction of the image dominate the residual noise. These results, along with practical considerations, clearly lead to the suggestion that ``blind'' differential imaging searches for cool companions (i.e., when no a priori information is available) are presently best performed with specialized instruments such as NACO-SDI, whereas follow-up observations, where more qualitative information is desired, are better suited for integral field units such as SINFONI. The results are highly relevant for the design of the next-generation instruments for planet detection.  While the collapsed broad-band images of L 449-1 imply the existence of a low-contrast candidate companion, the absence of a clear spectral signature associated with this feature suggests that it is a PSF artefact rather than a physical object.  The spectral type of L449-1, as determined from the H- and K-band spectra, was found to be in the range of M3.0-M4.0, in close agreement with previous results (M4, Scholz et al. 2005) in the visual range."
    },
    "0710/0710.0961_arXiv.txt": {
        "abstract": "The transient X-ray binary pulsar A0535+262 was observed with $Suzaku$ on 2005 September 14 when the source was in the declining phase of the August-September minor outburst. The $\\sim$103 s X-ray pulse profile was strongly energy dependent, a double peaked profile at soft X-ray energy band ($<$ 3 keV) and a single peaked smooth profile at hard X-rays. The width of the primary dip is found to be increasing with energy. The broad-band energy spectrum of the pulsar is well described with a Negative and Positive power-law with EXponential (NPEX) continuum model along with a blackbody component for soft excess. A weak iron K$_\\alpha$ emission line with an equivalent width $\\sim$ 25 eV was detected in the source spectrum. The blackbody component is found to be pulsating over the pulse phase implying the accretion column and/or the inner edge of the accretion disk may be the possible emission site of the soft excess in A0535+262. The higher value of the column density is believed to be the cause of the secondary dip at the soft X-ray energy band. The iron line equivalent width is found to be constant (within errors) over the pulse phase. However, a sinusoidal type of flux variation of iron emission line, in phase with the hard X-ray flux suggests that the inner accretion disk is the possible emission region of the iron fluorescence line. ",
        "introduction": "Be X-ray binaries consist of a neutron star in an eccentric orbit around a Be star companion. The orbit of the Be X-ray binaries is generally wide and eccentric with orbital periods in the range of 16 days to 400 days (Coe 2000). Mass transfer from the Be companion to the neutron star takes place through the circumstellar disk. When the neutron star passes through the disk or during the periastron passage, it shows strong outbursts with an increase in X-ray luminosity by a factor $\\geq$ 100 (Negueruela 1998).  A0535+262 is a 103 s Be/X-ray binary pulsar discovered by $Ariel~V$ during a large outburst in 1975 (Coe et al. 1975). The binary companion HDE~245770 is an O9.7-B0 IIIe star in a relatively wide eccentric orbit ($e$ = 0.47) with orbital period of $\\sim$ 111 days and at a distance of $\\sim$ 2 kpc (Finger et al. 1996; Steele et al. 1998). The pulsar shows regular outbursts with the orbital periodicity. Occasional giant X-ray outbursts are also observed when the object becomes even brighter than the Crab. The pulsar shows three typical intensity states, such as quiescence with flux level of below 10 mCrab, normal outbursts with flux level in the range 10 mCrab to 1 Crab, and giant outbursts during which the object becomes the brightest X-ray source in the sky with the flux level of several Crab (Kendziorra et al. 1994).  103 s pulsations were detected during past X-ray observations of A0535+262 in quiescence, outburst, and giant outbursts. The pulse profile was single peaked in quiescence (in 1-10 keV range; Mukherjee \\& Paul 2005, 3-20 keV; Negueruela et al. 2000), double peaked with a clearly asymmetric ''main'' and a more symmetric ''secondary'' pulse during the X-ray outbursts (Mihara 1995; Kretschmar et al. 1996, Maisack et al. 1997). The X-ray spectrum of the pulsar has been studied at soft and hard X-rays with various instruments at different luminosity levels. The spectrum of the object, during outbursts, shows cyclotron resonant scattering features at higher energies than that of the other pulsars. Two harmonic features at around 50 keV and 100 keV were detected in its 1989 outburst with the HEXE/TTM instrument on Mir/Kvant (Kendziorra et al. 1994). The CGRO/OSSE observations of 1994 outburst of the pulsar showed a significant absorption feature at 110 keV (Grove et al. 1995). These detections did not resolve whether the magnetic field of the pulsar is $\\sim$ 5 $\\times$ 10$^{12}$ G (when the fundamental occurs at 55 keV) or $\\sim$10$^{13}$ G (for the 110 keV fundamental).  The most recent major X-ray outburst was detected in 2005 May/June with the BAT instrument on Swift when the 15-195 keV count rate was greater than 3 times that of the Crab Nebula (Tueller et al. 2005). Following the detection of the recent outburst, the pulsar was observed by INTEGRAL, RXTE and the recently launched $Suzaku$. A cyclotron resonance feature at $\\sim$45 keV was detected in the Hard X-ray Detector (HXD) spectrum of $Suzaku$ with estimated magnetic field of the pulsar as $\\sim$ 4 $\\times$ 10$^{12}$ Gauss (Terada et al. 2006). The detection of the absorption feature at $\\sim$45 keV and its first harmonic at $\\sim$100 keV is reported from the INTEGRAL and RXTE observations of the pulsar during the 2005 August/September outburst (Caballero et al. 2007). Using the same Suzaku observation used for the analysis of the cyclotron resonance feature (Terada et al. 2006), we study the broad-band spectral properties of the pulsar in the present paper. ",
        "conclusions": "Apart from two RXTE observations and three BeppoSAX observations during the quiescence, all other observations reported before 2005 August were made during the giant outbursts when the source flux was about equal to or more than that of the Crab nebula. The RXTE and INTEGRAL observations of the pulsar, during the 2005 August-September minor outburst, confirmed the earlier discovery of a cyclotron feature at $\\sim$45 keV along with the detection of its first harmonic at $\\sim$100 keV (Caballero et al. 2007). Suzaku observation of A0535+262 during the declining phase of the 2005 August-September normal outburst provided new data on spectral properties of the pulsar up to $\\sim$100 keV and how they might evolve throughout the outburst.  \\subsection{Pulse Profile}  During the present observation of the minor outburst, the pulse profile of A0535+262 is different from that in the quiescence or during the earlier reported outbursts. The pulse profile in the hard energy band is known to show energy and luminosity dependence in outbursts, which is described in Bildsten et al. (1997) based on the BATSE observations. The profile tends to be double peaked when the luminosity is high, eg.\\ $>\\!10^{37}$ erg s$^{-1}$ (20--100 keV), whereas it becomes a single peak at $4\\times10^{36}$ erg s$^{-1}$. The current $Suzaku$ observation had even lower luminosity ($2\\times10^{35}$ erg s$^{-1}$ in 20--100 keV band), and showed almost sinusoidal profile above 25~keV\\@. The $Suzaku$ hard band observation found a single peaked profile, as was seen by BATSE.  Hard X-ray pulse profiles of A0535+262, obtained from the RXTE and INTEGRAL observations of the same 2005 outburst with a peak luminosity of $\\sim$10$^{37}$ erg s$^{-1}$, appear to be double-peaked (Caballero et al. 2007). The two peaks with a shallow valley constitute a trapezoid-like profile above 35 keV separated by a wide trough. Though the observations were carried out during the peak and the decay of the outburst, the added pulse profile may be dominated by the profile during the peak of the outburst. Therefore, the profile presented by Caballero et al. (2007) probably reflect that at the higher luminosity observations (at the peak of the outburst). The double-peaked hard X-ray profiles at a source luminosity of $\\sim$10$^{37}$ erg s$^{-1}$ agrees with the luminosity dependence of pulse profiles observed with BATSE. The profiles obtained from Suzaku observation, however, appear to be different from that of the RXTE and INTEGRAL observations. This is interpreted due to the luminosity dependent effect as the Suzaku observation was at a luminosity level of two orders of magnitude less compared to the peak luminosity.  On the other hand, the Suzaku pulse profile in the soft energy band ($<\\!10$ keV) is found to be similar to that obtained from the $Ginga$ observation of the 1989 outburst (Mihara 1995). The dip-like structure, prominent only at soft X-rays is also recognized in the $Ginga$ profile. Because the $Ginga$ observation was made during the outburst when the source was $\\sim$80 times brighter than the present observation, the profile may be interpreted that A0535+262 in outbursts generally contains a dip-like structure at soft X-rays as seen in 2005 $Suzaku$ observation. This type of structure is also seen in the RXTE observation of the pulsar during the same minor outburst in 2005 August/September (Caballero et al. 2007). This dip-like structure is absent in the pulse profiles when the pulsar was in quiescence (Mukherjee \\& Paul 2005). As other observations of A0535+262 in giant outbursts focused on the hard X-ray energy ranges, the dip-like structure could have been missed in the pulse profiles.  \\subsection{Phase-Averaged Spectroscopy}  The X-ray spectrum of A0535+262 has been described by different continuum models at different energy ranges. Hard X-ray observations during outbursts, meant for the study of cyclotron absorption features and hence the pulsar magnetic field, were fitted by different continuum models such as optically thin thermal bremsstrahlung, power-law, power-law with an exponential cut-off, Wien's law or different combinations of above components (dal Fiume et al.\\ 1988; Kendziorra et al.\\ 1994; Grove et al.\\ 1995), whereas the 2--37 keV spectrum, obtained from the $Ginga$ observation of the 1989 outburst, was well fitted with NPEX continuum model (Mihara 1995). Current analysis of the $Suzaku$ data showed that the NPEX continuum model fits well to the source spectrum.  Selection of an appropriate continuum model may be crucial to investigate the broad features in the spectrum, such as the soft excess represented by a blackbody component. The spectral fitting to the 0.3--10~keV BeppoSAX spectra, during quiescence, yielded a blackbody component of temperature $\\sim$1.3~keV along with a power-law component of index $\\sim$1.8 (Mukherjee \\& Paul 2005). On the other hand, the Suzaku data when fitted with NPEX continuum model yielded a much lower temperature of $\\sim$0.16 keV, which is common for most of the X-ray pulsars. We suspect that the difference arises from the different selection of the continuum model. In fact, when the Suzaku wide band spectra were fitted by a power-law with an exponential cutoff, we did indeed obtain a blackbody temperature of 1.36 keV. This demonstrates the importance of the continuum selection to quantify the blackbody component. However, the selection criteria for the continuum choice have yet to be rigorously defined.  Most of the transient Be/X-ray binary pulsars undergo periodic outbursts due to the enhanced mass accretion when the neutron star passes through the dense regions of the circumstellar disk. During the passage, there is every possibility of the increase in the absorption column density. In case of A0535+262, we found that the absorption column density was higher than the Galactic column density towards A0535+262 (5.9 $\\times$ 10$^{21}$ atoms cm$^{-2}$) during the minor outburst in 2005 August-September. Note that the $Suzaku$ observation was made near the apastron at the orbital phase of 0.42--0.43. From BeppoSAX observations of the pulsar in quiescence (orbital phases 0.77, 0.05, 0.42), though, the estimated value of the absorption column density was found to be similar to that of the Galactic column density (Mukherjee \\& Paul 2005). During the 1998 RXTE/PCA observations (in quiescence, orbital phases of 0.0 and 0.77), the value of the absorption column density was found to be a factor of about 10 higher than that of the Galactic column density (Negueruela et al. 2000). During the $Ginga$ observation of the pulsar in the 1989 outburst (orbital phase of 0.96), the estimated absorption column density was found to be 4.8$^{+1.7}_{-1.2}$ $\\times$ 10$^{21}$ atoms cm$^{-2}$ (Mihara 1995), which is compatible with the Galactic value.  These are only three instances where the absorption column density of the pulsar is reported so far. Although the available data are scarce, the absorption column density does not show clear correlation with the source intensity or the orbital phase. In this sense, it is note worthy that the observations of the large absorption column with RXTE/PCA were made when the circumstellar disk around the Be companion (Be disk) was absent (Negueruela et al. 2000).  Haigh et al. (2004) suggest that the X-ray activity of the pulsar is correlated with the change of the truncation radius of the Be disk. The truncation of the disk by the tidal interaction with the neutron star, which defines the truncation radius, could explain the observed X-ray outbursts in Be/X-ray binary systems. When the truncation radius reduces, the material from the outer portions of the disk are expelled and may be found elsewhere in the binary system. This may trigger the giant X-ray outburst. When the Be disk disappears, as the case of RXTE/PCA observations, all the disk material is considered to be accreted and/or ejected throughout the binary system. The presence of the expelled material in the line of sight may have caused a moderately large absorption column. Haigh et al. (2004) argues that the reduction of the Be disk continues for several orbital periods before/after the giant outburst. Because the Suzaku observation was made about an orbital period after the giant outburst in May/June 2005, the reduced Be disk may have caused moderately large absorption column.  Broad-band spectroscopy of A0535+262 shows the presence of a weak iron K$_\\alpha$ emission line with $\\sim$25~eV equivalent width. Other than the 1989 $Ginga$ observation and recent RXTE and INTEGRAL observations of the same minor outburst, none of the previous observations of the pulsar could detect the weak emission line at $\\sim$6.4~keV\\@. The emission line is interpreted as due to the fluorescent line from the cold ambient matter in the vicinity of the neutron star. If the cold matter is distributed spherically symmetric around the neutron star, the excess column density over the Galactic value, $\\sim\\!5\\times10^{21}$ atoms cm$^{-2}$, would produce an iron emission line with an equivalent width of only a few eV (Makishima 1986). The larger equivalent width we detected may be interpreted as the cold matter having an asymmetric distribution, and more matter is distributed outside the line of sight. The fluorescence of the atmosphere of the binary companion could be a possible source of the iron emission line. However, we consider that it does not have a major contribution to the observed equivalent width, because the solid angle subtended by the companion, which is estimated as $\\Omega/4\\pi$ $\\leq$ 10$^{-4}$ at apastron, is very small due to the long-orbital period of A0535+262. Instead, we conjecture that the accretion disk is the probable reprocessing site to produce the iron fluorescent line. Because the pulse phase with higher line flux covers more than a half of the pulse period, the reprocessing site should subtend a large solid angle against the neutron star. The accretion disk conforms to this condition.  \\subsection{Pulse Phase Resolved Spectroscopy}  Pulse phase resolved spectroscopy of A0535+262 showed that the blackbody component, used to describe the soft excess, is pulsating as seen in some other accreting X-ray pulsars, such as LMC~X-4, SMC~X-1 etc. (Naik \\& Paul 2004a, Naik \\& Paul 2004b, Paul et al.\\ 2002 and references therein). The pulsation is found to be in phase with that of the middle-band X-ray emission (eg. 2.5--5.0 keV band; see Figure~3). However, in case of Her~X-1, the pulsating soft component is found to be shifted by about 230$^\\circ$ in phase from the power-law continuum component (Endo et al. 2000). A systematic and detailed study of the accreting X-ray pulsars which show soft excess revealed that the soft excess is a common feature in the X-ray pulsars (Hickox et al. 2004). Both the pulsating and non-pulsating natures of the soft excess seen in various X-ray pulsars suggest a different origin of emission of the soft components. The possible origins of the soft excess in accretion powered X-ray pulsars are (a) emission from accretion column, (b) emission by a collisionally energized cloud, (c) reprocessing by a diffuse cloud, and (d) reprocessing by optically thick material in the accretion disk (Hickox et al. 2004). The soft excess emission from the accretion column and by the reprocessing of harder X-rays by optically thick material in the accretion disk, are expected to show pulsations whereas in other cases, it is non-pulsating in nature. The pulsation of the soft component (blackbody component) in phase with the middle-band X-ray emission in A0535+262 suggests that the most probable region of soft X-ray emission is the accretion column and/or the inner accretion disk.  The absorption column seems to have significant contribution to produce the dip like structure at soft X-rays at phases 0.65-0.80 in the count rate profile (left top panel of Figure~\\ref{phrs}). Column density takes large values, $\\sim\\!1.2\\times10^{22}$ atoms cm$^{-2}$, at phase 0.7-0.8. This may contribute to produce the dip like structure in the profiles below 1 keV which disappears from the profiles at higher energies. On the other hand, the primary dip exists at all the energies. The observed primary dip in the pulse profile can be understood by the change in the NPEX model parameters over pulse phase. The NPEX continuum model is an approximation of the unsaturated thermal Comptonization in hot plasma (Makishima et al. 1999). At low energies, it reduces to the ordinary power-law with negative slope. The lower value of the power-law index ($\\alpha_1$) at the primary dip phase implies large optical depth to the Compton scattering. This means that many photons are scattered-out from the line of sight. It explains the presence of the primary dip in the soft and hard X-ray pulse profiles at same phase. In addition to the primary dip, the deepening of the dip at the phases 0.0-0.17 (Figure~\\ref{erpp}) can also be understood by the changes of the parameters of the Comptonizing plasma. The power-law index ($\\alpha_1$) tends to be larger and the exponential cut-off energy is smaller in this phase range. This means that the Comptonizing plasma has a relatively smaller optical depth and a lower temperature. This reduces the efficiency of Compton up-scattering and the number of hard photons, which causes the deepening of the dip at phases 0.0-0.17 above 5 keV."
    },
    "0710/0710.4545_arXiv.txt": {
        "abstract": "We present novel evidence for a fine structure observed in the net-circular polarization (NCP) of a sunspot penumbra based on spectropolarimetric measurements utilizing the Zeeman sensitive \\FeI\\ 630.2\\,nm line. For the first time we detect a filamentary organized fine structure of the NCP on spatial scales that are similar to the inhomogeneities found in the penumbral flow field. We also observe an additional property of the visible NCP, a zero-crossing of the NCP in the outer parts of the center-side penumbra, which has not been recognized before. In order to interprete the observations we solve the radiative transfer equations for polarized light in a model penumbra with embedded magnetic flux tubes. We demonstrate that the observed zero-crossing of the NCP can be explained by an increased magnetic field strength inside magnetic flux tubes in the outer penumbra combined with a decreased magnetic field strength in the background field. Our results strongly support the concept of the uncombed penumbra. ",
        "introduction": "\\label{sec:intro}  The asymmetry of the Stokes parameters with wavelength is a powerful diagnostic tool to probe the relation between flows and the magnetic field in the atmosphere of sunspots. Line asymmetries convey information about lateral (within one resolution element) and the line-of-sight (LOS) gradients and discontinuities of physical parameters. Since the first systematic characterization and interpretation of these asymmetries and their spatial distribution throughout a sunspot by \\citet[and references therein]{sanchezalmeida+lites1992} the complementary improvements achieved in both observational techniques and theoretical modeling, have advanced our understanding of particularly the penumbral inhomogeneities. It is general consent that a combination of non-trivial magnetic field and flow topologies along the LOS are the cause of the observed penumbral Stokes asymmetries. A simple and instructive way to characterize the asymmetry of Stokes V profiles is through the net-circular polarization (NCP), defined as $\\mathcal{N}=\\int V(\\lambda) d\\lambda$, where the integral is performed over one spectral line only. The sum of the NCP of individual lines than leads to the well-known broad-band circular polarization (BBCP) first noticed by \\cite{illing+landmann+mickey1974a, illing+landmann+mickey1974b}. The appearance of the NCP is very distinct in the penumbra and shows a peculiar difference when between observations in the visible or in the near infrared wavelength region: in the visible at \\FeI\\,630.2\\,nm the NCP maps tend to be symmetric about the disk center and sun center connection line \\citep{sanchezalmeida+lites1992, westendorp+etal2001b, mueller+etal2002} while infrared NCP maps at \\FeI\\,1564.8\\,nm reveal an antisymmetric distribution \\citep{schlichenmaier+collados2002, bellot+balthasar+collados2004}. \\cite{schlichenmaier+etal2002} and \\cite{mueller+etal2002} demonstrate that the uncombed penumbral atmosphere \\citep{solanki+montavon1993} mimicked by the \\textsl{moving tube model} \\citep{schlichenmaier+jahn+schmidt1998a} reproduces the azimuthal behavior of the NCP and, furthermore, convincingly show that the incisive discrepancy between the visible and the near infrared can be understood when the effects of anomalous dispersion \\citep{landolfi+landi1996} are included.  So far penumbral NCP maps of the \\FeI\\,630.2\\,nm line have been only presented before by \\cite{sanchezalmeida+lites1992} (in form of a contour line map), \\cite{westendorp+etal2001b}, \\cite{mueller+etal2002} and \\cite{mueller+etal2006}. These maps visualize very well the symmetric behaviour of the NCP but lack the spatial resolution to reveal the presence of any potential fine structure.  In this Letter we focus on high-spatial resolution spectropolarimetric measurements of a sunspot to investigate inhomogeneities in the penumbral magnetic field and flow field. Instead of utilizing inversion techniques we calculate the NCP and line-of-sight (LOS) velocities and elaborate on the properties of the same. In order to interpret the observational results we use the 3D geometric flux tube model VTUBE \\citep{mueller+etal2006} and perform subsequent radiative transfer calculations from which NCP and Doppler maps can be synthesized. ",
        "conclusions": "\\label{sec:conclusions}  In this paper, we have presented novel evidence for a spatial fine structure in the visible NCP of the \\FeI\\,630.25\\,nm line. We find for the first time that the NCP is structured in radial filaments very much alike maps of other line parameters, e.g. the intensity, equivalent width and line width \\citep[see e.g.][]{johannesson1993, rimmele1995a, rimmele1995b, tritschler+etal2004}. The spatial scales of the NCP filaments are similar to those found in the observed LOS velocity map although we cannot find a good overall correlation between the signals of the NCP and the LOS velocity. We conjecture that higher spatial resolution is needed to settle this issue. \\cite{martinezpillet1997} find a correlation between the azimuthal variation of the NCP and the magnetic field inclination on angular scales of $10^{\\circ}$. Although we cannot provide information about the inclination at this point we find that our observations show azimuthal variations of the NCP that are only of the order of $4^{\\circ}$.  Furthermore, we have observed a zero-crossing of the NCP on the outer center-side penumbra and have demonstrated that this signature can be explained by an increased magnetic field strength inside magnetic flux tubes in the outer penumbra, while the magnetic field strength is decreased in the background field ($\\Delta B$-effect). To this end, we have compared high-resolution observations with maps of the NCP and the LOS velocity that have been synthesized from a 3D geometric sunspot model that has been coupled with a 1D radiative transfer code \\citep[see][ for details]{mueller+etal2006}. We point out that the space of possible atmospheric models that can  reproduce the observations is significantly constrained by  taking  into account the information contained in Doppler maps in addition to  the NCP maps. This paper extends the work by \\cite{solanki+montavon1993}, \\cite{martinezpillet2000}, \\cite{mueller+etal2002}, and \\cite{mueller+etal2006} and clearly demonstrates the need to model the \"uncombed\"  penumbra of sunspots in greater detail. This would include also taking into account the effect of field line curvature around flux tubes as considered by \\cite{borrero+bellot+mueller2007}."
    },
    "0710/0710.1010_arXiv.txt": {
        "abstract": "The evolution of single stars at low metallicity has attracted a large interest, while the effect of metallicity on binary evolution remains still relatively unexplored. We study the effect of metallicity on the number of binary systems that undergo different cases of mass transfer.  We find that binaries at low metallicity are more likely to start transferring mass after the onset of central helium burning, often referred to as case C mass transfer. In other words, the donor star in a metal poor binary is more likely to have formed a massive CO core before the onset of mass transfer. At solar metallicity the range of initial binary separations that result in case C evolution is very small for massive stars, because they do not expand much after the ignition of helium and because mass loss from the system by stellar winds causes the orbit to widen, preventing the primary star to fill its Roche lobe.  This effect is likely to have important consequences for the metallicity dependence of the formation rate of various objects through binary evolution channels, such as long GRBs, double neutron stars and double white dwarfs. ",
        "introduction": "During their life stars can expand to a radius which is up to 1000 times bigger than their initial radius. In close binary systems, they start to transfer mass if their radius exceeds a critical radius, the Roche lobe radius. It is usually the initially most massive star, the primary star, which ``fills its Roche lobe'' first.  At solar metallicity binary evolution has been studied by various groups, while their evolution at low metallicity is relatively unexplored. Heavy elements, such as carbon, oxygen, nitrogen and iron, are important contributors to the opacity of stellar material in metal rich stars. Due to the lower opacity, metal poor stars are generally hotter and more compact. As it is the evolution of the radius of a star in a binary which determines if and when mass transfer occurs, we expect that the evolution of binaries in metal poor environments is significantly different from binaries in the solar neighborhood.  In this work we study the radius evolution of stars with different masses in order to infer the frequency of various cases of binary evolution as a function of metallicity. In a second contribution to this proceedings \\citep{acc07} we discuss the effect of metallicity on the expansion of accreting main sequence stars and the potential consequences for binary evolution.    \\begin{figure} \\includegraphics[bb=190 0 900 500, clip, width=\\columnwidth]{deminkea_poster2_fig1.eps} \\caption{ Evolution track of a 6 \\Msun star at solar (Z = 0.02) and low metallicity (Z =0.00001). \\label{hrd} } \\end{figure}    \\begin{figure*} \\includegraphics[ bb=50 10 680 480, width=\\columnwidth]{deminkea_poster2_fig2.ps} \\includegraphics[  bb=50 10 680 480, width=\\columnwidth]{deminkea_poster2_fig3.ps} \\caption{ The initial orbital periods that lead to case A, B and C mass transfer are given as function of the primary mass assuming an initial mass ratio of 0.75 \\citep[inspired by][]{Webbink79}. Widening of the orbit due to angular momentum loss in the form of stellar winds is taken into account. The hatching indicates approximately in which cases the donor has a convective envelope. \\label{cases} } \\end{figure*} ",
        "conclusions": "We find that at low metallicity the donor star in a binary is more likely to have developed a massive CO core before the onset of mass transfer. For binaries with primary masses above 30\\Msun Case C mass transfer may occur almost exclusively at low metallicity.  This effect potentially has important consequences for the metallicity dependence of the formation rate of various objects through binary evolution channels such as long Gamma Ray Burst progenitors \\citep[see discussion section of][]{wolf+podsiadlowski07} but also double neutron stars and double white dwarfs. More work on binary evolution at low metallicity is needed.  The precise fraction of Case C binaries as function of metallicity depends on assumptions about the distribution of initial binary parameters, the description of mixing and overshooting and the (metallicity dependence) of the mass loss rate. We find that the maximum radius can vary by more than a factor of 10 when we compare calculations of the STARS code to calculations done with the code described in \\cite{Yoon+ea06} and when we vary the amount of overshooting. In a forth coming paper we will present the range case A, B and C mass transfer for a range of metallicities for different assumptions for the uncertain parameters \\citep{CASEABC08}.       \\begin{theacknowledgments} We would like to thank Arend-Jan Poelarends, Matteo Cantiello, and Philip Podsiadlowski for interesting discussions and useful comments and suggestions. \\end{theacknowledgments}"
    },
    "0710/0710.3223_arXiv.txt": {
        "abstract": "Recent advances have made it possible to obtain two-dimensional line-of-sight magnetic field maps of the solar corona from spectropolarimetric observations of the \\ion{Fe}{13} 1075 nm forbidden coronal emission line. Together with the linear polarization measurements that map the azimuthal direction of the coronal magnetic field projected in the plane of the sky containing Sun center, these coronal vector magnetograms allow for direct and quantitative observational testing of theoretical coronal magnetic field models. This paper presents a study testing the validity of potential-field coronal magnetic field models. We constructed a theoretical coronal magnetic field model of active region AR 10582 observed by the SOLARC coronagraph in 2004 by using a global potential field extrapolation of the synoptic map of Carrington Rotation 2014. Synthesized linear and circular polarization maps from thin layers of the coronal magnetic field model above the active region along the line of sight are compared with the observed maps. We found that the observed linear and circular polarization signals are consistent with the synthesized ones from layers located just above the sunspot of AR 10582 near the plane of the sky containing the Sun center. ",
        "introduction": "\\label{sec:intro}  Understanding the static and dynamic properties of the solar corona is one of the great challenges of modern solar physics. Magnetic fields are believed to play a dominant role in shaping the solar corona. Current theories also attribute reorganization of the coronal magnetic field and the release of magnetic energy in the process as the primary mechanism that drives energetic solar events. However, direct measurement of the coronal magnetic field is a very difficult observational problem. Early experiments have demonstrated the feasibility of the measurement of the orientation of the coronal magnetic fields by observation of the linear polarization of forbidden coronal emission lines (CELs) in the visible and at IR wavelengths \\citep{eddy1967,mickey1973,arnaud1982, querfeld1982,tomczyk_et_al_2007}. Radio observations have also been successful in measuring the strength of the coronal magnetic field near the base of the solar corona \\citep[e.g.,] [and references therein] {brosius_2006}. Direct measurement of the coronal magnetic field strength at a higher height by IR spectropolarimetry of the CELs was achieved only recently \\citep{lin_et_al_2000, lkc_2004}. Without direct measurements, past studies involving coronal magnetic fields have relied on indirect modeling techniques to infer the coronal magnetic field configurations, including coronal intensity images observed in the EUV and X-ray wavelength ranges and numerical methods that reconstruct the three-dimensional coronal magnetic field structure by extrapolation and MHD (magnetohydrodynamics) simulation based on photospheric magnetic field measurements. Since experimental verification of theories and models is one of the cornerstones of modern science, the lack of observational verification of these indirect magnetic field inference methods that are in widespread use is a very unsatisfactory deficiency in our field.  Our 2004 observations were obtained above active region AR 10582 right before its west limb transit. We have obtained the first measurement of the height dependence of the strength of the line-of-sight (LOS) component of the coronal magnetic field in this data, and it showed an intriguing reversal in the direction of the LOS magnetic field at a height of approximately $0.15\\ R_{\\odot}$ above the solar limb, as shown in Figures 4 and 5 of \\cite{lkc_2004}. This feature and the observed linear polarization map have presented us with our first opportunity to carry out a comprehensive observational test of our coronal magnetic modeling methods in which the strength and direction of the  magnetic fields predicted by the models can be directly checked by the observations.  Force-free extrapolation of photospheric magnetic fields is currently the primary tool for the modeling of coronal magnetic fields. However, it is not without limitations or uncertainties. For example, the force-free assumption does not hold true in the photosphere and low chromosphere, and possibly in the high corona above 2 $R_{\\odot}$ \\citep{gary_2001}. Moreover, a different assumption (current-free, linear and nonlinear force free) about the state of the electric current in the corona can lead to substantially different extrapolation results. Without direct magnetic field measurements, many fundamental questions concerning the basic assumptions and the validity of our tools cannot be addressed directly. To date, questions like ``Is a potential field approximation generally an acceptable approximation for coronal magnetic fields?\" or ``Do linear or nonlinear force-free extrapolations provide a more accurate description of the coronal magnetic field?\" can only be addressed by visual comparison between the morphology of selected field lines of the extrapolated magnetic field model and observational tracers of coronal magnetic fields such as the loops seen in EUV images. However, these visual tests are qualitative and subjective. Furthermore, they assume the coalignment between the magnetic field lines and the loops in EUV images, which has not been verified observationally.  As our first test, we attempted to address the question ``Is the potential field extrapolation generally an acceptable approximation for the coronal magnetic field?\" This test was conducted by comparing the observed polarization maps of AR 10582 to those derived from a coronal magnetic field model constructed from the potential field extrapolation of photospheric magnetic field data. However, before we present our study, we should point out that because of the nature of our modeling tools and observational data, the results and conclusions of this research are subject to certain limitations and uncertainties. One of the intrinsic limitations of the observational data used in this study is that because of our single sight line from Earth to the Sun, the photospheric and coronal magnetic field observations cannot be obtained simultaneously. In the case of global coronal magnetic field models, the whole-Sun photospheric magnetic field data used as the boundary condition of the extrapolation can be obtained only over an extended period of time. Although the large-scale magnetic structure of nonflaring active regions may appear stable over a long period of time, high-resolution EUV observations have shown that the small-scale coronal structures are constantly changing. Thus, studies such as ours that compare coronal magnetic field observations and models constructed from photospheric magnetic fields inevitably are subject to uncertainties due to the evolution of the small-scale structures in the regions, and we should not expect a precise match between the observed and synthesized polarization maps.  Another deficiency due to the single sight line of our observations is the lack of knowledge of the source regions of the coronal radiation. This is perhaps the most limiting deficiency of the observations and models of this research. The uncertainty of the location of the source regions associated with coronal intensity observations due to the low optical density of the coronal plasma and the resulting long integration path length is familiar. In the case of the coronal magnetic observations, the LOS integration problem prevents us from performing an inversion of the polarization data to reconstruct the three-dimensional magnetic field structure of the corona for direct comparison with those derived from extrapolations or MHD simulations.  On the other hand, since extrapolation techniques do not include the thermodynamic properties of the plasma in the construction of the coronal magnetic field models, they do not include information about the location of the CEL source regions either. Therefore, they cannot predict the intensity and polarization distribution of the CEL projected on the plane of the sky that are needed for direct comparison with the polarimetric observations.  Without the information about the location of the source regions from the observational data and the extrapolated models, we adopted a trial-and-error approach in which synthesized linear and circular polarization maps were derived using empirical source functions and the extrapolated potential magnetic field model, and were compared directly with those obtained from observations. Obviously, if acceptable agreement can be achieved with any of the models tested, then we can argue with a certain degree of confidence that these models are plausible models of the observed corona and that the potential field approximation is a reasonable approximation of the coronal magnetic fields. Nevertheless, we should emphasize that this trial-and-error approach is not an exhaustive search of all the possible source functions and therefore cannot provide a clear-cut true-or-false answer. In other words, even if no acceptable agreement can be found with all the model source functions we have considered, the potential field approximation still cannot be dismissed completely. ",
        "conclusions": "This research examines observationally the validity of current-free, force-free potential-field approximation for coronal magnetic fields. We conducted a study comparing observed and synthesized spatial variations of the linear and circular polarization maps in the corona above active region AR 10582, which after a week of extensive flaring activities should have settled into a minimum energy configuration that could be adequately modeled by a potential field model. The coronal magnetic field model used for this study was constructed from a global potential field extrapolation of the synoptic photospheric magnetogram of Carrington cycle 2014 obtained by the \\SOHO/MDI instrument. Because the most important source of error of this type of study is the uncertainty of the location of the source function of the coronal radiation, we first tested three analytical, but empirically determined, source functions.  These simple source functions are based on a gravitationally stratified atmospheric density model with a uniform temperature in the entire modeled volume, supplemented by magnetic weighting functions based on the observational impression that, at least at the length scale of the typical active region size, CEL radiation seems to be correlated with the strength of the photospheric magnetic fields. We found that none of these empirical source functions can adequately reproduce the observed linear polarization maps, although it seems that the source functions that include both density and magnetic fields produced slightly better results. Based again on the observational impression that the coronal intensity structures have a spatial scale much smaller than that of the active regions, we then compared the observed polarization maps with those constructed from thin (56 Mm FWHM) layers along the LOS. In this analysis, we found that polarization maps originating from layers located near the sunspot of the region are in reasonable agreement with the observed ones. However, the best fit for linear and circular polarization did not occur at the same layer. They are separated by a distance of about 50 Mm.  Does the small discrepancy between the best-fit locations of the linear and circular polarization weaken the support for the potential field extrapolation? As we have discussed in \\S\\ref{sec:intro}, many uncertainties conspire to limit the precision of this study. For example, the difference may be due to the evolution of the small-scale photospheric magnetic field of the active region, and we do not have any observational information that we can use to test this possibility.  The assumption of uniform temperature distribution in our source function is certainly not a physically realistic assumption. Therefore, it is not possible to assess the significance of the small difference in the location of the source regions of the linear and circular polarization.  Furthermore, in addition to density and temperature, the source functions of CEL linear and circular polarization depend on different components of the coronal magnetic field. So it is in fact physically reasonable that we would find the linear and circular polarization signals originate from slightly different locations. Finally, because the three parameters ($\\sigma_p$, $\\sigma_\\chi$, and $H_0$) we used to evaluate the quality of the fit were obtained independently, the statistical significance that all three parameters reached minimum at approximately the same location near the strongest photospheric magnetic feature of the active region cannot be dismissed as pure coincidence. These considerations lead us to conclude that potential field extrapolation can be used to provide a zero-order approximation of the real solar corona if the active region is in a relatively simple and stable configuration. Additionally, this study suggests that, at least for isolated active regions, CEL radiation may originate from a region close to the strongest photospheric magnetic feature in the active region with a small spatial scale comparable to the characteristic size of the coronal loops seen in the intensity images. If this is confirmed, then a single-source inversion to infer the magnetic field directly from the polarimetric observation such as that proposed by \\cite{judge_2007} may be justified.  Our conclusion about the viability of potential field extrapolation as a coronal magnetic field modeling tool for stable active regions is supported by a study by \\cite{riley_et_al_2006}, in which the coronal magnetic field configuration derived from a potential field model was found to closely match that derived from a MHD simulation in the case of untwisted fields. Nevertheless, we should emphasize that our conclusion is derived from a single observation of a simple and stable active region. Clearly, more observations and model comparison are needed for a more comprehensive test of this result.  Can the potential field approximation be used to model more complicated active regions? Using radio observations, \\cite{lee_et_al_1999} found that a force-free-field model yields better agreement between the temperatures of two isogauss surfaces connected by the modeled field lines of an active region with strong magnetic shear. This study thus provides observational evidence against the use of potential field approximations for the modeling of complex active regions. Therefore, linear and nonlinear force-free extrapolations should be employed in future testing of theoretical coronal magnetic field models using the IR spectropolarimetric observations to study if these models can offer a better description of the observed coronal fields.  Can we distinguish the potential coronal magnetic field configurations from the non-potential ones with the spectropolarimetric observations of the coronal emission lines? In a numerical study, \\cite{judge_2007} has demonstrated the sensitivity of LOS-integrated coronal polarization measurements to the electric current in the corona using theoretical coronal magnetic field models as input.  Therefore, we should expect to find better agreement between observed and synthesized polarization maps for more complex active regions with linear or nonlinear force-free magnetic field models. Work to model AR10582 using the force-free extrapolation method is already underway, and we should be able to address this question in the near future. Since all extrapolation methods are subject to the ambiguities problem of the source regions, we will also employ MHD simulations that include both the magnetic and thermodynamic properties of the corona in the calculation to help resolve this problem. These are research activities that we will be pursuing in the near future as the solar cycle evolves toward the next solar maximum and more coronal magnetic field data become available.  The greatest difficulty of this study is the uncertainty of the location of the source function due to the long integration path along the LOS. However, this is not a difficulty affecting only the interpretation of coronal magnetic field measurements. It affects the intensity observation as well, and is the primary reason that years into the operation of \\SOHO/EIT and \\TRACE, we still cannot deduce 3-D intensity and temperature structure of the corona using data from these instruments. Fortunately, this deficiency in our observing capability may finally be removed with the recent launch of the \\STEREO\\ mission (Solar TErrestrial RElations Observatory).  For the resolution of the LOS integration problem in polarimetric observations, \\cite{kramar_et_al_2006} have demonstrated the promising potential of vector tomography techniques. While stereoscopic coronal magnetic field observations will not be realized any time soon, this method can be applied to observations obtained over periods of several days during the limb transit of active regions, provided that the active regions are in a stable condition. This is perhaps the best observational tool available for the resolution of the LOS integration problem in the near future."
    },
    "0710/0710.5225_arXiv.txt": {
        "abstract": "{Maser emission from the  \\hzo\\ molecule probes the warm,   inner circumstellar envelopes of oxygen-rich red giant and supergiant stars. Multi-maser transition studies can be used to put constraints on the density and    temperature of the emission regions.} {A number of known \\hzo\\ maser lines were observed toward the long period variables R Leo and W Hya and the red supergiant VY CMa. A search for a new, not yet detected line near 475 GHz was conducted toward these stars.} {The Atacama Pathfinder Experiment telescope was used for a multi-transition observational study of submillimeter \\hzo\\ lines.} {The $5_{33} - 4_{40}$ transition near 475 GHz  was clearly detected toward VY CMa and W Hya. Many other \\hzo\\ lines were detected toward all three target stars. Relative line intensity ratios and velocity widths were found to vary significantly from star to star. } {Maser action is observed in all but one line for which it was theoretically predicted. In contrast, one of the strongest maser lines, in R Leo by far \\textit{the} strongest, the 437 GHz $7_{53}-6_{60}$  transition, is not predicted to be inverted. Some other qualitative predictions of the model calculations are at variance with our observations. Plausible reasons for this are discussed. Based on our findings for W Hya and VY CMa, we find evidence that the \\hzo\\ masers in the AGB star W Hya arise from the regular circumstellar outflow, while shock excitation in a high velocity flow seems to be required to excite masers far from the red supergiant VY CMa.} ",
        "introduction": " ",
        "conclusions": "Discussion} \\subsection{\\label{excitation}Modeling the excitation of water maser lines} Earlier excitation studies aimed at explaining stellar (and interstellar) \\hzo\\ maser pumping only considered a limited number of energy levels, both, due to limitations of computational capabilities and because collision rates were not known for excitation to and from higher energy states \\citep[see, e.g., ][]{deJong1973, Deguchi1977, CookeElitzur1985}. Nevertheless, all these studies predicted strong inversion of the 22.2 GHz line, which was the only \\hzo\\ maser line known at the time. These studies also predicted maser action in a number of other transitions that were later indeed found to be masing, despite the fact that in some cases the collisional excitation rate coefficients used were vastly different from the ``modern'' values computed by \\citet{Palma_etal1988}. The more recent,  comprehensive, study by NM91 considered all 349 states of ortho and para water up to an energy of 7700 K above the ground state and extrapolated the Palma et al. rates to high enough temperatures.  \\citet{NeufeldMelnick1990} and NM91 performed statistical equilibrium calculations modeling vibrational ground state \\hzo\\ excitation and radiative transfer for plane parallel media under the large velocity gradient (LVG) assumption. \\citet{NeufeldMelnick1990} specifically explore conditions under which, both, the the 22.2 GHz and the 321 GHz lines are masing for the case of an O-rich evolved star with a mass-loss rate ($3~10^{-5}$ \\Mspy) intermediate between the values of our objects. They find that the pumping for these lines is dominated by collisions and that radiation plays a minor role. The same was also found for other transitions in circumstellar envelopes \\citep[][ NM91]{Deguchi1977, CookeElitzur1985}.  NM91 explored larger regions of (essentially) density and temperature parameter space in which  the energy levels of certain transitions become inverted. They calculate  maser emissivities and optical depths as functions of temperature and a variable, $\\zeta'$, which is equal to a geometric factor, $G$, divided by the velocity gradient, $d{\\rm v}/dr$, times the product of the hydrogen nuclei, $n$, and the water density, $n({\\rm H}_2{\\rm O})$, i.e., $n\\times n({\\rm H}_2{\\rm O})$ $ = n^2 x({\\rm H}_2{\\rm O})$, where $x({\\rm H}_2{\\rm O})$ is the \\hzo\\ abundance. $\\zeta'$ is, thus, given by:   $$\\zeta' = G{n^2 x({{\\rm H}_2{\\rm O})}\\over{d{\\rm v}/dr}}$$ Assuming a plausible, constant, value for the \\hzo\\ abundance and the velocity gradient of $10^{-4}$ and $10^{-8}$~cm~s$^{-1}$~cm$^{-1}$, respectively, makes $\\zeta'$ dependent on the H density squared alone. Measuring the H density in units of $10^9$ \\ccm\\ in the parametrization  of NM91, values of $log_{10}\\zeta' = -1, 0,$ and 1 correspond to H densities of 0.32, 1, and $3.2\\times10^9$ \\ccm.  \\subsection{\\label{comparison}Comparison with the APEX observations} NM91 predicted maser emission in all of the lines listed in Table \\ref{lines}, except for the 437 GHz $7_{53}-6_{60}$ and 443 GHz  $7_{52} - 6_{61}$ transitions in, roughly, this range of values quoted above. Maser emission in all their predicted maser lines has indeed been observed in star-forming regions and/or red giant stars, except for the 355 GHz $17_{4,13} - 16_{7,10}$ transition. They present opacities and emissivities (which are proportional to isotropic photon luminosities) for assorted lines calculated for two values of the kinetic temperature, 400 K and 1000 K. In particular, they find the very widespread 22.2 GHz \\kbw\\ transition to be by far the most luminous and the most robust maser line in the sense that it is inverted over the widest range of physical conditions in particular at extreme values of the density. A general trend is that for 1000 K the maximum emissivities in all the lines are reached for a factor of a few higher densities than at 400 K.  In the calculations of NM91, the peak emissivities of all lines but the 22.2 GHz line are similar within a factor of a few (although they occur at different values of $\\zeta'$). At, both, 400 and 1000 K, the peak emissivity of the 22.2 GHz line is much higher than that of any of the others, by a factor of $\\sim10$ at 400 K and $\\sim5$ at 1000 K. The peak emissivity is also reached at an order of magnitude higher value of $\\zeta'$ than for other lines. As can be seen from Table \\ref{lines} and Fig. \\ref{lineratios} in none of our stars does the 22.2 GHz line show the behavior predicted by these calculations, i.e. that it should be much more luminous than all the other lines.  While it is among the strongest lines in VY CMa, it is anomalously weak in R Leo. We note that R Leo showed an even more extreme 321/22.2 GHz line ratio in 1989/1990 \\citep{MentenMelnick1991}. In this star, the 22 GHz line is known to undergo extreme and erratic variations changing from not or barely detectable to hundreds of Jy within dozens of days \\citep{Rudnitskij1987}.  \\subsection{The ''new'' 475 GHz transition} For the newly discovered 475 GHz $5_{33} - 4_{40}$ transition NM91 predict, as for the 439 and 471 GHz lines, maser action only for their 1000 K calculation. The presence of maser action in this line is, thus, an important indicator of a high temperature environment.  \\subsection{\\label{predictions}Non- and mispredictions} NM91\\textit{do not} predict maser emission for the 437 GHz $7_{53}-6_{60}$ transition, which we find to be by far the strongest line in R Leo and the second strongest in VY CMa. This line has only been found toward evolved stars and not toward SFRs \\citep{Melnick_etal1993}. Pumping by the strong IR field in these objects, possibly via line overlaps, not considered by NM91, may be responsible for its excitation. NM91 \\textit{do}, however, predict maser action  in the extremely high excitation (5764.3 K) $17_{4,13} - 16_{7,10}$  transition near 355 GHz high with peak emissivities at $\\zeta'$ values 1--2 order of magnitude higher than values for which other lines show their peak emissivities. For constant water abundance and velocity gradient this translates into H$_2$ densities of order $10^{10-11}$ cm$^{-3}$, which is 1--2 orders of magnitudes  higher than the values over which most of the other lines show maser action. Our non-detection of this line (see Table \\ref{lineresults} and \\S\\ref{comparison}) seems to indicates that under such extreme conditions the necessary requirements for producing detectable maser emission may not be met.   \\subsection{A note of caution} Guided by the described model calculations, observations of combinations of maser lines in principle should put constraints on the physical parameters of the masing regions. One naive hope might be: The more lines one finds masing the tighter the physical parameters of the masing region can be constrained. One might even hope that the observed relative luminosities of the lines might provide yet tighter constraints.  Unfortunately at present, in the absence of more realistic models of the maser regions such expectations are most likely overly optimistic for several reasons: Comparing the emissivities or their ratios  predicted by the generic calculations of NM91 for a \\textit{plane-parallel(!)} medium with measured data probably does not make much sense. Moreover, an inherent assumption in the values calculated by NM91 is that all the lines are saturated throughout the region in which they are inverted. There is good observational evidence that this is not true for the 22.2 GHz line around W Hya \\citep{ReidMenten1990} although this issue is unclear (and difficult to address) for other transitions and other stars.  Another caveat when using line ratios as diagnostics is time variability. While all our APEX and Effelsberg data were taken within 1--2 weeks, at least the 22.2 GHz line in some sources (including R Leo) is known to show week-to-week variability (see \\S\\ref{comparison}). Information on variability on such small time scales is not available for the submillimeter lines, but presumably all of them are variable. This has been directly shown from a comparison of line profiles taken at different epochs 1--2 months apart for the 321 and 325 (and 22.2) GHz lines \\citep{MentenMelnick1991,YatesCohen1996}  Furthermore, to use multi-line data as diagnostics, an inherent assumption is that all lines arise from the same region (whose physical parameters are to be constrained). However, for VY CMa, the different shapes of the lines and the different velocity ranges they cover suggest that this is not true and that at least a significant portion of the lines' emission arises from disjoint regions. Clearly, high resolution interferometric observations are necessary to get a clearer picture of this. Given the high expected brightness temperatures, such observations with the Atacama Large Millimeter Array (ALMA) will afford studies of R Leo and W Hya with a resolution better than a stellar radius and a few stellar radii for VY CMa \\citep[see \\S4.4 of ][]{Menten2000}. All the lines discussed here have frequencies within  ALMA's 275--370 and 385--500 GHz ``First Light'' Receiver bands \\citep[``Bands 7 and 8'', ][]{Wilson_etal2005}. Observations of the 321 and 325 GHz lines can already be made today with the Submillimeter Array \\citep{Ho_etal2004}.  It is clear that current models for \\hzo\\ maser excitation in circumstellar envelopes fail to explain the observational picture. Velocity-resolved observations of a large number of maser and non-maser lines from a wide range of energies above the ground state with the Heterodyne Instrument for the Far Infrared (HIFI) aboard the Herschel satellite to be launched in the near future will soon deliver very tight constraints on \\hzo\\ excitation. Together with high spatial resolution data for high excitation (mostly maser) lines (see above), which will \\textit{only} be attainable with ground-based interferometers  and more sophisticated radiative transfer modeling these observations will lead to an understanding of water, which dominates the thermal balance in these regions.   \\subsection{\\label{differentnature}Different natures of AGB star and red supergiant water masers?} Let us come back to the fact, mentioned in \\S\\ref{spectralappearance}, that toward W Hya, the 22.2 GHz \\hzo\\ maser emission arises from a ring with a diameter of 24 AU centered on the star. In contrast, toward VY CMa, the maser emission arises from a $770\\times440$ AU region and proper motion observations appear to indicate that the masers partake in a bipolar outflow \\citep{RichardsCohen1998}, although we notice that this flow would be directed perpendicular to the larger-scale CO outflow imaged by \\citet{Muller_etal2007}. Are these distributions consistent with the masers arising from a ''normal'', i.e. steadily expanding circumstellar outflow?  In the following, we investigate whether over the whole of the maser distributions in both objects the temperatures (calculated from the luminosities of the stars) and the densities (derived from the mass-loss rates) are conducive for maser excitation.  For W Hya \\citet{Justtanont_etal2005} modeled CO data to derive a temperature, $T$, profile, i.e., $T$ vs. distance from the star, $r$. For the maser region ($r = 1.8~10^{14}$ cm), one finds $T \\approx 1000$ K from their Fig. 3. Also, assuming, like these authors, a luminosity of 5400 \\Lsun\\ and an effective temperature of 2500 K, respectively, one calculates a comparable 970 K at this $r$ from the Stefan-Boltzmann law.  The H$_2$ density  in a spherically symmetric outflow at distance $r$ from the central star is given by $$n({\\rm H}_2) = 1.13~10^{43}\\Mdot(\\Mspy)r^{-2}({\\rm cm}){\\rm v}_{\\rm exp}^{-1}(\\kms)$$ where \\Mdot\\ and ${\\rm v}_{\\rm exp}$ are the mass-loss rate and the expansion velocity, respectively, \\citep[see, e.g. ][]{Millar1988,MentenAlcolea1995}. If we assume for the masing region $n({\\rm H}_2) = 10^9$ \\ccm\\ (see \\S\\ref{comparison}), $r = 1.6~10^{14}$ cm and an expansion velocity of 2 \\kms\\ and invert this equation to calculate the mass-loss rate, we obtain $\\Mdot = 6~10^{-6}$ \\Mspy, which is half the rate derived by \\citet{ZubkoElitzur2000} and in between the estimates of \\citet{Neufeld_etal1996} and \\cite{Justtanont_etal2005} and, thus, hopefully, a plausible value.  \\footnote{\\citet{Neufeld_etal1996} derive a large value of $0.3$--$2~10^{-5}$ \\Mspy\\ for \\Mdot\\ from modeling emission of thermally excited \\hzo\\ lines observed by the Infrared Space Observatory (ISO). In contrast, \\citet{Justtanont_etal2005}, also using data from ISO as well as the Odin satellite, derive a much smaller mass loss rate of  $(2.5\\pm0.5)~10^{-7}$ \\Mspy, comparable to values derived from CO lines. Their low \\Mdot\\ comes ``at the expense'' of an extremely high \\hzo\\ abundance of [\\hzo/H$_2$ $=2~10^{3}$, which is significantly higher than the cosmic abundance of O. To explain it they invoke an influx of water from evaporating icy bodies in an Oort cloud analog, a mechanism that had earlier been advocated to explain the presence of water in IRC10216's envelope \\citep{Melnick_etal2001}. Although the calculations of Justtanont et al. make a strong case in favor of a low \\Mdot\\ one wonders whether a higher \\Mdot\\ and a lower \\hzo\\ abundance  in the range of the Neufeld et al. values might also be consistent with the measurements.}  Doing the same exercise for the RSG VY CMa with a luminosity of  $\\approx 2~10^5$ \\Lsun, we calculate for a region with an average radius of 300 AU (see \\S\\ref{spectralappearance}) a temperature of 480 K, which would be adequate at least for 22.2 GHz maser production. However, if we calculate the mass-loss rate required for a spherically symmetric outflow to maintain a density of $10^9$ \\ccm\\ at 300 AU, we obtain a value of 0.045 \\Mspy. This is $\\approx 200$ times higher than the generally quoted (already extremely high) mass-loss rate for this object (see \\S\\ref{intro}). We conclude that, in contrast to W Hya, at least the 22.2 GHz \\hzo\\ masers far from VY CMa cannot arise from the general material in a spherically symmetric outflow, but rather from special regions (e.g., shock fronts) within such a flow or from a collimated flow. This situation is reminiscent of \\hzo\\ masers in star-forming regions modeled by \\citet{Elitzur_etal1989}. These authors invoke a temperature (400 K) similar to the value we derive for VY CMa.  At this temperature, the highest excitation lines' maser emissivities are small (NM91), restricting their emission to regions closer to the star. This is in fact consistent with the smaller velocity ranges covered by the 321 and the 437 GHz lines."
    },
    "0710/0710.2633_arXiv.txt": {
        "abstract": "The near Main Sequence B stars show a sharp drop-off in their X-ray to bolometric luminosity ratio in going from B1 to later spectral types. Here we focus attention on the subset of these stars which are also Oe/Be stars, to test the concept that the disks of these stars form by magnetic channeling of wind material toward the equator. Calculations are made of the X-rays expected from the Magnetically Torqued Disk (MTD) model for Be stars discussed by Cassinelli \\etal (2002), by Maheswaran (2003), and by Brown \\etal (2004). In this model, the wind outflow from Be stars is channeled and torqued by a magnetic field such that the flows from the upper and lower hemispheres of the star collide as they approach the equatorial zone. X-rays are produced by the material that enters the shocks above and below the disk region and radiatively cools and compresses in moving toward the MTD central plane. It differs from the Babel \\& Montemerle (1997) model in having a weaker B field and a large centrifugal effect. The dominant parameters in the model are the $\\beta$ value of the velocity law, the rotation rate of the star, $S_o$, and the ratio of the magnetic field energy density to the disk gravitational energy density, $\\gamma$.  The model predictions are compared with the $ROSAT$ observations obtained for an O9.5 star $\\zeta$ Oph from \\Berghofer\\ \\etal (1996) and for 7 Be stars from Cohen \\etal (1997). Two types of fitting models were used to compare predictions with observations of X-ray luminosities versus spectral types. In the first model we choose an estimate of the spin rate parameter $S_o$ from observed $\\vsini$ values, and choose $\\gamma$ using the threshold magnetic field value derived in Cassinelli \\etal (2002), then the $\\beta$ value is adjusted to fit the observations. For all but one case, the $\\beta$ value was found somewhat larger than unity, a typical value derived for radially streaming stellar winds. This value, appropriate to a slowly accelerating wind, might be an indication that the magnetic field modifies the dynamics of the outflow from the star. In the second fitting model, we choose $\\beta$ to be unity, $\\gamma$ as in the first model, and adjust $S_o$ to fit the observations. For these comparisons with the X-ray observations we find that $S_o$ is in the range 0.49 to 0.88, which agrees with traditional estimates of the rotation rate of Be stars, but is below the breakup values that are required in recent non-magnetic models for Be star disks (Townsend \\etal 2004).  Extra considerations are also given here to the well studied Oe star $\\zeta$ Oph for which we have $Chandra$ observations of the X-ray line profiles of the triad of He-like lines from the ion Mg XI. We find that a reasonably good fit is made to the observed Mg XI line profiles. The lines are predicted to form primarily at radial distance of about two stellar radii in the disk, and the ratio of the forbidden to intercombination lines (i.e., $f/i$) that is a diagnostic of source distances, agrees with this prediction. In addition, the lines are broad, with HWHM of about 400 $\\kmsec$. Again this is in compatibility with the model predictions for the disk rotation of this star.  Thus the X-ray properties add to the list of observables which can be explained within the context of the MTD concept. This list already includes the \\Halpha\\ equivalent widths and white light polarization of Be stars. We do not include here the disk density correction due to gravity, neglected in Cassinelli \\etal (2002) but included in Maheswaran (2003). This will likely increase the \\Halpha\\ and polarization predictions but have little or no effect on the X-rays which are generated in the upstream region. Thus the X-rays are not sensitive to the cooled disk region where Keplerization of the disk material might be occurring. Nonetheless the process by which matter and angular momentum are added to the disk are equally important and this study indicates that X-ray properties are consistent with the overall MTD concept. ",
        "introduction": "In a survey of the X-ray emission of near Main-Sequence B stars (or B V stars), Cassinelli \\etal (1994) and Cohen \\etal (1997) found a departure from the canonical ``law\" relating X-ray luminosity to the bolometric luminosity for hot star X-rays: $L_x/L_B=10^{-7}$. This relation holds for stars throughout the O spectral range and extends to about B1 V. However, beyond that there is a sharp drop in the ratio values in going to B3 V by about 2 orders of magnitude. Cohen \\etal (1997) investigated if the sharp decrease could be explained merely by the reduction of the wind outflow from these stars, and found that this could indeed explain the initial decrease in the X-ray luminosity, but that in going to even later B V stars another problem arose. The emission measures of the X-ray producing material at spectral type B3 V and later become larger than the predicted wind emission measures for B stars for the smooth wind case. Cohen \\etal (1997) suggested that this could be the first indication that the late B stars lie at the transition to the outer atmospheric structure of cool stars, for which surface magnetic fields control the X-ray properties. These ideas were based on a spherical radial outflow picture for B stars.   In fact, B stars are known to be rather rapid rotators. Bjorkman \\& Cassinelli (1993) developed the Wind Compressed Disk (WCD) model, which suggested the wind from a rapidly rotating star would orbit towards the equatorial region where it would shock and compress the incident gas. The model had success in explaining the polarization properties of emission line Be stars (Wood \\etal 1997). WCD was also supported, initially, by hydrodynamic simulations performed by Owocki, Cranmer \\& Blondin (1994). However, in a more detailed consideration of the flow to the equator idea, Owocki, Cranmer \\& Gayley (1996) found that non-radial line forces in a rotating and distorted star tend to impede the flow to the equator and produce a bipolar flow instead. Hence, the cause of the {\\it disks} that exist around Be stars as opposed to {\\it polar plumes}, has become a topic of much debate among theorists.   Observers also found problems with the WCD idea. Hanuschik \\etal (1996) found in their observations of {\\it equator-on} Be stars, that the mass outflow speeds {\\it detectable} in the disks were negligible compared with the steady but slow equatorial outflows predicted by the WCD model, though no one seems ever to have predicted whether in fact the WCD material is dense enough to be detectable in this way. Even more interesting was their observational conclusion that the azimuthal speed of the inner disk material was larger than the angular speed of the star from which the disk presumably originated (as it must be to remain in Keplerian orbit unless supported by other forces.) Specifically, observations of the equator-on Be star $\\beta^1$ Mon showed that the Fe II line emission arising from the disk is broader than the $\\vsini$ value derived from photospheric lines.  Hanuschik \\etal\\ (1996) suggested that the equatorial disk material was in Keplerian motion about the star.  In the context of any Keplerian paradigm, it is especially important to note that in order to form a Keplerian disk, there needs to be an increase in the specific angular momentum of the matter after it leaves the star. The mechanism for providing that additional angular momentum is currently an unresolved subject of debate (e.g. Baade \\& Ud-Doula 2005, and Brown \\& Cassinelli 2005).   Transfer of mass and torquing of the outflow could be produced by magnetic fields rooted in the star's surface, as is well-known from magnetic rotator theory (Lamers \\& Cassinelli 1999 Ch. 9).  The existence of magnetic fields in hot stars is now well established (Donati \\etal 2001, 2002). Thus, the Magnetically Torqued Disk (MTD) model was proposed by Cassinelli et al. (2002) for the disks around Be stars. In this model, the star is pictured as having a co-aligned dipolar field that both channels and torques the wind from the star towards a disk.  Being that the detailed structure of the magnetic fields of these stars has not yet been determined, it seems reasonable to use the pure dipole as the most conservative hypothesis. Minimal magnetic fields were derived from the need to torque the outflow and the much denser disk material to Keplerian speeds, and the fields required were compared with upper and lower limits for hot star fields that had been derived by Maheswaran \\& Cassinelli (1988, 1992). For stars of spectral class B2 V (which correspond to the most common class of Be stars), the field required to torque the dense disk is about 300 Gauss while that required to torque the wind is only around 10 Gauss. This difference is because the field needed to torque a wind is proportional to the square root of the density, and the density of the wind is about 3 orders of magnitude lower than the density at the equatorial plane in the disk. The fact that the higher figure is comparable to the fields that have now been derived from multi-line Zeeman effect measurements of the very slowly rotating star $\\beta$ Cephei (Donati \\etal 2001) shows that both the flow and the disk will in reality undergo MTD torquing  for this star.  In other stars, the magnetic fields (whose strengths have not yet been measured) may lie in an intermediate regime in which the wind material would be torqued to high specific angular momentum while being channeled, but have entirely Keplerian dynamics in the denser disk (Owocki \\etal 2005). The MTD model was shown to produce the \\Halpha\\ emission observed in Be stars (Doazan \\etal 1991), and also the level of intrinsic polarization seen (Quirrenbach \\etal 1998). It is important to note that only a fraction of B stars are Be stars, and those stars identified as Be stars only spend a fraction of their time in a state with identifiable Be-star features. Therefore it is not necessary for a theoretical paradigm to cause a magnetically torqued disk for all possible sets of stellar and magnetic parameters, it is only necessary for a disk model to encompass a wide enough range of parameter space so as to make it reasonable that some stars show disks some of the time.  Detailed comparison with observations will only be possible when the ``duty cycles'' of Be stars and the Be star fraction have been better determined observationally.   In the original MTD model, the stellar wind mass flux and wind speed distribution were taken to be uniform over the stellar surface. However, in the case of a rapidly rotating star, this assumption is invalid since the rotation results in gravity darkening (Von Zeipel 1924) and reduces the wind mass flux and the terminal speed in the equatorial region (Owocki \\etal 1998). By incorporating Gravity Darkening into the MTD model (MTDGD), Brown \\etal (2004) derived several important disk properties such as, the dependence of the disk mass density distribution on its extent, the total number of disk particles, and the functional dependencies of the emission measure and polarization on the rotation rate ($S_o$) and wind velocity law ($\\beta$). In contrast to what had been expected, they found that the critical rotation (or $S_o= 1$) is not optimal for creation of hot star disks.  One important omission in the basic MTD formulation, as well as in MTDGD to date was the erroneous neglect of gravity in the disk density structure - see Brown \\& Cassinelli (2005). Though $g_z \\propto z$ is zero in the equatorial plane, its increase with $z$ enhances the density in the disk and causes it to grow with time. It will likely increase the \\Halpha\\ and polarization predictions. In turn, this will result eventually in radial escape of disk material, though whether this is slow and steady or episodic as claimed by Owocki \\& ud-Doula (2003) is not yet clear.   Other computational and analytic attempts to model disks similar to those envisioned here have met with mixed results. Owocki \\& ud-Doula (2003) criticized the MTD idea because the model was found to be unstable in their MHD simulations.  However, the numerical simulations by Keppens \\& Goedbloed (1999, 2000), and Matt \\etal (2000) demonstrated the existence of disks around some hot stars like post-AGB stars. Also Maheswaran (2003) has studied this scenario using analytic MHD and finds results that the magnetically torqued disks are likely to be persistent, which cast doubt on the numerical simulations of Owocki \\& ud-Doula (2003).  Thus further work is needed to test the basic ideas of the MTD and MTDGD models, either thorough numerical and analytic MHD calculations, or using observational diagnostics from radio to X-ray wavelengths.   In this paper, however, we are primarily concerned with X-ray emission from the MTDGD models, emission which occurs well upstream of the dense disk, and explain the X-ray anomalies associated with Be stars. This will be little affected by the density in the disk itself though our use of MTD to find the outer disk radius will make our X-ray source emission measure estimates a little too high. A successful model should be able to explain the drop-off at B2 Ve, the apparently excessive X-rays of late BV stars, while using disk parameters consistent with theory and the entire set of observational data. This process can then be inverted to use X-ray properties of a star to derive limits on its wind, disk, and rotational properties.  In Section 2, we describe how X-ray emission is produced by the model. The effects of model parameters on X-ray emission are discussed in Section 3.  Comparisons of model predictions with both $ROSAT$ and $Chandra$ observations are presented in Section 4. The discussion and conclusions are presented in Section 5. ",
        "conclusions": "The magnetically torqued disk model including gravity darkening effects or MTDGD, has here been tested to see if it can explain the basic X-ray properties of Oe/Be stars for which the model was developed. We had already found in the original papers on the MTD concept by Cassinelli \\etal (2002) and Brown \\etal (2004) that the idea of mass in Be disks is channeled by magnetic fields is consistent with the \\Halpha\\ luminosities of Be stars and that the mass in the disks as derived from polarization observations is also explainable, though we again note that inclusion of $g_z$ will likely increase these in the model.  We have used a model that explains how matter can enter a disk with sufficient angular momentum to explain the quasi-Keplerian disks of Be stars, and the model uses field strengths that are comparable to those being found for other B stars.  However an essential property of the MTDGD model is that it requires that X-rays be produced owing to the abrupt braking of the wind at the shock fronts. In summary, the paper contains several interesting results. (a) The paper tests the prediction that X-rays should be produced by the impact of channeled winds onto a disk. These X-rays would probably not be predicted from other current Be star models such as those in which the disk is produced by an extraction of angular momentum from the surface of a critically rotating star. (b) The model was based on the assumption that Be stars are rotating at their traditional values of about 70 percent critical and these rotation rates were found to be sufficient to explain the X-ray emission within the context of the MTD picture. (c) The model was found to require fields of order $10^2$ Gauss for Be stars, and our results show that these are adequate for the broad band X-ray production. (d) Broad band $ROSAT$ X-ray fluxes can be produced from B1 to B8 with MTDGD model parameters. (e) Fits are achievable even for the late B8 stars without invoking the presence of a dwarf M companion. (f) The model predicted that the Helium-like ion Mg XI had an $\\rm R_{fir}$ about 1.8 stellar radii, which agrees well with where that line emission is expected to arise in $\\zeta$ Oph with $Chandra$ observations, i.e., the $fir$ lines are predicted to occur there and to be broad, with a half width of about 400 $\\kmsec$ and the temperature structure is consistent with the formation of the Mg XI line.   Our discussion thus far has dealt with understanding the fundamental properties of Be stars as revealed by the $ROSAT$ and $Chandra$ observations. We have raised several questions during this paper that can now be addressed. (1) For our latest star B7 IVe, we found that the observed $ROSAT$ level of X-rays could be produced if the velocity law had value $\\beta =0.77$. This small value for $\\beta$ means that the channeled wind is colliding with the disk at a larger fraction of terminal wind speed than is the case for our other stars, which have $\\beta$ values ranging from 1.3 to 2.8. The other required parameters for this star seem plausible: $B \\sim 50$ Gauss, $S_o \\sim 0.6$, (in the first of our two fitting procedures). Cohen \\etal (1997) suggested that the emission measure needed to explain X-rays from late B stars were excessive, and that perhaps X-rays from magnetically confined region at the base of the wind are needed. So from our model of this star it appears that a magnetically confined X-ray formation region at the base is not needed. A bipolar magnetic field could instead just be torquing and channeling the wind toward the disk via X-ray emitting shocks. (2) The sharp drop-off of the X-ray luminosity beyond about B2 V, also seems to be explainable with a plausible range of our $\\beta$, $S_o$, and magnetic field values of about 685 to 130 Gauss for our B2 and B3 stars.   It appears that the MTDGD model has the ability to answer two of the more difficult questions concerning existing X-ray observations of Be stars, with plausible parameters. The model can explain the X-ray luminosity across the B spectral band and it can explain reasonably well the observed line profile results from $Chandra$. Finally, it is important to note that our model results are not dependent on the nature of the high density regions on the equatorial plane for which there is controversy regarding field wrapping and magnetic breakouts. This is because our models are in effect providing information only about the X-ray formation regions at the boundaries of the disk, and the flow through these boundaries occurs before the matter reaches the cooled compressed region near the equatorial plane. We are not requiring that the gas be controlled by a strong field all the way to the equatorial plane, and in fact think that the gas is no longer dominated by the field in that region and is free to acquire a quasi-Keplerian orbital motion. The needed angular momentum had been transferred to the gas in the pre-shock magnetically torqued and channeled MTD flow."
    },
    "0710/0710.5822_arXiv.txt": {
        "abstract": "Magnetar's persistent emission above 10 keV was recently discovered thanks to the imaging capabilities of the IBIS coded mask telescope on board the \\int~ satellite. The only two sources that show some degree of long term variability are \\zerosei~ and \\rxs. We find some indications that variability of these hard tails could be the driver of the spectral variability measured in these sources below 10 keV.  In addition we report for the first time  the detection at 2.8 $\\sigma$ level of pulsations in the hard X-ray tail of \\zerosei. ",
        "introduction": "Anomalous X-ray Pulsars (AXPs) and Soft Gamma-Ray Repeaters (SGRs) are two small classes of sources that are believed to be magnetar candidates, namely isolated neutron stars with magnetic fields larger than the quantum critical value B$_{QED}\\sim$4.4$\\times$10$^{13}$G. For a review of this class of objects see \\cite{woodsrew}.  They share some common properties like a long spin period in the range of $P$=2--12 s, a large period derivative of $\\dot P$=10$^{-13}$--10$^{-10}$ s s$^{-1}$, and a typical X-ray luminosity of $L_{X}\\sim 10^{34}$--10$^{36}$ erg s$^{-1}$. This luminosity is well above the rotational energy losses, and is believed to be powered by the decay of the huge magnetic field.  Since the spectra below $\\sim$10 keV are rather soft, the first \\int~ detections above 20 keV of very hard high-energy tails associated with these objects came as a surprise \\citep{kuiper,denhartog,rev,mere05,molkov,dg06}. The spectra flatten ($\\Gamma\\sim$ 1, where $\\Gamma$ is the photon index) above 20 keV and the pulsed fraction of some of them reaches up to 100\\% \\citep{kuiper}. The discovery of these hard tails provides new constraints on the emission models for these objects since their luminosities might well be dominated by hard, rather than soft, X-rays.  In this paper we will focus on the timing properties and long term variations of two magnetar candidates, \\rxs~ and \\zerosei, as measured by IBIS/ISGRI \\cite{ibis,isgri} on board the \\int~ satellite \\cite{integral}. ",
        "conclusions": ""
    },
    "0710/0710.2543_arXiv.txt": {
        "abstract": "Transonic accretion flow with self-consistent treatment of the magnetic field is presented. My website \\url{http://www.cfa.harvard.edu/~rshcherb/}. ",
        "introduction": "The averaged quantities can be obtained in two different ways in magnetohydrodynamics. The first way is to solve 3D MHD equations and then average the results. The second way is to solve some system of equations on averages. Combination of numerical simulations and averaged theory brings phenomenology that can describe observations or experimental data.  The problem of spherically symmetric accretion takes its origin from Bondi's work \\citep{bondi}. He presented idealized hydrodynamic solution with accretion rate $\\dot{M}_B.$ However, magnetic field $\\vec{B}$ always exists in the real systems. Even small seed $\\vec{B}$ amplifies in spherical infall and becomes dynamically important \\citep{schwa}.  Magnetic field inhibits accretion \\citep{schwa}. None of many theories has reasonably calculated the magnetic field evolution and how it influences dynamics. These theories have some common pitfalls. First of all, the direction of magnetic field is usually defined. Secondly, the magnetic field strength is prescribed by thermal equipartition assumption. In third, dynamical effect of magnetic field is calculated with conventional magnetic energy and pressure. All these inaccuracies can be eliminated.  In Section 2\\ref{section_method} I develop a model that abandons equipartition prescription, calculates the magnetic field direction and strength and employs the correct equations of magnetized fluid dynamics. In Section 3\\ref{results} I show this accretion pattern to be in qualitative agreement with Sgr A* spectrum models. I discuss my assumptions in Section 4 \\ref{discussion}. ",
        "conclusions": "\\label{discussion} The presented accretion study self-consistently treats turbulence in the averaged model. This model introduces many weak assumptions instead of few strong ones.  I take dissipation rate to be that of collisional MHD simulations. But flow in question is rather in collisionless regime. Observations of collisionless flares in solar corona \\citep{noglik} gives dissipation rate $20$ times smaller than in collisional simulations \\citep{biskamp03}. However, flares in solar corona may represent a large-scale reconnection event rather than developed turbulence. It is unclear which dissipation rate is more realistic for accretion.  Magnetic field presents another caveat. Magnetic field lines should close, or $\\vec{\\nabla}\\cdot\\vec{B}=0$ should hold. Radial field is much larger than perpendicular in the inner region. Therefore, characteristic radial scale of the flow is much larger than perpendicular. If radial turbulence scale is larger than radius, freezing-in condition does not hold anymore. Matter can freely slip along radial field lines into the black hole. If matter slips already at the sonic point, the accretion rate should be higher than calculated.  Some other assumptions are more likely to be valid. Diffusion should be weak because of high Mach number that approaches unity at large radius. Magnetic helicity was found to play very small dynamical role. Only when the initial turbulence is highly helical, magnetic helicity conservation may lead to smaller accretion rate. Neglect of radiative cooling is justified a posteriori. Line cooling time is about $20$ times larger that inflow time from outer boundary.  The study is the extension of basic theory, but realistic analytical models should include more physics. The work is underway. \\begin{theacknowledgments} I thank my advisor Prof. Ramesh Narayan for fruitful discussions. \\end{theacknowledgments}"
    },
    "0710/0710.0546_arXiv.txt": {
        "abstract": "Estimates of velocities from time series of photospheric and/or chromospheric vector magnetograms can be used to determine fluxes of magnetic energy (the Poynting flux) and helicity across the magnetogram layer, and to provide time-dependent boundary conditions for data-driven simulations of the solar atmosphere above this layer. Velocity components perpendicular to the magnetic field are necessary both to compute these transport rates and to derive model boundary conditions.  Here, we discuss some possible approaches to estimating perpendicular flows from magnetograms.  Since Doppler shifts contain contributions from flows parallel to the magnetic field, perpendicular velocities are not generally recoverable from Doppler shifts alone. The induction equation's vertical component relates evolution in $B_z$ to the perpendicular flow field, but has a finite null space, meaning some ``null'' flows, e.g., motions along contours of normal field, do not affect $B_z$.  Consequently, additional information is required to accurately specify the perpendicular flow field.  Tracking methods, which analyze $\\partial_t B_z$ in a neighborhood, have a long heritage, but other approaches have recently been developed.  In a recent paper, several such techniques were tested using synthetic magnetograms from MHD simulations.  Here, we use the same test data to characterize: 1) the ability of the induction equation's normal component, by itself, to estimate flows; and 2) a tracking method's ability to recover flow components that are perpendicular to $\\mathbf{B}$ and parallel to contours of $B_z$.  This work has been supported by NASA Heliophysics Theory grant NNG05G144G. ",
        "introduction": "\\label{sec:intro}  The large length scales and relatively high conductivity of the plasma in the solar corona imply that, to a good approximation, magnetic flux is frozen to the plasma there.  Consequently, the coronal magnetic field is ``line-tied'' to the plasma in lower atmospheric layers where hydrodynamic forces can be stronger than Lorentz forces --- the photosphere and lower chromosphere --- and coronal evolution is strongly coupled to evolution in these layers.  Accordingly, observations of magnetic field evolution below the Sun's corona --- typically, sequences of photospheric or chromospheric magnetograms --- provide crucial tools understand coronal evolution.  Usually, vector magnetograms are more useful than line-of-sight (LOS) magnetograms for studying coronal evolution, because information derived from LOS measurements alone will not, in general, be consistent with the actual magnetic field, which has field components both parallel and transverse to the LOS.  Although time series of vector magnetograms have historically been rare, SOLIS \\cite{BW_Henney2002}, the Solar Optical Telescope (SOT, Tarbell 2006) \\nocite{BW_Tarbell2006} on Hinode, and the Solar Dynamics Observatory's Helioseismic and Magnetic Imager (HMI) \\cite{BW_Scherrer2005}, should dramatically improve photospheric vector magnetogram spatial and temporal coverage in the near future.  Several techniques have been developed to derive flows from time series of magnetograms \\cite{BW_Chae2001,BW_Kusano2002,BW_Welsch2004, BW_Longcope2004,BW_Georgoulis2006,BW_Schuck2006}. Estimated flows at the base of the corona can be used to derive the fluxes of magnetic helicity, energy, and free energy into the corona, \\cite{BW_Demoulin2003,BW_Pariat2005,BW_Welsch2006}.  Further, flow estimates can be used to provide time-dependent boundary conditions for data-driven simulations of coronal magnetic field evolution.  This paper briefly reviews progress on estimating surface flows from magnetogram sequences, and demonstrates some aspects of the problem with test data. ",
        "conclusions": "\\label{sec:discuss}  We have briefly reviewed central concepts regarding surface flow estimation from magnetograms.  In addition, we have presented the results of simple tests which demonstrate that the induction equation's normal component, equation (\\ref{eqn:normal}), can be used to quantitatively estimate flux transport rates.  Fluxes of magnetic energy and helicity derived from these estimates are, however, likely to possess significant systematic errors.  We have also demonstrated that flow estimation techniques that depend upon evolution in $B_z$ alone are insensitive to flux transport along contours of $B_z$, compared to flux transport along $\\nabla_h B_z$.  Equations (\\ref{eqn:dSdt}) and (\\ref{eqn:dhdt_bf}), which depend on the dot products of $\\mathbf{u} B_z$ with $\\mathbf{B}_h$ and with $\\mathbf{A}_P$ respectively, combined with the sensitivity of $\\mathbf{u} B_z$ to $\\nabla_h B_z$, suggest that if $\\mathbf{B}_h$ and/or $\\mathbf{A}_P$ lie primarily along $\\nabla_h B_z$, then the fluxes of magnetic energy and/or helicity can, in principle, be recovered accurately from $\\Delta B_z/\\Delta t$.  If, in contrast, $\\mathbf{B}_h$ and/or $\\mathbf{A}_P$ lie primarily along contours of $B_z$, then the fluxes of magnetic energy and/or helicity probably cannot be recovered accurately from evolution in $B_z$ alone.  We note that, as discussed in \\cite{BW_Welsch2007}, the ANMHD data used in the tests presented here differ from actual magnetograms in significant ways, so the properties of flows estimated from actual magnetograms will probably differ substantially from the properties of flows estimated from ANMHD data."
    },
    "0710/0710.2858_arXiv.txt": {
        "abstract": "GRB 021206 is one of the brightest GRBs ever observed. Its prompt emission, as measured by RHESSI, shows an unexpected spectral feature. The spectrum has a peak energy of about 700~\\keV\\ and can be described by a Band function up to 4.5~\\MeV. Above 4.5~\\MeV, the spectrum hardens again, so that the Band function fails to fit the whole RHESSI energy range up to 17~\\MeV. Nor does the sum of a blackbody function plus a power law, even though such a function can describe a spectral hardening. The cannonball model on the other hand predicts such a hardening, and we found that it fits the spectrum of GRB 021206 perfectly. We also analysed other strong GRBs observed by RHESSI, namely GRBs 020715, 021008, 030329, 030406, 030519B, 031027, 031111. We found that all their spectra can be fit by the cannonball model as well as by a Band function. ",
        "introduction": "\\label{sec:intro}  The exact mechanism which produces  $\\gamma$-ray bursts (GRBs) has not yet been definitively established. Their prompt $\\gamma$-ray spectra can be used to distinguish between different models. Several mathematical functions have been used for parametrizing the prompt $\\gamma$-ray emission. Most commonly used is the empirical Band function \\citep{Band93}, which is not motivated by a physical model.  There have been attempts to distinguish between spectral models analysing the low energy part of the spectrum. \\citet{Ghirlanda2003},  \\citet{Ryde2004}, and more recently \\citet{Ghirlanda2007} searched for blackbody components in GRB spectra with varying degrees of success. \\citet{Preece2002}, using BATSE GRB spectra, tested the synchrotron shock model and conclude that it ''does not account for the observed spectra during the GRB phase''.   Spectral studies above the peak energy are rare, one reason being the poor data quality because of lack of statistics. Combining BATSE and EGRET spectra, \\citet{Nature2003} report a high energy component for GRB 941017. They find a photon index of about 1.0 at energies above 5~\\MeV.  In this paper we report a high energy component in GRB 021206 \\citep{GCN_021206_IPN,GCN_021206_final}, observed with the Reuven Ramaty High Energy Solar Spectroscopic Imager RHESSI \\citep{RHESSI}. Having a peak energy of about 700~\\keV, the spectrum of this burst can be described by a Band function from 70~\\keV\\ up to 4.5~\\MeV, with a high energy photon index $\\beta \\approx 3.2$. Above 4.5~\\MeV, the spectrum hardens again, and can be described with a photon index  $\\beta' \\approx 2.2$. This significant hardening around 4.5~\\MeV\\ can not be described with a Band function. But it seems to differ from the spectral hardening in GRB 941017 as well.   There is one model that fits the entire RHESSI spectrum of GRB 021206: the cannonball model \\citep{Dar2004,Dado2002,Dado2003}. The cannonball model predicts a spectral hardening at several times the peak energy with a high energy photon index reaching $\\beta \\approx 2.1$.  The question immediately arises whether the cannonball model can improve our description of other GRB spectra. The difference between the Band function and the cannonball model arises only at the high energy part of the spectrum, where data usually suffer from low statistics. Therefore, we choose the strongest GRBs registered by RHESSI in the years 2002 to 2004. We find that they all  can be fit by the cannonball model as well as by the Band function.  The outline of the paper is the following: We first present shortly the instrument, the GRB selection, the spectrum extraction, and the fit method (\\S \\ref{sec:observations}). In the next section (\\S \\ref{sec:models}), many spectral functions are given. In \\S \\ref{sec:results}, the fit results for GRB 020715, GRB 021008, GRB 021206, GRB 030329, GRB 030406, GRB 030519B, GRB 031027, and GRB 031111 are presented. The fits are discussed and, if possible, compared to other measurements. The more general discussion, including an outlook, follows in \\S \\ref{sec:discussion}. We end with a short summary in \\S \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{The spectral functions}  What is an acceptable $\\chi^2$? In the limit of many degrees of freedom ($n_{DoF} > 30$), $\\chi^2$ is normal distributed with an expectation value of $n_{DoF}-0.5$ and a variance of $\\sigma_{\\chi^2} = \\sqrt{2n_{DoF}-1}$. A fit is acceptable if $\\chi^2$ is close to its expectation value {\\em and} if the residuals scatter around zero over the whole fit range, i.e.\\ if the fit ``looks good''.  From Table~\\ref{tab:chi2} we conclude that the \\CB\\ gives acceptable $\\chi^2$ for {\\em all} GRBs studied. And they also look good, as can be seen in Figs.~\\ref{fig:020715} to \\ref{fig:031111}. Except for GRB 021206, rear (Fig.~\\ref{fig:spec_rear}), the same can be said for the Band function.  In many cases, a cut off power law (CPL) fits the spectrum up to high energies, e.g.\\ GRB 021008, GRB 030329 or GRB 031027. In these cases, Band and \\CB\\ improve the goodness-of-fit slightly, but all three spectral shapes fit the data acceptably.  A broken power law fits sometimes, but usually not well.  BBPL and BBmPL do not fit in general, BBPL worse than BBmPL. However, it should be mentioned that a blackbody component is expected---if at all---only at the beginning of a GRB (see e.g.\\ \\citet{Ryde2006} and references therein), whereas we fitted the entire duration of the bursts. When using BBPL, we often find that the PL component fits either at high energies or at low energies. This is also discussed by \\citet{Ghirlanda2007} who studied six BATSE GRBs in detail, where low energy data from the WFC instrument (on board BeppoSAX) are available. They find that the WFC data fit the Band function or CPL extrapolation, but not the BBPL extrapolation to low energies. Arguing that the PL contribution is too simple, they try to fit a blackbody spectrum plus CPL.  % We suggest to use our BBmPL function instead. Its modified PL component describes a spectral break from $dN/dE \\propto E^{-(\\beta-1)}$ at low energies to $dN/dE \\propto E^{-\\beta}$ at high energies.   \\subsection{\\CB\\ function }  The present work is, to our knowledge, the first systematic attempt to fit the \\CB\\ function to prompt GRB spectra. The two terms in \\refeq{eq:CB} have a simple meaning. According to the cannonball theory, all GRBs are associated with a supernova. The ambient light is Compton up-scattered by the cannonball's electrons, producing the prompt GRB emission. Some electrons are simply comoving with the cannonball, giving rise to the CPL term in \\refeq{eq:CB}. Since the photon spectrum of the ambient light can be described by a thin thermal bremsstrahlung spectrum, $\\alpha$ is expected to be $\\approx 1$. The second term (mPL) is caused by a small fraction of electrons accelerated to a power law distribution, resulting in a photon index of $\\beta \\approx 2.1$. See e.g.\\ \\citet{Dado2005}, \\S 3.8. or \\citet{Dado2007}, \\S 2 and 4.1 for a summary.  In our study, the observed values for $\\alpha$ are all approximately 1, as predicted by the \\CB. Because of the low count statistics at high energies, we could not always fit $\\beta$. We then fixed it to its theoretical value of 2.1 in order to make the fit converge and to obtain a value for the parameter $b$. In the cases where we could fit $\\beta$, we found values close to 2.1 (Table~\\ref{tab:CB_pars}).  For the factor $b$ of the modified PL component in the \\CB\\ function we typically found values of the order $0.1\\,$. An exception is GRB 031111, where $b$ is of the order 1.0, but with a large error ($0.4$).  Our values for $\\beta$ and $b$ are similar to the ones found by \\citet{Dado2005} (fit of GRB 941017) and by \\citet{Dado2004} (fit  of $X$-ray flashes XRF 971019, XRF 980128, and XRF 990520 using BeppoSAX/WFC and CGRO/BATSE data). The authors of the \\CB\\ hypothesize that XRFs are simply GRBs viewed further off the jet axis.  The different time development of the CPL- and the mPL-fluences, as reported in Table~\\ref{tab:dt_grb021206}, possibly point to a different time dependence of the two underlying electron distributions within a cannonball.   \\subsection{Fitting \\CB\\ function versus CPL and Band function} Both the Band function and the \\CB\\ function are extensions of a CPL, the Band function with one additional parameter, the \\CB\\ function with two. For cases where a CPL fits the data well, also a Band function with $\\beta = \\infty$ or a \\CB\\ function with $b=0$ (and $\\beta=2.1$ or any other value) fits. This is the case for GRB 031027.  Whether additional parameters are necessary in a fit, can be tested with the $F$-test. For GRB 030329, the extra parameters are barely needed. For GRB 020715, 021008, 021206, 030406, 030519B, 031111 additional parameters are required at a confidence level of at least 90\\%.  Concerning the question of whether the high energy power law parameter $\\beta$ in the \\CB\\ should be treated as free parameter, the answer is 'yes' from a theoretical point of view, but in practice, see Table~\\ref{tab:chi2}, the improvements in $\\chi^2$ are marginal or small for all bursts we studied. Our practice to freeze $\\beta$ at its theoretically predicted value in cases of bad convergence seems to be acceptable.  It is more difficult to compare the goodness of fit using the Band function compared to using the \\CB\\ function. The two functions are not independent, because they both are dominated by a CPL up to the peak energy and higher. In most cases of our study, the two functions fit the observed spectrum equally well with a slight preference for the Band function. At high energies however (typically above several times the peak energy) the two functions are different, the spectral hardening being a unique feature of the \\CB\\ function. There is only one case, namely GRB 021206, where this hardening is observed. For the rear data going up to high energies, a Band function fit gives $\\chi^2 / \\mbox{dof} = 133.3 / 74$ (see Table~\\ref{tab:chi2}). This is not acceptable at $<0.01$\\% probability of being accidentally so high. The \\CB\\ fit on the other hand gives  $\\chi^2 / \\mbox{dof} = 82.7 / 73$, which is fully acceptable at a $20$\\% level.  We would like to stress again that, while the \\CB\\ gives acceptable fits for {\\em all} cases, the Band function fails in one case. This seems enough to us to give some credit to the \\CB.  But it is, of course, no proof that the \\CB\\ is the only theory capable of describing the spectrum of GRB 021206. For example, a Band function plus a PL with $\\gamma\\approx 1.5$ would also fit. But there is no theory to predict such a shape. To our knowledge, \\CB\\ is to date the only existing GRB model % that explains the prompt GRB spectra from first principles.  At this place we also would like to note the the mean $\\alpha$-value found for the BATSE catalogue is 1 \\citep[see][]{BATSEcatalog}). We cite from their summary: ``{\\em We confirmed, using a much larger sample, that the most common value for the low-energy index is $\\approx -1\\,$}\\footnote{ this corresponds to $+1$ in our notation} {\\em \\citep{Preece2000, Ghirlanda02}. The overall distribution of this parameter shows no clustering or distinct features at the values expected from various emission models, such as $-2/3$ for synchrotron \\citep{Katz94, Tavani96}, $0$ for jitter radiation \\citep{Medv00}, or $-3/2$ for cooling synchrotron \\citep{GhisCel99}.}'' They do not mention the \\CB\\ which would explain $\\alpha \\approx 1$.  Note that the $\\beta$ values of the \\CB\\ are systematically lower than the $\\beta$ values of the Band function, compare Tables \\ref{tab:CB_pars} and \\ref{tab:Band_pars}. From Band function fits to BATSE GRBs, it is known that $\\beta$ is clustered around 2.3, with a long tail towards higher values,   see \\citet{BATSEcatalog}. For \\CB\\ we would expect $\\beta$ to cluster at slightly lower values.  For criticism of the \\CB, see e.g.\\ \\citet{Hillas}, but see also the answer by \\citet{Dar2006}.   \\subsection{The spectral hardening}  The difference of a \\CB\\ spectrum and Band function is the hardening at high energies. This becomes visible---for the GRBs studied here---in the few \\MeV\\ region, but it depends on the peak energy and the factor $b$. For $\\alpha=1.0$ and $b=0.10$ the hardening typically appears at several times the peak energy and the second term dominates at 10 times the peak energy. For the spectral fit of XRFs done by \\citet{Dado2004}, the spectral hardening is expected in the few hundred \\keV\\ region, just where the number of photons detected runs low. Most of our GRBs also suffer from this lack of statistics at high energies, preventing the detection of a hardening.  A spectral coverage of two decades and good detection efficiency at high energies is necessary to experimentally observe the full shape of the \\CB\\ function. In the case of GRB 021206 we were able to detect this hardening, thanks to RHESSI's broad energy range (30~\\keV\\ to 15~\\MeV), and because this is one of the brightest GRBs ever observed.  There is a GRB observed by SMM from 20 keV up to 100 MeV, namely GRB 840805. As reported by \\citet{Share84}, the spectrum of this burst shows emission up 100 MeV. In order to fit the spectrum, ``a classical thermal synchrotron function plus a power law'' was used. The power law component was required to fit the data above about 6 MeV. This is a hint of a spectral hardening around 6 MeV, and we suppose that the spectrum of this GRB can be fit by a \\CB\\ function.  The spectral hardening observed in GRB 941017 \\citep{Nature2003} seems to be different. The photon index of GRB 941017 above a few \\MeV\\ is $\\approx 1.0$. This case is discussed by \\citet{Dado2005} as a possible additional feature in the \\CB\\ spectrum.   \\subsection{Outlook} \\label{sec:outlook}  In order to find more GRB spectra that show the hardening characteristic for the \\CB\\ function, strong GRBs have to be observed over a broad enough energy range. With the forthcoming GLAST mission, we expect that more such spectra will be observed. But also joint analyses with more than one instrument could reveal this hardening. We therefore suggest:   $\\bullet$ to search for \\CB\\ spectrum candidates among joint Swift/RHESSI GRBs and XRFs, and joint Swift/Konus GRBs.  $\\bullet$ to reanalyse some BATSE bursts. Looking at the BATSE spectra published by \\citet{Ghirlanda2007}, we suppose that the \\CB\\ can possibly improve the fits of GRB 980329, GRB 990123, and GRB 990510. The same can be said for GRB 911031 as published by \\citet{Ryde2006}. And GRB 000429, as published in Fig.\\ 19 of \\citet{BATSEcatalog}, looks like a promising candidate as well.  $\\bullet$ to search in KONUS data for suitable GRBs.  $\\bullet$ to add the \\CB\\ function to XSPEC in order to make it more accessible to the astronomical community."
    },
    "0710/0710.4153_arXiv.txt": {
        "abstract": "It is shown that nearly-flat  3+1D spacetime emerging from  a dual quantum field theory in 2+1D   displays quantum fluctuations from classical Euclidean geometry on macroscopic scales.  A covariant holographic mapping is assumed,  where plane wave states with wavevector $\\vec k$  on a  2D surface map onto classical null trajectories  in the emergent third dimension at an angle $\\vec\\theta=l_P\\vec k$ relative to the surface element normal, where $l_P$ denotes the Planck length. Null trajectories in the 3+1D world then display   quantum uncertainty of angular orientation,  with standard deviation $\\Delta\\theta=\\sqrt{l_P/z}$ for longitudinal propagation distance $z$ in a given frame.   The quantum complementarity of transverse position at macroscopically separated events along null trajectories corresponds to  a geometry that is not completely classical, but  displays observable  holographic quantum noise. A statistical estimator of the fluctuations  from Euclidean behavior is given for a simple thought experiment based on measured sides of   triangles. The effect can be viewed as sampling noise due to the limited degrees of freedom of such a theory,  consistent with   covariant bounds on entropy. ",
        "introduction": "It is not understood  how a fundamentally quantum world creates the appearance, to internally constituted observers,  of an approximately classical 3+1D classical spacetime, within which quantum fields operate.  There are however indications that general relativity  is in some sense ``holographic'' and that spacetime somehow emerges from a quantum theory with lower dimensionality.  Examples exist where a conformal field theory with  D spatial dimensions appears to be dual to a (supersymmetric) quantum field theory with gravity in D+1 spatial dimensions\\cite{Maldacena:1997re,Witten:1998zw,Aharony:1999ti,Alishahiha:2005dj,Horowitz:2006ct}. Arguments based on black  hole thermodynamics and evaporation\\cite{Bekenstein:1972tm,Bardeen:gs,Bekenstein:1973ur,Bekenstein:1974ax,Hawking:1975sw,Hawking:1976ra,'tHooft:1985re,Susskind:1993if},   string theory\\cite{Strominger:1996sh},  thermodynamics  in nearly-flat space\\cite{Jacobson:1995ab},  and classical relativity\\cite{Padmanabhan:2006fn,Padmanabhan:2007en,Padmanabhan:2007xy, Padmanabhan:2007tm}, have led many authors to suspect that a region of  3+1D spacetime with gravity  and the fields within it may have a  dual description in terms of null surfaces swept out by a 2+1D quantum    theory with a cutoff at the Planck length $l_P$ \\cite{Padmanabhan:2006fn,Padmanabhan:2007en,Padmanabhan:2007xy, Padmanabhan:2007tm,'tHooft:1993gx,Susskind:1994vu,'tHooft:1999bw,Bigatti:1999dp,Bousso:2002ju}.  This paper shows that such holographic theories generally display directly measurable quantum departures from classical general relativity.  Recently, the author has argued\\cite{Hogan:2007rz,Hogan:2007hc,Hogan:2007ci} that  a nearly-flat  holographic  spacetime   approaches classical behavior  on macroscopic scales   in a distinctive way that  displays  small but detectable quantum fluctuations from  Euclidean geometry even at macroscopic separations. Specifically, the   classical metric at widely separated points along any null trajectory displays a quantum complementarity of transverse position.  Conversely, at small separations, angles become increasingly poorly defined  until at the Planck scale the third dimension melts into the behavior of a 2+1D theory.  This paper shows how this behavior arises simply from the  geometry of the holographic mapping of the 2+1D  states into 3+1D. This exercise helps to define the class of theories that have general relativity as an approximate classical limit, and that display the phenomena of  holographic indeterminacy and noise. ",
        "conclusions": ""
    },
    "0710/0710.3545_arXiv.txt": {
        "abstract": "{ For some time, the Draco dwarf spheroidal galaxy has garnered interest as a possible source for the indirect detection of dark matter.  Its large mass-to-light ratio and relative proximity to the Earth provide favorable conditions for the production of detectable gamma rays from dark matter self-annihilation in its core. The Solar Tower Atmospheric Cherenkov Effect Experiment (STACEE) is an air-shower Cherenkov telescope located in Albuquerque, NM capable of detecting gamma rays at energies above 100 GeV.  We present the results of the STACEE observations of Draco during the 2005-2006 observing season totaling 10 hours of livetime after cuts. }  \\begin{document} ",
        "introduction": "Dark matter is thought to be an important component of the Universe and research into its nature is actively pursued using a variety of techniques.  Dark matter may be weakly interacting massive particles (WIMPs) which would tend to accumulate at the bottom of gravitational potential wells, such as galaxies, where they could undergo self-annihilation processes. Depending on WIMP mass and branching ratios, a measurable flux of high energy gamma rays could result.  The Draco dwarf spheroidal galaxy has long garnered interest as a potential source of concentrated dark matter~\\cite{Tyler}. Draco has one of the highest known mass-to-light ratios ($M/L$), perhaps as high as $500 M_\\odot/L_\\odot$~\\cite{Wu}.  Current observations are consistent with a cuspy density profile~\\cite{Lokas}, which would enhance the gamma-ray production rate. Furthermore, since Draco is a satellite of the Milky Way, its relative proximity to the Earth ($d \\sim 75~kpc$)\\cite{Bonanos} might allow for the detection of such gamma rays. ",
        "conclusions": "STACEE does not detect a significant signal from Draco, and sets upper limits on cross-sections for WIMP with rest-mass energy greater than about 150 GeV.   {\\bf Acknowledgments:} Many thanks go to the staff of the National Solar Tower Test Facility, who have made this work possible. This work was funded in part by the U.S. National Science Foundation, the Natural Sciences and Engineering Research Council of Canada, Fonds Quebecois de la Recherche sur la Nature et les Technologies, the Research Corporation, and the University of California at Los Angeles."
    },
    "0710/0710.4962_arXiv.txt": {
        "abstract": "We have explored the galaxy disk/extended halo gas kinematic relationship using rotation curves (Keck/ESI) of ten intermediate redshift galaxies which were selected by {\\ggkMgII} halo gas absorption observed in quasar spectra.  Previous results of six edge--on galaxies, probed along their major axis, suggest that observed halo gas velocities are consistent with extended disk--like halo rotation at galactocentric distances of 25--72~kpc. Using our new sample, we demonstrate that the gas velocities are by and large not consistent with being directly coupled to the galaxy kinematics. Thus, mechanisms other than co--rotation dynamics (i.e., gas inflow, feedback, galaxy--galaxy interactions, etc.)  must be invoked to account for the overall observed kinematics of the halo gas.  \\noindent In order to better understand the dynamic interaction of the galaxy/halo/cosmic web environment, we performed similar mock observations of galaxies and gaseous halos in $\\Lambda$--CDM cosmological simulations. We discuss an example case of a $z=0.92$ galaxy with various orientations probing halo gas at a range of positions. The gas dynamics inferred using simulated quasar absorption lines are consistent with observational data. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.1630_arXiv.txt": {
        "abstract": "We derive new constraints on the Hubble function $H(\\phi)$ and subsequently on the inflationary potential $V(\\phi)$ from WMAP 3-year data combined with the Sloan Luminous Red Galaxy survey (SDSS-LRG), using a new methodology which appears to be more generic, conservative and model-independent than in most of the recent literature, since it depends neither on the slow-roll approximation for computing the primordial spectra, nor on any extrapolation scheme for the potential beyond the observable e-fold range, nor on additional assumptions about initial conditions for the inflaton velocity.  This last feature represents the main improvement of this work, and is made possible by the reconstruction of $H(\\phi)$ prior to $V(\\phi)$.  Our results only rely on the assumption that within the observable range, corresponding to $\\sim$ 10 e-folds, inflation is not interrupted and the function $H(\\phi)$ is smooth enough for being Taylor-expanded at order one, two or three.  We conclude that the variety of potentials allowed by the data is still large. However, it is clear that the first two slow-roll parameters are really small while the validity of the slow-roll expansion beyond them is not established. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.4934_arXiv.txt": {
        "abstract": "{} {We present the results of the analysis of an archival observation of LMC X--2 performed with XMM/Newton. The spectra taken by high-precision instruments have never been analyzed before.} {We find an X--ray position for the source that is inconsistent with the one obtained by ROSAT, but in agreement with the Einstein position and that of the optical counterpart. The correlated spectral and timing behaviour of the source suggests that the source is probably in the normal branch of its X-ray color-color diagram. The spectrum of the source can be fitted with a blackbody with a temperature 1.5 keV plus a disk blackbody at 0.8 keV. Photoelectric absorption from neutral matter has an equivalent hydrogen column of $4 \\times 10^{20}~\\mathrm{cm}^{-2}$. An emission line, which we identify as the O VIII Lyman-$\\alpha$ line, is detected, while no feature due to iron is detected in the spectrum.} {We argue that the emission of this source can be straightforwardly interpreted as a sum of the emission from a boundary layer between the NS and the disc and a blackbody component coming from the disc itself. Other canonical models that are used to fit Z-sources do not give a satisfactory fit to the data. The detection of the O VIII emission line (and the lack of detection of lines in the iron region) can be due to the fact that the source lies in the Large Magellanic Cloud. } {} ",
        "introduction": "Low  Mass   X-ray  Binaries  (LMXBs)   are  binary  systems harboring a compact object accreting mass from a late-type  star. Systems hosting a weakly magnetized neutron star (NS) are  usually divided into  two classes: the Z  sources, with  luminosities close  to  the Eddington luminosity, $L_{\\rm  edd}$, and  Atoll sources,  usually with  lower luminosities  of $\\sim  0.01-0.1\\ L_{\\rm edd}$.  This classification relies  upon the correlated X-ray spectral and timing properties, namely the  pattern traced out by individual sources in the X-ray  color-color diagram  (CD,  Hasinger \\&  van der  Klis 1989). The seven  known  (Galactic) Z sources usually describe  a complete Z-track in the  CD on timescales  of a  few days, while Atoll sources cover their pattern on the CD on a longer timescale (usually several weeks).  The correlated spectral and timing variability of LMC X--2 has suggested its classification as a Z-source (Smale \\& Kuulkers 2000). An extensive study of these characteristics, using data taken with the PCA on board RXTE, has allowed to observe the complete Z-track in the CD which is completed in about 1 day (Smale, Homan \\& Kuulkers 2003). Smale \\& Kuulkers (2000) found evidence of a periodicity of $8.16~\\mathrm{h}$ which they tentatively ascribe to an orbital period. A similar orbital period modulation has been found by Cornelisse et al. (2007) in the VLT optical lightcurve of the companion. From the shape of the Z-pattern described in the CD, it was inferred that LMC X-2 falls squarely into the category of the so-called ``Sco-like'' Z-sources due to its common characteristics with Sco X-1 (Smale et al. 2003). Its very high luminosity ($\\sim 0.5-2 L_\\mathrm{edd}$; Markert \\& Clark 1975; Johnston, Bradt \\& Doxsey 1979; Long, Helfand \\& Grabelsky 1981; Bonnet-Bidaud et al. 1989; Christian \\& Swank 1997; Smale \\& Kuulkers 2000) makes it the brightest LMXB known together with Sco X--1. There are other analogies with Sco X--1: the optical counterpart is a faint, $M_V\\sim 18.8$, blue star (similar to the optical counterpart of Sco X--1), and both sources show a similar correlation between the optical and the X-ray lightcurve during flares (McGowan et al. 2003). In some ways, we can consider LMC X--2 an extragalactic twin to Sco X--1.  Measurement of Doppler effect on emission lines in the Bowen region allowed Cornelisse et al. (2007) to constrain the mass ratio of LMC X-2 to be $\\le 0.4$.  However, there is no recent X-ray spectral study of this source: a spectral analysis has been carried out by Bonnet-Bidaud et al. (1989) using EXOSAT data, where the spectrum of the source was fitted with a model consisting of a blackbody $kT\\sim 1.3~\\mathrm{keV}$ plus a thermal bremsstrahlung ($kT\\sim 5~\\mathrm{keV}$), or alternatively with a comptonized thermal model with $kT\\sim 3~\\mathrm{keV}$. They found no detectable feature in the Fe K-$\\alpha$ range. In their survey of Low Mass X-Ray Binaries with Einstein, Christian \\& Swank (1997) report that a fit in the 1-20 keV band with an unsaturated Comptonized model (i.e. a cutoff power-law) gave a reasonable fit with $\\Gamma = 1.4$ and $kT = 6.4~\\mathrm{keV}$. Schulz (1999), using ROSAT data in the 0.1--2.4 keV band, fitted the data using a blackbody with a temperature $kT = 1.5~\\mathrm{keV}$ plus a thermal bresstrahlung model with a temperature of $kT=5~\\mathrm{keV}$. Also, a gaussian emission line at 0.9 keV was evident in the spectrum. The accuracy of this study is obviously limited by the narrow band and the relatively poor spectral resolution of ROSAT. Smale \\& Kuulkers (2000) studying RXTE/PCA data of the source found that the data can be well described by a cutoff power-law (e.g. a completely Comptonized component) with a cutoff temperature of $2.8~\\mathrm{keV}$, or by a blackbody plus bremsstrahlung model with $kT_\\mathrm{bb}\\sim 1.5$ keV, $kT_\\mathrm{brems} \\sim 4.5$ keV. In this case the blackbody accounts for about 20\\% of the total emission: as the authors point out, however, this model is not physically realistic given the enormous emitting volumes required for the bremsstrahlung emission.  These low-resolution spectra differ noticeably from the  X-ray spectra  of other Z-sources, which are usually described in terms of a two-component model. The spectral models of Z-sources usually follow one of two main paradigms: the Eastern model (Mitsuda et al. 1989), in which the spectrum of the source is interpreted as a sum of an emission coming directly from the compact object (Comptonized or not) plus a blackbody emission from the disc, and the Western model (White et al. 1986), where the emission is due to a hot blackbody coming from the central source, plus a Comptonized component that is due to the emission from an accretion disc corona surrounding the disc (no direct thermal emission from the disc is seen as photons emitted there are reprocessed and Comptonized in the corona).   Broad emission  lines (FWHM  up to  $\\sim 1$ keV)  at energies  in the range 6.4 -- 6.7 keV have been observed in the spectra of all galactic Z-sources, with the noticeable exception of GX~5-1 (Asai et al. 1994). These lines are identified with  the K$\\alpha$ radiative transitions of iron at different  ionization states. Sometimes an iron  absorption edge at energies  $\\sim 8$ keV  has been detected (see e.g. Di Salvo et al. 2001).  In this paper we analyze an archival XMM/Newton observation of LMC X--2, which gives us the opportunity to perform the first high resolution spectral study available to date on this source. ",
        "conclusions": "We have analyzed an archival XMM/Newton observation of LMC X--2, which allowed us to redetermine the position of the source with respect to the value reported in the Rosat All Sky Survey catalog. The source was probably in the normal or in the flaring branch, as its timing characteristics and high luminosity (exceeding the Eddington limit) exclude that the source is in the horizontal branch. Analyzing the spectra of the 3 active instruments we found that they could be well fitted by a two component model (blackbody plus disk blackbody) plus an emission line at 653 eV, which is identified as an O VIII Ly-$\\alpha$ emission line. On the contrary, the iron region lacks any emission line. As discussed, this can be caused by the low metallicity of the stars in the LMC."
    },
    "0710/0710.4872_arXiv.txt": {
        "abstract": "We have observed 13 methanol maser sources associated with massive star-forming regions; W3(OH), Mon~R2, S~255, W~33A, IRAS~18151$-$1208, G~24.78$+$0.08, G~29.95$-$0.02, IRAS~18556$+$0136, W~48, OH~43.8$-$0.1, ON~1, Cep~A and NGC~7538 at 6.7~GHz using the Japanese VLBI Network (JVN). Twelve of the thirteen sources were detected at our longest baseline of $\\sim $50~M$\\lambda $, and their images are presented. Seven of them are the first VLBI images at 6.7~GHz. This high detection rate and the small fringe spacing of $\\sim$4~milli-arcsecond suggest that most of the methanol maser sources have compact structure. Given this compactness as well as the known properties of long-life and small internal-motion, this methanol maser line is suitable for astrometry with VLBI. ",
        "introduction": "\\label{section:introduction} The class II methanol masers are well-known as tracers of early stages of high-mass star formation (\\cite{1998MNRAS.301..640W}; \\cite{2001A&A...369..278M}; \\cite{2006ApJ...638..241E}). Classes I and II are defined on the basis of associated sources (\\cite{1987Natur.326...49B}; \\cite{1991ASPC...16..119M}); different pumping mechanisms are proposed for them (\\cite{1992MNRAS.259..203C}, \\cite{1997MNRAS.288L..39S}). The class II masers are represented by 6.7 and 12.2~GHz lines (\\cite{1991ApJ...380L..75M}; \\cite{1987Natur.326...49B}). The spot size of class II masers is several AU (\\cite{1992ApJ...401L..39M}; \\cite{1999ApJ...519..244M}), while the class I masers are resolved out with Very Long Baseline Interferometric (VLBI) technique \\citep{1998AAS...193.7101L}.  VLBI observations for astrometry with hydroxyl, water and methanol masers have been made using the phase referencing technique. The accuracy of measuring with hydroxyl masers is $\\sim$1~milli-arcsecond (mas) \\citep{2003A&A...407..213V}, while that achieved with water masers is up to a few tens of micro-arcsecond ($\\mu $as) \\citep{2006ApJ...645..337H}. \\citet{2006Sci...311...54X} also have achieved the positional accuracy of $\\sim $10~$\\mu $as with methanol maser at 12.2~GHz. This observation has been the only one astrometric observation with methanol masers. Internal proper motions are often measured typically at a few mas per year for water masers at 22~GHz (e.g., \\cite{1981ApJ...247.1039G}, \\yearcite{1981ApJ...244..884G}) and this motion sometimes prevents the separation of the annual parallax from the internal proper motion. The internal proper motion for class II methanol masers are small and have been measured only for W3(OH) at 12.2~GHz \\citep{2002ApJ...564..813M}. The lifetime of 22~GHz water masers is sometimes too short for measuring annual parallax. For example, several spots disappear over timescales of 1 month for Cepheus~A and IRAS~21391$+$5802 (\\cite{2001ApJ...560..853T}, \\yearcite{2001Natur.411..277T}; \\cite{2000ApJ...538..268P}). A monitoring of variability of 6.7~GHz methanol line for four years showed that each spectral feature survives regardless of some variability \\citep{2004MNRAS.355..553G}, namely the lifetime of this maser is usually long enough for measuring annual parallax. There are 519 sites of 6.7~GHz methanol maser in a list compiled by \\citet{2005A&A...432..737P}. Recently, new 48 sources have been detected using the 305~m Arecibo radio telescope \\citep{2007ApJ...656..255P}. For the above reasons, i.e., compactness, small internal proper motion, long-life and large number of known sources, class II methanol masers may also be a useful probe for astrometry.  The VLBI Exploration of Radio Astrometry (VERA; \\cite{Kobayashi_etal.2003}) is a dedicated VLBI network for astrometry, which mainly observes galactic water masers for measuring their distance and motion. Methanol masers at 6.7~GHz would also contribute to revealing the Galactic structure as well as water masers. The spot size of methanol maser should be small enough for high precision astrometry. The study of spot size, however, have been made only a few cases so far. The number of sources imaged with baseline of $\\geq $50~M$\\lambda$ at 6.7~GHz were only four (\\cite{1992ApJ...401L..39M}; \\cite{2005A&A...442L..61B}; \\cite{2006A&A...448L..57P}; \\cite{2007A&A...461.1027G}). \\citet{2002A&A...383..614M} discussed the size and structure of individual masing regions in detail based on their 12.2 and 6.7~GHz observations. They showed that the majority of the masing regions consist of a compact maser core surrounded by extended emission (halo), and derived the core size of 2 to 20~AU.  We have started investigations of methanol masers using the Japanese VLBI Network (JVN) for astrometry. The network is a newly-established one with 50--2560~km baselines across the Japanese islands \\citep{2006astro.ph.12528D} and consists of ten antennas, including four radio telescopes of the VERA. We have installed 6.7~GHz receivers on three telescopes of the JVN and have made a snap-shot imaging survey toward thirteen sources associated with massive star-forming regions. The aims of this observation are to investigate the detectability of 6.7~GHz masers with our longest baseline of $\\sim$50~M$\\lambda$ corresponding fringe spacing of $\\sim$4~mas, and to verify the imaging capability of this network.  In this paper, we describe the detail of this observation and data reduction in section \\ref{section:observation}. In section \\ref{section:result} we present the images of the detected sources and report the results on each individual source. Finally, we discuss the properties of this maser for astrometry in section \\ref{section:discussion}, based on the results of this observation.     \\begin{table*} \\caption{The sample of 6.7~GHz methanol masers}\\label{tab:table1} \\begin{center} \\begin{tabular}{lllcrcc} \\hline\\hline Source & \\multicolumn{2}{c}{Coordinates(J2000)}                & Ref. & \\multicolumn{1}{c}{$S_{\\mathrm{p}}$} &   $d$ & \\multicolumn{1}{c}{VLBI obs.} \\\\ & \\multicolumn{1}{c}{RA}      & \\multicolumn{1}{c}{Dec} & &                                      & & \\multicolumn{1}{c}{ }         \\\\ & \\multicolumn{1}{c}{(h m s)} & \\multicolumn{1}{c}{($ ^{\\circ} $ $ ^{\\prime} $ $ ^{\\prime\\prime} $)} &         & \\multicolumn{1}{c}{(Jy)}  & (kpc) & \\\\\\hline W3(OH)              & 02 27 03.820 & \\hspace{1.5mm} 61 52 25.40  & 10 & 3294  & \\hspace{0.8mm} 1.95 &    1   \\\\ Mon~R2              & 06 07 47.867 &             $-$06 22 56.89  &  4 &  104  & \\hspace{0.8mm} 0.83 & 6, 7   \\\\ S~255               & 06 12 54.024 & \\hspace{1.5mm} 17 59 23.01  &  6 &   79  & 2.5  & 6, 7   \\\\ W~33A               & 18 14 39.52  &             $-$17 51 59.7   & 13 &  297  & 4.0  &  \\ldots \\\\ IRAS~18151$-$1208   & 18 17 58.07  &             $-$12 07 27.2   & 13 &  119  & 3.0  &  8   \\\\ G~24.78$+$0.08 & 18 36 12.57 & $-$07 12 11.4   & \\footnotemark[$*$] &   84  & 7.7  & \\ldots \\\\ G~29.95$-$0.02      & 18 46 03.741 &             $-$02 39 21.43  &  6 &  182  & 9.0  & \\ldots \\\\ IRAS~18556$+$0136   & 18 58 13.1   & \\hspace{1.5mm} 01 40 35     & 12 &  191  & 2.0  & \\ldots \\\\ W~48                & 19 01 45.5   & \\hspace{1.5mm} 01 13 28     &  3 &  733  & 3.4  &   6   \\\\ OH~43.8$-$0.1       & 19 11 53.987 & \\hspace{1.5mm} 09 35 50.308 &  2 &   51  & 2.8  & \\ldots \\\\ ON~1                & 20 10 09.1   & \\hspace{1.5mm} 31 31 34     & 12 &  107  & 1.8  & \\ldots \\\\ Cep~A               & 22 56 17.903 & \\hspace{1.5mm} 62 01 49.65  & 13 &  371  & \\hspace{0.8mm} 0.73 & \\ldots \\\\ NGC~7538            & 23 13 45.364 & \\hspace{1.5mm} 61 28 10.55  &  6 &  256  & 2.8  &  5, 6, 7, 9, 11 \\\\\\hline \\multicolumn{7}{l}{\\hbox to 0pt{\\parbox{136mm}{\\footnotesize Col~(1)~source name; Col~(2)--(3)~coordinates in J2000; Col~(4)~reference of coordinates; Col~(5)~peak flux densities from the single-dish observations by Yamaguchi 32~m; Col~(6)~source distance; Col~(7)~Published VLBI observations at 6.7~GHz.  References --- (1)~\\cite{1992ApJ...401L..39M}; (2)~\\cite{1994ApJS...91..659K}; (3)~\\cite{1995MNRAS.272...96C}; (4)~\\cite{1998MNRAS.301..640W}; (5)~\\cite{1998A&A...336L...5M}; (6)~\\cite{2000A&A...362.1093M}; (7)~\\cite{2001A&A...369..278M}; (8)~\\cite{2002evn..conf..213V}; (9)~\\cite{2004ApJ...603L.113P}; (10)~\\cite{2005MNRAS.360.1162E}; (11)~\\cite{2006A&A...448L..57P}; (12)~\\cite{1994yCat.2125....0J}; (13)~Fringe rate mapping in our observations (accuracy of position is from 100 to 300~mas).  \\footnotemark[$*$] The coordinate of this source is not used for this data reduction. This coordinate is obtained from the observations made in the following year. It is coincident with the coordinate in the catalog listed by \\citet{2005A&A...432..737P}.}\\hss}} \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "\\label{section:discussion} We have presented maps of twelve sources of methanol maser emission, and the seven of them are the first VLBI results at 6.7~GHz. The spatial distributions of maser spots show various morphology, such as linear (Cep~A, NGC~7538), ring like structure (W~48), largely separated clusters (W3(OH), W~33A, G~24.78$+$0.08, G~29.95$-$0.02, IRAS~18556$+$0136, OH~43.8$-$0.1, ON~1, Cep~A).  We discuss the properties of methanol maser at 6.7~GHz as a possible probe for astrometry comparing with the methanol maser at 12.2~GHz and the water maser at 22~GHz. We have detected twelve out of thirteen methanol masers at 6.7~GHz with the longest baseline of 50~M$\\lambda $ of our array. The integrated flux density of the correlated spectra account for $\\sim$50~{\\%} of the total flux, and some of which account for more than 90~{\\%}. This high detection rate, flux recovery, and the small fringe spacing of 4~mas suggest that most of the methanol maser emission have compact structure. This result is consistent with a previous study for W3(OH) by \\citet{1992ApJ...401L..39M}. We also showed that the correlated flux density at 50~M$\\lambda$ is typically 20~{\\%} of the total flux density. The size of maser spot inferred from the flux ratio varies from 2 to 30~AU (at a distance from 0.73 to 9~kpc of the sources). This is consistent with the core size of 2 to 20~AU obtained by \\citet{2002A&A...383..614M}. The velocity range of typically 10~km~s$^{-1}$ \\citep{1995MNRAS.272...96C} is narrow compared to that of water masers. Given this compactness and narrow velocity range as well as the known properties of long-life and small internal-motion, this methanol maser line is suitable for astrometry with VLBI.  From the viewpoint of observation, atomospheric fluctuation is a significant problem for observations at higher frequency. Strong absorption by water vapor in the atomosphere makes observations relatively difficult at 22~GHz in summer. This would be a potential problem in measurement of annual parallax. This is not the case for 6.7~GHz observation.  The astrometric VLBI observation uses continuum reference sources which are usually distant quasars. Such continuum sources show power-law, decreasing spectra, and the flux density is larger at 6.7~GHz than that of 12.2~GHz or 22~GHz. This property makes the astrometric observations easier in terms of detectability of reference sources.  The 6.7~GHz line might have some disadvantages due to its lower frequency in comparison with the 12.2~GHz line: The interstellar scintillation broadens the size of maser and reference sources. The interstellar broadening is stronger at lower frequency, and might affect to measure the precise position of the sources. Also the ionospheric density fluctuation changes path-length, consequently affects phase measurement at lower frequency. These effects depend on frequency as $\\nu^{-2}$, i.e., affect 3.3 times larger for 6.7~GHz than for 12.2~GHz. Although \\citet{2002A&A...383..614M} discussed that the interstellar broadening is not significant at a scale larger that 1~mas, we have to take account of these effects for astrometry.  We have made a phase-referencing VLBI observation with the JVN and achieved that the positional accuracy of $\\sim$50~$\\mu $as at 8.4~GHz and the image dynamic range of $\\sim$50 on a target \\citep{2006PASJ...58..777D}. The Bigradient Phase Referencing (BPR) was used for this observation, and a reference calibrator was separated by a \\timeform{2D.1} from the target source. The following equation (\\ref{equ:equ1}) can be used to derive an accuracy $\\Delta \\pi$ of annual parallax, \\begin{eqnarray} \\Delta \\pi = \\frac{\\theta_{\\mathrm{beam}}}{D \\cdot \\sqrt{N_{\\mathrm{spot}} \\cdot N_{\\mathrm{obs}}}} \\label{equ:equ1} \\end{eqnarray} where $\\theta_{\\mathrm{beam}}$ is the minimum fringe spacing, $D$ is Dynamic range of the image, $N_{\\mathrm{spot}}$ is the number of spots used in measuring annual parallax, and $N_{\\mathrm{obs}}$ is the number of observations. If we make a series of observations for nine methanol maser spots for five epochs as in the cases of the observation by \\citet{2006Sci...311...54X}, it is expected that an accuracy of $\\sim$12~$\\mu $as in annual parallax would be achieved.  The flux density at 6.7~GHz is typically $\\sim $10 times larger than that of 12.2~GHz \\citep{1995MNRAS.274.1126C}, and the number of observable sources of 6.7~GHz is much larger than that of 12.2~GHz. This practical reason let us choose 6.7~GHz line than 12.2~GHz one.  We have started to improve the JVN at 6.7~GHz in terms of sensitivity and the number of stations to detect weak masers and reference sources. Usuda 64~m is one of the newly participating telescopes. Assuming an aperture efficiency of 50~{\\%} and $T_{\\mathrm{sys}}$ of 50~K for all telescopes, it is expected that sensitivity for fringe detection (7~$\\sigma $) would be less than 5~Jy. We showed that the correlated flux density at 50~M$\\lambda $ is typically 20~{\\%} of the total flux density. Hence, sources with total flux density of larger than 25~Jy would be potential targets for astrometric observation. The number of such sources found in the catalog by \\citet{2005A&A...432..737P} is larger than 150.    \\bigskip   The authors wish to thank the JVN team for observing assistance and support. The JVN project is led by the National Astronomical Observatory of Japan~(NAOJ) that is a branch of the National Institutes of Natural Sciences~(NINS), Hokkaido University, Gifu University, Yamaguchi University, and Kagoshima University, in cooperation with Geographical Survey Institute~(GSI), the Japan Aerospace Exploration Agency~(JAXA), and the National Institute of Information and Communications Technology~(NICT)."
    },
    "0710/0710.1666_arXiv.txt": {
        "abstract": "Although supernova explosions and stellar winds happens at scales bellow 100 pc, they affect the interstellar medium(ISM) and galaxy formation. We use cosmological N-body+Hydrodynamics simulations of galaxy formation, as well as simulations of the ISM to study the effect of stellar feedback on galactic scales. Stellar feedback maintains gas with temperatures above a million degrees. This gas fills bubbles, super-bubbles and chimneys.  Our model of feedback, in which 10\\%-30\\% of the feedback energy is coming from runaway stars, reproduces this hot gas only if the resolution is better than 50 pc. This is 10 times better than the typical resolution in cosmological simulations of galaxy formation. Only with this resolution, the effect of stellar feedback in galaxy formation is resolved without any assumption about sub-resolution physics. Stellar feedback can regulate the formation of bulges and can shape the inner parts of the rotation curve. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.3349_arXiv.txt": {
        "abstract": "{\\small After a period of inflationary expansion, the Universe reheated and reached full thermal equilibrium at the reheating temperature $\\treh$. In this work we point out that, in the context of effective low-energy supersymmetric models, LHC measurements may allow one to determine $\\treh$ as a function of the mass of the dark matter particle assumed to be either an axino or a gravitino. An upper bound on their mass and on $\\treh$ may also be derived.} ",
        "introduction": "\\label{sect:intro}  Dark matter (DM) remains an unknown component of the Universe. While it has so far escaped detection, its existence has been convincingly inferred from gravitational effects that it imparts on visible matter through rotational curves of spiral galaxies, gravitational lensing, etc,~\\cite{susy-dm-reviews}. The effects of dark matter also can be seen on large structure formation and on anisotropy of the cosmic microwave background (CMB).  The CMB in particular provides a powerful tool for determining the global abundance of cold DM (CDM).  Recently, the Wilkinson Microwave Anisotropy Probe (WMAP) has performed a high-accuracy measurement of several cosmological parameters. In particular, the relic density of CDM has been determined to lie in the range~\\cite{Spergel:2006hy} \\beq \\abundcdm =0.104 \\pm 0.009. \\label{Oh2WMAP} \\eeq Since DM has to be electrically and (preferably) color-charge neutral, from the particle physics point of view, a natural candidate for DM is some weakly interacting massive particle (WIMP). Within standard cosmology, the WIMP is produced via a usual freeze-out mechanism from an expanding plasma.  Among specific, well-motivated particle candidates for the WIMP, the by far most popular choice is a stable lightest neutralino $\\chi$ of effective low-energy supersymmetry (SUSY) models. Most efforts have gone to exploring the lightest neutralino $\\chi$ of the Minimal Supersymmetric Standard Model (MSSM) as the lightest supersymmetric particle (LSP) that, in the presence of R-parity, is stable and makes up all, or most of, the CDM in the Universe.  In addition to an impressive experimental effort of direct and indirect searches for cosmic WIMPs, the Large Hadron Collider (LHC) will soon start exploring the TeV energy scale and is expected to find several superpartners and to determine their properties. In particular, some authors have explored the feasibility of determining the neutralino's relic abundance $\\abundchi$ from LHC measurements~\\cite{oh2atlhc-cmssm,oh2atlhc-mssm}.  Their conclusion was that, under favorable circumstances, this should be possible with rather good accuracy, of order 10\\% or better, although this may be challenging~\\cite{bk05}. An analogous study has also been done in the context of the Linear Collider, where accuracy of a similar determination would be much better~\\cite{oh2atilc}.  If a WIMP signal is detected in one or more DM detection experiments, and also at the LHC a large missing-mass and missing-energy signature, characteristic of the stable neutralino, is measured and implies a similar mass, and if perhaps additionally eventually its relic abundance is determined from LHC data, even if with limited precision, and agrees with the ``WMAP range''~(\\ref{Oh2WMAP}), then the DM problem will most likely be declared solved, and for a good reason.  However, such an optimistic outcome is by no means guaranteed. One realistic possibility is that the neutralino, even if it is found as an apparently stable state in LHC detectors, may not be the true LSP and therefore DM in the Universe. Instead, it could decay in the early Universe into an even lighter, and (possibly much) more weakly interacting, state, the real LSP, outside the MSSM (or some other low-energy SUSY model) spectrum. In this case, current cosmic WIMP searches will prove futile, even after improving the upper limit on the spin-independent interaction cross section on a free proton $\\sigsip$ from the current sensitivity of $\\sim10^{-7}\\pb$~\\cite{silimit-07} down to $\\sim10^{-10}\\pb$, which is as far down as experiment can probably go given background from natural radioactivity.  Moreover, the neutralino relic abundance, as determined at the LHC, may come out convincingly outside the range~(\\ref{Oh2WMAP}).  In fact, a value of $\\abundchi$ above about 0.1 can be easily explained in terms of a lighter LSP into which the neutralino, after its freeze-out, decayed in the early Universe~\\cite{ckr}. This is because the relic abundance of the true LSP is in this case related to $\\abundchi$ by the ratio of the LSP mass $\\mlsp$ to $\\mchi$, which is less than one. If $\\abundchi$ comes out below~(\\ref{Oh2WMAP}), several solutions have been suggested which invoke non-standard cosmology, e.g. quintessence-driven kination~\\cite{lowoh2-quint}, while preserving the neutralino as the DM in the Universe.  However, if at the same time DM searches bring null results, this will provide a strong indication against the neutralino nature of DM.  In contrast to the above attempts, we will consider a whole range of possible values of $\\abundchi$ at the LHC, both below and above 0.1. We will work instead within the framework of standard cosmology but  will not assume the neutralino to be the DM in the Universe.  In this context we point an intriguing possibility of determining at the LHC the reheating temperature of the Universe. The framework we consider will therefore give us an opportunity to probe some crucial features of the early period of the Universe's history.  The reheating temperature $\\treh$ is normally thought of as the temperature at which, after a period of rapid inflationary expansion, the Universe reheated (or, more properly, defrosted), and the expanding plasma reached full thermal equilibrium. Determining $\\treh$ cannot be done with the neutralino as DM since it freezes out at $\\tf\\simeq \\mchi/24$, which is normally thought to be much below $\\treh$.\\footnote{The possibility of $\\treh\\lsim\\tf$ has been explored in~\\cite{gkr00} and more recently in~\\cite{ggetal06}. In this case one could think of determining experimentally $\\treh$ even with the neutralino as DM.}  Here instead we assume a different candidate for the LSP (assuming R-parity) and cold DM, whose relic density depends, at least in part, on $\\treh$. This is the case for either an axino or a gravitino.  The spin-$1/2$ axino (the fermionic superpartner of an axion) and the spin-$3/2$ gravitino (the fermionic superpartner of a graviton) are both well-motivated. The former arises in SUSY extensions of models incorporating the Peccei-Quinn solution to the strong CP problem. The latter is an inherent ingredient of the particle spectrum of supergravity models. Both, like the axion, form a subclass of extremely weakly interacting massive particles (E-WIMPs)~\\cite{Choi:2005vq}.  The characteristic strength of their interactions with ordinary matter is strongly suppressed by a large mass scale, the Peccei-Quinn scale $\\fa\\sim 10^{11}\\gev$ in the case of axinos and the (reduced) Planck scale $\\mplanck\\simeq 2.4\\times 10^{18} \\gev$ for gravitinos. The mass of the axino $\\maxino$ is strongly model-dependent and can take values ranging from keV up to TeV~\\cite{ckn}. The mass of the gravitino $\\mgravitino$ in gravity-mediated SUSY breaking schemes is given by $\\mgravitino\\sim{\\ms^2/\\mplanck}$, where $\\ms\\sim10^{11}\\gev$ is the scale of local SUSY-breaking in the hidden sector, and is expected in the GeV to TeV regime. In other schemes of SUSY breaking $\\mgravitino$ can be (much) smaller. In this work we want to remain as model-independent as possible and will treat $\\maxino$ and $\\mgravitino$ as free parameters.  Both axinos and gravitinos can be produced in decays of the neutralino (or another ordinary superpartner, e.g. the stau) after freeze-out, as mentioned above, or in thermal scatterings and decay processes of ordinary particles and sparticles in hot plasma at high enough reheating temperatures $\\treh$ - hence their relic density dependence on $\\treh$. The possibility of axinos in a KSVZ axion framework~\\cite{ksvz} as cold DM was pointed out in~\\cite{ckr,ckkr} and next studied in several papers~\\cite{crs02,crrs04,bs04,rs06}, while axinos as warm DM was considered in~\\cite{rtw}. The gravitino as a cosmological relic was extensively studied in the literature, starting from~\\cite{gravitino-early,ekn84}, more recent papers on gravitino CDM include~\\cite{mmy93,bbp98,bbb00,fengetal,rrc04,ccjrr}.   For definiteness, in this work we will first assume the lightest neutralino to be the lightest ordinary superpartner and the next-to-lightest particle (NLSP), although below we will also consider the case of the stau. Our main result is that, assuming the axino or gravitino as the true LSP and CDM, and that the (apparently stable) neutralino is discovered at the LHC and its ``relic abundance'' $\\abundchi$ is determined from LHC data, then one should be able to determine the reheating temperature in the Universe as a function of the LSP mass, or at least place an {\\em upper bound} on it, as we show below. In the regime where thermal production dominates, in the axino case we find $\\treh\\propto \\fa^2/\\maxino$ while in the gravitino case $\\treh\\propto \\mgravitino$. Alternatively, in the non thermal production dominated regime, we find an {\\em upper bound} on the allowed mass range of a CDM particle. Furthermore, $\\abundchi$ at the LHC can be expected to come out either below or above (or for that matter even accidentally agree with) the ``WMAP range''~(\\ref{Oh2WMAP}).  The same holds true also in the even more striking case of the lighter stau taking the role of the NLSP instead. In this case the very discovery of an (apparently stable) charged massive particle at the LHC will immediately imply that DM is made up of some state outside the usual spectrum of low-energy superpartners. In this case, dedicated studies of the differential photon spectrum may allow one to distinguish between the axino and the gravitino LSP~\\cite{bchrr05}.  We stress that, a detection of a (seemingly) stable neutralino at the LHC will not be sufficient to prove that the neutralino is the LSP and the CDM.\\footnote{Note that, for the neutralino, or any other superpartner, to appear stable in an LHC detector, it is sufficient that its lifetime is longer than a microsecond or so. Of course, if it is unstable, then it will be immediately clear that the neutralino is not cosmologically relevant.}  Even establishing an apparent agreement of $\\abundchi$, as to be determined at the LHC, with the range~(\\ref{Oh2WMAP}) will not necessarily imply the neutralino nature of DM, although admittedly it will be very persuasive.  For this, a signal in DM searches must also be detected and the resulting WIMP mass must be consistent with LHC measurements of the neutralino mass.  The paper is organized as follows. In section~\\ref{production}, we review the thermal and non-thermal production of axinos and gravitinos and next explain our strategy for determining $\\treh$ at the LHC.  In sections~\\ref{axino} and~\\ref{gravitino}, we discuss the reheating temperature determination with axino LSP and gravitino LSP respectively, assuming the neutralino to be the NLSP. In section~\\ref{staunlsp} we consider instead the lighter stau as the NLSP. We make final remarks and summarize our findings in section~\\ref{summary}. ",
        "conclusions": "\\label{summary}  If the axino or the gravitino E-WIMP is the lightest superpartner and the dominant component of CDM in the Universe, then, under favorable conditions, it will be possible to determine the reheating temperature after inflation from LHC measurements. Basically, one will need to measure the neutralino mass and other parameters determining its ``relic abundance'' - a non-trivial and highly exciting possibility in itself. The quantity may be found to be different from the ``WMAP range''~(\\ref{Oh2WMAP}), thus suggesting a non-standard CDM candidate. Additionally, one will need to know the mass of the gluino, either through a direct measurement, or via the gaugino unification mass relation, since it plays an important role in E-WIMP production at high $\\treh$.  The neutralino is likely to be sufficiently long-lived to appear stable in LHC detectors, but in the early Universe it would have decayed into the axino or the gravitino LSP. Since each of them is extremely weakly interacting, it would be hopeless to look for them in usual WIMP dark matter searches.  In the axino LSP case, if thermal production dominates, a consistency of its relic abundance with the ``WMAP range''~(\\ref{Oh2WMAP}) implies $\\treh\\propto \\fa^2/\\maxino$. If non-thermal production is dominant instead, then, for a known value of $\\mchi$, the value of $\\abundchi$ will imply an upper bound on $\\maxino/\\fa^2$, which in turn will imply a lower bound on $\\treh$.  For the gravitino LSP, in the thermal production dominated case the analogous dependence is $\\treh\\propto \\mgravitino$. If non-thermal production dominates then the value of $\\abundchi$ will imply an upper bound on $\\mgravitino$. For both axino and gravitino one can express this as \\beq \\mlspmax=\\min\\{\\mnlsp, \\frac{\\abundcdm}{\\abundnlsp}\\mnlsp\\}, \\label{eq:mlspmax} \\eeq which is independent of $\\treh$.  For both relics, one can also derive an upper bound on $\\treh$, although in the axino case has to further assume that it is cold DM. On the other hand, for both axino and gravitino, it will be difficult to distinguish between the TP and NTP regimes but some information may be helpful. Typically, when $\\mchi$ is large and $\\abundnlsp$ is also large then NTP is dominant, while if $\\abundnlsp$ is small then TP typically dominates.  Similar conclusions as for the neutralino (for both E-WIMPs) apply if the NLSP is the stau, instead. Additionally, the very discovery of a charged, massive and (seemingly) stable state at the LHC will immediately imply that dark matter in the Universe is not made up of the usual neutralino.  With so many similarities between both E-WIMPs, a natural question arises whether it would be possible to distinguish them in a collider experiment. Unfortunately, with neutralino NLSP this is unlikely to be possible. On the other hand, with the stau NLSP being electrically charged, enough of them could possibly be accumulated. By comparing their 3-body decays, it may be possible to only tell whether Nature has selected the axino or the gravitino as the lightest superpartner, but to actually also determine their mass with some reasonable precision~\\cite{Buchmuller:2004rq,bchrr05}. Needless to say, this would allow one to determine also the reheating temperature.  In this work, we focused on demonstrating the principle of determining $\\treh$ at the LHC, and did not concern ourselves with uncertainties of experimental measurements. These are likely to be substantial at the LHC, the case we have mostly focused on here, unless one considers specific models~\\cite{bk05}. A more detailed study will be required to assess their impact. An analogous study can be performed in the context of the planned Linear Collider. We also implicitly assumed that only one relic dominates the CDM component of the Universe, which in principle does not have to be the case. For example, the axion and the axino could co-exist and contribute comparable fractions to the CDM relic density. (Note that this could be the case also in the ``standard'' case of the stable neutralino.) This would introduce an additional uncertainty and would lead to replacing the derived values of $\\treh$ with an upper bound on the quantity.   \\medskip {\\bf Acknowledgements} \\\\ This work is supported by an STFC grant. K.-Y.C. has been partially supported by a PPARC grant, by the Ministerio de Educacion y Ciencia of Spain under Proyecto Nacional FPA2006-05423 and by the Comunidad de Madrid under Proyecto HEPHACOS, Ayudas de I+D S-0505/ESP-0346. L.R is partially supported by the EC 6th Framework Programmes MRTN-CT-2004-503369 and MRTN-CT-2006-035505. R.RdA is supported by the program ``Juan de la Cierva'' of the Ministerio de Educaci\\'{o}n y Ciencia of Spain.  The author(s) would like to thank CERN for hospitality and support and the European Network of Theoretical Astroparticle Physics ENTApP ILIAS/N6 under contract number RII3-CT-2004-506222 for support. This project benefited from the CERN-ENTApP joint visitor's programme on dark matter, 5-9 March 2007 and the CERN Theory Institute ``LHC-Cosmology Interplay'', 25 June - 10 Aug 2007.  \\newcommand\\jcap[3] % {{\\it J.\\ Cosmol.\\ Astrop.\\ Phys.\\ }{\\bf #1} (#2) #3} \\newcommand\\mnras[3]   { % {{\\it Mon.\\ Not.\\ R.\\ Astron.\\ Soc.\\ }{\\bf #1} (#2) #3}} \\newcommand\\apjs[3]   { {{\\it Astrophys.\\ J.\\ Supp.\\ }{\\bf #1} (#2) #3}} \\newcommand\\aipcp[3]   { {{\\it AIP Conf.\\ Proc.\\ }{\\bf #1} (#2) #3}}"
    },
    "0710/0710.5759_arXiv.txt": {
        "abstract": "\\noindent  We present low resolution UV-blue spectroscopic observations of red giant stars in the globular cluster M53 ([Fe/H]$=-1.84$), obtained to study primordial abundance variations and deep mixing via the CN and CH absorption bands.  The metallicity of M53 makes it an attractive target: a bimodal distribution of 3883 \\AA\\ CN bandstrength is common in moderate- and high-metallicity globular clusters ([Fe/H] $\\geq -1.6$) but unusual in those of lower metallicity ([Fe/H] $\\leq -2.0$).  We find that M53 is an intermediate case, and has a broad but not strongly bimodal distribution of CN bandstrength, with CN and CH bandstrengths anticorrelated in the less-evolved stars.  Like many other globular clusters, M53 also exhibits a general decline in CH bandstrength and [C/Fe] abundance with rising luminosity on the red giant branch. ",
        "introduction": "Star-to-star variations of the strength of the 3883 \\AA\\ CN band are nearly universal in moderate-metallicity ([Fe/H] $\\geq -1.6$) globular clusters of the Milky Way, but unusual in low-metallicity globular clusters ([Fe/H] $\\leq -2.0$) (see Gratton, Sneden \\& Carretta 2004 for a thorough review).  Typically a CN bandstrength is measured as a magnitude difference between the flux within the 3883 \\AA\\ CN band and the flux in a nearby comparison band chosen to be free of CN absorption.  The index $S(3839)$ introduced by Norris et al. (1981) is one such measure of CN band strength.  Not only does $S(3839)$ vary from star to star within individual globular clusters, its distribution is often bimodal in moderate-metallicity ([Fe/H] $\\geq -1.6$) clusters.  It is often anticorrelated with the strength of CH absorption in the G band at 4300 \\AA\\ among giants of similar absolute magnitude.  This anticorrelation indicates that CN-strong stars are depleted in carbon.  In order to have strong CN bands with low [C/Fe] abundances, the CN-strong stars must also be enhanced in nitrogen.  This abundance pattern has directed investigations into the origin of the bimodality toward stellar sources, where hydrogen shell burning can produce abundance enhancements and depletions such as those observed in globular clusters in C and N (and also O, Na, Mg, and Al).  Meanwhile, few CN-strong giants are known in the halo field (see, e.g., Pilachowski et al. 1996; Shetrone et al. 1999), indicating that the enrichment process responsible for CN bimodality only operated in cluster (or proto-cluster) environments.  This paper considers whether the apparent lack of CN bimodality in low-metallicity globular clusters is a true lack of abundance variation or simply an effect of low overall metal abundance.  Previous investigations into the behavior of CN abundance at low metallicity have been inconclusive as to whether large variations in bandstrength can be present.  Trefzger et al. (1983) and Shetrone et al. (1999) report several CN-strong red giants with large nitrogen abundances in M15 ([Fe/H]$=-2.26$), and Carbon et al. (1982) found variations of [C/Fe] and [N/Fe] but no strong 3883 \\AA\\ CN bands among the red giant stars of M92 ([Fe/H]$=-2.28$).  Briley et al. (1994) showed that the carbon and nitrogen abundances measured within M92 would result in a bimodal distribution of $S(3839)$ if the overall metallicity were scaled up to [Fe/H]$=-1.0$ dex.  Smith \\& Norris (1982) and Briley et al. (1993) both studied M55 ([Fe/H]$=-1.81$, Harris 1996), and found that it is quite homogeneous in CN bandstrength.  That lack of CN-strong stars could be the result of low overall metallicity inhibiting the formation of CN molecules, low [C/Fe] among the brighter red giants, or a sign that whatever enrichment process caused CN variations in moderate-metallicity globular clusters did not operate at lower metallicities.  A study of additional globular clusters with metallicities similar to M55 could help decide between these alternatives.  We chose M53 as the subject of this study because it is similar in metallicity to M55 (based on our isochrone fits discussed below), and will hopefully provide insight into whether low-metallicity globular clusters contain CN-enhanced stars.  While our study does not comment directly on the process that causes abundance inhomogeneities in globular clusters, we note here that a lack of abundance variation in the lowest-metallicity clusters would provide a constraint on all hypotheses about that process.  To understand the CN and CH index behavior, we compare spectroscopy of the 3883 \\AA\\ CN band for red giants in M53 to the Norris et al. (1981) study of NGC 6752, a moderate-metallicity globular cluster ([Fe/H]$=-1.56$) that shows clear CN bimodality.  We also convert CH bandstrengths into [C/Fe] abundances using synthetic spectra and look at the relations between carbon abundance and CN bandstrengths.  An intrinsic spread is found in CN bandstrength, but it is smaller than the range observed in NGC 6752.  This result is interesting in the light of work by D'Antona et al. (2002), who use stellar evolution models to draw connections between helium abundance and horizontal branch morphology.  They find that globular clusters with solar helium abundances have short, red horizontal branches, while extended blue horizontal branches result from the smaller core mass of helium-enhanced stars.  In their model, the helium enrichment comes from AGB feedback, implying that globular clusters with red horizontal branches, like M53, should exhibit only minor CNO-product enrichment.  \\begin{figure} \\plotone{f1.eps} \\caption{ Color-magnitude diagram for the red giant branch of M53.  The photometry is from Rey et al. (1998) for stars identified as red giants in Cuffey (1965).  Filled circles are for stars observed with Kast, open circles are stars we did not observe, crosses are for stars dismissed in $\\S$2 as nonmembers.} \\end{figure} ",
        "conclusions": ""
    },
    "0710/0710.5273_arXiv.txt": {
        "abstract": "{It is expected that Be stars are surrounded by circumstellar envelopes, and that a significant fraction have companions. Achernar ($\\alpha$\\,Eri) is the nearest Be star, and is thus a favourable target to search for their signatures using high resolution imaging.} {We aim at detecting circumstellar material or companions around Achernar at distances of a few tens of AU.} {We obtained diffraction-limited thermal IR images of Achernar using the BURST mode of the VLT/VISIR instrument.} {The images obtained in the PAH1 band show a point-like secondary source located 0.280\" north-west of Achernar. Its emission is 1.8\\% of the flux of Achernar in this band, but is not detected in the PAH2, SiC and NeII bands. } {The flux from the detected secondary source is compatible with a late A spectral type main sequence companion to Achernar. The position angle of this source (almost aligned with the equatorial plane of Achernar) and its projected linear separation (12.3\\,AU at the distance of Achernar) favor this interpretation.} ",
        "introduction": "The southern star \\object{Achernar} ($\\alpha$\\,Eridani, \\object{HD 10144}) is the brightest and nearest of all Be stars (V $=0.46$ mag). Depending on the author (and the technique used) its spectral type ranges from B3-B4IIIe to B4Ve (e.g., Slettebak~\\cite{slettebak82}, Balona et al.~\\cite{balona87}). The estimated projected rotation velocity $v \\sin i$ ranges from 220 to 270\\,km.s$^{-1}$ and the effective temperature $T_{\\rm eff}$ from $15\\,000$ to $20\\,000$\\,K (see e.g., Vinicius et al.~\\cite{vinicius06}, Rivinius \\textit{priv. comm.}, Chauville et al.~\\cite{chauville01}). It has been the subject of a renewed interest since its distorted photosphere was resolved by means of near-IR long-baseline interferometry (Domicano de Souza et al.~\\cite{domiciano03}). Further interferometric observations revealed the polar wind ejected from the hot polar caps of the star (Kervella \\& Domiciano~\\cite{kervella06}), due to the von Zeipel effect (von Zeipel~\\cite{vonzeipel24}).  Our observations aim at studying two aspects of the close environment of Achernar: the circumstellar envelope at distances of up to a few tens of AU, and binarity. Mid-infrared (hereafter MIR) imaging is prefectly suited for these two objectives. Firstly, this wavelength range corresponds to where a circumstellar envelope becomes optically thick (asuming the emission is caused by free-free radiation, following for instance Panagia \\& Felli~\\cite{panagia75}). Secondly, the contrast between Achernar and a cool companion will be significantly reduced.  We hereafter present the result of our imaging campaign with the VLT/VISIR instrument, to explore the environment of this star at angular distances of $\\approx 0.1$ to 10\". ",
        "conclusions": "We presented high resolution thermal infrared images of the close environment of Achernar. A point-like source is identified at an angular distance of 0.280\" from Achernar, contributing for 1.8\\% of the flux of Achernar in the PAH1 band. This source is not detected in our longer wavelength images, probably due to their insufficient sensitivity. Its location, close to Achernar and almost in the equatorial plane of the star, indicates that it is likely to be a physical companion orbiting Achernar. The measured flux corresponds to that of an A7V star (similar to Altair, for instance). A secure identification requires further observations, that will also determine if source B indeed shares the proper motion of Achernar ($\\simeq 97$\\,mas/yr)."
    },
    "0710/0710.0984_arXiv.txt": {
        "abstract": "{ The Alpha Magnetic Spectrometer (AMS), whose final version AMS-02 is to be installed on the International Space Station (ISS) for at least 3 years, is a detector designed to measure charged cosmic ray spectra with energies up to the TeV region and with high energy photon detection capability up to a few hundred GeV, using state-of-the-art particle identification techniques. Following the successful flight of the detector prototype (AMS-01) aboard the space shuttle, AMS-02 is expected to provide a significant improvement on the current knowledge of the elemental and isotopic composition of hadronic cosmic rays due to its long exposure time (minimum of 3 years) and large acceptance (0.5 m$^2$sr) which will enable it to collect a total statistics of more than $10^{10}$ nuclei. Detector capabilities for charge, velocity and mass identification, estimated from ion beam tests and detailed Monte Carlo simulations, are presented. Relevant issues in cosmic ray astrophysics addressed by AMS-02, including the test of cosmic ray propagation models, galactic confinement times and the influence of solar cycles on the local cosmic ray flux, are briefly discussed. }   \\FullConference{International Europhysics Conference on High Energy Physics\\\\ July 21st - 27th 2005\\\\ Lisboa, Portugal}  \\begin{document} ",
        "introduction": " ",
        "conclusions": "Data from AMS-02 will provide a major improvement with respect to existing results for the hadronic cosmic ray spectrum. A total statistics of more than $10^{10}$ events will be collected during its operation. Detector capabilities include charge separation up to $Z \\sim 30$, velocity reconstruction with $\\Delta \\beta / \\beta \\sim 10^{-3}$ for $Z=1$ and $\\Delta \\beta / \\beta \\sim 10^{-4}$ for $Z \\sim 10-20$, and isotopic separation of light elements. Results of AMS-02 will address key issues in cosmic ray astrophysics, namely, propagation of cosmic rays in the Galaxy, confinement times, and the effect of solar cycles on the flux in Earth's vicinity."
    },
    "0710/0710.5429_arXiv.txt": {
        "abstract": "We present $UBVI_C$ CCD photometry of the young open cluster Be 59 with the aim to study the star formation scenario in the cluster. The radial extent of the cluster is found to be $\\sim$ 10 arcmin (2.9 pc). The interstellar extinction in the cluster region varies between $E(B-V) \\simeq$ 1.4 to 1.8 mag. The ratio of total-to-selective extinction in the cluster region is estimated as $3.7\\pm0.3$. The distance of the cluster is found to be $1.00\\pm0.05$ kpc. Using near-infrared colours and slitless spectroscopy, we have identified young stellar objects (YSOs) in the open cluster Be 59 region. The ages of these YSOs range between $<1$ Myr to $\\sim$ 2 Myr, whereas the mean age of the massive stars in the cluster region is found to be $\\sim$ 2 Myr. There is evidence for second generation star formation outside the boundary of the cluster, which may be triggered by massive stars in the cluster. The slope of the initial mass function, $\\Gamma$, in the mass range $2.5 < M/M_\\odot \\le 28$ is found to be $-1.01\\pm0.11$ which is shallower than the Salpeter value (-1.35), whereas in the mass range $1.5 < M/M_\\odot \\le 2.5$ the slope is almost flat. The slope of the K-band luminosity function is estimated as $0.27\\pm0.02$, which is smaller than the average value ($\\sim$0.4) reported for young embedded clusters. Approximately $32\\%$ of H$\\alpha$ emission stars of Be 59 exhibit NIR excess indicating that inner disks of the T-Tauri star (TTS) population have not dissipated. The MSX and IRAS-HIRES images around the cluster region are also used to study the emission from unidentified infrared bands and to estimate the spatial distribution of optical depth of warm and cold interstellar dust. ",
        "introduction": "High-mass star forming regions have been known for many years as OB associations and HII regions and they  have been observed quite extensively on various aspects. However, the census of low mass stars in such regions has not been possible until recently. Recent advancement in detectors have permitted the detection of substantial population of low mass stars in OB associations. Recently relatively large number of low mass stars have been detected in a few OB associations (e.g. Upper Scorpios, the $\\sigma$ and $\\lambda$ Ori regions; Preibisch \\& Zinnecker 1999, Dolan \\& Mathieu 2002). Since this realization, surveys have demonstrated that the Initial Mass Functions (IMFs) are essentially the same in all star forming regions. The apparent difference may be due mainly to the  inherent low percentage of high mass stars and the incomplete survey of low mass stars in high-mass star forming regions (e.g. Kroupa 2002, Preibisch \\& Zinneker 1999, Hillenbrand 1997, Massey et al. 1995).  The massive stars of OB associations have strong influence and significantly affect the entire star-forming  region. Strong UV radiation from massive stars could evaporate nearby clouds and consequently terminate star formation. Herbig (1962) suggested that low and intermediate-mass stars form first and with the formation of most massive star in the region, the cloud gets disrupted and star formation ceases. Alternatively shock waves associated with the ionization front may squeeze molecular cloud and induce subsequent star formation. Elmegreen \\& Lada (1977) proposed the expanding ionization fronts play a constructive role to incite a sequence of star-formation activities in the neighborhood. The question however is: could both of these effects occur in different parts of the same star-forming regions?  Recent studies have shown cases where massive stars ionize adjacent clouds revealing embedded stars (e.g. the Eagle Nebula; Hester et al. 1996). Preibisch \\& Zinnecker (1999) show that star formation in the upper Scorpius OB association was likely triggered by a nearby supernovae. Recently Lee et al. (2005) have found evidence for triggered star formation in the bright-rimmed clouds (BRCs) in the vicinity of O stars. Pandey et al. (2001, 2005) found that star formation in few young clusters may be a continuous process. For example, In the case of NGC 663, star formation seems to have taken place non-coevally in the sense that formation of low mass stars precedes the formation of most massive stars. Whereas, in the case of NGC 654 and NGC 3603, formation of low mass stars did not cease after the formation of most massive stars in the clusters.   The Sharpless region S171 is a large HII region associated with the Cepheus OB4 stellar association (Yang \\& Fukui 1992). The star cluster Be 59 is located at the center of the nebulous region with nine  O7 - B3 stars. The physical properties of the ionized gas have been investigated in radio continuum and recombination lines. (e.g. Felli et al. 1977, Rossano et al. 1980, Harten et al. 1981). Yang \\& Fukui (1992) have mapped the molecular gas distribution using the $J=1-0$ lines of $^{12}$CO and $^{13}$CO emission, indicating that Be 59 generates the ionization front on the surface of two dense molecular clouds. They suggested that the dense gas is in contact with the continuum source and ionization front is deriving shocks into the clumps. As the shock penetrates into the molecular clumps, it forms a compressed gas layer which is unstable against self-gravity (Elmegreen 1989, 1998). One can expect that a new generation of stars may form from the compressed gas layer. Thus the S171 region provides an opportunity to make a comprehensive exploration of the effects of massive stars on low mass star formation. Therefore in this paper, we present a multi-wavelength study of the star forming region S171. We have also carried out slitless spectroscopy to identify pre-main sequence H$\\alpha$ emission stars in the region. In section 2 we present optical CCD photometric and slitless spectroscopic observations and brief description of data reduction. In sect 3, we discuss archival data set used in the present study. In the ensuing sections we discuss results and star formation scenario in the Be 59 region. ",
        "conclusions": ""
    },
    "0710/0710.5103_arXiv.txt": {
        "abstract": "{The dust component of the interstellar medium (ISM) has been extensively studied in the past decades. Late-type stars have been assumed as the main source of dust to the ISM, but recent observations show that supernova remnants may play a role on the ISM dust feedback.} {In this work, we study the importance of low and high mass stars, as well as their evolutionary phase, on the ISM dust feedback process. We also determine the changes on the obtained results considering different mass distribution functions and star formation history.} {We describe a semi-empirical calculation of the relative importance of each star at each evolutionary phase in the dust ejection to the ISM. We compare the obtained results for two stellar mass distribution functions, the classic Salpeter initial mass function and the present day mass function. We used the evolutionary track models for each stellar mass, and the empirical mass-loss rates and dust-to-gas ratio. } {We show that the relative contribution of each stellar mass depends on the used distribution. Ejecta from massive stars represent the most important objects for the ISM dust replenishment using the Salpeter IMF. On the other hand, for the present day mass function low and intermediate mass stars are dominant.} {We confirm that late-type giant and supergiant stars dominate the ISM dust feedback in our actual Galaxy, but this may not the case of galaxies experiencing high star formation rates, or at high redshifts. In those cases, SNe are dominant in the dust feedback process.} ",
        "introduction": "RGB and AGB stars are known as the major continuous dust producers in the Universe, but also SN remnants have shown fast grain growth and dusty shells of M$_d$ $\\sim 10^{-2} - 10^{1}$ M$_{\\odot}$ \\citep{noz03}. From the classical nucleation theory, the timescale for grain growth in the ISM can be estimated by $\\tau_g = 4 s a (f n_i m_i v_i)^{-1}$, where $a$ is the dust mean size, $s$ is the material density, $f$ the sticking probability, $n_i$ the gas phase density, $m_i$ the atom mass and $v_i$ the velocity of the i-th atom to be added onto the grain surface. Considering typical ISM parameters, $a \\sim 0.1 - 1 \\ \\mu$m, $s \\sim 2$ g cm$^{-3}$, $n_i \\sim 1$ cm$^{-3}$, $T \\sim 10$ K, and assuming an efficiency $f = 0.1 - 1$, we find $\\tau_g < 10^9$ yr.  Dust particles are likely to be destroyed by shock waves \\citep{draine79, mckee89}. Grain-grain or ion-grain collisions will lead to the shattering process, reducing or destroying dust particles. \\citet{jones94, jones96} described the shattering process of ISM dust induced by SN blasts, obtaining destruction timescales of $\\tau_{d} \\sim 4 - 6 \\times 10^8$ yr. Actually, as mentioned in these works, the accurate derivation of destruction timescales depends on the velocity of the shock waves, the frequency of SNe and the physical properties of the ISM, as density and temperature. Hence, the destruction timescales are shorter than the dust growth scale, the observed dust could not be explained by nucleation in loco, but had to be recently ($< 10^9$ yr) injected into the ISM \\citep{tielens98}.  It is commonly suggested that the ISM dust should be mostly originated from evolved low and intermediate mass stars but recent observations showed the presence of large quantities of dust ($M_d = 10^8$ M$_{\\odot}$) in the early Universe (z $> 5$) \\citep{hughes98,arc01,dunne03,bert03,maio04}. At that age low and intermediate mass stars were not evolved yet and, therefore, SNe are recognized as responsible for such material ejections \\citep{tod01,mor03,sug06,dwek07}.  In the post-shock phase of the SN remnant, the gas is generally cool and dense enough to allow dust formation and growth \\citep{fal03,fal05}. Observationally, it is confirmed for SN1987A, which shows a $10^{-3}$ M$_{\\odot}$ dust shell with a dust to gas ratio for heavier elements $\\sim 0.3$. The same result was obtained for several other galactic and extragalactic SNe \\citep{bar05,gomez07}. Therefore, SNe seem to play an important part on the ISM dust replenishment process \\citep{dwek98}. However, if the short dust destruction timescale is taken into account, SNe could only be the main source of dust of the Galaxy if it presented a high star formation rates in its recent history.  In this work we present a semi-empirical model, in which we study the role of different stellar mass functions in the output of dust ejected to the ISM. The model is described in Sect. 2, in Sect. 3 we show the main results and present a brief discussion, followed by the conclusions. ",
        "conclusions": "It is still unclear what is the main source of the dust feedback in our Galaxy. The dust must be formed in stars and ejected to the ISM and some authors have argued favoring cool late-type stellar winds, which present high mass-loss rates and dust is proved to be formed in these sites by observations. On the other hand, other plausible sources are the SNe ejecta. During the final evolutionary phases, high mass stars explode and supersonically eject a very rich gas to the ISM.  Firstly, to determine the relative amount of dust ejected to the ISM for each stellar mass at each evolutionary phase we calculated the dust ejection during each evolutionary phase of the stars for different stellar mass distributions. As main result we showed that SNe are the main source of ISM dust feedback if a classic Salpeter IMF distribution is assumed. On the other hand, if we use the present day mass function, we show that the main sources of dust to the present Galaxy are the low and intermediate mass stars, representing more than 90\\% of the total dust mass.  Secondly, we studied the dust feedback process along the galactic time evolution, as done by \\cite{mor03}, but including the effects of dust destruction by SN blasts. For simplicity we used a constant destruction rate, consistent with the current galactic physical parameters. During previous ages the SNe ejecta are dominant, in agreement with previous works. We showed that, considering the dust destruction, low and intermediate mass stars are dominant for a galactic age of $t > 10^9$yr, in a much higher proportion. The total dust mass of $\\sim 10^7$ M$_{\\odot}$ is obtained for a star formation rate of 5 M$_{\\odot}$yr$^{-1}$.  The dust destruction timescale depends on the SNe frequency, as well as the ISM density and temperature. It is probable that the destruction rate was higher earlier during the galactic evolution. In this case, the presented conclusions will stand, and the role of low and intermediate mass stars in later stages will be even higher."
    },
    "0710/0710.3934_arXiv.txt": {
        "abstract": "A review of recent studies on a new mechanism of generation of large-scale magnetic field in a sheared turbulent plasma is presented. This mechanism is associated with the shear-current effect which is related to the ${\\bf W} {\\bf \\times} {\\bf J}$-term in the mean electromotive force. This effect causes the generation of the large-scale magnetic field even in a nonrotating and nonhelical homogeneous sheared turbulent convection whereby the $\\alpha$ effect vanishes (where ${\\bf W}$ is the mean vorticity due to the large-scale shear motions and ${\\bf J}$ is the mean electric current). It is found that turbulent convection promotes the shear-current dynamo instability, i.e., the heat flux causes positive contribution to the shear-current effect. However, there is no dynamo action due to the shear-current effect for small hydrodynamic and magnetic Reynolds numbers even in a turbulent convection, if the spatial scaling for the turbulent correlation time is $\\tau(k) \\propto k^{-2}$, where $k$ is the small-scale wave number. We discuss here also the nonlinear mean-field dynamo due to the shear-current effect and take into account the transport of magnetic helicity as a dynamical nonlinearity. The magnetic helicity flux strongly affects the magnetic field dynamics in the nonlinear stage of the dynamo action. When the magnetic helicity flux is not small, the saturated level of the mean magnetic field is of the order of the equipartition field determined by the turbulent kinetic energy. The obtained results are important for elucidation of origin of the large-scale magnetic fields in astrophysical and cosmic sheared turbulent plasma. ",
        "introduction": "\\label{}  Turbulence with a large-scale velocity shear is a universal feature in astrophysical plasmas. It has been recently recognized that in a sheared turbulent plasma with high hydrodynamic and magnetic Reynolds numbers a mean-field dynamo is possible even in a nonhelical and nonrotating homogeneous turbulence whereby a kinetic helicity and $\\alpha$ effect vanish (see Rogachevskii and Kleeorin 2003; 2004; 2007; Brandenburg 2005; Brandenburg and Subramanian 2005c; Rogachevskii et al. 2006a; 2006b). The large-scale velocity shear produces anisotropy of turbulence with a nonzero background mean vorticity ${\\bf W} = \\bec{\\nabla} {\\bf \\times} {\\bf U}$, where ${\\bf U}$ is the mean velocity. The dynamo instability in a sheared turbulent plasma is related to the ${\\bf W} {\\bf \\times} {\\bf J}$-term in the mean electromotive force, and it can be written in the form $\\bec{\\cal E}^\\delta \\propto - l_0^2 \\, {\\bf W} {\\bf \\times} (\\bec{\\nabla} {\\bf \\times} {\\bf B}) \\propto l_0^2 \\, ({\\bf W} {\\bf \\cdot} {\\bf \\Lambda}^B) {\\bf B}$, where $l_0$ is the maximum scale of turbulent motions (the integral turbulent scale) and $ {\\bf \\Lambda}^B = \\bec{\\nabla} {\\bf B}^2 / 2{\\bf B}^2 $ determines the inhomogeneity of the mean original magnetic field ${\\bf B}$. In a sheared turbulent plasma the deformations of the original magnetic field lines are caused by the upward and downward turbulent eddies, and the inhomogeneity of the original mean magnetic field in the shear-current dynamo breaks a symmetry between the influence of upward and downward turbulent eddies on the mean magnetic field. This creates the mean electric current ${\\bf J}$ along the mean magnetic field and produces the mean-field dynamo due to the shear-current effect.  The goal of this communication is to review recent studies on the new mechanism of generation of large-scale magnetic field due to the shear-current effect in a sheared turbulent plasma. The mean-field dynamo instability is saturated by the nonlinear effects. There are two types of the nonlinear effects caused by algebraic and dynamic nonlinearities. The effects of the mean magnetic field on the motion of fluid and on the cross-helicity result in quenching of the mean electromotive force which determines the algebraic nonlinearity. The dynamical nonlinearity in the mean-field dynamo is caused by the evolution of small-scale magnetic helicity, and it is of a great importance due to the conservation law for the magnetic helicity in turbulent plasma with very large magnetic Reynolds numbers (see, e.g., Kleeorin and Rogachevskii 1999; Brandenburg and Subramanian 2005a; Rogachevskii et al. 2006b, and references therein). The combined effect of the dynamic and algebraic nonlinearities saturates the growth of the mean magnetic field.  The shear-current effect has been studied by Rogachevskii and Kleeorin (2003; 2004; 2007) for large hydrodynamic and magnetic Reynolds numbers using two different approaches: the spectral $\\tau$ approximation (the third-order closure procedure) and the stochastic calculus (the path integral approach in a turbulence with a finite correlation time). A justification of the $\\tau$ approximation for different situations has been performed in numerical simulations and analytical studies by Blackman and Field (2002); Field and Blackman (2002); Brandenburg at al. (2004); Brandenburg and Subramanian (2005a, 2005b); Sur et al. (2007). ",
        "conclusions": ""
    },
    "0710/0710.4513_arXiv.txt": {
        "abstract": "Extragalactic jets are visualized as dynamic erruptive events modelled by time-dependent magnetohydrodynamic (MHD) equations. The jet structure comes through the temporally self-similar solutions in two-dimensional axisymmetric spherical geometry. The two-dimensional magnetic field is solved in the finite plasma pressure regime, or finite $\\beta$ regime, and it is described by an equation where plasma pressure plays the role of an eigenvalue. This allows a structure of magnetic lobes in space, among which the polar axis lobe is strongly peaked in intensity and collimated in angular spread comparing to the others. For this reason, the polar lobe overwhelmes the other lobes, and a jet structure arises in the polar direction naturally. Furthermore, within each magnetic lobe in space, there are small secondary regions with closed two-dimensional field lines embedded along this primary lobe. In these embedded magnetic toroids, plasma pressure and mass density are much higher accordingly. These are termed as secondary plasmoids. The magnetic field lines in these secondary plasmoids circle in alternating sequence such that adjacent plasmoids have opposite field lines. In particular, along the polar primary lobe, such periodic plasmoid structure happens to be compatible with radio observations where islands of high radio intensities are mapped. ",
        "introduction": "Collimated jets with high terminal velocities appear to be universal phenomena in astrophysics. They are often associated with young stellar objects, compact galactic objects, and active galactic nuclei [Livio 1997]. These jets are always accompanied by accretion disks on the equatorial plane. The magnetosphere of the accretion disk contains a plasma that follows the accretion of the materials in the disk. The plasma density and pressure get higher as the accretion approaches to the center which drive the magnetic field. The dynamics of this system was first described in magnetohydrodynamic (MHD) model by Blandford and Payne [1982]. In this landmark paper, jets are considered as a spatial structure in stationary state. Ideal MHD equations in cylindrical coordinates $(r,\\phi,z)$ are solved for time independent solutions. The central mass $M$ is replaced by a linear mass along the cylindrical axis. Self-similar solutions in space with a scale invariance $z/r$ are sought. Such self-similar solutions are compatible to a Keplerian disk plasma rotation velocity field superimposed by an Alfv\\'{e}nic plasma velocity. The complete MHD instability spectrum with such Keplerian profile are analysed by Keppens, Casse, and Goedbloed [2002]. This model predicts that the magnetospheric disk plasma would be ejected towards the polar direction magnetocentrifugally should the magnetic field lines be at an angle less than $\\pi/3$ or more that $2\\pi/3$ on the $rz$ plane with respect to the outward radius of the disk. Collimating action of this plasma outflow would be provided by the hoop force of the azimuthal magnetic field and its associated parallel current in a force-free configuration. The interaction of the magnetic field with the plasma disk generates a MHD Poynting flux [Ferreira and Pelletier 1995] that can be converted into kinetic energy of the jet plasma [Zanni et. al. 2004]. Variants of jet formation model are proposed by Contopoulos and Lovelace [1994], Contopoulos [1995], and Cao and Spruit [1994]. This accretion-ejection model provides the basic framework of current investigations of accretion disk and jets as is reviewed by Balbus and Hawley [1998]. Dissipative MHD effects are examined by Casse and Ferreira [2000 a,b] and Casse [2004], and relativistic jets are analyzed by Vlahakis [2004]. These stationary state analytic studies are often complemented by numerical works in cylindrical geometry to simulate time evolutions of the jets [Ustyugova et. al. 1995, Ouyed and Pudritz 1997, Krasnopolski et. al. 1999], and the disk-jet system [Matsumoto et. al. 1996, Kato et. al. 2002].  Here, instead of a stationary model, we take the view that jets are a time dependent spatial structure. What we see is only a snapshot of their state at this particular moment. Due to their galactic dimensions, the time scale of these structures is believed to be extraordinarily large which gives the impression of a stationary structure. This implies that jets are results of an erruption originating from the galactic nucleus. They could be dissipated in time before another erruption takes place due to pressure built-up from accretion. Or one erruption could be superimposed on an earlier event. To model the jet system, we will do a self-similar analysis on the full time-dependent ideal MHD equations in spherical coordinates $(r,\\theta,\\phi)$ with a mass $M$ at the nucleus. In particular, we consider axisymmetric solutions. This type of self-similar solutions were pioneered by Low [1982a,b, 1984a,b] for astrophysical applications and solar corona mass ejections with pure radial plasma velocity flow. Variants of these solutions include cases where the plasma domain lies outside the mass $M$ such as interplanetary magnetic ropes in one-dimensional [Osherovich, Farrugia, and Burlaga 1993, 1995] and in two-dimensional [Tsui and Tavares 2005] cylindrical geometry, interplanetary magnetic clouds [Tsui 2006a], and also atmospheric ball lightnings [Tsui 2006b] in spherical geometry. In these descriptions where $M$ is outside the spherical domain of interest, the magnetic field is axisymmetric force-free and contains regions of closed field lines, while the plasma is spherically symmetric decoupled from the magnetic field.  For the present case of extragalactic jets described by the mechanism of accretion-ejection, we will follow the self-similar solutions of Low with the polytropic index $\\gamma=4/3$, but with particular emphasis on the finite plasma pressure. This current approach differs from the Keplerian disk plasma in that the radial flow is not tied to the Alfv\\'{e}n velocity as a priori. In this dynamic model, we consider jets as a manifestation of mass ejection on a galactic scale. The plasma pressure, in this self-similar MHD model, proves to have an important role in collimating the magnetic fields and the jet plasmas. The time evolution function gives a dynamic description of the high radial flow especially in the jets. The self-similar solutions, that converge at the center and at infinity, give small regions of closed axisymmetric two-dimensional magnetic field lines where plasma density and pressure are much higher. These regions along the jets could correspond to the high intensity islands in radio frequency maps.  \\newpage ",
        "conclusions": "One of the main objections of self-organized plasmoid representation in free space in the absence of an adequate boundary is that it apparently violates the Virial theorem which states that [Schmidt 1966] \\\\ $${1\\over 2}{d^2 I\\over dt^2} +\\int x_{k}{\\partial G_{k}\\over\\partial t}dV\\, =\\,2(T+U)+W^{E}+W^{M} -\\int x_{k}(P_{ik}+T_{ik})dS\\,\\,\\,,\\eqno(28)$$ \\\\ where $I$ is the moment of inertia of the plasmoid, $G$ is the momentum density of the electromagnetic field, $T$ and $U$ are the kinetic and thermal energies of the plasma, $W^{E}$ and $W^{M}$ are the electric and magnetic energies in the volume, $P_{ik}$ and $T_{ik}$ are the plasma and electromagnetic stress tensors. Taking the volume to cover the entire plasma and field, the surface term on the right side vanishes. In laboratory plasmas, this surface could be the machine vessel. Should the plasmoid be stationary, the volume term on the left side would be null, and the moment of inertia would be accelerating since the terms on the right side are all positive definite. While this statement has no conflict with the unbounded solutions, it apparently contradicts the stationary state of the bounded solutions. Nevertheless, this argument has overlooked the asymptotically bounded nature of the $H=-|H|<0$ plasmoid state. In this asymptotic case, we have $dI/dt=(dI/dy)(dy/dt)=0$ so that $I$ is stationary. The fact that $d^2I/dt^2>0$ implies that $I$ is at an asymptotic minimum, not an acceleration of $I$, which complies with the Virial theorem.  The classical accretion-ejection model of Blandford and Payne [1982] is a time-independent stationary state model with spatial self-similar MHD solutions in cylindrical geometry $(r,\\phi,z)$ with Alfv\\'{e}nic plasma flow velocity plus a rotation, all with Keplerian scaling. In this model, the jets are formed by convecting the magnetospheric disk plasmas from the disk plane to the axis by magnetocentrifugal action through the magnetic field lines with low inclination angles to the disk plane. The angular momentum of the plasma on the disk plane is focused to the axis. The jets were put in place in the distant past according to the spatial self-similar solutions in $z/r$, and are maintained there by continuously transporting disk plasmas to the axis to sustain the axial outflow. Collimation to the axis is accomplished by magnetic hoop force.  Here, we have taken a dynamic view where jets are the consequences of erruptive events, based on time-dependent MHD equations in spherical geometry $(r,\\theta,\\phi)$. The disk plasmas are accreted to the galactic nucleus where plasma pressure is built up to cause an erruption. The radially symmetric expanding velocity interacts the plasma with the magnetic field self-consistently through the MHD equations to generate spatial structures. Due to the existence of multiple quadratic invariants in the absence of dissipations, MHD systems have the tendency of developing self-organized and self-similar states through dissipative processes. The force-free configuration and the vortex nature of the magnetic field are akin to the quadratic invariants of the magnetic helicity and cross helicity of the MHD system. For these reasons, although temporally self-similar solutions are only a subset of general time-dependent MHD solutions, these self-similar configurations are prone to develope in natural phenomena.  We, therefore, describe these spatial structures by self-similar temporal solutions in $\\eta=r/y$. In this self-similar spherical model, consistent self-similar representations of plasma pressure, mass density, and magnetic fields are worked out for the $\\gamma=4/3$ BC Low model, with special emphasis on the finite $\\beta$ case. The spatial distribution of the magnetic field is described by an equation where the plasma pressure acts as the eigenvalue. The nature of this eigenvalue equation is as such that the magnetic field gets converged to the polar axis as a response to the high plasma pressure directly related to the eigenvalue. Since the separation constant $n(n+1)$ in this eigenvalue equation is a free parameter that as yet to be specified, there will be an adequate $n$ for almost any given plasma pressure. Although the radial plasma velocity is isotropic, the spatial structures are not. They could be highly collimated along the axial direction, and expand into the previously void space as time progresses.  Although other types of solutions are permitted, we have specifically examined the solutions that bear resemblance with jet features with $n\\gg 1$ and $a\\eta_{0}\\gg 1$. Since plasma and magnetic field are frozen into each other, plasma outflow is also collimated to the polar axis. The existence of regions of closed field lines along the primary polar magnetic lobe permits secondary plasmoids be embedded in it. These secondary plasmoids appear to be compatible to the observed islands of radio intensities. The time evolution function of the radial velocity consistent to the temporal self-similar solutions has different types of solutions according to the sign of $H$. Although the equatorial disk magnetospheric plasma is not addressed here, the accelerated accretion of plasma influx could be modelled with $H=0$ as the boundary condition at infinity. As for the polar jets, $H<0$ gives a bounded oscillating, or pulsating, solution. This bounded stage, we believe, nourishes the self-organized and self-similar states. With $H>0$, it gives an unbounded expanding solution with a high terminal velocity. It is apparent that plasma pressure due to accretion is the prime driving force that determines the value of $H$. The bounded oscillation would make a transition to the unbounded expansion as $H$ goes from negative to positive. According to our model, jet structures are, therefore, considered as a spatial configuration that has been expanding continuously into space.  \\newpage"
    },
    "0710/0710.4039_arXiv.txt": {
        "abstract": "{The X-ray source \\source\\/ is a very weak Low Mass X-ray Binary discovered in 1999 with \\sax\\/ and located in the Galactic Center. This region has been deeply investigated by the \\integral\\/ satellite with an unprecedented exposure time, giving us an unique opportunity to study the hard X-ray behavior also for weak objects. } {We analysed all available \\integral\\/ public and private Key-Program observations with the main aim of studying the long-term behavior of this Galactic bulge X-ray burster. } {The spectral results are based on the systematic analysis of all \\integral\\/ observations covering the source position performed between February 2003 and October 2006. \\source\\/ did not shows any flux variation along this period as well as compared to previous \\sax\\/ observation. Hence, to better constrain the physical parameters we combined both instrument data. } {Long \\integral\\/ monitoring reveals that this X-ray burster is a weak persistent source, displaying a X-ray spectrum extended to high energy and spending most of the time in a low luminosity hard state. The broad-band spectrum is well modeled with a simple Comptonized model with a seed photons temperature of $\\sim$ 0.5 keV and an electron temperature of $\\sim$24 keV. The low mass accretion rate ($\\sim$ 2$\\times$10$^{-10}$ M$_{sun}$/yr), the long bursts recurrence time, the small sizes of the region emitting the seed photons consisting with the inner disk radius and the high luminosity ratio in the 40--100 keV and 20--40 keV band, are all features common to the Ultra Compact source class. } {  We report, for the first time, on the X-ray behavior of this source: observations with unprecedent wide energy band and sensitivity revealed this X-ray burster is a persistent source with an hard spectrum extending up to energies $>$ 100 keV. } ",
        "introduction": "\\source\\/ was discovered in 1999 during the monitoring campaign of the Galactic Centre region performed by the \\sax\\//WFC (Cocchi et al. 1999 and in 't Zand et al. 1999). This instrument observed a type-I X-ray burst with exponential decays on August 1999, identifying the compact object as a weakly magnetized neutron star in a low mass X-ray binary system with a derived distance of 7 kpc (Cocchi et al. 2001). The source was followed-up by the \\sax\\//instruments to investigate the spectral properties: \\source\\/ shows a spectrum extended up to about 60 keV, with an intensity of $\\sim$ 6 mCrab, both in the MECS (1--10 keV) and in the PDS (15--60 keV). The spectrum was fitted by an absorbed power law ($\\Gamma\\sim$2.2) and a column density N$_H$$\\sim$2$\\times$10$^{22}$cm$^{-2}$ or by a Comptonised emission model with electron temperature of $\\sim$ 50 keV and a lower column density N$_H$$\\sim$1$\\times$10$^{22}$cm$^{-2}$ (Cocchi et al. 2001). The ROSAT source J171237.1-373834 is just 0\\farcm6 from the \\sax\\/ position with a flux of 1.6 mCrab in the energy range 0.1--2.4 keV (in 't Zand et al. 1999). No optical counterpart has been identified yet within the 13\\arcsec\\ (1$\\sigma$) ROSAT error circle radius. \\source, monitored during the PCA/RXTE bulge scan program, is continuously active, apparently in two states: a slowly changing state, and a quicker one (in 't Zand et al. 2007). It was included in the IBIS/ISGRI soft Gamma-Ray survey catalog (Bird et al. 2007) as a transient X-ray buster with fluxes corresponding to $4.7\\pm0.1$ mCrab and $4.0\\pm0.2$ mCrab in the energy bands 20--40 keV and 40--100 keV, respectively. Recently, Chelovekov et al. (2006) reported two further burst detections with \\integral\\//IBIS in 2003-2004, both with the same peak flux. ",
        "conclusions": "\\integral\\/ observations reveal for the first time that this X-ray burster is a persistent source, displaying a very hard X-ray spectrum and spending most of the time in a low luminosity and hard state. A broad-band spectrum obtained combining the \\integral\\/ together with the \\sax\\/ data displays properties which fits well into the classification of low luminosity ($\\sim$0.01  $ L_{\\rm Edd}$) weakly magnetized neutron stars with \"hard spectra'' (Barret et al. 2001): a broad-band spectrum extending up to 100 keV, dominated by a Comptonized component, with a optical depth of the Comptonizing corona of $\\sim$3, seed photons temperature of $\\sim$0.5 keV, an electron temperature of $\\sim$24 keV. A few LMXBs appear to spend most of the time in this \"hard state'' (Di Salvo \\& Stella 2002), e.g. 4U 0614+091 (Piraino et al. 1999) or 4U1850-087 (Sidoli et al. 2006).  The emission spectrum could be also fitted with a two component model consisting of a blackbody, or multi color blackbody component which represents emission from an optically thick accretion disk or from the neutron star surface, together with a  Comptonized component which is interpreted as emission from a hot inner disk region or a boundary layer between the disk and the neutron star. We have estimate a black body radius of $\\sim 5$ km for the black body component and the inner radius of the disk $< 11 km$ for the multi disk component. Although the data statistics does not allow us to distinguish between the two physical models, both values of the radius give an indication of the system compactness.  We have also estimated the sizes of the emitting seed photons region, assuming their emission as blackbody with temperature T$_0$ and radius R$_{seed}$ following in 't Zand et al. (1999). We obtain $R_{\\rm seed}\\simeq 8.3\\,{\\rm km}\\, (L_{\\rm C}/10^{37}\\,{\\rm erg\\, s}^{-1})^{1/2} (1+y)^{-1/2} (kT_0/1\\,{\\rm keV})^{-2}$, where $L_{\\rm C}$ is the luminosity of the Comptonization component. Here, we have estimated the amplification factor of the seed luminosity by the Comptonization as $(1+y)$, where $y$ is the Comptonization parameter, $y=4kT_{\\rm e}\\max(\\tau,\\tau^2)/m_{\\rm e} c^2$. The obtained values are $\\sim$ 1 km and $\\sim$ 10 km for the model with the blackbody component and with multicolor disk blackbody component, respectively. The parameters obtained using the latter model allows for a more reasonable interpretation.  \\begin{table*} \\vspace{1cm} \\caption{Results of the \\source\\/ fit, using different models: a simple comptonized model, a blackbody plus a Comptonized component and a multicolor disk blackbody plus a Comptonized component. Uncertainties are at the 90\\% confidence level for a single parameter variation.} \\label{fitsax} \\vspace{0.3cm} \\centering \\begin{tabular}{ccccccccc} \\hline\\hline && & & & &  &&\\\\ Model 1&{N$_{\\rm H}$} &...&...&  {T$_{\\rm o}$}& {T$_{\\rm e}$} &$\\tau$ & $n_{\\rm COMPTT}$ & $\\chi^2$/d.o.f \\\\ &($10^{22}~cm^{-2}$)&...&... &($keV$) &(keV)&...&($10^{-3}$)&...\\\\ && & & & & & &\\\\ $comptt$&$1.3\\pm0.1$&...&... &$0.47\\pm0.02$ &$24^{+16}_{-10}$  &$2.7\\pm1.3$ &$2.0^{+1.2}_{-0.8}$&142/134\\\\ && & & & & & &\\\\ \\hline                                                                              && & & & & & &\\\\ Model 2&{N$_{\\rm H}$} & {T$_{\\rm in}$}& ${\\rm r_{in}({\\cos}i)^{0.5}}$& {T$_{\\rm o}$}& {T$_{\\rm e}$} &$\\tau$ & $n_{\\rm COMPTT}$ & $\\chi^2$/d.o.f \\\\ &($10^{22}~cm^{-2}$) &($keV$)&(km)&($keV$) &(keV)&...&($10^{-3}$)&...\\\\ && & & & & & &\\\\ $diskbb+comptt$&$1.4\\pm0.3$&$0.7\\pm0.3$  &$<11$ &$0.5\\pm0.2$ &$24^{+22}_{-9}$  &$3\\pm2$ &$1.6^{+1.6}_{-0.8}$&120/132\\\\ && & & & & & &\\\\ \\hline && & & & & & &\\\\ Model 3 &{N$_{\\rm H}$} & {T$_{\\rm bb}$}& {R$_{\\rm bb}$}& {T$_{\\rm o}$}& {T$_{\\rm e}$} &$\\tau$ & $n_{\\rm COMPTT}$ & $\\chi^2$/d.o.f \\\\ &($10^{22}~cm^{-2}$) &($keV$)&(km)&($keV$) &(keV)&...&($10^{-3}$)&...\\\\ && & & & & & &\\\\ $bbody+comptt$&$1.2\\pm0.2$&$0.6\\pm0.1$&$5\\pm1$&$1.2\\pm0.2$ &$34^{+54}_{-14}$ &$2.0\\pm1.5$ &$0.5^{+0.6}_{-0.2}$&119/132\\\\ && & & & & & &\\\\ \\hline \\hline \\end{tabular}\\\\ \\end{table*}  These results show that the radius of the region providing the seed photons for the Comptonization emission is consistent with the inner radius of the accretion disk, showing that the seed photons are bounded in small region spatially coincident with the inner part of the disk. Moreover the seed photons temperature $T_0$ is consistent with the inner disk temperature $T_{in}$. By imposing $T_0$=$T_{in}$, the fit values correspond to a disk temperature of $0.57\\pm0.08$ keV with a $\\chi^2/dof$ of 121/133. This behavior suggests to identify the inner disk as the region providing the seed photons of the Comptonization emission. This result was observed in most of ultra compact X-ray source: 4U~1850-087, 4U~1820-303 and 4U~0513-40 (Sidoli et al. 2001), XTE J1751-305 (Gierli\\'nski \\&  Poutanen 2005), XTE J1807-294 (Falanga et al. 2005).  Assuming the persistent flux represents the average mass accretion rate and an accretion efficiency of $\\eta = 0.2$ (corresponding, e.g., to $M_{\\rm NS}=1.4 {\\rm M}_{\\sun}$ and $R_{\\rm NS}=10$ km) and using multi color disk model luminosity, $L_{0.1-200 keV}=1.6\\times10^{36} erg s^{-1}$, we get $\\dot{M}\\simeq 2.8 \\times 10^{-10} {\\rm M}_{\\sun}$ yr$^{-1}$ for \\source. This low mass accretion rate is consistent with the hard spectrum of \\source\\/, according to our present understanding on the X-ray bursters (for review see van der Klis 2006). Moreover, the burst recurrence time is dependent on the mass accretion rate on the neutron star: the faster new fuel is provided from the donor star, the shorter is the burst recurrence time. Long \\integral\\/ monitoring shows clearly that the frequency of bursts in this sources is low, as previously reported by in 't Zand et al. (2007), who estimated a recurrence time of 345--6507 hr.  In 't Zand et al. (2007) have selected all 31 persistent X-ray bursters as reported in the \\integral\\/ survey (Bird et al. 2006) and demonstrated that the low mass accretion rates are accompanied by the hard X-ray spectra of the persistent emission. Almost all the Ultra Compact X-ray binaries have the highest values of the (40--100)keV/(20--40)keV hardness ratio. \\source\\/ was not included as supposed to have a transient nature. We now know it is a persistent and weak source, showing a high ratio L$_{(40-100)keV}$/L$_{(20-40)keV}$$\\sim$0.90, in agreement with the value of 0.85$\\pm$0.06 as derivated from recent data in Bird et al. 2007. Finally, the low mass accretion rate, the long bursts recurrence time, the small sizes of the region emitting the seed photons consisting with the inner disk radius and the high luminosity ratio in the 40--100 keV and 20--40 keV support the idea of In 't Zand et al. (2007) that this source could be a candidate Ultra Compact Binary."
    },
    "0710/0710.3105_arXiv.txt": {
        "abstract": "We present HST/WFPC2 Linear Ramp Filter images of high surface brightness emission lines (either [OII], [OIII], or H$\\alpha$+[NII]) in 80 3CR radio sources. We overlay the emission line images on high resolution VLA radio images (eight of which are new reductions of archival data) in order to examine the spatial relationship between the optical and radio emission. We confirm that the radio and optical emission line structures are consistent with weak alignment at low redshift ($z < 0.6$) except in the Compact Steep Spectrum (CSS) radio galaxies where both the radio source and the emission line nebulae are on galactic scales and strong alignment is seen at all redshifts. There are weak trends for the aligned emission line nebulae to be more luminous, and for the emission line nebula size to increase with redshift and/or radio power. The combination of these results suggests that there is a limited but real capacity for the radio source to influence the properties of the emission line nebulae at these low redshifts ($z < 0.6$). Our results are consistent with previous suggestions that both mechanical and radiant energy are responsible for generating alignment between the radio source and emission line gas. ",
        "introduction": "\\label{sec:Introduction}  Radio galaxies are an important class of extragalactic objects:  they represent one of the most energetic astrophysical phenomena; they may be used as probes of their environments; and they are unique probes of the early Universe \\citep{McCarthy93}. As the 3CR sample of radio galaxies is radio flux-density selected \\citep{Bennett62}, it provides an unbiased optical sample to study the host galaxies of these radio-loud active galactic nuclei (AGN). The sample has been well studied at multiple wavelengths and previous ground-based studies have investigated the spatial coincidence of optical emission nebulae and the radio source \\citep{Baum88,Baum89a,McCarthy87,Chambers87}.  \\citet{Chambers87} and \\citet{McCarthy87} first demonstrated the ``alignment effect'' at high redshift ($z \\geqslant 0.6$) where the continuum optical and/or IR emission is aligned along the radio axis. An ``alignment effect'' for emission line gas was also shown by \\citet{McCarthy87} and \\citet{McCarthy95}.  In this paper we focus on the alignment of  the fine scale high surface brightness emission line gas imaged by the Hubble Space Telescope (HST) with the radio emission. The primary mechanisms for the emission line gas alignment are thought to be shocks induced by the radio jet and photoionization from the central AGN \\citep[e.g.,][]{Baum89a,McCarthy93,Best00}.  \\citet{McCarthy89} demonstrated the redshift dependence of the alignment effect - at $z < 0.1$ there is no alignment, at $0.1 < z < 0.6$ some alignment is seen, and for $z > 0.6$ essentially all the powerful radio galaxies show strong alignment. \\citet{Baum89b} found alignment at low $z$ if one compares the emission in the quadrants containing the radio source to quadrants without radio emission. Studies which attempt to break the redshift-radio power degeneracy suggest that the alignment depends on both redshift and radio power \\citep[e.g.,][]{Inskip02b}. Thus, the alignment depends on both the presence of extended gas in the environment as well as the ability of the radio source to influence its environment.  In this paper we examine the relationship of the high surface brightness emission line gas and radio emission over a wide range of redshift ($0.017 < z < 1.406$), focusing on the transition redshift range $z < 0.6$ where the alignment effect starts to turn on. Using the Wide Field and Planetary Camera 2 (WFPC2) on HST, we obtained images with resolutions of 0.$\\arcsec$05 -- 0.$\\arcsec$1, similar to or better than the synthesis imaging resolution obtained with the Very Large Array \\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} (VLA) at 5 GHz in the ``A'' configuration. The higher resolution allows us to probe denser gas in the center of these galaxies. In addition, the tunable Linear Ramp Filter (LRF) permits narrow band imaging of a selected spectral feature for a wide range of redshifts, enabling studies of a specified emission line for a sample such as the 3CR covering a large range of redshift.  \\citet{Biretta02} conducted an initial study using this data, and published a subset of the data.  Here we expand on their analysis, presenting the complete set of detected images as well as statistical results.  The paper is organized in the following fashion: we discuss the properties of the sample in $\\S$\\ref{sec:Sample}. $\\S$\\ref{sec:Observations} discusses the observations made and the data reduction methods used. In $\\S$\\ref{sec:Data Analysis} we outline the analysis of the data. $\\S$\\ref{sec:Properties} contains a discussion of our sources and the results of our study. In $\\S$\\ref{sec:Conclusions} we conclude by discussing the implications of our results. Appendix \\ref{Appendix} contains notes for individual sources. ",
        "conclusions": "\\label{sec:Conclusions}  We present WFPC2 LRF images of selected emission lines in a sample of 3CR radio galaxies with bright emission lines. We detect clumpy, high surface brightness emission line gas with a wide range of morphologies. We examine the properties of the emission line gas and its relationship to the radio source, and in particular the nature of the alignment effect at $z < 0.6$ where the effect is beginning to become important.  We find there is a weak alignment effect in our low $z$ sample ($z < 0.6$). The effect is clearly weaker than seen at high redshifts and is also somewhat weaker than seen at similar redshifts in ground-based data. This suggests that the alignment effect is dominated by the more extended lower surface brightness emission line gas which falls below our detection threshold in WFPC2.   The alignment at low redshift is strong when the radio source and the nebula are both of order galactic scales (e.g., CSS sources). This implies there is mechanical energy input from the radio source as it propagates through the ambient gas in the host galaxy \\citep[e.g.,][]{deVries99,Axon00}.  There is a weak trend for the nebulae in the aligned sources (at all radio/optical size ratios) to have higher line luminosity. This suggests that the radio source preferentially provides energy to gas which is along the radio source axis. This is likely to be predominantly mechanical energy when the radio source and nebulae have similar sizes, and predominantly photoionization when the radio source is much larger than the nebula \\citep[e.g.,][]{Best00,Inskip02a,Moy02,ODea02,Labiano05}.  There is a weak trend for the size of the high surface brightness emission line nebulae to increase with increasing radio power and/or redshift over the range $z < 0.6$. This may be due to both the increasing presence of  gas on large scales and the ability of the more powerful radio sources  to provide additional mechanical energy and/or ionizing photons as a function of redshift.  Thus, there is evidence for a modest excess of aligned sources at low redshift ($z < 0.6$). The alignment is strong when the radio source and emission line nebula are both on galactic scales. There are also weak trends for the aligned emission line nebulae to be more luminous, and for the emission line size to increase with redshift and radio power. The combination of these results  suggests that there is a limited but real capacity for the radio source to influence the properties of the emission line nebulae at these low redshifts."
    },
    "0710/0710.3823.txt": {
        "abstract": "We study the transition from inspiral to plunge in general relativity by computing gravitational waveforms of non-spinning, equal-mass black-hole binaries. We consider three sequences of simulations, starting with a quasi-circular inspiral completing 1.5, 2.3 and 9.6 orbits, respectively, prior to coalescence of the holes. For each sequence, the binding energy of the system is kept constant and the orbital angular momentum is progressively reduced, producing orbits of increasing eccentricity and eventually a head-on collision. We analyze in detail the radiation of energy and angular momentum in gravitational waves, the contribution of different multipolar components and the final spin of the remnant, comparing numerical predictions with the post-Newtonian approximation and with extrapolations of point-particle results. We find that the motion transitions from inspiral to plunge when the orbital angular momentum $L=L_{\\rm crit}\\simeq 0.8M^2$. For $L<L_{\\rm crit}$ the radiated energy drops very rapidly. Orbits with $L\\simeq L_{\\rm crit}$ produce our largest dimensionless Kerr parameter for the remnant, $j=J/M^2\\simeq 0.724\\pm0.13$ (to be compared with the Kerr parameter $j\\simeq 0.69$ resulting from quasi-circular inspirals). This value is in good agreement with the value of $0.72$ reported in \\cite{Hinder2007}.  These conclusions are quite insensitive to the initial separation of the holes, and they can be understood by extrapolating point particle results.  Generalizing a model recently proposed by Buonanno, Kidder and Lehner \\cite{Buonanno:2007sv} to eccentric binaries, we conjecture that (1) $j\\simeq 0.724$ is close to the maximal Kerr parameter that can be obtained by any merger of non-spinning holes, and (2) no binary merger (even if the binary members are extremal Kerr black holes with spins aligned to the orbital angular momentum, and the inspiral is highly eccentric) can violate the cosmic censorship conjecture. ",
        "introduction": "The research area of gravitational wave (GW) physics has reached a very exciting stage, both experimentally and theoretically.  Earth-based laser-interferometric detectors, including LIGO \\cite{Abramovici:1992ah}, GEO600 \\cite{Luck:1997hv} and TAMA \\cite{Ando:2001ej}, are collecting data at design sensitivity, searching for GWs in the frequency range $\\sim 10-10^3$~Hz. VIRGO \\cite{Caron:1997hu} should reach design sensitivity within one year, and the space-based interferometer LISA is expected to open an observational window at low frequencies ($\\sim 10^{-4}-10^{-1}$~Hz) within the next decade \\cite{Danzmann:1998}.  The last two years have also seen a remarkable breakthrough in the simulation of the strongest expected GW sources, the inspiral and coalescence of black-hole binaries \\cite{Pretorius2005a,Campanelli2006, Baker2006}.  Several groups have now generated independent numerical codes for such simulations \\cite{Pretorius2006, Baker2006a, Campanelli2006a, Sperhake2006, Scheel2006, Bruegmann2006a, Herrmann2007, Pollney2007, Etienne2007} and studied various aspects of binary black hole mergers. % including the resulting gravitational recoil %\\cite{Herrmann2007,Gonzalez2007,Koppitz2007,Gonzalez2007a,Campanelli2007v2, %  Baker2007,Tichy2007,Herrmann2007c,Schnittman:2007ij,Lousto:2007db,Sopuerta:2006wj,Sopuerta:2006et} %spin precession and spin flips \\cite{Campanelli2007b,Campanelli2007v2}, In the context of analyzing the resulting gravitational waveforms, these include in particular the comparisons of numerical results with post-Newtonian (PN) predictions \\cite{Baker2006c, Schnittman2007, Hannam2007, Buonanno2006, Berti:2007fi, Boyle2007}, multipolar analyses of the emitted radiation \\cite{Buonanno2006, Berti:2007fi, Schnittman:2007ij}, the use of numerical waveforms in data analysis \\cite{Ajith:2007qp, Pan2007, Vaishnav2007, Ajith:2007kx} and gravitational wave emission from systems of three black holes \\cite{Campanelli:2007ea}.  Despite this progress, a comprehensive analysis of binary black hole inspirals remains a daunting task, mainly because of the large dimensionality of the parameter space.  In geometrical units, the total mass of the binary is just an overall scale factor. The source parameters to be explored by numerical simulations (sometimes called ``intrinsic'' parameters in the GW data analysis literature) include the mass ratio $q=M_2/M_1$, the eccentricity $e$ of the orbit and six parameters for the magnitude of the individual black hole spins and their direction with respect to the binary's orbital angular momentum.  In this paper we present results from numerical simulations of non-spinning, equal-mass black-hole binaries, and we focus on the effect of the orbital eccentricity on the merger waveforms. We consider three sequences, starting with quasi-circular inspirals that complete $\\sim 1.5$, $\\sim 2.3$ and $\\sim 9.6$ orbits, respectively, prior to coalescence of the holes. By fixing the binding energy of the system and progressively reducing the orbital angular momentum, we produce a sequence of orbits of increasing eccentricity and eventually a head-on collision. For each of these simulations we analyze in detail the radiation of energy and angular momentum in GWs, the contribution of different multipolar components and the final spin of the remnant, comparing numerical predictions with the PN approximation and with extrapolations of point-particle results.  Non-eccentric inspirals are usually considered the most interesting cases for GW detection. For an isolated binary evolving under the effect of gravitational radiation reaction, the eccentricity decreases by roughly a factor of 3 when the orbital semimajor axis is halved \\cite{Peters:1963ux}. For most conceivable formation mechanisms of solar-mass black hole binaries, the orbit will usually be circular by the time the GW signal enters the best-sensitivity bandwidth of Earth-based interferometers. % However, we wish to stress that our simulations could be of interest for GW detection. For example, according to some astrophysical scenarios, eccentric binaries may be potential GW sources for Earth-based detectors.  In globular clusters, the inner binaries of hierarchical triplets undergoing Kozai oscillations can merge under gravitational radiation reaction, and $\\sim 30\\%$ of these systems can have eccentricity $\\sim 0.1$ when GWs enter the detectors' most sensitive bandwidth at $\\sim 10$~Hz \\cite{Wen:2002km}. %In globular clusters, binaries with essentially a thermal distribution of %eccentricities may exist \\cite{Benacquista2002}. Massive black hole binaries to be observed by LISA could also have significant eccentricity in the last year of inspiral. %Analytical calculations and $N$-body simulations show that, in purely %collisionless spherical backgrounds, the expected equilibrium distribution of %eccentricities is skewed towards high $e\\simeq 0.6-0.7$, and that dynamical %friction does not play a major role in modifying such a distribution %(\\cite{1999ApJ...525..720C}, in particular Fig.~5). Recent simulations using smoothed particle hydrodynamics follow the dynamics of binary black holes in massive, rotationally supported circumnuclear discs \\cite{2006MNRAS.367..103D, 2007MNRAS.379..956D, Mayer:2007vk}. In these simulations, a primary black hole is placed at the center of the disc and a secondary black hole is set initially on an eccentric orbit in the disc plane. By using the particle splitting technique, the most recent simulations follow the binary's orbital decay down to distances $\\sim 0.1$~pc. Dynamical friction is found to circularize the orbit if the binary {\\it corotates} with the disc \\cite{2007MNRAS.379..956D}.  However, if the orbit is {\\it counterrotating} with the disc the initial eccentricity does not seem to decrease, and black holes may still enter the GW emission phase with high eccentricity \\cite{2006MNRAS.367..103D}.  Complementary studies show that eccentricity evolution may still occur in later stages of the binary's life, because of close encounters with single stars and/or gas-dynamical processes.  Three-body encounters with background stars have been studied mainly in spherical backgrounds. These studies find that stellar dynamical hardening can lead to an increase of the eccentricity, acting against the circularization driven by the large-scale action of the gaseous and/or stellar disc, possibly leaving the binary with non-zero eccentricity when gravitational radiation reaction becomes dominant \\cite{1996NewA....1...35Q,Aarseth:2002ie,Berczik:2005ff,Berczik:2006tz,Matsubayashi:2005eg}. It has also been suggested that the gravitational interaction of a binary with a circumbinary gas disc could increase the binary's eccentricity. The transition between disc-driven and GW-driven inspiral can occur at small enough radii that a small but significant eccentricity survives, typical values being $e\\sim 0.02$ (with a lower limit $e\\simeq 0.01$) one year prior to merger (cf.  Fig.~5 of \\cite{Armitage:2005xq}). If the binary has an ``extreme'' mass ratio $q\\lesssim 0.02$ the residual eccentricity predicted by this scenario can be considerably larger, $e\\gtrsim 0.1$. Numerical simulations should be able to test these predictions in the near future.  As shown by Sopuerta, Yunes and Laguna, eccentricity could significantly increase the recoil velocity resulting from the merger of non-spinning black-hole binaries \\cite{Sopuerta:2006et}.  Independently of the presence of eccentricity in astrophysical binary mergers, the problem we consider here has considerable theoretical interest.  Our simulations explore the transition between gravitational radiation from a quasi-circular inspiral (the expected final outcome in most astrophysical scenarios) and the radiation emitted by a head-on collision, where the binary has maximal symmetry. Our work should provide some guidance for analytical studies of the ``transition from inspiral to plunge''. The first analytical study of this problem in the context of PN theory was carried out by Kidder, Will and Wiseman \\cite{Kidder:1993}. The transition between the adiabatic phase and the plunge was studied in \\cite{Buonanno2000} using nonperturbative resummed estimates of the damping and conservative parts of the two-body dynamics, i.e. the so-called ``Effective One Body'' (EOB) model.  Ori and Thorne \\cite{Ori:2000zn} provided a semi-analytical treatment of the transition in the extreme mass ratio limit.  Waveforms comprising inspiral, merger and ringdown for comparable-mass bodies have also been produced using the EOB model (see eg.~\\cite{Buonanno:2005xu} for extensions of the original model to spinning binaries and for references to previous work).  Preliminary comparisons of EOB and numerical relativity waveforms showed that improved models of ringdown excitation \\cite{Buonanno2006,Pan2007,Berti:2006wq} or additional phenomenological terms in the EOB effective potential \\cite{Buonanno:2007pf} are needed to achieve acceptable phase differences between the numerical and analytical waveforms.  Our study is complementary to Ref.~\\cite{Pretorius:2007jn}, that considered sequences of eccentric, equal-mass, non-spinning binary black hole evolutions around the ``threshold of immediate merger'': a region of parameter space separating binaries that quickly merge to form a final Kerr black hole from those that do not merge in a short time. Similar scenarios have also been studied in Ref.~\\cite{Washik2008}, with particular regard to the maximal spin of the final hole generated in this way. The universality of the gravitational wave signal during the merger was analysed in Ref.~\\cite{Hinder2007}, where it was pointed out that binaries largely circularize after about 9 orbits when starting with eccentricities below about $0.4$. The first comparison between numerical evolutions of eccentric binaries with post-Newtonian predictions was presented in \\cite{Hinder2008}.  Our focus in this work is on the high-eccentricity region of the parameter space, which always leads to merger. In particular, the near-head-on limit of our study is of interest as a first step to compute the energy loss and production cross-section of mini-black holes in TeV-scale gravity scenarios (possibly at the upcoming LHC \\cite{Giddings:2007nr}), and trans-Planckian scattering in general \\cite{Dray:1984ha,Giddings:2007bw}.  Present semi-analytical techniques (including a trapped surface search in the union of Aichelburg-Sexl shock waves, close-limit approximation calculations and perturbation theory) only give rough estimates of the emitted energy and production cross-section \\cite{Cardoso:2005jq} and do not provide much insight into the details of the process (but see \\cite{Sperhake:2008ga} for a first numerical investigation).  Our main finding is that, for all sequences we studied, the motion radically changes character when the black holes' % initial momentum $P\\sim P_{\\rm crit}\\simeq 0.08-0.09M$ and the orbital angular momentum $L\\sim L_{\\rm crit}\\simeq 0.8 M^2$, turning from an eccentric inspiral into a plunge.  In particular, for $L\\lesssim L_{\\rm crit}$ we observe that:  \\begin{itemize} \\item The number of orbits $N_{\\rm waves}$ (as estimated using the gravitational wave cycles) or $N_{\\rm punc}$ (as computed from the punctures' trajectories) becomes less than one, so the motion effectively turns into a plunge (see Table \\ref{tab: models} and Fig.~\\ref{fig: traj} below);  \\item The energy emission starts decreasing exponentially (Fig.~\\ref{fig: El});  \\item PN-based eccentricity estimates yield meaningless results (Table \\ref{tab: models});  \\item The polarization becomes linear rather than circular (Fig.~\\ref{fig: pol});  \\item The final angular momentum starts {\\it decreasing}, rather than increasing, as $P$ and $L$ decrease (Fig.~\\ref{fig: jfin}). \\end{itemize}  Binary mergers with $L\\simeq L_{\\rm crit}$ are those producing the largest Kerr parameter for the final black hole observed in our simulations, $j_{\\rm fin}\\simeq 0.724$. One is led to suspect that for maximally spinning holes having spins aligned with the orbital angular momentum, a large orbital eccentricity may lead to violations of the cosmic censorship conjecture. Using arguments based on the extrapolation of point-particle results (see also \\cite{Buonanno:2007sv}), we conjecture that (1) the maximal Kerr parameter that can be obtained by any merger of non-spinning holes is not much larger than $j\\simeq 0.724$, and (2) cosmic censorship will {\\it not} be violated as a result of any merger, even in the presence of orbital eccentricity. Further numerical simulations are needed to confirm or disprove these conjectures.  The paper is organized as follows. We begin in Sec.~\\ref{eccentricity} discussing to what extent the Newtonian concept of eccentricity can be generalized to characterize orbiting binaries in general relativity. For this purpose, we introduce and compare various PN estimates of the orbital eccentricity, and we show that these eccentricity estimates break down when the motion turns from inspiral to plunge.  Sec.~\\ref{numerics} is a brief introduction to the numerical code used for the simulations. After a discussion of the choice of initial data and of the code's accuracy, we show how reducing the orbital angular momentum affects the gravitational waveforms, the puncture trajectories and the polarization of the waves.  In Sec.~\\ref{energyj} we study the multipolar energy distribution of the radiation and the angular momentum of the final Kerr black hole. In Sec.~\\ref{BKL} we show that the salient features of our simulations can be understood using extrapolations of point-particle results.  Sec.~\\ref{QNMs} is devoted to fits of the ringdown waveform and to estimates of the energy radiated in ringdown waves. We conclude by considering possible future extensions of our investigation.  %\\clearpage  %============================================================================= ",
        "conclusions": "We have presented a study of the gravitational waveforms produced by sequences of equal-mass, non-spinning black hole binaries. For each sequence, the binding energy of the system is kept constant and the orbital angular momentum is progressively reduced to zero, producing orbits of increasing eccentricity and eventually a head-on collision. We find that the motion transitions from inspiral to plunge when the orbital angular momentum $L=L_{\\rm crit}\\simeq 0.8M$. For $L<L_{\\rm crit}$ the binary always completes less than $\\sim 1$ orbit, and PN estimates of the orbital eccentricity are no longer meaningful. As the initial momentum of the holes $P/M\\to 0$ the polarization quickly becomes linear, rather than circular, and the radiated energy drops (roughly) exponentially. For equal-mass, non-spinning binaries, orbits with $L\\simeq L_{\\rm crit}$ produce the largest dimensionless Kerr parameter for the remnant, $j_{\\rm fin}=J/M^2\\simeq 0.724\\pm0.13$ (to be compared with the Kerr parameter $j_{\\rm fin}\\simeq 0.69$ resulting from quasi-circular inspirals). These results are quite insensitive to the initial separation of the holes, and they can be understood using extrapolations from black hole perturbation theory. % It will be interesting to improve the accuracy of our predictions by %using higher resolution, larger extraction radii and especially larger initial %orbital separations. Larger separations, as used in sequence 3, will be necessary to perform accurate comparisons with PN predictions for the evolution of eccentric binaries (see \\cite{Hinder2008}). Such an analysis is beyond the scope of this work, however, and a corresponding analysis of the sequence 3 data will be presented elsewhere.  For equal masses, we found that eccentric binary mergers with $L\\simeq L_{\\rm crit}$ maximize the Kerr parameter of the final black hole. Using arguments based on point-particle extrapolations, we proposed two conjectures: (1) $j_{\\rm fin}\\simeq 0.724\\pm 0.13$ should be close to the largest Kerr parameter that can be produced by {\\it any} non-spinning black hole binary merger, independently of the binary's mass ratio; and (2) even if we consider maximally spinning holes with spins aligned with the orbital angular momentum, orbital eccentricity should {\\it not} lead to violations of the cosmic censorship conjecture. It will be interesting to check these conjectures using numerical simulations of eccentric binaries with unequal masses and non-zero spins.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0710/0710.0795_arXiv.txt": {
        "abstract": "Two aspects of our recent N-body studies of star clusters are presented:\\\\ 1) What impact does mass segregation and selective mass loss have on integrated photometry?\\\\ 2) How well compare results from N-body simulations using NBODY4 and STARLAB/KIRA? ",
        "introduction": "``Mass segregation'' describes the effect that high-mass stars are preferentially found in the center of clusters, while the outskirts are preferentially occupied by lower-mass stars.  (Primordial) mass segregation is studied observationally in a number of clusters. Some examples are: the Orion Nebula  Cluster (\\cite{1998ApJ...492..540H}), 6 LMC clusters (\\cite{2002MNRAS.331..245D}).  Dynamical mass segregation is also found from N-body simulations (e.g. \\cite{1985ApJ...292..339I,1997MNRAS.286..709G}), caused by two-body encounters and energy equipartition. A number of authors also point at the preferential mass loss of low-mass stars when the cluster evolves in a tidal field (e.g. \\cite{1975ApJ...201..773S,1997MNRAS.286..709G}), as mass segregation populates the cluster outskirts preferentially with low-mass stars where they are most easily stripped from the cluster potential.  The resulting changes in the overall stellar mass function inside the cluster are quantified by \\cite{2003MNRAS.340..227B}.  In \\cite{2006A&A...452..131L} we use the results from \\cite{2003MNRAS.340..227B} to incorporate the changing mass function slope in a simplified manner into the GALEV code (see \\cite{2003A&A...401.1063A} and references therein) to calculate its impact on the integrated photometry of star clusters.  Our main findings are: \\begin{itemize}  \\item at 0 -- 40 per cent of the cluster's lifetime: the cluster colours are comparable to standard models (i.e. without preferential loss of low-mass stars)  \\item at 40 -- 80 per cent of the cluster's lifetime: a cluster appears too blue/young (compared to standard models) due to the loss of lower main-sequence stars  \\item at 80 -- 100 per cent of the cluster's lifetime: a cluster appears too red/old (compared to standard models) due to the loss of main sequence turn-off stars  \\item when interpreting photometry of mass-segregated clusters, that have preferentially lost low-mass stars, with standard photometric models (without taking mass-segregation effects into account) the derived ages can be wrong by 0.3 -- 0.5 dex  \\end{itemize} ",
        "conclusions": ""
    },
    "0710/0710.0089_arXiv.txt": {
        "abstract": "{The Very Energetic Radiation Imaging Telescope Array System (VERITAS) is an array of four 12m diameter Imaging Atmospheric Cherenkov Technique (IACT) telescopes operated at the base of Mt. Hopkins in southern Arizona. The four-telescope experiment started operation in April, 2007. GeV and TeV gamma-ray observations of blazars can be used to probe the structure and composition of their jets, and to contribute to our understanding of how supermassive black holes accrete matter. In this contribution, we present first VERITAS blazar results obtained with three and four telescopes.} \\begin{document} ",
        "introduction": "The EGRET (Energetic Gamma Ray Experiment Telescope) detector on board the Compton Gamma-Ray Observatory discovered strong MeV $\\gamma$-ray emission from 66 Active Galactic Nuclei (AGNs), mainly from Flat Spectrum Radio Quasars and Flat Spectrum Radio Sources \\cite{1999yCat..21230079H}. As of writing these proceedings in May 2007, ground-based Cherenkov telescopes have discovered TeV $\\gamma$-ray emission from 17 AGB \\cite{2004NewAR..48..367K}. Sixteen of the 17 sources are blazars and one is the radio galaxy M 87 \\cite{2006Sci...314.1424A,2003A&A...403L...1A}. The blazars are mainly high energy peaked BL Lac objects, with BL Lac itself (an intermediate peaked BL Lac) being the only exception \\cite{Albe:07}. The redshifts of the TeV $\\gamma$-ray sources range from $z\\,=$ 0.031 for Mrk 421 \\cite{1992Natur.358..477P} to $z\\,=$ 0.188 for 1ES 0347-121 \\cite{2006Natur.440.1018A}.  In this contribution, we will give an overview of the blazar observations performed with the VERITAS experiment. VERITAS is an array of four 12~m diameter Cherenkov telescopes located at an altitude of 1268 m above sea level on Mt. Hopkins, Az (31$^{\\circ}$ 40' 30.21\" N, 110$^{\\circ}$ 57' 07.77\" W) \\cite{2002APh....17..221W}. The experiment started operation with two telescopes in spring 2006, and with four telescopes in winter 2006. The telescope system achieves an angular resolution of 0.16$^{\\circ}$ and a 250 GeV-1 TeV $\\nu F_{\\nu}$ sensitivity of 10$^{-12}$ ergs cm$^{-2}$ s$^{-1}$ for 10 hours of integration. For a detailed description of the status and performance of the telescope system the reader is referred to the contributions of Meier et al. \\cite{Meie:07} and Celik et al. \\cite{Celi:07} in this volume. The most important blazar detections are described in dedicated contributions, please see Fortin et al. \\cite{Fort:07} for the 1ES 1218+304 results, Cogan et al. \\cite{Coga:07} for the 1ES 0806+524 and 1ES 647+250 result, Fegan et al. for Mrk 421 and Mrk 501 results \\cite{Fega:07}, and Colin et al. \\cite{Coli:07} for the M 87 results. The Whipple 10 m Cherenkov telescope is used to monitor blazars on a regular basis. The results of the observations taken in 2006 are described by Steele et al. \\cite{Stee:07}. \\begin{table}[t] \\begin{center} \\begin{tabular}{p{1.4cm}p{1.4cm}p{3.1cm}} \\hline Observatory \t& Wavelength \t& Contact\\\\ \\hline { Owen V.}\t& Radio \t& A. Readhead\\\\ Metsahovi \t& Radio \t& A. Lahteenmaki\\\\ WEBT\t\t& Radio/IR/O \t& M. Villata\\\\ Abastumani\t& Opt. \t\t& O. Kurtanidze\\\\ Antipodal\t& Opt.  \t& J. Buckley\\\\ Bell \t\t& Opt.  \t& M. Carini\\\\ Boltwood  \t& Opt.  \t& P. Boltwood\\\\ Bordeaux\t& Opt.\t     \t& P. Charlot\\\\ Tuorla   \t& Opt. \t\t& A. Sillanpaa\\\\ WIYN 0.9m  & Opt.          & T. Montaruli\\\\ Swift\t\t& X-ray \t& H. Krimm\\\\ AGILE\t\t& $\\gamma$-ray \t& M. Tavani\\\\ GLAST\t\t& $\\gamma$-ray \t& J. McEnery\\\\ MAGIC\t\t& $\\gamma$-ray  & D. Mazin\\\\ H.E.S.S.\t& $\\gamma$-ray\t& S. Wagner\\\\ IceCUBE\t\t& Neutrino \t& T. Montaruli\\\\ \\hline \\end{tabular} \\caption{List of VERITAS multiwavelength collaborators. Only one contact person is given for each observatory.} \\end{center} \\end{table} ",
        "conclusions": "The VERITAS AGN program is fully underway. The program includes intensive multiwavelength observations of blazars in a flaring state, deep observations of blazars to determine their energy spectra with high accuracy, and the search for TeV $\\gamma$-ray emission from a wide range of different types of blazars. The VERITAS collaboration is working together with a large number of observers to sample the spectral energy distribution of blazars along the entire electromagnetic energy spectrum, and to obtain complementary information through the detection of high-energy neutrinos. Agreements have been reached to assure a fruitful collaboration between the three Cherenkov telescope experiments VERITAS, MAGIC, and H.E.S.S..  The first observations have resulted in the highly significant detection of the blazars Mrk 421, Mrk 501, 1ES 1218+304, and M 87. In this contribution, we have described 3-telescope observations of a number of HBL, IBL, and FSRQs. The flux upper limits are between 2.2\\% and 8.6\\% of the flux from the Crab Nebula. We anticipate exciting results with the full VERITAS system of 4 telescopes."
    },
    "0710/0710.4157.txt": {
        "abstract": "We report new L and T dwarfs found in a cross-match of the SDSS Data Release 1 and 2MASS.   Our simultaneous search of the two databases effectively allows us to relax the criteria for object detection in either survey and to explore the combined databases to a greater completeness level. We find two new T dwarfs in addition to the 13 already known in the SDSS DR1 footprint.  We also identify 22 new candidate and bona-fide L dwarfs, including a new young L2 dwarf and a peculiar L2 dwarf with unusually blue near-IR colors---potentially the result of mildly sub-solar metallicity.  These discoveries underscore the utility of simultaneous database cross-correlation in searching for rare objects.  Our cross-match completes the census of T dwarfs within the joint SDSS and 2MASS flux limits to the $ \\approx$97\\% level.  Hence, we are able to accurately infer the space density of T dwarfs.  We employ Monte Carlo tools to simulate the observed population of SDSS DR1 T dwarfs with 2MASS counterparts and find that the space density of T0--T8 dwarf systems is $0.0070_{-0.0030}^{+0.0032} $~pc$^{-3}$ (95\\% confidence interval), i.e., about one per 140~pc$^3$.  Compared to predictions for the T dwarf space density that depend on various assumptions for the sub-stellar mass function, this result is most consistent with models that assume a flat sub-stellar mass function $d N/d M \\propto M^ {0.0}$.   No $>$T8 dwarfs were discovered in the present cross-match, though less than one was expected in the limited area (2099~deg$^2$) of SDSS DR1. ",
        "introduction": "Our knowledge of the properties of ultra-cool L and T dwarfs has increased dramatically over the past decade as a result of the completion of several large-area optical and near-IR imaging surveys, and the implementation of fast computerized access to survey databases.  L and T dwarfs are readily identified in imaging surveys by their characteristic red optical minus near-infrared (near-IR) colors. There are now hundreds of L dwarfs and over 100 T dwarfs known\\footnote{A database of known L and T dwarfs is maintained at http://DwarfArchives.org \\citep{kirkpatrick03, gelino_etal04}.}, the vast majority of which have been found in the Two-Micron All-Sky Survey \\citep[2MASS;][]{skrutskie_etal06} and in the Sloan Digital Sky Survey \\citep[SDSS;][]{stoughton_etal02}.  The large number of L and T dwarfs identified in these two uniform and well-characterized data sets allows detailed investigations of the population properties of sub-stellar objects, namely their mass and luminosity functions and their multiplicity.  A detailed investigation focusing on a flux-limited sample of field L dwarfs has already been presented in \\citet{cruz_etal07}.  However, a similarly comprehensive empirical investigation of field T dwarfs has not been performed yet.  The most detailed study of T dwarfs to date is the 2MASS T5--T8 dwarf survey of \\citet{burgasser02}.  \\citeauthor{burgasser02}'s focus on the T5-- T8 sub-range was driven by their characteristic {\\sl blue} near-IR colors ($J-K_S\\sim0$~mag) that set them apart from the majority of main-sequence stars in 2MASS.  T0--T4 dwarfs, on the other hand, have red to neutral near-IR colors (2.0~mag~$\\gtrsim J-K_S \\gtrsim0.5$~mag) and searches for them face a vast contamination by background low-mass stars.  As a result, our understanding of the field T0--T4 population has lagged.  Although a number of T0--T4 dwarfs have been identified in the optical in SDSS \\citep[][and references therein]{geballe_etal02, knapp_etal04, chiu_etal06}, an adequate analysis of the population of early T dwarfs is still lacking.  Accurate knowledge of the number density of early T dwarfs relative to those of late L and mid T dwarfs is important for studies aimed at constraining the time scale of dust sedimentation and cloud formation in sub-stellar photospheres at the L/T transition.  Completing the census of known T dwarfs to allow such studies is the primary science motivation for the present study.  With hundreds of L and T dwarfs now known, a small number of peculiar L and T dwarfs have also emerged from the larger sample.  These unusual and rare objects are set apart from their counterparts either by having abnormal surface gravities \\citep[e.g.,][]{kirkpatrick_etal06, burgasser_etal06, cruz_etal07} or lower metallicities \\citep{burgasser_etal03b, burgasser04b}.  The recognition of such variety among the known L and T dwarfs has revealed a necessity for dimensional expansion of the present L and T dwarf classification schemes to include the effects of surface gravity and metallicity \\citep{kirkpatrick05}.   However, the number of known peculiar objects is presently too small to enable their accurate characterization; a larger sample will be needed to adequately anchor an expanded classification scheme.  The defining photometric characteristics of peculiar ultra-cool dwarfs that set them apart from the normal population, e.g., redder near-IR colors for young L dwarfs \\citep {kirkpatrick_etal06}, or redder optical and bluer near-IR colors for metal-poor ultra-cool sub-dwarfs \\citep{lepine_etal03b, burgasser_etal03b, cruz_etal07}, are only now being recognized.  Targeted photometric searches for such objects in the existing databases may be more fruitful in the near future. As a by-product of the present study, we remark on the characteristics of two peculiar L dwarfs discovered in our search.  Finally, the analysis of the late-T dwarf population of \\citeauthor{burgasser02} (\\citeyear{burgasser02}; see also \\citealt{burgasser04, burgasser07}, \\citealt{allen_etal05}) has shown that the number density of sub-stellar objects monotonically increases until the cool end of the present spectral type sequence (at T8; $T_{\\rm eff}\\approx750$~K), and is expected to continue increasing for even cooler objects. That is, brown dwarfs with spectral types $>$T8 are likely numerous, but have eluded detection in present large-area surveys because of being intrinsically faint.  The photospheres of extremely cool brown dwarfs, with effective temperatures below 400~K, are expected to have undergone a chemical transformation that is similar to the one occurring at the transition between the L and T spectral types, with the dominant source of opacity in the near-IR becoming water clouds, as opposed to methane clouds \\citep{burrows_etal03}.  Even cooler ($\\lesssim$200~K) brown dwarfs may have ammonia- dominated photospheres that are very similar to those of giant planets in the Solar System.  At very low effective temperatures, the emergent spectral energy distribution (SED) may be such that these objects may require a new spectral type (``Y'') for classification.  The discovery and characterization of such extremely cool brown dwarfs are among the primary science drivers for present and future deep large area surveys, e.g., with UKIRT \\citep[The UKIRT Infrared Deep Sky Survey;][]{lawrence_etal06}, with the Panoramic Survey Telescope and Rapid Response System \\citep[Pan-STARRS;][] {kaiser_etal02}, or with the Wide-field Infrared Survey Explorer \\citep[WISE;][]{mainzer_etal06}.   These large sensitive projects will undoubtedly dramatically expand our knowledge of sub-stellar objects at the bottom of the main sequence.  Nevertheless, it is possible that a small population of such extremely cool objects may already be present in the current generation of sky surveys.  Among the existing surveys, SDSS and 2MASS offer the best chance for finding brown dwarfs later than spectral type T8 because they cover the most volume.  Given their anticipated faintness, very red optical colors, and potentially blue near-IR colors, $>$T8 dwarfs may be present only as low signal-to-noise ($S/N$) single-band detections in SDSS (at $z$) and 2MASS (at $J$).  As such, they are more likely to have been overlooked or flagged as artifacts in either survey.  A combined consideration of the optical and near-IR data from SDSS and 2MASS may improve the chance for their discovery.  That is, a cross- correlation of the SDSS and 2MASS databases may allow us to not only probe deeper, but also cooler, than is possible in either survey alone.  Such a cross-correlation is the underlying approach of the present work.  The ability to cross-correlate large astronomical databases is one of the main technological goals of the National Virtual Observatory (NVO). In this paper we present results from a pilot project to test an implementation of this approach, focusing on the search of new brown dwarfs from a rapid cross- match of the 2MASS All-Sky Point Source Catalog (PSC) and SDSS Data Release 1 (DR1). The project was selected by the NVO as one of three demonstration research projects that would inform of the long-term hardware and software technology needs of the NVO.  The brown dwarf project in particular was aimed at identifying the technologies that will be needed to cross-match source catalogs at scale.  In the present paper we describe the implementation of our cross-matching technique (\\S~\\ref{sec_xmatch}), and report first results from the project, including identifications of two previously overlooked T dwarfs, a new peculiar L dwarf, and a young L dwarf in SDSS DR1 and 2MASS (\\S~\\ref{sec_results}).  We demonstrate that our dual-database cross-correlation search is more sensitive to T dwarfs than previous searches performed on SDSS or 2MASS alone, and take advantage of the high degree of completeness to T dwarfs attained in our search to estimate the T dwarf space density in the solar neighborhood (\\S~\\ref{sec_density}).  We discuss reasons for the omission of the newly identified T dwarfs in previous SDSS and 2MASS searches and draw lessons from our experience in cross-correlating large imaging databases  in \\S~\\ref{sec_discussion}. Finally, we outline the improved prospects for finding brown dwarfs cooler than spectral type T8 in a future iteration of the SDSS/2MASS cross-match using the much expanded Fifth Data Release (DR5) of the SDSS imaging survey (\\S~\\ref{sec_conclusion}). ",
        "conclusions": "\\label{sec_conclusion}  Our pilot project to search for previously overlooked T dwarfs in 2MASS and SDSS DR1 demonstrates the feasibility and utility of large database cross-correlation in discovering rare interesting objects. Our simultaneous positional and color cross-match of the 2MASS and SDSS DR1 databases uncovered 2 more T dwarfs in addition to the 13 already known in the SDSS DR1 footprint.  Despite the great scrutiny with which this area has already been explored for T dwarfs, both of the new T dwarfs had previously been overlooked, probably because of suspect photometry flags in SDSS.  The discovery of the two new T dwarfs demonstrates the superior sensitivity to ultra-cool dwarfs that can be attained by simultaneously cross-correlating large optical and near-IR databases, compared to searches based on individual optical or near-IR databases alone.  As a by-product of our search, which focused on objects with very red optical minus near-IR colors, we also report the discovery of two new peculiar L dwarfs: an L2 dwarf with unusually blue near-IR colors, potentially linked to mildly sub-solar metallicity, and another young L2 dwarf.  We took advantage of the high degree of completeness attained through our approach to obtain a flux- limited estimate of the local T dwarf space density.  We used Monte Carlo analysis to reproduce the observed T dwarf population in the overlap area of SDSS DR1 and 2MASS, and found that the local space density of T dwarfs is $0.0070_{-0.0030}^{+0.0032}$~pc$^{-3}$ (95\\% confidence interval), i.e., about one per 140 pc$^3$.  This empirical result is the first empirical estimate of the number density of T dwarfs over the entire range of T0--T8 spectral type range and extends earlier work by \\citet {burgasser02} that focused on T5--T8 dwarfs.  In the context of various predictions for the local sub- stellar population \\citep{burgasser04, burgasser07, allen_etal05}, we find that our result is most consistent with model-dependent estimates that assume a flat sub-stellar mass function, $d N/d M \\propto M^{0.0}$.  Given the success of the 2MASS/SDSS-DR1 cross-match, we expect that the approach will be instrumental for the identification of brown dwarfs cooler than the coolest ones presently known, with spectral types $>$T8.  While no such brown dwarfs were identified in the present cross-match covering the 2099~deg$^2$ area of SDSS DR1, we anticipate with a 86\\% probability that at least one T9 dwarf will be detectable in a similar cross-comparison of the entire 8000~deg$^2$ SDSS DR5 footprint with 2MASS."
    },
    "0710/0710.2335_arXiv.txt": {
        "abstract": "Gravitational waveforms from the inspiral and ring-down stages of the binary black hole coalescences can be modelled accurately by approximation/perturbation techniques in general relativity. Recent progress in numerical relativity has enabled us to model also the non-perturbative merger phase of the binary black-hole coalescence problem. This enables us to \\emph{coherently} search for all three stages of the coalescence of non-spinning binary black holes using a single template bank. Taking our motivation from these results, we propose a family of template waveforms which can model the inspiral, merger, and ring-down stages of the coalescence of non-spinning binary black holes that follow quasi-circular inspiral. This two-dimensional template family is explicitly parametrized by the physical parameters of the binary. We show that the template family is not only \\emph{effectual} in detecting the signals from black hole coalescences, but also \\emph{faithful} in estimating the parameters of the binary. We compare the sensitivity of a search (in the context of different ground-based interferometers) using all three stages of the black hole coalescence with other template-based searches which look for individual stages separately. We find that the proposed search is significantly more sensitive than other template-based searches for a substantial mass-range, potentially bringing about remarkable improvement in the event-rate of ground-based interferometers. As part of this work, we also prescribe a general procedure to construct interpolated template banks using non-spinning black hole waveforms produced by numerical relativity. ",
        "introduction": "\\label{sec:intro}  A network of ground based gravitational-wave (GW) detectors (LIGO, Virgo, GEO~600, TAMA) is currently collecting data, which a world-wide scientific collaboration is involved in analyzing.  Among the most promising sources detectable by these observatories are coalescing compact binaries consisting of black holes (BHs) and/or neutron stars spiraling toward each other as they lose orbital energy and angular momentum through gravitational-wave emission.  The gravitational-wave signal from coalescing binaries is conventionally split into three parts: inspiral, merger and ring down.  In the first stage, the two compact objects, usually treated as point masses, move in quasi-circular orbits (eccentricity, if present initially, is quickly radiated away). This part of the waveform is described very well by the post-Newtonian (PN) approximation of general relativity. In this approximation the Einstein equations are solved in the near zone (which contains the source) using an expansion in terms of the (small) velocity of the point masses. In the far zone, the vacuum equations are solved assuming weak gravitational fields, and these two solutions are matched in the intermediate region~\\cite{fock,Blanchet:1985sp,Wagoner:1976am}.  The PN approximation breaks down as the two compact objects approach the ultra-relativistic regime and eventually merge with each other. Although various resummation methods, such as Pad\\'e~\\cite{DIS98} and effective-one-body (EOB) approaches~\\cite{Buonanno:1998gg}, have been developed to extend the validity of the PN approximation, unambiguous waveforms in the merger stage must be calculated numerically in full general relativity. Recent breakthroughs in numerical relativity \\cite{Pretorius:2005gq,Campanelli:2005dd,Baker05a} have allowed many groups \\cite{Pretorius:2005gq,Campanelli:2005dd,Baker05a,Herrmann2006, Sperhake2006,Bruegmann:2006at,Thornburg-etal-2007a,Etienne:2007hr} to evolve BH binaries fully numerically for the last several orbits through the plunge to single BH formation. The field is now rapidly developing the capability to routinely evolve generic black-hole binary configurations in the comparable-mass regime, and to accurately extract the gravitational-wave signal. Important milestones include simulations of unequal-mass binaries and calculations of the gravitational recoil effect and the evolution of black-hole binaries with spin \\cite{Baker2006b,Gonzalez06tr,Campanelli2006b,Campanelli2006c,Herrmann:2007ac, Koppitz-etal-2007aa,Gonzalez:2007hi,Campanelli2007,Campanelli:2006fy, Pollney:2007ss,Rezzolla-etal-2007}.  Comparisons with post-Newtonian results are essential for data analysis efforts, and several groups have published results showing good agreement of various aspects of non-spinning simulations with post-Newtonian predictions (see e.g.~\\cite{Baker:2006yw,Baker:2006ha,Buonanno:2006ui,Berti:2007fi, Pan:2007nw,Ajith:2007qp,Hannam:2007ik,Boyle:2007ft}), and first results for certain configurations with spin have also become available \\cite{Brugmann:2007zj,Lousto:2007db}. In order to overcome phase inaccuracies in long evolutions, significant progress has been made by the Caltech-Cornell group using spectral codes \\cite{Scheel-etal-2006:dual-frame,Pfeiffer:2007}, and by the Jena group using higher (sixth) order finite differencing \\cite{Husa2007a}. Methods to reduce the eccentricity to around $10^{-3}$ (so far only for equal-mass binaries) have been presented by the Caltech-Cornell group \\cite{Pfeiffer:2007}, and the Jena group \\cite{Husa:2007ec} (using initial parameters from PN solutions that take into account radiation reaction). Current numerical waveforms can be generated for the last ($\\lesssim 10$) orbits, and these waveforms can be joined continuously with analytic PN inspiral waveforms to obtain one full signal. This was done in \\cite{Buonanno:2006ui,Pan:2007nw,Ajith:2007qp,Buonanno:2007pf}. Indeed, there are no fundamental obstructions to generating the whole waveform, including long inspiral over hundreds of orbits, by solving the full Einstein equations numerically. But, not only would this be computationally prohibitive with current methods, it is also unnecessary: the PN formalism is known to work very well in the weak-field regime (when the BHs are well-separated), and is a low-cost and perfectly adequate substitute to fully general relativistic solutions in that regime.  The numerically generated part of the gravitational-wave signal from coalescing binaries also includes the final stage of the coalescence, when a single perturbed black hole is formed and it rapidly loses its deviations from a Kerr black hole via gravitational waves. This part of the signal can be decomposed as a superposition of exponentially damped modes, and is called quasi-normal mode `ring down', by analogy with the vibrations of a bell. The detectable part of the ring down is rather short and only a few modes (if not only the dominant one) are expected to be important/detectable by initial ground-based observatories. This will not be true, however, for the advanced detectors \\cite{Berti:2007zu} and certainly it is not the case for LISA, the planned space-borne gravitational-wave observatory. Indeed, the majority of the signal-to-noise ratio (SNR) comes from the quasi-normal mode ringing of binary systems with a total mass above a few $10^6\\, M_{\\odot}$ \\cite{Berti:2005ys}.  For LISA, and also perhaps for the next generation of ground based detectors, it will be possible to detect several quasi-normal modes and test the `no hair' theorem, according to which all modes are functions of a BH's mass and spin~\\cite{Dreyer:2003bv,Berti:2005qd, Berti:2005ys}.  Joining analytically modeled inspiral with numerically generated merger and ring down allows us to produce the complete gravitational-wave signal from coalescing binaries, and to use it in the analysis of detector data. There are several benefits to using the whole signal in searches. The most obvious one is the increase in SNR in a fully coherent matched filtering search~\\cite{FlanHugh98,BuonD00, DIS01,Ajith:2007qp}. Increase in SNR implies increase in the event rate and improvement in the parameter estimation. Including the inspiral, merger and ring down parts in a template waveform also means that the waveform has a more complex structure. This extra complexity will also bring about some improvement in the parameter estimation~\\cite{ParamEstim} and possibly also a reduction in the false alarm rate in analysis of the data from the ground-based network of detectors.  This is because it is in general harder for the noise to mimic a complex signal\\footnote{At least we expect this to happen for those binaries for which both the inspiral and the merger contribute significantly to SNR.}.  For LISA, the detection of inspiralling super-massive black holes is not a problem; the SNR is expected to be so large that we expect some signals to be visible by eye in LISA data.  However, using the full signal for LISA data analysis is equally important because the full signal is essential in estimating parameters of the binary with the required accuracy. This is important not only from the astrophysical point of view, but also because we need to subtract loud signals from the data in order to detect/analyze other signals.  Imperfect signal removal due to errors in the parameter estimation will result in large residuals and will adversely affect subsequent analyses. Improved parameter estimation will also enable GW observations (in conjunction with electromagnetic observations) to constrain important cosmological parameters, most importantly the equation of state of dark energy~\\cite{Schutz86,Markovic:1993cr,ChernoffFinn:1993,HolzHugh05, Arun:2007hu,ParamEstim}.  The numerical waveforms described above are still computationally expensive and cannot be used directly to densely cover the parameter space of the binary BHs that will be searched over by matched filtering techniques. A promising alternative is to use the post-Newtonian and numerical-relativity waveforms to construct an analytic model that sufficiently accurately mimics a true signal \\cite{Ajith:2007qp, Buonanno:2007pf}. In~\\cite{Ajith:2007qp} we have suggested a phenomenological family of waveforms which can match physical signals from non-spinning binaries in quasi-circular orbits with fitting factors above 99\\%.  In this paper we extend this formulation to propose a two-parameter family of template waveforms which are explicitly parametrized by the physical parameters of the binary. We show that this two-dimensional template family is not only \\textit{`effectual'} in detecting the signals from binary BH coalescences, but also \\textit{`faithful'} in estimating the parameters of the binary. This family of template waveforms can be used to densely cover the parameter space of the binary, thus avoiding the computational burden of generating numerical waveforms in each grid point in the parameter space. We compute the effectualness and faithfulness (see Section~\\ref{sec:DA} for definitions) of the template family in the context of three different ground-based detectors: namely, Initial LIGO, Virgo and Advanced LIGO. We also compare the sensitivity of a search which coherently includes all three (inspiral, merger and ring down) stages of the BH coalescence with other template-based searches which look for each stage separately.  Our `target signals' are constructed by matching the numerical-relativity waveforms to a particular family (\\emph{TaylorT1} approximant~\\cite{DIS01}) of post-Newtonian waveforms, but this choice is by no means necessary. Indeed, we expect that more robust ways of constructing post-Newtonian approximants, such as the effective one-body approach~\\cite{Buonanno:1998gg} or Pad\\'e resummation approach~\\cite{DIS98}, will give better agreement with numerical-relativity (NR) waveforms. But the purpose of the current paper is to explicitly prescribe a general procedure to produce hybrid and phenomenological waveforms, and to construct interpolated template banks using parametrized waveforms.  We show that, given the number of numerical wave cycles we employ, even a simple PN choice like TaylorT1 leads to very faithful and effectual templates, and significantly increases the possible range of gravitational-wave searches. The use of improved PN approximants will require a smaller number of NR cycles, thereby further reducing computational cost for template construction.  There are also other approaches for comparing analytic and numerical waveforms and for constructing hybrid waveforms (see, for example \\cite{Pan:2007nw}); it would be interesting to compare the results presented in this work with other approaches presented in the literature.  The paper is structured as follows. In Section~\\ref{sec:numrelintro} we summarize the methods of current numerical-relativity simulations, including a setup of the initial data that allows an unambiguous comparison with post-Newtonian results, and the wave extraction techniques. In Section~\\ref{sec:DA} we briefly outline the waveform generation using the \\emph{restricted} post-Newtonian approximation. There we briefly introduce the main data-analysis techniques and define notations that are used in the subsequent sections.  In Section~\\ref{sec:phenomtemplates} we construct a phenomenological template family parametrized only by the masses of the two individual black holes. First we combine restricted 3.5PN waveforms \\cite{BDEI04} with results from NR simulations to construct `hybrid' waveforms for the quasi-circular inspiral of non-spinning binaries with possibly unequal masses. Then, we introduce a phenomenological family of templates constructed in the frequency domain. Initially the template family is parametrized by 10 phenomenological parameters. We then find a unique mapping of these 10 parameters to the two physical parameters: namely, the total mass $M$ and the symmetric mass ratio $\\eta \\equiv M_1 M_2 / M^2$, so that the template family is just two-dimensional. The resulting templates have remarkably high fitting factors with target waveforms. Here we also compute the faithfulness of the templates and the bias in the estimation of the parameter of the binary. A comparison of the sensitivity of the search using the proposed template family with other existing template-based searches is also presented. Finally, we summarize our main results in Section~\\ref{sec:summary}.  Some details of the calculations involved are described in Appendices~\\ref{app:fitfactor} and~\\ref{app:horDist}. We adopt geometrical units throughout this paper: $G=c=1$.   \\label{sec:DA}  In this Section we will introduce notation that will be used later in the paper and describe briefly the main data-analysis techniques currently used in gravitational wave astronomy.  \\subsection{Restricted post-Newtonian waveforms}  We use the restricted  PN waveform at mass-quadrupole order, which has a phase equal to twice the orbital phase up to highest available order in the adiabatic approximation, and amplitude accurate up to leading order.  The corresponding $\\hc$ is given by \\begin{equation} \\hc=\\frac{\\eta M}{r}v^2(t) e^{2\\rmi\\phi}\\left[(1+\\cos\\iota)^2 e^{-\\rmi \\varphi(t)} + (1-\\cos\\iota)^2 e^{\\rmi \\varphi(t)}\\right] \\label{eq:PNhOptOrient} \\end{equation} where $M \\equiv M_1+ M_2$ is the total mass, $\\eta \\equiv M_1 M_2/M^2$ is the symmetric mass ratio, $r$ is the observation radius, $\\iota$ is the inclination angle;  the quantity $v(t)$ is an expansion parameter, defined by $v=(M \\dot{\\varphi}/2)^{1/3}$ with $\\varphi(t)$ equal to twice the adiabatic orbital phase.  The {\\it waveform} seen by the detector is given  by \\be s(t) = 4\\,\\eta\\, \\frac{M}{r} A \\, v^2(t) \\cos[\\varphi(t)+\\varphi_0], \\ee where, for short-lived signals (i.e., with duration much shorter than the earth rotation time, as well as de-phasing time scale due to Doppler shifts induced by earth motion and rotation),  $A$ and $\\varphi_0$ are numerical constants depending on the relative position and orientation of the source relative to the detector, as well as the antenna pattern functions of the detector.  In PN theory, the adiabatic phase $\\varphi(t)$ is determined by the following ordinary differential equations (also called the \\emph{phasing formula}): \\begin{equation} \\frac{{\\rm d}\\varphi}{\\dt} = \\frac{2v^3}{M},\\ \\ \\ \\ \\frac{{\\rm d}v}{\\dt} = -\\frac{{\\cal F}(v)}{ME'(v)}. \\label{eq:phasing1} \\end{equation} In these expressions, $E'(v)={\\rm d}E(v)/{\\rm d}v$ where $E(v)$ is the binding energy (per unit mass) of the system, and ${\\cal F}(v)$ is the GW luminosity. $E(v)$ and ${\\cal F}(v)$ are computed as post-Newtonian expansions in terms of $v$~\\cite{Blanchet:LivRev}. Currently, the binding energy function $E(v)$ has been calculated to $v^6$ (3PN) accuracy by a variety of methods~\\cite{DJS2000, DJSequiv,BF00,BFeom,ABF01,BDE04,itoh1,itoh2}. The flux function ${\\cal F}(v),$ on the other hand, has been calculated to $v^7$ (3.5PN) accuracy \\cite{BFIJ02, BDEI04} up to now only by the  multipolar-post-Minkowskian method and matching to a post-Newtonian source~\\cite{Blanchet:LivRev}.  The inspiralling phase is usually pushed up to the point where the adiabatic evolution of circular orbits breaks down due to the lack of further stable circular orbits. In the test-mass limit, the last (or innermost) stable circular orbit (ISCO) can be computed exactly (at 6$M$ in Schwarzschild coordinates). For comparable-mass binaries, on the other hand, the ISCO cannot always arise unambiguously from PN theories.  In adiabatic models, the maximum-binding-energy condition (referred to as MECO, or the maximum binding energy circular orbit,~\\cite{Blanchet:2002}) can be used in place of the ISCO.  This condition is reached when the derivative of the orbital binding energy with respect to orbital frequency vanishes. As a consequence, in this paper, the waveforms are evolved in time up to MECO: $E'(v)=0$. It may be noted that the ISCO and MECO may not be physically meaningful beyond the test-mass limit, but they make convenient cutoff criteria. The appropriate region of validity of PN waveforms can only be determined by comparison with fully general relativistic results, such as the numerical simulations that we discussed earlier.  Given $E(v)$ and ${\\cal F}(v)$, one can construct different, but equivalent in terms of accuracy, approximations to the phasing by choosing to retain the involved functions or to re-expand them. Indeed, the different PN models which describe the GW signal from inspiralling binaries agree with each other in the early stages of inspiral; but start to deviate in the late inspiral. The classification and explicit form of various models is nicely summarized in~\\cite{DIS01}. In this paper we use PN waveforms obtained by numerically solving Eqs. (\\ref{eq:phasing1}), called the {\\it TaylorT1} approximant, to construct the `hybrid waveforms' (see Section~\\ref{sec:Matching}).  \\subsection{Introduction to matched filtering} Since we can model the signal reasonably well, it is natural to employ matched filtering (which is the optimal detection strategy for a signal of known shape in the stationary Gaussian noise) to search for the gravitational-wave signal.  Suppose the detector's data $x(t)$ contains noise $n(t)$, and possible signal $s(t)$, i.e., $x(t) = n(t) + s(t)$.  Assuming $n$ to be stationary Gaussian noise, it is convenient to work in the Fourier domain, because the statistical property of the noise is completely characterized by its power spectral density $S_n(f)$, which is given by (here we use a {\\it single-sided} spectrum) \\begin{equation} \\langle\\tilde{n}(f)\\tilde{n}^*(f')\\rangle = \\frac1{2}S_n(f)\\,\\delta(f-f')\\,, \\end{equation} where $\\tilde n(f)$ is the Fourier Transform of $n(t)$ \\begin{equation} \\tilde{n}(f) \\equiv \\int_{-\\infty}^{\\infty}n(t) e^{-2\\pi \\rmi   ft}\\;\\dt\\,, \\end{equation} and $\\langle\\ldots\\rangle$ denotes taking the expectation value. Based on the detector noise spectrum, we introduce a Hermitian inner product: \\begin{equation} (g | h) \\equiv 2 \\int_0^{\\infty} \\frac{\\tilde{g}^*(f)\\tilde{h}(f) + \\tilde{g}(f)\\tilde{h}^*(f)}{S_n(f)}\\;\\df\\,. \\end{equation} For the data $x$ with known signal $s$, the optimal detection statistic is given by applying a template $h$ with the same shape as $s$, or $h = \\alpha s$: \\begin{equation} \\rho_{\\rm opt} \\equiv ( x | h)\\,. \\end{equation} The detectability of the signal is then determined by the SNR of $\\rho_{\\rm opt}$, \\begin{equation} \\frac{S}{N} = \\left. \\frac{(s|h)}{\\sqrt{\\langle(h|n)(n|h)\\rangle}}\\right|_{h=\\alpha s} = (s|s)^{1/2}. \\end{equation} (Note that the SNR does not depend on the overall normalization of $h$.) In case the template $h$ is not exactly of the same shape as $s$, the SNR will be reduced to \\begin{equation} \\frac{S}{N} = (s|s)^{1/2}\\mathcal{M}\\,, \\end{equation} where $\\mathcal{M}\\le 1$ is the {\\it match} of the template to the signal, given by \\begin{equation} \\mathcal{M}[s,h]\\equiv \\frac{(s|h)}{\\sqrt{(s|s)\\,(h|h)}} \\equiv (\\hat s|\\hat h)\\,, \\end{equation} and where a hat denotes a normalized waveform.  For more details, we refer the reader to Ref.~\\cite{Cutler:1994ys}.  \\subsection{Template banks, effectualness and faithfulness}  We now consider the more realistic problem of attempting to detect a family of waveforms $s(\\vectheta)$, parametrized by a vector of physical parameters $\\vectheta \\in \\Theta$, using a family of templates $h(\\veclambda)$ parametrized by a vector of parameters $\\veclambda\\in\\Lambda$.  We first introduce the concepts of {\\it physical template bank} and {\\it phenomenological template bank}. Roughly speaking, physical template banks are constructed from well-motivated physical models (e.g., approximation up to a certain order)~ \\cite{Babak:2006ty}, while phenomenological banks are constructed in an ad-hoc manner to mimic the desired physical signals with high accuracy.  For physical banks, the vectors $\\vectheta$ and $\\veclambda$ consists of the same set of {\\it physical parameters}, while for phenomenological banks, the vector $\\veclambda$ usually contains {\\it phenomenological parameters,} which can be larger or smaller in number than the physical parameters.  Two phenomenological template families~\\cite{BCV, BCV2} are used currently in the search for BH binaries in LIGO data~\\cite{Abbott:2007xi,Abbott:2005kq}.  They each represent a different motivation for introducing phenomenological banks: (i) when we have uncertainty in the signal model, we can produce a template bank with larger detection efficiency by introducing extra (phenomenological) parameters (BCV1, \\cite{BCV}) so that $\\textrm{dim}(\\Lambda) > \\textrm{dim}(\\Theta)$; (ii) when the true signal depends on too many parameters and is too difficult to search over, it is sometimes possible to come up with a model with fewer (phenomenological) parameters ($\\textrm{dim}(\\Lambda) < \\textrm{dim}(\\Theta)$) and still high fitting factors (BCV2, \\cite{BCV2}).  The detection efficiency of a template bank towards a specific signal $s(\\vectheta)$ can be measured by the threshold SNR above which the detection probability exceeds a certain minimum (usually 50$\\%$), while the false-alarm probability is kept below a certain maximum (usually 1$\\%$ for one-year data). The threshold value depends (logarithmically, in the case of Gaussian noise) on the number of statistically independent templates, and (inverse-proportionally) on the {\\it fitting factor} (FF)~\\cite{Apostolatos:1995pj}: \\begin{equation} {\\rm FF}[h;\\vectheta] \\equiv \\max_{\\veclambda}\\mathcal{M}[s(\\vectheta),h(\\veclambda)] \\equiv \\mathcal{M}[s(\\vectheta),h(\\veclambda_{\\rm max})]  \\,. \\end{equation} A bank with high FF is said to be \\emph{effectual}~\\cite{DIS98,DIJS03}. Typically, we require that the total mismatch between the template and true signal (including the effects of both the fitting factor and the discreteness of the template bank) to not exceed $3\\%$.  We shall see that this requirement is easily met by our template bank.   \\begin{figure*} \\begin{center}  \\psfrag{Pmap}{$P$} \\psfrag{PMap2D}{$P_{2\\rm D}$} \\psfrag{Pt}{$P(\\Theta)$} \\psfrag{Pint}{$P_{\\rm int}(\\Theta)$} \\psfrag{BT}{$\\vectheta$} \\psfrag{BLM}{$\\veclambda_{\\rm max}$} \\psfrag{BLMP}{$\\veclambda_{\\rm max^\\prime}$} \\psfrag{BLMInt}{$\\veclambda_{\\rm int}$} \\psfrag{Th}{$\\Theta$} \\psfrag{(i)}{(i)} \\psfrag{(ii)}{(ii)} \\psfrag{(iii)}{(iii)}  \\includegraphics[width=17.5cm]{TemplSpace_1.eps} \\caption{Construction of the phenomenological template bank: (i) mapping physical signals (solid curve) into a sub-manifold (dashed curve, with example templates marked by dots) of a larger-dimensional template bank (curved surface), (ii) obtaining a lower-dimensional phenomenological bank with the same number of parameters as physical parameters, through interpolation (solid curve on the curved surface, with example templates marked by triangles), and (iii) Estimating the bias of the lower-dimensional interpolated bank by mapping physical signals into the bank (with images of example signals marked by dots).} \\label{fig:templSpaceCartoon} \\end{center} \\end{figure*}   It is natural to associate every point $\\vectheta$ in the physical space $\\Theta$ with the best matched point $\\veclambda_{\\rm max}\\in\\Lambda$. This leads to a mapping $P:\\Theta \\mapsto \\Lambda$ defined by \\begin{equation} \\label{def:P} P (\\vectheta) = \\veclambda_{\\rm max} \\,. \\end{equation} This mapping will play a key role in the construction of our template bank.  We will assume the mapping $P$ to be single-valued, i.e., given a target signal, the best-matched template is unique.  We depict this mapping schematically in the left panel of Fig.~\\ref{fig:templSpaceCartoon}.  For a physical template bank with $\\vectheta$ and $\\veclambda$ the same set of parameters (which we use $\\vectheta$ to denote), it is {\\it most convenient} to identify the best-match parameter $\\vectheta_{\\rm max}$ as the estimation of the original parameter $\\vectheta$. In general this will lead to a systematic bias \\begin{equation} \\Delta\\vectheta = \\vectheta_{\\rm max}-\\vectheta = P(\\vectheta)-\\vectheta\\,. \\end{equation} A bank with a small bias (as defined above) is said to be \\emph{faithful}~\\cite{DIS98,DIJS03}.   However, if we assume no uncertainty in the true waveforms (thereby excluding the case of BCV1), then as long as $P$ is invertible, a non-faithful physical or phenomenological bank can always be {\\it converted into} a faithful bank by the re-parametrization \\begin{equation} \\label{faithful} h_{\\rm faithful}(\\vectheta) \\equiv h \\circ P(\\vectheta) \\end{equation} where we have used the standard notation $h \\circ P(\\vectheta) :=h (P(\\vectheta))$.  In other words, each template $\\veclambda$ in the image set of physical signals $P(\\Theta)$ is labeled by physical parameters $\\vectheta = P^{-1}(\\veclambda)$.  For this reason, we require $P$ to be invertible. It is quite conceivable that for physical banks, $P$ should be invertible, if the physical bank does not fail to describe the true waveforms too dramatically (and of course assuming the true waveform does contain independent information about the physical parameters $\\vectheta$).  In this way, {\\it all reasonable physical banks can be made faithful}.  By contrast, if for some phenomenological bank (e.g., BCV2 if we only take into account the intrinsic parameters of the bank), $P$ is a many-to-one map, with $P(\\vectheta_1)=P(\\vectheta_2)$ for some $\\vectheta_1 \\neq \\vectheta_2$. Then for a physical signal with parameter $\\vectheta_1$, the template bank $h_{\\rm faithful}$ would achieve the same best match at both $\\vectheta_1$ and $\\vectheta_2$, making physical parameter determination non-unique. In this case, we can simply keep using the phenomenological bank $h(\\veclambda)$; once a detection is made with $\\veclambda_{\\rm max}$, the a set of parameters $P^{-1} (\\veclambda_{\\rm max})$ would be the best knowledge we have about the physical parameters of the source. (In practice, statistical uncertainty also applies to $\\veclambda_{\\rm max}$.)  \\begin{figure*} \\begin{center} \\includegraphics[width=16cm]{MatchedWaves_AEIJena.eps} \\caption{NR waveforms (thick/red), the `best-matched' 3.5PN waveforms (dashed/black), and the hybrid waveforms (thin/green) from three binary systems. The top panel corresponds to $\\eta = 0.25$ NR waveform produced by the AEI-CCT group. The second, third and fourth panels, respectively, correspond to $\\eta = 0.25, 0.22$ and $0.19$ NR waveforms from produced by the Jena group. In each case, the matching region is $-750 \\leq t/M \\leq -550$ and we plot the real part of the complex strain (the `+' polarization).}. \\label{fig:timeDomWaveNRPNHyb} \\end{center} \\end{figure*} ",
        "conclusions": "\\label{sec:summary}  Making use of the recent results from numerical relativity we have proposed a phenomenological waveform family which can model the inspiral, merger and ring-down stages of the coalescence of non-spinning binary black holes in quasi-circular orbits. We first constructed a set of hybrid waveforms by matching the NR waveforms with analytical PN waveforms. Then, we constructed analytical phenomenological waveforms which approximated the hybrid waveforms. The family of phenomenological waveforms that we propose was found to have fitting factors larger than 0.99 with the hybrid waveforms.  We have also shown how this phenomenological waveform family can be parametrized solely in terms of the physical parameters ($M$ and $\\eta$) of the binary, so that the template bank is, in the end, two dimensional~\\footnote{It may be noted that, the mapping from the phenomenological to physical parameters might not be unique in the case of spinning binaries, because of the degeneracies of different spin configurations.}. This two dimensional template family can be explicitly expressed in terms of the physical parameters of the binary. We have estimated the `closeness' of this two-dimensional template family with the family of hybrid waveforms in the detection band of three ground-based GW detectors, namely Initial LIGO, Virgo and Advanced LIGO. We have estimated the effectualness (larger overlaps with the target signals for the purpose of detection) and faithfulness (smaller biases in the estimation of the parameters of the target signals) of the template family. Having both types of overlap always greater than 0.99, the two dimensional template family is found to be both effectual and faithful in the detection band of these ground-based detectors.  This phenomenological waveform family can be used to densely cover the parameter space, avoiding the computational cost of generating numerical waveforms at every grid point in the parameter space. We have compared the sensitivity of a search using this template family with other searches. For a substantial mass-range, the search using all three stages of the binary black hole coalescence was found to be significantly more sensitive than any other template-based searches considered in this paper. This might enable us to do a more sensitive search for intermediate-mass black holes using ground-based GW detectors.  A number of practical issues need to be addressed before we can employ this template family in an actual search for GW signatures. The first issue will be how to construct a bank of templates sufficiently densely spaced in the parameter space so that the loss in the event rate because of the mismatch between the signal and template is restricted to an acceptable amount (say, 10\\%). The explicit frequency domain parametrization of the proposed template family makes it easier to adopt the formalism proposed by Owen~\\cite{Owen:1995tm} in laying down the templates using a metric in the parameter space. Work is ongoing to compare the metric formalism adopted to the proposed template family and other ways of laying out the templates, for example a `stochastic' template bank~\\cite{StochTemplBank}. Also, this explicit parametrization makes it easier to employ additional signal-based vetoes, such as the `chi-square test'~\\cite{Allen:2004gu}. This will also be explored in a forthcoming work.  Since this template bank is also a faithful representation of the target signals considered, we expect that, for a certain mass-range, a search which coherently includes all three stages of the binary coalescence will bring about remarkable improvement in the estimation of parameters of the binary. This may be especially important for LISA data analysis in estimating the parameters of supermassive black hole binaries. This is also being explored in an ongoing work~\\cite{ParamEstim}.  It is worth pointing out that the family of target signals (the hybrid waveforms) that we have considered in this paper is not unique. One can construct alternate families of hybrid waveforms by matching PN waveforms computed using different approximations with NR waveforms. Also, owing to the differences in initial data and accuracy of numerical techniques, the NR waveforms from different simulations can also be slightly different. Thus, the coefficients listed in Tables~\\ref{tab:polCoeffsAmpParams} and~\\ref{tab:polCoeffsPhaseParams} have a unique meaning only related to this particular family of target waveforms. But we expect that the general parametrization that we propose  in this paper will hold for the whole family of non-spinning black hole coalescence waveforms from quasi-circular inspiral. As we have mentioned in the Introduction, the purpose of this paper is to explicitly prescribe a general procedure to construct interpolated template banks using parametrized waveforms which mimics actual signals from binary black hole coalescence (as predicted by numerical relativity and analytical methods).  Nevertheless, it may be noted that most of the PN waveforms constructed using different approximations are known to be very close to each other (see, for example,~\\cite{DIS01}). Also, we expect that NR waveforms from different simulations will converge as the accuracy of numerical simulations improves (see, for example,~\\cite{Baker:2007fb}). Thus, since different families of PN and NR waveforms, which are the `ingredients' for constructing our target signals, are very close to each other, we expect that the phenomenological waveform family proposed in this paper, in its present form, will be sufficiently close to other families of target signals for the purpose of detecting these signals. As a preliminary illustration of this, we have computed the fitting factors of the template waveforms with a different family of hybrid waveforms (constructed from longer and more accurate NR waveforms), and have shown that the overlaps are indeed very high. This will be explored in detail in a forthcoming work.  Also, we remind the reader that this paper consider only the leading harmonic of the GW signal ($\\ell=2,\\ m=\\pm 2$). We expect that the contribution from the higher harmonics become important for high mass ratios, which will be investigated in a forthcoming work."
    },
    "0710/0710.2103_arXiv.txt": {
        "abstract": "We present a new version of the fully 3D photoionization and dust radiative transfer code, {\\sc mocassin}, that uses a Monte Carlo approach for the transfer of radiation. The X-ray enabled {\\sc mocassin} allows a fully geometry independent description of low-density gaseous environments strongly photoionized by a radiation field extending from radio to gamma rays.  The code has been thoroughly benchmarked against other established codes routinely used in the literature, using simple plane parallel models designed to test performance under standard conditions. We show the results of our benchmarking exercise and discuss applicability and limitations of the new code, which should be of guidance for future astrophysical studies with {\\sc mocassin}. ",
        "introduction": "\\label{s:intro}  Photoionized environments characterize a wide range of astrophysical problems involving sources of X--radiation. With the advent of new technology used for instruments on board of (e.g.) {\\it XMM-Newton} and {\\it Chandra}, high resolution spectroscopy of such environments has become a reality. Paerels et al. (2000), for instance, observed the photoionized wind in Cygnus X--3 with {\\it Chandra} High Energy Transmission Grating Spectrometer (HETGS) showing the discrete emission to be excited by recombination in a tenuous X--ray--photoionized medium which is not symmetric with the source of the wind. Other examples include the detection of several recombination emission lines (from Fe~{\\sc xxvi} at 1.78~{\\AA} to N~{\\sc vi} at 29.08~{\\AA}), by Jimenez-Garate et al. (2005) in a 50~ks observation of the bright X-ray binary Hercules X-1 with {\\it Chandra} HETGS. We also note the {\\it Chandra} ACIS observations of the deeply eclipsing cataclysmic variable DQ Herculis by Mukai et al. (2003), who were able to pin down the origin of the soft X-rays from this system as being due to scattering of the unseen central X-ray source, probably in an accretion disk wind.  A number of 1D photoionization codes, including G.~Ferland's {\\sc cloudy} (Ferland et al. 1998) and T.~Kallman's {\\sc xstar} (Kallman \\& McCray 1980, Kallman \\& Bautista 2001) continue to represent powerful analytical tools for the analysis of astrophysical spectra from the X-ray to the infrared regime. These codes are designed for diffuse, optically thin media, which may also be irradiated by a non-thermal X-ray continuum, and have been, for instance, applied to the modelling of Narrow Line Regions (NLRs) of Active Galactic Nuclei (AGNs). Several other codes have been developed which specialize in emission and reflection spectra from optically thick hot photoionized media, irradiated by a non-thermal continuum extending to the hard X-ray region, such as X--ray irradiated accretion disks (e.g. Ross \\& Fabian, 1993; Nayakshin et al. 2000; Dumont et al. 2000).  To date, these and most other photoionization codes have numerically solved the equations of radiative transfer (RT) under the assumption of spherical symmetry or in plane parallel geometries, whereupon the problem is reduced to a 1D calculation.  While very few real X-ray sources are spherically symmetric, this approximation has been driven by the available computing power and the complexity of the multi-dimensional case. Mauche et al. (2004) developed a Monte Carlo code to investigate the radiation transfer of Ly$\\alpha$, He$\\alpha$, and recombination continua photons of H- and He-like C, N, O and Ne produced in the atmosphere of a relativistic black hole accretion disk. This code, however, while accounting for Compton scattering and photoabsorption followed by recombination, does not calculate the ionization state of the plasma. To our knowledge, there are currently no general, self-consistent and publicly available X-ray photoionization and dust RT codes capable of working in 3D. Although the 1D codes mentioned above are powerful tools for the analysis of the pan-chromatic spectra of numerous astrophysical environments, their application is restricted to rather simplified geometries.  The computational demand of realistic 3D simulations has recently come within the reach of low-cost clusters.  Taking advantage of this, the first self-consistent, 3D photoionization and dust RT code was developed for the IR-UV regime using Monte Carlo techniques (Ercolano et al. 2003a, 2005). The code, {\\sc mocassin} (MOnte CArlo SimulationS of Ionized Nebulae) was designed to build realistic models of photoionized environments of arbitrary geometry and density distributions, and can simultaneously treat the dust RT. The code can also treat illumination from multiple point- or arbitrarily extended sources. The fully parallel version of {\\sc mocassin} (documented and publicly available) is well-tested for classical nebulae, according to standard photoionization benchmarks (P\\'equignot et al., 2001), and has been successfully applied to the modeling of H~{\\sc ii} regions (e.g. Ercolano, Bastian \\& Stasi\\'{n}ska, 2007) and planetary nebulae (e.g. Ercolano et al. 2003b,c, 2004; Gon\\c{c}alves et al., 2006; Schwarz \\& Monteiro, 2006; Wright, Ercolano \\& Barlow, 2006). We now present a significantly improved version of the {\\sc mocassin} code (version 3.00) which extends to the X-ray regime. In the tradition of previous releases of the {\\sc mocassin} code, the X-ray version will also be made publicly available to the scientific community shortly after publication of this article\\footnote{Previous versions are available upon request from BE.}.  In Section~\\ref{s:code} we summarize the physical processes and atomic data added/modified in the new implementation and discuss the applicability and limitations of the code in its current form. In Section~\\ref{s:bench} we compare our code to established 1D codes for benchmark tests, comprising emission line spectra from model NLR and from three thin low-density slab models illuminated by a hard continuum. A brief summary is given in Section~\\ref{s:summary}. ",
        "conclusions": "\\label{s:summary}  We have presented a new version of the fully 3D Monte Carlo photoionization code, {\\sc mocassin}. The code was extended to allow the modelling of plasma irradiated by a hard continuum spanning from radio to gamma rays. The atomic data set of the code was also significantly updated and it is now synchronized with the latest release of the {\\sc chianti} database. The applicability and limitations of the new code were discussed, and the results of a thorough benchmarking exercise presented.  No major problems were found by the benchmark tests, although some minor differences have been found. We have highlighted a number of significant improvements in the atomic datasets available today compared to those available at the time the original benchmark tables were compiled. We provide here updated values and we emphasize the need of a new benchmarking exercise to be undertaken by the plasma modelling community.  The good performance of the {\\sc mocassin} code in all benchmark tests demonstrates that it is ready for application to real astrophysical problems. The public version of the X-ray enhanced {\\sc mocassin} code is available on request from the author."
    },
    "0710/0710.0324_arXiv.txt": {
        "abstract": "{Emission lines in quasars are believed to originate from a photoionized plasma. There are, however, some emission features which appear to be collisionally excited, such as the \\FeII\\ multiplet bands. Shortward of \\Lya, there also are a few permitted lines of species from low to intermediate ionization.} {\\ton\\ ($\\zq=1.928$) exhibits the steepest far-UV continuum decline known ($\\Fnu \\propto \\nu^{-5.3}$) shortward of 1050\\,\\AA. This object also emits unusually strong low to intermediate excitation permitted lines shortward of the Lyman limit.}{Using archive spectra of \\ton\\ from HST, IUE and Palomar, we measure the fluxes of all the lines present in the spectra and compare their relative intensities with those observed in composite quasar spectra.} {Our analysis  reveals unusual strengths with respect to \\Lya\\ of the following low to intermediate excitation permitted lines: \\OII+\\OIII\\ (835\\,\\AA), \\NIII+\\OIII\\ (686--703\\,\\AA) and \\NIII+\\NIV\\ (765\\,\\AA). We compare the observed line spectrum with both photoionization and shock models.}{Photoionization cannot reproduce the strengths of these far-UV lines. Shocks with $\\Vs \\simeq 100$\\,\\kms\\ turn out to be extremely efficient emitters of these lines and are favored as excitation mechanism. } ",
        "introduction": "\\label{sec:intro}   In this work, we analyze the emission lines of an unusual quasar, \\ton, which is alternatively named PG\\,1017+280 or J1019+2745 with redshift $\\zq=1.928$. It is severely deficient in ionizing photons, since its Spectral Energy Distribution (\\sed) shows a remarkable steepening of the continuum in the rest-frame far-UV, shortward of 1100\\,\\AA\\ (Binette \\& Krongold 2007, hereafter BK07; Binette et\\,al. 2007). If the far-UV is fitted by a powerlaw ($\\Fnu \\propto \\nu^{+\\alpha}$), the index\\footnote{Among the 77 quasars whose far-UV indices could be measured by Telfer et\\,al. 2002, there were only 3 objects with a continuum steeper than $\\nu^{-3}$.} is as steep as $\\nu^{-5.3}$. BK07 suggest that the extreme-UV flux might undergo a recovery shortward of 450\\,\\AA.  While the near-UV emission-line spectrum appears to be `normal', the far-UV spectrum shows low to intermediate ionization species with unusual strengths. Using the UV \\sed\\ constructed by BK07 from archive data, we will quantify this statement and present photoionization and shock models for comparison. The aim is to understand how the extreme deficiency of ionizing photons in \\ton\\ might be impacting the emission line spectrum.  The emission-line spectrum of quasar and Seyfert\\,I galaxies is generally believed to originate from gas photoionized by a nuclear UV source. State of the art photoionization models of the Broad Emission Line Region (BELR) such as those developed by Baldwin et\\,al. (1995) and dubbed `locally optimally emitting clouds' (LOC) models can successfully reproduce most of the emission lines observed in quasars. A grid of such models can be found in Korista et\\,al. (1997, hereafter KO97) and more recently in Casebeer et\\,al. (2006 and references therein). There are, however,  a few exceptions to the success of pure photoionization. In particular, photoionization models require micro-turbulences in order to reproduce the shape and intensity of the \\FeII\\ UV-band (Baldwin et\\,al. 2004). A possible alternative is that the region producing \\FeII\\ is collisionally ionized, as proposed by Grandi (1981, 1982), Joly (1987), V\\'eron-Cetty et\\,al. (2004, 2006) and Joly et\\,al. (2007). In this work, we present evidence that photoionization might not be sustainable in the case of some of the far-UV permitted lines reported in this paper. ",
        "conclusions": ""
    },
    "0710/0710.3866_arXiv.txt": {
        "abstract": "We present here X-ray spectra of the HMXB SMC~X-1 obtained in an observation with the \\xmm{} observatory beginning before eclipse and ending near the end of eclipse.  With the Reflection Grating Spectrometers (RGS) on board \\xmm{}, we observe emission lines from hydrogen-like and helium-like ions of nitrogen, oxygen, neon, magnesium, and silicon.  Though the resolution of the RGS is sufficient to resolve the helium-like $n$=2$\\to$1 emission into three line components, only one of these components, the intercombination line, is detected in our data.  The lack of flux in the forbidden lines of the helium-like triplets is explained by pumping by ultraviolet photons from the B0 star and, from this, we set an upper limit on the distance of the emitting ions from the star.  The lack of observable flux in the resonance lines of the helium-like triplets indicate a lack of enhancement due to resonance line scattering and, from this, we derive a new observational constraint on the distribution of the wind in SMC~X-1 in velocity and coordinate space. We find that the solid angle subtended by the volume containing the helium-like ions at the neutron star multiplied by the velocity dispersion of the helium-like ions must be less than 4$\\pi$\\,steradians\\,\\kms{}.  This constraint will be satisfied if the helium-like ions are located primarily in clumps distributed throughout the wind or in a thin layer along the surface of the B0 star. ",
        "introduction": "In isolated early-type stars winds are driven as ultraviolet photons from the stellar surface impart their outward momentum to the wind in resonance line transitions.  In a high-mass X-ray binary (HMXB), the wind is ionized by X-rays from the compact object.  If the X-radiation is intense enough, the resulting ions will not have transitions in the ultraviolet and this greatly affects the dynamics of the wind \\citep[see, e.g., ][ and references therein]{blo94}.  The SMC~X-1/Sk~160 system, which consists of a 0.71 second X-ray pulsar and a B0I star together in a 3.9-day orbit, is the most X-ray luminous known HMXB and, therefore, presumably, an extreme example of wind disruption by X-ray ionization.  The behavior of the wind of an early-type star under the influence of ionization from an X-ray emitting companion has been the subject of many theoretical studies.  For SMC~X-1, the most relevant and detailed such study is, arguably, the one by \\citet{blo95} which included a numerical hydrodynamic simulation.  The features that appeared in that simulation included a regular, UV-driven wind on the X-ray shadowed side of the star, a thermal wind on the X-ray illuminated side, transition regions between the two types of winds and dense, finger-like structures in the equatorial plane.  However, even this simulation includes significant approximations: the gravity of the compact object is not included and X-ray photoionization and its dynamical effects are treated in a very approximate way.  Because of the complexity of these systems and the difficulty of accounting for all of the relevant physics, observations are critical to characterizing the behavior of HMXB winds.  Though previous X-ray observations of SMC~X-1 have had spectral resolving powers ($\\equiv\\lambda/\\Delta\\lambda=E/\\Delta{}E$) of $\\sim$50 or less, X-ray spectroscopy has already revealed much about the wind.  X-ray spectroscopic observations of this resolving power can, through measurements of line fluxes, indicate the quantity and ionization level of X-ray emitting material.  In observations of SMC X-1 with \\asca{}, emission lines were not detected and upper limits were set on the quantity of material of intermediate ionization level in the wind of SMC~X-1, excluding the presence of finger-like structures as they appeared in the simulation of \\citeauthor{blo95} \\citep{woj00}.  Further observations of SMC~X-1 with the Advanced Camera for Imaging Spectroscopy (ACIS) on-board the {\\it Chandra X-ray Observatory}, however, have detected emission lines and, therefore, the presence of material of intermediate ionization in the wind of SMC~X-1 \\citep{vrt01}.  These results indicate that overdense regions do exist in the wind of SMC~X-1 and that while the simulations of \\citeauthor{blo95} may not be accurate in detail, the general form of the structure predicted by \\citeauthor{blo95} may, in fact, be present.  To develop further constraints on models of the wind in SMC~X-1, it is desirable to have constraints in addition to the quantity and ionization level of the X-ray emitting plasma.  High-resolution spectroscopic observations with resolving powers $\\gtrsim$200 have the potential to provide useful constraints on the kinematics of the high-ionization wind in HMXBs through direct measurements of Doppler line shifts and broadening.  In addition, the $n=2\\to1$ triplets of He-like ions are resolved at high resolution and measurements of the fluxes of the individual lines of these triplets provide a constraint on the structure and kinematics of the X-ray emitting material.  Of the three components of the helium-like triplet, only one --- the resonance line --- has a large oscillator strength and can be enhanced by resonant line scattering.  Because resonant line scattering saturates and because this saturation depends on the distribution of the scatterers in physical and velocity space, constraints on the distribution of the wind in physical and velocity space can be derived from the relative fluxes of these three lines \\citep{woj03}.  The flux ratios of the helium-like triplets are also affected by ultraviolet radiation \\citep*{blu72} and, therefore, the flux ratios measured from an HMXB constrain the distance from the photosphere of the high-mass star to the emission region.  Because of the location of SMC X-1 outside of the plane of the Galaxy, the column density of interstellar material to it is small and, unlike the HMXBs in the Galaxy, it can be observed in the wavelength range 15--35\\,\\AA.  The $n=2\\to1$ transitions of the hydrogen-like and helium-like ions of oxygen and nitrogen as well as hydrogen-like carbon lie in this range.  Therefore, unlike with the Galactic HMXBs, it is possible to derive observational constraints on the density and velocity distributions of the regions of moderate ionization in SMC~X-1 where these ions exist.  Because of the high spectral resolution (FWHM of $\\sim$65\\,m\\AA, corresponding to resolving powers in the range 230--540) and large effective area (total of $\\sim$80\\,cm$^2$) of the Reflection Grating Spectrometer on \\xmm{} in this band, the RGS on \\xmm{} is particularly well-suited to spectroscopic study there.  We present here observations of SMC~X-1 with \\xmm{} beginning before an eclipse and ending near the end of that eclipse.  We present an analysis in which we focus on the data from the RGS and derive constraints on the space and velocity distribution of material in the wind.  In \\S\\ref{sec:empir} we describe our observations and an empirical analysis of the line emission in which we measure line fluxes, shifts, and widths.  In \\S\\ref{sec:phys} we derive constraints on the distribution of the wind in physical and velocity space from the line fluxes we measure.  In \\S\\ref{sec:discuss}, we discuss the implications of our results for models of HMXB winds. ",
        "conclusions": "\\label{sec:discuss}  We have observed SMC~X-1 in eclipse with \\xmm{} and, with the RGS, resolved the helium-like triplets of nitrogen, oxygen, neon, and magnesium.  To our knowledge, this is the first time helium-like triplets from SMC~X-1 have been resolved and the first time the helium-like triplets of nitrogen, oxygen, or neon have been resolved for any HMXB.  In all cases, only one component of the triplet, the intercombination line, is observably present.  The lack of observable forbidden line fluxes in the triplet is easily explained by photoexcitation pumping of the forbidden line into the intercombination line by the ultraviolet radiation of the B0 star Sk~160 and allows upper limits to be set on distance of the emitting helium-like ions from Sk~160.  The absence of observable fluxes in the resonance lines is consistent with what we expect from recombination in the photoionized wind.  However, the fact that the resonance lines are not enhanced by resonant scattering implies a lower limit on the optical depth of the wind in the resonance line and, therefore, constraints on the structure and kinematics of the wind are also implied.  We have modeled the helium-like line emission as recombination and scattering emission from a region that is partly obscured by the companion star and is described by a single solid angle ($\\Omega$) and column density ($N$) and the velocities along lines of sight from the neutron star have a boxcar distribution with width $\\Delta{}v$ for each of the ions.  We have inferred infer the quantity $\\Omega\\Delta{}N$ from the sum of the flux of the intercombination and forbidden lines which are not affected by resonant line scattering.  We infer the quantity $N/\\Delta{}v$ from flux of the resonance line relative to the sum of the other two (the inverse of the $G$ ratio).  From the constraints on these two quantities, we also obtain a constraint on the quantity $\\Omega\\Delta{}v$.  Because we do not detect the resonance line, we obtain only lower limits on $N/\\Delta{}v$ and only upper limits on $\\Omega{}\\Delta{}v$.  As previously mentioned, in our analysis we have assumed that photons escape isotropically along lines from the location of their production or first scattering.  While we do expect recombination emission to be isotropic, resonant line scattering is not isotropic.  The angular distribution for resonant line scattering is a linear combination of an isotropic distribution and an dipole distribution \\citep[see, e.g.,][Ch.\\ 19]{cha60}.  Furthermore, for large optical depths that are comparable to or greater than unity such as we infer, photons may undergo several scatterings before escaping.  In this case, photons will escape preferentially along directions where the optical depth is least.  In general, the calculation of photon escape direction is also complicated by the fact that photons may travel large distances across the plasma before escaping.  While it is difficult to include this possibility in analytic calculations, it is straightforward using Monte Carlo techniques.  However, if the region that emits the helium-like line emission has supersonic velocity differentials, as we know that the bulk of the wind does, the transfer of resonant line photons is essentially local (the Sobolev approximation, \\citealt{cas70}) and can be approached analytically.  For a supersonically expanding wind, the optical depth in a direction given by the coordinate $x$ is inversely proportional to $dv_x/dx$.  If the flow moves along lines from a single point and $y$ is the distance from that point, then these derivatives, in directions parallel and perpendicular to the flow, respectively, are $dv/dy$ and $v/y$.  While these quantities will not generally be equal, they will generally be of the same order over the bulk of the wind and, therefore, if we assume complete redistribution in frequency and direction for individual scatterings, then the radiation will escape isotropically \\citep[see][]{cas70}.  In fact, individual scattering events do not redistribute frequency and direction completely and so even with $v/s=dv/ds$, radiation does not escape isotropically.  However the difference from the case of complete redistribution is no greater than 10\\%--20\\% (\\citealt{car72}, see also \\citealt{mih80}).  In light of these facts, we proceed to explore the implications of our results. However, it must be kept in mind that the approximations we have made for the radiative line transfer may quantitatively affect our conclusions.  The implication of our results on fundamental wind parameters, such as the mass-loss rate and the terminal wind velocity, is complex owing to factors such as the complex dependence of the ion fractions on density in photoionization balance.  Therefore, interpreting these results in terms of fundamental wind parameters is difficult to do without computing line fluxes for complete wind models which is beyond the scope of this work.  However, the upper limits on $(\\Omega/4\\pi)\\Delta{}v$ that we obtain, approximately 1\\,\\kms, are quite small compared to 1000\\,\\kms{}, the order of the terminal wind velocities of massive stars and the wind velocity in a simulation of the wind of SMC~X-1 by \\citet{blo95}.  Therefore, the part of the wind that emits the helium-like lines cannot be homogeneously distributed around the entire wind.  Instead, we believe that our results indicate either that the volume (or volumes) containing the helium-like ions of the various elements subtend a small solid angle at the neutron star or that the helium-like ions exist primarily in a region where the wind is at a small fraction of its terminal velocity or, perhaps, some combination of both.  One possible explanation of our results is that the helium-like ions exist primarily in dense clumps throughout the wind.  Dense clumps in the wind of an HMXB have previously been invoked in the case of Vela~X-1 by \\citet{sak99} in order to explain the strong fluorescence lines observed from that system.  Another possible explanation is that helium-like ions exist mainly near the surface of the star where the density is higher, the ionization less, and the velocity lower than in the outer parts of the wind \\citep[see, e.g.,][]{lie01}.  Furthermore, the part of a volume along the surface of the companion star with a thickness significantly less than one stellar radius that is visible during eclipse would subtend a solid angle significantly less than $4\\pi$ from the neutron star.  In Figure~\\ref{fig:clump_face} we illustrate both of these possibilities. \\begin{figure} \\begin{center} \\includegraphics[angle=-90,width=2.4in]{f6a.eps}\\hspace{0.5in} \\includegraphics[angle=-90,bb=-83 -27 295 263,width=2.4in]{f6b.eps} \\end{center} \\caption{In this figure we illustrate to possible distributions of the helium-like ions consistent with the helium-like line fluxes we measure.  In the first panel we illustrate the possibility that the helium-like ions are in clumps distributed throughout the wind.  In the second panel we illustrate the possibility that the helium-like ions are near the surface of the companion star.} \\label{fig:clump_face} \\end{figure} Definitive tests of these hypotheses would require calculations of the line emission from detailed wind models and that is beyond the scope of this work. However, we expect that our results, including the remarkably small values of $\\Omega\\Delta{}v$ are consistent with current expectations about the nature of winds in HMXBs."
    },
    "0710/0710.4488_arXiv.txt": {
        "abstract": "We construct evolutionary models of the populations of active galactic nuclei (AGN) and supermassive black holes, in which the black hole mass function grows at the rate implied by the observed luminosity function, given assumptions about the radiative efficiency and the luminosity in Eddington units. We draw on a variety of recent X-ray and optical measurements to estimate the bolometric AGN luminosity function and compare to X-ray background data and the independent estimate of Hopkins et al.\\ (2007) to assess remaining systematic uncertainties. The integrated AGN emissivity closely tracks the cosmic star formation history, suggesting that star formation and black hole growth are closely linked at all redshifts. We discuss observational uncertainties in the local black hole mass function, which remain substantial, with estimates of the integrated black hole mass density $\\rho_{\\bullet}$ spanning the range $3 - 5.5 \\times 10^5\\, {\\rm M_{\\odot}\\, Mpc^{-3}}$. We find good agreement with estimates of the local mass function for a reference model where all active black holes have a fixed efficiency $\\epsilon = 0.065$ and $L_{\\rm bol}/L_{\\rm Edd} \\approx 0.4$ (shifting to $\\epsilon = 0.09$, $L_{\\rm bol}/L_{\\rm Edd} \\approx 0.9$ for the Hopkins et al.\\ luminosity function). In our reference model, the duty cycle of $10^9 M_\\odot$ black holes declines from 0.07 at $z=3$ to 0.004 at $z=1$ and $10^{-4}$ at $z=0$. The decline is shallower for less massive black holes, a signature of ``downsizing'' evolution in which more massive black holes build their mass earlier. The predicted duty cycles and AGN clustering bias in this model are in reasonable accord with observational estimates. If the typical Eddington ratio declines at $z<2$, then the ``downsizing'' of black hole growth is less pronounced. Models with reduced Eddington ratios at low redshift or black hole mass predict fewer low mass black holes ($M_{\\bullet}\\lesssim 10^8\\, {\\rm M_{\\odot}}$) in the local universe, while models with black hole mergers predict more black holes at $M_{\\bullet}>10^9\\, {\\rm M_{\\odot}}$. Matching the integrated AGN emissivity to the local black hole mass density implies $\\epsilon = 0.075 \\times (\\rho_{\\bullet} / 4.5\\times 10^5\\, {\\rm M_{\\odot}\\, Mpc^{-3}})^{-1}$ for our standard luminosity function estimate, or 25\\% higher for Hopkins et al.'s estimate. It is difficult to reconcile current observations with a model in which most black holes have the high efficiencies $\\epsilon \\approx 0.16-0.20$ predicted by MHD simulations of disk accretion. We provide electronic tabulations of our bolometric luminosity function and our reference model predictions for black hole mass functions and duty cycles as a function of redshift. ",
        "introduction": "\\label{sec|intro}  The long-standing hypothesis that quasars are powered by accretion onto supermassive black holes (Salpeter 1964; Lynden-Bell 1969; Rees 1984) is now strongly supported by many lines of evidence, including the apparent ubiquity of remnant black holes in the spheroids of local galaxies (Richstone et al. 1998). The strong correlations between the masses of central black holes and the luminosities, dynamical masses, and velocity dispersions of their host spheroids (e.g., Magorrian et al. 1998; Ferrarese \\& Merritt 2000; McLure \\& Dunlop 2002; Marconi \\& Hunt 2003; H\\\"{a}ring \\& Rix 2004; Ferrarese \\& Ford 2005; Greene \\& Ho 2006; Graham 2007; Hopkins et al. 2007b; Shankar \\& Ferrarese 2008; Shankar et al. 2008) imply that the processes of black hole growth and bulge formation are intimately linked. Theoretical models typically tie black hole growth to episodes of rapid star formation, perhaps associated with galaxy mergers, and ascribe the black hole-bulge correlations to energy or momentum feedback from the black hole or to regulation of black hole growth by the bulge potential (e.g., Silk \\& Rees 1998; Kauffmann \\& Haehnelt 2000; Cavaliere \\& Vittorini 2000; Granato et al. 2004, 2006; Murray, Quataert, \\& Thompson 2004; Cattaneo et al. 2005; Miralda-Escud\\`{e} \\& Kollmeier 2005; Monaco \\& Fontanot 2005; Croton et al. 2006; Hopkins et al. 2006a, Malbon et al. 2006). These correlations also make it possible to estimate the mass function of black holes in the local universe (e.g., Salucci et al. 1999; Yu \\& Tremaine 2002; Marconi et al. 2004; Shankar et al. 2004, S04 hereafter). This local mass function provides important constraints on the co-evolution of the quasar and black hole populations. The most general and most well known of these constraints is the link between the integrated emissivity of the quasar population, the integrated mass density of remnant black holes, and the average radiative efficiency of black hole accretion (So{\\l}tan 1982; Fabian \\& Iwasawa 1999; Elvis et al. 2002).  In the paper we construct self-consistent models of the quasar population using a method that can be considered a ``differential'' generalization of So{\\l}tan's (1982) argument. Given assumed values of the radiative efficiency and the Eddington ratio $L/L_{\\rm Edd}$, the observed luminosity function of quasars at a given redshift can be linked to the average growth rate of black holes of the corresponding mass, and these growth rates can be integrated forward in time to track the evolution of the black hole mass function. This modeling approach has been developed and applied in a variety of forms by numerous authors, drawing on steadily improving observational data (e.g. Cavaliere et al. 1973; Small \\& Blandford 1992; Salucci et al. 1999; Cavaliere \\& Vittorini 2000; Yu \\& Tremaine 2002; Steed \\& Weinberg 2003, hereafter SW03; Hosokawa 2004; Yu \\& Lu 2004; Marconi et al. 2004; S04; Merloni 2004; Vittorini, Shankar \\& Cavaliere 2005; Tamura et al. 2006; Hopkins et al. 2007). The consensus of recent studies is that evolutionary models with radiative efficiencies of roughly 10\\% and mildly sub-Eddington accretion rates yield a reasonable match to the observational estimates of the remnant mass function. However, uncertainties in the bolometric luminosity function of active galactic nuclei (AGN, a term we will use to describe both quasars and less luminous systems powered by black hole accretion) and in the local black hole mass function remain an important source of uncertainty in these conclusions.  We begin our investigation by constructing an estimate of the bolometric AGN luminosity function. Our estimate starts from the model of Ueda et al. (2003, hereafter U03), based on data from several X-ray surveys, but we adjust its parameters based on more recent optical and X-ray data that provide more complete coverage of luminosity and redshift. Comparison to the X-ray background and to the independent luminosity function estimate of Hopkins, Richards \\& Hernquist (2007; HRH07 hereafter) gives an indication of the remaining uncertainties, associated mainly with bolometric corrections and the fraction of obscured sources. In agreement with Marconi et al. (2004) and Merloni (2004), we find similar trends in the evolution of AGN emissivity and the cosmic star formation rate, supporting models in which the growth of black holes occurs together with the formation of stars in their hosts (e.g., Sanders et al. 1988; Granato et al. 2004; Hopkins et al. 2006a). We also reassess current estimates of the local black hole mass function, finding that different choices and calibrations of the black hole-bulge correlation lead to significantly different results, with integrated mass densities that span nearly a factor of two.  Our simplest models of the evolving black hole and AGN populations assume that all active black holes have a single radiative efficiency $\\epsilon$ and a single accretion rate $\\dot{m}$ in Eddington units. The work of Kollmeier et al. (2006) suggests that a constant value of $\\dot{m}$ may be a reasonable approximation for luminous quasars, but the observational evidence on this point is mixed (e.g., McLure \\& Dunlop 2004; Vestergaard 2004; Babic et al. 2007; Bundy et al. 2007; Netzer \\& Trakhtenbrot 2007; Netzer et al. 2007; Rovilos \\& Gorgantopoulos 2007; Shen et al. 2007b), and for lower luminosity, local AGN there is clearly a broad range of Eddington ratios (e.g., Heckmann et al. 2004). We also consider models in which $\\dot{m}$ depends on redshift or black hole mass, and we consider a simple model of black hole mergers to assess their potential impact on the mass function. We will examine models with distributions of $\\dot{m}$ values and a more realistic treatment of mergers in future work.  Our modeling allows a consistency test of the basic scenario in which the observed luminosity of black hole accretion drives the growth of the underlying black hole population, and comparison to the local black hole mass function yields constraints on the average radiative efficiency and the typical accretion rate. For specified $\\epsilon$ and $\\dot{m}$, the model predicts the black hole mass function as a function of time, and combination with the observed luminosity function yields the duty cycle as a function of redshift and black hole mass. These predictions can be tested against observations of AGN host galaxy properties and AGN clustering, and they can be used as inputs for further modeling. While significant uncertainties remain, we find that a simple model of the black hole and AGN populations achieves a good match to a wide range of observational data. ",
        "conclusions": "\\label{sec|DiscuConclu}  We have constructed self-consistent models for the evolution of the supermassive black hole population and the AGN population, in which black holes grow at the rate implied by the observed luminosity function given assumed values of the radiative efficiency $\\epsilon$ and the characteristic accretion rate $\\dot{m}_0=\\dot{M}_{\\bullet}/\\dot{M}_{\\rm Edd}$ (see Eqs.~\\ref{eq|conteq}, ~\\ref{eq|mdotav}, ~\\ref{eq|POsingle}). These models can be tested against the mass distribution of black holes in the local universe, and they make predictions for the duty cycles of black holes as a function of redshift and mass, which can be tested against observations of quasar hosts and quasar clustering if one assumes an approximately monotonic relation between the masses of black holes and the masses of their host galaxies and halos. Our method is similar to that used previously by Cavaliere et al. (1971), Small \\& Blandford (1992), Yu \\& Tremaine (2002), SW03, Marconi et al. (2004), and S04. However, we have drawn on more recent data on the AGN luminosity function and the local black hole mass function, and we have considered a broader spectrum of models. This approach can be considered a ``differential'' generalization of the So{\\l}tan (1982) argument relating the integrated emissivity of the quasar population to the integrated mass density of the local black hole population.  Our model for the bolometric AGN luminosity function starts from the model of U03, based on X-ray data from a variety of surveys, but we adjust its parameters as a function of redshift in light of more recent measurements and data at other wavelengths. In agreement with previous studies we find that the LF of optical AGNs is roughly consistent with that of X-ray AGN that have absorbing column densities $\\log N_H/{\\rm cm^{-2}}\\le 22$ and that unobscured AGN dominate the bright end of the LF. We show that the latest constraints on the hard X-ray background ($E\\sim 10-100$ keV) from \\emph{INTEGRAL} and from the \\emph{PDS} instrument on BeppoSax support a reduced normalization relative to extrapolations from other missions at lower energies. They therefore favor a lower contribution from very highly obscured AGN ($\\log N_H/{\\rm cm^{-2}}\\gtrsim 24.5$) than some previous estimates (but see also Gilli et al. 2007). Our estimate of the bolometric AGN LF is independent in implementation but similar in spirit to that of HRH07. The most significant differences for purposes of this investigation are that HRH07 have a higher LF normalization at $z\\sim 1-3$, principally because of their choice of bolometric correction, and they have a steeper bright-end slope at $z\\sim 1-2.5$, where they adopt the slopes measured by Richards et al. (2006) from the SDSS and we use the slopes inferred by U03 from X-ray data. We regard the differences between our estimate and that of HRH07 as a reasonable indication of the remaining systematic uncertainties in the bolometric LF of AGNs.  With our LF estimate, the bolometric emissivity of the AGN population tracks recent estimates of the cosmic star formation rate as a function of redshift. (Comparisons based on the space density of high luminosity quasars [e.g., Richards et al. 2006, Osmer 2004 and references therein] reach a different conclusion because at low redshifts the bright end of the AGN LF drops much more rapidly with time than the overall emissivity.) The integrated black hole mass density implied by this emissivity is $\\sim 8\\times 10^{-4}$ of the stellar mass at all redshifts, or about half of the estimated ratio of black hole mass to bulge stellar mass in local galaxies. This tracking favors scenarios in which black holes and the stellar mass of bulges grow in parallel, with about $50\\%$ of the star formation linked to black hole growth at all redshifts. These findings are hard to reconcile with any models where black hole growth substantially precedes stellar mass buildup, or with recent claims that the ratio of black hole mass to stellar mass is much larger at high redshift than the local value. However, our finding refers to integrated densities, so it does not indicate the relative timing of black hole and bulge growth on an object-by-object basis.  Observational estimates of the local black hole mass function still show substantial discrepancies among different authors, depending on the correlation used to derive it (e.g., $M_{\\bullet}-\\sigma$, $M_{\\bullet}-L_{\\rm bulge}$, $M_{\\bullet}-M_{\\rm star}$, $M_{\\bullet}$-S\\'{e}rsic index, or fundamental plane), the calibration of the correlation, and the intrinsic scatter of the correlation. Above $M_{\\bullet}\\sim 10^9\\, {\\rm M_{\\odot}}$, estimates depend on extrapolation of the observed scaling relations, and below $M_{\\bullet}\\sim 10^{7.5}\\, {\\rm M_{\\odot}}$ they are sensitive to the treatment of spiral bulges. The grey band in Figure~\\ref{fig|LocalMFs} (and subsequent figures) encompasses most estimates, but the fundamental plane (Hopkins et al. 2007b) and S\\'{e}rsic index (Graham et al. 2007) methods imply more sharply peaked mass functions. The integrated mass densities of all of these estimates are in the range $\\rho_{\\bullet}\\sim 3-5.5\\times 10^5\\, {\\rm M_{\\odot}\\,Mpc^{-3}}$.  Our simplest models assume a single characteristic Eddington accretion rate $\\dot{m}_0$, independent of redshift and black hole mass, and all of our models assume a single value of the radiative efficiency $\\epsilon$. Matching the local black hole mass density requires $\\epsilon=0.075\\times(\\rho_{\\bullet}/4.5\\times 10^5\\, {\\rm M_{\\odot}\\,Mpc^{-3}})^{-1}$ for our standard estimate of the AGN LF, or $\\epsilon \\sim 0.094(\\rho_{\\bullet}/4.5\\times 10^5\\, {\\rm M_{\\odot}\\,Mpc^{-3}})^{-1}$ for the HRH07 luminosity function.\\footnote{More precisely, it is $\\epsilon/(1-\\epsilon)$ that is proportional to $\\rho_{\\bullet}^{-1}$, but the difference from $\\epsilon\\propto \\rho_{\\bullet}^{-1}$ is tiny over the allowed range.} With $\\epsilon$ thus fixed, the value of $\\dot{m}_0$ determines the peak of the predicted local black hole mass function in the $M_{\\bullet}\\Phi(M_{\\bullet})$ vs $M_{\\bullet}$ plane. Note that with our definitions the Eddington luminosity ratio is $\\lambda=L/L_{\\rm Edd}\\approx \\dot{m}_0(\\epsilon/0.1)$ (equation~\\ref{eq|Lambda}). Matching the observed peak at $\\log M_{\\bullet}/{\\rm M_{\\odot}}\\sim 8.5$ implies $\\dot{m}_0\\approx 0.6$ ($\\lambda \\approx 0.45$) for our standard LF estimate and $\\dot{m}_0\\approx 1$ ($\\lambda \\approx 0.95$) for the HRH07 LF. Lower values, $\\dot{m}_0=0.1-0.3$, shift the peak location to untenably high masses, $\\log M_{\\bullet}/{\\rm M_{\\odot}}\\sim 8.9-9.3$ or, for HRH07, $\\log M_{\\bullet}/{\\rm M_{\\odot}}\\sim 9.1-9.6$. The single-$\\dot{m}_0$ models achieve a reasonable match to our ``grey-band'' observational estimates of the width and asymmetry of the local mass function, though for our LF estimate the predicted mass function is too high at $M_{\\bullet}>10^9\\, {\\rm M_{\\odot}}$ and is therefore somewhat too broad. Single-$\\dot{m}_0$ models cannot reproduce the more sharply peaked local mass functions estimated by Graham et al. (2007) or Hopkins et al. (2007b).  Our reference model, which has $\\epsilon=0.065$, $\\dot{m}_0=0.60$, and our standard LF estimate, predicts a duty cycle for activity of $10^9\\, {\\rm M_{\\odot}}$ black holes that declines steadily from 0.15 at $z=4$ to 0.07 ($z=3$), 0.035 ($z=2$), 0.004 ($z=1$), and $10^{-4}$ ($z=0$). The decline in duty cycle for lower mass black holes is much shallower. Massive black holes therefore build their mass relatively early while low mass black holes grow later, the phenomenon often referred to as ``downsizing''. Our results on mean radiative efficiency and duty cycle evolution are also in qualitative agreement with those found by Haiman et al. (2004). The predicted duty cycles seem in reasonable accord with observational estimates, though these estimates have considerable uncertainty and do not, as yet, probe mass and redshift dependence in much detail. The electronic tables described in \\S~\\ref{subsec|Tables} provide tabulations as a function of redshift of our AGN bolometric LF estimate, the HRH07 LF, and black hole mass functions and duty cycles for single-$\\dot{m}_0$ models that are in good agreement with the observed $z=0$ mass functions given these LF inputs.  We have examined models in which the Eddington ratio accretion rate $\\dot{m}_0$ is reduced at low redshift or at low black hole mass. Declining redshift evolution of $\\dot{m}_0$ damps ``downsizing,'' reducing the dependence of duty cycles on black hole mass and redshift. This model produces a typical duty cycle $P_0\\sim 10^{-2.5}$ at $z=0$, about two orders of magnitude higher than in the reference model and consistent with the local duty cycle inferred from observations by Greene \\& Ho (2007). In general terms, the observed luminosity-dependent density evolution of the AGN LF can be explained by preferential suppression of activity in high mass black holes at low redshift, by a decline in the typical accretion rate at low redshift, or by some combination thereof. The mass-dependent $\\dot{m}_0$ model associates more of the AGN emissivity to high mass black holes, so it predicts a $z=0$ mass function that is more sharply peaked, in better agreement with the estimates of Hopkins et al. (2007b) and Graham et al. (2007) but worse agreement with other estimates. This model predicts a stronger mass dependence of duty cycles than our reference model because it maps low mass black holes, whose abundance is already suppressed, to less luminous, and hence more common, AGN.  We have also considered a model in which each black hole has a $50\\%$ probability per Hubble time of merging with another black hole of equal mass. Merger-driven growth in this model has little impact on the black hole mass function until $z<1$, when accretion-driven evolution has slowed. Low redshift mergers slightly depress the low mass end of the $z=0$ mass function and significantly enhance the high mass tail, worsening the agreement with observational estimates. Models incorporating theoretically predicted merger rates can allow more realistic calculations of the impact of mergers on the black hole population; this impact will probably be smaller than in the simplified model considered here.  We have calculated the clustering bias of AGN as a function of luminosity and redshift for our reference model, assuming a monotonic relation between black hole mass and halo mass. The predictions are in reasonable accord with observational estimates. We will examine AGN clustering predictions in more detail in future work, with attention to what models can be excluded by the data and what can be learned by matching the full AGN correlation function in addition to an overall bias factor.  MHD simulations (e.g., Gammie et al. 2004; Shapiro 2005) show that disk accretion onto Kerr black holes spins them up to an equilibrium spin rate $a\\approx 0.95$ (where $a=1$ is the angular momentum parameter for a maximally rotating black hole). The radiative efficiency in these models is $\\epsilon\\approx 0.16-0.2$. These high efficiencies would lead to black hole mass densities a factor of two or more below our central estimate, and below our estimated lower bound. Furthermore, our results show that models with $\\epsilon$ in the range $0.06-0.11$ can achieve a good match to the overall shape of the black hole mass function near the peak in $M_{\\bullet}\\Phi(M_{\\bullet})$, not just the value of $\\rho_{\\bullet}$, given plausible choices of $\\dot{m}_0$. Systematic uncertainties in the AGN LF do not appear large enough to accommodate $\\epsilon\\gtrsim 0.15$. Accommodating these high efficiencies would instead require a substantial downward revision of recent estimates of the local black hole mass function, reducing the integrated mass density to $\\rho_{\\bullet}\\sim 2\\times 10^5\\, {\\rm M_{\\odot}\\,Mpc^{-3}}$. Our results are consistent with a scenario like the one of King \\& Pringle (2006) in which chaotic accretion spins down black holes because of counter-alignment with the accretion disk angular momentum, or with other mechanisms that reduce efficiencies below the MHD-simulation predictions.  The assumption that all active black holes at a given mass and redshift have the same $\\dot{m}_0$ is clearly an idealization, at least in the local universe where observations indicate a wide range of Eddington ratios (Heckman et al.\\ 2005; Greene \\& Ho 2007). Steed \\& Weinberg (2003) and Yu \\& Lu (2004) discussed continuity equation models evolved adopting a distribution of Eddington ratios. In particular, Yu \\& Lu (2004) have derived the relation between the integrated number of AGNs shining at all times at a given luminosity $L$, the mean light curve of black holes, and the local black hole mass function. Following their equation (18), we find that a good match between the cumulative number of AGNs and of relic black holes can be achieved for $\\epsilon\\sim 0.07$ (required by the So{\\l}tan argument) and a mean AGN light curve exponentially increasing with $\\lambda=0.6$ and a negligible declining phase, similar to our reference model. Alternatively, we find that a good match can be obtained by assuming that black holes grow rapidly in a Super-Eddington phase with $\\lambda\\gtrsim 2$ and then have a long declining phase, qualitatively resembling our $\\dot{m}(z)$ model discussed in \\S~\\ref{subsec|etamdot}. In future work we will investigate models that incorporate multiple $\\dot{m}_0$-values and accretion modes, including the addition of modest log-normal scatter in $\\dot{m}_0(M,z)$ (e.g., Kollmeier et al.\\ 2005; Netzer et al.\\ 2007; Shen et al.\\ 2007b) and sharper revisions in which some black holes accrete at super-Eddington or highly sub-Eddington rates, perhaps with reduced radiative efficiencies (Narayan, Mahadevan, \\& Quataert 1998 and references therein).  We will also incorporate mergers at the rates predicted by theoretical models of cold dark matter subhalos and their associated black holes (Yoo et al.\\ 2007). For appropriate parameter choices, we expect that many scenarios can be made consistent with the observed AGN LF and the local black hole mass function. However, direct measurements of Eddington ratio distributions and measurements of AGN clustering and host properties, all as a function of luminosity and redshift, should greatly narrow the field of viable models.  Within the (often substantial) uncertainties of existing data, a simple model in which all black holes grow by accreting gas at mildly sub-Eddington rates with a radiative efficiency $\\epsilon \\approx 0.06-0.1$ is surprisingly successful at reproducing a wide range of observations."
    },
    "0710/0710.1863_arXiv.txt": {
        "abstract": "We present the results of a search for molecular gas emission via the CO line in the far outer disk of the nearby spiral, NGC~6946. The positions targeted were chosen to lie on or near previously-identified outer disk HII regions. Molecular gas was clearly detected out to 1.3~R$_{25}$, with a further tentative detection at 1.4~R$_{25}$. The CO detections show excellent agreement with the HI velocities and imply beam-averaged column densities of $0.3-9\\times 10^{20}$~cm$^{-2}$ and molecular gas masses of (2-70)$\\times 10^{5}$~M$_{\\sun}$ per 21$''$ beam (560pc). We find evidence for an abrupt decrease in the molecular fraction at the edge of the optical disk, similar to that seen previously in the azimuthally-averaged areal star formation rate.  Our observations provide new constraints on the factors that determine the presence and detectability of molecular gas in the outskirts of galaxies, and suggest that neither the HI column, the metallicity or the local heating rate alone plays a dominant role. ",
        "introduction": "It has become increasingly apparent in recent years that star formation is not just confined to the optically-bright parts of galaxies.  Signposts of low-level star formation, such as HII regions and UV-bright clusters, have been discovered in the far outer regions of galactic disks, well beyond the classical R$_{25}$ radius (e.g. \\cite{ferg98a,thilker05}), and in the tidal HI filaments of some interacting systems (e.g. \\cite{ryanw04, Braine01}).  The existence of star formation in these unusual environments, characterised by low metallicity, low HI column density and low interstellar radiation field, provides important constraints on star formation in galaxies, as well as the necessary conditions for sustaining a multi-phase interstellar medium (ISM).  Much theoretical effort has been devoted to understanding the formation of a cold molecular phase in the ISM, a prerequisite for star formation.  Thermal instabilities are generally believed to be the primary mechanism for converting cool atomic gas into cold molecular gas once a minimum pressure or local threshold surface density is exceeded (e.g. \\cite{elme94,schaye04, blitz06}). Additional mechanisms for forming unstable cold clouds include shocks from spiral density waves and swing-amplifier and magneto-Jeans instabiities (e.g. \\citet{dobbs06, kim02}).  The threshold column densities predicted for molecular cloud formation are on average a few times smaller than those inferred for massive star formation in galaxies (e.g. \\cite{skill87,mk01}).  Direct detection of molecular clouds in extreme environments provides a means to test ideas about cloud formation and survival.  Carbon monoxide (CO) is the most abundant heteronuclear molecule (i.e. with permitted rotational transitions) and is excited by collisions at densities typical of molecular clouds, making CO the standard tracer of molecular gas.  Concerns have previously been raised \\citep[e.g. ][]{Pfeniger94} as to whether the physical conditions in the outskirts of disks are sufficient to excite CO to detectable temperatures and whether the metallicities in these regions are sufficient to form enough CO for detection.  With the discovery of CO emission out to 1.5~R$_{25}$ in the disk of NGC~4414 \\citep{braine04} and in high column density tidal HI filaments in the M81 Group \\citep{brou92,walter06}, these concerns have been proven somewhat unfounded.  In this Letter, we present the results of a search for molecular gas in the far outer disk of the nearby spiral NGC~6946.  While the study of molecular gas in NGC~4414 focused on high column HI regions \\citep{braine04}, we have chosen to search for CO emission in the vicinity of outer disk star-forming regions. NGC~6946 was previously identified through H$\\alpha$ observations to exhibit very extended massive star formation, out to $\\sim 2$R$_{25}$ \\citep{ferg98a}, and thus presents an excellent opportunity to attempt to measure and quantify molecular gas at extreme radii.  We also report on the null detection of molecular gas at two positions in the far outer disk of another spiral with extended star formation, NGC~1058. ",
        "conclusions": "We have presented evidence for the existence of molecular gas in the far outskirts of NGC~6946 and noted a rather abrupt decrease in the molecular gas fraction, as well as the PAH emission, as one crosses R$_{25}$ along an individual spiral arm.  An important question is whether this behavior could be explained by metallicity variations in the underlying gas disk. Indeed there is strong evidence for a radially increasing $\\ratioo$ value \\citep[e.g.][]{Sodroski95} as well as much higher values in low-metallicity systems \\citep[e.g.][]{Rubio91}.  The metallicity in the outer disk of NGC~6946 varies by a factor $\\sim 5$ from log(O/H)$=-3.2$ to $-3.9$ \\citep{ferg98b} while the I$_{\\rm CO}$/N(HI) ratio decreases by a factor of $\\sim30$. Focusing only on the positions which bridge R$_{25}$ (P12 to P2), the metallicity remains essentially constant while the I$_{\\rm CO}$/N(HI) ratio decreases by a factor of $\\sim6$.  In the absence of a very strong radiation field, the $\\ratioo$ factor should not vary more than linearly with metallicity and \\citet{Wilson95} estimates that $\\ratioo \\propto [O/H]^{2/3}$ from a study of molecular clouds in Local Group galaxies.  It is thus reasonable to expect that the $\\ratioo$ ratio might be a factor 2 -- 3 higher in the most distant points than at P9, where the \"standard\" value is probably appropriate.  Such a correction is not sufficient to explain the strong decrease in I$_{\\rm CO}$/N(HI) at R$_{25}$ and we conclude that there is a genuine drop in the molecular gas content at the edge of the optical disk as defined by R$_{25}$.  Although our pointings were selected to lie on or near outer disk HII regions, the H$\\alpha$ luminosity contained within each beam shows significant variation and corresponds to local star formation rates (SFRs) ranging from $10^{-2}-10^{-4}$ M$_\\odot$ yr$^{-1}$, using the \\citet{kenn98} calibration (SFR$=$L(H$\\alpha / 1.26 \\times 10^{41}$ erg s$^{-1}$). The H$\\alpha$ emission in some regions is predominantly diffuse (e.g. P12 and P7) and may more reflect ionization by photons produced elsewhere rather than {\\it in situ} star formation \\citep{ferg96}. The star formation efficiency (SFE), defined as SFR/M(H$_2$), ranges from $\\sim 10-0.07$ Gyr$^{-1}$; eliminating the two outlier points (P12 and P2), the range reduces to $2 - 0.5$ Gyr$^{-1}$.  The mere presence of a luminous HII region does not guarantee the detection of molecular gas.  Position P1 in NGC~1058 is the second most luminous star-forming region targeted in our study, and is our second most sensitive integration, yet CO was not detected there.  Intriguingly, there is a potential correlation between the detection of CO(2--1) emission and the presence of a luminous HII region in that the three non-detections of CO(2--1) are among the four lowest H$\\alpha$ fluxes.  With the exception of P4, the HI column densities at our observed positions within NGC~6946 are quite high, close to or greater than $10^{21}$cm$^{-2}$ on the scale of several hundred parsecs.  The HI column density may be an important factor in determining whether molecular gas will be detected but a high HI column is clearly not a sufficient condition \\citep[see also][]{Gardan07}.  It may be purely coincidental that N(H$_2$)/N(HI) changes abruptly across R$_{25}$ in NGC~6946. While R$_{25}$ provides a means to characterise the optical extents of galaxies, it is not expected to relate to any underlying physical properties. In NGC~6946, neither the metallicity, the HI surface density nor the surface brightness of the stellar disk exhibit unusual features at this location.  Interestingly, the azimuthally-averaged areal star formation rate in NGC~6946 does show a sharp decline near the edge of the optical disk ($\\sim 0.8$~R$_{25}$), although widespread low-level star formation is observed much further out \\citep{ferg98a,mk01}.  The rough correspondence between the star formation and molecular profile breaks could provide support for threshold models \\citep[e.g.][]{elme94,schaye04} although the physics that drives this behavior is still open to debate.  Pressure based \\citep{elme94,blitz06} and column density based \\citep{schaye04} models are very difficult to distinguish once the HI surface density exceeds that of the stars. Larger-scale mapping of the outer disk of NGC~6946 and sensitive observations of the outskirts of additional galaxies are required in order to more thoroughly understand the molecular content and distribution at large galactocentric radii."
    },
    "0710/0710.1114_arXiv.txt": {
        "abstract": "\ufeffWe present an \\xmm\\ detection of two low radio surface brightness SNRs, G85.4+0.7 and G85.9$-$0.6, discovered with the Canadian Galactic Plane Survey (CGPS). High-resolution \\xmm\\ images revealing the morphology of the diffuse emission, as well as discrete point sources, are presented and correlated with radio and \\cha\\ images. The new data also permit a spectroscopic analysis of the diffuse emission regions, and a spectroscopic and timing analysis of the point sources. Distances have been determined from \\ion{H}{1} and CO data to be $3.5\\pm 1.0$ kpc for SNR G85.4+0.7 and $4.8\\pm 1.6$ kpc for SNR G85.9$-$0.6.  The SNR G85.4+0.7 is found to have a temperature of $\\sim12-13$ MK and a 0.5--2.5~keV luminosity of $\\sim1-4\\times 10^{33}D^2_{3.5}$ erg/s (where $D_{3.5}$ is the distance in units of 3.5 kpc), with an electron density $n_e$ of $\\sim0.07-0.16(fD_{3.5})^{-1/2}$ cm$^{-3}$ (where $f$ is the volume filling factor), and a shock age of $\\sim9-49(fD_{3.5})^{1/2}$ kyr.  The SNR G85.9$-$0.6 is found to have a temperature of $\\sim15-19$ MK and a 0.5--2.5~keV luminosity of $\\sim1-4\\times 10^{34}D^2_{4.8}$ erg/s (where $D_{4.8}$ is the distance in units of 4.8 kpc), with an electron density $n_e$ of $\\sim0.04-0.10(fD_{4.8})^{-1/2}$ cm$^{-3}$ and a shock age of $\\sim 12-42(fD_{4.8})^{1/2}$ kyr.  Based on the data presented here, none of the point sources appears to be the neutron star associated with either SNR. ",
        "introduction": "In 2001, two new supernova remnants (SNRs) with low radio surface brightness, G85.4+0.7 and G85.9$-$0.6,  were discovered in Canadian Galactic Plane Survey (CGPS)\\citep{tay03} data and confirmed in X-rays with \\ro\\ data \\citep{kot01}. Both show distinct shells in the radio band with an extended region of X-ray emission in the centre. The radio surface brightness of G85.4+0.7 at 1 GHz is $\\Sigma_{1 {\\rm GHz}}\\le 1\\times 10^{-22}$ Watt ${\\rm m}^{-2}$ ${\\rm Hz}^{-1}$ ${\\rm sr}^{-1}$ and the radio data also indicate that the SNR has a non-thermal shell with angular diameter $\\approx0.4^{\\circ}$ which is surrounded by a thermal shell with an angular diameter of $\\approx0.6^{\\circ}$ and is located within an H I bubble. The bubble also contains two B stars which may have been part of the same association as the SNR's progenitor star. G85.9$-$0.6 has a radio surface brightness $\\Sigma_{1 {\\rm GHz}}\\le 2\\times 10^{-22}$ Watt ${\\rm m}^{-2}$ ${\\rm Hz}^{-1}$ ${\\rm sr}^{-1}$ and it has no discernible H I features, indicating that it is expanding into a low density medium, perhaps between the local and Perseus spiral arms. The most likely event which would produce an SNR in such a region would be a Type Ia supernova.   X-ray observations are important to the study of SNRs, particularly those with low surface brightness, because they provide information about the morphology and emission processes of these objects, which are indicators of both the nature of the supernova explosion which formed them and the properties of the progenitor star. SNRs with low surface brightness are expected to be formed after the core collapse of a massive star in a Type Ib/c or Type II explosion, because the stellar wind would have blown away much of the interstellar medium (ISM) surrounding it, leaving a low ambient density into which the shock from the supernova expands. A Type Ia supernova can also result in a low surface brightness SNR if the surrounding ISM has a low density, such as it would if it were located between two spiral arms of the galaxy. Thermal X-ray emission from a SNR arises as the blast wave of the explosion travels through and shocks the ISM, and as a reverse shock travels back into and shocks the ejecta. The X-ray spectrum of the SNR gives information about the temperature, the density, and the luminosity of the shocked material, while imaging data provides information about the size and morphology of the region.    The low angular and spectral resolutions as well as the small number of counts in the \\ro\\ data did not allow spectroscopy nor detailed imaging to be done, so \\xmm\\ observations, described in \\S\\ref{obs}, have been made in order to confirm the detection of the SNRs, and to perform imaging, spectroscopic and timing studies. Detailed X-ray imaging, described in \\S\\ref{imaging}, is used to map the diffuse emission and compare it to the location and size of the radio shells, and \\cha\\ data have been used to search for compact objects not resolved by XMM. Spectral parameters obtained in \\S\\ref{spec} lead to an estimate of such quantities as temperature, density, and luminosity of the SNRs. In \\S\\ref{ps} the point sources are catalogued and an attempt at identification is made by matching their positions to objects in other catalogues. Timing analysis is performed to identify any pulsar candidates. The distances to the SNRs are derived from \\ion{H}{1} and CO data in \\S\\ref{dist}. The results are discussed in \\S\\ref{disc}.   Preliminary results were presented in \\cite{jac06} and \\cite{saf06}. ",
        "conclusions": "\\label{disc}   \\subsection{Determination of shock age and mass of emitting gas}  To estimate $t$, the shock age, for the SNRs, the relation $t=\\tau/n_{\\rm e}$ is used, where $\\tau$ is the upper limit of the ionization timescale from the VPSHOCK model. The electron density $n_{\\rm e}$ is determined from the distance in cm $D_{\\rm cm}$, angular size in radians $\\alpha$, given by the diameter of the extraction region, and the relation $n_{\\rm H}=n_{\\rm e}/1.2$, where $n_{\\rm H}$ is the volume density of hydrogen within. Given that ${\\rm Norm}=\\frac{\\int n_{\\rm e}n_{\\rm{H}}dV}{10^{14}(4\\pi D_{\\rm cm}^2)}$, where Norm is the normalization of the VPSHOCK model, which is proportional to the emission measure, and the electron and hydrogen densities $n_{\\rm e}$ and $n_{\\rm{H}}$ are assumed to be uniform, a volume filling factor $f$ is employed so that $n_{\\rm e}=(\\frac{2.88\\times 10^{15} ({\\rm Norm})}{f \\alpha^3D_{\\rm cm}})^{1/2}$. From these calculations, it is determined that for G85.4+0.7, $n_e \\approx (0.11\\pm0.03)(fD_{3.5})^{-1/2}$ for the ESAS background (for the MOS instrument) or $(0.11^{+0.05}_{-0.04})(fD_{3.5})^{-1/2}$ for the observation background, and $t =(18^{+17}_{-9})(fD_{3.5})^{1/2}$ kyr for the ESAS background or $(23^{+26}_{-14})(fD_{3.5})^{1/2}$ kyr for the observation background. For G85.9$-$0.6, $n_e \\approx (0.07^{+0.03}_{-0.02})(fD_{4.8})^{-1/2}$ for the ESAS background or $(0.07\\pm0.03)(fD_{4.8})^{-1/2}$ for the observation background and $t = (30^{+12}_{-13})(fD_{4.8})^{1/2}$ kyr for the ESAS background or $(22^{+11}_{-10}) (fD_{4.8})^{1/2}$ kyr for the observation background. These are larger than the age estimates from the radio data.  The mass of the emitting gas is calculated based on spherical emitting regions with the size given by the extraction radius (given in Table~\\ref{dfit}), composed of 92\\% hydrogen and 8\\% helium, and the relation $n_{\\rm H}=n_{\\rm e}/1.2$ is used as for the above calculation of $t$. The mass of the emitting gas in G85.4+0.7 is $(2.8\\pm1.5) D_{3.5}^{5/2} M_{\\sun}$ for the ESAS background and $(2.8^{+1.8}_{-1.6})D_{3.5}^{5/2} M_{\\sun}$ for the observation background and that for G85.9$-$0.6 is $(2.7^{+1.8}_{-1.6}) D_{4.8}^{5/2} M_{\\sun}$ for the ESAS background and $(2.7\\pm1.8)D_{4.8}^{5/2} M_{\\sun}$ for the observation background.  \\subsection{Background subtraction}  The background subtraction was problematic, particularly for the spectrum of G85.4+0.7, which is fainter and required a different model to be fit to it depending on which background region on the observation was used, leading to a very careful selection of the background extraction region, the process of which is described in \\S\\ref{bg}. Instrumental lines, which are fortunately different for PN and MOS, nevertheless increased the value of $\\chi^2$ for the combined fits, even when the background was chosen very carefully, and needed to be explicitly fit when the ESAS background was used, as described in \\S\\ref{bg}. In addition, the soft spectra of the SNRs exhibit large amounts of noise in the high energy end of the spectra, allowing fits to be made only up to 2.5 keV. Future longer observations of these SNRs will help to resolve these difficulties, allow for better fits to the abundances, eliminate ambiguities in the spectral results, and perhaps allow for conclusive identifications of neutron star and pulsar candidates in addition to other point sources.    \\subsection{G85.4+0.7}\\label{disc85.4}  A comparison between the morphologies of the X-ray diffuse emission and the radio emission of G85.4+0.7 can be seen in Figure~\\ref{85.4imagens} in which the X-ray point sources have been removed and the X-ray image has been smoothed to 1$^{\\prime}$ to match the radio contours. The diffuse X-ray emission exhibits some structure, and it lies in the approximate centre of the radio shell. A spectrum was extracted from the central blob, the position and size of which is shown in the top panel of Figure~\\ref{85.4reg}, and it could not be adequately fit with a non-thermal power law model, indicating that a pulsar wind nebula origin for the emission can be ruled out.  The fact that the SNR does not exhibit any limb brightening and the X-ray emission is mostly concentrated in the centre, as projected in the plane of the sky and three dimensionally, can be interpreted as evidence that the X-ray emission is produced by the ejecta. The fact that there appears to be no emission from the swept up material can be explained if the remnant is evolutionarily young.  With the ejecta interpretation, if it is assumed that the free electrons are evenly distributed in the extraction area, the resulting ages are 18 and 23 kyr for the two backgrounds. In Figure~\\ref{85.4imagens} it can be seen that the outer radio shell has an approximately constant radius, whereas the inner shell has a smaller radius in the vertical centre, which expands as the latitude changes, and this indicates that the SNR is moving either toward us or away from us. The two shells appear to meet near the bottom of the image, so the radius of the SNR can be approximated as that of the outer shell, which is $\\sim0.3^\\circ$ on the image. The average velocity of the expanding SNR is then 910 or 730 km/s for the two ages. However, the morphology of the X-ray emission is not very smooth, indicating that the filling factor $f<<1$. Assuming $f=0.2$, the electron density $n_e$ would be 0.25 cm$^{-3}$ for either the ESAS and observation background, and the ejecta masses would be 1.3 $M_{\\sun}$ for either background. The shock age would then be 8000 or 10300 years, and the average expansion velocity would be 2150 or 1640 km/s. For a $10^{51}$ erg supernova explosion, the ejecta mass would be 22 or 37 $M_{\\sun}$, or 2.2 or 3.7 $M_{\\sun}$ for a $10^{50}$ erg explosion, indicating that the lower energy explosion would result in an ejecta mass consistent with a young freely expanding SNR.  The radio continuum emission indicates that there is some swept up material. Assuming that the density inside the stellar wind bubble was $\\sim 0.01$ cm$^{-3}$ before the explosion, the SNR would have swept up $\\sim 6.5$ $M_{\\sun}$ of material, which is on the order of the ejecta mass, and means that the SNR is in the transition between free expansion and Sedov expansion.  The slightly above-solar abundance of O in the diffuse spectra reinforces the hypothesis that the supernova most likely resulted from a core collapse, though the large error bars weaken the argument. The abundances of all elements other than O and Fe are below solar, but they would be enhanced if the X-rays were from the ejecta. However, it is possible that the spectral parameters for the abundances are affected by the poor quality of the spectra and by the uncertainties associated with the background subtraction.  Since the most likely origin of G85.4+0.7 was a core collapse supernova (see \\S\\ref{dist}), it is possible that a neutron star or pulsar which is associated with this SNR can be found. Given its power law X-ray spectrum, proximity to the centre of the remnant, and similar $N_{\\rm H}$ value to the diffuse emission, source 1 in Figure \\ref{85.4reg} is possibly the associated neutron star, but sources 4 and 9 are within the radio shell and are also candidates, given their spectral parameters. However, the low fit quality of source 4 ($\\chi^2\\sim2$) indicates that it is not likely to be a neutron star, and source 9 is situated far from the centre of the diffuse emission so is less likely to be the associated neutron star. \\cha\\ sources 10, 11, 12, 13, 15, or 16 in Figure~\\ref{cha85.4} can also be considered as neutron star candidates (provided the optical counterparts listed in Table~\\ref{pscat85.4} are not the true counterparts.   \\subsection{G85.9$-$0.6}\\label{disc85.9} Figure~\\ref{85.9imagens} shows X-ray and radio images of G85.9$-$0.6, produced in a similar way to Figure~\\ref{85.4imagens}. The diffuse X-ray emission appears to contain less structure than G85.4+0.7, and, as for G85.4+0.7, a spectrum extracted from the central blob, the position and size of which are shown in the top panel of Figure~\\ref{85.9reg}, indicates that a pulsar wind nebula origin for the emission can be ruled out.  Because Figure~\\ref{85.4imagens} shows no limb brightening and the emission is mostly from the central blob, a similar argument to that used for G85.4+0.7 can be employed here to suggest an ejecta interpretation for the X-ray emission. Assuming that the free electrons are evenly distributed in the extraction area the resulting age is 30 or 22 kyr for the two backgrounds. Using the average distance between the centre of the X-ray emission and the shell, $\\sim 0.2^\\circ$, the radius of the shell is 15.3 pc  The average expansion velocity would then be 500 or 680 km/s. The emission is concentrated in the inner $6^\\prime$, which indicates an electron density of 0.20 cm$^{-3}$ and an ejecta mass of 1.0 $M_{\\sun}$ which agrees with the predicted mass for a type Ia explosion, 1.4 $M_{\\sun}$. This would result in an age of 10600 or 7800 years and average expansion velocity of 1350 or 1850 km/s.  For a Type Ia supernova, the explosion energy is $\\sim10^{51}$ erg and the ejecta mass is 1.4 $M_{\\sun}$. Unlike for G85.4+0.7, the density in the interarm region is closer to 0.1 cm$^{-3}$, resulting in a swept up mass of nearly 50 $M_{\\sun}$. This indicates that the SNR is in the Sedov expansion phase, which is described by $R=14 (E_0/n_0)^{1/5}t^{2/5}$, where $R$ is the radius in pc, $E_0$ is the explosion energy in units of $10^{51}$ erg, $n_0$ is the ambient density in cm$^{-3}$, and $t$ is the age in units of $10^4$ years. For an explosion energy of 10$^{51}$ erg, a radius of 15.3 pc, and an age of 10600 or 7800 years, the resulting ambient density is 0.84 or 0.55 cm$^{-3}$ which are consistent with the density of the interarm region, but the swept up mass would be 360 or 240 $M_{\\sun}$, from which it should be possible to measure thermal X-ray emission, from the part of the shell that is included in the X-ray pointing. The current expansion velocity would be $dR/dt=0.4(R/t)$, which is 570 or 760 km/s.    The interpretation of G85.9$-$0.6 as having been produced by a Type Ia supernova is reinforced by the Fe abundance, which is well above solar. As with G85.4+0.7, the ejecta interpretation is called into question by the remaining abundances, which should be above solar, but are instead below solar. This could again be a result of poor quality spectra.  Given that the radio results described in \\cite{kot01}, as well as the distance presented here, indicated that G85.9$-$0.6 was most likely produced by a Type Ia supernova, it was not expected to find a neutron star associated with this SNR. The above solar Fe abundance for this SNR is consistent with a Type Ia explosion. An identification of sources 6 and 7 in Figure~\\ref{85.9reg} has not yet been made. They are bright X-ray emitting objects with no known optical or radio counterpart, making them good neutron star or radio-quiet AGN candidates, though if one of them were a neutron star, it would be unlikely that it is associated with the G85.9$-$0.6 SNR because they are both situated far outside the radio shell of the remnant, and furthermore, the value of $N_{\\rm H}$ for source 6 does not match that of the SNR itself, and the fit quality of source 7 ($\\chi^2\\sim2$) indicates that an absorbed power law is not a good fit. Their identification with possible 2MASS counterparts (see Table~\\ref{pscat85.9}) makes them more likely radio-quiet AGN than neutron stars, provided the 2MASS objects are the true counterparts. A future detailed deep X-ray or multiwavelength study of these objects should be undertaken to identify and further study them, even though they are probably not associated with the G85.9$-$0.6 SNR. Source 4 is on the edge of the radio shell of the SNR, but its identification as a neutron star is questionable because of its photon index and hardness ratio, as well as the fact that it is not expected that there is a neutron star associated with this remnant.   \\subsection{Mixed Morphology Interpretation} The centrally filled morphology of both SNRs and the thermal nature of their X-ray emission confined within the radio shells suggest that they belong to the class of mixed-morphology SNRs (also known as thermal composites; Rho \\& Petre 1997). The origin of the thermal X-ray emission interior to the radio shells in these SNRs has been attributed to several mechanisms which include: a) cloudlet evaporation in the SNR interior (White \\& Long 1991), b) thermal conduction smoothing out the temperature gradient across the SNR and enhancing the central density (Cox et al. 1999), c) a radiatively cooled rim with a hot interior (e.g. Harrus et al. 1997), and d) possible interaction with a nearby cloud (e.g. Safi-Harb et al. 2005). While modeling these SNRs in the light of the above mentioned models is beyond the scope of this paper and has to await better quality data, we can rule out the cloudlet evaporation model based on the low ambient densities inferred from our spectral fits (see \\S7.1 and Table~1).  Except for Fe and possibly O, the abundances inferred from our spectral fits are for the most elements consistent with or below solar values, as observed in most mixed-morphology SNRs. However, enhanced metal abundances have been observed in younger SNRs, e.g. 3C~397, estimated to be $\\sim$5.3 kyr--old and proposed to be evolving into the mixed-morphology phase (Safi-Harb et al. 2005).  The ages inferred for G85.4+0.7 ($\\sim$8--10 kyr; see Table~1 and \\S\\ref{disc85.4}) and G85.9$-$0.6 ($\\sim$8--11 kyrs; see Table~1 and \\S\\ref{disc85.9}) suggest a later evolutionary phase where only shock-heated ejecta from Fe (for G85.9$-$0.6) and possibly Oxygen are still observed."
    },
    "0710/0710.2737_arXiv.txt": {
        "abstract": "We measure the relative evolution of the number of bright and faint (as faint as 0.05 $L^*$) red galaxies in a sample of 28 clusters, of which 16 are at $0.50\\le z \\le 1.27$, all observed through a pair of filters bracketing the $4000$ \\AA \\ break rest-frame. The abundance of red galaxies, relative to bright ones, is constant over all the studied redshift range, $0<z<1.3$, and rules out a differential evolution between bright and faint red galaxies as large as claimed in some past works. Faint red galaxies are largely assembled and in place at $z=1.3$ and their abundance does not depend on cluster mass, parametrized by velocity dispersion or X-ray luminosity. Our analysis, with respect to previous one, samples a wider redshift range, minimizes systematics and put a more attention to statistical issues, keeping at the same time a large number of clusters. ",
        "introduction": "The evolution of faint red galaxies in clusters is a highly debated topic for two reasons: different observers have claimed controversial results, and clusters of galaxies are often claimed to be interesting laboratories where studying the effect of the environment. Red galaxies, in particular, have different assembly histories in halos of different masses, yet observationally the detection of a environmental dependence of their properties escapes a detection. For example, differences between cluster and field fundamental planes are small, if any (Pahre et al. 1998), so small that the Coma cluster fundamental plane (Jorgensen et al. 1996) is routinely used as zero-redshift reference for studying the evolution of field galaxies, and so small that previously claimed differences are probably due to having overlooked the difficulty of the statistical analysis (van Dokkum \\& van der Marel 2007). Similarly, the colour of the red sequence seems not to depend on clustercentric distance (Pimbblet et al. 2002; Andreon 2003) or galaxy number density (Hogg et al. 2004, Cool et al. 2006).  The red colour, by which red galaxies are defined and selected, induces a selection effect: at every redshift only galaxies whose stellar populations are red (i.e. old, modulo dust, of no interest here) enter the sample. It is not a surprise, then, to find old-selected populations to be old. A different question is whether galaxies that have an old stellar population were fully assembled at early or late times. Answering this question requires a measurement of the abundance of red galaxies as a function of look-back time. For clusters, there is a further complication: clusters have different richnesses, jeopardizing any look-back time trend if the richness dependence is not factored out. It is easy, furthermore, to qualitatively claim that the red sequence is built later (i.e. a lower redshift) in poor environments than it is in dense environments, but this might just be do to signal to noise issues, because in poorer environments the red population is a minority one, and its contrast with respect to other populations (e.g. background) noisier. A sound statistical assessment of the abundance of faint red galaxies is therefore compelling.   \\begin{table} \\caption{The ACS $z\\ge0.5$ cluster and control field samples} \\begin{tabular}{l r c c l l} \\hline Name & z & $N^{1}$ &\\multispan{2}{\\hfill Filters \\hfill} \\\\ &   &         & blue & red \\\\ \\hline Lynx W            & 1.27 &  3\t &  F775W &   F850LP \\\\ Lynx E            & 1.26 &  3\t &  F775W &   F850LP \\\\ RDCS J1252-2927   & 1.23 &  4\t &  F775W &   F850LP \\\\ RDCS J0910+5422   & 1.11 &  1\t &  F775W &   F850LP \\\\ GHO 1602+4329\t  & 0.92 &  1\t &  F606W &   F814W  \\\\ GHO 1602+4312\t  & 0.90 &  1\t &  F606W &   F814W  \\\\ 1WGA J1226.9+3332 & 0.89 &  6\t &  F606W &   F814W  \\\\ MACS J0744.8+3927 & 0.70 &  1\t &  F555W &   F814W \\\\ MACS J2129.4-0741 & 0.59 &  1\t &  F555W &   F814W \\\\ MACS J0717.5+3745 & 0.55 &  1\t &  F555W &   F814W \\\\ MACS J1423.8+2404 & 0.54 &  1\t &  F555W &   F814W \\\\ MACS J1149.5+2223 & 0.54 &  1\t &  F555W &   F814W \\\\ MACS J0911.2+1746 & 0.50 &  1\t &  F555W &   F814W \\\\ MACS J2214.9-1359 & 0.50 &  1\t &  F555W &   F814W \\\\ MACS J0257.1-2325 & 0.50 &  1\t &  F555W &   F814W \\\\ \\\\ CT344\t   &  & 1 &  F606W &   F814W  \\\\ B0455\t   &  & 1 &  F555W &   F814W  \\\\ GOODS+PAN  &  & $\\sim30$  &  F775W &   F850LP   \\\\ \\hline \\end{tabular} \\hfill \\break \\footnotesize{$^1$ number of ACS field of view per filter. \\hfill\\break All clusters have coordinates and redshift listed in NED, except for MACS clusters, listed in Ebeling et al. (2007).  MACS clusters have been also studied by Stott et al. (2007). \\hfill\\break} \\end{table}    Usually, the richness dependence of the  abundance of faint red galaxies is removed by normalizing it to the number of bright red galaxies, i.e. by computing the faint-to-luminous ratio, or any related quantity, like the faint-end slope $\\alpha$ of the luminosity function. The analysis of the faint-to-luminous ratio, performed by Stott et al. (2007), or its reciprocal, the luminous-to-faint ratio by De Lucia et al. (2007), both suggest an evolution of the relative abundance of faint red galaxies, in the sense that at high redshift there is a deficit of faint red galaxies per unit bright galaxy. On the other end, Andreon (2006a) suggests no deficit of red galaxies, using a very small cluster sample, and Andreon et al. (2006) discard a considerable build up of the red sequence on the basis of fossil evidence. Evidences presented in earlier works have been discussed in the mentioned papers and references therein.  In this paper, we aim to understand if the colour--magnitude relation has been build up at early or late times, by studying many galaxy clusters at several look-back times.   Throughout this paper we assume $\\Omega_M=0.3$, $\\Omega_\\Lambda=0.7$ and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. All results of our stochastical computations are quoted in the form $x\\pm\\sigma$ where $x$ is the posterior mean and $\\sigma$ is the posterior standard deviation. ",
        "conclusions": "The history of mass assembly of bright (massive) red galaxies in clusters is pretty well known: they were assembled at early times, as testified by the passive evolution of their characteristic magnitude (e.g. de Propris et al. 1998, 2007; Andreon et al. 2007), the constancy of their stellar mass function (e.g. Andreon 2006b) and of the halo occupation number, i.e. the number of galaxies per unit cluster mass (Lin et al. 2006, Andreon et al. 2008). We stress that all mentioned works favour the above scenario, but only one, (Andreon 2006b), excludes contender models, and we emphasize that most mentioned works have samples that are dominated, but not exclusively composed, by red galaxies.   The history of mass assembly of faint red galaxies is far less clear. The present paper shows that a non-evolution of the faint end slope $\\alpha$, or any related number such as the luminous-to-faint ratio, is fully compatible with the data. This implies that the history of mass assembly of faint red galaxies is strictly parallel to the one of their massive cousins, in order to keep the relative abundance constant. Therefore, the build up of the red sequence is largely complete by $z=1.3$ down to $0.05 L^*$, and, if a differential filling is envisaged, it should occur mostly at much larger redshift. Similarly, cluster mass, as parametrized by X-ray luminosity or velocity dispersion, seems not to play any role in shaping the relative abundance of faint galaxies, contrary to some previous claims. Our claims are based on one of the largest samples, spread over the wider redshift range studied thus far with a large cluster sample, with great attention to systematics. A recent ($z<1.3$) transformation of many blue galaxies in faint red galaxies would modify the faint-end slope of the LF, change the $F/L$ ratio and inflate the color scatter of the colour-magnitude relation, none of which have been observed. Yet, a redshift trend is expected because of the spiral morphology of some faint red galaxies in nearby clusters, but a larger sample of clusters (at $z \\gg 0.5 $) is needed to measure its small amplitude. The present sample is however large enough to discard the claimed steep trends previously suggested in literature."
    },
    "0710/0710.0168_arXiv.txt": {
        "abstract": "We present observations of the HCN and HCO$^+$ J=1-0 transitions in the center of the nearby spiral galaxy NGC 6946 made with the BIMA and CARMA interferometers. Using the BIMA SONG CO map, we investigate the change in the $I_{\\rm HCN}/I_{\\rm CO}$  and $I_{\\rm HCO+}/I_{\\rm CO}$ integrated intensity ratios as a function of radius in the central kiloparsec of the galaxy, and find that they are strongly concentrated at the  center. We use the 2MASS $K_S$ band image to find the stellar surface density, and then construct a map of the hydrostatic midplane pressure. We apply a PDR model to the observed $I_{\\rm HCN}/I_{\\rm HCO+}$ integrated intensity ratio to calculate the number density of molecular hydrogen in the dense gas tracer emitting region, and find that it is roughly constant at $10^5$ cm$^{-3}$ across our map. We explore two hypotheses for the distribution of the dense gas. If the HCN and HCO$^+$ emission comes from self-gravitating density peaks inside of a less dense gas distribution, there is a linear proportionality between the internal velocity dispersion of the dense gas and the size of the density peak. Alternatively, the HCN and HCO$^+$ emission could come from dense gas homogeneously distributed throughout the center and bound by ambient pressure, similar to what is observed toward the center of the Milky Way. We find both of these hypotheses to be plausible. We fit the relationships between $I_{\\rm HCN}$, $I_{\\rm HCO+}$, and $I_{\\rm CO}$. Correlations between the hydrostatic midplane pressure and $I_{\\rm HCN}$ and $I_{\\rm HCO+}$ are demonstrated, and power law fits are provided. We confirm the validity of a relation found by \\citet{BR2006} between pressure and the molecular to atomic gas ratio in the high hydrostatic midplane pressure regime ($10^6$-$10^8$ cm$^{-3}$ K). ",
        "introduction": "How is molecular gas organized in  galaxies? In the disk of the Milky Way, molecular gas as traced by CO is concentrated in giant molecular clouds (GMCs) with a characteristic size of 50 pc \\citep{B1987}. These clouds have masses around 10$^4$--10$^6$ $M_\\odot$ and contain 80\\% of the molecular mass \\citep{B1993}. The hydrostatic pressure from the self-gravity of a GMC this size is $\\sim10^5$ cm$^{-3}$ K \\citep{B1993}, which is an order of magnitude larger than the mean pressure of the local ISM \\citep{B1987a}; this leads to the conclusion that GMCs in the Milky Way disk are self-gravitating. GMCs in M31 \\citep{R2007} and M33 \\citep{RB2004,RKMW2007} have similar properties to those observed in the Milky Way.  The densest parts of GMCs are traced by molecules with larger dipole moments than CO, such as HCN, HCO$^+$, and CS. These dense gas tracers need a molecular hydrogen density near $10^5$ cm$^{-3}$ to be excited collisionally, 100 times greater than CO. In the disk, emission from these molecules  is limited to the star-forming cores of GMCs \\citep{PJE1992,WE2003,HB1997b}.  However, the molecular gas composition    of the central kiloparsec of a spiral galaxy is very different from that of  its disk. In the Milky Way, widespread CO emission in the inner kpc makes it  difficult to differentiate between a distribution of discrete clouds or one in which the molecular gas is organized homogeneously \\citep{DUCD1987}. The emission from  HCN and CS within $\\sim$250 pc of the Galactic center is  distributed similarly to the CO \\citep{JHPB1996,BSWH1987}, suggesting that high  density regions are common, at least compared to the average ISM  density of $\\sim$1 cm$^{-3}$  in the solar neighborhood. In other galaxies, dense gas tracers are more concentrated than CO in the central kpc  and have lower intensities elsewhere \\citep{GS2004,HB1997a}, consistent with observations of the Milky Way \\citep{HB1997b}. Furthermore, the centers of spiral galaxies tend to have high  densities and pressures;  bulges  have hydrostatic midplane pressures 2-3 orders of magnitude higher than those of their disks \\citep{SB1992}.  \\citet{BR2006} (hereafter BR06) found a simple relationship in disk galaxies between  the hydrostatic midplane pressure and the ratio of molecular to atomic gas that holds in disks and bulges.   In this paper, we use new  HCN and HCO$^+$ observations of the central region of NGC 6946 to study  molecular gas organization in this nearby spiral galaxy. Optically, this galaxy is classified as SAB(rs)cd \\citep{D1963}, it has a moderate starburst in its nucleus \\citep{TH1983}, and it is at a distance of 5.5 Mpc \\citep{T1988}. At this distance, 1'' corresponds to 27 pc. Observations of NGC 6946 in the R, I, and K bands reveal a small bulge (radius 15'') within a larger  bar (radius 63'')  \\citep{DTVR2003}. Recent studies revealed the rich details of its nuclear structure; \\citet{SBEL2006} used high resolution images of CO to map the molecular gas spiral in the inner 10'', likely caused by an inner stellar bar 15'' in length (see also NIR images in \\citealt{ECS1998} and \\citealt{KDSB2003}). A further study used HCN emission in the central 2'' as a proxy for star formation in young clusters still embedded in their parent clouds, and combined these data with observations of HII regions to map the total amount of star formation \\citep{SBED2007}. They concluded the nuclear star formation is being fed by an accumulation of dense gas driven by the inner stellar bar.  Using  HCN and HCO$^+$ observations in conjunction with CO and near-infrared data, we investigate the distribution of dense gas in the central kiloparsec of NGC 6946. After reducing our observations, we calculate  HCN/CO and HCO$^+$/CO integrated intensity ratios, and we estimate the pressure using two different methods. We then describe two hypotheses for the distribution of the dense molecular gas, and explore the relationship between dense gas tracers  and hydrostatic midplane pressure. ",
        "conclusions": "Using the BIMA and CARMA interferometers, we  produced high-resolution maps of the dense gas tracers HCN and HCO$^+$ in NGC 6946. The maps are rich with detail, and we have 13 independent data points across the inner $\\sim$750 pc of the galactic center in HCN and CO, and 7 in HCO$^+$. The integrated intensities of HCN and HCO$^+$ peak near the center of the galaxy, and fall off from there.  Using BIMA SONG CO 1-0 and 2MASS $K_S$ band images, we calculated an approximation to the hydrostatic midplane pressure in the center of NGC 6946. We used a PDR model to calculate the number density in the dense gas regions traced by HCN and HCO$^+$, and found that the number density in these regions does not vary as a function of radius or the surrounding hydrostatic midplane pressure. We explored two hypotheses for the distribution of the dense gas in the central kpc of NGC 6946. If the dense gas is concentrated in clumps, we showed that self-gravity implies that $v_{\\rm int}\\propto r$. However, we also explored another plausible dense gas distribution; the dense gas could instead be homogeneously distributed throughout the galactic center if thermal pressure provides enough support for hydrostatic equilibrium. We then demonstrated clear correlations between the HCN and HCO$^+$ integrated intensities and the midplane pressure.  Finally, we confirmed the validity of the relation between hydrostatic midplane pressure and the molecular to atomic gas ratio found in BR06 in the high pressure regime."
    },
    "0710/0710.0397_arXiv.txt": {
        "abstract": "Our understanding of Pulsar Wind Nebulae (PWNe), has greatly improved in the last years thanks to unprecedented high resolution images taken from the HUBBLE, CHANDRA and XMM satellites. The discovery of complex but similar inner features, with the presence of unexpected axisymmetric rings and jets, has prompted a new investigation into the dynamics of the interaction of the pulsar winds with the surrounding SNR, which, thanks to the improvement in the computational resources, has let to a better understanding of the properties of these objects. On the other hand the discovery of non-thermal emission from bow shock PWNe, and of systems with a complex interaction between pulsar and SNR, has led to the development of more reliable evolutionary models. I will review the standard theory of PWNe, their evolution, and the current status in the modeling of their emission properties, in particular I will show that our evolutionary models are able to describe the observations, and that the X-ray emission can now be reproduced with sufficient accuracy, to the point that we can use these nebulae to investigate fundamental issues as the properties of relativistic outflows and particle acceleration. ",
        "introduction": "Pulsar Wind Nebulae (PWNe) are bubbles of relativistic particles and magnetic field created when the ultra-relativistic wind from a pulsar interacts with the ambient medium, either SNR or ISM. The best example of a PWN, often considered the prototype of this entire class of objects, is the Crab Nebula. The first theoretical model of PWNe was presented by \\citet{ree74}, developed in more details by  \\citet{ken84a,ken84b} (KC84 hereafter), and is based on a relativistic MHD description. The ultra-relativistic pulsar wind is confined inside the SNR, and slowed down to non relativistic speeds in a strong termination shock (TS). At the shock the toroidal magnetic field of the wind is compressed, the plasma is heated and particles are accelerated to high energies. A bubble of high energy particles and magnetic field is produced where the post-shock flow expands at a non relativistic speeds toward the edge of the nebula.  Despite its simplicity the MHD model can explain many of the observed properties of PWNe. Acceleration at the TS accounts for the continuous, non-thermal, very broad-band spectrum, extending from Radio to X-rays, with spectral index in the range 0-1.2, steepening with increasing frequencies, and modeled as synchrotron emission \\citep{ver93,ban99,wei00,wil01,mor04}. The under-luminous region, centered on the location of the pulsar, is interpreted as the ultra-relativistic unshocked wind. Polarization measures \\citep{wil72,vel85,sch79,hic90,mic91} show that emission is highly polarized and the nebular magnetic field is mostly toroidal, as one would expect from the compression of the pulsar wind. This model also predicts that PWNe should appear bigger at smaller frequencies: high energy X-rays emitting particles have a short lifetime for synchrotron losses, and they are are present only in the vicinity of the TS; in contrast the synchrotron lifetime for Radio particles is longer than the age of the nebula, so they fill the entire volume. This increase in size at smaller frequencies is observed in the Crab Nebula \\citep{ver93,bie97,ban98}. If one considers the pressure anisotropy due to the compressed nebular toroidal magnetic field  \\citep{beg92,van03}, it is also possible to recover the elongated axisymmetric shape of many PWNe ({\\it i.e.} Crab Nebula, 3C58).  By comparing observations with the predictions of the simple spherically symmetric MHD model it is possible to  constraint some of the properties of the pulsar wind, at least at the distance of the TS. To explain the dynamics of the plasma, as well as the emission properties of the nebula, the Lorentz factor of the wind is estimated to be $\\sim 10^6$, and the ratio between Poynting flux and kinetic energy $\\sigma\\sim 0.003$. This shows that nebular properties can be used to derive informations on the conditions of the pulsar wind at large distances. ",
        "conclusions": "\\label{sec:concl}  In the last few years, the combination of high resolution observations, and numerical simulations, has improved our understanding of the evolution and internal dynamics of PWNe. We can reproduce the observed jet-torus structure and we can relate  the formation of the jet in the post shock flow to the wind magnetization. Simulated maps can reproduce many of the observed features, including the details of spectral properties. Results suggest that, the best agreement is achieved in the case of a striped wind, even if MHD simulations are not able to distinguish between dissipation of the current sheet in the wind or at the TS. Results also suggest that it is  possible to use X-rays imaging to constrain the pulsar wind properties; already the rings and tori observed in many PWNe have been used to determine the spin axis of the pulsar \\citep{rom05}. Interestingly in Crab the inner ring does not appear boosted, while the wisps (which are interpreted as its optical counterpart) are.  There are still however unsolved questions, and possible future developments for research in this field.  All present simulations are axisymmetric, and none is able to address the problem of the stability of the toroidal field, and cannot reproduce the observed variability in the jet.  It is not clear if small scale disordered field is present in the inner region (may be a residual of the dissipation in the TS of the striped wind). A combination of simulations and polarimetry might help to answer this question.    \\begin{theacknowledgments} This work was supported partly by NASA through Hubble Fellowship grant HST-HF-01193.01-A, awarded by the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., for NASA, under contract NAS 5-26555. \\end{theacknowledgments}"
    },
    "0710/0710.3394_arXiv.txt": {
        "abstract": "{}{We analyze systematics in the asteroseismological mass determination methods in pulsating \\pg\\ stars.}{We compare the seismic masses resulting from the comparison of the observed mean period spacings with  the  usually adopted  asymptotic period spacings, $\\Delta \\Pi_{\\ell}^{\\rm  a}$, and  the average  of the computed period spacings, $\\overline{\\Delta \\Pi_{\\ell}}$. Computations are based on full PG1159  evolutionary models with  stellar masses ranging from $0.530$ to  $0.741 M_{\\odot}$ that take into account the complete evolution of progenitor stars.} {We conclude that asteroseismology  is a precise and powerful technique that determines the masses to a high internal accuracy, but it depends on the adopted mass determination method. In particular, we find that in the case of pulsating \\pg\\ stars characterized by short pulsation periods, like  \\ppg\\ and  \\pp, the employment of the asymptotic period  spacings overestimates the stellar mass by about 0.06 \\msun\\ as compared with inferences from the average of the period spacings. In this case, the  discrepancy between asteroseismological and spectroscopical masses is markedly reduced when use is made of the mean period spacing $\\overline{\\Delta \\Pi_{\\ell}}$ instead of the asymptotic period spacing $\\Delta \\Pi_{\\ell}^{\\rm  a}$.}{} ",
        "introduction": "\\label{intro}  Pulsating PG1159 stars (or GW Virginis) are evolved hot stars that can pose constraints  to the  stellar evolution theory  of post-Asymptotic Giant Branch (AGB).  These variable  stars belong to the population of hydrogen--deficient  objects characterized by  surface layers  rich in helium, carbon and oxygen (Werner  \\& Herwig 2006) that are considered the  evolutionary  link  between   post-AGB  stars  and  most  of  the hydrogen-deficient  white dwarfs.  The  origin of  most \\pg\\  stars is traced  back  to  the   occurrence  of  post--AGB  thermal  pulses:  a born-again episode induced either by  a very late thermal pulse (VLTP) experienced  by  a hot  hydrogen-rich  white  dwarf  during its  early cooling phase ---  see Herwig et al. (1999),  Bl\\\"ocker (2001), Lawlor \\& MacDonald (2003),  Althaus et al.  (2005), Miller  Bertolami et al. (2006), or a late thermal pulse (LTP) during which hydrogen deficiency is a result of a dredge--up episode.  (see Bl\\\"ocker 2001). During the VLTP,  the convection  zone driven  by  the helium  flash reaches  the hydrogen--rich envelope of the star,  with the result that most of the hydrogen content is burnt.   About  a   third  of spectroscopic  \\pg\\  stars   exhibit multiperiodic  luminosity variations  with periods in  the range  $300-3000$ s,  attributable to global nonradial $g$-modes pulsation (e.g. Quirion et al. 2007). The presence of a pulsational  pattern in  many \\pg\\  stars has  allowed  researchers to infer structural parameters ---  particularly the stellar mass --- and the evolutionary status of  individual pulsators  --- e.g.  Kawaler \\& Bradley (1994), Kawaler  et al.  (1995),  O'Brien et al.  (1998),  Vauclair et al.  (2002)  and more recently  C\\'orsico \\& Althaus  (2006).  Stellar masses of PG1159 stars can  independently be assessed by comparing the values of  log $ g $  and \\lteff, as inferred  from detailed non--LTE model  atmospheres (Werner  et  al.  1991),  with  tracks coming  from stellar evolution  modeling, i. e.  the  spectroscopic mass (Dreizler \\& Heber 1998,  Werner \\& Herwig 2006).  These two different approaches enable us to compare the derived stellar masses.  Recently, considerable  observational and theoretical  effort has been devoted  to the  study  of some  pulsating  \\pg\\ stars.   Particularly noteworthy is the  work of Fu et al. (2007) who  have detected a total of 23  frequencies in \\pp\\ and Costa  et al.  (2007)  who have enlarged to 198 the total  number of pulsation modes in \\pgs, making  it the  star with  the largest  number  of modes detected   besides   the  Sun.    Parallel   to  these   observational breakthroughs, substantial progress in the theoretical modeling of \\pg\\ stars has been possible (Herwig et al. 1999, Althaus et al. 2005, Lawlor \\& Mac Donald 2006). In  this sense, the new generation of \\pg\\ evolutionary models recently developed  by Miller Bertolami \\& Althaus (2006) (hereinafter  MA06)  has  proved  to be  valuable  at  deriving structural  parameters of pulsating  \\pg\\ on  the basis  of individual period fits --- see C\\'orsico et al. 2007a and C\\'orsico et al. 2007b, respectively, for an application  to the hot pulsating \\rxj\\ and the coolest  member of the class, \\pp.  These evolutionary models  are   derived  from  the  complete   evolutionary  history  of progenitor  stars  with  different  stellar masses  and  an  elaborate treatment  of the  mixing and  extramixing processes  during  the core helium burning and born again  phases.  The success of these models at explaining both the spread in surface chemical composition observed in \\pg\\ stars  and the location  of the GW  Vir instability strip  in the $\\log  T_{\\rm eff}-  \\log g$  plane  (C\\'orsico et  al.  2006) renders reliability to the inferences drawn from individual pulsating \\pg.   As shown in  MA06 the employment of detailed  \\pg\\ evolutionary models yields  spectroscopical masses  that are  systematically lower  --- by about 0.05 \\msun\\ --- than those derived from hydrogen--rich post--AGB tracks  (Werner  \\& Herwig  2006).   Most  importantly, the  resulting asteroseismological masses (as inferred from the period spacings) are usually 10\\%  higher than the  new spectroscopical masses,  except for the hot pulsating PG1159 star  \\rxj, the spectroscopical mass of which is more  than 20\\% higher than the  asteroseismological one (C\\'orsico et al.   2007a).  The  mass discrepancy is  a clear indication  of the uncertainties weighting upon the mass determination methods, even though the spectroscopic uncertainties are of that order.  In  an attempt to  understand the  persisting discrepancy  between the asteroseismological  and spectroscopical  masses, Miller  Bertolami \\& Althaus (2007) have recently shown  that previous evolution is not the dominant factor in shaping hydrogen--deficient post--VLTP tracks. They conclude that the MA06 \\pg\\ tracks are robust enough as to be used for spectroscopical  mass   determinations  of  \\pg--type   stars,  unless opacities in  the intershell  region are strongly  subestimated. Their results   make   clear  that   the   systematic  discrepancy   between asteroseismological   and  spectroscopical   masses   should  not   be attributed to uncertainties in post--AGB tracks; rather, they call for the   need   of  an   analysis   of   possible   systematics  in   the asteroseismological mass determination methods.  This is precisely the core feature of the present work. Specifically, we will concentrate on the  usually adopted  asymptotic period  spacing approach  (Kawaler et al.  1995; O'Brien  et al.   1998; Vauclair  et al.   2002; Fu  et al. 2007) used in  most mass  determinations of individual  pulsating \\pg\\ stars. The advantage  of this approach lies in the  fact that the mean period spacing  of PG1159 pulsators  depends primarily on  the stellar mass (Kawaler  \\& Bradley 1994; C\\'orsico \\&  Althaus 2006).  However, the derivation  of the stellar  mass using the  asymptotic predictions may  not be  entirely reliable  because  they are  strictly valid  for chemically homogeneous stellar models, while PG1159 stars are expected to be  chemically stratified with  strong chemical gradients  built up during  the progenitor  star life.   We will  show that  this approach overestimates the seismic  mass for those pulsating \\pg\\  stars on the white dwarf  cooling track.   We will also  show that  the discrepancy between  asteroseismological  and  spectroscopic  masses  is  markedly alleviated if the average of  the computed period spacings, instead of the asymptotic  ones, is used. In  the next Section,  we summarize the seismological tools to  infer the stellar mass from  the observed mean period spacings. We also describe the evolutionary sequences employed. In Sect. 3 and \\ref{period-spacing} we present our results and compare them with  other mass  determinations methods. We  close the  paper in Sect. \\ref{conclusions} by summarizing our findings. ",
        "conclusions": "\\label{conclusions}  This paper explores the systematic discrepancy between spectroscopical and asteroseismological masses of pulsating \\pg\\ stars.  Our  motivation is the result  of Miller Bertolami \\& Althaus  (2007) that such  discrepancy should not be  attributed to uncertainties in post--AGB tracks,  but possibly to systematics in the asteroseismological mass determination  methods. Recently, Quirion has pointed to  one of us (M3B)  that a possible  opacity change resulting from the spread of He/C/O abundances in PG1159 stars could be a source of uncertainty in the location of the tracks.  We addressed this issue by calculating sequences in which helium and carbon are changed in the whole envelope above the helium burning  shell. We find that changing helium into carbon by an amount of 0.4  by mass shifts the track by only 0.02 dex in effective  temperature (being bluer if carbon  is higher). This translates  into  a shift  of  only 0.005  and  0.015  \\msun\\ for  the spectroscopic mass near the  0.51 and 0.6 \\msun\\ tracks, respectively. Thus, the  precise values  of the He/C/O  abundances do not  seem to introduce appreciable changes in the masses derived by MA06.  Specifically, we  have concentrated on the seismic  masses that result from a  comparison of  the observed period  spacings with  the usually adopted asymptotic period  spacings ($\\Delta \\Pi_{\\ell}^{\\rm a}$) used in  most mass  determination  of individual  pulsating  \\pg\\, and  the better    suited   average   of    the   computed    period   spacings ($\\overline{\\Delta \\Pi_{\\ell}})$. On  the basis  of  full  \\pg\\ evolutionary  models  that consider  the evolutionary  history  of  progenitor  stars  (MA06),and  the  ensuing internal chemical  profile, we have  shown that the derivation  of the stellar mass using the asymptotic period spacing is not appropriate in the case of \\pg\\ stars.   In particular, we demonstrate that for those pulsating \\pg\\ stars characterized  by short pulsation periods, i. e., the pulsating  \\pg\\ stars  on the hot  white dwarf regime  (DOVs), the asymptotic  $\\Delta \\Pi_{\\ell}^{\\rm a}$  differs appreciably  (by more than 1 s)  from the mean $\\overline{\\Delta \\Pi_{\\ell}}$.   Only in the case   of   variables   with    long   periods   (PNNVs),   like   the high--luminosity, log--gravity pulsating \\pg\\  stars, do the $g-$ mode period  spacings  given  by  asymptotic  $\\Delta  \\Pi_{\\ell}^{\\rm  a}$ resemble those predicted by mean $\\overline{\\Delta \\Pi_{\\ell}}$.  This is expected  because the asymptotic  conditions are approached  in the limit of very high radial order $k$.  For  quantitative  inferences,  we  have  computed  the  seismic  mass resulting from the employment of the asymptotic and the average of the computed  period  spacing  for  those  pulsating  \\pg\\  which  have  a sufficiently large number of detected modes to infer an observed value of the mean period spacing. Our  selected stars are listed in Table 1, together  with the  stellar  mass inferences. The employment of the asymptotic theory, in principle formally valid for chemically homogeneous stellar models at high radial index k, overestimates the seismic mass by about 0.06 \\msun\\ in the case of  very short period  pulsating \\pg\\  stars like \\ppg\\  and \\pp. Because \\pg\\ stars  are  expected to  be chemically  stratified, estimations  of the stellar  mass from mean  $\\overline{\\Delta \\Pi_{\\ell}}$ are  more  realistic than  those  inferred by  means of asymptotic $\\Delta \\Pi_{\\ell}^{\\rm    a}$.    Indeed,    stellar   mass derived   from the mean $\\overline{\\Delta  \\Pi_{\\ell}}$ are  in good  agreement with  the mass values obtained  from detailed  period fittings. The  discrepancy between  asteroseismological   and  spectroscopical  masses   is  markedly alleviated by  the employment  of the average  of the  computed period spacing instead of the asymptotic period spacings.  In  closing,  a  Fortran  program  to derive,  from  our  evolutionary sequences,  averages  of  the  period  spacing  for  arbitrary  period intervals      is       available      at      our       web      site http://www.fcaglp.unlp.edu.ar/evolgroup."
    },
    "0710/0710.4057_arXiv.txt": {
        "abstract": "{The shell-type supernova remnant RX J1713.7--3946 was observed during three years with the H.E.S.S. Cherenkov telescope system. The first observation campaign in 2003 yielded the first-ever resolved TeV gamma-ray image. Follow-up observations in 2004 and 2005 revealed the very-high-energy gamma-ray morphology with unprecedented precision and enabled spatially resolved spectral studies. Combining the data of three years, we obtain significantly increased statistics and energy coverage of the gamma-ray spectrum as compared to earlier H.E.S.S. results. We present the analysis of the data of different years separately for comparison and demonstrate that the telescope system operates stably over the course of three years. When combining the data sets, a gamma-ray image is obtained with a superb angular resolution of 0.06 degrees. The combined spectrum extends over three orders of magnitude, with significant gamma-ray emission approaching 100 TeV. For realistic scenarios of very-high-energy gamma-ray production, the measured gamma-ray energies imply efficient particle acceleration of primary particles, electrons or protons, to energies exceeding 100 TeV in the shell of RX J1713.7--3946.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.1444_arXiv.txt": {
        "abstract": "One of the promising methods to search for life on extra-solar planets (exoplanets) is to detect its signature in the chemical disequilibrium of exoplanet atmospheres. Spectra at the modest resolutions needed to search for methane, oxygen, carbon dioxide, or water will demand large collecting areas and large diameters to capture and isolate the light from planets in the habitable zones around the stars. Single telescopes with coronagraphs to isolate the light from the planet will have to be 8\\,m or more in diameter to generate sample sizes with a reasonable probability of finding at least one life-bearing planet; interferometers of smaller telescopes can overcome some of the limitations but will still need large similarly large collecting areas. Even larger telescopes will be needed to detect atmospheric signatures in transiting planets. In all cases, the sample sizes increase as the third power of telescope diameter. Direct observation using coronagraphs or interferometers are most sensitive to planets around stars with masses similar to that of the Sun, whereas transit observations favor low-mass stars near the nuclear burning limit. If the technical difficulties of constructing very large space telescopes can be overcome, they will be able to observe planets near hundreds to thousands of stars with adequate resolution and sensitivity to look for the signatures of life. ",
        "introduction": "The detection of more than 200 planets outside the Solar System is a powerful incentive to search for extra-terrestrial life. Although extra-terrestrial life could take on many guises, economy of hypothesis (and a practical approach) implies that we should first search for signs of life similar to those seen on Earth. An advantage of using the Earth as a proxy for analysis is that it bounds the problem and provides concrete examples of signatures subject to passive detection, i.e. not requiring signals broadcast by sentient beings.  In the Earth's present atmosphere, the chemical components have been altered by life. Evidence for life on Earth could be detected from afar in the spectral signatures of these molecules: oxygen, carbon dioxide, methane, and water vapor \\citep{sag93}. The rise of photoplankton and plants on Earth created an atmosphere with a large reservoir of oxygen today that requires steady production by photosynthesis to maintain its present level. If we could study the same spectral signatures in the atmospheres of exoplanets, we could search for signs of life similar to some of the earliest and most robust forms on our own planet. But the Earth's atmosphere has been altered by life several times over the last 4\\,Gyr, and there are potentially many signatures that could indicate the presence of life on exoplanets \\cite{tra02, sea02, kalt07}. Chemical signatures of life on other planets would revolutionize our thinking about Earth's uniqueness and provide tantalizing evidence that we are not alone in the universe.  Observing exoplanets directly is difficult owing to their proximity to the much brighter stars that keep them warm. Although technically challenging, this problem is well understood, and there are a variety of strategies that can reduce the brightness of the starlight without diminishing the light from the planet for direct detection \\citep{guy06,cas06} or use the star itself as a background source to probe the atmosphere when the exoplanet transits the face of the star \\citep{cha02, ehr06}. The first technique must overcome diffraction in the pupil of the telescope, a well understood phenomenon, and it is possible to characterize the detection problem in general terms to understand the kinds of instruments that will be needed to study exoplanets and search for signs of life. The second technique depends only on the photometric accuracy of an observation and is easy to calculate for any star.  Any planet supporting life as on Earth must satisfy two broad criteria: (1) it must have surface temperatures in the range 273 to 373\\,K, where water is in the liquid phase, and (2) it must have an atmosphere. The first criterion is met if the planet is in the {\\it habitable zone} (HZ) around the star, a range of orbital distances where the equilibrium temperature for a rotating body is between the freezing and boiling points of water. The second criterion is met if the planet is rocky and can retain an atmosphere; current estimates specify a mass between 0.5 to 10 Earth masses. Smaller planets will not retain their atmospheres, and larger planets accrete gas and become gas giants. These criteria are necessary but not sufficient to create Earth-like life. Although they are probably far too restrictive to encompass all the possibilities for other life forms in the universe---or even on Earth itself---they are the only ones amenable to remote observation with technology that we can foresee at present and thus provide a good basis for a targeted search.  There have been many calculations aimed at refining our ideas of the habitable zone \\citep[e.g.][]{kast93}, suitable samples of stars to search \\citep{tur03, tur04, tur06}, and the impact of specific telescopes on such a search \\citep{alg07}, and there is some disagreement about the likelihood of success depending on the different assumptions used. Most authors concentrate on photometric detection alone, ignoring the means to search for life (e.g. Agol 2007).  But the utility of photometric searches alone to identify exoplanets for subsequent study could be obviated by the difficulty of measuring the orbits accurately enough to recover them at a later time \\cite{bro05, bro07}. For apertures sizes under discussion for the Terrestrial Planet Finder (TPF) mission, habitable zones of most potential target stars are blocked by the coronagraph. The planets of interest will spend the majority of time behind the central obscuration of the imaging instrument. To ensure efficient recovery of a newly discovered planet at future observing epochs, it will be necessary to estimate the orbit to high accuracy from a small number of astrometric measurements. This requirement implies a lower limit to the aperture size based on operational requirements (Brown et al. 2007). Achieving adequate astrometric precision may demand apertures larger than any so far discussed for TPF (Brown 2007, personal communication). If true, discovery and immediate spectroscopy of candidate sources near the stars may be the most efficient means of identifying those most interesting for follow up observations to determine if life is, indeed, present, making it essential to understand the spectroscopic requirements at the outset.  The purpose of this article is to derive the main scaling parameters for the study of life-bearing exoplanets in known samples of stars to understand the size of the telescopes needed for a robust search. We adopt simple but optimistic assumptions to bound the problem and find the minimum size for survey telescopes. Using only the lowest order approximations and assuming ``best case'' observing conditions allows robust conclusions about the scale of facilities needed to tackle the search for life.  The main premise is that direct photometric detection of exoplanets in a band where the exo-planet atmosphere is free of chemical signatures cannot be the endpoint of any mission to search for life-bearing planets; spectra of the atmospheres will be the major advance of observing the planets directly. Moreover, the rapid increase in the number of exoplanets discovered to date suggests that finding terrestrial planets will be easiest with indirect methods, such as observing radial velocity, photometric or astrometric variations in the host stars, and the real thrust of direct observations will be to  search for signs of life. ",
        "conclusions": ""
    },
    "0710/0710.1734_arXiv.txt": {
        "abstract": "When two stars collide and merge they form a new star that can stand out against the background population in a starcluster as a blue straggler. In so called collision runaways many stars can merge and may form a very massive star that eventually forms an intermediate mass blackhole. We have performed detailed evolution calculations of merger remnants from collisions between main sequence stars, both for lower mass stars and higher mass stars. These stars can be significantly brighter than ordinary stars of the same mass due to their increased helium abundance. Simplified treatments ignoring this effect give incorrect predictions for the collision product lifetime and evolution in the Hertzsprung-Russell diagram. ",
        "introduction": "In star clusters stars can experience close encounters with other stars, which in some cases can lead to the collision and merger of two or more stars. This is a possible formation mechanism for blue straggler stars.  Blue stragglers are stars that appear on the extension of the main sequence in the colour-magnitude diagram (CMD) of star clusters (\\citet{article:piotto_freq_bss} and figure \\ref{fig:hrd_m67}). In the past, various mechanisms for their formation have been proposed (see \\emph{e.~g.~} \\citet{article:livio_blue_stragglers} for a list), but the most commonly accepted explanation is that they are stars that have gained mass long after they were formed, either through mass transfer in a close binary or by merging two stars. Such a merger can in turn be the result of normal binary evolution or of a collision with another star. These mechanisms are thought to operate and produce blue stragglers in different regions in a cluster \\citep{article:davies_bss_formatio}. We have studied mergers that result from collisions.  In some cases a sequence of collisions in the centre of a cluster can lead to a runaway where many stars merge together \\citep{article:portegies_zwart_intermediate_mass_blackhole}. Such runaways may lead to the formation of very massive stars and ultimately intermediate mass blackholes. We have studied the evolution of the outcome of a single collision.  One of the goals of the MODEST collaboration \\citep{article:modest1} is to study the evolution of stellar mergers by combining stellar evolution, stellar hydrodynamics and stellar dynamics in one software framework. This work is the first result from combining stellar hydrodynamics and stellar evolution in one framework. ",
        "conclusions": "Stellar collisions are a means of making stars that appear in unusual locations in the colour magnitude diagram, such as blue stragglers. Remnants from high mass stars that merge after the end of core hydrogen burning can become much brighter while crossing the Hertzsprung-gap and stay there much longer.  Approximating mergers with ordinary stellar models or fully mixed models can give significantly different results compared to proper detailed models."
    },
    "0710/0710.1028_arXiv.txt": {
        "abstract": "We present new spectroscopic observations of the stellar cluster population of region~B in the prototype starburst galaxy M82 obtained with the Gillett Gemini-North 8.1-metre telescope. By coupling the spectroscopy with \\emph{UBVI} photometry acquired with the Advanced Camera for Surveys (ACS) on the Hubble Space Telescope (HST), we derive ages, extinctions and radial velocities for seven young massive clusters (YMCs) in region B. We find the clusters to have ages between 70 and 200\\,Myr and velocities in the range 230 to 350\\kms, while extinctions  $A_V$ vary between $\\sim$1--2.5 mag. We also find evidence of differential extinction across the faces of some clusters which hinders the photometric determination of ages and extinctions in these cases. The cluster radial velocities indicate that the clusters are located at different depths within the disk, and are on regular disk orbits. Our results overall contradict the findings of previous studies, where region~B was thought to be a bound region populated by intermediate-age clusters that formed in an independent, offset starburst episode that commenced 600\\,Myr--1\\,Gyr ago. Our findings instead suggest that region B is optically bright because of low extinction patches, and this allows us to view the cluster population of the inner M82 disk, which probably formed as a result of the last encounter with M81. This study forms part of a series of papers aimed at studying the cluster population of M82 using deep optical spectroscopy and multi-band photometry. ",
        "introduction": "The extensively studied galaxy M82 is a local example of a nuclear starburst galaxy. \\citet[hereafter OM78]{om78} first catalogued the complex star-forming regions seen in ground-based images of the M82 disk, and introduced the nomenclature A--H. Region~B is the brightest region in the disk and is located 350--1050\\,pc north-east of the nucleus. OM78 and \\citet{marcum96} find that the integrated spectrum of this region is indicative of a fossil starburst region as it shows the characteristic `E$+$A' post-starburst spectrum, suggestive of a truncated burst of star formation occurring 100--1000\\,Myr ago. Moreover, they find that the intrinsic brightness of region~B is such that the burst must have been of comparable intensity to the present starburst in the nucleus.  More recently, de Grijs, O'Connell, \\& Gallagher (2001) studied the M82-B cluster population using photometry obtained with the Wide Field and Planetary Camera 2 (WFPC2) and the Near-Infrared Camera and Multi-Object Spectrometer (NICMOS) onboard the Hubble Space Telescope (HST). They identified 113 clusters and, by estimating ages and extinctions from $BVI$ photometry, they found that a concentrated episode of cluster formation occurred 400--1000\\,Myr ago with a peak at 600\\,Myr. Subsequent to this analysis, de Grijs, Bastian, \\&~Lamers~(2003a) re-derived ages and extinctions by combining $BVI$ and $JH$ photometry; they find that the peak of cluster formation occurred at 1.10\\,Gyr with an age spread of 500--1500\\,Myr. de Grijs, Bastian, \\&~Lamers~(2003b) used these new ages to derive the cluster luminosity function (CLF) for a fiducial age of 1.0\\,Gyr and find that it has a Gaussian shape, similar to globular clusters (GCs), rather than the standard power-law CLF found for populations of young massive clusters (YMCs) \\citep[e. g. ][]{larsen04,gieles06a}. This result is surprising but supports theoretical models advocating that an initial power-law distribution of cluster masses will be transformed into a Gaussian distribution because low mass clusters will be preferentially disrupted \\citep[e.g.][]{fall01, vesperini03, gieles06b}. \\citet{RdG03-2} therefore suggest that because M82-B is of intermediate age, it provides the `missing evolutionary link' between the power-law CLFs found for YMCs and the Gaussian CLFs of globular clusters. However, for an environment as turbulent as that of M82, one would expect a timescale of several Gyr for preferential disruption to produce this effect \\citep[e. g. ][where such a timescale is calculated for NGC 1316 and NGC 3610 respectively]{goud1316,goud3610}.  To date, spectroscopy of the region~B cluster population has been extremely limited; the only published spectra are those of \\citet{stis} for the two brightest members. The spectra were obtained with the Space Telescope Imaging Spectrograph (STIS) and, although they are of low quality, the derived ages of 350$\\pm$100\\,Myr are much lower than the photometrically-based ages of $\\sim$0.7--6~Gyr \\citep{stis, RdG03-1} for the same clusters. This discrepancy hints at the possibility that region~B may be younger and that the ages derived from the photometry are too high. It is important to determine accurate ages for the M82-B cluster population to verify its unusual CLF, and for understanding the cluster formation history of M82 and its relationship to encounters with its close neighbour M81.  We therefore acquired new spectroscopic and photometric data for the M82-B cluster population. By using both techniques we are able to obtain information for a large number of clusters and also overcome the degeneracy between age and extinction that presents a hurdle in the analysis of photometric data \\citep{gelys07a, gelys07b}. In this paper, we present optical spectroscopy for seven clusters obtained with the Gillett Gemini \\mbox{8.1-m} telescope, and in a companion paper \\citep{phot}, we present new HST imaging, including $U$-band photometry of the clusters. ",
        "conclusions": "We have presented Gemini GMOS spectroscopy and ACS \\emph{UBVI} photometry for seven clusters in region~B of the starbust galaxy M82. Our aims were to obtain accurate ages, extinctions and radial velocities for the clusters.  We have used both photometric and spectroscopic methods to determine cluster ages. Our main age determination method is fitting the Balmer spectral lines with all available models and finding the best fit (i.e. the lowest $\\chi^2$ on the overall line profile fit), in order to eliminate the known Balmer line strength degeneracy. We find our method to agree with photometric ages derived using the 3DEF method (based on $UBVI$ colours). This demonstrates that the photometrically-derived ages are accurate but they are not as precise as the ones obtained from spectroscopy and confirms that the only truly independently reliable method for age-dating clusters in extragalactic environments is spectroscopy. The inclusion of $U$-band photometry may enhance the accuracy of photometric measurements, but unless they can be cross-checked by spectroscopic means, the potential age degeneracy may not be broken.  We find cluster ages in the range of 80 to 200\\,Myr, a distribution that is consistent with the timescale for the last encounter between M82 and M81 as proposed by \\citet{yun99}, namely 220\\,Myr. These findings in combination with the larger photometric sample of \\citet{phot}, disagree with the `fossil starburst' scenario proposed by \\citet{RdG01}, where region B in M82 was identified with the remnant of an off-centre starburst that commenced about 600\\,Myr ago, a time scale based on their photometric determinations of star cluster ages. We suggest that the increased cluster formation rate in region~B is representative of the era of increased star formation across the galaxy disk triggered by the last encounter with M81.  The extinction along our line of sight in this extended region appears to vary greatly, between 1.1 and 2.5 magnitudes. As M82 is seen virtually edge-on, we find that \\mbox{M82-B} presents a view into various depths of the body of the galaxy through an arrangement of `windows' in the dust distribution \\citep[as hypothesised for cluster M82-F by][]{smith01}. We also find differential extinction across the face of cluster \\#91 (and also possibly \\#108), which sets a serious obstacle in the photometric determination of extinction and age in these cases. In fact, the effect of dust in this environment is in some cases so grave that photometry cannot be used at all as an age/extinction diagnostic.  We have also used the available spectroscopy to derive cluster kinematics. This allows us to reinforce our extinction-based findings and show that the seven clusters reside at different depths within the disk of M82. The large scatter of cluster velocities about the gently rising component of the rotation curve indicates that the clusters do not move in a co-ordinated fashion and that region B cannot be bound.  Previous studies of clusters F and its neighbour L on the western side of the disk show that they both have an age of $60\\pm20$ Myr \\citep{GS99,bastian07}. This age fits in well with the region B cluster age distribution presented in this paper and \\citet{phot}, and supports our suggestion that the increase in the cluster formation rate was not local to region B but part of a galaxy-wide burst. In addition, the very high masses of clusters F \\citep[$\\sim10^6$\\,\\Msun, ][]{smith01,bastian07} and L \\citep[$4\\times10^6$~\\Msun, ][]{mccrady07} imply that many lower mass clusters were also likely to have formed with F and L.  In summary, we find region~B to be optically bright owing to the presence of low internal extinction patches, thus offering a deep view into the M82 disk at radii between $\\sim$ 400--1200\\,pc. The range of cluster ages and other properties in this region then are typical of the evolution of the main body of M82 and reflect  the large increase in star formation that occurred about 220\\,Myr ago when M82 last passed close to M81. M82 region B stands out because it is representative of the mid-disk zone, and is relatively clear of dust, rather than being a special substructure.  This model also receives support from the larger sample of photometric ages in Smith et al. (2007), and will be further discussed in the Konstantopoulos et al. (in preparation) study of another three dozen M82 star clusters distributed across M82. \\\\"
    },
    "0710/0710.5520_arXiv.txt": {
        "abstract": "Virial mass is used as an estimator for the mass of a dark matter halo. However, the commonly used constant overdensity criterion does not reflect the dynamical structure of haloes. Here we analyze dark matter cosmological simulations in order to obtain properties of haloes of different masses focusing on the size of the region with zero mean radial velocity. Dark matter inside this region is stationary, and thus the mass of this region is a much better approximation for the virial mass. We call this mass the static mass to distinguish from the commonly used constant overdensity mass. We also study the relation of this static mass with the traditional virial mass, and we find that the matter inside galaxy-size haloes ($M\\approx 10^{12}{\\rm M}_{\\sun}$) is underestimated by the virial mass by nearly a factor of two. At $z\\approx 0$ the virial mass is close to the static mass for cluster-size haloes ($M\\approx 10^{14}{\\rm M}_{\\sun}$). The same pattern -- large haloes having $M_{\\rm vir} > M_{\\rm static}$ -- exists at all redshifts, but the transition mass $M_0 = M_{\\rm vir} = M_{\\rm static}$ decreases dramatically with increasing redshift: $M_0(z) \\approx 3\\times 10^{15}h^{-1}{\\rm M}_{\\sun} (1+z)^{-8.9}$. When rescaled to the same $M_0$ haloes clearly demonstrate a self-similar behaviour, which in a statistical sense gives a relation between the static and virial mass. To our surprise we find that the abundance of haloes with a given static mass, i.e. the static mass function, is very accurately fitted by the Press \\& Schechter approximation at $z=0$, but this approximation breaks at higher redshifts $z\\simeq 1$. Instead, the virial mass function is well fitted as usual by the Sheth \\& Tormen approximation even at $z\\lesssim 2$. We find an explanation why the static radius can be 2--3 times larger as compared with the constant overdensity estimate. The traditional estimate is based on the top-hat model, which assumes a constant density and no rms velocities for the matter before it collapses into a halo. Those assumptions fail for small haloes, which find themselves in environment, where density is falling off well outside the virial radius and random velocities grow due to other haloes. Applying the non-stationary Jeans equation we find that the role of the pressure gradients is significantly larger for small haloes. At some moment it gets too large and stops the accretion. ",
        "introduction": "The currently accepted paradigm of hierarchical clustering (\\citealt{Whi78}, \\citealt{Blu84}) provides a picture of assembly of dark matter haloes in which more massive haloes are formed through merging and accretion of smaller ones. This picture is supported for example by recent observations of ongoing mergers in clusters (e.g. \\citealt{vDok99}, \\citealt{Rin07}). The theoretical framework of the cold dark matter (CDM) has been proved successful in the description of the structure formation in the Universe from tiny fluctuations in the primordial density field (see \\citealt{Pri03} for a review). This model has received support from many different observations, in particular of the gravitational lensing effect (\\citealt{Smi01}, \\citealt{Guz02}, \\citealt{Kne03}, \\citealt{Hoe04}, \\citealt{She04}, \\citealt{Man06}), CMB \\citep{Spe07}, the abundance of clusters (\\citealt{Pie01}, \\citealt{Gla07}), and satellite dynamics (\\citealt{Zar94}, \\citealt{Pra03}). Cosmological simulations with ever increasing resolution play important role by making accurate predictions of different properties of dark matter haloes. Results of those simulations are used by other methods. For example, semi-analytical models of galaxy formation have been either incorporated into N-body simulations or use statistics such as halo mass function and merging trees, which were calibrated and tested using the simulations (e.g. \\citealt{Som99,Cro06}). The simulations reveal important information about the internal structure of dark matter haloes (e.g. \\citealt{Nav97}, \\citealt{Bul01b}, \\citealt{Tay01}) which is reflecting their underlying dynamics.  Many results of simulations implicitly use some definition of what is the size and mass of a collapsed dark matter halo. The problem is that there is no well-defined boundary of a halo: density field is smooth around the halo. The common prescription for this boundary (and hence the mass belonging to the halo) is defined through the spherical collapse model (\\citealt{Gun72}, \\citealt{Gun77}). The size of the halo at redshift $z$ is given by the half of the radius of the spherical shell at turnaround which is collapsing at that redshift. This is the \\textit{virial radius}. Thus, it is very common to measure the mass of haloes in cosmological simulations taking the particles inside a sphere of fixed spherical overdensity or to take those particles which are connected by a inter-particle separation below a given value (the friends-of-friends algorithm). We note that there is little justification for using the top-hat collapse model. Haloes do not collapse from perfect spherical homogeneous distribution. The environment of haloes is typically very non-spherical with most of accretion happening from few elongated filaments. The random velocities of the accreted matter also cannot be neglected. As the fluctuations collapse, the dark matter, which is being accreted, increases it rms velocities. This effective pressure should affect the accretion rate. The only motivation for using the top-hat model comes from simulations. Indeed, early simulations indicated that the radius of overdensity 178 is close to the virial radius \\citep{Col96}: \"the radius $r_{178}$ approximately demarcates the inner regions of haloes at $r\\lesssim r_{178}$ which are in approximate dynamical equilibrium from the outer regions at $r\\gtrsim r_{178}$ which are still infalling\". Thus, the radius of overdensity 200 ($\\approx 178$) became the virial radius. The reason why this was a good approximation is simply coincidental: the early simulations were mostly done for cluster-size haloes and, indeed, for those masses the virial radius is close to the radius of overdensity 200. The early models were flat models without the cosmological constant. Models with the cosmological constant have produced significant confusion in the community. The top-hat model must be modified to incorporate the changes due to the different rate of expansion and due to the different rate of growth of perturbations (see \\citealt{Pri97}). That path produced the so called virial radius, which for the standard cosmological model gives the radius of overdensity relative to matter of about 340 \\citep{Bul01}. Still, a large group of cosmologists uses the old overdensity 200 relative to the critical density even for the models with the cosmological constant.  In this paper, we cast some light on this subject by searching for a physical extent of dark matter haloes, which is related to the physical processes that occur around collapsed structures. \\citet{Pra06} provided first results, which indicated that spherically averaged, mean radial velocity profiles show an inner region in which there is no net infall or outflow. The size of this region in virial units is mass-dependent: for galactic-size haloes it may even reach three times the virial radius. We use this result to study the properties of the mass inside this region and in particular, to determine the redshift evolution of such mass. Recent effort (\\citealt{Wec02}, \\citealt{Zha03}) has been devoted to the analysis of the evolution of virial mass inside dark matter haloes, sometimes referred to as the mass accretion history. However, the picture presented in these works has been recently put into question. The analysis of the mass inside a fixed physical radius has revealed that galaxy-size haloes experience unphysical growth of their virial mass \\citep{Die07}. The mean background density decreases as the Universe expands, making the virial radius to increase even in the case of no accretion or mergers. This is clearly an artifact of the definition. Other works realized this issue, for example, the evolution of the spin parameter when the matter inside a fixed radius is taken into account differs from that when using the evolving $R_{\\rm vir}$ instead \\citep{Don07}. On the other hand, there is another effect apart from accretion and merging, which has not been taken into account: haloes may also grow via relaxation of the surrounding regions near them. This is an important effect as it is a reflection of the dynamical processes which are turning non-virialized mass in the outskirts of a halo into the mass associated to it. Moreover, the long-term evolution of dark matter haloes show an interesting feature for $\\Lambda$CDM cosmologies: their mass turns out to converge to an asymptotic value which depends on the definition of mass, as pointed out by \\citet{Bus05}. Thus, it is desirable that the mass of a halo is measured using a virialization-based criteria (see \\citealt{Mac03} for an interesting approach), instead of using boundaries of a given overdensity.  The interest of the measurement of the physical mass associated to dark matter haloes is not only theoretical. Indeed, it may have a great impact on the number of collapsed objects in a given range of mass, i.e. the mass function \\citep{Whi01}. Besides, it is of great importance for the process of formation and evolution of galaxies, and, hence, it is relevant for the results obtained from semi-analytical modelling of galaxy formation (e.g. \\citealt{Cro06}), to predict the main properties of observed galaxies. Thus, it turns out to be mandatory to bring attention on the mass belonging to a halo, if accurate predictions of the physics behind galaxies from cosmological simulations are to be drawn.   This paper is organized as follows: in Section~\\ref{sec:simul} we describe the set of cosmological simulations used in our analysis and the properties of the halo samples. In Section~\\ref{sec:static} we present the definition of static mass and how it is related to objects of different sizes. A brief analysis of equilibrium in this context is presented in Section~\\ref{sec:jeans}. The scaling properties of the static to virial mass relation with redshift are shown in Section~\\ref{sec:scaling}, together with a simple model derived from this scaling relation. We obtain the static mass function and compare it with analytical models in Section~\\ref{sec:mf}. In Section~\\ref{sec:evolution} we present the evolution of the static mass tracking the halo progenitors. We discuss our results in Section~\\ref{sec:discussion} and present our conclusions in Section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  In this paper we use a set of high resolution $N$--body cosmological simulations for the analysis of the mass inside the region of dark matter haloes with no net infall or outflow velocities, i.e. the static region. Haloes show a typical radial velocity pattern which depends on their mass: low-mass haloes tend to display a region with outflow in their surroundings, while cluster-size haloes show a prominent infall velocity pattern. Galaxy-size haloes show an intermediate pattern with a sharper transition between the static region and the Hubble flow. We find that the former virial radius tends to underestimate the size of the region with zero mean radial velocity for haloes with masses $10^{10}h^{-1}{\\rm M}_{\\sun}<M_{\\rm vir}\\lesssim 10^{14}h^{-1}{\\rm M}_{\\sun}$ at $z=0$. The mass inside this static region can be about a factor of two larger than the former virial mass, when the threshold defining the static radius is $5$ per cent of the virial velocity $V_{\\rm vir} =\\sqrt(GM_{\\rm vir}/R_{\\rm vir})$. Lower values of the threshold ($\\lesssim 1$ per cent) may lower this factor. However, we observe a clear trend in the $M_{\\rm static}/M_{\\rm vir}$ vs $M_{\\rm vir}$ relation which remains unchanged with this threshold: for low-mass haloes, this ratio is an increasing function of $M_{\\rm vir}$ due to reduction in the outflow. On the other hand, high-mass haloes show a larger infall for larger $M_{\\rm vir}$, making this ratio to decrease. The maximum occurs at $M_{\\rm vir}\\simeq 5\\times 10^{12}h^{-1}{\\rm M}_{\\sun}$, where the size of the static region is the largest in units of $R_{\\rm vir}$.  The mass function of objects whose mass is defined in the way presented here resembles that of Press \\& Schechter in the range of mass we have studied, but only at $z=0$. At higher redshifts the static mass function deviates significantly from it. This disagreement could be accounted for by the scatter in $M_0(z)$, although it seems more likely that the resemblance with Press \\& Schechter at $z=0$ is due to coincidence (in both shape and amplitude to within $\\sim 20\\%$) of the ratio $f_{PS}/f_{ST}$ with $M_{\\rm static}/M_{\\rm vir}$. Whatever the case may be, this resemblance is theoretically unexpected because the Sheth \\& Tormen function, which fits very well the mass function for virial masses, is derived from a more realistic approach (ellipsoidal collapse). In any case it seems clear that when the virial mass is used, the number density of dark matter haloes of a given mass at $z=0$ might be underestimated in the mass range $10^{10}h^{-1}{\\rm M}_{\\sun}<M_{\\rm vir}\\lesssim 10^{14}h^{-1}{\\rm M}_{\\sun}$.  The redshift evolution of the $M_{\\rm static}$--$M_{\\rm vir}$ relation turns out to be very weakly redshift-dependent from $z=2$ to $z=0$ when appropriate variables are used. This dependence on the redshift is encoded as an evolving mass scale $M_0$ which indicates the approximate mass at which $M_{\\rm static}=M_{\\rm vir}$ in the declining part of the relation. By doing this, it is straightforward to derive from this relation a differential equation for the evolution of the static mass, provided that the evolution of $M_{\\rm vir}$ and of $M_0$ are already known. As we have shown in this paper, the evolution of the mass scale has an acceptable fit to a power law: $ M_0\\simeq 2.8\\times10^{15} a^{8.9} h^{-1}{\\rm M}_{\\sun}$ . On the other hand, the evolution of the virial mass $M_{\\rm vir}$ of the major progenitor is well described by an exponential law in $z$, as claimed by \\citet{Wec02}. Once we have fixed the evolution of $M_{\\rm vir}$ and of $M_0$, we solve this differential equation and we find a simple model for the evolution of the static mass of the major progenitor, i.e. $M_{\\rm static}(z) =M_{\\rm static}(z=0)(1+z)^{-\\beta}e^{-\\alpha z}$, which turns out to be a very good fit to our data. In any case, there is a different behaviour for the evolution of the static mass in different mass bins. The static mass of low mass haloes is nearly constant from $z=0.5$ to $z=0$. Galaxy-size haloes keep growing at present, but the reason here is not the decreasing background density of the expanding Universe. Instead, the relaxation of the matter in the surroundings of these haloes is incorporating mass at even a higher rate than the unphysical growth of the virial mass. The same stands for clusters, although the static and virial mass are not very different at present.  The mean radial velocity profile (averaged over a halo mass bin) evolves in the way found by \\citet{Bus05} for cluster-size haloes: while the halo is forming, it accretes mass from its surroundings via infall, which is being reduced with time. This is followed by a period of outflow of unbound matter, and a final configuration thereafter consisting in a static region whose size is independent of time. Although the simulation used in that paper, was evolved until $a=100$, and our halo sample was instead tracked up to redshift $z=0$, we find a situation that is consistent with this picture. Moreover, we find that the timescale of this sequence is different for each mass bin: low-mass haloes at present are in the outflow phase, but cluster-size haloes are still in their infall phase. Interestingly, galactic-size haloes are experiencing a transition (infall to outflow) epoch at $z=0$.   A.J.C., F.P. thank the Spanish MEC under grant PNAYA 2005-07789 for their support. A. K. acknowledges support from NASA and NSF grants to NMSU. A.J.C. acknowledges the financial support of the MEC through Spanish grant FPU AP2005-1826. We thank Juan Betancort-Rijo for valuable discussion and useful comments in this work, and Douglas Rudd for suggesting some points which improved the text. We do appreciate the useful feedback from the anonymous referee. A.J.C. also thanks the Helmholtz Institute for Supercomputational Physics for their outstanding Summerschool 2006 in Supercomputational Cosmology. Computer simulations were done at the LRZ Munich, NIC Julich, and NASA Ames."
    },
    "0710/0710.2246_arXiv.txt": {
        "abstract": "The strong lensing modelling of gravitational ``rings'' formed around massive galaxies is sensitive to the amplitude of the external shear and convergence produced by nearby mass condensations. In current wide field surveys, it is now possible to find out a large number of rings, typically 10 gravitational rings per square degree. We propose here, to systematically study gravitational rings around galaxy clusters to probe the cluster mass profile beyond the cluster strong lensing regions. For cluster of galaxies with multiple arc systems, we show that rings found at various distances from the cluster centre can improve the modelling by constraining the slope of the cluster mass profile.  We outline the principle of the method with simple numerical simulations and we apply it to 3 rings discovered recently in Abell~1689. In particular, the lens modelling of the 3 rings confirms that the cluster is bimodal, and favours a slope of the mass profile steeper than isothermal at a cluster radius $\\sim 300 \\kpc$. These results are compared with previous lens modelling of Abell~1689 including weak lensing analysis.  Because of the difficulty arising from the complex mass distribution in Abell~1689, we argue that the ring method will be better implemented on simpler and relaxed clusters. ",
        "introduction": "In recent years, the modelling of cluster mass distribution has been progressively improved by (a) coupling the strong lensing (SL) analysis in cluster cores with weak lensing (WL) measurements at large radius (e.g Gavazzi 2005\\nocite{gavazzi05}; Limousin et al. 2007b\\nocite{limousin1689}; Cacciato et al. 2006 \\nocite{cacciato06}); (b) SZ measurements (e.g. Zaroubi et al. 2000 \\nocite{zaroubi01}; Dore et al.  2001\\nocite{dore01}; Seren 2007\\nocite{Seren07}), (c) joint modelling of the cluster X-ray gas distribution (e.g. Mahdavi et al.  2007\\nocite{Mahdavi07}), (d) dynamical analysis of the velocity distribution of the stars near the potential centre (e.g. Miralda-Escud\\'{e} 1995\\nocite{miraldababul95}; Sand et al. 2004 \\nocite{Sand04}; Gavazzi 2005\\nocite{gavazzi05} and Koopmans et al. 2006\\nocite{Koopmans06}).  Despite these important improvements, it is not fully proved that cluster Dark Matter (DM) distribution closely follows the ``universal'' profile (Navarro, Frenk \\& White 1996\\nocite{nfw}, hereafter NFW) predicted by N-body numerical simulations. For this profile, the DM space density is cuspy at the centre ($r \\propto r^{-1}$ for $r \\la r_s$) and behaves as $r^{-3}$ outward. Measuring with great accuracy these two main characteristics is challenging and no consensus has arisen yet with the present day lensing observations.  Recent analysis has shown that S\\'{e}rsic profile (Merritt et al.  2005\\nocite{Merritt05}) is also fitting the universal mass profile of numerical simulations. Importantly, the mass profile has only a weak dependence on the halo mass or cosmology, allowing to stack different measurements together to improve their significance.  At large radius, analyses of the hot gas distribution from the X-ray observations (Pointecouteau 2005\\nocite{Pointecouteau05}; Schmidt 2007\\nocite{Schmidt07}), as well as weak lensing (Kneib et al. 2003)\\nocite{kneib03} seem to favour an NFW-like profile for galaxy clusters.  For elliptical galaxies, Wilson et al. (2001\\nocite{Wilson01}) claimed that they were consistent with isothermal profile out to about $1 \\Mpc$. Also, for massive elliptical galaxies with gravitational rings, Gavazzi et al. (2007) \\nocite{Gavazzi07} found that the WL slope of the mass density could be $\\propto r^{-2}$ out to $300 \\kpc$.  In the core of mass concentrations, the actual existence of a DM cusp predicted by simple numerical simulations is more unclear.  For spiral galaxy halos, the existence of a singular density profile is still a debate because rotation curves are better explained by isothermal profiles with a core radius (Salucci 2003\\nocite{Salucci03}). Projection effects of non circular star orbits in triaxial halos has been invoked, but only in a few cases, to explain the linear increase of the velocity at the centres of galaxies (Hayashi et al. 2006 \\nocite{Hayashi06}). As a consequence, various mechanical processes, such as gas cooling, supernova and AGN feedback, binary super massive black holes, dynamical friction of the gas outflow during AGN activities, were investigated to explain the formation of a DM core radius (Peirani et al. 2006\\nocite{Peirani06} and references therein).  For galaxy clusters, it is often argued that isothermal ellipsoid mass distributions with a flat core could better match the gravitational arcs geometry than NFW profiles (Sand et al. 2004\\nocite{Sand04}; Gavazzi 2005\\nocite{gavazzi05}; Gavazzi et al. 2003\\nocite{Gavazzi03}).  X-ray observations generally do not help much for this issue because, due to the limited spatial resolution of X-ray telescopes, one can just place upper limits on the radius of the smallest DM core, around $30-50\\kpc$ (Chen et al. 2007\\nocite{Chen07}).  In summary, the DM density profile is still an open question, and high quality data are necessary to settle these questions.  The deviation of light by masses is well described by gravitational lensing effect deduced from general relativity theory. The exquisite {\\it Hubble Space Telescope} images (particularly from the now defunct ACS camera) are providing the necessary lensing constraints to model observed gravitational lensing systems and may directly probe the existence of a ``universal'' density mass profile. However, one has to struggle to fully take into account the many observational parameters entering a lens modelling.  Firstly, it is difficult to assess the stellar mass contribution because the stellar mass-to-light ratio ($M/L$) is generally badly determined, as well as the number of sub-halos and their galaxy occupation numbers (Wright et al.  2002\\nocite{Wright02}). Secondly, lens modelling only probes the projected mass distribution of lenses along the line of sight which introduces degeneracies in the 3D density profile if the mass profile is only determined on a small range of radius. When testing parametric models of mass distributions for a given DM condensation, these current difficulties are only partly alleviated if we can probe the projected mass at many different radius.  As an example, to disentangle between a flat core or a cusp with strong lensing it is not enough to detect and to analyse gravitational images very close to the centre, the so-called demagnified central images (Gavazzi et al. 2003\\nocite{Gavazzi03}) or inner radial arcs (Mellier et al. 1993\\nocite{Mellier93}; Comerford et al. 2006\\nocite{comerford06}). Even in such ideal cases, some information on the mass distribution beyond the Einstein radius is also critically needed.  Beyond giant arc radii, one can use the information in the distortion of singly highly magnified arclets, in an intermediate shear regime ($2 < \\mu < 3$), also called flexion regime (Bacon et al. 2006\\nocite{Bacon06}; Massey et al. 2007\\nocite{Massey07}). However, a generic difficulty similar to the one encountered in the weak lensing needs to be overcome. We do not know the shape of background sources and the flexion method must be used in a statistical way. Only with the most deepest space-based observations, it becomes possible to reach surface number density of background galaxies large enough to conduct such analysis on a single cluster (see recent work on Abell~1689 by Leonard et~al.(2007)\\nocite{Leonard07}, where they reach a density of background sources equal to $\\sim 200$ sources/arcmin$^2$).  From the above discussion, we understand the difficulty to conduct an accurate measurement of the slope  of the 3D density profile at large radial distances. Hence considering  new probes of cluster mass profile is important.  In this paper, we propose a method to investigate the slope of cluster mass profiles. Gravitational image systems, i.e. ``rings'' formed around galaxy cluster members, are used to analyse the slope of the cluster's density profile. Nowadays, such rings can be systematically searched with dedicated software (e.g. Gavazzi et al. 2007, in preparation; Cabanac et al. 2005). Here, we outline the method with three rings detected around Abell~1689.  The main goal is to provide some constraints on the cluster potential at the location of the rings (i.e. at $\\sim$ 100$\\arcsec$ from the centre of the brightest cluster galaxy).  The coordinate system in this work is centred on the brightest cluster galaxy: $\\alpha_{J2000}=13:11:29.52, \\delta_{J200}=-01:20:27.59$.  The paper is organised as follows.  First we rapidly summarise the properties of gravitational rings observed in the field around elliptical lenses.  In section~\\ref{simulation}, we illustrate the method by using simulated cluster profiles, lensing galaxies, and resulting images. These simulations show that we can put constraints on the local slope of the projected mass distribution. In section~\\ref{reallife}, the method is applied to Abell~1689. In this case we show that the three rings confirm that the cluster is dominated by a bimodal mass distribution and that the local slopes of both clumps are not much steeper than isothermal. Then, the results are discussed relatively to previous lensing models of Abell~1689, including a weak lensing analysis in the field of the rings. Finally, we conclude that this method should be better used on very relaxed clusters (single halo) with regular geometry to better probe the slope of the mass profile at various distances from the cluster centre.  Throughout this paper we assume a cosmological model with $\\Omega_{\\rm{m}}$=0.3, $\\Omega_\\Lambda$=0.7, $H_0=70$km s$^{-1} \\Mpc^{-1}$. At the redshift of the cluster Abell~1689 ($z=0.185$) $1 \\arcsec$ is equivalent to 3.089\\kpc. ",
        "conclusions": "This work revisits the importance of external shear perturbations for the modelling of galaxy lenses (see Dye et al. 2007\\nocite{Dye07}). But instead of focusing our attention on the improvement of the ring modelling by introducing an external shear, we probe the gravitational potential slope in the outskirts of groups and clusters of galaxies thanks to the high sensitivity of the ring modelling both to the local external shear and convergence of the cluster.  We test the principle of the method with simple simulations using power law models to describe the cluster mass clumps.  From this simulation we found that within the astrometric accuracy of ACS images a ring can be used to detect a logarithmic slope of the external mass density distribution up to about $n=2.8$. For an NFW profile with a typical concentration parameter $C_{\\rm{vir}} \\sim 5$ (Comerford \\& Natarajan 2007\\nocite{Comerford07}), this should allow to probe the cluster potential up to a large fraction of its virial radius.  It is now possible to detect rings in the outskirts of clusters with new ground based surveys like CFHT-LS, but the determination of the direction and amplitude of the external shear depends crucially on the high resolution HST images. On the other hand, the knowledge of the lens stellar velocity dispersion and of the redshift of the ring itself would increase the robustness to the ring modelling (Koopmans et al. 2006).  Using the 3 rings found in Abell~1689 we have shown that it is possible to estimate the mass density slope.  Hence, we have shown with the information given by rings alone that the potential of Abell~1689 is bimodal. With this limited number of rings it is however not possible to better trace other DM substructures within each clump. Therefore a global modelling of the 31 arcs systems and rings remains a good challenge.  We find a mass density slope a bit larger than $n=2$ for each DM clump at about $100\\arcsec$ from their centre, but the complexity of the cluster potential really weakens the result. Since the most immediate appealing application of the method is to probe the slope departure from a $r^{-2}$ mass density profile, we strongly suggest similar analysis on less complex clusters. The ideal lenses to probe the reality of a universal NFW profile would be a cluster or a fossil group with a bright cD galaxy having a single dominant halo, and which displays at least a distant ring, a multiple arc with known redshift, and a radial arc. Such a configuration would probe the potential at various distances from the centre. The discovery of such ``golden lenses'' requests a large survey of massive clusters. Most often clusters with a great number of arcs have multi-polar potentials (i.e. longer caustics) so that are more complex to analyse. However, rings are influenced by all their nearby environment (including foregrounds). In this sense, a statistical study of the environmental effect on the rings should be conducted in the future.  Many rings can be discovered now in the field of deep wide field surveys. It will be possible to improve the method and to use it more systematically. As an example, the CFHT-LS ``wide'' survey will cover 170 degree , and we expect to detect about 10 rings per square degree with an average redshift $z_{\\rm{lens}} = 0.65$ (Cabanac et al. 2007\\nocite{Cabanac07}). The clusters and groups in front of or at the same redshift as the lensing galaxies have approximately the same sky density (Oguri 2006\\nocite{Oguri06}). Thus we can expect that several rings per square degree will be influenced by the external shear of a nearby mass condensation in the field. Such cases are already observed in the SL2S survey (Cabanac et al. 2007\\nocite{Cabanac07}).  For Abell~1689, we have found three rings within the small HST field ($r \\sim 100\\arcsec$), which may be consider as surprisingly high number. In fact, the large number of ellipticals in this cluster and the extra convergence that it adds to the lensing effect most probably compensate for the small field of view imaged with the ACS camera\\footnote{After the submission of this paper, King (2007\\nocite{king07}) has shown with cluster simulation that the ring cross section is increased by a factor about 3 near critical lines. }.  We also found two possible rings closer to the critical region of Abell~1689 ({\\it i.e.} at radii smaller than 50$\\arcsec$). However, we are not considering these rings here, because these systems are in a crowded region and can not be described by a simple contribution of a galaxy and a cluster scale component. Any attempt of modelling these latter rings must involve all the multiple arc systems of Abell~1689. In wide field surveys, perturbing clusters may be found at a lower redshift than the ring-producing elliptical lenses (e.g. Smail et al. 2007\\nocite{smailcosmiceye}). In such cases, a proper analysis requires a multi-plan ring modelling.  In conclusion, rings seem to be a promising tool to constrain the mass distribution slope at large radii from the centres of groups and clusters provided dedicated observations are carried out on a few golden lenses. They can provide information on many structural parameters of halos of clusters and galaxies. In complement to the modelling of the multiple arcs in the cluster core, they can confirm or otherwise dispute the existence of the universal DM halo profile predicted by numerical simulation, as well as to study sub-halos and their cut-off radii. For these studies, it is crucial to measure the velocity dispersion of the lens galaxies which produce the gravitational ring images. It clearly appears from the study of Abell~1689 that this method should be implemented first on a sample of very relaxed clusters or (fossil) groups in order to analyse a single DM potential with the simplest possible geometry. Only when large sample of rings is available, it will become possible to start a systematic analysis of the external shear perturbation on ring shapes in correlation with more complex nearby (eventually foreground) environment."
    },
    "0710/0710.5466_arXiv.txt": {
        "abstract": "{A variety of formation scenarios have been proposed to explain the diversity of properties observed in bulges. Studying their intrinsic shape can help to constraint the dominant mechanisms at the epochs of their assembly.} {The structural parameters of a magnitude-limited sample of 148 unbarred S0--Sb galaxies were derived in order to study the correlations between bulges and disks, as well as the probability distribution function of the intrinsic equatorial ellipticity of bulges.} {We present a new fitting algorithm (GASP2D) to perform two-dimensional photometric decomposition of the galaxy surface-brightness distribution. This was assumed to be the sum of the contribution of a bulge and disk component characterized by elliptical and concentric isophotes with constant (but possibly different) ellipticity and position angles. Bulge and disk parameters of the sample galaxies were derived from the $J-$band images, which were available in the Two Micron All Sky Survey. The probability distribution function of the equatorial ellipticity of the bulges was derived from the distribution of the observed ellipticities of bulges and misalignments between bulges and disks.} {Strong correlations between the bulge and disk parameters were found.  About $80\\%$ of bulges in unbarred lenticular and early-to-intermediate spiral galaxies are not oblate but triaxial ellipsoids. Their mean axial ratio in the equatorial plane is $\\langle B/A \\rangle = 0.85$. Their probability distribution function is not significantly dependent on morphology, light concentration or luminosity. The possible presence of nuclear bars does not influence our results.} {The interplay between bulge and disk parameters favors scenarios in which bulges have assembled from mergers and/or have grown over long times through disk secular evolution. However, all these mechanisms have to be tested against the derived distribution of bulge intrinsic ellipticities.} ",
        "introduction": "\\label{sec:introduction}  The relative prominence of galactic bulges with respect to their disks is important in the definition of galaxy types. Therefore, understanding the formation of bulges is key to understanding the origin of the Hubble sequence.  Bulges are diverse and heterogeneous \\citep[see the reviews by][]{kormendy93,wyse97,kormendy04}. The big bulges of lenticulars and early-type spirals are similar to low-luminosity elliptical galaxies. Their surface-brightness radial profiles generally follow a De Vaucouleurs law \\citep[][hereafter MH01]{andredakis95,carollo98,mollenhoff01}. The majority of these bulges appear rounder than their associated disks \\citep{kent85}. Their kinematical properties are well described by dynamical models of rotationally flattened oblate spheroids with little or no anisotropy \\citep{kormendy82,davies83,cappellari06}. They have photometrical and kinematical properties, which satisfy the fundamental plane (FP) correlation \\citep{bender92,bender93,burstein97,aguerri05a}. On the contrary, the small bulges of late-type spiral galaxies seems to be reminiscent of disks. Their surface-brightness radial profiles have an almost exponential falloff \\citep{andredakis94,dejong96,macarthur03}. In some cases they have apparent flattenings that are similar or even larger than their associated disks \\citep{fathi03} and rotate as fast as disks \\citep{kormendy93,kormendy02}. Late-type bulges deviate from the FP \\citep{carollo99}.  Different formation mechanisms (or at least a variety of dominant mechanisms at the epochs of star formation and mass assembly) were proposed to explain the variety of properties observed in bulges. Some of these formation processes are rapid. They include early formation from the dissipative collapse of protogalactic gas clouds \\citep{eggen62,sandage90,gilmore98,merlinchiosi06} or later assembly from mergers between pre-existing disks \\citep{kauffmann96,baugh96,cole00}. In both scenarios the disk forms after the bulge as a consequence of either a long star-formation time compared to the collapse time or a re-accretion around the newly formed bulge.  Bulges can also grow over long timescales through the disk secular evolution driven by bars and/or environmental effects. Bars are present in more than half of disk galaxies in the local universe \\citep{eskridge00,menendez-delmestre07} and out to $z\\sim1$ \\citep{elmegreen04,jogee04}. They are efficient mechanisms for driving gas inward to the galactic center and feed the galactic supermassive black hole \\citep[see][and references therein]{corsini03a}. In addition, bar dissolution due to the growth of a central mass \\citep{pfenninger90}, scattering of disk stars at vertical resonances \\citep{combes90}, and coherent bending of the bar perpendicular to the disk plane \\citep{raha91,debattista04,athanassoula05,martinez-valpuesta06} are efficient mechanisms in building central bulge-like structures, the so-called boxy/peanut bulges. Moreover, the growth of the bulge out of disk material may also be externally triggered by satellite accretion during minor merging events \\citep{searlezinn78,aguerri01,eliche-moral06} and gas infall \\citep{thakarryden98}.  Traditionally, the study of the relations between the structural parameters of the galaxies have been used to understand the bulge formation processes, e.g., the correlation between the bulge effective radius and the scale length of the disk in many galaxy samples has always been interpreted as an indication that bulges were formed by secular evolution of their disks \\citep[see][]{macarthur03}. However, one piece lost in this study is the three-dimensional shape of the bulges. By studying this, one might be able to provide the relative importance of rapid and slow processes in assembling the dense central components of disk galaxies. A statistical study can provide a crucial piece of information for testing the results of numerical simulations of bulge formation for different galaxy type along the morphological sequence.  In this paper, we analyze a sample of unbarred early-type disk galaxies to derive the intrinsic ellipticity of their bulges in the galactic plane. The twisting of bulge isophotes \\citep{lindblad56,zaritsky86} and misalignment between the major axes of the bulge and disk \\citep{bertola91} are not possible if the bulge and disk are both oblate.  Therefore, they were interpreted as a signature of bulge triaxiality.  This idea is supported by the presence of non-circular gas motions \\citep[e.g.,][]{gerhardvietri86,bertola89,gerhard89,berman01} and a velocity gradient along the galaxy minor axis \\citep{corsini03b,coccato04,coccato05}. We improve the previous works in several aspects. First, we use near-infrared images to map the distribution of the mass-carrying evolved stars and avoid contamination of dust and bright young stars. Second, we retrieve the structural parameters of the bulge and disk by applying a new algorithm for two-dimensional photometric decomposition of the observed surface-brightness distribution.  Finally, we obtain the probability distribution function (PDF) of the intrinsic equatorial ellipticity of bulges by using a new mathematical treatment of the equations describing their three-dimensional shape.  The paper is organized as follow. The selection criteria of our sample galaxies and the analysis of their near-infrared images are described in Sect. \\ref{sec:sample}. Our new photometric decomposition method for deriving the structural parameters of the bulge and disk by analyzing the two-dimensional surface brightness distribution of galaxies is presented in Sect. \\ref{sec:decomposition}. The correlations between the structural parameters of the sample galaxies are discussed in Sect. \\ref{sec:correlations}. The PDF of intrinsic equatorial ellipticity of the studied bulges is derived in Sect. \\ref{sec:ellipticity}. Our conclusions and a summary of the results are given in Sect. \\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  The structural parameters of the bulge and disk of a magnitude-limited sample of 148 unbarred S0--Sb galaxies were investigated to constrain the dominant mechanism at the epoch of bulge assembly.  \\begin{itemize}  \\item We presented a new fitting algorithm (GASP2D) to perform two-dimensional photometric decomposition of galaxy images. The surface-brightness distribution of the galaxy was assumed to be the sum of the contribution of a S\\'ersic bulge and an exponential disk. The two components were characterized by elliptical and concentric isophotes with constant (but possibly different) ellipticity and position angles. GASP2D is optimized to deal with large image samples, and it adopts a robust Levenberg-Marquard fitting algorithm in order to obtain reliable estimates of the galaxy structural parameters.  \\item The bulge and disk parameters of the sample galaxies were derived from the $J-$band images, which were available in the Two Micron All Sky Survey.  \\item The bulges of the sample galaxies follow the same FP, FJ, and PP relationships found for elliptical galaxies. No statistically significant difference is observed when only bulges of lenticular and early-to-intermediate spiral galaxies were considered. This supports the idea that bulges and ellipticals formed in the same way.  \\item Tight correlations between the parameters of bulges and disks were found. In fact, the disk scale lengths increase with both the central velocity dispersion and bulge effective radius. Therefore, larger disks reside in galaxies with more massive and larger bulges. This was interpreted as an indication of the formation of bulges via secular evolution of their host disks.  \\item Our measurements of the exponential scale length of the bulge and disk, as well as of bulge shape parameter, were also fully consistent with numerical simulations of the effects of mergers on the mass distribution of the bulge and disk in galaxies formed in hierarchical clustering scenarios.  \\item These results indicate that the above relations are not enough to clearly distinguish between bulges formed by early dissipative collapse, merging or secular evolution. All these mechanisms could be tested against the intrinsic shape of bulges. Therefore, the PDF of the intrinsic equatorial ellipticity of the bulges was derived from the distribution of the observed ellipticities of bulges and their misalignments with disks.  \\item About $80\\%$ of bulges in unbarred lenticular and early-to-intermediate spiral galaxies are not oblate but triaxial ellipsoids. Their mean axial ratio in the equatorial plane is $\\langle B/A \\rangle = 0.85$. This is consistent with previous findings by Bertola et al. (1991) and Fathi \\& Peletier (2003). There is no significant dependence of the PDF on the morphology, light concentration, and luminosity of bulges. The derived PDF is independent of the possible presence of nuclear bars.  \\end{itemize}"
    },
    "0710/0710.0844_arXiv.txt": {
        "abstract": "We present a complete set of diagnostic tools aimed at reproducing synthetic non-thermal (synchrotron and/or Inverse Compton, IC) emissivity, integrated flux energy, polarization and spectral index simulated maps in comparison to observations. The time dependent relativistic magnetohydrodynamic (RMHD) equations are solved with a shock capturing code together with the evolution of the maximum particles energy. Applications to Pulsar Wind Nebulae (PWNe) are shown. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.2888_arXiv.txt": {
        "abstract": "Two-temperature, two-mass quasi-equilibrium plasmas may occur in electron-ion plasmas, nuclear-matter, as well as in electron-hole condensed-matter systems. Dense two-temperature hydrogen plasmas straddle the difficult partially-degenerate regime of electron densities and temperatures which are important in astrophysics,  in inertial-confinement fusion research, and other areas of warm dense matter physics. Results from Kohn-Sham calculations and QMC are used to benchmark the procedures used in classical molecular-dynamics simulations, HNC and CHNC methods to derive electron-electron and electron-proton pair-distribution functions. Then, nonequilibrium molecular dynamics for two-temperature, two-mass plasmas are used to obtain the pair distribution. Using these results, the correct HNC and CHNC procedures for the evaluation of pair-distribution functions in two-temperature two-mass two-component charged fluids are established. Results for a mass ratio of 1:5, typical of electron-hole fluids, as well as for compressed hydrogen are presented. ",
        "introduction": "The study of hot strongly-coupled dense charged fluids is a difficult task, especially near the regime of or molecular and atomic species\\cite{ceperleyH}, or excitons in electron-hole plasmas. The system is better understood for fully ionized systems, such as hydrogen, in the form of free electrons and protons, and fully-ionized electron-hole condensates. In fact, considerable headway has been made using methods based on density-functional theory (DFT), even for plasmas with multiple states of ionization. DFT methods have been used with molecular-dynamics based approaches\\cite{kwon,dejarlais,mazevet}, as well as within multi-component integral-equation approaches \\cite{pdw95,ilciacco}. Equilibrium properties of plasmas, as well as their linear transport properties, have been successfully studied in these papers, and excellent agreement between the molecular-dynamics based DFT and integral-equation based DFT has been found \\cite{lvm}.  On the other hand, laser-produced plasmas are initially formed as two-temperature plasmas, where the electrons have absorbed the laser energy and have self-equilibrated to some ``electron temperature'' $T_e$, while the ions remain cool, at some temperature $T_i$, with $T_i<T_e$. The opposite situation arises in shock-wave generated plasmas, where the ions absorb the shock energy and $T_i>T_e$. Such two-temperature plasmas also occur in astrophysical settings, affecting the time of termination of synthesis of light-nuclei to occur at different stages of cooling of the electrons\\cite{astrop}, and influencing the Coulomb nuclear-tunneling rates\\cite{dewitt}. The possibility of such well-defined two-temperature plasmas is largely a result of the extreme mass ratio $m_i/m_e\\ge 1836$ between ions and electrons. Similar, but less well defined situations can arise in electron-hole plasmas, where the masses are of the same order of magnitude (e.g, the electron and hole masses in GaAs are 0.067$m_e$ and 0.34$m_e$ respectively, with an electron-hole mass ratio of $\\sim 5$). GaAs is a direct bandgap material, and electron-hole plasmas are more easily studied in indirect-gap systems like Si where the density-of states mass ratio is $\\sim 3$. The simulation of such systems at two temperatures, using quantum Monte-Carlo methods is at present unavailable, even in regimes of densities and temperatures where bound states (or exciton formation in electron-hole systems) do not exist. Thus it is natural to look for analytical methods based on integral-equation techniques which are computationally simple and physically insightful. However, although $T_e,\\, T_i$ define the temperatures of each subsystem and the pair-distribution functions (PDFs) $g_{ee}$ and $g_{ii}$, the  `temperature' $T_{ei}$ entering into the cross-correlations $g_{ei}$,  as well as the the effects of electron spin, exchange etc., relevant to two-temperature systems need to be clarified. In this context we use $T_{ee}=T_e$, $T_{ii}=T_i$, and $T_{ei}$ to refer to the electron-, ion-, and electron-ion temperatures as they enter independently into the Ornstein-Zernike (OZ) and hypernetted chain (HNC) equations. In fact, some authors\\cite{seuf} have proposed to modify the well-established OZ equations in dealing with two-temperature (2T) two-mass (2M) systems.  The objective of this paper is to study such 2T-2M  plasmas using results from molecular-dynamics (MD)\\cite{mc}, HNC\\cite{hncref}, classical-map HNC (CHNC)\\cite{prl1,prb3d}, quantum Monte Carlo (QMC)\\cite{mc} and Kohn-Sham (KS)\\cite{ilciacco} methods to establish the proper implementation of quantum  effects and 2T, 2M situations in simulation studies. One of our main interests would be uniform hydrogenic plasmas free of bound states, in the regime of warm-dense matter. ",
        "conclusions": "The simplest classical rendering of quantum plasmas, based on the use of diffraction corrected potentials (Eq.~\\ref{zerothset}) was used with HNC calculations and MD simulations to resolve the ambiguities and difficulties in handling the two-temperature, two-mass system. We conclude that the modifications to the OZ equations proposed by Seuferling et al.\\cite{seuf}., are not needed. The classical mapping of quantum systems to the HNC equations, as used in the CHNC was confirmed by comparisons with Kohn-Sham DFT calculations as well as with available PIMC results for compressed hydrogen plasmas at finite temperatures. We conclude that the HNC and CHNC, together with the standard OZ equations provide excellent, accurate and simple analytical tools for the investigation of many-particle quasi-equilibrium systems for which direct quantum simulations continue to remain too prohibitive or unfeasible."
    },
    "0710/0710.3076_arXiv.txt": {
        "abstract": "Ca II triplet spectroscopy has been used to derive stellar metallicities for individual stars in four LMC fields situated at galactocentric distances of 3\\arcdeg, 5\\arcdeg, 6\\arcdeg\\@ and 8\\arcdeg\\@ to the north of the Bar. Observed metallicity distributions show a well defined peak, with a tail toward low metallicities. The mean metallicity remains constant until 6\\arcdeg\\@ ([Fe/H]$\\sim$-0.5 dex), while for the outermost field, at 8\\arcdeg, the mean metallicity is substantially lower than in the rest of the disk ([Fe/H]$\\sim$-0.8 dex). The combination of spectroscopy with deep CCD photometry has allowed us to break the RGB age--metallicity degeneracy and compute the ages for the objects observed spectroscopically. The obtained age--metallicity relationships for our four fields are statistically indistinguishable. We conclude that the lower mean metallicity in the outermost field is a consequence of it having a lower fraction of intermediate-age stars, which are more metal-rich than the older stars. The disk age--metallicity relationship is similar to that for clusters. However, the lack of objects with ages between 3 and 10 Gyr is not observed in the field population. Finally, we used data from the literature to derive consistently the age--metallicity relationship of the bar. Simple chemical evolution models have been used to reproduce the observed age--metallicity relationships with the purpose of investigating which mechanism has participated in the evolution of the disk and bar. We find that while the disk age--metallicity relationship is well reproduced by close-box models or models with a small degree of outflow, that of the bar is only reproduced by models with combination of infall and outflow. ",
        "introduction": "Despite decades of work, there are still significant gaps in our knowledge of the LMC's star formation and chemical enrichment histories \\citep{ol96}. This is motivated in part by the vastness of its stellar populations and to our limitations in observing sizable samples of stars. The age distribution of star clusters is relatively well known \\citep[e.g.][]{geisler97}: there is an age interval between 10 and 3 Gyr  with almost no clusters. This so-called age-gap may also correspond to an abundance gap \\citep{ol91}, since the old clusters are metal-poor while the young ones are relatively metal-rich. The star formation history (SFH) of field stars is much less precisely known. Studies using {\\itshape HST} data in small fields, suggest that the SFH of the LMC disk has been more or less continuous, with some increase in the star formation rate (SFR) in the last few Gyr \\citep[e.g.][]{Smecker-Hane02,castro01,holtzman99}. This contrasts with previous results (based on much shallower data), which found a relatively young age (a few Gyr) for the dominant LMC population \\citep{hardy84,bertelli92,westerlund95,vallenari96}. The SFH of the LMC can now be studied using sufficiently deep ground-based data (e.g.\\ reaching the oldest main-sequence turn-off with good photometric precision) in large areas and in different positions of the galaxy \\citep[e.g.][hereafter Paper I]{gshpz04,gshpz05}.  Less is known about the LMC chemical enrichment history. The age--metallicity relation (AMR) normally used for the LMC is defined by star clusters \\citep[e.g.][]{ol91,geisler97,bica98,dirsch00}. All works based on clusters obtain similar results.  The mean metallicity jumps from [Fe/H] $\\sim -1.5$ for the oldest clusters to [Fe/H] $\\sim -0.5$ for the youngest ones. There were no studies on the field population AMR until the last decade. \\citet{dopita97}, from a study of $\\alpha$-elements in planetary nebulae, obtained a result qualitatively similar to that  found in the clusters although only ten objects were used. Later, \\citet{bica98}, using Washington photometry, and \\citet{c00} and \\citet{dirsch00}, using Str\\\"ongrem photometry, obtained the metallicity of RGB stars in different positions of the LMC. \\citet{c00} also observed 20 stars with the infrared CaII triplet (CaT). All of them found that the age-gap observed in the clusters is not found in the field population. More recently, \\citet{c05} have obtained stellar metallicities for almost 400 stars in the bar of the LMC, also using CaT lines. The metallicity for each star has been combined with its position in the color--magnitude diagram (CMD) to estimate its age. The obtained AMR shows a similar behavior to that of the clusters for the oldest population. However, while for the clusters the metallicity has increased over the last 2 Gyr, this has not happened in the bar. Finally \\citet{gratton04} measured the metallicity of about 100 bar RR Lyrae, obtaining an average metallicity of [Fe/H] $=-1.48$.  Another point that it is necessary to investigate is the presence, or not, of an abundance gradient in the LMC. From observations of clusters, \\citet{ol91} and \\citet{santos99} found no evidence for the presence of a radial metallicity gradient in the LMC. The only evidence for a radial metallicity gradient in the LMC cluster system was reported by \\citet{kon93} based on six outer LMC clusters ($\\geq$ 8 kpc). \\citet{hill95} were the first to report evidences of a gradient in the field population from high-resolution spectroscopy, using a sample of nine stars in the bar and in the disk. They found that the bar is on average 0.3 dex more metal-rich than the disk population, but the disk stars studied are located within a radius of 2\\arcdeg. Subsequent work by \\citet{cioni03}, detected that the C/M ratio between number of asymptotic giant branch stars of spectral types C and M increased when moving away from the bar and within a radius  of 6$\\fdg$7. As the C/M ratio is anticorrelated with metallicity, this increment implies a decrease in metallicity. Finally, an outward radial gradient of decreasing metallicity was also found by \\citet{alves04} from infrared CMD using the 2MASS survey.  We have obtained deep photometry with the Mosaic II CCD Imager on the CTIO 4m telescope in four disk fields at different distances from the center of the LMC (RA=5$^h$23$^m$34$\\fs$5, $\\delta$=-69$\\arcdeg$45'22``) with a quality similar to that obtained by {\\itshape HST} in more crowded areas (see Paper I). The position of these fields, together with a description of their CMDs, are presented in Paper I, and detailed SFHs will be published in forthcoming papers. In the present investigation we focus on obtaining stellar metallicities for a significant number of individual RGB stars in these four fields using spectra obtained with the HYDRA spectrograph at the CTIO 4m telescope. In Section \\ref{targetselection} we present our target selection. The observations and data reduction are presented in Section \\ref{datareduction}. The radial velocities of the stars in our sample are obtained in Section \\ref{radialvelocities}. In Section \\ref{cat} we discuss the calculation of the CaT equivalent widths and the determination of metallicities. Section \\ref{agedetermination} presents the method used to derive ages for each star by combining the information on their metallicity and position on the CMD. The analysis of the data is presented in Section \\ref{analysis}, where the possible presence of a kinematically hot halo is discussed and the AMR and the SFH in each of our fields are described and compared. In the final section, the AMRs are compared to theoretical models and conclusions are drawn about the chemical evolution of the LMC. ",
        "conclusions": "\\label{conclusions}  Using infrared spectra in the CaT region, we have obtained metallicities and radial velocities for a sample of stars in four LMC fields. Metallicities have been calculated using the relationships between the equivalent width of the CaT lines, $\\Sigma Ca$, and metallicity derived in Paper II. In addition, we have estimated the age of each star using a relationship derived from a synthetic CMD which, from the color, magnitude and metallicity of a star, allow us to estimate its age. The main results of this paper are: \\begin{itemize} \\item The velocity distribution observed in each field agrees with the rotational thick disk kinematics of the LMC. The velocity dispersion is slightly larger for the most metal-poor stars. However, the  values obtained do not indicate the presence of a kinematically hot halo. \\item The metallicity distribution of each field has a well defined peak with a tail toward low metallicities. The mean metallicity is constant until the field at 6\\arcdeg\\@ ([Fe/H]$\\sim$-0.5 dex), and  is a factor two more metal-poor for the outermost field ([Fe/H]$\\sim$-0.8 dex). \\item The AMR observed in each disk field is compatible with a single global relationship for the disk. We conclude that the outermost field is more metal-poor on average because it contains a lower fraction of relatively young stars (age$\\leq$5 Gyr), which are also more metal-rich. \\item The disk AMR shows a prompt initial chemical enrichment. Subsequently, the metallicity increased very slowly until about 3 Gyr ago, when the rate of metal enrichment increased again. This AMR is similar to that of the cluster system, except for the lack of clusters with ages between 3 and 10 Gyr. The recent fast enrichment observed in the disk and  in the cluster system is not observed in the bar. \\item The $\\psi(t)$ of the three innermost fields show a first episode of star formation until about 10 Gyr ago, followed by a period with a low SFR until $\\sim$5 Gyr ago, when the SFR increases, and reaches its highest values $\\sim$2-3 Gyr ago. The outermost field does not show the recent increase of SFR. The second main episode is also observed, and is more prominent in the bar, where an increased SFR at old ages is not observed. The lower SFR between 5 and 10 Gyr ago is probably related to the age-gap observed in the clusters. \\item Under the assumption of a solar yield, the disk AMR is well reproduced either by a chemical evolution model with outflow with $\\lambda$ between 1 and 2, so the disk losses the same amount of gas that has taken part in the star formation, or by models combining infall and outflow with $\\alpha$=0.2 and $\\lambda$=0.05, which means that the galaxy is almost a closed-box system. With a smaller yield, the AMR could also be reproduced with a closed-box model. The bar AMR is well reproduced by models with a combination of inflow and outflow with $\\alpha$=1.2 and $\\lambda$ between 0.4 and 0.5. This suggests that the amount of infalling gas was larger than the amount that participated in the star formation in the bar, and also that the amount of gas that escaped the bar was 50\\% of the total that participated in star formation. \\end{itemize}"
    },
    "0710/0710.1073_arXiv.txt": {
        "abstract": "We determine the ratio of helium- to hydrogen-atmosphere white dwarf stars as a function of effective temperature from a model atmosphere analysis of the infrared photometric data from the Two Micron All Sky Survey combined with available visual magnitudes. Our study surpasses any previous analysis of this kind both in terms of the accuracy of the $\\Te$ determinations as well as the size of the sample. We observe that the ratio of helium- to hydrogen-atmosphere white dwarfs increases gradually from a constant value of $\\sim0.25$ between $\\Te=15,000$~K and 10,000~K to a value twice as large in the range $10,000 > \\Te > 8000$ K, suggesting that convective mixing, which occurs when the bottom of the hydrogen convection zone reaches the underlying convective helium envelope, is responsible for this gradual transition. The comparison of our results with an approximate model used to describe the outcome of this convective mixing process implies hydrogen mass layers in the range $M_{\\rm H}/M_{\\rm tot}=10^{-10}$ to $10^{-8}$ for about 15\\% of the DA stars that survived the DA to DB transition near $\\Te\\sim 30,000$~K, the remainder having presumably more massive layers above $M_{\\rm H}/M_{\\rm tot}\\sim10^{-6}$. ",
        "introduction": "The Two Micron All Sky Survey (2MASS) represents one the largest set of homogeneous infrared photometric data for all areas of the sky applicable to white dwarf stars of various spectral types. Indeed, the 2MASS survey contains $JHK_S$ magnitudes for almost all point sources of the sky up to the magnitude thresholds of $J \\sim 16.8$, $H \\sim 16.5$, and $K_S\\sim 16.2$ for detections with any formal uncertainty \\citep[see for instance Fig.~1 of][]{tremblay07}. This yields a nearly complete magnitude limited survey of white dwarfs that allows us to trace a portrait of the local population of white dwarfs, assuming we can properly identify the white dwarfs in this survey. To that effect, the most complete compilation is that of the Villanova White Dwarf Catalog\\footnote{http://www.astronomy.villanova.edu/WDCatalog/index.html} (WDC), which contains more than 5500 spectroscopically identified white dwarfs reported in the literature. The interest of such a sizeable sample of homogeneous infrared photometric observations is to reach an unprecedented level in the statistical analysis of white dwarf stars. In addition, \\citet{tremblay07} have demonstrated the overall reliability of the 2MASS Point-Source Catalog (PSC) photometry for the purpose of detailed comparisons with the predictions from white dwarf model atmospheres.  Studies of the spectral evolution of white dwarfs attempt to understand how the surface composition and the corresponding spectral signature evolve along the white dwarf cooling sequence (see \\citealt{fontaine87,fontaine91} for a review). Because of the intense gravitational field present in these stars, hydrogen will float to the surface in the absence of competing mechanisms, while heavier elements will sink below the photospheric regions.  Indeed, the majority of white dwarfs have optical spectra that are completely dominated by strong hydrogen Balmer lines --- the DA stars --- with hydrogen-dominated atmospheres. We also find that almost 25\\% of white dwarfs have helium-dominated atmospheres, mostly DO, DB, DQ, and DZ stars, and some of the DC stars as well. The fact that not all white dwarfs have hydrogen-rich atmospheres imply that either some stars are born with hydrogen-deficient atmospheres and remain that way throughout their evolution, or alternatively, that several physical mechanisms such as diffusion, accretion, convection, radiation pressure, and stellar winds are competing with gravitational settling in determining their surface composition as they cool off. This particular issue can be resolved by analyzing the distribution of spectral types as a function of effective temperature (or any temperature index) for a large sample of white dwarf stars, in a way similar to the analysis of \\citet{sion84} nearly twenty five years ago.  Before discussing the physical processes that will affect the spectral evolution of hydrogen- and helium-atmosphere white dwarfs, it is important first to consider the constraints brought about by the studies of the hottest white dwarfs. The canonical mass fractions of light elements in white dwarfs considering the mass threshold for residual nuclear burning are $M_{\\rm He}/M_{\\rm tot} \\sim 10^{-2}$ and $M_{\\rm H}/M_{\\rm tot} \\sim 10^{-4}$. However, multiple excursions on the AGB nuclear burning phase, during which very late helium flashes could remove essentially all of the hydrogen \\citep{werner06}, could also produce white dwarfs with total hydrogen masses much lower than the canonical value of $10^{-4}$. The first class would correspond to the progenitors of the hottest DA stars while the second class would represent the progenitors of hot white dwarf stars with hydrogen-deficient atmospheres --- the PG 1159 stars at very high effective temperatures, and the DO stars at slightly lower temperatures, whose spectra are dominated by He\\textsc{ii} lines. It has traditionally been believed that by the time DO stars cool off to $\\Te\\sim45,000$~K, residual amounts of hydrogen ($M_{\\rm H}/M_{\\rm tot} \\sim 10^{-16}$) thoroughly mixed in their helium-rich envelope would gradually accumulate to the surface, turning all helium-atmosphere white dwarfs into DA stars \\citep{fontaine87,fontaine91}. This scenario would account for the existence of the so-called DB gap, a range of temperature between $\\Te=45,000$~K and 30,000~K where all white dwarfs have a DA spectral type. However, the recent discovery in the Sloan Digital Sky Survey of several DB stars within the gap \\citep{eisenstein06a} implies that the total mass of hydrogen left in the envelope of DO stars can be even smaller than previously believed. It is then generally admitted that the significant increase in the number of helium-atmosphere DB white dwarfs below $\\Te\\sim30,000$~K can be explained in terms of the convective dilution of the superficial hydrogen atmosphere by the underlying helium convective envelope, provided that the hydrogen layer is sufficiently thin \\citep[$M_{\\rm H}/M_{\\rm tot} \\sim 10^{-15}$,][]{macdonald91}.  At even lower effective temperatures ($\\Te\\lesssim12,000$~K), hydrogen-atmosphere white dwarfs get a second opportunity to evolve into stars with helium-dominated atmospheres when the superficial hydrogen layer becomes convective over a significant fraction of its depth. This process is illustrated in Figure \\ref{fg:f1} where the extent of the hydrogen convection zone is displayed as a function of decreasing effective temperature for a 0.6 \\msun\\ DA white dwarf, based on evolutionary models with thick hydrogen layers similar to those described by \\citet{fon01} and kindly provided to us by G.~Fontaine and P.~Brassard. These calculations show that if the hydrogen envelope is thin enough, the bottom of the hydrogen convection zone may eventually reach the underlying and more massive convective helium layer, resulting in a mixing of the hydrogen and helium layers \\citep{strittmatter71,shipman72,baglin73,koester76,vauclair77}. Figure \\ref{fg:f1} also indicates that the effective temperature at which this mixing occurs will depend on the thickness of the hydrogen envelope. The thicker the envelope, the lower the mixing temperature; if the hydrogen layer is more massive than $M_{\\rm H}/M_{\\rm tot} \\sim 10^{-6}$, mixing will never occur.  The simplest physical model that can be used to describe this convective mixing process is to assume that hydrogen and helium are homogeneously mixed. Since the helium convection zone is much more massive ($M_{\\rm He-conv}/M_{\\rm tot}\\sim10^{-6}$) than the hydrogen layer when mixing occurs, it is generally assumed that a DA star would be transformed into a helium-atmosphere white dwarf of type DB, DQ, DZ, or DC with only a trace abundance of hydrogen. More detailed calculations discussed by \\cite{fontaine91} confirm these predictions. If such a process takes place in cool white dwarfs, we then expect the ratio of helium- to hydrogen-atmosphere white dwarfs to increase at low effective temperatures.  A comparison between the observed ratio as a function of $\\Te$ and the theoretical expectations would make it possible to estimate the thickness of the hydrogen layers in DA white dwarfs. In turn, these determinations could be compared to independent measurements of the hydrogen layer mass inferred from ZZ Ceti asteroseismology.  The best approach to study the occurrence of this convective mixing process begins with the statistical analysis of large samples of cool white dwarfs for which we can determine the main atmospheric constituent and the effective temperature, or any analogue temperature index such as the absolute $V$ magnitude. \\citet{sion84} was the first to study this problem by estimating the non-DA to DA ratio at low effective temperatures using a proper motion sample of 695 white dwarfs drawn from the catalog of spectroscopically identified white dwarfs available at that time. The results from Figure 1 of \\cite{sion84} are reproduced here in Figure \\ref{fg:f2} in terms of the non-DA to DA ratio (rather than the absolute number of DA and non-DA stars) as a function of the absolute visual magnitude $M_V$. Looking at the results, we notice a first increase in this ratio at $M_V>11.25$, which corresponds to $\\Te\\sim15,000$ K for a 0.6 \\msun\\ white dwarf, a temperature much in excess of the theoretically predicted value for the convective mixing process.  A second increase occurs at $M_V\\sim 12.5$, or $\\Te \\sim 10000$ K, but the observed ratio drops suddenly afterwards and this increase is probably not significant. Then a third increase occurs at $M_V> 13$, or $\\Te<8000$~K.  Clearly, the evidence based on these results are not exactly convincing, and our current picture of the situation below 12,000 K is at best sketchy.  The global portrait did not change significantly since the analysis of \\citet{sion84}, and this remains the only available proof quoted in the literature of the transformation of some DA stars into non-DA stars (see also \\citealt{greenstein86}).  Our goal is to improve upon the analysis of \\citet{sion84} by using the 2MASS photometric sample combined with the WDC database to identify white dwarfs and determine their atmospheric composition, and most importantly to obtain more accurate temperature determinations using a full model atmosphere analysis of the $VJHK_S$ photometry. We first describe in \\S~2 the 2MASS photometric white dwarf sample used in our analysis. The determination of the atmospheric parameters is then presented in \\S~3 and the uncertainties related to our approach are discussed at length in \\S~4. The evolution of the ratio of helium- to hydrogen-atmosphere white dwarfs as a function of effective temperature is then determined in \\S~5. Our conclusions follow in \\S~6. ",
        "conclusions": "We have presented a detailed study of the spectral evolution of cool white dwarfs following the pioneering effort of \\citet{sion84}. In particular, we have determined the ratio of helium- to hydrogen-atmosphere white dwarfs below $\\Te=15,000$~K as a function of effective temperature by increasing the size of the sample by a factor of 2, and by obtaining more precise measurements of the effective temperature of these stars through model atmosphere fits to Johnson $V$ (or Str\\\"omgren $y$) and 2MASS $JHK_S$ photometric observations. We have confirmed that in this temperature range, the evolution of hydrogen- and helium-atmosphere white dwarfs is not independent, and we have revealed in greater detail the role of the convective mixing process responsible for the coupling between these two white dwarf types. We have found that about 15$\\%$ of cool hydrogen-atmosphere DA white dwarfs between $\\Te=15,000$~K and 10,000~K are transformed into helium-atmosphere non-DA white dwarfs at lower temperatures. Using a basic model for convective mixing, our results imply in turn that 15\\% of the DA stars that have survived the DA to DB transition have hydrogen layer masses in the range $M_{\\rm H}/M_{\\rm tot}\\sim10^{-10}$ to $10^{-8}$, although the exact values depend on the assumed convective efficiency in the evolutionary model calculations. The remaining DA stars would presumably have more massive hydrogen layers, probably in excess of $M_{\\rm H}/M_{\\rm tot}\\sim10^{-6}$.  Further analysis of the convective mixing problem would require a significantly larger photometric sample, such as the {\\it ugriz} photometric sample of the Sloan Digital Sky Survey (SDSS). This sample covers a few areas of the sky with a completeness of around 95\\% up to $g=20$. However, most spectroscopic follow-ups have relied on color-color diagrams to identify white dwarfs \\citep{kl04,eisenstein06b}, and this method is strongly biased towards hot white dwarfs ($\\Te > 12000$~K). Therefore, the SDSS white dwarfs identified so far cannot be used to study the convective mixing process at low effective temperatures. One way to proceed to identify cooler white dwarfs in the SDSS is through proper motion-color diagrams. \\citet{harris06} have used this method to identify nearly 6000 white dwarfs but a complete spectroscopic analysis similar to that of \\citet{kilic06} has yet to be performed to take full advantage of this sample."
    },
    "0710/0710.1135_arXiv.txt": {
        "abstract": "The Cygnus Loop was observed from the northeast to the southwest with XMM-Newton.  We divided the observed region into two parts, the north path and the south path, and studied the X-ray spectra along two paths. The spectra can be well fitted either by a one-component non-equilibrium ionization (NEI) model or by a two-component NEI model. The rim regions can be well fitted by a one-component model with relatively low \\kTe~whose metal abundances are sub-solar (0.1--0.2). The major part of the paths requires a two-component model.  Due to projection effects, we concluded that the low \\kTe~($\\sim$0.2\\,keV) component surrounds the high \\kTe~($\\sim$0.6\\,keV) component, with the latter having relatively high metal abundances ($\\sim$5 times solar). Since the Cygnus Loop is thought to originate in a cavity explosion, the low \\kTe~component originates from the cavity wall while the high \\kTe~component originates from the ejecta.  The flux of the cavity wall component shows a large variation along our path. We found it to be very thin in the south-west region, suggesting a blowout along our line of sight.  The metal distribution inside the ejecta shows non-uniformity, depending on the element.  O, Ne and Mg are relatively more abundant in the outer region while Si, S and Fe are concentrated in the inner region, with all metals showing strong asymmetry. This observational evidence implies an asymmetric explosion of the progenitor star. The abundance of the ejecta also indicates the progenitor star to be about 15\\,\\Msun. ",
        "introduction": "A supernova remnant (SNR) reflects the abundance of the progenitor star when the remnant is young and that of the interstellar matter (ISM) when it becomes old.  In this way, we can study the evolution of the ejecta and the ISM.  The Cygnus Loop is a proto-typical middle-aged shell-like SNR. The angular diameter is about 2$^{\\circ}$.4 and it is very close to us (540\\,pc; Blair et al.\\ 2005), implying a diameter of $\\sim$23\\,pc. The estimated age is about 10000\\,years, less than half that based on the previous distance estimate of 770\\,pc \\cite{minkowski58}.  Since the Cygnus Loop is an evolved SNR, the bright shell mainly consists of a shock-heated surrounding material.  Its supernova (SN) explosion is generally considered to have occurred in a preexisting cavity \\cite{mccray79}.  Levenson et al. (1997) found that the Cygnus Loop was a result of a cavity explosion that was created by a star no later than B0.  It is almost circular in shape with a break-out in the south where the hot plasma extends out of the circular shape.  Miyata et al. (1994) observed the northeast (NE) shell of the Loop with ASCA and revealed the metal deficiency there \\cite{miyata94}.  Since Dopita et al. (1977) reported the metal deficiency of the ISM around the Cygnus Loop, they concluded that the plasma in the NE-shell is dominated by the ISM.  Due to the constraints of the detector efficiency, they assumed that the relative abundances of C, N and O are equal to those of the solar value \\cite{anders89}.  More recently, Miyata et al. (2007) used the Suzaku satellite \\cite{mitsuda07} to observe one pointing position in the NE rim.  They detected emission lines from C and N and determined the relative abundances \\cite{miyata07}.  They concluded that the relative abundances of C, N and O are consistent with those of the solar values whereas the absolute abundances show depletion from the solar values \\cite{anders89}.  Katsuda et al. (2007) observed four pointings in the NE rim and detected a region where the relative abundances of C and N are a few times higher than that of O.  Hatsukade \\& Tsunemi (1990) detected a hot plasma inside the Cygnus Loop that is not expected in the simple Sedov model \\cite{hatsukade90}.  They reported that the hot plasma was confined inside the Loop.  Miyata et al. (1998) detected strong emission lines from Si, S and Fe-L from inside the Loop \\cite{miyata98}.  They found that the metal abundance is at least several times higher than that of the solar value \\cite{anders89}, indicating that a few tens of higher than that of the shell region.  They concluded that the metal rich plasma was a fossil of the SN explosion. The abundance ratio of Si, S and Fe indicated the progenitor star mass to be 25\\Msun.  Miyata \\& Tsunemi (1999) measured the radial profile inside the Loop and found a discontinuity around 0.9\\,R$_\\mathrm{s}$ where R$_\\mathrm{s}$ is the shock radius.  They measured the metallicity inside the hot cavity and estimated the progenitor mass to be 15\\Msun.  Levenson et al. (1998) estimated the size of the cavity and the progenitor mass to be 15\\Msun.  Therefore, the progenitor star of the Cygnus Loop is a massive star in which the triple-$\\alpha$ reaction should have dominated rather than the CNO cycle.  If the surrounding material of the Cygnus Loop is contaminated by the stellar activity of the progenitor star, it may explain the C abundance inferred for this region with Suzaku \\cite{katsuda07}.  In order to study the plasma condition inside the Cygnus Loop, we observed it from the NE rim to the south-west (SW) rim with the XMM-Newton satellite.  We report here the result covering a full diameter by seven pointings. ",
        "conclusions": "We have observed the Cygnus Loop along the diameter from the NE rim to the SW rim employing XMM Newton.  The FOV is divided into two paths: the north path and the south path.  Then it is divided into many small annuli so that each annulus contains a similar number of photons to preserve statistics.  The spectra from the rim regions can be expressed by a one-\\kTe~component model while those in the inner region require a two-\\kTe~component model.  The low \\kTe~plasma shows relatively low metal abundance and covers the entire FOV.  It forms a shell that originates from the preexisting cavity.  The high \\kTe~plasma shows high metal abundance and occupies a large part of the FOV. The origins of these two components are different: the high \\kTe~plasma with the high metal abundance must come from the ejecta while low \\kTe~plasma with low metal abundance must come from the cavity material.  We find that the thickness of the shell is very thin in the south west part where, we guess, the ejecta plasma is blow out in the direction of our line of sight.  We estimate the mass of the metals.  Based on the relative metal abundance, we find that the Cygnus Loop originated from a 15\\,\\Msun~star.  The distribution of the ejecta is asymmetric, suggesting an asymmetric explosion."
    },
    "0710/0710.3130_arXiv.txt": {
        "abstract": "We have considered precession in accretion disks in which a second moment of inertia relative to an axis perpendicular to the axis of rotation may be very important. This formalism, that takes into account the precession contribution to the angular momentum, is based on the existence of a parameter $\\it p$ which  determines three characteristic densities resulting from the averaging process and imposes constraints on the actual disk density. It is shown that the precession velocity will lie in a three branch solution, and depends on how large is the disk actual density as compared to the characteristic densities. Besides the large spread on the solution for the precession velocity, depending on the density strength, it may be prograde and retrograde. It is shown that the keplerian thin disk, with very large density values compared to characteristic ones  , only precesses very far away from the primary object, which implies very large precession periods. For other models, the disk will thicken, with large deviations from the keplerian approximation. Constraints on the density only will be effective for very large values of the ratio ${{\\dot M} \\over M_{p}}$, respectively, accretion rate and mass of the primary. Under this condition, the structure of the precessing region is found. Lower bounds on the precession period are found for not so large values of this ratio. Deviations from the mean precessional motion are considered. It is shown that these deviations result in periodic motions as long as the time scales associated to them are comparable to the remaining time scales. Otherwise, they result in misalignment motions, forcing the plane of the disk to become normal to the orbital plane of the secondary. ",
        "introduction": "Precessional activity has long been invoked to explain time variations occurring in the spectra of  galactic X-ray binaries like LMC X-4, Her X-1, SS 443, Cygnus X-1, etc. Observational evidences give support to the idea that these systems, and quite a lot of suspected others,  consist of a disk-like system plus a third body. This is very suggestive  of precession as being the effect  due to the perturbing torque of a distant companion star in the the disk or, even,  as the result in the disk of the secondary star precessing in the tidal field of the central compact object. As a matter of fact, the scenario is not unique, and according to \\citet{prie87}, there are, at least , five possibilities. This is so because there is uncertainty concerning the reason why the disk precesses, where and the extent on it that precesses. The possibilities  explicitly dependent upon tidal torques under which the disk will precess  can be summarized by the following two scenarios:  a- the disk has a permanently tilted edge that freely precesses in the tidal field of the companion star. This kind of scenario has been first exploited by \\citet{ka73}, trying to explain the 35 day period on the light curve of Her X-1.  This same model was used by \\citet{ka80}  to explain observational features in SS 443;  b- the companion star precesses in the tidal field of the central compact star, slavishly followed by the the outer tilted disk. This kind of scheme was proposed by \\citet{rob74}, also to explain the 35 day period observed in Her X-1.  However, precession is not a exclusiveness of galactic X -ray binaries sources, its presence being suspected in a lot of other astrophysical systems. Jet structures observed in AGN, and also in protostars seem to be associated with a precessing disk ( \\citet{li99}). In the context of disks around supermassive objects, the reason why the disk precesses is still more unclear. \\citet{ka97}  suggested that observational data in OJ 287 could be explained by a tidal torque due to the presence of another massive companion star, in a way similar to that employed to Her X-1 and SS 443.  \\citet{prin96} and \\citet{prin97} suggests precession as a result of  a warping instability caused by the central source irradiation.  \\citet{rom00} argue, in analogy with \\citet{ka82}and \\citet{ka97},, aiming to explain the pattern of ejection observed in the Quasar 3C 273, that  precession is due to a massive secondary black hole tidal torque on the disk,a model, subsequently,  applied to 3C 345 by \\citet{ca02}.  It should be argued, however, that precession is a kind of solid-like response of the disk. A fluid-like response would be a more appropriate one for a fluid system like the disk. Besides,   in the case of fluid-like response, the tidal torque acts as a perturbation, which main consequence is the generation of different sort of waves. Quite a variety of physical issues, suitable for a wave approach, such as several tidally driven instabilities, disk tidal deformation (warping), horizontally  and vertically driven resonances, angular momentum transport, and so on (\\citep{lu91, lua92, lu93, olga02, olgb03, vis92}) have been tackled quite successfully.  In addition, it has been shown  that, under this approach,  for disks in systems with extreme mass ratio, precession may occur due to a coupling between an eccentric instability and Lindblad resonances (\\citet{lua92}).  \\citet{pat95}, extending \\citet{lua92}    result, argue that the axisymmetric part of the tidal potential is related to the disk solid-like response, while the disk fluid-like response is related to the non-symmetric part of the potential. In these works, the the disk solid body-like response comes as an integrabil;ity condition.   \\citet{ka82} proposal to compute frequencies treats the disk as a ring, with no allowance given for the width in the plane of the ring, nor to the thickness perpendicular to that plane. By employing a suitable averaging procedure, he obtains for the for the precession angular velocity  $${\\Omega}_{p}=-\\frac{3}{4} \\, \\frac{{{\\omega}_{s}}^2}{{\\Omega}_{d}} \\, cos{\\delta},$$  where ${\\omega}_{s}$ is the keplerian angular velocity of the secondary, ${\\Omega}_{d}$ is the ring angular velocity, and $\\delta$ is the inclination angle. This formalism, besides treating the disk as a rigid body, not taking into account its structure, assumes the disk angular velocity parallel to the disk  angular momentum . Though keeping $\\vec{{\\Omega}_{p}}$ parallel to ${\\vec L}$, assuming thin disks obeying polytropic equation of state, the structure has been considered by \\citet{pat95}, \\citet{laral96}, \\citet{larp97} and \\citet{lar97}. Decomposing the secondary tidal potential into odd and even z parts, with subsequent Fourier decomposition, and arguing that the odd axisymmetric part is responsible for rigid precession, these authors,assuming ${\\ell} \\over r$ constant, obtain  $${\\Omega}_{p}=-\\frac{3}{4} \\,   \\, {\\left( \\frac{7-2 \\, n}{5-n} \\right)} \\, \\frac{{{\\omega}_{s}}^2}{{\\Omega}_{d}} \\, cos{\\delta},$$  as an integrability condition. In this expression, n is the polytropic index and ${\\Omega}_{d}$ is the disk angular velocity evaluated at the outer edge of the disk, $R_{0}$. For $n=2$, they recover \\citet{ka82} result. According to them, the validity of this result holds whenever the disk doesn't thicken, the perturbative tidal potential is weak, and, above all, the sound crossing time  is shorter than the precession period. It should be argued, however, that the disk angular momentum and the precession velocity are not parallel. Besides, heat generation, cooling and angular momentum transport do occur in accretion disks, a fact that makes not realistic the structure obtained  from a polytropic equation of state. This imposes severe restrictions on the model of disk we want to consider.It is reasonable to expect the polytropic index varying from $n=\\frac{3}{2}$, for a (monoatomic) gas pressure dominated disk, to $n=3$ for a radiation pressure dominated one. Disks with negligible temperature gradients would require  a polytropic index much greater than 3: $\\gamma =1$ is appropriate for an isothermal atmosphere ( $n \\rightarrow \\infty $). In that situation, the mass content in the disk, the moment of inertia and the precession velocity are determined by values close to the inner radius rather than close to the outer radius, the disk structure being very susceptible to relativistic corrections and gravitational radiation losses    may become important.  Besides,  if one is interested in a region in parameter space $M-{\\dot M}$ where the disk thickens, the dynamics can not be dissociated from the structure, as is the case for the keplerian thin disk. In other words, the azimuthal velocity depends on the disk height scale      . In addition, for a given tidal potential, its relative strength increases as the disk thickens. For actual radiation pressure dominated, the height scale is constant, the polytropic index is negative, implying negative specific heat. Besides this, in this case, one can not satisfy null boundary condition at the surface of the disk  Finally, the precession period diminishes relative to the sound crossing time as the disk gets thick. A final possible question concerning the applicability of of these results to disks is related to the precession velocity dependence on the outer radius of the disk, which yields proportionality with the mass of the disk. Despite implying larger inertia moments, the larger the mass the faster the disk will precess.   \\citet{rom00} work consists essentially of the application of \\citet{ka82} and \\citet{ka97} to a binary system composed by a disk around a primary supermassive black hole, plus a supermassive secondary orbiting the primary in a keplerian way, not coplanar to the disk. It should be argued, however, that the formulation followed by these authors, besides suffering from the drawbacks, we just mentioned, of any precessing disk model, relies heavily on the above cited authors results for the precession angular velocity, which was obtained without considering  the precession contribution to the total angular momentum. Besides, this model is supposed to apply to disks eventually dominated by radiation pressure.  In addition, account of the precession contribution to the angular momentum introduces some non-linearity in the problem, which, in turn, introduces conditions on the averaging procedure, imposing constraints on the density, on the disk unperturbed angular velocity, on the viscosity parameter as well as on the height scale of the disk. In that case, there will be a contribution from the precession velocity to a motion in a plane perpendicular to the plane containing the normal to the disk and the normal to the plane of the orbit of the secondary. This motion, depending on the magnitude of its time scale, will be periodic or not .  In the following we shall address the question of precession in accretion disks in systems similar to that considered by \\citet{rom00}. In our formulation, however, we will be concerned to treat the secondary tidal torque on the disk as a perturbation, which means that the secondary never crosses the disk, being far away from any point in the disk. Granted this, we will take into account the precession contribution to the total angular momentum and will discuss the conditions under which we may take an average of the Euler equations, in such a way as to keep algebraic solutions to the precessional motion. Deviations from the precessional motion will be treated by solving the equation for the time evolution of $\\delta$, the angle that measures the misalignment between the normal to the plane of the disk and the normal to the plane of the secondary orbit. The time evolution will be obtained assuming the time scale for this motion is large or comparable to the remaining time scales. We obtain results for the precession velocity for different disk models, taking into account deviations from geometrical thinness due to a relation between the angular velocity of the disk and the height scale. One of the main results of this work is to show that, if precession is to occur in the inner region of disks around a supermassive black hole, the disk is not keplerian, being very thick and luminosity deficient. ",
        "conclusions": "We have treated the problem of precession in accretion disks taking into account its contribution to the total angular momentum. We have looked for conditions under which the problem may be treated by solving the algebraic Euler equations, i.e., constant angular velocities. We have found that the problem is characterized by the parameter p, given by  $$p=  {\\frac{256}{27}} {{cos^2{\\delta} \\, cot^2{\\delta}} \\over {\\beta}^4} \\, {\\left( {M_{p} + M_{s}}\\over {M_{p}} \\right)}^2 \\, {\\left( {r \\over d} \\right)}^4 ,$$  and the densities  $${\\rho}_{min}=  12 cos{\\delta} cot{\\delta} {\\left({{M_{p} + M_{s}}\\over {M_{p}}}\\right)} {\\frac{K}{{\\beta}^2 \\, {d}^3}}, $$  $${\\rho}_{c}= {K \\over r^3} ,$$  $${\\rho}_{*}=  512 \\,  cos^3{\\delta}  cot^3{\\delta}   {\\left( {M_{p} + M_{s}}\\over {M_{p}} \\right)}^3  {\\left( {r \\over {{\\beta}  d}} \\right)}^6 {K \\over r^3} .$$  For $p \\leq 1$, the only constraint is $\\rho \\geq {\\rho}_{min} $, and the solutions, as a function of the density,  lie in three branches: the upper one, with prograde precession velocities, at the upper left of $\\rho ={2 \\over 3} {\\rho}_{c} $; the lower branch, with retrograde velocities, at the lower right of $\\rho ={2 \\over 3} {\\rho}_{c} $; and the middle branch, between the upper and the lower, in which the solution is retrograde for $\\rho < {r \\over d} {\\rho}_{c} $, and prograde for $\\rho  > {r \\over d} {\\rho}_{c} $. For $ p \\geq 1$, part of the $\\rho $-space, $ {\\rho}_{c} \\leq {\\rho} \\leq {\\rho}_{*} $, is not allowed. Using simple energy and angular momentum transport arguments, together with constraints obtained in that formulation, it is shown that, for $p >1$, $ \\alpha \\rightarrow 1$ and ${\\ell} \\rightarrow r$, leading to the breakdown of the keplerian and thin disk approximations. In that situation the disk is luminosity defficient.  As a matter of fact, a very important contribution of this paper is to show the incompatibility between keplerian disk and thick disk. The procedure we have adopted is suitable for treating deviations from the mean precessional motion, the misalignment motion, as well. For different limiting expressions for the precession velocity, we have obtained the time evolution for $ \\delta$, the angle that measures the misalignement of the plane of the disk and the orbital plane of the secondary star, the misalignment  angle. It has been shown that the time evolution of ${\\delta}$ depends on the magnitude of the time scale associated to this motion. Periodic motions only occur if this time scale is comparable to the remaining ones. Otherwise, there will be a tendency to force the plane of the disk to become normal to the secondary orbital plane.   It is shown that, if the disk is keplerian, full misalignment, i.e. ${\\delta}= 0.5 \\, \\pi$, is reached in a time comparable to the precession period, which means that the solution for the system is no longer given algebraically. We should make resort to the solutions of the differential Euler equations, a formidable task far beyond our goal in this paper. Finally, for different disk models we have shown that the constraints on the density will hold as long as ${{\\dot M}_{1} \\over M_{9}} >> 1$. In that situation, we have found the properties of the precessing region as a function of ${\\dot M}_{1}$ and $ M_{9}$. An expression for the separation distance between the primary and the secondary is also found. Otherwise, if ${{\\dot M}_{1} \\over M_{9}} < 1$, lower bounds on the precession period and on the separation distance are found. These bounds depend on ${\\dot M}_{1} \\over M_{9}$.  Application of this formalism to different disk models, may be summarized as follows: 1- for the keplerian thin $\\alpha$ standard disk model, the precession period (in years) will be  $$ T_{p} \\geq  0.08 \\, z_{d}^{1.5}  \\,  M_{p}, $$  For accretion disks in binary systems, we may take the size of the disk as the truncation radius, the point where the Roche equipotentials first intersect. In our systems we are not allowed to do so but, in any case, we found in the literature  that values for $ z \\approx 100-1000 $ are not unusual. Besides, expecting precessiuon periods of about hundred years, the system will reach full misalignment in a time of the order of the precession period.  The disk is not affected by constraints on the density.  2- For disk in which radiation dominates pressure and cooling, if ${{\\dot M}_{1} \\over M_{9}} > 75.76$ for a Schwarzschild black hole and ${{\\dot M}_{1} \\over M_{9}} > 12.62$ for a Kerr black hole, density in the disk may come close to the critical density ${\\rho}_{*}$. In that situation, ${\\ell} \\approx r$, and the size of the disk can not exceed  $$z_{d}= 0.0132 \\, {{\\dot M}_{1} \\over M_{9}}.$$  Besides, if the system is not to be affected by gravitational radiation, ${{\\dot M}_{1} \\over M_{9}}$ should be much larger than the values given above. Precession period will be  $$T_{p}= 2.81\\times 10^{-3} \\, cos{\\delta} \\, {\\left( \\frac{{{\\dot M}_{1}}^3}{{M_{9}}^5} \\right)}^{1/2} .$$  A reasonable value for $z_{d}$ would imply a very large ratio ${{\\dot M}_{1} \\over M_{9}}$. We are, therefore, led to the suspicion that the assumption $ {\\ell} \\approx r$, or $p_{0} \\approx 1$, is very strong for this model. Abandon of this assumption leads to a lower bound on the precession period, given by  $${T_{p}}^2=0.02 \\, cos^2{\\delta} \\, \\frac{{z_{d}}^3}{{M_{9}}^2} \\, \\frac{a^2}{{\\left(4 \\, a^2-{\\left( {\\left(1+4 \\, a^2 \\right)}^{1/2}-1\\right)}^2 \\right)}} .$$  3-Finally, for a radiation pressure and advective cooling dominated disk, $ {\\ell} \\approx r$, and its size  $$z_{d}= 0.75 \\, {{\\dot M}_{1} \\over M_{9}}.$$  Again, if the system is not to be affected by gravitational radiation, otherwise it will live for a short period of time, ${{\\dot M}_{1} \\over M_{9}} >> 1$. If ${{\\dot M}_{1} \\over M_{9}}$ is not too large, ${\\ell} << r$,  and the constraints on the density will not be effective. Again, there will be a lower bound on the precession period given by  $${T_{p}}^2=0.031 cos^2{\\delta} \\, \\frac{z^3}{{M_{9}}^2} \\, {{\\left( {\\left(1-4.4 \\, a+8.84 \\, a^2 \\right)}^{1/2}+1-2.2 \\, a \\right)}^2\\over {\\left( {\\left( {\\left(1-4.4 \\, a+8.84 \\, a^2 \\right)}^{1/2}+1-2.2 \\, a \\right)}^2-4 \\, a^2 \\right)}}.$$  To finalize, we would stress the main contribution of this work as pointing out the possibility of considering the disk structure, when studying disk precession, by means of an alternative procedure."
    },
    "0710/0710.4526_arXiv.txt": {
        "abstract": "Halo coronal mass ejections (HCMEs) originating from regions close to the center of the Sun are likely to be geoeffective. Assuming that the shape of HCMEs  is a cone and they propagate with constant angular widths and velocities, at least in their early phase, we have developed a technique (Michalek et al. 2003) which allowed us to obtain the space speed, width and source location. We apply this technique to obtain the parameters of all full HCMEs observed by the Solar and Heliospheric Observatory (SOHO) mission's Large Angle and Spectrometric Coronagraph (LASCO) experiment until the end of 2002. Using this data we examine which parameters determine the geoeffectiveness of HCMEs. We show that in the considered period of time only fast halo CMEs (with the space velocities higher than $\\sim 1000{km\\over s}$ and originating from the western hemisphere close to the solar center could cause the severe geomagnetic storms. We illustrate how the HCME  parameters can be used for space weather forecast. It is also demonstrated that the strength of a geomagnetic storm does not depend on the determined width of HCMEs. This means that HCMEs do not have to be very large  to cause major geomagnetic storms. ",
        "introduction": "Coronal mass ejections (CMEs)  originating from regions close to the central meridian of the Sun and directed toward  Earth cause the most severe geomagnetic storms (Gosling, 1993; Kahler, 1992; Webb et al., 2001). Many of these Earth-directed CMEs  appear as an enhancement surrounding the  occulting disk of coronagraphs. We call them halo CMEs (Howard et al. 1982).  The measured properties of CMEs include their occurrence rate, direction of propagation in the plane of the sky, angular width, and speed (e.g. Kahler, 1992; Webb, 2000; St. Cyr et al., 2000, Gopalswamy et al., 2003a; Gopalswamy, 2004; Yashiro et al., 2004). It is well known  that the geoffective CMEs originate mostly within a latitude $\\pm30^o$ (Gopalswamy et al., 2000a, 2001; Webb et al., 2000, 2001; Wang et al., 2002; Zhang et al. 2003). Srivastava and Venkatakrishan (2002) showed that the initial speed of the CMEs is correlated with the $D_{ST}$ index strength of the geomagnetic storm, although their conclusion was based only on the study of four events. This tendency was also suggested earlier by Gosling et al. (1990) and Tsurutani $\\&$ Gonzalez (1998). On the other hand, Zhang et al. (2003) demonstrated that both slow and fast HCMEs can cause major geomagnetic disturbances. They  showed that geoeffective CMEs are more likely to originate from the western hemisphere than from the eastern hemisphere. They also demonstrated a lack of correlation between the size of X-ray associated with a given CME and the importance of geomagnetic storms. Unfortunately, these studies were based on the sky plane speeds of CMEs without consideration of the projection effects. The parameters describing properties of CMEs, especially for HCMEs, are affected by projection effects (Gopalswamy et al., 2000b). Assuming that the shape of HCMEs is a cone and they propagate with constant angular widths and speeds, at least in their early phase of propagation, we have developed a technique (Michalek et al. 2003) which allows us to determine the following parameters: the linear distance $r$ of source location measured from the solar disk center,  the angular distance $\\gamma$ of source location measured from the plane of sky,  the angular width $\\alpha$ (cone angle =$0.5\\alpha$) and  the space velocity $V$ of a given HCME. A similar cone model was used recently by Xie et~al.~(2004) to determine the angular width and orientation of HCMEs.  The present paper is divided into two parts. First, in the Section~2 we applied the cone model (Michalek et al. 2003)  to obtain the space parameters of all HCMEs observed by the Solar and Heliospheric Observatory (SOHO) mission's Large Angle and Spectrometric Coronagraph (LASCO) until the end of 2002. In the Subsection~2.2 a short statistical analysis, based on the derived parameters, of HCMEs is presented  (Fig.~1~-~Fig.~4). In the Section~3,  we use these parameters to identify the most important factors determining geoeffectiveness of HCMEs and how they could be used for space weather forecast (Fig.~5~-~Fig.~17) ",
        "conclusions": "In this study we considered the geoeffectiveness of all full HCMEs observed by SOHO/LASCO coronagraphs from the launch in 1995 until the end of 2002. For $101/144$ ($70\\%$) of full HCMEs we were able to find the source location, width and space velocity using the cone model (Michalek et al., 2003). We must be aware  that the cone model is only rough simplification of real events. We know that not all CMEs are perfectly symmetric (Moran and Davila, 2004; Jackson et al., 2004). Most of CMEs could be approximate using cone model but probably for some of them this assumption is unrealistic. Fortunately technique presented by Michalek et al. does not demand  perfect symmetry for CMEs. This approach requires measurements of sky-plane speeds and the moments of the first appearance of the halo CMEs above limb at only two opposite points. We are able to determine,  with good accuracy, the space velocity and with of a given CME at least in the plane symmetry crossing CMEs at these points. When a given CME could be approximated by the cone model these derived parameters are  valid for the entire CME. HCMEs originating very close to the disk center (mostly within a latitude of $\\pm40^o$), are very wide (the average angular width $=120^o$) and are very fast (the average space speed $=1291km/s$). We find significant ($40\\%$) increase in the average space velocities of HCMEs during the maximum of solar activity. These results could suggest that the HCMEs represent a special class of CMEs which are very wide and fast. It is important to note that this \"class\" of CMEs is defined due to artificial effect caused by coronagraphic observations. Events originating close to the disk center (from SOHO/LASCO point of view) must be wide and fast to appear as HCMEs in LASCO observations. This is not due to localization on the solar disk but due to  oculting disk which not only blocks bright photospheric  light but also eliminates some narrow and slow events. We have to emphasize  that this effect mostly depends on the dimension of oculting disk but in less degree on the sensitivity of instrument. More sensitive instrument can record some poorer events (halos and also not halos so statistic will be similar) but could not register these events which never appear behind occulting disk. Potentially more sensitive instrument could register less energetic events (narrower and slower) and the average velocities and widths (for halos and whole population of CMEs) could be slightly lower but the main relation between the halos and   whole population of events will be the same. Fortunately, poor events do not cause a big concern because they are not geoeffective. We do not expect, in the near future, any  special programs devoted to looking for less energetic CMEs. Next scientific mission (STEREO) will be mostly dedicated to recognize 3D structure of CMEs. Such fast and wide CMEs are known to be associated with electron and proton acceleration by driving fast mode MHD shocks (e.g., Cane et al., 1987; Gopalswamy et al., 2002a). Using observations from Wind spacecraft, interplanetary magnetic clouds (MC) and geomagnetic disturbances associated to HCMEs were identified. The strength of geomagnetic storms, described by $D_{ST}$ and $Ap$ indices, is highly correlated with the source location and space velocity of a given event. Only HCMEs originating in the western hemisphere, close to the solar center and very fast (space velocity $\\geq 1100km/s$) are likely to cause major geomagnetic storms ($D_ST<-150nT$). Slow HCMEs (space velocity $\\leq 1100km/s$), even originating close to the solar center, may not cause severe geomagnetic disturbances. We have to note that there was one event (04 April 2000), which originated far from disk center and produced a severe geomagnetic storm ($D_{ST}=-288nT$). Probably this storm was not due to an ICME. It was caused by the sheath region ahead of the CME as was reported by Gopalswamy (2002b). We illustrated, using contour maps, how the derived HCME parameters can be useful for space weather forecast. We have to note that geoeffectiveness of events does not depend on their widths.  During our study period we recognized $56/144$  ($30\\%$) FHCMEs without any geomagnetic signature at Earth.  This is significant population of FHCMEs. To distinguish them from the geoeffective events we considered the source locations and space velocities of HCMEs. When both the parameters are available, it becomes easier to assess the geoeffectiveness of HCMEs. We may say that fast FHCMs ($V>1200km/s$) originating close to the disk center ($|longitude|<30^o$) must be geoeffective. For such events there were no false alarms. But, even very fast events originating far from the disk center can be non-geoeffective."
    },
    "0710/0710.1303_arXiv.txt": {
        "abstract": "We propose to study cosmic reionization using absorption line spectra of high-redshift Gamma Ray Burst (GRB) afterglows. We show that the statistics of the dark portions (gaps) in GRB absorption spectra represent exquisite tools to discriminate among different reionization models.  We then compute the probability to find the largest gap in a given width range $[W_{\\rm max}, W_{\\rm max}+dW]$ at a flux threshold $F_{\\rm th}$ for burst afterglows at redshifts $6.3 \\le z \\le 6.7$. We show that different reionization scenarios populate the $(W_{\\rm max},F_{\\rm th})$ plane in a very different way, allowing to distinguish among different reionization histories. We provide here useful plots that allow a very simple and direct comparison between observations and model results. Finally, we apply our methods to GRB~050904 detected at $z=6.29$. We show that the observation of this burst strongly favors reionization models which predict a highly ionized intergalactic medium at $z\\sim 6$, with an estimated mean neutral hydrogen fraction $x_{\\rm HI}=6.4\\pm 0.3 \\times 10^{-5}$ along the line of sight towards GRB~050904. ",
        "introduction": "In the last few years, our knowledge of the high-$z$ Universe and in particular of the reionization process has been enormously increased mainly owing to the observation of quasars by the SDSS survey (Fan 2006) and CMB data (Hinshaw et al. 2007, Page et al 2007). Long gamma ray bursts (GRB) may constitute a complementary way to study the reionization process avoiding the proximity effects and possibly probing even larger redshifts. This has now become clear after the detection of five GRBs at $z\\gsim 5$, over a sample of about 200 objects observed with the {\\it Swift} satellite (Gehrels et al. 2004). The current record holder is GRB~050904 at $z = 6.29$ (Tagliaferri et al. 2005, Kawai et al. 2006). Totani et al. (2006) have used this object to constrain the ionization state of the intergalactic medium (IGM) at high redshift by modeling its optical afterglow spectrum. They report the evidence that the IGM was largely ionized already at $z=6.3$. The best-fit neutral hydrogen fraction is consistent with zero with upper limit $x_{\\rm HI}<0.17$ ($<0.6$) at 68\\% (95\\%) C.L.  Several authors (Natarajan et al. 2005, Daigne, Rossi \\& Mochkovitch 2006, Bromm \\& Loeb 2006, Salvaterra et al. 2007a) have computed the number of high--$z$ GRBs detectable by {\\it Swift}. In spite of model details, all these different studies consistently predict that a non-negligible fraction (up to $\\sim 10$\\%) of all observed GRBs should lie at very high redshift. On the basis of these results, several GRBs at $z\\ge 6$ should be observed by {\\it Swift} in the near future. Depending on the {\\it Swift}/BAT trigger sensitivity and on model details, Salvaterra \\& Chincarini (2007) found that $\\sim 2-8$ GRBs can be detected above this redshift during every year of mission. From the observational point of view, Ruiz-Velasco et al. (2008) find that the fraction of {\\it Swift} GRBs at $z>6$ should not exceed the conservative upper limit of $19\\%$. Moreover, future X-- and Gamma--ray missions will increase rapidly the sample of high--$z$ GRBs, possibly up to $z\\sim 10$ (Salvaterra et al. 2007b). Once detected, spectroscopic follow-up observations of high--$z$ GRBs require a rapid trigger of  8-meter, ground based telescopes. This can be done by pre-selecting reliable candidates on the basis of some promptly-available information provided by {\\it Swift}, such as burst duration, photon flux, the lack of detection in the UVOT $V$-band, and the low Galactic extinction (Campana et al. 2007, Salvaterra et al. 2007a). We note that cosmological time dilation helps keeping the flux bright, since observations will sample the afterglow early phases, even a few days after trigger.  The spectra of high-redshift sources (as QSOs and GRBs) bluewards of the Ly$\\alpha$ are characterized by dark portions (gaps) produced by intervening neutral hydrogen along the line of sight. The use of various gap statistics in QSO spectra has been recently recognized as a very powerful tool to constrain the IGM ionization state (Paschos \\& Norman 2005; Fan et al. 2006; Gallerani, Choudhury \\& Ferrara 2006, hereafter G06; Gallerani et al. 2007, hereafter G07). For example, by comparing the statistics of these spectral features in a sample of 17 observed QSOs with Ly$\\alpha$ forest simulations, G07 concluded that the HI fraction, $x_{\\rm{HI}}$, evolves smoothly from $10^{-4.4}$ at $z=5.3$ to $10^{-4.2}$ at $z=5.6$, with a robust upper limit $x_{\\rm{HI}} < 0.36$ at $z=6.3$. These results encourage the application of such analysis to GRBs. There are several advantages promised by such an attempt. First, GRBs are soon expected to be found at redshifts higher than those typical of QSOs; second, and contrary to the massive hosts of QSOs, they reside in ``average\" cosmic regions only marginally affected by local ionization effects or strong clustering; finally, they are bright and their afterglow spectra closely follows a power-law, making continuum determination much easier  In agreement with the WMAP3 results (Spergel et al. 2007) we assume a flat universe with $\\Omega_m=0.24$, $\\Omega_{\\Lambda}=0.76$, $\\Omega_bh^2=0.022$, $h=0.73$. The parameters defining the linear dark matter power spectrum are $n=0.95$, $\\de n/\\de \\ln k=0$, $\\sigma_{8}=0.82$. Mpc are physical unless differently stated. ",
        "conclusions": "We have proposed to investigate cosmic reionization using absorption line spectra of high--$z$ GRB afterglows. The evolution of the largest gap width (LGW) as a function of the flux threshold  $F_{\\rm th}$ used to define gaps shows marked differences in the two reionization models favored by present data. In particular, (i) the LGW is typically $\\sim 2$ times wider in the LRM than in the ERM, for a fixed $F_{\\rm th}$; (ii) the difference between ERM and LRM in terms of $(W_{\\rm max}, F_{\\rm th})$ increases with $z$; (iii) an overall shift towards larger $W_{\\rm max}$ values for a fixed $F_{\\rm th}$ is found in both models towards higher redshifts. This analysis has shown that we can robustly distinguish among different reionization histories: improved results can be obtained if data is collected promptly after burst detection and by using GRBs at the highest available redshifts.  A direct comparison of the model with data can be carried on in terms of the probability to find the largest gap in a given width range for burst afterglows at $z=z_{\\rm GRB}$.  When applied to the only known GRB at $z>6$, i.e. GRB~050904 at $z_{\\rm GRB}=6.29$, a clear indication is obtained that reionization must have occurred well before $z=6$.  We find that the observed LGW in the GRB~050904 afterglow spectrum are consistent with $x_{\\rm HI}=6.4\\pm 0.3\\times 10^{-5}$. This result is in agreement with previous measurements by Totani et al. (2006), who find that $x_{\\rm HI}$ is bound to be $x_{\\rm HI}<0.17$ ($<0.6$) at 68\\% (95\\%) C.L.  Some gaps/peaks in absorption spectra could be due to DLAs/HII regions intervening along the LOS towards the background source (e.g. the DLAs/transverse proximity effect detection by Totani et al. 2006/Gallerani et al. 2007). Such contaminants could affect neutral hydrogen measurements, if the foreground sources are not properly removed.  In this work we have applied the gap statistics to the GRB 050904 observed flux. However, the same analysis can be done by using the transmitted flux $F_{\\rm obs}/F(\\nu)$, i.e. normalizing the observed flux to the continuum. Following this approach, we provide useful plots (Fig. 2) which allow a straightforward comparison between our model results and future observations of $z>6$ GRBs afterglow spectra.  It is worth noting that the gap statistics with a variable flux threshold can be readily applied to other background sources, as QSOs. We plan to apply this analysis to the Fan et al. (2006) data, thus improving the results obtained by Gallerani et al. (2007). The advantage of varying the flux threshold is that this technique allows to study the response of the transmitted flux both to the high and low end tail of the IGM density field. In fact, high (low) $F_{\\rm th}$ values probe regions corresponding to low (high) overdensities.  In spite of the fact that QSO data are presently more abundant, GRBs are expected to be found at higher redshift, as they should result from the death of (early) massive stars (e.g. Abel et al. 2002, Schneider et al. 2002). Larger samples of GRB afterglow spectra at $z\\ge 6$, likely available in the near future, will allow a statistically significant analysis and possibly to reconstruct the cosmic reionization history. At redshifts approaching reionization epoch, gaps in afterglow absorption spectra are expected to become as large as the spectral region between the Ly$\\alpha$ and Ly$\\beta$ emission lines, thus obscuring completely the GRBs optical counterpart."
    },
    "0710/0710.5300_arXiv.txt": {
        "abstract": "Using big bang nucleosynthesis and present, high-precision measurements of light element abundances, we constrain the self-gravity of radiation pressure in the early universe. The self-gravity of pressure is strictly non-Newtonian, and thus the constraints we set provide a direct test of this prediction of general relativity and of the standard, Robertson-Walker-Friedmann cosmology. ",
        "introduction": "Certain aspects of general relativity are well tested.  For example, the Schwarzschild metric has been quantitatively verified in the weak-field limit on small scales, e.g., the Solar system \\citep{Shapiro64, Bertotti03} and binary radio pulsars \\citetext{e.g., \\citealp{Hulse75, Taylor79, Weisberg84}}; and on galaxy scales \\citep[e.g.,][]{Bolton06}.  In another fundamental test of general relativity, the existence of gravity waves has been established  \\citetext{e.g., \\citealp{Taylor82}, \\citealp{Weisberg84}}. General relativity theory, utilizing the Robertson--Walker metric \\citetext{\\citealp{Friedmann22}, \\citealp{Robertson35}, \\citealp{Walker36}} leads to the Friedmann equations \\citetext{\\citealp{Friedmann22, Lemaitre27}} which govern the expansion behavior of a homogeneous, isotropic Universe. However, it is probably fair to say that the Friedmann equations, while providing a self-consistent and highly successful framework for cosmology, have not been subjected to extensive, independent testing.  In this paper we show that one particular aspect of the Friedmann equations, the self-gravity of pressure, can be tested quantitatively.  The development of big-bang nucleosynthesis (BBN) codes \\citep{Wagoner67, Kawano92}, coupled with measurements of the relevant nuclear reaction rates \\citep{Caughlan88, desc04}, have allowed observations of light element abundances to become powerful tools with which to investigate the early evolution of the universe. Computational predictions over a wide range of parameter space, when compared with primordial abundances inferred from observations, have yielded constraints on the current-epoch baryon density \\citep{Wagoner73, Yang79, KS04, Steigman07}, neutrino physics \\citep{SSG77, Yang79, KS04, Steigman07}, the fine structure constant \\citep{Bergstrom99}, the gravitational constant \\citep{Steigman76, Yang79, Copi04, Steigman07}, primordial magnetic fields \\citep{Kernan96}, the universal lepton asymmetry \\citep{Wagoner67, KS04, Steigman07} and other parameters of astrophysical interest.  Increasingly accurate measurements of element abundances, as well as improved understanding of the processes (i.e., stellar and galactic nucleosynthesis) which have altered the original abundances, allow these restrictions to be continually refined. Deuterium abundances \\citep{GeissReeves72, Omeara06}, helium abundances \\citep{HoyleTayler64, IT04, Izotov07}, and lithium abundances \\citep{RNB00, Asplund06} have all been well measured, although the inferred primordial abundances are subject to large and often difficult to quantify systematic uncertainties. More recently, observations of the cosmic microwave background (CMB) have yielded an independent estimate of $\\eta$, the baryon to photon ratio at a much later epoch in the evolution of the Universe \\citep{Spergel07}. ",
        "conclusions": "As illustrated in Figure 1, the combined constraints are, within the uncertainties, consistent with the general relativity prediction of $\\chi = 1$ and the independent (of BBN) WMAP constraint on $\\eta_{10}$ of $6.1 \\pm 0.2$ \\citep{Spergel07} which corresponds to $\\chi = 1$.  For the \\citet{PLP07} choice of Y$_{P}$, $\\chi = 1.00\\pm 0.14$, while for the Steigman \\cite{Steigman07} helium abundance, $\\chi = 0.84\\pm 0.25$.  Note that the data strongly exclude $\\chi = 0$.  The current light element observations and BBN computations have provided a test of the general relativistic self-gravity of pressure.  Since the modification of GR we are testing corresponds, for the radiation-dominated evolution appropriate for BBN, to an overall multiplicative factor of the product of Newton's gravitational constant and the radiation density, $G\\rho \\rightarrow G\\rho \\left({\\frac{1 + \\chi}{ 2}}\\right)$, our result is equivalent to the BBN constraint on the variation of Newton's constant or, alternatively, to a modification of the radiation energy density as parameterized by the effective number of neutrinos (see \\S1 for references). \\begin{equation} \\frac{1 + \\chi}{2} = 1 + \\frac{\\Delta G}{G} = 1 + \\frac{7\\Delta N_{\\nu}}{ 43} \\end{equation} Assuming that these other parameters take on their standard-model values ($\\Delta G = \\Delta N_{\\nu} = 0$), the self-gravity of the radiation (photons and neutrinos) pressure during the BBN epoch has been constrained quantitatively."
    },
    "0710/0710.1715_arXiv.txt": {
        "abstract": "AA~Dor is an eclipsing, close, post common-envelope binary (PCEB). We present a detailed spectral analysis of its sdOB primary star based on observations obtained with the Far Ultraviolet Spectroscopic Explorer (FUSE). Due to a strong contamination by interstellar absorption, we had to model both, the stellar spectrum as well as the interstellar line absorption in order to reproduce the FUV observation well and to determine the photospheric parameters precisely. ",
        "introduction": "The eclipsing binary system AA~Dor consists of a sdOB with \\mbox{$T_{\\rm{eff}}=42\\,\\rm{kK}$} and an unseen low-mass companion. Recent studies of its primary star have shown discrepancies in the determination of the surface gravity $g$. \\citet{Rauch00} obtained \\mbox{$\\log g\\hspace{-0.5mm} =\\hspace{-0.5mm}  5.21$} by a NLTE spectral analysis of optical and ultraviolet data. The results of radial-velocity and lightcurve analyses by \\citet{Hilditch96,Hilditch03} indicate a higher value of \\mbox{$\\log g\\hspace{-0.5mm} =\\hspace{-0.5mm}  5.45\\,-\\,5.51$}. Twelve high S/N-ratio FUSE observations were performed in August 2003 and June 2004. With about 200 secs each, the exposure times were chosen short in order to reduce the smearing by orbital motion. The FUV range includes the hydrogen Lyman series which in general is sensitive to changes in $\\log g$. However, a precise determination of the surface gravity and further stellar parameters is hampered by strong interstellar absorption lines. We used our NLTE code TMAP \\citep[T\\\"ubingen Model Atmosphere Package, ][]{Werner03} to calculate plane-parallel, line-blanketed model atmospheres in radiative, hydrostatic and statistical equilibrium. We included calcium and all elements of the iron group (Sc, Ti, V, Cr, Mn, Fe, Co, Ni) into our calculations. Due to the large amount of transitions, those elements were treated in a statistical approach \\citep{rd2003}. Atomic data were taken from NIST (National Institute of Standards and Technology, \\texttt{http://physics.nist.gov}), the Opacity Project \\citep{Seaton94}, and the Iron Project \\citep{Hummer93}. In case of the iron-group elements, energy levels and oscillator strengths stem from Kurucz's line lists \\citep{Kurucz91}. Only weak lines of those elements are detected in the FUSE spectra, which form absorption troughs due to rotational broadening (Fig. \\ref{fig:POSLIN}).  \\begin{figure}[ht!] \\begin{center} \\includegraphics[width=0.7\\textwidth]{fleig_01.eps}\\vspace{-2mm} \\caption[]{ Identification of iron-group lines. \\citet{Kurucz91} gives only a few lines with reliably measured wavelengths (POS lines, blue, shifted to the top) which makes the identification of isolated lines difficult. The addition of iron-group lines with theoretical positions (Kurucz's LIN lines, red) improves the overall-fit to the observation. The synthetic spectra are convolved with a rotational profile of $v_{\\rm{rot}}=35\\,\\rm{km\\,s^{-1}}$.}\\vspace{-5mm} \\label{fig:POSLIN} \\end{center} \\end{figure} ",
        "conclusions": "The effective temperature was derived using the ionization equilibrium of N\\,{\\scriptsize III} and N\\,{\\scriptsize IV}. The best fit is achieved for $T_{\\rm{eff}}\\hspace{-0.5mm}=\\hspace{-0.5mm}40 \\pm 3\\,\\mathrm{kK}$ (Fig\\@. \\ref{fig:Teff}). Due to the interstellar contamination of the spectra (Fig\\@. \\ref{fig:ISM_MOD}), a precise determination of the surface gravity was not possible (Fig\\@. \\ref{fig:logg}). However, a hint towards a higher value of $\\log g$ is given by the rotational velocity of the primary (Fig\\@. \\ref{fig:vrot}). With $v_{\\rm{rot}}=35 \\pm 5\\,\\rm{km\\,s}^{-1}$, the primary rotates slower than bound ($v_{\\rm{rot}}=45.7\\,\\rm{km\\,s}^{-1}$). \\citet{Rauch00} assumed bound rotation for the calculation of the primary's spectrum and thus, one can expect a slightly higher value of $\\log g$. Calculating the surface gravity directly via $\\log\\,g=\\log\\, [(GM_1)/r_1^2]$, $r_1\\propto v_{\\rm{rot}}\\cdot P$, returns a surface gravity of 5.44, which is in agreement with the value found by \\citet{Hilditch03}.  \\begin{figure}[ht!] \\begin{center} \\includegraphics[width=0.7\\textwidth]{fleig_02.eps}\\vspace{-2mm} \\caption[]{Determination of $T_{\\rm{eff}}$ from the \\mbox{N\\,{\\scriptsize{III}} / N\\,\\scriptsize{IV}} ionization equilibrium. While the \\mbox{N\\,\\scriptsize{IV}} line (left panel) appears saturated within the relevant $T_{\\rm{eff}}$, the \\mbox{N\\,\\scriptsize{III}} multiplet (right) matches the observation best at $T_{\\rm{eff}} = 40\\,\\mathrm{kK}$. } \\label{fig:Teff} \\end{center} \\end{figure}  \\begin{figure}[ht!] \\begin{center} \\includegraphics[width=0.7\\textwidth]{fleig_03.eps}\\vspace{-2mm} \\caption[]{Comparison of a pure photospheric model spectrum (blue) and the combined photospheric + ISM model (red) within a section of the FUSE observation. Interstellar lines are labeled with an asterisk.} \\label{fig:ISM_MOD} \\end{center} \\end{figure}  \\begin{figure}[ht!] \\begin{center} \\includegraphics[width=0.9\\textwidth]{fleig_04.eps}\\vspace{-2mm} \\caption[]{Log $g$ - dependence of the first four lines of the Lyman series. Lyman $\\alpha$ is compared to IUE observations, Lyman $\\beta - \\delta$ to FUSE observations. An interstellar H\\,{\\scriptsize I} column density of $\\log N_\\mathrm{H\\,{\\footnotesize I}} = 20.3$ and an extinction of $E_\\mathrm{B-V}=0.04$ are considered. Interstellar absorption prevents a precise determination of the continuum and the course of the line wings. All continua are normalized to that of the synthetic spectra with $\\log g=5.20$.} \\label{fig:logg} \\end{center} \\end{figure}  \\begin{figure}[ht!] \\begin{center} \\includegraphics[width=0.7\\textwidth]{fleig_05.eps}\\vspace{-2mm} \\caption[]{ Comparison of the \\mbox{C\\,{\\scriptsize III} $\\lambda\\lambda 1175\\,\\mathrm{\\AA}$} multiplet with the FUSE observation at three different rotational velocities $v_{\\rm{rot}}$. The best fit is achieved with $v_{\\rm{rot}}=35\\,\\rm{km\\,s}^{-1}$.} \\label{fig:vrot} \\end{center} \\end{figure}   \\begin{figure}[ht!] \\begin{center} \\includegraphics[width=0.9\\textwidth]{fleig_06.eps}\\vspace{-2mm} \\caption[]{Detection of phosphorus and sulfur. The resonance doublets of \\mbox{S\\,{\\scriptsize VI} $\\lambda\\lambda 933.38, 944.52\\,\\mathrm{\\AA}$} and \\mbox{P\\,{\\scriptsize V} $\\lambda\\lambda 1117.98, 1128.01\\,\\mathrm{\\AA}$} are well reproduced at solar abundances. Since interstellar absorption contributions cannot be excluded, these abundances are upper limits.} \\label{fig:doublets} \\end{center} \\end{figure}"
    },
    "0710/0710.3704_arXiv.txt": {
        "abstract": "{} {  % We study the evolution of the galaxy population up to $z\\sim 3$ as a function of its colour properties. In particular, luminosity functions and luminosity densities have been derived as a function of redshift for the blue/late and red/early populations.} {We use data from the GOODS-MUSIC catalogue which have typical magnitude limits $z_{850}\\leq 26$ and $Ks\\leq 23.5$ for most of the sample. About 8\\% of the galaxies have spectroscopic redshifts; the remaining have well calibrated photometric redshifts derived from the extremely wide multi-wavelength coverage in 14 bands (from the U band to the Spitzer $8 \\mu$m band). We have derived a catalogue of galaxies complete in rest-frame B-band, which has been divided in two subsamples according to their rest-frame U-V colour (or derived specific star formation rate, SSFR) properties.} {We confirm a bimodality in the U-V colour and SSFR of the galaxy sample up to $z\\sim 3$. This bimodality is used to compute the LFs of the blue/late and red/early subsamples. The LFs of the blue/late and total samples are well represented by steep Schechter functions evolving in luminosity with increasing redshifts. The volume density of the LFs of the red/early populations decreases with increasing redshift. The shape of the red/early LFs shows an excess of faint red dwarfs with respect to the extrapolation of a flat Schechter function and can be represented by the sum of two Schechter functions. Our model for galaxy formation in the hierarchical clustering scenario, which also includes external feedback due to a diffuse UV background, shows a general broad agreement with the LFs of both populations, the larger discrepancies being present at the faint end for the red population. Hints on the nature of the red dwarf population are given on the basis of their stellar mass and spatial distributions.} {} ",
        "introduction": "The evolution of galaxy Luminosity Function (LF) is one of the main tools to study the structure evolution through the cosmic time. The advent of large surveys has allowed the analysis of sub-samples of galaxies selected as a function of their morphological, spectroscopic or colour properties \\citep[][elsewhere G05]{strateva2001,norberg2002,madgwick2002,wolf2003,willmer,bell2004,baldry2004,hogg2004,weiner2005,blanton2005,ilbert2006,marchesini2006,baldry2006,driver2006,cucciati06,cirasuolo2006,arnouts2007,giallo2005}. In fact these kind of studies allow us to probe the evolution of galaxies having different star formation histories.  Of special interest are the studies concerning the statistical properties of galaxies selected on the basis of their intrinsic colour distribution. This distribution appears bimodal up to $z\\sim 2-3$ \\citep{baldry2004,blanton2005,giallo2005} and separates the galaxies in two populations, red early types vs. blue late types. It has been shown that this spectral classification is roughly consistent with the correspondent morphological classification (bulge vs. disc dominated) at least at low and intermediate redshifts \\citep{strateva2001,weiner2005}.  This bimodal colour distribution can find a natural explanation in hierarchical models for galaxy formation \\citep{menci2005,menci2006} where two distinct populations arise in the colour distribution based on two different star formation histories affected by the feedback effects produced by the SN and AGN activities \\citep{menci2006}. However the effect of environmental density on the paths of galaxy evolution can have a fundamental role. In this context it is not clear whether the evolutionary history of galaxies is originated by a {\\it nurture} senario (galaxy properties are affected by environment through physical mechanisms acting on galaxies) or by a {\\it nature} scenario  \\citep[the evolution is driven by the initial condition established during the formation epoch of galaxies, e.g. ][]{mateus2007,cooper2007}.  Recent studies have estimated the shape and evolution of the LF of galaxies selected according to their bimodal colour distribution using both the large local Sloan survey \\citep{baldry2004,blanton2005}, and other surveys at intermediate and high redshifts \\citep{bell2003,giallo2005,faber,willmer,ilbert2006}. Their results show the red LF evolving mildly in density up to $z \\sim 1$ with a quite flat shape at the faint-end, although the evaluation of the faint-end slope of the red LF remains an open issue especially at intermediate and high redshifts where the present surveys do not constrain the faint slope very well \\citep{faber,bell2003}.  In G05 we studied the red and blue LFs, using the properties of bimodality in colour and in specific star formation rate (SSFR), with a complete but relatively small sample of galaxies selected in the rest-frame B-band from low to high redshifts. We showed that the bimodality extends at least up to $z\\sim 2.5$. We also found that the red/early galaxies decrease in their luminosity density by a factor $\\sim 5-6$ from $z\\sim 0.5 $ to $z\\sim 2.5-3$ in broad agreement with the hierarchical cold dark matter model. These results provided a first picture of the evolution of the red and blue LFs up to high redshifts relaying on a relatively deep but small sample. For a more reliable picture a wider sample at high redshift is clearly needed. For this purpose larger areas with deep near-IR imaging are required.  Thanks to the wide area ($\\sim 140\\ arcmin^2$) and to the deep near-IR observations, the GOODS-South survey provides a good starting point for the study of the galaxy properties at high redshift. In particular, the inclusion of the deep IR observations obtained with the Spitzer telescope represent a useful constraint for the estimate of the physical properties of galaxies at high redshift. Last but not least the extensive spectroscopic follow up obtained in this field provides a wide set of spectroscopic redshifts. From this public data set we have obtained a multi-colour catalogue of galaxies we named GOODS-MUSIC \\citep[GOODS MUlticolour South Infrared Catalog, ][]{grazian}. This catalog, where galaxies are selected both in the $z_{850}$ and $Ks$ bands, contains information in 14 bands from the U to the Spitzer 8 $\\mu m$ band, and all the available spectroscopic information. For all the objects without spectroscopic information we have obtained well calibrated photometric redshifts by means of our photometric redshift code \\citep{fontana00}.  The GOODS-MUSIC catalogue has been already used to investigate the physical and clustering properties of high redshift galaxies \\citep{fontana06,grazian2006_eros,grazianc,pentericci2007,castellano2007}. With the present paper we study the galaxy LFs of the red and blue populations, enlightening evolutionary features which are characteristic of the two populations.  The paper is organised as follows: in section 2, we describe the basic features of our dataset. In section 3, we describe the bimodality properties of the sample and we define the loci of minimum for the selection of the red/early and the blue/late sub-samples as a function of $z$. In section 4, we compute the shape and evolutionary properties of the LFs and the luminosity density of both populations. Section 5 is devoted to the analysis of the physical properties of the faint early population.  All the magnitudes used in the present paper are in the AB system. An $\\Omega_\\Lambda=0.7$, $\\Omega_M=0.3$, and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$ cosmology is adopted. ",
        "conclusions": "We have used a composite sample of galaxies selected in deep NIR images, obtained from the GOODS public survey, to study the evolution of the galaxy LFs of red/early and blue/late galaxy populations selected using the colour and SSFR statistical properties of the sample.  The  observed  $U-V$  colour  and  SSFR  distributions  show  a clear bimodality up to $z\\sim 3$, confirming the results obtained in G05  at a  higher  level of  statistical  confidence. We  found  a trend  with redshift  for  the  colour magnitude  distribution  with an intrinsic blueing of about 0.15 mag.  in the redshift interval $z=0.4-2.0$  for both populations. This observed bimodality can be explained  in  a  hierarchical clustering scenario  as  due to the different  star formation histories of the red/early  and blue/late galaxies, see e.g. \\cite{menci2006}.  For the total and the blue/late sample the LF is well described by a Schechter function and shows a mild luminosity evolution in the redshift interval $z=0.2-1$ (e.g. $\\Delta M^* \\sim 0.7$ for the total sample; $\\Delta M^*\\sim 1$ for the blue/late fraction), while at higher redshifts the LFs are consistent with no evolution. A comparison with our hierarchical CDM model shows a good agreement at bright and intermediate magnitudes. A better agreement of the model has also been found at fainter magnitudes due to the suppression of star formation in small objects by the action of an ionising UV background.  The shape of the red/early luminosity function is better constrained only at low and intermediate redshifts and it shows an excess of faint red dwarfs with respect to the extrapolation of a flat Schechter function. In fact a minimum around magnitude $M_B(AB)=-18$ is present together with an upturn at fainter magnitudes. This peculiar shape has been represented by the sum of two Schechter functions.  We found that the bright one is constant up to $z\\sim 0.7$ beyond which it decreases in density by a factor $\\sim 5$ (10 for the early galaxies) up to redshift $z\\sim 3.5$. The comparison with our hierarchical CDM model shows that, although the predicted LF has a slight flattening at intermediate luminosity, the model still overpredicts the LF at faint magnitudes. The bright end of blue and red LFs at low and intermediate  redshifts is  in   good  agreement with  recent estimates  from   the  DEEP2 spectroscopic survey. As a consequence  of   this  complex  evolutionary  behaviour, the luminosity  densities  of  the  relatively   bright ($M_B(AB)<-20.2$) red/early and  blue/late galaxies  show a bifurcation beyond redshift $z\\sim  1$. Indeed  the  LD of  the blue/late  population  keeps  increasing  up  to    $z\\sim  3.5$, while the   luminosity density  of red/early galaxies decreases  by a factor  $\\sim 2-3$ respectively  in the $z=1-3.5$ interval.  To derive  hints on  the nature  of the  galaxies responsible  for the peculiar shape of the red/early  LF, we have performed an analysis of their stellar masses and spatial distribution. We found that the early galaxies have systematically higher stellar masses with respect to the late ones for a given B band luminosity. Brighter early galaxies have  a spatial distribution  more concentrated  in higher  density regions if compared  to  the late ones  of the  same  luminosity class.  On  the contrary, fainter early and late  galaxies  show a very similar spatial distribution. Thus, the  different environmental properties  do not seem  to  be the  main   responsible  for the difference  in shape at intermediate magnitudes between the blue and red  LFs. The latter seem  to   stem   from  the different star formation and feedback histories corresponding to different possible merging trees (evolutionary paths)  leading  to the final assembled galaxy; this specific history, driving  the evolution of the   star formation, leads to the different  $M/L$ ratios characterising the different properties  of blue/late  and red/early galaxies. In summary, the peculiar shape of the red LF is mainly driven by the {\\it  nature} of the galaxy merging tree rather than by the {\\it nurture} where the galaxy has grown."
    },
    "0710/0710.1537_arXiv.txt": {
        "abstract": "{The rapidly varying non-thermal X-ray emission observed from Sgr A$^{\\star}$ points to particle acceleration taking place close to the supermassive black hole. The TeV $\\gamma$-ray source HESS\\,J1745$-$290 is coincident with Sgr A$^{\\star}$ and may be closely related to the X-ray emission. Simultaneous X-ray and TeV observations are required to elucidate the relationship between these two objects. Here we report on joint H.E.S.S./Chandra observations in July 2005, during which an X-ray flare was detected. Despite a factor $>10$ increase in the X-ray flux of Sgr~A$^{\\star}$, no evidence is found for an increase in the TeV $\\gamma$-ray flux. We find that an increase of the $\\gamma$-ray flux of a factor 2 or greater can be excluded at a confidence level of 99\\%. This finding disfavours scenarios in which the bulk of the $\\gamma$-ray emission observed is produced close to Sgr A$^{\\star}$. }  \\begin{document}  \\newcommand{\\astar}{Sgr~A$^{\\star}$} ",
        "introduction": "The existence of a supermassive ($3.6\\pm0.3 \\times 10^{6}$ solar mass) black hole at the centre of our galaxy has been inferred using measurements of stellar orbits in the central parsec (see e.g. ~\\cite{GC:Eisenhauer05}). The supermassive black hole (SMBH) is coincident with the faint radio source: Sgr~A$^{\\star}$. The compact nature of Sgr~A$^{\\star}$ has been demonstrated both by direct VLBI measurements~\\cite{GC:Shen05} and by the observation of X-ray and near IR flares with timescales as short as a few minutes (see for example~\\cite{GC:Eckart06,GC:Porquet03}).  Variability on such short timescales limits the emission region (via causality arguments) to within $<10$ Schwarzchild radii of the black hole. X-ray flares from \\astar\\ have reached fluxes of $4\\times10^{35}$ erg s$^{-1}$, two orders of magnitude brighter than the quiescent flux~\\cite{GC:Porquet03, GC:Baganoff03}, and exhibit a range of spectral shapes~\\cite{GC:Porquet03}.  Several models exist for the origin of this variable emission, all of which invoke non-thermal processes close to the event horizon of the central black hole to produce a population of relativistic particles.  Model independent evidence for the existence of ultra-relativistic particles close to Sgr~A$^{\\star}$ can be provided by the observation of TeV $\\gamma$-rays from this source. Indeed, TeV $\\gamma$-ray emission has been detected from the Sgr~A region by several ground-based instruments~\\cite{GC:Whipple04,GC:CANGAROO,HESS:gc04, GC:MAGIC06}. The most precise measurement of this source, HESS\\,J1745$-$290, are those made using the H.E.S.S. telescope array. The centroid of the source is located $7'' \\pm 14_{\\mathrm{stat}}'' \\pm 28_{\\mathrm{sys}}''$ from Sgr~A$^{\\star}$, and has an rms extension of  $<1.2'$~\\cite{HESS:gcprl}.  TeV emission from \\astar\\ is expected in several models of particle acceleration in the environment of the black hole. In some of these scenarios \\cite{Levinson, AhNer1} TeV emission is produced in the immediate vicinity of the SMBH and variability is expected. In alternative scenarios particles are accelerated at \\astar\\ but radiate in within the central $\\sim$10~parsec region~\\cite{AhNer2}, or are accelerated at the termination shock of a wind driven by the SMBH \\cite{AtDer}. However, several additional candidate objects exist for the origin of the observed $\\gamma$-ray emission. The radio centroid of the supernova remnant (SNR) Sgr~A East lies $\\sim 1'$ from Sgr~A$^{\\star}$, only marginally inconsistent with the position of the TeV source given in \\cite{HESS:gcprl}. Shell-type SNR are now well established TeV $\\gamma$-ray sources~\\cite{HESS:rxj1713,HESS:velajnr} and several authors have suggested Sgr A East as the origin of the TeV emission (see for example \\cite{Melia}). However, recent improvements in the statistical and systematic uncertainties of the centroid of HESS\\,J1745$-$290 effectively exclude Sgr~A East as the dominant $\\gamma$-ray source in the region~\\cite{VanEldik}. The recently discovered pulsar wind nebula candidate G\\,359.95-0.04~\\cite{GC:Wang06} lies only 9 arcseconds from \\astar\\ and can plausibly explain the TeV emission~\\cite{GC:Hinton07}. Particle acceleration at stellar wind collision shocks within the central young stellar cluster has also been hypothesised to explain the $\\gamma$-ray source~\\cite{GC:Quataert05}. Finally, an origin of this source in the annihilation of WIMPs in a central dark matter cusp has been extensively discussed~\\cite{GC:Hooper04,GC:Profumo05,HESS:gcprl}.  Given the limited angular resolution of current VHE $\\gamma$-ray telescopes, the most promising tool for identification of the TeV source is the detection of \\emph{correlated variability} between the $\\gamma$-ray and X-ray and/or NIR regimes. A significant increase of the flux of HESS\\,J1745$-$290 simultaneous with a flare in wavebands with sufficient angular resolution to isolate \\astar, would provide an unambiguous identification of the $\\gamma$-ray source. Therefore, whilst not all models for TeV emission from Sgr~A$^{\\star}$ predict variability of the VHE source, coordinated IR/keV/TeV observations can be seen as a key aspect of the ongoing program to understand the nature of this enigmatic source. ",
        "conclusions": "For the first time simultaneous TeV $\\gamma$-ray observations have been presented for a period of X-ray activity of \\astar. The non-detection of an increase in the TeV flux provides an important constraint on scenarios in which the source HESS\\,J1745-290 is associated with the supermassive black hole."
    },
    "0710/0710.4254_arXiv.txt": {
        "abstract": "The Atacama Large Millimeter Array (ALMA) will consist of up to 64 state-of-the-art sub-mm telescopes, subject to stringent performance specifications which will push the boundaries of the technology, and makes testing of antenna performance a likewise challenging task. Two antenna prototypes were evaluated at the ALMA Test Facility at the Very Large Array site in New Mexico, USA. The dynamic behaviour of the antennas under operational conditions was investigated with the help of an accelerometer system capable of measuring rigid body motion of the elevation structure of the antenna, as well as a few low-order deformation modes, resulting in dynamic performance numbers for pointing stability, reflector surface stability, path length stability, and structure flexure. Special emphasis was given to wind effects, one of the major factors affecting performance on timescales of seconds to tens of minutes.  Though the accelerometers could not directly measure antenna performance on timescales longer than 10 seconds, it was possible to use them to derive antenna properties which allowed extrapolation of the wind-affected performance to timescales of 15 minutes and longer. This paper describes the accelerometer system, its capabilities and limitations, and presents the dynamic performance results of the two prototype antennas investigated.  In addition to verification of the performance requirements, we investigated the vibration environment on the antennas, relevant for vibration-sensitive equipment for use on the ALMA antennas, the lowest eigenfrequencies for the antennas, and the sensitivity to vibration generated by similar antennas operating nearby.  This work shows that seismic accelerometers can successfully be used for characterisation of antenna dynamics, in particular when complemented with simultaneous wind measurements and measurements for static performance. The versatility of the accelerometers also makes them a valuable tool for troubleshooting of unforeseen antenna features. ",
        "introduction": "\\label{intro}   The enormous growth of radio astronomy in the millimeter and submillimeter wavelength regime (frequencies from 100 - 1000 GHz) over the last 25 years has been made possible both by the emergence of sensitive detectors like SIS-diodes and HEB-devices and the construction of ever larger and more accurate radio telescopes. Reflector antennas for this wavelength region have been built from relatively small (up to 15 m diameter), extremely accurate (reflector accuracy \\mbox{15 - 25 $\\mu$m}) submillimeter telescopes to larger (20 - 45 m) millimeter antennas with surface accuracy from \\mbox{75 - 150 $\\mu$m}.  Current projects in this area are the 50 m diameter Large Millimeter Telescope (LMT) under construction in Mexico \\cite{Kaercher2000}, which is aiming to reach \\mbox{75 $\\mu$m} rms surface and 1 arcsec pointing accuracy and the Atacama Large Millimeter Array (ALMA). ALMA is a global collaboration of North America (USA and Canada) and Europe (European Southern Observatory) with contributions from Japan to build a powerful millimeter wavelength aperture synthesis array at a 5000 m high plateau in Northern Chile. The instrument will consist of up to 64 high accuracy Cassegrain reflector antennas of 12 m diameter with a surface rms accuracy of \\mbox{20 $\\mu$m} and a pointing and tracking accuracy and stability of 0.6 arcseconds, all under the severe operational conditions of the high site. This telescope will operate over the entire frequency range from 30 to 950 GHz.  These specifications are among the most severe ever realised in radio telescopes and force the designer and manufacturer to push the boundaries of the technology. At the same time, it is becoming increasingly difficult for the contractor and the customer to quantitatively and reliably evaluate the performance characteristics of these instruments. At the longer wavelengths radio astronomers have developed and used methods of antenna evaluation which are based on the use of strong cosmic radio sources and astronomical observing techniques \\cite{Baars1973}. These measurements are only of limited use at the short millimeter wavelengths at frequencies above 100 GHz. The number of suitable cosmic test sources is severely limited because of their intrinsic weakness and the sensitivity limitations of the relatively small antennas.  \\pubidadjcol  For the ALMA Project the partners decided early to obtain two prototype antennas from different design and fabrication groups to increase the chance of achieving the desired performance. The prototype antennas tested here, one designed and constructed by VertexRSI, the other by a consortium of Alcatel and European Industrial Engineering, hereafter referred to as AEC, are similar in overall design and built to meet identical requirements, though significant differences exist in the approaches taken to meet these requirements. The antennas, together with a third one from Japan, are located at the ALMA Test Facility (ATF), on the site of the Very Large Array (VLA) in New Mexico, USA (Figures~\\ref{fig:vertexrsi} and \\ref{fig:atf}). For more information on the antennas and the evaluation program, see \\cite{Mangum2006}.  \\begin{figure} \\resizebox{\\hsize}{!}{ \\includegraphics[scale=0.48,angle=0]{Vertex-astroph.jpg}} \\caption{The VertexRSI ALMA prototype antenna.} \\label{fig:vertexrsi} \\end{figure}  \\begin{figure} \\resizebox{\\hsize}{!}{ \\includegraphics[scale=0.65,angle=0]{DSC01896-astroph.jpg}} \\caption{The ALMA Test Facility, with the AEC antenna (left), VertexRSI antenna (middle) and Mitsubishi antenna (right).} \\label{fig:atf} \\end{figure}    An international team of radio astronomers was formed to subject the antennas to an extensive evaluation program. In preparing their tasks this group determined that additional measurements and test methods and instruments beyond the usual astronomical testing would be needed to check the very stringent specification of the antennas. A particularly important, but difficult to measure quantity is the accuracy with which the antenna can be pointed at arbitrary positions on the sky and the stability with which such a pointing can be maintained under the influence of variations in temperature and wind forces. Given our need to check these parameters independently of the availability of celestial radio sources, we looked into the possibility of using accelerometers on the antenna structure to establish its dynamical behaviour.  The use of seismic accelerometers for performance characterization has been successfully demonstrated on optical telescopes \\cite{Smith2004} and mm-antennas \\cite{Ukita2002}, \\cite{Ukita2004}. Using a set of 10 seismic accelerometers, installed on the antenna back-up structure (BUS), subreflector support structure (apex), and receiver cabin, we have measured accelerations allowing determination of rigid body motion of the elevation structure, and a few low-order distortions of the BUS.  The nature of the accelerometers used in this work limits accurate displacement measurement to time scales of at most 10 seconds or frequencies of at least 0.1 Hz. Since this is well below the lowest eigenfrequencies of the antennas, this is sufficient to determine dynamic antenna behaviour. For the frequency range covered accurately by the accelerometers, approximately 0.1 to 30 Hz, it is possible to check the following antenna specifications:  \\begin{enumerate} \\item Variations in surface shape, restricted to large scale effects like focal length and astigmatism \\item Variations in pointing in elevation and cross-elevation direction \\item Translation of apex structure with respect to the BUS \\item Variations in path length along the boresight direction \\end{enumerate}  Using detailed long term wind studies, and wind measurements simultaneous with accelerometer measurements, antenna performance can be extrapolated to longer time scales under the assumption that wind effects dominate antenna performance on these time scales.  Antenna performance should be met for all modes in which the antenna will be used to perform astronomical observations. For this paper, we have considered  \\begin{enumerate} \\item pointing where the antenna is commanded to remain fixed at an azimuth and elevation position, \\item sidereal tracking, where azimuth and elevation are updated continuously, \\item fast switching mode, where the antenna is switched quickly between two neighbouring points, \\item interferometric mosaicking, in which areas of sky are mapped at slow speed (0.05 deg/s), \\item on-the-fly mapping (OTF), in which large areas of sky are mapped at high scan speed (0.5 deg/s). \\end{enumerate} ",
        "conclusions": "\\subsection{Accelerometers}  Seismic accelerometers mounted on the back-up structures of large reflector antennas are capable of characterising all relevant BUS rigid body motions, and a few of its low-order deformation modes. Independently of any external sources to the antenna, and even for antennas without receivers, it is possible to derive the dynamical behaviour of the performance parameters of an antenna, such as pointing accuracy, primary reflector surface stability, and path length stability. The accuracy at which performance parameters can be measured is a function of frequency, and typically at the sub-$\\mu$m and sub-arcsecond level for frequencies above 0.1 Hz.  In combination with wind measurements, antenna performance for windy conditions can safely be extrapolated to frequencies well below the limit of 0.1 Hz imposed by noise in the accelerometer data. The validity of this extrapolation has been demonstrated with the use of an optical pointing telescope, not limited by the low-frequency noise.  In addition to performance testing, the robust accelerometer system has proven to be well suited for troubleshooting of unexpected antenna behaviour, such as large and off-axis apex structure rotation for the AEC antenna, and servo-tuning issues for the VertexRSI antenna.  \\subsection{Wind-Driven Performance}  Performance of the antennas for windy conditions was a major design driver. In spite of the very different wind conditions at the antenna test site and the ALMA site, it was possible to accurately predict performance for the ALMA site based on measurements performed at the ATF. The key to this extrapolation is careful characterisation of the wind at the sites, in particular at the ATF site. Wind-driven performance as measured on the antennas was combined with the wind characteristics during the time of measurement, which allowed calculation of wind-independent antenna properties, allowing calculation of antenna performance for any wind condition.  As a spin-off of the investigations, antenna wake turbulence at a typical distance of 30 - 50 m was determined; which is useful information for calculation of antenna performance in the compact configuration.   \\subsection{Antenna Performance}  Antenna performance as measured with the accelerometers complements other performance measurements, such as  optical and radio pointing, and holography. The full performance of the antennas could not be determined with any of the individual methods, but a combination of them gives confidence in the completeness of the test results.  During design, antenna performance was calculated from the sum of individual contributions to the total error budget. The performance numbers presented in this paper are best interpreted in the context of the corresponding error budget contributions used for antenna design. Table~\\ref{tab:pointing} gives the measured performance for each antenna, and the calculated error budget entry as taken from the design documentation where available \\cite{Mangum2006}. %  \\begin{table*} \\centering \\caption{Wind Pointing} \\begin{tabular}{lll} & VertexRSI & AEC \\\\ \\\\ Pointing accuracy (wind only) & 0.81 arcsec& 0.45 arcsec\\\\ Pointing accuracy (wind only) error budget & 0.035 arcsec& 0.35 arcsec\\\\ Pointing accuracy (wind + tracking) & 0.94 arcsec& 0.50 arcsec\\\\ Pointing accuracy offset pointing requirement& 0.6 arcsec& 0.6 arcsec\\\\ \\\\ Primary reflector surface stability, wind effects (astigmatism + % AAM, at edge of BUS) & 5.3 + 2.2 $\\mu$m& 6 + 5 $\\mu$m\\\\ Primary reflector surface stability, wind effects error budget& 8.4  $\\mu$m& 2.1  $\\mu$m\\\\ Primary reflector surface stability, overall requirement& 25  $\\mu$m& 25  $\\mu$m\\\\ \\\\ Path length stability, wind effects & 6 $\\mu$m& 6 $\\mu$m\\\\ Path length stability, wind effects error budget& 7.6 $\\mu$m& 3.5 $\\mu$m\\\\ Path length stability, requirement total non-repeatable residual delay & 15 $\\mu$m& 15 $\\mu$m\\\\  \\\\ Structure flexure cross-elevation & 2.1 arcsec/(deg/s$^2$)& 1.6 arcsec/(deg/s$^2$)\\\\ Structure flexure elevation & 2.8 arcsec/(deg/s$^2$)& 2.1 arcsec/(deg/s$^2$)\\\\ \\\\ Lowest eigenfrequencies & 5.57 Hz & 6.8 Hz \\\\  \\end{tabular} \\label{tab:pointing} \\end{table*}"
    },
    "0710/0710.2191_arXiv.txt": {
        "abstract": "We have studied non-axisymmetric standing accretion shock instability, or SASI, by 3D hydrodynamical simulations. This is an extention of our previous study on axisymmetric SASI. We have prepared a spherically symmetric and steady accretion flow through a standing shock wave onto a proto-neutron star, taking into account a realistic equation of state and neutrino heating and cooling. This unperturbed model is supposed to represent approximately the typical post-bounce phase of core-collapse supernovae. We then have added a small perturbation ($\\sim$1\\%) to the radial velocity and computed the ensuing evolutions. Not only axisymmetric but non-axisymmetric perturbations have been also imposed. We have applied mode analysis to the non-spherical deformation of the shock surface, using the spherical harmonics. We have found that (1) the growth rates of SASI are degenerate with respect to the azimuthal index $m$ of the spherical harmonics $Y_{l}^{m}$, just as expected for a spherically symmetric background, (2) nonlinear mode couplings produce only $m=0$ modes for the axisymmetric perturbations, whereas $m \\neq 0$ modes are also generated in the non-axisymmetric cases according to the selection rule for the quadratic couplings, (3) the nonlinear saturation level of each mode is lower in general for 3D than for 2D because a larger number of modes are contributing to turbulence in 3D, (4) low $l$ modes are dominant in the nonlinear phase, (5) the equi-partition is nearly established among different $m$ modes in the nonlinear phase, (6) the spectra with respect to $l$ obey power laws with a slope slightly steeper for 3D, and (7) although these features are common to the models with and without a shock revival at the end of simulation, the dominance of low $l$ modes is more remarkable in the models with a shock revival. ",
        "introduction": "Many efforts have been made for the multi-dimensional modeling of core-collapse supernovae (see \\cite{tomas,kotake_rev} for reviews), urged by accumulated observational evidences that core-collapse supernovae are globally aspherical commonly \\citep{wang96,wang01,wang02}. Various mechanisms to produce the asymmetry have been discussed so far: convection (e.g., \\citet{herant_94,burrows_95,jankamueller96}), magnetic field and rapid rotation (see, e.g., \\cite{kotake_rev} for collective references), standing (/stationary/spherical) accretion shock instability, or SASI, \\citep{blondin_03,scheck_04,blondin_05,ohnishi_1,ohnishi_2,fog06}, and g-mode oscillations of protoneutron stars \\citep{burr_new}. Note, however, that most of them have been investigated with two-dimensional (2D) simulations.  Recently a 3D study on SASI was reported by \\citet{blondin_nat}. In 2D, the shock deformation by SASI is described with the Legendre polynomials $P_{l}(\\theta)$, or the spherical harmonics $Y_{l}^{m}(\\theta, \\phi)$ with $m=0$. Various numerical simulations have demonstrated unequivocally that the $l=1$ mode is dominant and a bipolar sloshing of the standing shock wave occurs with pulsational strong expansions and contractions along the symmetry axis \\citep{blondin_03,scheck_04,blondin_05,ohnishi_1,ohnishi_2}. In 3D, on the other hand, \\citet{blondin_nat} observed the dominance of a non-axisymmetric mode with $l =1$, $m=1$, which produces a single-armed spiral in the later nonlinear phase. They claimed that this \"spiral SASI\" generates a rotational flow in the accretion (see also \\cite{blondin_shaw} for 2D computations in the equatorial plane) and that it may be an origin of pulsar spins.  There are many questions on 3D SASI remaining to be answered yet, however. For example, we are interested in how the growth of SASI differs between 3D and 2D among other things. In particular, the change in the saturation properties should be made clear. This will be the focus of this paper. Another issue will be the generation of rotation in the accretion flow by SASI~\\citep{blondin_nat}. Its efficiency and possible correlation with the net linear momentum should be studied more in detail and will be the subject of our forthcoming paper~\\citep{iwa07}. Incidentally, it is noted that the neutrino heating and cooling were entirely ignored and the flow was  assumed to be isentropic in \\citet{blondin_nat}, but the implementation of these physics is helpful for considering the implications not only for the shock revival but also for the nucleosynthesis~\\citep{kifo}.  In this paper, we have performed 3D hydrodynamic simulations, employing a realistic equation of state \\citep{shen98} and taking into account the heating and cooling of matter via neutrino emissions and absorptions on nucleons as done in our 2D studies \\citep{ohnishi_1,ohnishi_2}. The inclusion of the neutrino heating allows us to discuss how the critical luminosity for SASI-triggered explosion could be changed in 3D from those in 2D. To answer the questions we raised above, we have varied the initial perturbations as well as the neutrino luminosity and compared the growth of SASI between 2D and 3D in detail, conducting a mode analysis not only for the linear phase but also for the nonlinear saturation phase.  The plan of this paper is as follows: In Section \\ref{sec2}, we describe the models and numerical methods. We show the main numerical results in Section \\ref{sec3}. We conclude the paper in Section \\ref{sec4}. ",
        "conclusions": "}  In this paper we have studied the non-axisymmetric SASI by 3D hydrodynamical simulations, taking into account the realistic EOS and neutrino-heating and cooling. We have added various non-axisymmetric perturbations to spherically symmetric steady flows that accrete through a standing shock wave onto a proto-neutron star, being irradiated by neutrinos emitted from the proto-neutron star. The mode analysis has been done for the deformation of the shock surface by the spherical harmonics expansion. After confirming that our 3D code is able to reproduce for the axisymmetric perturbations the previous results on 2D SASI that we obtained in \\citet{ohnishi_1}, we have done genuinely 3D simulations and found the followings.  First, the model of the initial perturbation with $(l, |m|)=(1,0)$ and $(1,1)$ have demonstrated that the non-axisymmetric SASI proceeds much in the same manner as the axisymmetric SASI: the linear phase, in which the initial perturbation grows exponentially, lasts for about $100$ms and is followed by the nonlinear phase, in which various modes are produced by nonlinear mode couplings and their amplitudes are saturated. It has been found that the critical neutrino luminosity, above which the shock revival occurs, is not much different between 2D and 3D. For the neutrino luminosities lower than the critical value, the SASI is settled to a quasi-steady turbulence. We have found that the saturation level of each mode in the non-axisymmetric SASI is lower in general than that of the axisymmetric counterpart. This is mainly due to the fact that the number of the modes contributing to the turbulence is larger for the non-axisymmetric case. The sequence of the mode generation, on the other hand, strongly suggests that the nonlinear mode coupling is chiefly of quadratic nature.  Second, the simulations with the random multi-mode perturbations being imposed initially have shown that the dynamics in the linear phase is essentially a superimposition of those of single-modes. Toward the end of the linear phase, high entropy blobs are generated continuously and grow, starting the nonlinear phase. We have observed that these blobs repeat expansions and contrations, being merged and splitted from time to time. In the non-explosion models, these nonlinear dynamics lead to the saturation of mode amplitudes and the quasi-steady turbulence. For the explosion models, on the other hand, the production of blobs proceeds much more quickly, followed by the oligarchic evolution, with a relatively small number of great blobs absorbing smaller ones, and as a result the shock radius increases almost monotonically. The spectral analysis has clearly demonstrated that low $l$ modes are predominant in the nonlinear phase just as in the axisymmetric case. We have also shown that the equi-partition is nearly established among different $m$ modes in the quasi-steady turbulence and that the spectrum summed over $m$ in the non-axisymmetric case is quite similar to the axisymmetric counterpart. This implies that the larger number of modes in the non-axisymmetric case is the main reason why the amplitude of each mode is smaller in 3D than in 2D. The power spectrum is approximated by the power law for $l \\gtrsim 10$. Although the slope is slightly steeper for the non-axisymmetric models, whether the difference is significant or not is unknown at present.  We have seen in the explosion models, on the other hand, that the oscillation period of each mode becomes longer in the late nonlinear phase as the shock radius gets larger and the advection-acoustic cycle becomes longer. What is more interesting is the fact that in this late phase the dominance of low $l$ modes becomes even more remarkable. Although this may be related with the volume-filling thermal convection, the mechanism is yet to be revealed. This spectral evolution leads to the global deformation of the shock surface to an ellipsoidal shape, whose major axis is not necessarily aligned with the symmetry axis, however.  Finally, we have presented the neutrino heating in the 3D SASI. It has been seen that the volume-integrated heating rate is affected mainly in the nonlinear phase. The comparison with the spherically symmetric counterparts has confirmed that the SASI enhances the neutrino heating also in the non-axisymmetric case. Although the critical neutrino luminosity in the non-axisymmetric SASI is not much changed from that for the axisymmetric case, the spatial distribution of neutrino heating is different in the non-linear phase. Relatively narrow regions surrounding high entropy blobs are efficiently heated for the non-axisymmetric case, while wider regions near the symmetry axis are heated strongly in accord with the sloshing motion of shock wave along the symmetry axis for the axisymmetric case. For the non-explosion models, the high entropy blobs produced by neutrino heatings occupy the inside of shock wave, repeating expansion and contraction and being splitted and merged intermittently, but the flow is finally settled to a quasi-static state. For the explosion models, on the other hand, the high entropy blobs are generated much more quickly and extend the heating region, pushing the shock outwards and narrowing the cooling region near the proto-neutron star.  In the present study we have not observed a persistent segregation of angular momentum in the accretion flow such as found by \\citet{blondin_nat, blondin_shaw}) although instantaneous spiral flows are frequently seen. As discussed above, the equi-partition is nearly established among different $m$ modes in our models. It should be emphasized here again, however, that we have added the initial perturbations only to the radial velocity in this study and, as a result, the modes with $m=\\pm1$ are equally contributing. We will defer the analysis of the models with the initial non-axisymmetric perturbations being added also to the azimuthal component of velocity to our forthcoming paper~\\citep{iwa07}, in which we will also discuss a possible correlation between the kick velocity and spin of neutron stars if they are produced by the 3D SASI indeed. A lot of issues on SASI are still remaining to be studied. In particular, the influences of rotation and magnetic field are among the top priorities and will be addressed elsewhere in the near future."
    },
    "0710/0710.0477_arXiv.txt": {
        "abstract": "We reexamined the gravitational time delay of light, allowing for various models of modified gravity. We clarify the dependence of the time delay (and induced frequency shift) on modified gravity models and investigate how to distinguish those models, when light propagates in static spherically symmetric spacetimes. Thus experiments by radio signal from spacecrafts at very different distances from Sun and future space-borne laser interferometric detectors could be a probe of modified gravity in the solar system. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.5778_arXiv.txt": {
        "abstract": "We validate the performance and accuracy of the current SEGUE (Sloan Extension for Galactic Understanding and Exploration) Stellar Parameter Pipeline (SSPP), which determines stellar atmospheric parameters (effective temperature, surface gravity, and metallicity) by comparing derived overall metallicities and radial velocities from selected likely members of three globular clusters (M~13, M~15, and M~2) and two open clusters (NGC~2420 and M~67) to the literature values. Spectroscopic and photometric data obtained during the course of the original Sloan Digital Sky Survey (SDSS-I) and its first extension (SDSS-II/SEGUE) are used to determine stellar radial velocities and atmospheric parameter estimates for stars in these clusters. Based on the scatter in the metallicities derived for the members of each cluster, we quantify the typical uncertainty of the SSPP values, $\\sigma (\\rm [Fe/H])$ = 0.13 dex for stars in the range of 4500 K $\\le T_{\\rm eff} \\le 7500$ K and $2.0 \\le \\log g \\le 5.0$, at least over the metallicity interval spanned by the clusters studied ($-2.3 \\le {\\rm [Fe/H]} < 0 $). The surface gravities and effective temperatures derived by the SSPP are also compared with those estimated from the comparison of the color-magnitude diagrams with stellar evolution models; we find satisfactory agreement. At present, the SSPP underestimates [Fe/H] for near-solar-metallicity stars, represented by members of M~67 in this study, by $\\sim$ 0.3 dex. ",
        "introduction": "The Sloan Extension for Galactic Understanding and Exploration (SEGUE) is one of three key projects (LEGACY, SUPERNOVA SURVEY, and SEGUE) in the current extension of the Sloan Digital Sky Survey, known collectively as SDSS-II. The SEGUE program is in the process of obtaining $ugriz$ imaging of some 3500 square degrees of sky outside of the SDSS-I footprint (Fukugita et al. 1996; Gunn et al. 1998, 2006; York et al. 2000; Stoughton et al. 2002; Abazajian et al. 2003, 2004, 2005; Pier et al. 2003), with special attention being given to scans of lower Galactic latitudes ($|b|$ $<$ 35$^{\\circ}$) in order to better probe the disk/halo interface of the Milky Way. SEGUE is also obtaining $R$ $\\simeq$ 2000 spectroscopy over the wavelength range 3800 $-$ 9200\\,{\\AA} for some 250,000 stars in 200 selected areas over the sky available from Apache Point, New Mexico.  The SEGUE Stellar Parameter Pipeline (hereafter, SSPP) processes the wavelength- and flux-calibrated spectra generated by the standard SDSS spectroscopic reduction pipeline (Stoughton et al. 2002), obtains equivalent widths and/or line indices for 77 atomic or molecular absorption lines, and estimates $T_{\\rm eff}$, log $g$, and [Fe/H] through the application of a number of approaches. The current techniques employed by the SSPP include a minimum distance method (Allende Prieto et al. 2006), neural network analysis (Bailer-Jones 2000; Willemsen et al. 2005; Re Fiorentin et al. 2007), auto-correlation analysis (Beers et al. 1999), and a variety of line index calculations based on previous calibrations with respect to known standard stars (Beers et al. 1999; Cenarro et al. 2001a,b; Morrison et al. 2003). The SSPP employs five different methods for estimation of $T_{\\rm eff}$, eight for estimation of log $g$, and nine for estimation of [Fe/H]. Details of the methods used are discussed in detail by Lee et al. (2007a, hereafter Paper I). The use of multiple methods allows for empirical determinations of the internal errors for each parameter, based on the range of reported values -- typical internal errors for stars in the temperature range $4500$~K $\\le$ $T_{\\rm eff}$ $\\le$ 7500~K are $\\sim$ 73~K, $\\sim$ 0.19 dex, and $\\sim$ 0.10 dex, in $T_{\\rm eff}$, log $g$, and [Fe/H], respectively. Allende Prieto et al. (2007, hereafter Paper III) point out that the internal uncertainties provided by the SSPP underestimate the typical random errors at high signal-to-noise ($S/N$) ratios because most methods in the SSPP make use of similar parameter indicators (e.g., hydrogen lines for effective temperature) and similar atmospheric models. Paper III empirically determines empirically external uncertainties of $\\sim$ 130~K, $\\sim$ 0.21 dex, and $\\sim$ 0.11 dex, for $T_{\\rm eff}$, log $g$, and [Fe/H], respectively, by comparison with high-resolution spectroscopy ($7000 < R < 45,000$) of brighter SDSS-I/SEGUE that have been obtained with 8m$-$10m class telescopes. Somewhat larger errors apply to stars with temperatures near the extremes of the range above. The present study of Galactic globular and open cluster stars tests the SSPP's ability to derive accurate results for stars with a wide range of temperatures and gravities appropriate for metal-poor and near-solar-metallicity stellar populations in the Galaxy, and demonstrates that the derived metallicity scale is identical for dwarfs and giants.  Although the SSPP will continue to evolve in the near future, it has been frozen for now at the version used for obtaining results for stars with suitable data from SDSS Data Release 6 (DR-6; Adelman-McCarthy et al. 2007b). Previous versions of the SSPP have already been used for the analysis of SDSS-I observations. For example, Allende Prieto et al. (2006) report on the application of one of the methods included in the SSPP to some 20,000 F- and G-type stars from SDSS-I DR-3 (Abazajian et al. 2005). Beers et al. (2006) have compiled a list of over 6000 stars with [Fe/H] $< -2.0$ (including several hundred with [Fe/H] $< -3.0$), based on application of the present SSPP to some 200,000 stars from SDSS-I DR-5 (Adelman-McCarthy et al. 2007a). Carollo et al. (2007) reports on an analysis of the kinematics of relatively bright stars from SDSS-I that have been used as calibration objects during the main survey.  In this paper, the second in the SSPP series, we show that estimates of the atmospheric parameters and radial velocities obtained by the SSPP for stars with a reasonable likelihood of membership in previously studied Galactic globular and open clusters are sufficiently accurate to justify the use of the present SSPP parameters for carrying out detailed studies of the halo and thick-disk populations of the Milky Way. In deriving the overall iron abundance for each cluster, we assume it comprises a chemically homogenous population.  In \\S 2, the photometric and spectroscopic data obtained for M~13, M~15, M~2, NGC~2402, and M~67 are described. Section 3 presents the methods used to separate likely cluster members from field stars in the directions toward these clusters. Best estimates of the overall [Fe/H] and radial velocity of each cluster are derived in \\S 4. In \\S 5 we compare the SSPP determinations of $T_{\\rm eff}$ and log $g$ for selected member stars in each cluster with their expected positions on color-magnitude diagrams. A summary and brief conclusions are provided in \\S 6. ",
        "conclusions": "Based on photometric and spectroscopic data reported in SDSS-I and SDSS-II/SEGUE, we have examined estimates of stellar atmospheric parameters and heliocentric radial velocities obtained by the SEGUE Stellar Parameter Pipeline (SSPP) for likely members of three Galactic globular clusters, M~13, M~15, and M~2, and two open clusters, NGC~2420 and M~67, and compared them with those obtained by external estimates for each cluster as a whole.  From the derived scatters in the metallicities and radial velocities obtained for the likely members of each cluster, we quantify the typical external uncertainties of the SSPP-determined values, $\\sigma (\\rm [Fe/H])$ = 0.13 dex, and $\\sigma (\\rm RV)$ = 7.7 km~s$^{-1}$, respectively. These uncertainties apply for stars in the range of 4500 K $\\le T_{\\rm eff} \\le 7500$ K and $2.0 \\le \\log g \\le 5.0$, at least over the metallicity interval spanned by the clusters studied ($-2.3 \\le {\\rm [Fe/H]} < -0.4$). Therefore, the metallicities and radial velocities obtained by the SSPP appear sufficiently accurate to be used for studies of the kinematics and chemistry of the metal-poor and moderately metal-rich stellar populations in the Galaxy. We have also confirmed that $T_{\\rm eff}$ and log $g$ are sufficiently well-determined by the SSPP to distinguish between different luminosity classes through a comparison with theoretical predictions.  A comparison of the analysis of the available high-resolution spectroscopy of SDSS-I/SEGUE stars (Paper III) with the SSPP predictions indicates that the uncertainty in radial velocities adopted by the SSPP is no more than 5 km~s$^{-1}$ (after adjusting for an empirical offset of +7.3 km~s$^{-1}$). The empirically determined precisions in estimated atmospheric parameters are $\\sim$ 130~K for effective temperature, $\\sim$ 0.21 dex for surface gravity, and $\\sim$ 0.11 dex for [Fe/H]. These errors apply to the brightest stars obtained by SDSS-I/SEGUE observations, on the order of $14.0 \\le g \\le 15.5$, and are expected to degrade somewhat for fainter stars. We also found that the SSPP tends to underestimate [Fe/H] for near-solar-metallicity stars (represented by members of M~67 in this study), by $\\sim$ 0.3 dex.  In future papers we will compare the predictions of the SSPP with intermediate-metallicity clusters ([Fe/H] $\\sim$ $-$0.7) and with additional near-solar-metallicity populations, as sampled by metal-rich globular clusters and nearby open clusters. Additionall metal-poor clusters will also be examined. Further refinements in the SSPP, which hopefully will be better able to recover accurate abundances for near-solar-metallicity stars, are anticipated."
    },
    "0710/0710.2702_arXiv.txt": {
        "abstract": "The effect of viscosity and of converging flows on the formation of blobs in the slow solar wind is analysed by means of resistive MHD simulations. The regions above coronal streamers where blobs are formed \\citep{sheeley} are simulated using a model previously proposed by \\citet{einaudijgr}. The result of our investigation is two-fold. First, we demonstrate a new mechanism for enhanced momentum transfer between a forming blob and the fast solar wind surrounding it. The effect is caused by the longer range of the electric field caused by the tearing instability forming the blob. The electric field reaches into the fast solar wind and interacts with it, causing a viscous drag that is global in nature rather than local across fluid layers as it is the case in normal uncharged fluids (like water). Second, the presence of a magnetic cusp at the tip of a coronal helmet streamer causes a converging of the flows on the two sides of the streamer and a direct push of the forming island by the fast solar wind, resulting in a more efficient momentum exchange. ",
        "introduction": "The understanding of slow solar wind genesis has progressed considerably in the last few years following the observational discovery of plasma inhomogeneities (called {\\em blobs}) formed and expelled from regions above coronal streamers~\\citep{sheeley, wang}. The observational evidence has been provided by the Large Angle and Spectrometric Coronagraph (LASCO) instrument on the Solar and Heliospheric Observatory (SOHO).   Approximately 4 blobs per day were observed in a relatively quiet period for the solar corona (year 1997, near solar minimum). After formation the blobs are carried away with the ambient slow solar wind. Generally, the plasma blobs move radially outward with their speed increasing from 0-250 km s$^{-1}$ in the region 2-6 ${\\rm R}_\\odot$ (C2 field of view) to 200-450 km s$^{-1}$ in the region 3.7-30 ${\\rm R}_\\odot$ (outer portion of the C3 field of view).   Models of the formation of blobs have focused on mechanisms involving   reconnection between open field lines and closed field lines \\citep{einaudijgr, einaudi, endeve,fisk, vanaalst, wang, wu}. Through field line reconnection, plasma from the denser regions within the helmet streamer is liberated and becomes tied to open field lines and forms the slow solar wind \\citep{einaudi}.  The present work is particularly centred around the model proposed by \\citet{einaudijgr}.  In the region downstream of the cusp in the streamer belt, \\citet{einaudijgr}  approximate the configuration as a 1D  reversed field configuration with a 1D wake velocity profile. Under these assumptions, they can reproduce both the acceleration of the slow solar wind and the formation of blobs, formulating a complete picture of the genesis of the slow solar wind, in accordance with the observations \\citep{einaudijgr}. The scenario has recently been further extended~\\citep{rappazzo} to include the so called {\\em melon seed effect} due to the diamagnetic force caused by the overall magnetic field radial gradients in the spherical geometry of the Sun~\\citep{schmidt} and to include the effects of the cusp magnetic configurations atop the closed field lines in coronal streamers~\\citep{lapenta-wind}.  The fundamental conclusion of the model by \\citet{einaudijgr} and in successive refinements is that a blob is formed by the tearing instability and that by virtue of the interaction with the fast solar wind the blob is accelerated and ejected, carrying with it the plasma that forms the slow solar wind.  The fundamental question posed by the present work is: ``what is the basic physical process for the plasma acceleration''? We consider specifically two possibilities. First, that the momentum transfer causing the blob acceleration is due to viscous drag. Second, that the acceleration is due to the insurgence of the tearing instability and its non-linear effects. The acute reader will not find the suspense spoiled by learning that, as usual, Nature turns out not to be black or white. We discover that both mechanisms are present in a interplay mediated by the electric field. Furthermore, the presence of the converging flows at the cusp of the coronal streamer causes a direct momentum transfer by the action of the plasma  flowing against the forming island and pushing  it upward from the solar surface.  The paper is organised as follows. Section 2 summarises the initial configurations considered distinguishing the 1D initial configuration proposed by \\citet{einaudijgr} and its 2D extension to include the presence of a magnetic cusp~\\citep{lapenta-wind}. Section 3 very briefly provides the methodology for our investigation: the resistive MHD model as implemented in the FLIP3D-MHD code~\\citep{brackbill}. Section 4 and 5 present the non-linear evolution, respectively, of the case of the initial 1D reversed field and of the 2D configuration with a magnetic cusp. Conclusions are reached in Section 6. ",
        "conclusions": ""
    },
    "0710/0710.1584_arXiv.txt": {
        "abstract": "Recent observations with atmospheric Cherenkov telescope systems such as \\HESS\\ and MAGIC have revealed a large number of new sources of very-high-energy (VHE) $\\gamma$-rays from 100~GeV -- 100~TeV, mostly concentrated along the Galactic plane. At lower energies (100 MeV -- 10 GeV) the satellite-based instrument EGRET revealed a population of sources clustering along the Galactic Plane. Given their adjacent energy bands a systematic correlation study between the two source catalogues seems appropriate. Here, the populations of Galactic sources in both energy domains are characterised on observational as well as on phenomenological grounds. Surprisingly few common sources are found in terms of positional coincidence and spectral consistency. These common sources and their potential counterparts and emission mechanisms will be discussed in detail. In cases of detection only in one energy band, for the first time consistent upper limits in the other energy band have been derived. The EGRET upper limits are rather unconstraining due to the sensitivity mismatch to current VHE instruments. The VHE upper limits put strong constraints on simple power-law extrapolation of several of the EGRET spectra and thus strongly suggest cutoffs in the unexplored energy range from 10~GeV -- 100~GeV. Physical reasons for the existence of cutoffs and for differences in the source population at GeV and TeV energies will be discussed. Finally, predictions will be derived for common GeV--TeV sources for the upcoming GLAST mission bridging for the first time the energy gap between current GeV and TeV instruments. ",
        "introduction": "In recent years the knowledge of the Galactic VHE $\\gamma$-ray sky above 100~GeV has been greatly improved through the detection and subsequent study of many sources, mostly by means of ground-based Imaging Atmospheric Cherenkov telescope systems such as the High Energy Stereoscopic System (\\HESS) or the Major Atmospheric Gamma-ray Imaging Cherenkov Observatory (MAGIC). Currently known Galactic VHE $\\gamma$-ray emitters include shell-type Supernova remnants (SNRs)~\\citep{HESS1713_II, HESS1713_III, HESSVelaJr}, Pulsar Wind Nebulae (PWNe)~\\citep{HESSMSH, HESS1825II, HESSKooka}, $\\gamma$-ray binaries~\\citep{HESSLS5039II, MAGICLSI}, Molecular clouds~\\citep{HESSGCDiffuse} and possibly also clusters of massive stars~\\citep{HESSWesterlund}. These various source classes were discovered both in pointed observations using \\HESS\\ and MAGIC as well as in a systematic survey of the inner Galaxy performed with the \\HESS\\ instrument. The highest energy photons detected from these source classes reach $\\sim 100$~TeV~\\citep{HESS1713_III}, currently representing the end of the observable electromagnetic spectrum for astrophysical objects. It is natural to investigate the relationship of these TeV sources to sources at lower energies as will be done in this work focusing on Galactic sources. The closest energy band for which data exist is that studied by the Energetic Gamma Ray Experiment Telescope (EGRET) aboard the Compton Gamma-Ray Observatory with an energetic coverage from 100~MeV -- 10~GeV~\\citep{EGRET}.  The GeV sky has a distinctively different overall appearance compared to TeV energies. In particular focusing on our Galaxy, the most prominent feature of the GeV sky is the dominant diffuse emission from cosmic ray (CR) interactions in the Galaxy, while the TeV sky due to the steeply falling energy spectrum of the diffuse component is dominated by individual sources. However, several prominent $\\gamma$-ray sources are known to emit at both GeV and at TeV energies, the Crab Nebula being the most prominent example~\\citep{WhippleCrab,EGRETCrab, HEGRACrab, HESSCrab, MAGICCrab}.  In this paper the relationship between Galactic EGRET and VHE $\\gamma$-ray sources will be assessed in a systematic way. For cases with a positional coincidence between a VHE and an EGRET source (in the following called ``coincident sources'') all currently known Galactic objects will be considered. For cases in which a source is detected only in one band -- the ``non-coincident sources'' -- we focus on the region covered by the \\HESS\\ Galactic plane survey (GPS) during 2004 and 2005~\\citep{HESSScan, HESSScanII} (Galactic longitude $\\pm 30^{\\circ}$, Galactic latitude $\\pm 3^{\\circ}$) so that a statistical assessment of the ``non-connection'' can be made. EGRET was unable to perform detailed studies of the $\\gamma$-ray sky above 10~GeV, partly due to back-splash of secondary particles produced by high-energy $\\gamma$-rays causing a self-veto in the monolithic anti-coincidence detector used to reject charged particles. The upcoming Gamma Ray Large Area Space Telescope (GLAST) Large Area Telescope (LAT) will not be strongly affected by this effect since the anti-coincidence shield was designed in a segmented fashion~\\citep{GLASTACD}. Moreover, the effective area of the GLAST-LAT will be roughly an order of magnitude larger then that of EGRET. The GLAST-LAT mission will therefore for the first time fully bridge the gap between the energy range of EGRET and current VHE instruments. Part of the study presented here can be seen as preparatory work for GLAST-LAT studies of sources in the largely unexplored energy band between 10 and 100~GeV.  \\begin{figure}[ht] \\begin{center} \\includegraphics[width=0.9\\textwidth]{f1.eps} \\end{center} \\caption{{\\bf{Left:}} Integral sensitivities for current, past and future $\\gamma$-ray instruments (5-$\\sigma$ sensitivity for $E>E_0$ multiplied with $E_0$ assuming a spectrum of $E^{-2}$). The solid lines show the nominal instrument sensitivities (for a typical observation time as specified below), the dashed curves show the actual sensitivities for the Inner Galaxy as appropriate for this work. INTEGRAL's (IBIS/ISGRI) sensitivity curve (solid green) shows the sensitivity for an observation time of $10^5$s, a typical value in the Inner Galaxy. The EGRET curves (brown) are shown for the whole lifetime of the mission (periods 1--9) for the Galactic anti-centre (solid) which received the largest exposure time and has a lower level of diffuse $\\gamma$-ray emission than the Inner Galaxy and for the position of RX\\,J1713.7--3946 (dashed), a typical position in the Inner Galaxy dominated by diffuse $\\gamma$-ray background emission. The GLAST curves in red (taken from http://www-glast.slac.stanford.edu/software/IS/glast\\_lat\\_performance.html) show the 1-year sky-survey sensitivity for the Galactic North pole -- again a position with low diffuse emission (solid), and for the position of RX\\,J1713.7--3946 (dashed). The H.E.S.S.\\ curves (blue) are shown for a 50-hour pointed observation of a point-like source (solid) and for a 5-hour observation of a somewhat extended source as is typical for the Galactic Plane survey (angular cut of $\\sqrt{0.05}$). The MAGIC curve (light blue) represents a 50-hour observation of a point-source. {\\bf{Right:}} Energy-dependence of the angular resolution for selected $\\gamma$-ray instruments expressed by the 68\\%-containment radius of the point-spread function (PSF). As can be seen, the angular resolution of GLAST becomes comparable with current VHE instruments at high energies, whilst at the lower energy end GLAST and EGRET have comparable resolutions.}\\label{fig::Sensitivity} \\end{figure}  From the 2004 and 2005 \\HESS\\ GPS 22 VHE $\\gamma$-ray sources have been reported in the Inner Galaxy. The third EGRET catalogue~\\citep{EGRET} represents the companion to the VHE source catalogue above an energy threshold of 100 MeV (with peak sensitivity between 150 and 400 MeV, depending on the $\\gamma$-ray source spectrum). It lists 271 sources, 17 of which are located within the \\HESS\\ GPS region. Whilst the EGRET range currently represents the nearest energy band to VHE $\\gamma$-rays, for the very few EGRET sources detected all the way up to $\\sim 10$~GeV, there is still an unexplored energy band of roughly one decade before the VHE $\\gamma$-ray energy range begins at $\\sim 100$~GeV (it should be noted that EGRET does have some sensitivity beyond 10~GeV: \\citet{EGRET10GeV} reported the detection of $\\sim$ 1500 photons above that energy with 187 of these photons being found within $1^{\\circ}$ of a source listed in the third EGRET catalogue). Comparing the instrumental parameters of VHE instruments and EGRET there is a clear mismatch both in angular resolution and in sensitivity as can be seen in Figure~\\ref{fig::Sensitivity}. In a $\\sim 5$ hour observation (as a typical value in the GPS region) \\HESS\\ is a factor of $\\sim 50-80$ more sensitive (in terms of energy flux $E^2 dN/dE$) than EGRET above 1~GeV in the Galactic Plane for the exposure accumulated between 1991 and 1995 (corresponding to the third EGRET catalogue). Assuming a similar energy flux output in the two different bands this mismatch implies at first sight that \\HESS\\ sources are not likely to be visible in the EGRET data set. Conversely (again under the assumption of equal energy flux output), VHE $\\gamma$-ray instruments should be able to detect the majority of the EGRET sources, as has been suggested in the past. Figure~\\ref{fig::EnergyFluxDistribution} compares the energy fluxes $\\nu F \\nu$ for EGRET sources and \\HESS\\ sources in the inner Galaxy. Clearly, the EGRET sources do not reach down as low in energy flux as the \\HESS\\ sources, a picture that will change once the GLAST-LAT is in orbit as depicted by the GLAST-LAT sensitivity (dashed line). In reality the \\naive\\ expectation of equal energy flux output in the GeV and TeV band can easily be wrong in Galactic $\\gamma$-ray sources for various reasons: EGRET sources may not emit comparable energy fluxes in the VHE $\\gamma$-ray band but rather exhibit cut-offs or spectral breaks in the energy band between EGRET and \\HESS\\ \\citep[this is certainly the case for pulsed magnetospheric emission from pulsars, see for example][]{HESSPulsar}. Furthermore, \\HESS-like instruments are typically only sensitive to emission on scales smaller than $\\sim 1^{\\circ}$. If any of the EGRET sources are extended beyond $1^{\\circ}$ without significant sub-structure on smaller scales (not precluded given the poor angular resolution of EGRET), current Imaging Cherenkov instruments may not be able to detect them since these sources would completely fill the field of view (FoV) and be removed by typical background subtraction methods ~\\citep[see for example][]{BergeBackground}. Given the upcoming launch of GLAST and the recent \\HESS\\ survey it seems timely to study the relationship between GeV and TeV emitting sources in more detail.  \\begin{figure}[ht] \\begin{center} \\includegraphics[width=0.6\\textwidth]{f2.eps} \\end{center} \\caption{Distribution of integrated energy flux $\\nu F \\nu$ for sources in the Inner Galaxy discussed here. For EGRET the energy flux between 1~GeV and 10~GeV, for the \\HESS\\ sources, the energy flux between 1~TeV and 10~TeV is shown. Also shown is the sensitivity prediction for the GLAST-LAT for a typical location in the Inner Galaxy (l=10, b=0).}\\label{fig::EnergyFluxDistribution} \\end{figure}  Section~\\ref{sec::analysis} describes the data and analysis methods used in this study, section~\\ref{sec::connect} describes the sources detected in both energy bands, and section~\\ref{sec::nonconnect} focuses on sources detected in only one of the two energy regimes. In section~\\ref{sec::interpretation} astrophysical implications of the study are discussed. ",
        "conclusions": "The main results of the study of the relationship between GeV and TeV sources are:  \\begin{enumerate} \\item There are rather few spatially coincident GeV-TeV sources for the considered Galactic region. \\item Those few positional coincident GeV-TeV sources could occur by chance, the chance probability of detecting two coincident sources within the \\HESS\\ GPS region is $\\sim 40$\\%, thus no strong hint for a common GeV/TeV source population is detected. \\item Spectral compatibility (based on a power-law extrapolation) seems present for most of the positionally coincident sources, but again, this is expected to occur by chance (as described in the text) given the sensitivity mismatch and the different energy bands. \\item Dedicated \\HESS\\ limits at the position of the EGRET sources are constraining for a power-law extrapolation from the GeV to the TeV range for several of the EGRET sources, strongly suggesting cutoffs in the energy spectra of these EGRET sources in the unexplored region below 100~GeV. Power-law extrapolation of EGRET spectra seem to be ruled out for most of the EGRET sources investigated in this study. \\item Dedicated EGRET limits at the position of the \\HESS\\ sources are not constraining for a power-law extrapolation from the TeV to the GeV range. This picture will dramatically change once the GLAST-LAT with its improved sensitivity over EGRET is in orbit. \\item Several important mechanisms for cutoffs in the energy spectra of GeV sources have been discussed. There are well motivated physical reasons why the population of GeV and of TeV sources might be distinct. \\item If a source can be detected with both GeV and TeV instruments, the huge energy ``lever arm'' over 5-6 decades in energy will undoubtedly provide stringent constraints on the $\\gamma$-ray emission mechanism in these Galactic particle accelerators. \\end{enumerate}  Summarising, the study presented here shows that the GLAST-LAT will tremendously advance the study of the relationship between GeV and TeV sources by improving the sensitivity over EGRET by an order of magnitude and in particular by bridging the currently uncovered energy range between 10~GeV and 100~GeV."
    },
    "0710/0710.1067_arXiv.txt": {
        "abstract": "% We present an analysis of the magnetic-field fluctuations in the magnetoionic medium in front of the radio galaxy 3C 31 derived from rotation-measure (RM) fits to high-resolution polarization images.  We first show that the Faraday rotation must be due primarily to a foreground medium.  We determine the RM structure functions for different parts of the source and infer that the simplest form for the power spectrum is a power law with a high-frequency cutoff. We also present three-dimensional simulations of RM produced by a tangled magnetic field in the hot plasma surrounding 3C 31, and show that the observed RM distribution is consistent with a spherical plasma distribution in which the radio source has produced a cavity. ",
        "introduction": "Our knowledge of the structure and origin of magnetic fields in elliptical galaxies, groups and clusters is still rudimentary, but Faraday rotation of linearly-polarized radio emission can be used to probe the fields in ionized foreground gas.  Here, we present an analysis of the magnetic-field fluctuations in the magnetoionic medium in front of the FR\\,I radio galaxy 3C\\,31 (z = 0.0169) derived from rotation-measure (RM) fits to high-resolution polarization images.  Our analysis is based on VLA observations at 6 frequencies in the range 1.4 -- 8.4\\,GHz with resolutions of 5.5 and 1.5\\,arcsec FWHM (Figs~\\ref{fig:rmlo} and \\ref{fig:rmhi}). We show images of normalized polarization gradient $p^\\prime(0)/p(0)$ from fits to $p(\\lambda^2) = p(0) + p^\\prime(0)\\lambda^2$, where $p(\\lambda^2)$ is the degree of polarization at wavelength $\\lambda$ and a prime denotes differentiation with respect to $\\lambda^2$, together with RM images derived from fits to $\\chi(\\lambda^2) = \\chi(0) + {\\rm RM}\\lambda^2$ at 4 -- 6 wavelengths, where $\\chi$ is the {\\bf E}-vector position angle.  The residual depolarization at 1.5\\,arcsec resolution is very small and the rotation is accurately proportional to $\\lambda^2$, indicating almost completely resolved foreground rotation. There is a large asymmetry across the nucleus: the lobe with the brighter jet shows a much smaller RM fluctuation amplitude than the counter-jet lobe on all scales, qualitatively as expected from relativistic jet models \\citep{L88}. ",
        "conclusions": ""
    },
    "0710/0710.3062_arXiv.txt": {
        "abstract": "We systematically study the effects of collisions on the overall dynamical evolution of dense star clusters using Monte Carlo simulations over many relaxation times. We derive many observable properties of these clusters, including their core radii and the radial distribution of collision products.  We also study different aspects of collisions in a cluster taking into account the shorter lifetimes of more massive stars, which has not been studied in detail before. Depending on the lifetimes of the significantly more massive collision products, observable properties of the cluster can be modified qualitatively; for example, even without binaries, core collapse can sometimes be avoided simply because of stellar collisions. ",
        "introduction": "In dense stellar systems like massive young clusters, galactic centers, and old globular clusters (GC), the high densities of stars can give rise to many direct, single--single physical collisions, in addition to collisions mediated by dynamical interactions of binaries (\\cite{2004MNRAS.352....1F}).  In these systems stellar evolution can also be significantly modified through physical collisions.  For example, in young star clusters, collisions can give rise to runaway growth of merger products, producing many exotic stellar populations, such as intermediate-mass black holes (\\cite{2006ApJ...640L..39G}; \\cite{2006astro.ph.12040T}).  The cores of typical old Galactic GCs can also attain high enough densities so that most core stars undergo collisions during their lifetime (\\cite{1976ApL....17...87H}; \\cite{2004MNRAS.352....1F}). These collisions not only change the evolution of individual stars, but the increased stellar masses  also change the overall GC properties like the core radius ($r_c$) and the half-mass radius ($r_h$).  Moreover, at least some of the observed exotic stellar populations in dense clusters, like blue straggler stars (BSS) (\\cite{2007ApJ...663..277F}; \\cite{2001ApJ...548..323S}; \\cite{1997ApJ...487..290S}), and compact binaries like ultracompact X-ray binaries (UCXB), are likely created through collisions (\\cite{2006ApJ...640..441L}).  A dense cluster of stars naturally evolves towards eventual core collapse through relaxation. At their present ages, most Galactic GCs are expected to have collapsed cores.  However, observations show that the measured values of $r_c/r_h$ for most Galactic GCs are higher than predicted by theoretical models (\\cite{1996ApJ...458..178V}).  Many scenarios have been proposed to explain this apparent discrepancy.  For example, the core can be supported against deep collapse by dynamically extracting the binding energy of hard primordial binaries (\\cite{2007MNRAS.374..857T}; \\cite{2007ApJ...658.1047F}). However, it is hard to explain most of the high observed values of $r_c/r_h$ in Galactic GCs purely through this ``binary burning'' process. Other mechanisms to halt core collapse like ejection of stellar mass BHs from the core (\\cite{2007MNRAS.379L..40M}) or stellar captures by a central intermediate-mass black hole (\\cite{2007MNRAS.374..857T}; \\cite{2006astro.ph.12040T}) have also been discussed at this meeting (see contributions by Mackey and Trenti in this volume). It has also been suggested that $r_c/r_h$ could simply keep increasing over long timescales in clusters having fairly long relaxation times via ongoing mass segregation (\\cite{2004ApJ...608L..25M}).  Here we study stellar collisions as a possible mechanism for supporting clusters against core collapse. In a regime where collisions are important, they can produce many stars significantly more massive than those in the background population.  Thus the subsequent evolution of collision products will be much faster than for normal stars in the cluster.  Although the stellar evolution and observable properties of collision products have been extensively studied before (\\cite{2001ApJ...548..323S}; \\cite{1997ApJ...487..290S}), the feedback effects of collisions on the overall dynamical evolution of GCs has received less attention.  The shorter evolution timescales of massive collision products may support the core against collapse even without any primordial binaries or other mechanisms for energy production, simply via indirect heating through stellar evolution mass loss (\\cite{1991ApJ...378..637G}; \\cite{1989Natur.339...40G}; \\cite{1987ApJ...319..801L}).  Using $N$-body simulations, we have begun studying numerically how collisions and the subsequent evolution of collision products can alter the overall properties of GCs. We use the Northwestern group's H\\'enon-type Monte-carlo code, which provides a detailed, star-by-star representation of clusters with up to $N\\sim 10^6-10^7$ stars (\\cite{2007ApJ...658.1047F}; \\cite{2001ApJ...550..691J}; \\cite{2000ApJ...540..969J}). In \\S2 we will first present two simple limiting cases, bracketing reality and illustrating the dramatic changes in global cluster properties depending on how the evolution of collision products is treated in the models. We also present the evolution of a more realistic GC model with a conventional ``rejuvenation'' prescription for determining the lifetimes of collision products.  We discuss the implications of our study and planned future work in \\S3. \\begin{table}\\def~{\\hphantom{0}} \\begin{center} \\caption{Initial Conditions of Simulated GCs} \\label{tab:ic} \\begin{tabular}{cccccccc}\\hline &  Model  &  IMF   &  N  &  $n_{binary}$  &  $r_c$ (pc) &  Virial radius (pc) &  $\\rho_c$ ($M_{\\odot}/pc^3$) \\\\ \\hline% Cases 1,2  &  Plummer  &  Single-Mass  &  $10^6$  &  $0$  &  $0.3$  &  $0.85$  &  $10^6$  \\\\ Case 3  &  King, $w_0 = 6$  &  Kroupa ($0.1$ -- $2.0 M_{\\odot}$)  &  $10^6$  &  $0$  &  $0.87$  &  $2.89$  &  $1.7\\times 10^4$  \\\\ \\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "Although determination of stellar evolution after a collision is a subject of continuing research, the effects of collisions on cluster dynamics have not been studied in detail previously.  This is a first report of an ongoing systematic study on this effect. Using typical initial conditions for old GCs and simple assumptions for rejuvenation we show that collisions between single stars not only alter the stellar properties and produce exotic stellar populations like some of the BSS, they also affect the overall GC properties.  Most importantly, collisions can support the core of a cluster against collapse in typical old Galactic GCs.  Even with our simple but reasonable assumptions, the values of $r_c/r_h$ obtained for our simulated GCs compare well with the observed $r_c/r_h$ values for Galactic GCs.  Furthermore, we obtain a population of BSS candidates contained within the core, also consistent with observations (\\cite{2007ApJ...661..210L}).  Note that we do not find any BSSs well outside the core, consistent with the current understanding that those BSSs are most likely formed via primordial binary mergers (\\cite{2006MNRAS.373..361M}).  We have adopted many extreme simplifications for this first look at the problem. For example, at this stage of our study, once a star evolves off the MS, it is removed from the simulation, leaving no remnant.  Since remnants are normally only a few percent of the total progenitor mass, we expect that this approximation will not affect the overall GC properties significantly, so far as the increase in energy is concerned from mass loss. However, some remnants from very massive stars at a very young age can remain in the cluster and sink into the core through dynamical friction. Dynamical interactions, including collisions, of these massive remnants can also alter the core properties of the GCs in certain regimes (\\cite{2007MNRAS.379L..40M}; \\cite{2007MNRAS.374..857T}; \\cite{2006astro.ph.12040T}).  Another possibly important effect left out of this study for now is the role of primordial binaries (\\cite{2007ApJ...658.1047F}; \\cite{2004MNRAS.352....1F}).  On the one hand, the presence of primordial binaries will increase the rate of collisions through resonant encounters (\\cite{2004MNRAS.352....1F}), on the other hand, binaries will provide further support of a cluster against collapse and hence may prevent the core from reaching high enough densities for significant collisions. \\begin{center} \\begin{figure} \\includegraphics[height=2.7in,width=2.7in,angle=0]{chatterjee_f3.eps} \\includegraphics[height=2.7in,width=2.7in,angle=0]{chatterjee_f4.eps} \\caption{ {\\em a)} Position vs mass scatter plot for all collision products still at their MS life at $14$ Gyr.  Two times the present age turn-off mass ($1.6 M_{\\odot}$) is shown as a horizontal solid line to guide the eye.  The position of the core is shown with the vertical dotted line. {\\em b)} Histogram showing the positions of the same population as in $(a)$. {\\em c)} Histogram showing the positions of the collision products still in their MS life at present age having masses $\\geqslant 1.6 M_{\\odot}$.  } \\label{fig:king_BS} \\end{figure} \\end{center}"
    },
    "0710/0710.2855_arXiv.txt": {
        "abstract": "{One of the ways to determine the contribution of the dark halo to the gravitational potential of a galaxy is the study of non-circular (streaming) motions and the associated gas shocks in the bar region.  These motions, determined by the potential in the inner parts, can break the disk-halo degeneracy. Here, two main fluid dynamical approaches have been chosen to model the non-circular motions in the bar region; a 2-D Eulerian grid code for an isothermal gas (FS2) and a 3-D smoothed particle hydrodynamic code (N-body/SPH)} % {The aim of this paper is to compare and quantify the differences of the gas flows in rotating barred potential obtained using two different fluid dynamical approaches. We analyse the effect of using  2-D and a 3-D codes in the calculation of gas flow in barred galaxies and to which extend the results are affected by the code. To do this, we derive the velocity field and density maps for the mass model of NGC~4123 using a 3-D N-body/SPH code and compare the results to the previous 2-D Eulerian grid code results.} {Numerical modelling, 3-D N-body/SPH simulations} {The global velocity field and the gas distribution is very similar in both models. The study shows that the position and strength of the shocks developed in the SPH simulations do not vary significantly compared to the results derived from the 2-D FS2 code. The largest velocity difference across the shock is 20\\kms between the 2-D and 3-D fluid dynamical models.} {The results obtained in the studies deriving the dark matter content of barred galaxies using the bar streaming motions and strength and position of shocks are robust to the fluid dynamical model used. The effect of 2-D and 3-D modelling can be neglected in this type of studies.} ",
        "introduction": "There have been several studies addressing the distribution of dark matter in galaxies using the non-axisymmetry of the potential of barred galaxies: a) the fluid dynamics approach with the work by~\\citet[][hereafter WSW01]{weiner01} for the modelling of NGC~4123; ~\\cite{perez3} and \\cite[][hereafter PFF04]{perez4} for the study of the dark matter content of 5 barred galaxies (NGC~5505, NGC~7483, NGC~5728, and NGC~7267); and b) the sticky-particle approach by~\\cite{rautiainen} modelling the dynamics of  ESO~566-24, and~\\cite{salo} who modelled the H$\\alpha$ velocity field of IC~4214 to derive the halo contribution.  The streaming motions and the associated gas shocks in the bar region are determined by the potential in the inner parts of galaxies. The velocity field of an axisymmetric galaxy does not allow to uniquely disentangle the contribution of the halo from that of the disk.The signatures of non-circular motions, however, can break the disk-halo degeneracy and be used to obtain the contribution of the dark halo to the potential. Common to all the modelling, the stellar potential is generated directly from the broad band galaxy images with some recipe to treat the vertical distribution and some assumptions about the mass-to-light (M/L) of the stellar populations. The velocity fields obtained in this way are then compared to the observed kinematic information to determine if they can be reproduced by the models.  The main result, common to all the modelling carried out up to now, is the fact that the gravitational field in the inner region is mostly provided by the stellar luminous component. The bar pattern speeds found by the different groups are consistent with fast rotators, and the best fitting M/L ratios obtained are compatible with M/L ratios derived from current population synthesis models. Although it is a very powerful method, only a handful of galaxies have been modelled in this way.  Two main fluid dynamical approaches have been chosen for this type of modelling: an Eulerian grid code for an isothermal gas (FS2) and a smoothed particle hydrodynamic code (N-body/SPH). SPH is a Lagrangian method for solving Euler's equation of motion. In this technique, the system is represented by a set of particles and the gas properties are calculated by a weighted average over the neighbouring particles.  On the other hand, the Eulerian grid is fixed in space and time,  the grid nodes and cells remain spatially fixed while the materials flow through the mesh. In the modelling carried out by WSW01, the scale height is introduced as a smoothing length for the potential while the approach chosen by PFF04 uses a full 3-D calculation of the forces. For future work and as a consistency check, it is important to understand whether the results obtained are dependent a) on the dynamical code used and b) 2D versus 3-D. No significant differences are expected between the SPH and the Eulerian grid approach. Previous simulations~\\cite[][hereafter PA00]{englmaier,patsis00}, running 2-D SPH under the same conditions as used by~\\cite{athanassoula92} with the FS2 Eulerian grid code indicated that the main features, such as the strength of the shocks in the bar are quite similar and small differences arise due to the statistical nature of SPH or to numerical and artificial viscosities used in the different codes. Both studies pointed out the importance of certain parameters used in the simulations for the outcome of the modelling, such as the sound speed and the way the non-axisymmetric part of the potential is introduced.  The use of a 2-D versus a 3-D code might introduce large variations in the radial forces due to the different treatment of the vertical forces, which in turn might vary the outcome of the analysis and the conclusions derived from this methodology. In this paper, we investigate the effect of 2-D vs 3-D numerical modelling in the gas flows of rigidly rotating bar potentials. In order to test this effect, we have modelled the gas flow in NGC~4123 using the mass model and images provided by B.J. Weiner but using the 3-D N-body/SPH code used in PFF04. The parameters chosen for the SPH simulations have been taken from WSW01.  In future work we will address the comparison between the sticky particle codes and the fluid dynamical methods. The modelling is presented in Section~\\ref{modelling}, the results and comparison are presented in Section~\\ref{comparison}. Finally, discussion and conclusions are presented in Section~\\ref{conclusions} ",
        "conclusions": "\\label{conclusions}  The position and strength of the shocks in barred galaxies derived from dynamical models is robust to the chosen model. Here, we have compared 2-D to a 3-D code and despite the differences in the two codes and  the small differences in the mass model, the resulting velocity and density maps for both models is very similar. However, for similar hydrodynamical parameters, the SPH code develops a central ring, associated to the x$_{2}$ family of orbits while the gas distribution in the FS2 model does not present this central ring. This difference seems to be due to differences between the two codes and not to the difference on the treatment of the vertical forces.  In any case, the position and the strength of the shocks in both models remains very similar. We can therefore conclude that the methodology used to derive the halo contribution to the potential using dynamical modelling of the gas flow in the bar region of spiral galaxies is robust to the codes used and to wether they are 2 or 3-D codes."
    },
    "0710/0710.0257_arXiv.txt": {
        "abstract": "\\hspace{0.6cm} The thickness of the equilibrium isothermal gaseous layers and their volume densities $\\rho_{gas}(R)$ in the disc midplane are calculated for 7 spiral galaxies (including our Galaxy) in the frame of self-consistent axisymmetric model. Local velocity dispersions of stellar discs were assumed to be close to marginal values necessary for the discs to be in a stable equilibrium state. Under this condition the stellar discs of at least 5 of 7 galaxies reveal a flaring. Their volume densities decrease with $R$ faster than $\\rho_{gas}$, and, as a result, the gas dominates by the density at the disc periphery. Comparison of the azimuthally averaged star formation rate $SFR$ with the gas density shows that there is no universal Schmidt law $SFR \\sim \\rho_{gas}^n$, common to all galaxies. Nevertheless, $SFR$ in different galaxies reveals better correlation with the volume gas density than with the column one. Parameter $n$ in the Schmidt law $SFR \\sim \\rho_{gas}^n$, formally calculated by the least square method, lies within 0.8\\,--\\,2.4 range and it's mean value is close to 1.5. Values of $n$ calculated for molecular gas only are characterized by large dispersion, but their mean value is close to 1. Hence the smaller $\\rho_{gas}$ the less is a fraction of gas actively taking part in the process of star formation.  To be published in Astronomy Reports, 2008. ",
        "introduction": "\\hspace{0.6cm}The key question of galaxy evolution is the dependence of star formation rate $SFR$ on the gas density and other interstellar medium parameters, averaged by large enough area or  volume for smoothing random fluctuations in gas and young stars distributions.  Scmidt~\\cite{Schmidt59}  suggested  a simple form of star formation rate parameterization: $SFR_v\\sim\\rho_{gas}^n$, usually called the Schmidt law. From the analysis of gas and young objects distribution in the solar vicinity, Schmidt~\\cite{Schmidt59} obtained $n\\approx 2$. Later a large number of papers were published aimed to check and to interpret the Schmidt law, based on the gas and $SFR$ distributions in the discs of  spiral galaxies. In particular, it was found that both local, azimuthally averaged $SFR$ values as well as the total star formation rate in a galaxy correlate (although not too tightly) with the gas content in its disc (e.g. Madore et al., ~\\cite{Madore74}, Kennicutt, ~\\cite{Kennicutt89}, Wong\\&Blitz, ~\\cite{WB02}, Boissier et al., ~\\cite{Bois03}, Shuster et al.,~\\cite{Schu06}). Being empirical by it sence, the Schmidt law and its modifications open a possibility to calculate evolution models of galaxies, parameterizing the star formation history. This allows us to understand better  mechanisms which regulate the rate of star formation.  A situation is entangled by the fact that the dependence $SFR_v(\\rho_{gas})$ (i.e. the local Schmidt law) and, in particular, the power $n$ can't be found directly from observations of other galaxies, because in order to estimate the volume-related values of $SFR_v$ and $\\rho_{gas}$ it is necessary to know the gas layer thickness, which may vary significantly both along the galaxy radius and from one galaxy to another. Therefore in practice the Schmidt law is often replaced by the other one, outwardly alike empirical law $SFR_s\\sim\\sigma_{gas}^N$ (sometimes it is called Kennicutt\\,--\\,Schmidt law), where the compared values are scaled to unit disc surface area, or by the law $SFR_t\\sim M_{gas}^N$ for the total values $SFR_{t}$ and gas mass $M_{gas}$ (so-called the global Schmidt law). These dependencies are more complicated for interpretation, because the compared parameters are the integrals of heterogeneous functions of distribution along of the line-of-sight (in the first case) or allover a whole disc (for the global law). In  general case, parameters $n$ and $N$ must not coincide (see the discussion of the question in papers Madore et al., ~\\cite{Madore74}, Tutukov, ~\\cite{tut06}). Only if at any distance from disc plane $n=1$, the power $N$ is also equal to unit. In most cases values of $N$ that obtained for different galaxies lay within the limits of $1<N<2$ (Wong and Blitz, ~\\cite{WB02}, Boissier et al.,~\\cite{Bois03}, Shuster et al.,~\\cite{Schu06}), but for some galaxies they proves to be more steep. For example, for M~33 $N>$3 (Heyer et al., ~\\cite{HeyerAll04}). For the case of the global Schmidt law when different galaxies are compared, $N\\approx 1.4-1.5$ (Kennicutt, ~\\cite{Kennicutt98}, Li et al., ~\\cite{li+06}). The scatter  of $N$ remains large: for galaxies with similar gas masses $SFR_{t}$ may differ by an order of magnitude.  Apparently, it is possible to reproduce the Schmidt law for different star formation models using some simplifying assumptions on mechanisms of self-regulation of large-scale star formation, or considering conversion of neutral gas into dense molecular clouds and formation of stars (see e.g., \\cite{tut06,li+06,elm02,krumholz05,Gerristen97}). Note however that usually  variation of the thickness of gas layer along galactic radius as well as from one galaxy to another is ignored, and $\\rho_{gas}$ is accepted to be much smaller than  the volume stellar density, that is not always correct.  In the first part of this paper, the gas volume density in the plane of an equilibrium disc is calculated as a function of $R$ for several nearby galaxies. In the second part, the relation between the gas density and star formation rate is analyzed. Chosen galaxies strongly differs by their properties. They include: M~33 and M~101 --- multiarmed late-type galaxies (Sc), the first of them is rather small; interacting galaxy M~51 and galaxy M~100 which are distinguished by high molecular gas content; Seyfert galaxy M~106 where the ejections from nucleus and star formation burst in the inner region are observed; massive early-type Sab galaxy M~81 which possesses a large bulge and regular spiral structure, and, finally, our Galaxy. The main simplifications we use are the following. The gaseous layers in galaxies are assumed to be axisymmetric and being in hydrostatic equilibrium. The pressure of the gas is determined by its turbulent motion: $P_{gas}=\\rho_{gas}\\,C_z^2$, where $C_z$ is one-dimensional velocity dispersion, which is assumed to be constant (although different for atomic and molecular gas). These simplifying suggestions are definitely too tough for  regions enveloped by the intensive star formation, for the inner discs deep inside dense bulges or in the neighborhood of active nucleus, and also for the far periphery of discs. Note however that gas velocity dispersion, although may slowly vary with the distance from the galaxy center, remains high enough even at large distances (see the discussion in paper Dib et al., ~\\cite{dib06}). Magnetic field pressure gradient and thermal gas pressure play significantly lesser role in the formation of gas layer thickness, at least for the case of our Galaxy (see  discussion of this question in Cox review,~\\cite{Cox05}). It is essential that within the limits of approximations mentioned above the observed distribution of atomic (HI) and molecular (H$_2$) hydrogen disc thicknesses along the Galaxy's radius can be sufficiently explained (Narayan and Jog,~\\cite{NJ02}). ",
        "conclusions": "\\hspace{0.6cm}As it follows from this paper, the assumption of marginally stable stellar discs leads to conclusion that their thickness, at least in some galaxies, changes significantly (usually increases) along the  radius. The thickness of equlibrium gas layer increases in all cases --- either nearly linearly within a wide range of $R$ (M~33, M~51, M~81 and M~101), or nonlinearly, with a positive second derivative (M~100 and the Galaxy). Volume gas density in the midplane in all cases decreases within the considered range of $R$ down to several units of $10^{-25}$ g/cm$^3$ . In all cases the stellar density decreases steeper, than the gas density, so at the peripheries of galaxies the gas midplane density becomes comparable with the stellar density or even exceeds it.  To compare the half thicknesses (HWHM) with the densities of stellar discs and gaseous layers in the galaxies it is convenient to consider their parameters at a fixed radius, say, $R\\approx 2R_0$ (see Table~\\ref{tab5}). \\begin{table}[h] \\caption{Stellar disc and gas layers thicknesses and the densities of gas at $R\\approx 2R_0$. \\emph{Columns}: (1) --- a galaxy, (2) --- stellar disc thickness to radial scalelength ratio, (3) --- atomic gas disc thickness, (4) --- molecular gas disc thickness, (5) --- volume gas density (including helium).}\\label{tab5} \\begin{center} \\begin{tabular}{ccccc} &&&&\\\\ Galaxy&$h_*(2\\,R_0)/R_0$&$h_{\\rm HI}(2\\,R_0)$\\,,\\,pc& $h_{\\rm H2}(2\\,R_0)$\\,,\\,pc&$\\rho_{gas}(2\\,R_0)$\\,,\\,g/cm$^3$\\\\ &&&&\\\\ \\hline &&&&\\\\ (1)&(2)\\,&(3)\\,&(4)\\,&(5)\\,\\\\ &&&&\\\\ \\hline &&&&\\\\ M~33&0,18&87,8&57,6&$6,3\\cdot10^{-24}$\\\\ &&&&\\\\ M~51&0,23&123,6&81,3&$3,5\\cdot10^{-24}$\\\\ &&&&\\\\ M~81&0,34&129,2&85,7&$1,1\\cdot10^{-24}$\\\\ &&&&\\\\ M~100&0,20&108,8&71,0&$6,1\\cdot10^{-24}$\\\\ &&&&\\\\ M~101&0,19&173,8&114,8&$1,5\\cdot10^{-24}$\\\\ &&&&\\\\ M~106$^*$&0,27&150,0&99,2&$2,2\\cdot10^{-24}$\\\\ &&&&\\\\ Galaxy&0,11&100,3&61,4&$4,4\\cdot10^{-24}$\\\\ &&&&\\\\ \\hline \\end{tabular} \\end{center} $^*$For M~106 all the values are given for R\\,=\\,1,5\\,$R_0$. \\end{table} As it follows from the Table, the relative thickness of stellar disc $h_*/R_0$ at $R=2R_0$ lies within the range of 0.1\\,--\\,0.3. Our Galaxy is the most thin ($h_*/R_0 \\approx 0.1$) in our sample, whereas the stellar discs of two giant early type spiral galaxies (M~81 and M~106) are almost three times thicker. The thicknesses of gas layers differ less than stellar ones in the galaxies. The most thin gas layer (in M~33) is less than twice thinner than the most thick one (in M~101). The gas density (at distance $2R_0$) in all cases is equal to several units of 10$^{-24}$\\,g/cm$^{-3}$ with no obvious dependence on the morphological type: in M~101 (Sc) it is approximately the same as in M~81 (Sab), and in our Galaxy (Sb\\,--\\,Sbc) it is close to that in M~33.  The comparison of the gaseous density with the star formation rates (figures\\,\\ref{fig9}a,~\\ref{fig9}b, and~\\ref{fig10}a) gives an evidence that the replacement of surface gas density (Figure~\\ref{fig9}b) by the gas volume density leads to more tight dependencies both for the ``surface'' and ``volume'' star formation rates.  The reason for it is rather evident: star formation takes place in a narrow layer near the midplane, hence this process is sensitive namely to the volume gas density, whereas the directly measured surface gas density depends, in addition, to the gas layer thickness (more precisely --- to the vertical density profile), which differs from galaxy to galaxy.  It is worth to remind that the regions that are very close to the center (in all galaxies but M~33), or are located at the disc peripheries (beyond the well defined spiral arms) need a special study and were not considered in this paper. Indeed, here we ignored the influence of the bulge on the vertical density profiles. The accepted value of the turbulent velocity may also be not suitable for all radii. A formal inclusion of the most inner and outer regions into our diagrams leads to the resulting dependencies ``gas density\\,--\\,$SFR$''\\,\\, which are less regular and more different between galaxies.  Of the galaxies we consider, M~106 stands out by almost constant density $\\rho_{gas}$ between 4 and 10\\,kpc. The increasing of stellar disc thickness along the radius in this galaxy accompanies by the increasing of the surface gas density out of the centre within the wide range of $R$. By this reason this galaxy differs from the others in the diagrams ``$SFR$\\,--\\,gas density\\,''. The similar abnormal behavior of this galaxy  on the Kennicutt-Schmidt diagram was also revealed by Boissier et al.~\\cite{Bois03}. However the molecular gas density in M~106 behaves in the common manner i.e. steeply decreases with $R$, that explains its ``normal'' position at the diagram $SFR_v\\,-\\,\\rho_{\\rm H2}$. Note, that this galaxy is not typical in other aspects: it differs from the others by the presence of the active nucleus, as well as by the very active star formation in the central part of the disc.  The values of $n$ in the Schmidt law $SFR_v\\,\\sim\\rho_{gas}^n$ for the galaxies we consider are rather different (see Table~\\ref{tab4}). For all galaxies except M~101 the exponent $n>1$, and its mean value is close to ``standard''\\,\\,value $\\approx 1.4$ for ``global''\\,\\,Kennicutt-Schmidt law (column (3) of the Table). However, the mean value of $n$ approaches unit, if to compare the $SFR_v$ with the molecular gas density (column (4) in the Table).  The most ``steep''\\,\\, relationships $SFR(\\rho_{gas})$ are obtained for M~33, M~81 and the Galaxy. In the case of M~33 it is caused by the steep decreasing of the $SFR$ with the galactocentric distance (see Figure~\\ref{fig8}). In the case of M~81, it is the consequence of the atypically slow decrease of the gas volume density along $R$ (see Figure~\\ref{fig3}) when the radial gradient of $SFR$ is moderate. For these three galaxies, the exponent $n$ is close to the ``classical''\\,\\,value $n=2$, suggested earlier by Schmidt.  In conclusion, parameter $n$ by no means can be considered as a universal one. The  Schmidt law  has a very approximate nature, and it fulfils much better when the  volume instead of the surface gas density is used. In this case the exponent $n<2$ (with some possible exceptions). It is important that in spite of a large dispersion its value becomes closer to unit in the mean if to replace the total gas density $\\rho_{gas}$ with the molecular gas density $\\rho_{\\rm H_2}$.  It is evident that relationship between the volume gas density and $SFR$ for galaxies should not be too tight because the process of star formation depends on a number of parameters besides the mean gas density. The latter seems to be a crucial factor only for the most dense gas such as the gas in the nuclei of molecular clouds where the HCN radiation comes from. Indeed, as observations show, the star formation rate depends linearly  on the gas mass determined by the HCN line intensity, i.e for the most dense molecular gas $n\\approx 1$ (Gao and Solomon,~\\cite{gao04a,gao04b}). Although the number of galaxies we considered is small, the results allow to propose that the value of $n$ is on average close to unit even for the less dense molecular gas, that reveals itself in CO line. Since $n>1$ for the total ($HI+H_2$) gas density, one may conclude that not only star formation rate, but also the efficiency of star formation ($SFR$ per gas mass unit) decreases along with $\\rho_{gas}$. In other words, the less is the gas density, the longer time the gas remains in the rarified atomic state (that is a time scale of gas consumption is larger). It agrees with general conception, according to which the fraction of the interstellar gas which takes active part in the process of a star formation, decreases when the mean total density of gas becomes lower (see the discussion in~\\cite{WB02,elm03}).  The other important factor, besides the gas density, that determines the star formation rate at a given radius $R$ is the surface density of the old stellar disc $\\sigma_{*}$ (see papers~\\cite{ZasAbr06,BdeJ00,DR94}). As it was shown by Zasov and Abramova~\\cite{ZasAbr06} on the example of four well-studied galaxies, local star formation efficiency defined as $SFE\\,=\\,SFR$/$\\sigma_{gas}$  changes approximately as $\\sigma_{*}^{\\footnotesize 0,7}$ in a wide interval of $R$. Partially this relationship may be explained by higher volume gas density (for a given surface density) in those regions, where $\\sigma_{*}$ is higher (that is closer to galaxy center). Nevertheless this relationship can not be reduced to the simple exponential Schmidt law (neither ``volume'' nor ``surface'' one) because the clearly defined relationship between the gas density and surface disc density is absent. It evidences the existence of more deep connection between the present-day star formation on the one hand and already formed stellar disc and the gas density on the other hand, which cannot be described by simple empiric laws.  \\bigskip  The authors thank Igor V. Abramov for assistance in the numerical solutions and D.Bizyaev for helpful discussion.  \\bigskip  This work was supported by the Russian Fond of Basic Researches grant 07-02-00792."
    },
    "0710/0710.2252_arXiv.txt": {
        "abstract": "{}{To   reexamine  the  implications   of  the   recent  HESS observations   of  the   blazar  1ES0229+200   for   constraining  the extragalactic  mid-infrared   background  radiation.}{We  examine  the effect of \\gray ~absorption by the extragalactic infrared radiation on predicted intrinsic  spectra for this  blazar and compare  our results with  the observational  data.}{We  find agreement  with our  previous results on  the shape  of the IR  spectral energy  distribution (SED), contrary  to the  recent assertion  of  the HESS  group. Our  analysis indicates that 1ES0229+200  has a very hard intrinsic  spectrum with a spectral index between  1.1 $\\pm$ 0.3 and 1.5 $\\pm$  0.3 in the energy range between $\\sim$0.5 TeV and $\\sim$15 TeV.}  {Under the assumptions that  (1) the  SED  models of  Stecker,  Malkan \\&  Scully (2006)  are reasonable as derived from  numerous detailed IR observations, and (2) spectral indexes in the range $1 < \\Gamma < 1.5$ have been shown to be obtainable   from   relativistic    shock   acceleration   under   the astrophysical conditions  extant in blazar flares  (Stecker, Baring \\& Summerlin 2007), the fits to the observations of 1ES0229+200 using our previous  IR  SEDs  are  consistent   with  both  the  IR  and  \\gray\\ observations.   Our  analysis presents  evidence  indicating that  the energy spectrum  of relativistic particles in  1ES0229+200 is produced by  relativistic  shock  acceleration,  producing an  intrinsic  \\gray ~spectrum with index $1 < \\Gamma < 1.5$ and with no evidence of a peak in the SED up to energies $\\sim$ 15 TeV.}  { ",
        "introduction": "Shortly after the first strong  \\gray ~ blazar 3C279 was discovered by the EGRET  detector aboard the  Compton Gamma Ray  Observatory (CGRO), Stecker, De  Jager \\&  Salamon (1992) proposed  that the study  of the spectra  of such  sources could  be  used to  probe the  intergalactic infrared  radiation.   At  that  time,  there were  no  actual  direct observations of  the diffuse extragalactic IR  background or extensive observations of  the sources of such  radiation. The idea  was to look for  the  effects  of  photon-photon  annihilation  interactions  into electron-positron pairs. The cross section for this process is exactly determined; it can be  calculated using quantum electrodynamics (Breit \\&  Wheeler 1934).   Thus, in  principle,  if one  knows the  emission spectrum  of an  extragalactic source  at  a given  redshift, one  can determine the  column density  of photons between  the source  and the Earth.  In the last  15 years, great advances have  been made in extragalactic infrared astronomy. The diffuse  background at wavelengths not totally dominated by  galactic or zodiacal  emission has been measured  by the Cosmic  Background  Explorer (COBE).   In  addition,  there have  been extensive observations of  infrared emission from galaxies themselves, whose  total  emission is  thought  to  make  up the  cosmic  infrared background (see review by Hauser  \\& Dwek 2001).  The latest extensive observations  have been  made by  the Spitzer  satellite.  It  is thus appropriate  to  use  a  synoptic  approach combining  the  TeV  \\gray ~observations with the extragalactic infrared observations in order to best  explore both  the  TeV  emission from  blazars  and the  diffuse extragalactic infrared radiation.  Aharonian et  al.  (2007) have recently observed  the spectrum of the BLLac  object 1ES0229+200 up to  an energy $\\sim$~15  TeV with the High Energy  Spectroscopic System (HESS).  Then, by  assuming that the intrinsic spectral index of this source is greater than 1.5, they drew conclusions  regarding wavelength  dependence and  flux of  the mid-IR extragalactic background  radiation.  Their conclusions  regarding the mid-IR  extragalactic  background  radiation  appear to  disfavor  the results   of  the   extensive  semi-empirical   calculations   of  the extragalactic  IR  background spectrum  given  by  Stecker, Malkan  \\& Scully (2007) (SMS).  In this paper, we  will reexamine the assumptions and conclusions of Aharonian et  al. (2007) and show that  the observations of 1ES0229 are fully consistent with  the diffuse IR background spectrum obtained by SMS.  Furthermore, we will show that 1ES0229+200 is an example of a set of blazars which exhibit  very hard spectra that are indicative of relativistic shock acceleration. ",
        "conclusions": "Aharonian et al. (2007), by assuming that blazar spectra have spectral indexes  $\\Gamma >  1.5$, concluded  that the  mid-IR  spectral energy distribution  (SED) must  have  a wavelength  dependence steeper  than $\\lambda ^{-1}$, in their terminology $\\lambda^{-\\alpha}, \\alpha > 1.1 \\pm 0.25$ in the wavelength range  between 2\\mic ~ and 10\\mic . On the contrary, galaxy  emission models which take into  account emission by warm dust  and PAH  and silicate  emission in the  mid-IR, as  well as direct mid-IR observations of  galaxy spectra give values for $\\alpha$ in the range between $\\sim$ 0.7  and $\\sim$ 0.8 (Spinoglio et al 1995, Xu  et al.   2001).  These  values  are consistent  with the  2-10\\mic ~diffuse background SED given  in the semi-empirical SMS models.  Such values lead to an energy dependence for the \\gray ~optical depth which is consistent with that obtained by both Totani \\& Takeuchi (2002) and SMS. This  result is also supported  by the lower limit  on the 15\\mic ~diffuse background flux obtained by galaxy counts of 3.3 $\\pm$ 1.3 nW m$^{-2}$sr$^{-1}$ (Altieri et al.  1999) which, under the conservative assumption that 80\\% of the  mid-IR flux is resolved out (SMS), yields a value for the total diffuse background flux at 15\\mic ~ of 4.1 $\\pm$ 1.6  nW m$^{-2}$sr$^{-1}$,  higher  than  the upper  limit  of 3.1  nW m$^{-2}$sr$^{-1}$  derived  by  Aharonian  et al.   (2007)  under  the assumption that $\\Gamma > 1.5$.  Under the  assumptions that (1) the  SED models of  Stecker, Malkan \\& Scully (2006)  (SMS) are reasonable as derived  from numerous detailed IR observations, and  (2) spectral indexes in the range  $1 < \\Gamma < 1.5$ have been  shown by Stecker, Baring \\& Scully  (2007) (SBS) to be obtainable   from   relativistic    shock   acceleration   under   the astrophysical conditions extant in blazar flares, the fits to the HESS TeV  observations  of 1ES0229+200  using  the  SMS  infrared SEDs  are consistent with both the IR and \\gray\\ observations.  The  SBS simulations indicate  with specific  test runs  that electron spectra with asymptotic spectral indexes  between 1.26 and 1.62 can be obtained from  acceleration by  relativistic shocks with  bulk Lorentz factors between 10 and 30 and these electrons can then Compton scatter to produce \\gray~ spectra with  indexes $\\sim$1.1 and $\\sim$ 1.3.  The simulations  show that  larger bulk  Lorentz factors  lead  to flatter spectra.  Such  Lorentz factors of 50  or more have  been implied from studies of specific flares in  BL Lac objects (Konopelko et al.  2003; Bagelman,  Fabian  \\& Rees  2007)  so  that  the existence  of  highly relativistic shocks in such sources with  indexes as flat as $\\sim$1 is not unreasonable.  For the SMS fast evolution SED,  the \\gray~ index obtained is 1.11 and for the baseline  SED, the index obtained is  1.45.  Our analysis thus presents evidence  for relativistic shock  acceleration in 1ES0229+200 that results in a very hard intrinsic \\gray ~spectrum with no evidence of  a peak  in the  \\gray ~  SED up  to energies  $\\sim$ 10  TeV. Unfortunately,  there  are   no  simultaneous  observations  at  other wavelengths  that can  be used  to model  the flare  that  occurred in 1ES0229+200 and, in any case,  TeV orphan flares have been observed in other BL Lac  objects.  However, we note that  hard X-ray spectra have been seen  in other sources as  discussed in Section 3.   We also note that  SBS have  derived  spectra for  three  other BL  Lac objects  at redshifts between 0.18 and 0.19  with indexes between 1 and 1.5, viz., 1ES1218+30, 1ES1101-232, and 1ES0347-121.  It therefore appears that a whole class of blazars or  blazar flares may exhibit the hard-spectrum characteristics of relativistic shock acceleration."
    },
    "0710/0710.2314_arXiv.txt": {
        "abstract": "Deep \\textit{Spitzer} IRAC images of L1157 reveal many of the details of the outflow and the circumstellar environment of this Class 0 protostar. In IRAC band 4, 8 $\\mu$m, there is a flattened structure seen in absorption against the background emission. The structure is perpendicular to the outflow and is extended to a diameter of $\\sim$2$\\arcmin$. This structure is the first clear detection of a flattened circumstellar envelope or pseudo-disk around a Class 0 protostar. Such a flattened morphology is an expected outcome for many collapse theories that include magnetic fields or rotation. We construct an extinction model for a power-law density profile, but we do not constrain the density power-law index. ",
        "introduction": "\\label{dist}  The L1157 dark cloud in Cepheus (IRAS 20386+6751) conceals a young protostar, a so-called Class 0 source \\citep{andre1993}, which is deeply embedded within a large circumstellar envelope \\citep{gueth2003,beltran2004}.  L1157 has a large powerful molecular outflow that is the prototype of chemically active outflows \\citep{bachiller2001}. Despite the attention that L1157 has received at radio wavelengths, few observations have been made in the near to mid-infrared outside of observations of the outflow in H$_2$ and K-band \\citep[e.g.][]{davis1995,cabrit1998}. Only recently have sensitive instruments been available to observe these objects shortward of 10 $\\mu$m \\citep[e.g.,][]{tobin2007}. The outflow carves cavities in the circumstellar envelope, which allow photons from the embedded central source to escape and scatter off dust in the cavity at NIR wavelengths. The morphology of the scattered light can be used to probe many of the fundamental properties of the source such as opening angle, envelope mass, etc. \\citep[e.g.,][]{whitney2003a,whitney2003b,tobin2007,robitaille2007,seale2008}.  In this letter, we present new, deep \\textit{Spitzer Space Telescope} observations of L1157. The IRAC continuum emission is dominated by molecular line emission in the outflow. Near the source there is a small amount of emission that may be attributed to scattered light and perhaps molecular line emission that is highly excited by the outflow jet. Other than the enormous outflow ($\\sim$0.5 pc) to the north and south, the most prominent feature observed is a large, flattened absorption feature at 8.0 $\\mu$m and less defined at 5.8 $\\mu$m. This absorption feature is a flattened circumstellar envelope observed in silhouette against the Galactic infrared background.  The distance to L1157 is important to any physical interpretation, but the distance is highly uncertain. The molecular clouds in Cepheus have three characteristic distances, 200, 300, 450 pc \\citep{kun1998}; L1157 has a similar galactic latitude as the 200 pc and 300 pc absorbing clouds. Due to this, we disagree with the current accepted distance of 440 pc. This value was based upon a study of NGC 7023 in \\citet{viotti1969}; L1157 is not in clear association with this cluster.  In this letter, we use a distance of 250 pc. ",
        "conclusions": "The shape of the absorption feature is especially intriguing, as it looks like a disk structure perpendicular to the outflow axis. This is the first clear detection of a flattened envelope or pseudo-disk in a Class 0 object. \\cite{galli1993a,galli1993b} have shown that a modest magnetic field structure modifies infall from the initial spherical cloud to form a so-called ``pseudo-disk''; a flat thin structure in the equatorial plane that is not rotationally-supported, thus collapsing. This type of structure is also seen from simple flat sheet models of collapse \\citep[e.g.,][]{hartmann1994,hartmann1996}, as well as detailed ambipolar diffusion models \\citep[e.g.,][]{fiedler1993}. On the other hand, this structure is large $\\sim$15,000-30,000 AU, depending on the background used. That size is somewhat larger than the inner envelope size estimated from interferometric models of the dust continuum \\citep{looney2003}.  However, the single-dish dust emission \\citep{gueth2003} is extended along the same axis as the absorption, which argues that the inner envelope in the equatorial plane either has higher density, and/or different dust opacity properties.  Our modeling results show that the properties of the structure, as determined by the absorption model, are consistent with the above theoretical constructs, i.e. flattened envelopes and density profiles. Although we model a range of indexes ranging from p = 0.5 to 3, only the 0.5 to 2 provide acceptable fits at the 90\\% confidence level with the vast majority of fits being p=1.5.  To better explore the physical meaning the model, we assume dust opacities (dust plus gas) from \\citet{lidraine2001} of $\\kappa_{8.0\\mu m} =5.912~cm^2~g^{-1}$. Although using interstellar dust opacities for $\\kappa_\\lambda$ is probably incorrect, as Class 0 sources are thought to have already experienced some grain growth \\citep[e.g.,][]{looney2003,natta2007}, it is still a useful approximation.  Using the assumed $\\kappa_{8.0\\mu m}$, the derived range for the density reference, $\\rho_0$, or 1 pixel (1$\\farcs$2) from the center of the envelope (i.e. 300 AU at 250 pc) is listed in Table \\ref{fits}.  In addition, we can estimate the absorbing mass of the flattened envelope component for each ``best fit'' model of Figure \\ref{fit} and a height of 3 pixels (the size of the box we averaged over). We assume that the vertical density profile is constant for the mass estimate, even though the observed absorption falls off vertically with scale heights of $\\sim$3-4 pixels, using a Gaussian vertical structure. Our mass estimate, listed in Table \\ref{fits}, ranges from 0.08 to 0.16 M$_\\odot$. Without using any model, we can also estimate the mass required for the observed extinction. We use the above value for $\\kappa_{8.0\\mu m}$ and a background of 0.54 MJy/sr to calculate the mass necessary for the extinction of all pixels in the central region below 0.458 MJy/sr (577 pixels). The total absorbing mass required for those pixels is 0.19 M$_{\\odot}$. It is important to note that this mass is only in the absorbing pseudo-disk, but the mass is comparable to the 140$\\arcsec$ extended envelope detected in the millimeter continuum as discussed in \\citet{gueth2003} with an estimated mass of 0.7 M$_{\\odot}$ at an assumed distance of 250 pc.  This implies that although a large fraction of the mass of the envelope is in the flattened structure, most of the mass is more diffuse. Due to the log nature of the mass absorption, the density contrast from the center of the absorption feature to slightly offset from the absorption feature is approximately an order of magnitude. This is more than expected from the numerical models of \\cite{galli1993b}, but somewhat consistent with the ambipolar models of \\cite{fiedler1993} and the sheet collapse model of \\cite{hartmann1996}.  As can be seen in Table \\ref{fits} and Figure \\ref{fit}, these data can not constrain the model density profiles directly; multiple power-laws are allowed.  On the other hand, although many theoretical models do suggest a flattened envelope structure, the power-law index in the pseudo-disk is expected to evolve.  In that case, we would not expect a single power-law to well describe the envelope density. In fact, the ambipolar models \\citep[e.g.][]{tassis2005a,tassis2005b} suggest that the power-law can be episodic in the flattened envelope. Further studies with increased sensitivity need to be compared directly to the theoretical density profiles in the flattened objects to say anything more completely."
    },
    "0710/0710.0311_arXiv.txt": {
        "abstract": "We derive probability density functions for the projected axial ratios of the real and mock 2PIGG galaxy groups, and use this data to investigate the intrinsic three dimensional shape of the dark matter ellipsoids that they trace. As well as analysing the raw data for groups of varying multiplicities, a convolution corrected form of the data is also considered which weights the probability density function according to the results of multiple Monte-Carlo realizations of discrete samples from the input spatial distributions. The important effect observed is that the best fit distribution for all the raw data is a prolate ellipsoid with a Gaussian distribution of axial ratios with $\\bar{\\beta}=0.36$ and $\\sigma=0.14$, whilst for the convolved data the best fit solution is that of an oblate ellipsoid $\\bar{\\beta}=0.22$ and $\\sigma=0.1$. Previously only prolate distributions were thought compatible with the data, this being interprated as evidence of filamentary collapse at nodes. We also find that even after allowing for the sampling effects, the corrected data is better fit using separate multiplicity bins, which display a trend towards more spherical halos in higher multiplicity groups. Finally, we find that all results in the real data are in good agreement with the mock data from $\\Lambda$CDM simulations, KS tests showing that all comparative data have been drawn from the same distributions within the $1\\sigma$ confidence limits. ",
        "introduction": "The spatial distribution of galaxies traces the shape of the dark matter potential in which they are embedded. Simulations show that dark matter halos are not spherical, as one would naively expect because of dark matter's non-dissipational nature, but are strongly flattened triaxial ellipsoids \\citep{dubinski91}. Although prolate and oblate shapes are equally likely in the simulations, dissipative infall of baryonic gas eventually forces the halo shapes toward pronounced oblateness with axial ratios $b/a > 0.7$ \\citep{katz91, dubinski94,combes02}.  Because dark matter structures evolve self-similarly it is possible to derive the general shape of dark halos from analysis of single test objects. The first such attempts were carried out on rich clusters, whose shapes were found to be triaxial and strongly prolate \\citep{plionis01}: however the non-linear evolution of these objects may play a role in shaping their density distributions \\citep[e.g][]{binney79}. Indeed there is evidence that the more high-density systems are more spherical than low-density objects \\citep[random gaussian field work of][]{bardeen86}. The observed trends are the opposite of what is observed in simulations \\citep[e.g.][]{kasun05,allgood06}.  The distribution of Milky Way satellites shows a strongly oblate and flattened distribution \\citep{ruzicka07}, which is also consistent with analysis of the orbital planes of the Sagittarius dwarf \\citep{johnston05} and the Monoceros stream \\citep{penarrubia05}. However, star counts \\citep{lemon04} and analysis of the stellar stream of the Sagittarius dwarf \\citep{ibata01} support a spherical halo. It is also uncertain whether the dwarf satellites can be used as test particles, as they may not originate from cosmological sub-structures \\citep{kroupa05}.  Groups of galaxies are likely to be the best testbeds to study the shapes of dark matter halos. Large group catalogs are now available from redshift surveys (e.g., the 2dF Percolation Inferred Galaxy Groups [2PIGG] of \\citealt{eke04} and the group catalog of \\citealt{merchan05} from Data Release 3 of the Sloan Digital Sky Survey) and these allow a statistical approach to the study of the shapes of groups and the shape dependence on richness, multiplicity and dynamical evolution. Early studies marginally favoured prolate shapes \\citep{fasano93,orlov01}, but not at the exclusion of oblate solutions. The most recent study of 2PIGG groups by \\cite{plionis06} favours prolate groups, but is obtained by means of a multiplicity cut and as such represents a different method to the one presented here. Prolate results indicate that the original (oblate ?) shapes may have been strongly modified during gravitational collapse, or that filamentary collapse is in fact being witnessed.  In this paper we carry out a detailed analysis of the shape of 2PIGG groups, coupled with extensive Monte Carlo simulations and comparison with groups extracted from the 2dF mock catalogs used by \\cite{eke04} to optimize group selection. In section 2 we present our method for analysing the shape of observed groups and in section 3 apply this to the real and mock 2PIGG catalog \\citep{eke04}. Comparisons are made between the raw data and data that is corrected for the finite (indeed, often sparse) sampling by factors which we determine from Monte-Carlo simulations of suitably populated groups. Furthermore, the mock and real data are considered separately and KS tests are used to confirm their distributional similarities. ",
        "conclusions": "We found good agreement between our results for raw data with a simple 20+ multiplicity cut and findings by other authors, indicating a strongly prolate distribution with a mean $\\bar{\\beta}=0.46$ and $\\sigma_{\\beta}=0.14$. However, when fits were applied to convolution corrected distributions which allow for the error introduced by finite sampling we found that for the higher multiplicity cuts both oblate and prolate distributions could produce reasonable fits. More significantly, evidence suggests that these large multiplicity populations share the same underlying distribution (if oblate: $\\bar{\\beta}\\sim0.3$ $\\sigma_{\\beta}=0.1$, if prolate $\\bar{\\beta}\\sim0.44$ $\\sigma_{\\beta}=0.14$). There is evidence in the data that low multiplicity groups have a highly elliptical sub population that is missing in other multiplicity cuts: for 5-9 multiplicity cuts oblate distributions produce a better quality of fit ($\\bar{\\beta}=0.2$, $\\sigma_{\\beta}=0.1$), this is despite the effect of the convolution correction being to reduce this signal. This is seen as evidence of high multiplicity groups being more spherical, and is consistant with them being located in nodes, or at least in a less filamentary structure than the extremely low multiplicity groups. Large multiplicity groups are better fit by a prolate distribution in the corrected data, but oblate fits are not as strongly rejected as in earlier work which did not account for the sampling bias. The effect of interlopers was extensively considered, and the results of this work imply all measurements obtained should be considered upper limits for both the under lying axial ratios (i.e. the true non-interloper distribution will be more elliptical than that measured) and the standard deviations of the populations.  KS tests were utilised where appropriate and indicate that the real and mock catalogs are in very good agreement over all ranges. For all multiplicity comparisons the real and mock data share the same underlying distribution to within $1\\sigma$ expectations, and the most self similar populations are for multiplicities 10-19 and 20+, which appear consistent with being part of the same overall population once the corrections have been applied.  Further exploration of projected group shapes will be presented in a future paper which will consider the differences we find when colour cuts and different grouping algorithms are used."
    },
    "0710/0710.1691_arXiv.txt": {
        "abstract": "Dynamical mass estimates of ultra-compact dwarfs galaxies and massive globular clusters in the Fornax and Virgo clusters and around the giant elliptical Cen\\,A have revealed some surprising results: 1) above $\\sim10^6 M_\\odot$ the mass-to-light ($M/L$) ratio increases with the objects' mass; 2) some UCDs/massive GCs show high $M/L$ values (4 to 6) that are not compatible with standard stellar population models; and 3) in the luminosity-velocity dispersion diagram, UCDs deviate from the well defined relation of ``normal'' GCs, being more in line with the Faber-Jackson relation of early-type galaxies. In this contribution, we present the observational evidences for high mass-to-light ratios of UCDs and discuss possible explanations for them. ",
        "introduction": "The so-called ultra-compact dwarf galaxies (UCDs) are very massive ($10^6 M_\\odot<M<10^8 M_\\odot$), old, compact stellar systems that were discovered in nearby galaxy clusters about a decade ago (\\cite{hilk99}, \\cite{drin00}). Their nature is unknown yet. Maybe they are remnant nuclei of disrupted galaxies, or maybe they are merged stellar super-clusters formed in interacting galaxies. Regardless of what UCDs actually are, some properties divide them from ``ordinary'' globular clusters (GCs). The half-light radii of UCDs scale with luminosity reaching $\\sim90$ pc for the most massive UCDs. Unlike for GCs, their densities within the half-light radii are not increasing with mass but stay at a constant level or even decrease. Thus UCDs are not that compact at all when compared to $10^6 M_\\odot$ GCs, but certainly much denser than dwarf ellipticals of comparable mass. ",
        "conclusions": ""
    },
    "0710/0710.4282_arXiv.txt": {
        "abstract": "{Charge transfer (or exchange) reactions between hydrogen atoms and protons in collisionless shocks of supernova remnants (SNRs) are a natural way of producing broad Balmer, Lyman and other lines of hydrogen.} {We wish to quantify the importance of shock-induced, non-thermal hydrogen emission from SNRs in young galaxies.} {We present a method to estimate the luminosity of broad ($\\sim 1000$ km s$^{-1}$) Ly$\\alpha$, Ly$\\beta$, Ly$\\gamma$, H$\\beta$ and P$\\alpha$ lines, as well as the broad and narrow luminosities of the two-photon (2$\\gamma$) continuum, from existing measurements of the H$\\alpha$ flux.  We consider cases of $\\beta=0.1$ and 1, where $\\beta \\equiv T_e/T_p$ is the ratio of electron to proton temperatures.  We examine a modest sample of 8 proximate, Balmer-dominated SNRs from our Galaxy and the Large Magellanic Cloud.  The expected broad Ly$\\alpha$ luminosity per object is at most $\\sim 10^{36}$ erg s$^{-1}$.  The 2$\\gamma$ continuum luminosities are comparable to the broad H$\\alpha$ and Ly$\\alpha$ ones.  We restrict our analysis to homogenous and static media.} {Differences in the Ly$\\alpha$/H$\\alpha$ and Ly$\\beta$/H$\\alpha$ luminosity ratios between the $\\beta=0.1$ and 1 cases are factors $\\sim 2$ for shock velocities $1000 \\lesssim v_s \\lesssim 4000$ km s$^{-1}$, thereby providing a direct and unique way to measure $\\beta$.  In principle, broad, ``non-radiative'' Ly$\\alpha$ from SNRs in young galaxies can be directly observed in the optical range of wavelengths.  However, by taking into consideration the different rates between core collapse and thermonuclear supernovae, as well as the duration we expect to observe such Ly$\\alpha$ emission from SNRs, we expect their contribution to the total Ly$\\alpha$ luminosity from $z \\sim 3$ to 5 galaxies to be negligibly small ($\\sim 0.001 \\%$), compared to the radiative shock mechanism described by Shull \\& Silk (1979).  Though broad, non-thermal Ly$\\alpha$ emission has never been observed, these photons are produced in SNRs and hence the non-radiative Ly$\\alpha$ luminosity is a part of the intrinsic Ly$\\alpha$ spectrum of young galaxies.} {} ",
        "introduction": "Observations of galaxies at high redshifts have revealed a broad class of Ly$\\alpha$-emitting galaxies at $z \\sim 3$ to 5 (e.g., Tapken et al. 2007).  The Ly$\\alpha$ emission from these objects is reaching us as light in the visible spectral band, enabling their study using large, ground-based optical telescopes, which in turn permits detailed spectroscopic studies of these galaxies.  Observations of quasars at $z \\sim 6$ (e.g., Fan et al. 2006) have revealed heavy elemental abundances exceeding solar values.  We know that at least some of the galaxies at $z \\sim 3$ to 5 have high abundances of heavy elements, facilitating the formation of dust.  In homogenous and static media, the dust particles impede the escape of Ly$\\alpha$ emission from gas-rich galaxies, due to the small mean free paths of the photons, low temperatures of the gas and ultimately high probabilities of absorption.  In clumpy media, dust can enhance the escape of Ly$\\alpha$ photons relative to the continuum (Neufeld 1991; Hansen \\& Oh 2006).  Broadening of Ly$\\alpha$ lines due to multiple scatterings is a slow process requiring a long diffusion time (though velocity fields in the interstellar medium may broaden the Ly$\\alpha$ lines and reduce the diffusion time).  Hence, there is special interest in the physical processes that are able to naturally produce extremely broad wings in Ly$\\alpha$ lines, which may permit the photons to leave the host galaxy without requiring many scatterings (but see \\S\\ref{sect:discussion}).  Among obvious mechanisms is the one at work in the unique massive binary SS433 (for a recent review, see Fabrika [2004]), with strongly blue- and redshifted H$\\alpha$ and H$\\beta$ lines, due to cooling and recombination of hydrogen in the baryon-dominated, precessing jet moving with velocity $\\sim 0.26 c$.  Such objects are very rare --- SS433 is the only such example in our Galaxy.  More well-known Galactic sources of H$\\alpha$ emission with broad line wings are the supernova remnants (SNRs) of Type Ia, emitting due to charge transfer (or ``charge exchange'') reactions between hydrogen atoms and protons in the blast wave penetrating the low-density ($\\sim 1$ cm$^{-3}$), ambient gas.  The widths of the H$\\alpha$ lines correspond to Doppler broadening with velocities up to $\\sim 5000$ km s$^{-1}$.  The same process should produce not only H$\\alpha$ emission, but photons in the Lyman series of hydrogen as well.  Recently, some of these SNRs were observed in Ly$\\beta$ using the {\\it FUSE} spacecraft (Korreck et al. 2004; Ghavamian et al. 2007, hereafter G07).  Knowledge of the cross sections of charge transfers to excited levels and excitation of the fast-moving hydrogen atoms permit us to find simple formulae relating the luminosities of SNRs in the broad H$\\alpha$ and Ly$\\alpha$ lines.  The Ly$\\alpha$ line should have a similar spectral distribution to the observed H$\\alpha$ one in the broad wings, because the optical depth of the SNR for broad photons is negligibly small and the optical depth for coherent scattering (in the distant Lorentzian wings) in interstellar gas is low.  We compile the existing data for core collapse and thermonuclear SNRs, including SNR 1987A (where the reverse shock is bright in the broad H$\\alpha$ line), and present their theoretically expected, broad Ly$\\alpha$ and Ly$\\beta$ luminosities.  For two objects, we present their expected broad Ly$\\gamma$, H$\\beta$ and P$\\alpha$ luminosities. Taking into account the supernova (SN) rates, the luminosities of the SNRs in H$\\alpha$ and the duration of their active phase (for the charge transfer mechanism described), we find that --- even without discussing the cosmological evolution of the SN rates --- the expected broad Ly$\\alpha$ is several orders of magnitude lower than the estimate of Shull \\& Silk (1979), who treated fully radiative SNRs with low metallicities and velocities (20 to 120 km s$^{-1}$).  We come to the conclusion that the contribution of both core collapse and thermonuclear SNRs to the Ly$\\alpha$ luminosity of young galaxies is negligibly small.  In \\S\\ref{sect:obs}, we gather a modest sample of 8 Galactic and Large Magellanic Cloud (LMC) remnants, and use them as a template for estimating the expected Ly$\\alpha$, Ly$\\beta$, Ly$\\gamma$, H$\\beta$ and P$\\alpha$ production.  In \\S\\ref{sect:ratios}, we compute the Ly$\\alpha$/H$\\alpha$, Ly$\\alpha$/Ly$\\beta$, Ly$\\beta$/H$\\alpha$, Ly$\\gamma$/H$\\alpha$, H$\\beta$/H$\\alpha$ and P$\\alpha$/H$\\alpha$ luminosity ratios.  We present our results in \\S\\ref{sect:results} and discuss their implications in \\S\\ref{sect:discussion}. ",
        "conclusions": "\\label{sect:discussion}  SNR 1987A is a unique example of a Balmer-dominated SNR.  By virtue of adiabatic expansion cooling, the SN ejecta comprises mostly neutral hydrogen; it rushes out at velocities $\\gtrsim 12,000$ km s$^{-1}$ (Michael et al. 2003; Heng et al. 2006).  The non-radiative H$\\alpha$ and Ly$\\alpha$ result from the interaction of the ejecta with the {\\it reverse shock} and not the blast wave (Heng 2007).  As SNR 1987A has a Type II origin, it is possible to produce Balmer and Lyman lines via this mechanism; this is obviously not possible with Type Ia's. Smith et al. (2005) have predicted that the H$\\alpha$ and Ly$\\alpha$ emission from the reverse shock of SNR 1987A is shortlived ($\\sim$ 2012 to 2014) and will be extinguished by the increasing flux of extreme ultraviolet (EUV) and X-ray photons traveling into the pre-shock region and ionizing the atoms --- pre-ionization.  This is marginal evidence that broad Ly$\\alpha$ from SNRs of a core collapse origin will be short-lived, i.e., $\\lesssim 100$ years.  In general, for this scenario to work, some interaction of the blast wave with the ambient material is needed, but if it is too strong the pre-shock gas becomes ionized (R. Chevalier 2007, private communication).  To further investigate the viability of the short-lived, non-radiative Ly$\\alpha$ hypothesis, we examine the sample of optically identified SNRs by Matonick \\& Fesen (1997), who studied an ensemble of 12 SNR samples from different galaxies, including the Small Magellanic Cloud (SMC), LMC, M31 and M33, with distances up to 7 Mpc.  In galaxies like NGC 2403, M81 and M101, the SNRs are associated with star-forming regions and most of them probably have a Type Ib/c origin.  In most cases, the measured H$\\alpha$ flux is $\\sim 10^{-15}$ erg cm$^{-2}$ s$^{-1}$ and the inferred luminosity is $\\sim 10^{36}$ erg s$^{-1}$.  Since Matonick \\& Fesen (1997) did not provide H$\\alpha$ line profiles, it is impossible to estimate the proportion of the H$\\alpha$ emission that is non-radiative.  Furthermore, their selection criterion is based on picking out objects with [S~{\\sc ii}]/H$\\alpha \\ge 0.45$, which will not detect SNRs with predominantly non-radiative H$\\alpha$ emission.  Shull \\& Silk (1979) computed the temporally-averaged Ly$\\alpha$ luminosity from radiative shocks of a population of Type II SNRs, assuming low metallicities, to be \\begin{equation} L_{\\rm{SS79}} = 3 \\times 10^{43} \\mbox{ erg s}^{-1} E^{3/4}_{51} n^{-1/2}_0 \\dot{N}_{\\rm{SN}}, \\end{equation} where $\\dot{N}_{\\rm{SN}}$ is the number of supernovae (SNe) a year.  They considered SNRs in both the ST and the PDS stages, and $v_s = 20$ to 120 km s$^{-1}$.  Charlot \\& Fall (1993) remark that the numerical coefficient in the preceding equation is about 40\\% lower if one assumes solar metallicity.  A very conservative upper limit on the broad Ly$\\alpha$ from the Matonick \\& Fesen (1997) samples can be obtained if one generously allows for all of the H$\\alpha$ to be broad, for the shock velocities to be low ($\\sim 500$ km s$^{-1}$) such that $\\Gamma_{\\rm{Ly}\\alpha/\\rm{H}\\alpha} \\sim 100$, and for the non-radiative emission to last $\\sim 10^4$ years.  Even in this very unlikely scenario, $L_{\\rm{Ly}\\alpha} \\sim 10^{42}$ erg s$^{-1}$ is only about $0.1 L_{\\rm{SS79}}$.  Hence, our charge transfer mechanism is not energetically competitive.  There is the possibility a SNR can produce both radiative and non-radiative components of H$\\alpha$.  Well-known examples are Kepler (Fesen et al. 1989; Blair, Long \\& Vancura 1991) and RCW 86 (Long \\& Blair 1990; Smith 1997).  There is also the possibility that the non-radiative emission from the SNR is inhibited.  For example, Foster (2005) observed and studied the Galactic SNR 3C 434.1 ($t_{\\rm{age}} \\approx 25,000$ yr; $d = 4.5 \\pm 0.9$ kpc; possible Type Ib/c), which formed inside the eastern portion of a pre-existing stellar-wind bubble of interior density $\\sim 0.1$ cm$^{-3}$.  Strong H$\\alpha$ emission ($6.1 \\pm 0.4 \\times 10^{36}$ erg s$^{-1}$) is measured from the eastern side; it is believed to be from a radiative shock.  Being farther away from the western wall of the bubble, the shock on the western side is essentially still in free expansion and produces no measurable, non-radiative H$\\alpha$.  Our SNR sample and the considerations of SNR 1987A lead us to believe that if the short-lived emission contribution from Type Ib/c and Type II SNRs in young galaxies exists, it has a luminosity of \\begin{equation} L_{\\rm{Ly}\\alpha,\\rm{CC}} \\sim 10^{38} \\mbox{ erg s}^{-1} ~t_{\\rm{emit},2} \\dot{N}_{\\rm{SN}}, \\end{equation} where $t_{\\rm{emit}} = t_{\\rm{emit},2} 100$ years is the length of time we expect core collapse SNRs to produce shock-induced Ly$\\alpha$ emission.  On the other hand, thermonuclear SNRs are expected to have  $t_{\\rm{emit}} = t_{\\rm{emit},4} 10^4$ years $\\sim t_{\\rm{PDS}}$.  However, they are also believed to be much scarcer at high redshifts.  For example, Dahlen et al. (2004) estimate that only 5\\% to 7\\% of available progenitors explode as Type Ia SNRs.  Therefore, the expected luminosity is \\begin{equation} L_{\\rm{Ly}\\alpha,\\rm{Ia}} \\sim 10^{38} \\mbox{ erg s}^{-1} ~t_{\\rm{emit},4} \\dot{N}_{\\rm{SN},-2}, \\end{equation} where $\\dot{N}_{\\rm{SN},-2}$ is the number of SN per year in units of 0.01.  We conclude that for both core collapse and thermonuclear SNRs, the expected luminosity from broad Ly$\\alpha$ is only a $\\sim 0.001\\%$ effect, compared to the mechanism of Shull \\& Silk (1979).  Ly$\\alpha$ line luminosities from $z \\sim 3$ to 5 galaxies have been observationally determined to be $\\sim 10^{42}$ to $10^{43}$ erg s$^{-1}$ (e.g., Saito et al. 2007), in general agreement with theoretical expectations.  In addition, the lifetime of an emitting atom is approximately the length of time corresponding to one atomic length scale, and is only $t_{\\rm mfp} \\sim l_a/v_s \\sim 10^7 n^{-1}_0 v^{-1}_{s,8}$ s, where $v_{s,8} = v_s/1000$ km s$^{-1}$ (H07).   We have restricted our analysis to homogeneous and static media.  Though broad, non-thermal Ly$\\alpha$ emission has never been observed, these photons {\\it are} produced in SNRs and hence the non-radiative Ly$\\alpha$ luminosity is a part of the intrinsic Ly$\\alpha$ spectrum of young galaxies.  The optical depth for a broad photon in the line wings is (Verhamme, Schaerer \\& Maselli 2006) \\begin{equation} \\tau \\sim 0.26 T^{-1/2}_4 N_{{\\rm H},20} v^{-2}_{w,8} b_{12.85}, \\end{equation} where $T = 10^4 T_4$ K and $N_{\\rm H} = N_{{\\rm H},20} 10^{20}$ cm$^{-2}$ are the temperature and obscuring hydrogen column density of the medium, respectively.  The turbulent velocity in the interstellar medium is $b = 12.85 b_{12.85}$ km s$^{-1}$ (Verhamme, Schaerer \\& Maselli 2006), while $v_w = 1000 v_{w,8}$ km s$^{-1}$ is the velocity of the emitting atom in the line wings.  Multiple scattering is important for $\\tau \\gtrsim 0.3$ (Chevalier 1986) and any realistic treatment of non-thermal Ly$\\alpha$ lines in a young galaxy has to include radiative transfer effects, which we have neglected in our analysis."
    },
    "0710/0710.3141_arXiv.txt": {
        "abstract": "We present a study of the globular cluster (GC) systems of nearby elliptical and S0 galaxies at a variety of wavelengths from the X-ray to the infrared. Our analysis shows that roughly half of the low mass X-ray binaries (LMXBs), that are the luminous tracers of accreting neutron star or black hole systems, are in clusters. There is a surprisingly strong correlation between the LMXB frequency and the metallicity of the GCs, with metal-rich GCs hosting three times as many LMXBs as metal-poor ones, and no convincing evidence of a correlation with GC age so far. In some of the galaxies the LMXB formation rate varies with GC color even within the red peak of the typical bimodal cluster color distribution, providing some of the strongest evidence to date that there are metallicity variations within the metal-rich GC peak as is expected in a hierarchical galaxy formation scenario. We also note that any analysis of subtler variations in GC color distributions must carefully account for both statistical and systematic errors. We caution that some published GC correlations, such as the apparent 'blue-tilt' or mass-metallicity effect might not have a physical origin and may be caused by systematic observational biases. ",
        "introduction": "High resolution Chandra X-ray images of nearby ellipticals and S0s have resolved large numbers of point sources,  confirming a long-standing suggestion that the hard X-ray emission in many of these galaxies is predominantly from X-ray binaries. In early type galaxies most of the bright, L$_X$$\\gtrsim$10$^{37}$ erg s$^{-1}$ sources seen in typical Chandra observations must be low mass X-ray binaries, binary systems comprising a neutron star or black hole accreting via Roche lobe overflow from a low mass companion, since they generally have stellar populations that are at least a few Gyrs old.  An important characteristic of LMXBs is that they are disproportionately abundant in globular clusters. Even though GCs account for $\\lesssim$0.1\\% of the stellar mass in the Galaxy, they harbor about 10\\% of the L$_X$$\\gtrsim$10$^{36}$ erg s$^{-1}$ LMXBs , indicating  a probability of hosting a LMXB that is at least two orders of magnitude larger than for field stars. This has long been attributed to efficient formation of LMXBs in clusters due to dynamical interactions in the core. Early type galaxies are ideal for studies of the LMXB-GC link as they are particularly abundant in globular clusters. The identification of LMXBs with these simple stellar systems that have well defined properties such as metallicity and age provides a unique opportunity to probe the effects of these parameters on LMXB formation and evolution.       \\begin{figure}[b] \\begin{center} \\includegraphics[width=5in]{kundu_fig1.ps} \\caption{Top: The V-I colors of GCs vs. distance from the center of NGC 4472 and GC color distribution. LMXB-GC matches are indicated by filled circles/bins.  Bottom: V magnitude of globular clusters vs. half light radius and the globular cluster luminosity function. LMXBs are preferentially located in the brightest, most metal-rich GCs. There is a weak anti-correlation with GC half-light radius and no obvious correlation with galactocentric distance. Each of these broad correlations (or lack thereof) have been confirmed in other galaxies. } \\label{fig1} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0710/0710.1528_arXiv.txt": {
        "abstract": "We present the general relativistic calculation of the energy release associated with a first order phase transition (PT) at the center of a rotating neutron star (NS). The energy release, $E_{\\rm rel}$, is equal to the difference in mass-energies between the initial (normal) phase configuration and the final configuration containing a superdense matter core, assuming constant total baryon number and the angular momentum. The calculations are performed with the use of precise pseudo-spectral 2-D numerical code; the polytropic equations of state (EOS) as well as realistic EOSs (Skyrme interactions, Mean Field Theory kaon condensate) are used. The results are obtained for a broad range of metastability of initial configuration and size of the new superdense phase core in the final configuration. For a fixed ``overpressure'', $\\delta\\overline{P}$, defined as the relative excess of central pressure of a collapsing metastable star over the pressure of the equilibrium first-order PT, the energy release up to numerical accuracy {\\it does not depend} on the stellar angular momentum and coincides with that for nonrotating stars with the same $\\delta\\overline{P}$. When the equatorial radius of the superdense phase core is much smaller than the equatorial radius of the star, analytical expressions for the $E_{\\rm rel}$ can be obtained: $E_{\\rm rel}$ is proportional to $(\\delta\\overline{P})^{2.5}$ for small $\\delta\\overline{P}$. At higher $\\delta\\overline{P}$, the results of 1-D calculations of $E_{\\rm rel}(\\delta\\overline{P})$ for non-rotating stars reproduce with very high precision exact 2-D results for fast-rotating stars. The energy release-angular momentum independence for a given overpressure holds also for the so-called ``strong'' PTs (that destabilise the star against the axi-symmetric perturbations), as well as for PTs with ``jumping'' over the energy barrier. ",
        "introduction": "Many theories of dense matter predict that at some density larger than the nuclear saturation density, a phase transition (PT) to some ``exotic'' state (i.e. boson condensate or quark deconfinement) occurs; for review see e.g., \\cite{NS1,Glend.book,weber99}). A first-order PTs are particularly interesting from the astrophysical and observational point of view, because are associated with a meta-stable state of dense matter. One can thus expect, in the case of PTs occurring in the interior of NSs, the release of non-negligible amount of energy. Here we will focus on the basic features of the energy release-angular momentum independence; for complete description we refer the reader to \\cite{erot,erots}. The text is arranged as follows: in Sect.~\\ref{calculations} we briefly describe the results of calculations. Sect.~\\ref{discussion} contains conclusions and open questions.  \\begin{figure}[t] \\begin{center} \\psfig{file=bejger1.eps,width=5.in} \\end{center} \\caption{ Left panel: Strong and weak PTs on mass-radius diagram (dotted line marks unstable configurations). Right panel: the energy release $E_{\\rm rel}$ vs the over-pressure parameter $\\delta\\bar{P}$ for strong (two upper curves) and weak (lowest curve) PTs. Points are colored differently for different total angular momenta $J=(0,0.1,...,1.3)\\times GM^2_\\odot/c$ and follow the curve for the $J=0$ (non-rotating) configurations.} \\label{reffig1} \\end{figure} ",
        "conclusions": "\\label{discussion} Our numerical calculations show the $E_{\\rm rel}(J)$ independence during the weak as well as strong first-order PTs, for large rotation rates and large oblatnesses of the stars: it is therefore not an property of slow rotating stars only. The independence holds also for PTs with negative over-pressure $\\delta\\bar{P}$ i.e. when the configuration ``jumps'' over the energy barrier to reach another stable configuration. For small positive $\\delta\\bar{P}$ analytical relations were found: $E_{\\rm rel} \\propto (\\delta\\overline{P})^{2.5}$. The energy release $E_{\\rm rel}\\sim 10^{51}-10^{52}$ erg is an absolute upper bound on the energy which can released in such PT. In astrophysical situation, the energy can be distributed between stellar pulsations, gravitational radiation, heating of stellar interior, X-ray emission from the neutron star surface, and even a gamma-ray burst.  Currently there is no mathematical proof of the $J$-independence of $E_{\\rm rel}$. It may be interesting to look at the problem from the ``thermodynamical'' point of view. We write the total energy of the star (gravitational mass $M$) as \\be \\label{e:M_A_J_p} M = \\frac{\\mu}{u^t} A + 2 \\, \\Omega J + 2 \\int_{\\Sigma_t} P\\,  N\\sqrt{\\gamma}\\,  d^3x , \\ee where $\\mu$ is the baryon chemical potential, $u^t$ the time component of the fluid 4-velocity, $A$ the number of baryons in the star, $\\Omega$ the angular velocity, $P$ the fluid pressure, $N$ the lapse function and $\\sqrt{\\gamma}\\, d^3x$ the covariant volume element in the constant $t$ hypersurface $\\Sigma_t$. Eq.~(\\ref{e:M_A_J_p}) is valid for any axi-symmetric stationary and rigidly rotating fluid star which obeys a barotropic EOS, as established by Bardeen \\& Wagoner in 1971 \\cite{BardeW71}. Since $\\mathcal{C}$ and $\\mathcal{C}^*$ have the same baryon number $A$ and the same angular momentum $J$, we get the energy release \\be \\label{e:DE} E_{\\rm rel} = \\Delta E = A \\, \\Delta\\! \\left( \\frac{\\mu}{u^t} \\right) + 2\\,  J \\, \\Delta\\Omega + 2 \\left[ \\int_{\\mathcal{C}} P\\,  N\\sqrt{\\gamma}\\,  d^3x - \\int_{\\mathcal{C}^*} P\\,  N\\sqrt{\\gamma}\\,  d^3x \\right] , \\ee with \\be \\Delta\\! \\left( \\frac{\\mu}{u^t} \\right) := \\left. \\frac{\\mu}{u^t} \\right| _{\\mathcal{C}} - \\left. \\frac{\\mu}{u^t} \\right| _{\\mathcal{C}^*} ,~~~ \\Delta \\Omega := \\left. \\Omega \\right| _{\\mathcal{C}} - \\left. \\Omega \\right| _{\\mathcal{C}^*} . \\ee The terms on the right-hand-side of Eq.~\\ref{e:DE} are of comparable magnitude in their contribution to the energy release, so the possible ``smallness'' of some of them in relation to the others is not responsible for the $E_{\\rm rel}(J)$ independence. We will address the problem of mathematically proving this property in the near future."
    },
    "0710/0710.2237_arXiv.txt": {
        "abstract": "{A wealth of observations of CO in absorption in diffuse clouds has accumulated in the past decade at $uv$ and mm-wavelengths} {Our aims are threefold: a) To compare the  $uv$ and mm-wave results; b) to interpret  \\coth\\ and \\cotw\\ abundances in terms of the physical processes which separately and jointly determine them; c) to interpret observed J=1-0 rotational excitation and line brightness in terms of ambient gas properties.} {A simple phenomenological model of CO formation as the immediate descendant of quiescently-recombining \\hcop\\ is used to study the accumulation, fractionation and rotational excitation of CO in more explicit and detailed models of \\HH-bearing diffuse/H I clouds} {The variation of N(CO) with N(\\HH) is explained by quiescent recombination of a steady fraction n(\\hcop)/n(\\HH) = $2\\times 10^{-9}$. Observed N(\\cotw))/N(\\coth) ratios generally do not require a special chemistry but result from competing processes and do not provide much insight into the local gas properties, especially the temperature. J=1-0 CO line brightnesses directly represent N(CO), not N(\\HH), so the CO-\\HH\\ conversion factor varies widely; it attains typical values at N(\\cotw) $\\la 10^{16}\\pcc$. Models of CO rotational excitation account for the line brightnesses and CO-\\HH conversion factors but readily reproduce the observed excitation temperatures and optical depths of the rotational transitions only if excitation by H-atoms is weak -- as seems to be the case for the very most recent calculations of these excitation rates.} { Mm-wave and $uv$ results generally agree well but the former show somewhat more enhancement of \\THC\\ in \\coth . In any case, fractionation may seriously bias \\TWC/\\THC\\ ratios measured in CO and other co-spatial molecules.  Complete C$\\rightarrow$CO conversion must occur over a very narrow range of \\AV\\ and N(\\HH) just beyond the diffuse regime.  For  N(\\HH) $< 7\\times10^{19}\\pcc$ the character of the chemistry changes inasmuch as CH is generally undetected while CO suffers no such break.  } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.2371_arXiv.txt": {
        "abstract": "We analyze \\WMAP 3 year data using the one-point distribution functions to probe the non-Gaussianity in the Cosmic Microwave Background (CMB) Anisotropy data. Computer simulations are performed to determine the uncertainties of the results. We report the non-Gaussianity parameter $\\fnl$ is constrained to $26<\\fnl<82$ for Q-band, $12<\\fnl<67$ for V-band, $7<\\fnl<64$ for W-band and $23<\\fnl<75$ for Q+V+W combined data at 95\\% confidence level (CL). ",
        "introduction": "Non-Gaussianity is one of the most important tests of models of the inflation. Among the various theoretical models on the inflation, slow-roll inflation is currently most lively being studied. There are various predictions on the magnitude of non-Gaussianity based on the simple model of slow-roll inflation and its extensions, ranging from undetectably tiny values to large enough values to be detectable with currently available data \\cite{barnabycline,battefeldeasther,calcagni,creminelli,bartolo.et.al}. On the other hand, observational works have claimed both detection and non-detection of non-Gaussianity (for reviews on recent works, see \\cite{komatsu.et.al,Spergel.et.al,troia.et.al,creminelli.et.al,gaztanagawagg}). Among the popular techniques for detecting non-Gaussianity are one-point distribution function fitting, bispectrum, trispectrum and Minkowski functionals. Here, we investigate the one-point distribution functions to probe primordial non-Gaussianity in the CMB anisotropy data. An observed CMB anisotropy at a direction ($\\delta T_{obs}$) can be regarded as the superposition of three parts: physical fluctuation of cosmic origin ($\\delta T_p$), instrumental noise ($\\delta T_n$), and foreground emissions ($T_{fg}$). Since the foreground templates are separately prepared, we start with foreground-removed data of which the CMB anisotropy can be decomposed into two uncorrelated components, \\begin{equation}\\label{eq1} \\delta T=\\delta T_{obs}-T_{fg}=\\delta T_{p}+\\delta T_n. \\end{equation} The primary source for the cosmic fluctuation of CMB at the large scale is attributed to the Sachs-Wolfe effect which is again triggered by the primordial curvature perturbation. The curvature perturbation $\\Phi$ by primordial seed during the inflation is transferred to CMB anisotropy with the relation \\begin{equation}\\label{eq2} \\frac{\\delta T_p\\fpr{\\bfx}}{T_0}=\\eta_{t}\\Phi\\fpr{\\bfx} \\end{equation} where $T_0=2.725$ K, the thermodynamic temperature of the CMB today, and $\\eta_{t}$ is the radiation transfer function. For the super-horizon scale, we take $\\eta_{t}=-1/3$ from the Sachs-Wolfe effects. At the first-order of perturbation, we may replace $\\Phi =\\Phi_g$, where $\\Phi_g$ is an auxiliary random Gaussian field with its mean $\\fang{\\Phi_g}=0$ and its variance denoted by $\\fang{\\Phi_g^2}$. When the second-order perturbation is considered, it is conventional to prescribe the nonlinear coupling of the curvature perturbation as \\cite{komatsuspergel} \\begin{equation}\\label{eq3} \\Phi (\\bfx )\\simeq\\Phi_g(\\bfx )+\\fnl\\fpr{\\Phi_g^2(\\bfx )-\\fang{\\Phi_g^2}} \\end{equation} where $\\fnl$ is the non-Gaussianity parameter. The second term in \\eqref{eq3} is responsible for the non-Gaussianity of the primordial fluctuation. Then, the probability distribution function of the non-Gaussian field $\\Phi$ can be derived as \\begin{eqnarray} f_{\\Phi}(\\Phi)&=&\\int f_G(\\Phi_g)\\delta_D \\fsq{\\Phi-\\Phi_g-\\fnl\\fpr{\\Phi_g^2-\\fang{\\Phi_g^2}}}d\\Phi_g\\nonumber\\\\ &=&\\frac{1}{\\sqrt{2\\pi\\fang{\\Phi_g^2}\\fnl^2\\fpr{\\Phi_+-\\Phi_-}^2}}\\nonumber\\\\ & &\\times\\fsq{\\exp\\fpr{-\\frac{\\Phi_+^2}{2\\fang{\\Phi_g^2}}} +\\exp\\fpr{-\\frac{\\Phi_-^2}{2\\fang{\\Phi_g^2}}}}\\label{eq4} \\end{eqnarray} where $\\Phi_{\\pm}$ are defined by \\begin{equation}\\label{eq5} \\Phi_{\\pm}=\\frac{1}{2\\fnl} \\fsq{-1\\pm\\sqrt{1+4\\fnl\\Phi+4\\fnl^2\\fang{\\Phi_g^2}}} \\end{equation} and $\\Phi$ has to be limited by the reality of $\\Phi_{\\pm}$ as \\begin{equation}\\label{eq6} \\fnl\\Phi > -\\frac{1}{4}-\\fnl^2\\fang{\\Phi_g^2}. \\end{equation} $\\fang{\\Phi_g^2}$ can be expressed in terms of $\\eta_t$, $T_0$ and $\\sigma_{CMB}$, \\begin{equation}\\label{eq7} \\fang{\\Phi_g^2}=\\frac{1}{4\\fnl^2}\\fsq{-1+\\sqrt{1+8 \\fpr{\\frac{\\fnl\\sigma_{CMB}}{\\eta_t T_0}}^2}}. \\end{equation} For a pixelized CMB anisotropy data set, the probability distribution function for Gaussian instrumental noise becomes \\begin{equation}\\label{eq8} f_N(\\delta T_n)=\\frac{1}{N_{pix}}\\sum_{i=1}^{N_{pix}} \\frac{1}{\\sqrt{2\\pi\\sigma_0^2/n_i}} \\exp\\fsq{-\\frac{\\delta T_n^2}{2\\sigma_0^2/n_i}} \\end{equation} where $n_i$ is the effective number of measurements at the $i_{th}$ pixel and $\\sigma_0$ represents the dispersion of the instrumental noise per observation ($\\sigma_0$=2.1898, 3.1249, 6.5112 mK for Q, V, W-band, respectively \\cite{Limon.et.al}). Now, it is straightforward to express the probability density function for $\\delta T$ in an integral form, \\begin{eqnarray} f(\\delta T)&=&\\int f_{\\delta T_p}(\\delta T_p)f_N(\\delta T_n)\\nonumber\\\\ & &\\times\\delta_D\\fpr{\\delta T-\\delta T_p-\\delta T_n} d\\delta T_p d\\delta T_n\\nonumber\\\\ &=&\\int f_{\\Phi}(\\Phi )f_N(\\delta T_n)\\nonumber\\\\ & &\\times\\delta_D\\fpr{\\delta T-\\eta_t T_0\\Phi-\\delta T_n}d\\Phi d\\delta T_n. \\label{eq9} \\end{eqnarray} The probability density function derived in \\eqref{eq9} explicitly contains the non-Gaussianity parameter $\\fnl$, and it can serve as the prediction of one-point distribution function with a given $\\fnl$ for a (ideally) foreground-removed CMB anisotropy data set to estimate the magnitude of deviation from Gaussian distribution in a quantitative manner. ",
        "conclusions": "We developed an algorithm that uses the one-point distribution function to investigate the non-Gaussianity of CMB anisotropy data, and applied it to \\WMAP 3 year data. We found that the null result ($\\fnl$=0) is manifestly excluded at 95\\% CL. The estimated magnitude of non-Gaussianity parameter is $23<\\fnl<75$ at 95\\% CL and $9<\\fnl<88$ at 99\\% CL for the (Q+V+W)-combined map. Since the quadratic term in \\eqref{eq3} takes a generic form of Taylor series for a perturbative expansion, it is a good possibility that the observed non-Gaussianity in this work is a combined effects of various physical processes, while the primordial seeds are very likely to be the leading one. There are two premises we have taken in developing the algorithm, which, provided they are not precise enough, could cause non-Gaussianity of not cosmic but systematic origin: (1) the probability distribution function of the instrumental noise for each pixel is centered at zero, and (2) the foreground emissions are removed efficiently enough in the foreground-removed maps. The first condition can be broken when the thermal and radiation environments of the \\WMAP satellite in its orbit are taken into account, while the \\WMAP team assessed they are insufficient to influence the science data \\cite{Limon.et.al}. So, we tested the effects of the alternative noise distributions with a random mean in each of the Gaussian distribution in \\eqref{eq8} and the algorithm was not misled to show non-Gaussianity within the statistical error. It is difficult to directly estimate how much residual foreground emissions after foreground subtraction would affect the one-point distribution function. We solely rely on the quality of foreground templates and it is remarkably successful, showing that the observed total Galactic emission matches the model to less than 1\\% \\cite{bennett.et.al2,Hinshaw.et.al}. We also analyzed simulated maps which are (Gaussian map + foreground templates), and all the templates for Q, V and W-channel showed negative values of the non-Gaussianity parameter with $\\fabs{\\fnl}\\sim\\mathcal{O}\\fpr{10^1}$ at the resolution $N_{\\mathrm{side}}=512$."
    },
    "0710/0710.1552_arXiv.txt": {
        "abstract": "The accreting millisecond pulsar SAX J1808.4-3658 may be a transition object between accreting X-ray binaries and millisecond radio pulsars.  We have constrained the thermal radiation from its surface through XMM-Newton X-ray observations, providing strong evidence for neutrino cooling processes from the neutron star core.  We have also undertaken simultaneous X-ray and optical (Gemini) observations, shedding light on whether the strong heating of the companion star in quiescence may be due to X-ray irradiation, or to a radio pulsar turning on when accretion stops. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.0762_arXiv.txt": {
        "abstract": "We use low-degree acoustic modes obtained by the BiSON  to estimate the main-sequence age $t_\\odot$ of the Sun. The calibration is accomplished by linearizing the deviations from a standard solar model the seismic frequencies of which are close to those of the Sun.  Formally, we obtain the preliminary value $t_\\odot=4.68\\pm0.02\\,$Gy, coupled with an initial heavy-element abundance $Z=0.0169\\pm0.0005$.  The quoted standard errors, which are not independent, are upper bounds implied under the assumption that the standard errors in the observed frequencies are independent. ",
        "introduction": "Seismological calibration of stellar models against observed frequencies of low-degree modes was first discussed more than two decades ago \\citep{jcd84, jcd88, ulrich86, dog87}, and can be regarded as a means of determining the main-sequence age of the Sun \\citep{guenther89, dog_nov90, guenther_demarque97, weiss_schlattl98, wd99, dog01, bon_schlat_pat02}. The procedure is to match certain appropriate seismic signatures of theoretical frequencies determined on a grid of stellar models with corresponding signatures obtained from the observations.  The signatures are chosen to reflect principally the properties of the energy-generating core, where nuclear transmutation leaves behind an augmenting concentration of helium, lowering the sound speed relative to the environs and thereby providing a diagnostic of age.  But the signatures are also susceptible to other properties of the stellar interior, which must be eliminated before a robust outcome can be achieved.  For example, although the so-called small frequency separation is sensitive predominantly to the evolving stratification of the core, its dependence on the zero-age chemical abundances plays a significant contaminating role.  Therefore it behoves us to seek an additional diagnostic to attempt to measure abundance separately. For given relative abundances of the heavy elements, the total absolute heavy-element abundance $Z$ and the $^4\\rm{He}$ abundance $Y$ are related by the requirement that the model has the observed luminosity and radius, principally the former.  Therefore we need to   % aim at detecting only one of them. Here we use a signature indicative of the abrupt variation of the first adiabatic exponent $\\gamma_1$ induced by the ionization of helium.  \\begin{figure} \\includegraphics[height=.42\\textheight]{hg_fig1.eps} \\caption{Top left: The symbols (with error bars obtained under the assumption that the raw frequency errors are independent) represent second differences, $\\Delta_2\\nu$, of low-degree solar frequencies from BiSON. Top right: The symbols are second differences $\\Delta_2\\nu$ of adiabatic pulsation eigenfrequencies of solar Model S. The solid curve in both panels is the diagnostic\\, (\\ref{eq:delnu}) -- (\\ref{eq:secdiff}), whose eleven parameters have been adjusted to fit the data optimally. Bottom: The symbols denote contributions $\\delta\\nu$ to the frequencies produced by the acoustic glitches of the Sun (left panel) and Model S (right panel).} \\end{figure} ",
        "conclusions": ""
    },
    "0710/0710.2290_arXiv.txt": {
        "abstract": "We study the possibility to extract the multipolar moments of an underlying distribution from a set of cosmic rays observed with non-uniform or even partial sky coverage. We show that if the degree is assumed to be upper bounded by $L$, each multipolar moment can be recovered whatever the coverage, but with a variance increasing exponentially with the bound $L$ if the coverage is zero somewhere. Despite this limitation, we show the possibility to test predictions of a model without any assumption on $L$ by building an estimate of the covariance matrix seen through the exposure function. ",
        "introduction": "Anisotropy in the arrival directions of cosmic rays is a major observable to understand their origin. Magnetic fields bend their trajectories in such a way that transport of cosmic rays is mainly diffusive up to high energies: this makes their angular distribution isotropic. Nevertheless, above the so-called knee of cosmic rays up to the ankle, there are predictions for small but increasing anisotropies with energy, predictions which of course depend on the regular and the turbulent components of the assumed galactic magnetic field, as well as the assumed distribution of sources and composition of cosmic rays \\cite{ptuskin,candia}. Further, at ultra-high energies, cosmic ray arrival directions are expected to be less and less smeared out by galactic and extragalactic magnetic fields, leading to a possible extraction of informations about the position of the sources \\cite{isola,sigl,eric,dolag,demarco}. Hence, it is clear that any evidence for an anisotropy, or any limit on anisotropies in the cosmic ray locations observed by experiments are among the most important constraints upon models.  The multipole expansion up to a given order $L$ is a powerful tool to study the structures standing out the noise down to an angular scale $\\approx \\pi/L$, whatever the shape of the underlying celestial pattern. In practice, the number of significant coefficients is limited by the angular resolution of the detector and, in the other hand, by the available statistics of observation. However, ground based experiments cover a limited range in declination, so that it is impossible to apply off the shelf the formalism of multipole moments: anyone of the coefficients may be modified in an unpredictable way by the unseen part of the sky. Methods have been developped to study the CMB with an incomplete coverage \\cite{gorski,wright,tegmark,mortlock}, but here we are faced to a different problem: we cannot suppose a priori that the distribution of cosmic rays is described by a power spectrum, because we want to detect possible non-isotropic structures, a priori unknown. In other terms, the information carried by the $a_{\\ell m}$ cannot be reduced to the only knowledge of the $C_{\\ell}$.   One purpose of this paper is to study the possibility of estimating the multipole moments of a distribution of points over a sphere in case of a non-uniform or even a partial coverage of the sky, together with the limitations of such an approach. The estimation of dipoles and quadrupoles was studied in \\cite{sommers,julien,silvia}. Here, we use the moments of the observed distribution on a set of orthogonal functions: either the spherical harmonics themselves, or a set of functions tailored on the coverage function. With these two different methods, we show that the interference between the modes induced by the the non-uniformity or the hole of the coverage can be removed assuming a bounded expansion in the conjugate space, allowing to recover the underlying multipole moments. However, in accordance with the simple intuition that it is impossible to describe the unseen part of the sky, we point out that the uncertainty on the recovered coefficients increases with the assumed bound $L$ of the expansion. We show that the larger the hole in the coverage of the sky, the faster  the increase of uncertainty with $L$. After some general considerations about the description of point processes on a sphere in Section 2, Sections 3 and 4 are dedicated to these methods whereas Section 5 illustrates them with some examples.  Because of the incomplete knowledge of the distribution of cosmic ray sources, and the stochastic nature of the propagation through magnetic fields, the anisotropies we want to characterize are not reducible to explicit models: they may be interpreted as a particular realization of a random process. This means that some model predictions are better expressed as average values of the coefficients, with their covariance matrix. This matrix is not necessarily diagonal to describe the physics we are interested in, contrary to the case of a power spectrum. We show in Sect.6 that under reasonable assumptions, an estimate can be performed with a partial sky coverage, evading the problem of setting a bound to the expansion. ",
        "conclusions": "To cope with a partial sky coverage, a formalism using the computation of moments on orthogonal functions was developed to recover the angular distribution of the incident flux from a sample of $N$ observed points. If the multipolar expansion is assumed to be upper bounded by $\\ell\\leq L$, the coefficients $a_{\\ell m}$ may be estimated with a variance proportional to $1/N$ as usually, with a penalty factor increasing exponentially with $L$ if there is a hole in the coverage (but stabilizing rapidly if the coverage is nowhere vanishing, even highly non-uniform). Two methods were tested, giving similar results, and practically the same variances.  Statistical tests based on likelihood ratios may be built to check an hypothesis on the distribution, for example a given bound $\\ell \\leq L$. In any case, it is possible to express predictions of a model in terms of coefficients which can be computed without any assumption on $L$, and tested against the moments found with a sample of observed points.  The methods presented in this paper may be applied any cosmic ray dataset, provided that the arrival directions and the coverage of the sky are known within a reasonable precision."
    },
    "0710/0710.0540_arXiv.txt": {
        "abstract": "We present the discovery of a silicate disc at the centre of the planetary nebula Mz3 (the Ant). The nebula was observed with MIDI on the Very Large Telescope Interferometer (VLTI). The visibilities obtained at different orientations clearly indicate the presence of a dusty, nearly edge-on disc in the heart of the nebula. An amorphous silicate absorption feature is clearly seen in our mid-IR spectrum and visibility curves. We used radiative transfer Monte Carlo simulations to constrain the geometrical and physical parameters of the disc. We derive an inner radius of 9 AU ($\\sim$6mas assuming D=1.4kpc). This disc is perpendicular to, but a factor of $10^{3}$ smaller than the optical bipolar outflow. ",
        "introduction": "\\label{sec:1} Complex phenomena perturb the mass-ejection in the late stages of stellar evolution: stellar magnetic fields, rotation or binarity are often invoked. Those mechanisms can lead to the creation of a circumstellar, dusty disc. The dust is created in the outer parts of the former stellar envelope. The Ant planetary nebula, or Mz3 is a bipolar nebula that has one of the most ``pinched'' waists. Its central object is suspected to be a binary but there are no stringent constraints. The temperature of the ionization source is about 30,000K and its distance is between $1-2$ kpc. Different expansion phases have shaped the current form of the nebula as seen in the Hubble Space Telescope images and long-slit spectroscopy \\cite{gue,san}. Former studies \\cite{smi1,smi2} have proposed the existence of a circumstellar disc, but due to its small size it could only be observed with high-angular resolution interferometry and only in the infrared. ",
        "conclusions": ""
    },
    "0710/0710.2635_arXiv.txt": {
        "abstract": "We have mapped the proto-binary source IRAS 16293--2422 in CO 2--1, $^{13}$CO 2--1, and CO 3--2 with the Submillimeter Array (SMA).  The maps with resolution of 1\\farcs5--5\\arcsec\\ reveal a single small scale ($\\sim$3000 AU) bipolar molecular outflow along the east-west direction. We found that the blueshifted emission of this small scale outflow mainly extends to the east and the redshifted emission to the west from the position of IRAS 16293A. A comparison with the morphology of the large scale outflows previously observed by single-dish telescopes at millimeter wavelengths suggests that the small scale outflow may be the inner part of the large scale ($\\sim$15000 AU) E--W outflow.  On the other hand, there is no clear counterpart of the large scale NE--SW outflow in our SMA maps.  Comparing analytical models to the data suggests that the morphology and kinematics of the small scale outflow can be explained by a wide-angle wind with an inclination angle of $\\sim$30\\degr--40\\degr~with respect to the plane of the sky.  The high resolution CO maps show that there are two compact, bright spots in the blueshifted velocity range.  An LVG analysis shows that the one located 1\\arcsec\\ to the east of source A is extremely dense, n(H$_2$)~$\\sim$10$^{7}$ cm$^{-3}$, and warm, T$_{\\rm kin} >$~55 K.  The other one located 1\\arcsec\\ southeast of source B has a higher temperature of T$_{\\rm kin} >$~65 K but slightly lower density of n(H$_2$)~$\\sim$10$^{6}$ cm$^{-3}$.  It is likely that these bright spots are associated with the hot core-like emission observed toward IRAS 16293. Since both two bright spots are blueshifted from the systemic velocity and are  offset from the protostellar positions, they are likely formed by shocks. ",
        "introduction": "Outflow phenomena have been recognized as an important phase in the star formation processes.  It is generally accepted that when stars are formed by gravitational infall, outflows transfer excess angular momentum out of the system  \\citep[e.g.,][]{shu1987}.  Outflows are observed at various wavelengths from the radio to the ultraviolet, among which the millimeter/submillimeter bands probe the molecular outflows.  Molecular outflows, recognized as high-velocity wings in CO and other molecular lines, are considered to be the ambient gas entrained or pushed by protostellar winds.  By observing molecular outflows, we can determine their physical, kinematic and chemical properties, trace the mass-loss history of protostars, and probe the early phases of the star formation process (\\citealt{arce2007}, and references therein).  IRAS 16293--2422 (hereafter referred to as I16293) is a well-studied proto-binary system located in the nearby molecular cloud complex L 1689. The distance to L 1689 is often assumed to be 160 pc \\citep{whittet1974}, however some recent studies suggested a distance of $\\sim$120 pc \\citep{degeus1989, knude1998}.  In this paper we adopt 120 pc as the distance.  The projected separation of the two protostars is $\\sim$5\\arcsec~\\citep{mundy1992}, which then corresponds to 600 AU.  I16293 is classified as a Class 0 young stellar object (YSO) since this system is undetectable at wavelengths shorter than 10 $\\mu$m and shows a spectral energy distribution (SED) with very low temperature \\citep{andre1994}.  The bolometric luminosity of I16293 is estimated to be $\\sim$36 $L_\\odot$ \\citep{correia2004}.  I16293 is known to have a quadrupolar outflow \\citep{walker1988, mizuno1990}; one bipolar pair extends along the east (blue) and west (red) (hereafter the E--W pair) directions, while the other one has an orientation from northeast (red) to southwest (blue) (hereafter the NE--SW pair) with an inclination angle of 30$^\\circ$--45$^\\circ$ \\citep{hirano2001} or 65$^\\circ$ \\citep{stark2004}.  The NE--SW pair is more collimated, and its axis crosses the position close to the SE continuum source IRAS 16293A (hereafter source A).  Therefore this pair has been thought to be powered by source A \\citep[e.g.,][]{walker1988, mundy1990}.  On the other hand, the E--W pair is not as well collimated and its axis, though not well defined, seems to pass to the north of source A, where the continuum source IRAS 16293B (hereafter source B) is the only possible driving YSO. Hence source B has often been assumed as the driving source of the E--W pair.  \\citet{stark2004} showed that source B has very narrow line widths, a low luminosity, and is not evidently associated with high-velocity gas. They interpreted source B as a T--Tauri star that drove the E--W pair and that the E--W pair is now a fossil flow.  However, these interpretations are based on observations at low angular resolution (i.e., $>$5\\arcsec), that are insufficient to resolve the two protostars and reveal the relation between the quadrupolar outflow and the binary.  In fact, using higher angular resolution observations \\citet{chandler2005} proposed an alternative interpretation.  They resolved source A into four components, two at 1~mm and another two at 7~mm, and suggested that two of them are responsible for driving the NE--SW and the E--W pair, respectively. In addition, \\citet{chandler2005} proposed that source B is actually much younger than previously interpreted and may have not yet begun the phase of mass loss.  Their proposition was, however, made based on (1) the dust continuum morphology and proper motions of source A, and (2) the velocity structures of emission of H$_2$CO, H$_2$S, and SO, rather than CO, which is a much better tracer of the overall outflow structure free from chemical peculiarities and shocks.  In order to study the molecular outflows in the vicinity of I16293 in more detail, we have carried out CO 2--1, $^{13}$CO 2--1, and CO 3--2 observations at higher angular resolutions using the Submillimeter Array\\footnote{The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics and is funded by the Smithsonian Institution and the Academia Sinica.} (SMA).  Our SMA observations provide sufficient angular and spectral resolution in the CO lines to resolve the kinematics and the structure of the outflow in the central region of I16293.  In this paper, we summarize the observational details in \\S2; present our SMA CO 2--1, $^{13}$CO 2--1, and CO 3--2  results in \\S3; discuss the outflow driving mechanism, physical parameters of the outflow, and the kinematics in the vicinity of I16293 in \\S4. ",
        "conclusions": "We have carried out CO 2--1 $^{13}$CO 2--1, and CO 3--2 observations of the proto-binary system IRAS 16293-2422 with a resolution of 1\\farcs5--5\\arcsec~using the Submillimeter Array (SMA), which are about a factor of 5 better than previous observations in these lines.  Our observations reveal the detailed structures of molecular outflows close to the binary system.  Our main results are summarized as follows:  1. The high resolution images of CO 2--1 and CO 3--2 show a compact bipolar outflow along the east-west direction in the vicinity of the binary.  The center of this small scale outflow is close to source A, suggesting that source A is  the driving source.  If we interpret this small scale outflow as the inner part of the large scale E--W pair of the quadrupolar outflow, the E--W pair is a currently active outflow rather than a fossil flow.  On the other hand, there is no clear counterpart of the large scale NE--SW pair in our interferometric maps.  2. The shape and velocity structure of the western redshifted lobe are well described by an analytical model suggesting that the small scale E--W pair could be driven by a wide-angle wind with an inclination angle of $\\sim$30$^\\circ$--40$^\\circ$.  3. There are two compact and bright spots at blueshifted velocities.  One is located 1\\arcsec\\ east of source A (labeled b1) and the other is 1\\arcsec\\ southeast of source B (labeled b2).  An LVG analysis shows that b1 is extremely dense, n(H$_2$)~$\\sim$10$^{7}$ cm$^{-3}$ and warm T$_{\\rm kin} >$~55 K, while component b2 has a higher temperature of T$_{\\rm kin} >$~65 K but slightly lower density of n(H$_2$)~$\\sim$10$^{6}$ cm$^{-3}$. It is likely that these bright spots are produced by means of shocks and are associated with hot core-like activity observed toward both source A and source B."
    },
    "0710/0710.0630_arXiv.txt": {
        "abstract": "Some aspects of disk-halo interactions for models of in and out of equilibrium disk galaxies are reviewed. Specifically, we focus on disk-halo resonant interaction without and in the presence of a gas component. Another issue is the disk growth within an assembling triaxial dark matter halo. We argue that while the triaxiality is the result of the merger process and the radial orbit instability, it is the developing chaos that damps the first generation of bars and washes out the halo prolateness. This chaos is triggered by the gravitational quadrupole interaction(s) in the system and supported by a number of other processes which are characteristic of baryons. ",
        "introduction": "The current paradigm of galaxy formation necessitates that galactic disks grow and evolve within dark matter (DM) halos. As such, the baryonic disks serve as a test bed for studying the DM properties, its dynamics and morphology, as well as mass, momentum and energy exchange within the disk-halo system. Interaction between the galactic disks and DM halos can in principle involve an exchange of these quantities. For example, observations of the low angular momentum H\\,I gas, deep in the halo of NGC~891, require cold extragalactic gas influx (e.g., Fraternali et al. 2007), maybe through filaments (Dekel \\& Birnboim 2006). On the other hand, the X-ray halos of starbursts and some normal galaxies can be explained by the supernovae-heated gas driven from the disk or by the shock-compressed gas in the halo (e.g., Strickland et al. 2004).  The situation is much more straightforward with the angular momentum ($J$) flow in the system. Galactic disks are rotationally supported, while the DM halos have low $J$, and, therefore, a low spin parameter $\\lambda$ (e.g., Barnes \\& Efstathiou 1987; Frenk, White \\& Efstathiou 1988). Most of the disk galaxies are barred, especially in the NIR (e.g., Knapen, Shlosman \\& Peletier 2000), with bar properties which remain steady up to the redshift of $\\sim 1$ at least (Jogee et al. 2004; Elmegreen, Elmegreen \\& Hirst 2004), and except for the very early Hubble types, they exhibit a spiral structure. This prevailing {\\it disk asymmetry} is extremely important for enhancing the $J$ transfer between different morphological components --- the alternative way can be achieved, e.g., by dynamical friction of baryons against the DM. The current understanding of galactic bar formation relies on the spontaneous breakup of the axial symmetry in the disk, the so-called classical bar instability (e.g., Hohl 1971), originally applied to isolated disks. Paradoxically, while Lynden-Bell \\& Kalnajs (1972) have shown that bars and spirals facilitate the $J$ transfer, DM halos have been considered as the main stabilizers against the bar instability for years (Ostriker \\& Peebles 1973), though they appear rather to be sinks of the disk momentum (e.g., Athanassoula \\& Misiriotis 2002).  Within the framework of the classical bar instability, a stellar bar can develop only if the unstable region loses its $J$. Gravitational torques serve as the mechanism that drives this process (Lynden-Bell \\& Pringle 1974). Their action can be described in terms of a non-local viscosity (Lin \\& Pringle 1987; Shlosman 1991) which shortens the characteristic timescale of the $J$ transfer dramatically, making this a dynamical rather then secular process. In principle, $J$ can flow across the corotation radius (CR) to the outer disk, to the DM halo, or to a flyby galaxy in case of a tidal galaxy interaction. The first option limits the $J$ flow because the mass of the outer disk is typically $\\sim 20\\%$ of the disk mass and its ability to absorb $J$ quickly saturates (Fig.~1, left). On the other hand, the halo is massive, its inner part has a comparable mass to the inner disk, and its overall $J$ is small --- the halo is not supported by rotation. Fig.~1 shows also that the outer halo, beyond the disk radius, is fully susceptible to the $J$ transfer, after the inner halo efficiency has decreased.  \\begin{figure}[!t] \\plotfiddle{fig01a.ps}{3.5cm}{-90}{45}{45}{-170}{140} \\plotfiddle{fig01b.ps}{3.5cm}{0}{31}{31}{-10}{-16.5} \\caption{ \\underline{\\it Left:} Evolution of angular momentum, $J$, in the disk (upper panel) and the halo (lower panel). Note, $J$ saturates in the disk outside the bar CR (from Martinez-Valpuesta, Shlosman \\& Heller 2006). \\underline{\\it Right:} The $m=2$ amplitude, $A_2$, of the stellar and DM ghost bars (upper panel) evolution. The fast dissolution of the ghost bar (lower) after its stellar counterpart is axisymmetrized. Note the very short time period for the lower frame (I. Berentzen \\& I. Shlosman, unpublished)} \\label{fig1} \\end{figure}  The relatively local density response in the DM halo to the growing bar in the disk is to generate a shadow `ghost' bar in the DM (Athanassoula 2006, 2007; Berentzen \\& Shlosman 2006). The ghost has the pattern speed of the stellar bar but its mass distribution differs and is more centrally concentrated. The particle orbits in the DM ghost, therefore, are in resonance with the stellar orbits. The principal difference between these two bars is that the DM bar is not strictly speaking a bar at all --- it is not self-gravitating {\\it per se} and represents a gravitational wake induced in the DM by its stellar counterpart. When the stellar bar is axisymmetrized, the DM bar dissolves in a fraction of a crossing time (Fig.~1, right). ",
        "conclusions": "We have summarized some aspects of disk-halo interactions in equilibrium and forming disk galaxies. Because the overall problem is of a high complexity, progress in this direction has been rather slow and dependent on numerical simulations. Nevertheless, certain trends have emerged. The baryonic disk and DM halo exchange mass, angular momentum and energy during the initial formation phase as well as during the subsequent quiescent evolution. A number of processes, like star formation as well as formation of the central SBH, require a much deeper understanding before they can be modeled beyond the phenomenological approach. On the other hand, nonlinear dynamics and the role of chaos in the evolution of disks and halos can be quantified already at present."
    },
    "0710/0710.4964_arXiv.txt": {
        "abstract": "We write the averaged Einstein equations in a form suitable for use with Newtonian gauge linear perturbation theory and track the size of the modifications to standard Robertson-Walker evolution on the largest scales as a function of redshift for both Einstein de-Sitter and $\\Lambda$CDM cosmologies. In both cases the effective energy density arising from linear perturbations is of the order of $10^{-5}$ times the matter density, as would be expected, with an effective equation of state $w_\\mathrm{eff}\\approx -1/19$. Employing a modified Halofit code to extend our results to quasilinear scales, we find that, while larger, the deviations from Robertson-Walker behaviour remain of the order of $10^{-5}$. ",
        "introduction": "\\label{intro} Observations of the cosmic microwave background (CMB) \\cite{WMAP,WMAP(2),WMAP3}, large-scale structure (LSS) (e.g. \\cite{2df,SDSS}) and supernovae type Ia data \\cite{RiessEtAl98,PerlmutterEtAl98,RiessEtAl06,WoodVaseyEtAl07} consistently indicate the presence in the universe of significant ``dark'' components. Data from the LSS and from nucleosynthesis indicates that the density of baryonic matter should not exceed about 5\\% of the critical density, while data from the CMB suggests that the universe should be flat, implying that around 95\\% of the matter-energy content of the universe is unobserved. In combination with numerical studies (e.g. \\cite{SpringelEtAl05}) and observations of the baryonic acoustic oscillations in the LSS \\cite{2df,SDSS} this has lead to the ``concordance model'' in which roughly 5\\% of the universe is in the form of baryonic matter, 20\\% of the universe is in the form of some ``dark matter'' that interacts only gravitationally, and about 75\\% is in the form of a ``dark energy'' fluid with a negative equation of state today $w\\approx-1$.  The most popular alternatives to a cosmological constant involve exotic fluids that violate the strong energy condition in the late universe, of which scalar-field models such as quintessence are the most popular (e.g. \\cite{Wetterich88,RatraPeebles88,CaldwellEtAl97} for pioneering works and \\cite{CopelandEtAl06,Linder07} for recent reviews.) There are also many other modifications of both the matter and gravitational sectors. Another approach is to consider the impact of local inhomogeneities on the luminosity distances in the local universe (see for example \\cite{EnqvistMattson07,Kasai07,Marra07} for some recent work on this topic). The modifications that the local inhomogeneities introduce can, it has been suggested, account for many of the features exhibited by dark energy.  However, the assumptions used to build up our standard model include an implicit averaging procedure, assuming that the averaged Einstein tensor is equivalent to the Einstein tensor built from an average metric -- and such is not the case \\cite{Ellis84,Futamase89}. The Einstein equations are non-linear and local; the correct approach is to average the local equations across some domain rather than to assume ab initio ``averaged'' equations. Studies into averaged cosmologies include \\cite{Kasai93,Futamase96,BuchertEhlers95,Boersma97,RussEtAl97} and the impact of inhomogeneities were applied to account for a dark energy in \\cite{Buchert99}. In the years since, this ``backreaction'' and the impact of inhomogeneities on observables has been relatively well studied in a variety of models and remains an active field \\cite{Buchert01,Wetterich01,BuchertCarfora02,Rasanen03,Rasanen04,KolbEtAl04,EllisBuchert05,AlnesEtAl05,Rasanen06,IshibashiWald06,BuchertEtAl06,LarenaBuchert06,LarenaEtAl06,BiswasEtAl06,KasaiEtAl06,VanderveldEtAl06,TanakaEtAl06,Brouzakis06,LiSchwarz07,Enqvist07,Wiltshire07,Mattsson07,Ishak07,VanderveldEtAl07,Wiltshire07_2,KhosraviEtAl07,Hossain07}; see also \\cite{BuchertReview07} for a recent status report. The attraction of this approach is that it can recover quantities that -- in principle -- could act as a dark energy, without the need for exotic matter components or modifications to general relativity. Moreover, one can intuitively state that as the universe grows increasingly non-linear, the deviations on averaging from a truly homogeneous and isotropic model should increase, potentially providing an appealing solution to the coincidence problem. The current universe is significantly inhomogeneous at scales of up to 100-300Mpc/$h$ \\cite{PercivalEtAl06}. While we might still expect to recover an FLRW-like metric when averaging across scales larger than 100Mpc/$h$ (see \\cite{Lu07,McClure07} for recent studies) there is no automatic guarantee that it will obey the familiar Friedmann and Raychaudhuri equations. It has been shown in many studies that structure -- both at the linear order and for highly inhomogeneous models -- can introduce modifications to the large scale effective Friedmann equations (though this is still somewhat debated; e.g. \\cite{AlnesEtAl05_2,KasaiEtAl06,TanakaEtAl06,EnqvistMattson07}). Even should the effect be relatively minor, with current and future experiments a divergence from Robertson-Walker behaviour on the order of $10^{-3}$ could be significant.  In this paper we calculate numerically as a function of redshift the size of the deviations from standard Robertson-Walker behaviour arising from linear and quasilinear perturbations. Such a calculation complements the recent studies by \\cite{VanderveldEtAl07}, in which the authors reconstruct the impact of backreaction effects from the observational data, and \\cite{KhosraviEtAl07} wherein the authors evaluate the size of the effective density of the backreaction as a function of redshift in a structured Robertson-Walker model. Much of the current literature concerns exact inhomogeneous rather than perturbative models, not least because the impact in a perturbative approach is not expected to be large. However, the magnitude and nature of the backreaction and other corrections associated with averaging that arise naturally within a standard Einstein de-Sitter or $\\Lambda$CDM cosmology remain open questions and this should be addressed. We work in conformal Newtonian gauge, in contrast to much of the literature which concerns itself with synchronous gauge, for two main reasons. Firstly Newtonian gauge, unlike synchronous gauge, is well-defined and should yield unambiguous conclusions; and, moreover, the use of Newtonian gauge enables us to employ the easily extensible cmbeasy \\cite{Doran03} code to track the size of the perturbations. We find that linear perturbations for a model Einstein de-Sitter universe introduce an effective energy density $\\rho_\\mathrm{eff}\\approx (4\\times 10^{-5})\\bkr_m$ with an equation of state $p_\\mathrm{eff}/\\rho_\\mathrm{eff}=w_\\mathrm{eff}\\approx -1/19$, while those for the $\\Lambda$CDM concordance model introduce $\\rho_\\mathrm{eff}\\approx (1.3\\times 10^{-5})\\bkr_m$ with a marginally lower equation of state. Employing a modified Halofit \\cite{SmithEtAl02} code to include quasilinear scales increases the effective energy densities to $\\rho_\\mathrm{eff}\\approx (5.6\\times 10^{-5})\\bkr_m$ and $\\rho_\\mathrm{eff}\\approx(1.6\\times 10^{-5})\\bkr_m$ for the EdS and $\\Lambda$CDM cases respectively, with unaltered equations of state. These are in good agreement with other evaluations of $\\rho_\\mathrm{eff}$ (e.g. \\cite{Wetterich01,Rasanen04,VanderveldEtAl07,KhosraviEtAl07}) which typically remain of the same order.  We begin with a brief overview of the 3+1 formalism of general relativity and in \\S\\ref{Averaging} briefly discuss averages in cosmology before applying a simple average in \\S\\ref{Averages}. In \\S\\ref{NewtGauge} we calculate the forms of the backreaction terms for conformal Newtonian gauge and in \\S\\ref{Backfast} we present the deviations from an FLRW model from linear perturbations employing a modified Boltzmann code for both Einstein de-Sitter and $\\Lambda$CDM universes, making it more general than the similar studies in \\cite{Rasanen04,KasaiEtAl06}, who did not consider $\\Lambda$CDM universes, and \\cite{TanakaEtAl06}, who worked in a rather narrow range of domain sizes. We follow by estimating the impact from quasilinear scales in \\S\\ref{Halofit}. ",
        "conclusions": "We have presented quantitive estimates of the corrections in Newtonian gauge to a standard perturbative FLRW model from an explicit averaging procedure, which can be separated into four distinct quantities: the kinematic backreaction, which remains strictly negligible across the scales of interest, the dynamic backreaction, the impact of spatial curvature and a correction to the FLRW energy density. In line with expectation, the impact from linear perturbations is insignificant and of the order of $10^{-5}$ for both $\\Lambda$CDM and an Einstein de-Sitter model. Specifically, the effective energy density arising at the background level from inhomogeneities is $\\rho_\\mathrm{eff}\\approx (4\\times 10^{-5})\\bkr_m$ and $\\rho_\\mathrm{eff}\\approx(1.3\\times 10^{-5})\\bkr_m$ for the EdS and $\\Lambda$CDM models respectively. Moreover, we have presented estimates arising from halo model corrections on quasilinear scales which are not much larger than those from the linear perturbations, the largest contribution arising from the quasilinear modes in an EdS universe remaining below $10^{-4}$; we find $\\rho_\\mathrm{eff}\\approx(5.6\\times 10^{-5})\\bkr_m$ for the EdS model and $\\rho_\\mathrm{eff}\\approx(1.6\\times 10^{-5})\\bkr_m$ for $\\Lambda$CDM. This is in broad agreement with recent calculations from Vanderveld \\emph{et. al.} estimating the total possible backreaction from observations of Type Ia supernovae and agrees with a growing consensus that the impact of backreaction across all scales remains at the level of $10^{-5}-10^{-4}$. Additionally, the total effective equation of state arising from the different modifications does not act as a dark energy and instead as a dark matter with a slowly-varying equation of state $w_\\mathrm{eff}\\approx -1/19$ for $z\\in(0,10)$. The impact is to decelerate the universe and, contrary to other studies, the corrections impact positively on the Friedmann equation implying a positive effective pressure. (The discrepancy arises due to our definitions of the effective energy density and pressure of the corrections taking into account the differences in definition of energy density in an FLRW frame and with respect to our foliation.)  This does not, however, necessarily imply that there are no observational consequences arising from inhomogeneities, particularly on small scales, as our analysis is limited to the impacts in very large volumes and breaks down when mode-coupling and virialisation become significant. In particular, one can readily imagine situations in which the local Hubble rate is significantly larger than the global average (the so-called ``Hubble bubble''). The impact of local inhomogeneities directly on luminosity distances has also long been well studied and remains an active area of interest. Moreover, the issue of averaging in cosmology and general relativity in particular remains very much an active field. The average we have employed in this study is certainly not perfect; even were we to extend our approach to second-order perturbations it cannot cope with the vector and tensor perturbations that necessarily become significant.% Improving the average and applying it to nonlinear situations is an immediate aim.  The prospect for future study remains large. Our model here is necessarily simple and there are many avenues both analytical and numerical that can be pursued. Much interest is currently focusing on the exact swiss-cheese models and particularly on the impact on direct observables from these models. The prospect of directly calculating both the direct backreaction in different domain sizes and the impact on null geodesics (and hence luminosity distances) in different matter configurations from numerical simulations is also an exciting one. More immediately, one can imagine employing second-order perturbation theory in which the ambiguity between averages taken across the assumed zeroth-order manifold and those across the linearly-perturbed manifold may become significant; depending on the surface one defines a vanishing $\\langle\\phi_{(2)}\\rangle$ across, one would expect a non-vanishing contribution on the other. In particular, if one defined second-order perturbations to vanish on average on the background then there would be a residual contribution to the linear perturbations even before any other backreaction is taken into account. This is an interesting issue that deserves more study.  More ambitiously, the fully non-linear approaches of Vernizzi and Langlois \\cite{LangloisVernizzi05_1,LangloisVernizzi05_2,LangloisVernizzi06_1} and Enqvist \\emph{et. al.} \\cite{EnqvistEtAl06} might provide an interesting tool for studying the impacts of averaged nonlinear perturbations. The direct backreaction from more complicated systems than our linearly-perturbed FLRW universe might also be considerable; an immediate example, first considered in \\cite{Wetterich01}, would be the backreaction from inhomogeneities in a cosmological scalar field which were shown to potentially have an impact of order unity. Such systems might also be considered in a fully non-linear manner \\cite{LangloisVernizzi06_2}.  In summary, one may comment that backreaction and the impact of inhomogeneities are direct physical effects and should certainly be taken into account in any full treatment of the observables. However, the direct impact when evaluated for the concordance cosmology on very large scales is too small to provide an alternative to the dark energy."
    },
    "0710/0710.1770_arXiv.txt": {
        "abstract": "While masers in the 1720~MHz transition of OH are detected toward many supernova remnants (SNRs), no other OH transition is seen as a maser in SNRs. We present a search for masers at 6049~MHz, which has recently been predicted to produce masers by pure collisional excitation at conditions similar to that required for 1720~MHz masing. The Effelsberg 100 m telescope was used to observe the excited-state 6016, 6030, 6035, and 6049~MHz lines of OH toward selected SNRs, most of which have previously-detected bright 1720~MHz masers. No excited-state masers are found toward SNRs, consistent with previous observations of the 6049~MHz and other excited-state transitions.  We do not see clear evidence of absorption toward SNR target positions, although we do see evidence of absorption in the molecular cloud at $+50$~km\\,s$^{-1}$ near Sgr~A East.  Weak absorption is detected at 6016~MHz toward W3(OH), while stronger, narrower emission is seen at 6049~MHz, suggesting that the 6049~MHz emission is a low-gain maser. We conclude that conditions in SNRs are not conducive to excited-state maser emission, especially in excited-state satellite lines. ",
        "introduction": "About two dozen supernova remnants (SNRs) are known to host 1720~MHz OH masers \\citep[][and references therein]{green06}.  The excitation mechanism for 1720~MHz masers in SNRs is commonly believed to be collisional excitation from a C-type shock \\citep{lockett99,wardle99}. Other ground-state OH transitions are sometimes seen in absorption \\citep[e.g.,][]{goss68}, which can be helpful for modelling the physical conditions in the 1720~MHz masing region \\citep{hewitt06}. However, to date no other OH transition, nor a transition of any other molecule, has been detected as a maser associated with an SNR.  The detection of a second maser transition in SNRs would place strong constraints on physical conditions (density, temperature, OH fraction, ortho-to-para H$_2$ ratio, etc.) in the masing region, especially if spatial coincidence with a 1720~MHz maser was observed.  Excited-state transitions, which occur at higher frequencies, would be especially useful if detected toward the Galactic center, where angular scattering is large \\citep[see \\S~1 of][]{pihlstrom07}.  Another motivation for finding a second masing transition in SNRs is to confirm the Zeeman interpretation of splitting seen between left and right circular (LCP, RCP) components at 1720~MHz.  Differences in LCP and RCP velocities are used to compute magnetic field strengths in SNRs \\citep[e.g.,][]{brogan00}, important for quantifying magnetic pressures.  However, \\citet{elitzur96,elitzur98} shows that circular polarization can be generated by non-Zeeman mechanisms when the maser is saturated and splitting is small compared to a line width, as is the case in SNR 1720~MHz masers.  While there are theoretical reasons to believe the Zeeman interpretation of SNR 1720~MHz maser splitting \\citep[see \\S~4.1 of][]{brogan00}, direct confirmation of the magnetic field strengths from a second OH transition would lend greater confidence in the derived magnetic field strengths.  Theoretical modelling predicts that collisions alone can excite several transitions of OH.  In the low-density ($n \\sim 10^5$~cm$^{-3}$) regime, only the 1720~MHz transition is inverted, but at higher number and/or column densities the 6049~MHz transition also becomes inverted, with some overlap in parameter space allowing simultaneous inversion of the 1720 and 6049~MHz lines \\citep{pavlakis00,wardle07,pihlstrom07}.  At still higher densities ($n \\sim 5~\\times~10^6$~cm$^{-3}$), the 4765 and 1612~MHz transitions become inverted \\citep{pavlakis96,pihlstrom07}.  The 6035~MHz line may also become inverted at high temperatures \\citep[$T \\sim 200$~K;][]{pavlakis00}.  Regardless of the details of the pump mechanism, weak inversion (emission) is predicted in the 6049~MHz and weak anti-inversion (enhanced absorption) in the 6016~MHz line from the structure of the OH level diagram alone when infrared trapping becomes important \\citep{litvak69,elitzur77}.  Few observations have sought excited-state OH masers in SNRs. \\citet{mcdonnell07} looked for 6030, 6035, and 6049~MHz OH maser emission toward southern SNRs but did not detect any 6049~MHz emission.  Although they did find three new 6030 and 6035~MHz maser sources, \\citet{mcdonnell07} believe that these masers are associated with \\ion{H}{2} regions and therefore represent emission from star-forming regions (SFRs), not SNRs.  In states of higher excitation, \\citet{pihlstrom07} searched for emission in the 4.7~GHz $\\Lambda$-doubling triplet and the 7.8 and 8.2~GHz quadruplets toward four SNR complexes (as well as the 23.8~GHz quadruplet toward Sgr~A~East) but did not detect any emission.  \\citet{pihlstrom07} also cross-correlated positions of single-peaked and irregular-spectrum 1612~MHz masers from blind surveys in the literature with positions of known SNRs but found no probable associations.  For small velocity differences along the amplification path, the 6049~MHz transition is theoretically the most promising transition in which to find new OH SNR masers, given that the collisional excitation conditions require only slightly higher densities than those that produce 1720~MHz masers.  We report on targeted observations of the four 6.0~GHz lines mostly toward SNRs with bright 1720~MHz masers. ",
        "conclusions": "Though our original interest was on the 6.0~GHz OH transitions in SNRs, we start off the discussion with our findings on our reference source, the SFR region W3(OH).  \\subsection{Satellite Lines in W3(OH)} \\label{satellite}  The 6049~MHz ($F = 3 \\rightarrow 2$) emission in W3(OH) is stronger than the 6016~MHz ($F = 2\\rightarrow 3$) absorption, as noted in the calculations by \\citet{guilloteau84}, who did not detect the 6016~MHz absorption at all.  This effect, in which 6016~MHz absorption is enhanced relative to the expected LTE value (from 6030 and 6035~MHz absorption) and the 6049~MHz line is seen in weak absorption or even emission, was seen in other SFRs by \\citet{whiteoak76} and \\citet{gardner83}.  This behavior was predicted by \\citet{elitzur77}, who showed that anti-inversion at 6016~MHz and inversion at 6049~MHz should occur whenever the 120~$\\mu$m radiation linking the first-excited $J = 5/2$ state with the ground $J = 3/2$ state in the $^2\\Pi_{3/2}$ ladder is optically thick, regardless of the details of the pump mechanism.  Interestingly, a similar phenomenon is seen in the next highest state in the $^2\\Pi_{3/2}$ ladder.  Absorption in the 13442~MHz ($F = 3 \\rightarrow 4$) transition is always stronger than at 13433~MHz ($F = 4 \\rightarrow 3$) and always enhanced from LTE values compared to the main lines \\citep{fish05}.  The \\citet{elitzur77} result is generalizable: whenever intra-ladder decays dominate and are optically thick, the $\\Delta F = +1$ line will be seen in enhanced absorption and the $\\Delta F = -1$ line will be seen in emission or reduced absorption.  For the 13.4~GHz transitions, infrared overlaps in the 84~$\\mu$m transitions linking the $J = 7/2$ and $5/2$ states may also be important, as proposed by \\citet{matthews86}.  We find that a Gaussian fit to the 6049~MHz emission in W3(OH) is also \\emph{narrower} than the fit to the 6016~MHz absorption, even though the full width at zero power of the two lines is the same.  This is consistent with weak inversion and suggests that the 6049~MHz emission should be understood as a low-gain maser, as discussed by \\citet{baudry97b}.  Further sensitive observations at higher spectral resolution will be required to confirm the maser nature of the 6049~MHz emission.  High angular resolution interferometric observations would be particularly useful to determine the size and brightness temperature of the 6049~MHz emission region, as well as to isolate the region of emission and determine whether the inversion exists over the entire source or merely in cluster C \\citep[in the nomenclature of][]{fish07}, which must be a region of very high excitation based on the presence of highly-excited OH \\citep{baudry93,baudry97}.  In general, satellite-line masers are seen almost exclusively in the lowest states of the $^2\\Pi_{3/2}$ and $^2\\Pi_{1/2}$ ladders.  In the $^2\\Pi_{3/2}, J = 3/2$ rotational state, masers are seen in the 1612~MHz transition in evolved stars and some SFRs, and 1720~MHz masers are seen in SNRs and some SFRs.  Masers in the $^2\\Pi_{1/2}, J = 1/2$ state are seen only in SFRs, with 4765~MHz ($F = 1 \\rightarrow 0$) masers much more common than 4660~MHz ($F = 0 \\rightarrow 1$) masers \\citep[e.g.,][]{gardner83,cohen95}, in agreement with \\citet{elitzur77}.  Even though the brightest main-line 6030 and 6035~MHz masers in SFRs can have flux densities of several hundred~Jy \\citep[e.g.,][]{fish07}, the only other satellite-line maser found to date is the low-gain 6049~MHz maser in W3(OH).  It thus appears to be a general result that it is extremely difficult to invert satellite lines except in the lowest rotational state in each ladder.   \\subsection{Sgr~A East} \\label{ae-discussion}  We see 6030 and 6035~MHz absorption near $+50$~km\\,s$^{-1}$ in pointings on the eastern side of Sgr~A East.  The absorption at 6035~MHz is deeper than at 6030~MHz, as would be expected for thermal absorption.  This absorption is likely neither part of the interaction of the SNR with the molecular cloud nor located at the interaction region \\citep[with most 1720~MHz masers in the slightly higher velocity range $V_\\mathrm{LSR} = 53$--$68$~km\\,s$^{-1}$;][]{pihlstrom06}, although interferometric confirmation, as with the Expanded Very Large Array (EVLA) will be required to establish the location of the absorption.  The molecular material is denser to the east and partially obscured by the circumnuclear disk to the west \\citep{mcgary01,herrnstein05}, consistent with our absorption detections as well as the ground-state absorption detected by \\citet{karlsson03}.  The detection of excited-state absorption toward Sgr~A East but not toward any other SNR may find its explanation in a chance alignment, the special conditions in the Galactic center region, or a combination thereof.  The difference with the other SNRs may be due to the geometry of a dense cloud in the line of sight with a relatively intense radio continuum background in the case of Sgr~A East, resulting in a larger column density and absorption of OH.  Or, the difference may be due to a difference in heating of the cloud: the +50~km\\,s$^{-1}$ cloud toward Sgr~A East may be heated more, not by the impact of the SNR into the cloud \\citep{herrnstein05} but, e.g., due to a strong local radiation field absorbed by high-metallicity material, dissipation of kinetic energy or collisions of local clumps in a steep gravitational potential, etc., providing a larger column density of OH in excited states.  As this absorption is likely not due to the SNR, further studies are required to investigate its origin.  Interferometric observations would be very useful in order to understand the distribution of excited-state OH in the molecular cloud near Sgr~A East.  Higher angular resolution than that provided by Effelsberg will be required to understand the origin of the 6.0~GHz absorption in Sgr~A East.  This absorption is an obvious target for reobservation with the EVLA when a sufficient number of antennas are equipped with the new C-band (4--8~GHz) receivers.   \\subsection{Other Sources}  No emission was seen in any 6.0~GHz transition toward SNRs.  To date, a total of 23 $\\Lambda$-doublet transitions have been observed toward SNRs: the quadruplets at 1.6, 6.0, 7.7, 8.1, and 23.8~GHz and the triplet at 4.7~GHz.  Only the 1720~MHz transition produces a detectable maser.  The 6049~MHz transition is the next one after 1720~MHz predicted to go into inversion as densities are increased \\citep{pihlstrom07,wardle07}, yet searches for 6049~MHz masers have not uncovered any detections \\citep[this work, as well as][]{mcdonnell07}.  Searches for SNR OH masers at 4765 and 1612~MHz, the next two transitions expected to produce masers via collisional excitation, have also produced only non-detections \\citep{pihlstrom07}.  It is probable that no OH masers, other than in the 1720~MHz transition, will be detected toward SNRs with current instrumentation.  Absorption at 6016~MHz is not seen toward any SNR, despite predictions that it should be anti-inverted.  We do not see absorption toward SNRs with the exception of 6030 and 6035~MHz absorption toward Sgr~A East. Much more sensitive observations will be required to test whether excited-state $\\Delta F = +1$ transitions are in enhanced absorption in SNRs.  No emission was seen toward the evolved stars AU Gem and NML Cyg. Excited-state OH maser emission has previously been reported in the 4750~MHz transition toward AU Gem \\citep{claussen81} and the 6035 (and possibly 6030) MHz transition toward NML Cyg \\citep{zuckerman72}, although subsequent observations have failed to confirm these detections \\citep[e.g.,][]{sjouwerman07}.  It is likely that excited-state OH maser emission from evolved stars is temporary and extremely rare \\citep{sjouwerman07}.  Our continued non-detection of excited-state emission toward these stars supports this hypothesis."
    },
    "0710/0710.1293_arXiv.txt": {
        "abstract": "{We investigate, by means of numerical simulations, the possibility of forming counter-rotating old stellar components by major mergers between an elliptical and a spiral galaxy. We show that counter-rotation can appear both in dissipative and dissipationless retrograde mergers, and it is mostly associated to the presence of  a disk component, which preserves part of its initial spin. In turn, the external regions of the two interacting galaxies acquire part of the orbital angular momentum, due to the action of tidal forces exerted on each galaxy by the companion. } ",
        "introduction": "Simulations show that stellar counter-rotation in galaxies could emerge thanks to two different processes: dissipative and dissipationless mergers.\\\\ Confirming the suggestion  of \\citet{korm84}, \\citet{balquin90} showed that unequal mass mergers of elliptical galaxies % can produce counter-rotation if the orbit of the encounter is retrograde with respect to the spin of the primary. They also pointed out that the rotation seen in the counter-rotating component is a tracer of the orbital angular momentum and that  both primary and secondary stars counter-rotate at the core.\\\\ \\citet{hb91} showed that counter-rotating central gas disks can form as a result of retrograde mergers between  two gas rich spiral galaxies, discussing the possibility that star formation in such disks could produce components with decoupled kinematics, as in the core of some elliptical galaxies.\\\\ Some years later, \\citet{balgon98} showed that kinematically peculiar cores may be generated also in retrograde stellar spiral-spiral mergers: in this picture, the central bulges transport orbital angular momentum inward to the center of the remnant, while the outer parts keep the spin signature of the precursor disks. Also \\citet{bendob00} put in evidence the possibility to form counter-rotation at large radii, simulating mergers between equal mass disk galaxies. Finally, \\citet{jes07}, studying the 2D kinematics of a sample of simulated disk merger remnants, showed that counter rotating cores, made of old stellar populations, are almost exclusively formed in equal mass mergers where a dissipative component is included. Evidently, both  mechanisms (the dissipative and dissipationless one) can occur in real systems, producing a variety of kinematically decoupled components, of different ages and physical extensions \\citep{mc06}.  In this paper, we want to present a new scenario, according to which stellar counter-rotation can form both in dissipative and dissipationless retrograde\\footnote{i.e. the spin of the spiral galaxy is initially anti-parallel to the orbital angular momentum.} major mergers of elliptical-spirals galaxies\\footnote{Note that all the previous  numerical works refer to mergers between two ellipticals or two disk galaxies. So far, the possibility of forming counter-rotating cores by mergers of galaxies of different morphologies has never been exploited.}. In this scenario, stars in the external regions of the galaxies involved in the encounter acquire orbital angular momentum, at first pericenter passage and, mostly, in the last phases of the merging process, while those in the most inner regions can maintain some of the initial spin (anti-parallel to the orbital one), producing decoupled counter-rotating central components. In the case of dissipative mergers, the central decoupled core could be composed of two different populations: the old stellar population, which has preserved its initial spin, and a new stellar population, born \\emph{in situ} from the kinematically decoupled gas component. The redistribution of the angular momentum between the different galactic components (gas, stars and dark matter) is  discussed. ",
        "conclusions": "We have presented a new scenario to form counter-rotating central components in early-type galaxies, by dissipative and dissipationless mergers of elliptical-spiral systems in retrograde orbits. In the case of dissipative mergers, the central decoupled core could be composed of two distinct populations: the old stellar population, which has preserved part of its initial spin, and a new stellar population, born \\emph{in situ} from the kinematically decoupled gas component. Counter-rotating cores are also found in the remnants of non-coplanar mergers, at least for inclinations $i \\le 20^0$ (being $i$ the angle between the spiral disk and the orbital plane). We plan to realise a wider and  systematic study of the role played by orbital parameters and morphology of the interacting systems in the near future, varying the morphological parameters of the interacting galaxies (bulge-to-disk ratio of the spiral galaxy as well as different initial models for the elliptical), and the orbital initial conditions. %"
    },
    "0710/0710.5897_arXiv.txt": {
        "abstract": "\\vspace{0.5cm} Modern cosmology has now emerged  as a testing ground for theories beyond the standard model of particle physics. In this paper, we consider quantum fluctuations of the inflaton scalar field on certain noncommutative spacetimes and look for noncommutative corrections in the cosmic microwave background (CMB) radiation. Inhomogeneities in the distribution of large scale structure and anisotropies in the CMB radiation can carry traces of noncommutativity of the early universe. We show that its power spectrum becomes direction-dependent when spacetime is noncommutative. (The effects due to noncommutativity can be observed experimentally in the distribution of large scale structure of matter as well.) Furthermore, we have shown that the probability distribution determining the temperature fluctuations is not Gaussian for  noncommutative spacetimes. ",
        "introduction": "\\label{sec:intro} The CMB radiation shows  how the universe was like when it was only $400, 000$ years old. If photons and baryons were in equilibrium before they decoupled from each other, then the CMB radiation we observe today should have a black body spectrum indicating a smooth early universe. But in 1992, the Cosmic Background Explorer (COBE) satellite detected anisotropies in the CMB radiation, which led to the conclusion that the early universe was not smooth: There were small perturbations in the photon-baryon fluid.  The theory of inflation  was introduced \\cite{guth, Linde, Albrecht} to resolve the fine tuning problems associated with the standard Big Bang cosmology. An important property of inflation is that it can generate irregularities in the universe, which may lead to the formation of structure. Inflation is assumed to be driven by a classical scalar field that accelerates the observed universe towards a perfect homogeneous state. But we live in a quantum world where perfect homogeneity is never attained. The classical scalar field has quantum fluctuations around it and these fluctuations act as seeds for the primordial perturbations over the smooth universe. Thus according to these ideas, the early universe had inhomogeneities and we observe them today in the distribution of large scale structure and anisotropies in the CMB radiation.  Physics at Planck  scale could be radically different. It is the regime of string theory and quantum gravity. Inflation stretches a region of Planck size into cosmological scales. So, at the end of inflation, physics at Planck region should leave its signature on the cosmological scales too.  There are indications  both from quantum gravity and string theory that spacetime is noncommutative with a length scale of the order of Planck length. In this paper we explore the consequences of such noncommutativity for CMB radiation in the light of  recent developments in the field of noncommutative quantum field theories relating to deformed Poincar\\'e symmetry.  The early universe and CMB in the noncommutative framework have been addressed in many places \\cite{Greene, Lizzi, Brandenberger1, Huang, Brandenberger2, BalBB, Fatollahi2, Fatollahi1}. In \\cite{Greene}, the noncommutative parameter $\\theta_{\\mu \\nu} = -\\theta_{\\nu \\mu} =\\textrm{constants}$ with $\\theta_{0i} =0$, ($\\mu, \\nu = 0, 1, 2, 3$, with $0$ denoting time direction), characterizing the Moyal plane is scale dependent, while \\cite{Brandenberger1, Brandenberger2, Huang} have considered noncommutativity based on stringy space-time uncertainty relations. Our approach differs from these authors since our quantum fields obey twisted statistics, as implied by the deformed Poincar\\'e symmetry in quantum theories.  We organize the paper as follows: In section II, we discuss how noncommutativity breaks the usual Lorentz invariance and indicate how this breaking can be interpreted as invariance under a  deformed Poincar\\'e symmetry. In section III, we write down an expression for a scalar quantum field in the noncommutative framework and show how its two-point function is modified. We review the theory of cosmological perturbations and (direction-independent) power spectrum for $\\theta_{\\mu \\nu}=0$ in section IV. In section V, we derive the power spectrum for the noncommutative Groenewold-Moyal plane ${\\cal A}_{\\theta}$ and show that it is direction-dependent and breaks statistical isotropy. In section VI, we compute the angular correlations using this power spectrum and show that there are nontrivial ${\\cal O}( \\theta^{2})$ corrections to the CMB temperature fluctuations.  Next, in section VII, we discuss the modifications of the $n$-point functions for any $n$ brought about by a non-zero $\\theta^{\\mu \\nu}$ and show in particular that the underlying probability distribution is not Gaussian. The paper concludes with section VIII. ",
        "conclusions": "In this paper, we have shown that the introduction of spacetime noncommutativity gives rise to  nontrivial contributions to the CMB temperature fluctuations. The two-point correlation function in momentum space, called the power spectrum, becomes direction-dependent. Thus spacetime noncommutativity breaks the rotational invariance of the CMB spectrum. That is, CMB radiation becomes statistically anisotropic. This can be measured experimentally to set bounds on the noncommutative parameter. Currently, we \\cite{numerical} are making numerical fits to the available CMB data to put bounds on $\\theta$.  We have also shown that the probability distribution governing correlations of fields on the Groenewold-Moyal algebra ${\\cal A}_{\\theta}$ are non-Gaussian. This affects the correlation functions of temperature fluctuations. By measuring the amount of non-Gaussianity from the four-point correlation function data for temperature fluctuations, we can thus set further limits on $\\theta$.  We have also discussed the signals of non-causality of non-commutative field theories in the temperature fluctuations of the CMB spectrum. It will be very interesting to test the data for such signals."
    },
    "0710/0710.2872.txt": {
        "abstract": "This contribution intends to give a pedagogical introduction to the topic of dark energy (the mysterious agent supposed to drive the observed late time acceleration of the Universe) and to various observational tests which require only assumptions on the geometry of the Universe. Those tests are the supernovae luminosity, the CMB shift, the direct Hubble data, and the baryon acoustic oscillations test. An historical overview of Cosmology is followed by some generalities on FRW spacetimes (the best large-scale description of the Universe), and then the test themselves are discussed. A convenient section on statistical inference is included as well. ",
        "introduction": "Cosmology is a branch of physics which is experiencing a tremendously fast development lately triggered by the arrival of many new observational data of ever more exquisite precision. These findings  have been crucial for the improvement in  our understanding of the Universe, but the vast amount of knowledge on our Cosmos which we have nowadays would never have been possible without the concerted effort by experimentalists and theoreticians. One of the results of this fantastic intellectual pursue is the puzzling discovery that there is in the Universe a manifestation of the repulsive side of gravity. This is the topic to which this lectures are devoted, and more specifically I wish to address some methods which the community believes are useful for building a deeper understanding of why, how and when our universe began to accelerate. Put in  more modest words,  these lectures will dissert about how one can take advantage of various observational datasets to describe some basic features of  the geometry of the Universe with the hope they will reveal us something on the nature of the agent causing the  observed accelerated expansion.  As this is quite an advanced topic in Physics (Astronomy), this contribution will build tougher as we proceed, so I hope the readers will enjoy to start off from a little historic stroll in the science of surveying the skies. For wider historical overviews than the one presented here, the two main sources I recommend to you, among the so many available, are \\cite{conv} and \\cite{wiki}.   Historical records tells us that Astronomy was born basically because of the need by agrarian societies to predict seasons and other yearly events and by  the esoteric need to place humanity in the Universe. This discipline developed in many ancient cultures (Egyptian, Chinese, Babylonian, Mayan), but it was only Greek people who cared understanding their observations and spreading their knowledge unlike in other cultures. Their influence was crucial as the modern scientific attitude of relying in empiricism (Aristotlean school) and translating physical phenomena into the mathematical language (Pithagorean school) built on their way of approaching science. Unfortunately, Aristotle was far so influential that his erroneous geocentric view of the Universe was not questioned aloud until the XVIth century.  At the beginning of that century Copernicus started to spread quietly his heliocentric cosmological model, and even though he did not make much propaganda about his ideas, they deeply influenced other figures, such as Galileo, who introduced telescopes into astronomy and found solid evidence against the geocentric model. Another very important contribution to the subject was made by Kepler, who gave accurate characterization of the motions of the planets around the Sun. These findings were later on synthetically explained by Newton in his theory of gravitation, which built on Galileo's developments on dynamics (the law of inertia basically).  The XVIIIth and XIXth century brought advances in the understanding of the Universe as a whole made of many parts which lie not necessarily in the Solar system. Stars began to be regarded as far-away objects in motion, and other objects like nebulae were discovered, so the idea of the existence of complicated structures outside the Solar system gained solidity. A discovery very much related to the main topic of this lectures was the realization that the combination of the apparent brightness of a star and its distance get combined to give its intrinsic brightness; basically the amount of energy reaching us in the form of light from a distant source decreases is inversely proportional to the square of the distance between the source an us (in an expanding the universe definition of distance is not the same as in a static one but this broad way of speaking applies all the same) This finding let observers realize that stars were objects very much like the Sun, or if you prefer they realized the Sun was nothing but yet another star.  The most important next breakthrough was  Einstein's theory of special relativity, which generalized Galileo's relativity to introduce light. On the conceptual realm, this was a very revolutionary theory at is warps the notions of space and time, and so it set the foundations for the best description of gravity so far: Einstein's theory of General Relativity. This second theory was the fruition of Einstein endeavors to unify the interactions known to him. In this theoretical framework matter/energy modifies the geometry, and in turn geometry tells matter/energy how to move/propagate (paraphrasing Wheeler's renowned quotation). Among other predictions this theory made a couple which are cornerstones of modern astronomy: gravitational redshift (light gets redder as it moves away from massive objects), and gravitational lensing (light gets bent as it  passes close to massive objects).   The next important advance come from the side of observations. In 1929 Hubble, and after having collected data carefully for almost a decade,  presented the surprising  conclusions that galaxies on average move away from us. This effect is encoded in the scale-invariant relation known as Hubble's law: $ v=H d$, where $v$ is the galaxy's velocity and $d$ its distance from us. The positiveness of the  quantity $H$ as measured by Hubble is precisely what told him the Universe is expanding, this being a discovery which gets accommodated nicely in Einstein's theory of general relativity. Interestingly, on learning about this finding, Einstein discarded his idea of the necessity of  some exotic fluid with negative pressure to counteract the attractive effect of usual matter (which would make the universe contract). It cannot look but funny from today's perspective that Einstein's idea has come to life again, at the end of the day the exotic fluid he imagined is one of the possible flavors of what  the  community calls dark energy \\cite{turner} these days.  At the risk of not giving everyone the credit they merit, I will just say that recognition for the concept of a expanding Universe is due to both theoreticians (de Sitter, Friedmann, Lema\\^itre, Gamow,...) and experimentalists (Slipher, Hubble,..). However, the  trampolin for this idea to jump into orthodoxy was put by Penzias and Wilson \\cite{penzias} who first detected the cosmic microwave background, and by Dicke, Peebles, Roll and Wilkinson \\cite{peebles} who were responsible for the not the less important interpretation of those observations. The existence of this radiation and its characteristic black body spectrum are a prediction of the Big-Bang theory, it is fair to say that if one combines it with other sources of evidence, it is almost impossible to refute it.   The CMB  in an invaluable source of cosmological information (visit \\cite{hu} for a higly recommended site on the subject). It has a temperature of of around $2.73\\,K$, with tiny temperature differences of about $10^{-5}$ between different patches of the sky. These anisotropies inform us that the photons forming the CMB where subject to an underlying   gravitational potential which had fluctuations, and this is just and indication of density irregularities that seeded the structure one observes today.  On the other hand, the image of the features of the CMB in different angular scales strongly favors a universe with no spatial curvature (flat or Euclidean on constant time slices) that is, it supports the theory of inflation. In addition, the CMB  has a saying about the fraction of the different fluids filling the Universe.  Fortunately, the mine of cosmic surprises was far from being exhausted, and it stored a diamond of many carats in the form of a discovery which has changed greatly the mainstream view of the Universe. Up to the late 1990, no repulsive manifestation of gravity had been spotted in Nature. In 1998 astronomical measurements provided evidence that gravity can not only push, but pull as well, and the repulsive side of gravity got unveiled. Very refined observations of the brightness of distant supernovae \\cite{Riess:1998cb,Perlmutter:1998np} seemed to hint the presence of a negative pressure component in the Universe which would make it accelerate, and then a new revolution started.  One may wonder at this stage what is the relation of supernovae luminosity and cosmic speed up. Objects in the Universe with a well calibrated intrinsic luminosity (supernovae, for instance) can be used to determine distances on cosmological scales. Supernovae are very bright objects, and so they result particularly attractive for this purpose (they can be $10^9$ times more luminous than the Sun so there is hope they will be visible from up to perhaps $1000$ Mpc \\cite{physto}). As we anticipated in the last paragraph, in  1998 two independent teams reported evidence that some distant supernovae were fainted than expected. This involved tracing the expansion history of the Universe by combining measurements of the recession velocity, apparent brightness and distance estimations. The most compelling explanation was (and keeps being as far I am concerned) that their light had traveled greater distances than assumed. The orthodox view up to then was that the expansion pace of the Universe was barely constant, but supernovae seemed to contradict this. This unexpected and exciting discovery obliged researchers to broaden their mind and accept the Universe is undergoing accelerated expansion.  I should have been able to have convinced the readers by now that these are exciting times to be working on Cosmology, as cosmic speed up is such and intriguing phenomenon with major open questions such as whether dark energy evolves with time, how much of it  is there, and if is rather not a manifestation of extradimensional physics.  The answer to these questions requires cannot be dissociated from the response to a perhaps more fundamental question: what is the Universe made of?  The combination of various astronomical observations  tell us our universe is basically made of thee major components (see for instance \\cite{wmap}). The most abundant one is dark energy \\cite{turner}, so this makes it even more interesting to find out whatever we can about it. At the other end the by far least abundant component is baryonic matter, and in the middle (as abundance is concerned) we have dark matter, in a proportion comparable to that of dark energy. There are various sources of astrophysical giving evidence in favor of it.  Hints  of it is existence are provided by the motion of stars, galaxies and clusters, but it is known it also played a crucial role in the amplification of the primordial density fluctuations which seeded the large scales structure we observe today and dark matter imprints can be found in the CMB as well. Dark matter represents quite a challenge as its nature remains a mystery; nevertheless if it were baryonic we know from big bang nucleosynthesis helium-4 would get converted to deuterium much easily and CMB calculations indicate indicate in addition anisotropies would be much larger so the odds are most of its is not baryonic.  Up to here we have made a very broad introduction to our topic with a little bit of history and a little bit of physics, but we must not forget Maths are key to Cosmology \\cite{Tegmark}. We only know how to study the Universe  using numbers and equations, but of course progress in this direction is done with as many reasonable simplifications as possible (if they do not compromise rigor, of course). Einstein equations relate geometry and matter/energy content of the Universe. Cosmologists are concerned by this relation on a large scales picture so as to understand  the expansion of the Universe. Those equations are non-linear, so studying them is painstaking unless one exploits the observational evidence the regularity of the Universe. Non-linearity of those equations makes their study a really hard task so simplifications are a must. The two basic ones are that galaxies are homogeneously distributed on galaxies  larger than $50$ {\\rm Mpc} \\cite{Yadav}, and that the Universe is isotropic around us on angular scales larger than about $10$ degrees \\cite{Souradeep}. But this is not enough, those two simplifications, which come from observational evidence must be completed with  the assumption we occupy no special place in the Universe.  This puts on the track of what we could call the parameterized Universe, cosmologists  work with a greatly a simplified geometric description of the Universe which emerges from the latter assumptions. The models for the sources are of reduced complexity too, the most common being perfect fluids and fields with known dynamics. This we have on the side of theory, but on the side of observations we have to make our own life easier too. When doing observations oriented Cosmology, one uses  as a variable the redshift $z$ of the electromagnetic radiation received, as it encodes information of how much the Universe has expanded between emission and reception. Observational tables of geometrical quantities can be given, and then one can test different theoretical values of the same quantities corresponding to models of interest. The theoretical predictions will depend on the sources assumed, and  ultimately it will be possible to estimate the suitability of a given dark energy model to observations, but which are the available  probes of dark energy?  Basically dark energy can be scrutinized observationally from two main perspectives  \\cite{trotta}: one possility is doing it through its effect of the growth of structures, another one is through its impact on geometrical quantities. We will concentrate on geometrical constraints/tests in these lectures as in a way they are those which can perhaps be applied with less difficulty, although their simplicity does not mean they are the least interesting, for instance, the supernovae test is the only test giving a direct indication of the need of a repulsive component in the Universe, whereas the baryon acoustic oscillations test is thought to have much information in store. Discussion on these two tests will be given later on in these lectures.  Finally, there is one more direction in which these topic is related to Maths apart from the geometrical side of it, statistics plays an important role too.  Physics is attractive because of its ability to ``tame'' natural phenomena in the sense that laws of physics bring order to the apparent chaos of Nature, Astrophysics is even more attractive because it evidences how laws of physics apply outside Earth, which is certainly surprising given the manifest differences between Earth and every other locations in the Universe of which we are aware. However, since these phenomena occur in places far, far away, and sometimes they can only be observed indirectly, there is typically an important degree of uncertainty. Thus, research is Astrophysics requires understanding  not only Physics, but also inference.  Inferential statistics  cares for the identification of patterns in the data taking in account the randomness and uncertainties in the observations. In contrast,  descriptive statistics is concerned with giving a summary of the data either numerically or graphically. Both will be needed toward two goals: we need to know  optimal ways to extract information from the astronomical data, but we also need to know rigorous approaches to compare theoretical predictions to observations.  Our approach will be that of Bayesian inference, and we will justify jut below my preference, but it must be admitted there is an old vivid controversy on the definition of probability between the two main schools: frequentists and Bayesians. Frequentists use a definition based on the possibility of repeating the experiment, but this does not apply to the Universe, and this sounds like the reason why the number of Bayesian astronomers grows every day, but you should not care for this battle right now; just stick to the idea that observations related Cosmology needs to resort to Statistics and that Inference is vital for constraining dark energy cosmological models.  By now you should be able to guess what to expect in the next sections. I will present you some basics of FRW (Friedmann-Robertson-Walker) cosmologies, then I will devote a great deal of this text to details of geometrical tests, and leave for the last but one section a convenient primer on statistics. ",
        "conclusions": "In this contribution I have tried to review some of the background on geometrical tests of dark energy models. A historical review of the development of Cosmology has been followed by an account of the importance of the discovery of the late-time acceleration, which is commonly attributed to the existence of an exotic component in the cosmic budget. Other possible explanations have been attempted, which required being open minded enough to admit the existence of extra dimensions.  Examples of both conventional (if that adjective can be used) and extradimensional models have been mentioned here with respect to their geometrical features. All these models can be subject to different observational tests which only require postulating a parametrization of the Hubble factor of the Universe or quantities derived from it. Here I have been concerned with four tests only: luminosity of supernovae, direct $H(z)$ measurements, the CMB shift and baryon acoustic oscillations. I cannot deny this selection is biased in the sense these are tests I have made research on, but the first and the last one are definitely very important and much effort is doing by research groups all over the world in finding the data their application requires.  This contribution has also tried to maintain a pedagogical tone as it has been prepared for the Advanced Summer School 2007 organized by Cinvestav in Mexico DF. I just hope it will be useful to any reader which happens to come across it."
    },
    "0710/0710.5918_arXiv.txt": {
        "abstract": "{} {We present the results of near-infrared, follow-up imaging and spectroscopic observations at VLT, aimed at characterizing the long-period companions of the exoplanet host stars HD\\,196885, HD\\,1237 and HD\\,27442. The three companions were previously discovered in the course of our CFHT and VLT coronographic imaging survey dedicated to the search for faint companions of exoplanet host stars.} {We used the NACO near-infrared adaptive optics instrument to obtain astrometric follow-up observations of HD\\,196885\\,A and B. The long-slit spectroscopic mode of NACO and the integral field spectrograph SINFONI were used to carry out a low-resolution spectral characterization of the three companions HD\\,196885\\,B, HD\\,1237\\,B and HD\\,27442\\,B between 1.4 and 2.5~$\\mu$m.} {We can now confirm that the companion HD\\,196885\\,B is comoving with its primary exoplanet host star, as previously shown for HD\\,1237\\,B and HD\\,27442\\,B. We find that both companions HD\\,196885\\,B and HD\\,1237\\,B are low-mass stars of spectral type M$1\\pm1$V and M$4\\pm1$V respectively. HD\\,196885\\,AB is one of the closer ($\\sim23$~AU) resolved binaries known to host an exoplanet. This system is then ideal for carrying out a combined radial velocity and astrometric investigation of the possible impact of the binary companion on the planetary system formation and evolution. Finally, we confirm via spectroscopy that HD\\,27442\\,B is a white dwarf companion, the third one to be discovered orbiting an exoplanet host star, following HD\\,147513 and Gliese~86. The detection of the broad Br$\\gamma$ line of hydrogen indicates a white dwarf atmosphere dominated by hydrogen.} {} ",
        "introduction": "The radial velocity (RV) technique is without contest nowadays the most successful method for detecting and characterizing the properties of exo-planetary systems. Since the discovery of 51 Peg (Mayor \\& Queloz 1995), more than 200 exo-planets have been identified featuring a broad range of physical (mass) and orbital (P, $e$) characteristics (Udry et al. 2003). Major progress has been made in improving detection performance and data analysis and has enabled us to explore the mass regime down to neptunian and telluric masses (Santos et al. 2004; McArthur et al. 2004). Hitherto, most surveys have been focused on solar-type stars because these stars show more and thinner lines than their more massive counterparts and less activity than their less massive ones, which ensures in both cases comparatively higher RV precision. Only recently have been planet-search programs devoted to late-type stars with the use of high precision spectrograph and repeated measurements (Delfosse et al. 1998; Endl et al. 2003; Wright et al. 2004; Bonfils et al. 2005) as well as to early-type stars with the development of new methods for radial velocity measurements (Chelli 2000; Galland et al. 2005).  Despite the success of this technique, the time span explored limits the study to the close ($\\le4-5$~AU) circumstellar environment. To understand the way exo-planetary systems form and evolve, it is therefore clearly worthwhile using complementary techniques such as pulsar timing, micro-lensing, photometric transit or direct imaging to fill out our knowledge. For solar analogs, the current deep-imaging capability is limited to the detection of massive brown dwarf (BD) companions. Typical separations larger than $50-100$~mas (i.e. $\\geq 3-5$~AU from a star at 50~pc) can be explored in the stellar regime and $0.5-1~\\!''$ (i.e. $\\geq 30-50$~AU from a star at 50~pc) in the substellar regime. The recent discovery of a T7 dwarf companion at 480~AU from the exoplanet host star HD\\,3651 (Mugrauer et al. 2006) illustrates the scope of deep, near-infrared imaging for resolving ultra-cool companions at large separations. Recent efforts have been also devoted to systematic search for stellar companions to nearby stars with and without planets, aimed at studying the impact of stellar duplicity on planet formation and evolution (Eggenberger et al. 2007).  Since 2003, we have conducted a deep-coronographic imaging survey of 26 exoplanet host stars, using PUEO-KIR at CFHT, and NACO at VLT (Chauvin et al. 2006). Three probable companions were detected around the stars HD\\,196885, HD\\,1237 and HD\\,27442 (see Tables~\\ref{tab:tab1} and \\ref{tab:tab2}). Follow-up observations were obtained for HD\\,1237\\,B and HD\\,27442\\,B, confirming their companionship (Chauvin et al. 2006; Raghavan et al. 2006). Based on their photometry, HD\\,196885\\.B and HD\\,1237\\,B are likely to be low-mass stars and HD\\,27442\\,B is probably the third white dwarf companion known to date orbiting an exoplanet host star. In this paper, we report new imaging and spectroscopic observations of these three companions, using NACO and SINFONI at VLT. In section~2, the instrument set-up and the data analysis are described. In section~3, we discuss the astrometric and spectroscopic results that confirm and refine their previous suggested nature.  \\begin{table}[t] \\caption{Characteristics of the observed exoplanet host stars HD\\,196885\\,A, HD\\,1237\\,A and HD\\,27442\\,A.} \\label{tab:tab1} \\centering \\begin{tabular*}{\\columnwidth}{@{\\excs}lllllll}     % \\hline\\hline Name         &  SpT         & d         &  Age$^a$  &   $\\rm{M}_{2}\\rm{sin}i$$^b$   &  $P\\,^b$  &   $e\\,^b$     \\\\ &              & (pc)      &  (Gyr)    &   $(M_{\\rm{Jup}})$        &  (days)  &            \\\\  \\\\ \\hline HD\\,196885\\,A   & F8IV         & 33.0      & 0.5       &   1.84                 & 386.0  & 0.30 \\\\ HD\\,1237\\,A     & G6V          & 17.6      & 0.8       &   1.94                 & 311.29 & 0.24\\\\ HD\\,27442\\,A    & K2IV         & 18.2      & 10        &   1.28                 & 423.84 & 0.07 \\\\ \\hline \\end{tabular*} \\begin{list}{}{} \\item[\\scriptsize{($^a$) AGE REFERENCES:}] \\scriptsize{Naef et al. 2001; Randich et al. 1999; Lambert et al. 2004} \\item[\\scriptsize{($^b$) RADIAL VELOCITY REFERENCES:}]\\scriptsize{Naef et al. 2001; Butler et al. 2001; http://exoplanets.org/esp/hd196885/hd196885.shtml} \\end{list} \\end{table} ",
        "conclusions": "We present the follow-up imaging and spectroscopic characterization of three long-period companions to the exoplanet host stars HD\\,196885, HD\\,1237 and HD\\,27442. The three objects were discovered during a previous deep-imaging survey carried out at CFHT and at VLT. Our new observations confirm their companionship unamiguously as well as their nature, which had previously been inferred from their photometry. Whereas HD\\,196885 and HD\\,1237 are two stellar companions of spectral type M$1\\pm1$V and M$4\\pm1$V respectively, HD\\,27442 is the second confirmed white dwarf companion of an exoplanet host star. The detection of the broad Br$\\gamma$ hydrogen line indicates a white dwarf atmosphere dominated by hydrogen.  HD\\,196885\\,AB is one of the closer resolved binaries known to host an exoplanet. The presence of this long-period companion may have played a key-role in the formation and the evolution of the inner planetary system, but how? The two main mechanisms of planetary formation, core accretion and disk instability, do not lead to the same predictions while the impact of a binary companion, depending on the separation and the mass ratio, is still debated. A complete dynamic characterization of nearby, tight binaries with planets and a dedicated imaging survey to study the multiplicity among stars with and without planets are clearly required to throw new light on the mechanisms of planetary formation and evolution."
    },
    "0710/0710.5129_arXiv.txt": {
        "abstract": "{ Unparticle $\\U$ with scaling dimension $d_\\U$ has peculiar thermal properties due to its unique phase space structure. We find that the equation of state parameter $\\omega_\\U$, the ratio of pressure to energy density, is given by $1/(2d_\\U +1)$ providing a new form of energy in our universe. In an expanding universe, the unparticle energy density $\\rho_\\U(T)$ evolves dramatically differently from that for photons. For $d_\\U >1$, even if $\\rho_\\U(T_{\\rm D})$ at a high decoupling temperature $T_{\\rm D}$ is very small, it is possible to have a large relic density $\\rho_\\U(T^0_\\gamma)$ at present photon temperature $T^0_\\gamma$, large enough to play the role of dark matter. We calculate $T_{\\rm D}$ and $\\rho_\\U(T^0_\\gamma)$ using photon-unparticle interactions for illustration. \\PACS{ {98.80.Cq}{Particle-theory and field-theory models of the early Universe} \\and {11.15.Tk}{Other nonperturbative techniques} \\and {11.25.Hf}{Conformal field theory, algebraic structures} \\and {14.80.-j}{Other particles} } % } % ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.1961_arXiv.txt": {
        "abstract": "In this paper, we present a detailed hydrodynamical study of the properties of the flow produced by the collision of a pulsar wind with the surrounding in a binary system. This work is the first attempt to simulate interaction of the ultrarelativistic flow (pulsar wind) with the nonrelativistic stellar wind. Obtained results show that the wind collision could result in the formation of an \"unclosed\" (at spatial scales comparable to the binary system size) pulsar wind termination shock even when the stellar wind  ram pressure exceeds significantly the pulsar wind kinetical pressure. Moreover, the post-shock flow propagates in a rather narrow region, with very high bulk Lorentz factor ($\\gamma\\sim100$). This flow acceleration is related to adiabatical losses, which are purely hydrodynamical effects. Interestingly, in this particular case, no magnetic field is required for formation of the ultrarelativistic bulk outflow. The obtained results provide a new interpretation for the orbital variability of radio, X-ray and gamma-ray signals detected from binary pulsar system PSR~1259-63/SS2883. ",
        "introduction": "Pulsars lose their rotation energy through  relativistic winds, the collision of which with the Interstellar Medium results in the formation of the  Pulsar Wind Nebulae (regions of nonthermal synchrotron radiation of ultrarelativistic electrons accelerated at the termination of  pulsar winds  \\citep{rees74,kc}). The Crab Nebula is the most famous example of such an object. The recent X-ray and TeV gamma-ray observations \\citep{GenSlane} show that this is a common phenomenon.  A very interesting situation arises when a pulsar is located in a binary system. In this case the pulsar wind interacts with the wind from the companion star.  This case, in particular, is realized in the binary system PSR1259-63/SS2883 which consists of a  $\\sim 48 \\rm ms$ pulsar in an elliptic orbit around a massive B2e optical companion \\citep{johnston1}. The density and  velocity of the stellar wind depend on the separation distance between two stars. Thus, the processes related to the interaction of two winds, in particular particle acceleration and radiation proceed under essentially different physical conditions depending on the orbital phase. This makes this object a unique laboratory for the study of nonthermal processes in ``on-line'' regime, due to the short acceleration and cooling time-scales characterizing this system, especially close to the periastron \\citep{khangoulian_psr}. Observations show that this system is indeed a strong source of nonthermal time-dependent emission extending from radio to TeV gamma-rays (see e.g. \\cite{neronov}).  Two variable TeV galactic gamma-ray sources, LS 5039 and LSI~61~303 (see e.g. \\cite{paredes_rev}), are discussed as possible, although less evident, candidates representing ``binary pulsar'' source population \\citep{dubus,mirabel,neronov2}. LS 5039 consists of O6.5V star and an unidentified compact object in a 3.9 day orbit.  This object has been detected as the source of gamma rays by EGRET (source 3EG J1824-1514) \\citep{ls5039EGRET} and by HESS \\citep{ls5039}. LSI~61~303 is a binary system with a B0Ve star in a 26.5 day orbit. This source presumably associates with a low energy gamma-ray source 2CG 135+01/3EG J0241+6103 \\citep{lsiaspulsar,tavani_lsi}. Recently, it has been detected in TeV gamma-rays  \\citep{lsi61303}. The nature of the compact objects (black hole or a neutron star) in both sources is not yet firmly established.   The collision of supersonic winds from two stars located in a binary system results in the formation of two terminating shock fronts and a tangential discontinuity separating relativistic and nonrelativistic parts of the shocked flow. It is believed that the shocked flow should propagate into a limited solid angle with rather high velocity. However the properties  of the flow downstream the shock  is unknown.  The collision of the pulsar wind with the supersonic flow of the nonrelativistic plasma has been recently numerically modeled by several groups \\citep{bucciantini,swaluw,vigelius,romero} using different nonrelativistic versions of hydrodynamical codes. However,  the dynamics of the relativistic plasma is rather specific, and the use of nonrelativistic codes cannot be \\textit{a priori} justified. The  modeling of  collision of a relativistic pulsar wind with a nonrelativistic one demands an adequate treatment of distinct features of relativistic outflows.  Moreover, in many previous studies the stellar wind has been approximated as plane parallel, while both winds initially expand radially. However, this purely geometric difference in the formulation of the problem, leads, in fact, to significantly different results and conclusions.  In this paper we present the results of our studies conducted in the hydrodynamical limit, i.e. the role of the magnetic field in dynamics has been ignored. The impact of the magnetic field will be published elswhere. ",
        "conclusions": "In this paper, we have conducted a detailed numerical study of the interaction of the relativistic and nonrelativistic winds assuming isotropic, radially expanding winds and ignoring the role of the magnetic field. In fact, the cold ultrarelativistic pulsar winds are generally believed to be highly anisotropic \\citep{bogovalov99,lyubarsky}. The anisotropy of the energy flux in the wind can result in a nonaxisymmetric form of the termination shock wave at the vicinity of the symmetry axis, as it follows from the studies of \\cite{vigelius}.  Whether the anisotropy has a noticeable impact on  the result and conclusions of this paper  is a subject of further investigations and will be discussed elsewhere. This concerns also the role of the magnetic field. Generally, as it follows from calculations of the interaction of the pulsar wind with the interstellar medium \\citep{crab,lk,bucciantini2}, the role of the magnetic field is rather important in the post shock region almost independent of the strength  of the field. This is explained by fast amplification of the magnetic field in the post shock region due to deceleration of the flow \\citep{Khangoulian}.  However, as shown in this paper, the flow is not decelerated. Just the opposite - it can be accelerated to large bulk motion Lorentz factor. Thus although one  should expect an impact  of the magnetic field on the post shocked flow, this affect most likely will not  be so strong as in the case of the interaction of the pulsar wind with interstellar medium.  Detailed MHD calculations are needed to to clarify the  role of the magnetic field.  The effect of reacceleration of the post shock flow to relativistic bulk motion Lorentz factors has direct implication to the interpretation of observations of high  energy  $\\gamma$- and X-rays from  binary pulsar systems like  PSR 1259-63/SS2883. This effect strongly modifies the  relationship between the synchrotron X-ray and inverse Compton gamma-ray fluxes produced  by the same population of relativistic electrons. It is well known that in the  pulsar wind nebulae (plerions) which are  formed, by the interaction of pulsar winds with the interstellar medium, the ratio between  the X-ray  and VHE $\\gamma$-ray fluxes are defined by the ratio of the energy density of the magnetic field to the energy density of soft radiation field, provided that  Compton scattering takes  place in the Thompson regime \\citep{atoyan}.   In binary systems inverse Compton scattering proceeds in the Klein-Nishina regime which changes the relationship between the  X-ray and VHE gamma-ray fluxes \\citep{khangulyan2}. Due to our calculations it becomes clear that  a significant deviation from the standard relations  should be expected also from hydrodynamics.  Let us consider these processes in more deatils  assuming that the magnetic field is present in the wind. Relativistic particles  moving in the magnetic field  usually produce synchrotron radiation. However,  if the wind is cold (in this case the velocity of particles  coincides with the bulk motion velocity)  these particles do not produce  synchrotron radiation. However, they can produce gamma-radiation with a specific sharp spectral feature  through the IC scattering  \\citep{bogovalov00}. An example of such a system is a cold pulsar wind which does not produce synchrotron radiation in the pre-shock region. In such a system the  ``standart\" relation between synchrotron and IC radiation components is violated. Remarkably,  the  possibility for the formation of relativistic flows in the post-shock region, as revealed in this paper, shows that in binary pulsar systems  we should expect a ``non-standard'' relation between synchrotron and IC radiation components in post-shock region as well.  Moreover, due to the large bulk motion Lorentz factor we should expect strong modulation of the observed nonthermal radiation of electrons. Indeed  in the case of binary-pulsar systems the direction of the post shock flow varies with the motion of the pulsar along the orbit around the  star. This implies significant changes of the Doppler factor, $\\delta$, given the large value of the Lorentz factor. Namely, $\\delta \\ll 1$ for large viewing angles $\\phi$ (e.g. close to $90^\\circ$) and $\\delta \\gg 1$ for small viewing angles. Correspondingly this will have a strong impact on the lightcurve of nonthermal radiation of electrons, $F_\\gamma \\propto \\delta^{n}$ where typically $n \\geq 3$. This effect, in particular,  can naturally explain the interesting feature of the nonthermal emission of PSR~1259/SS2883, the both synchrotron and inverse Compton components of which disappear during the periastron passage of the pulsar (see e.g. \\cite{neronov}). This interesting issue will be discussed elsewhere."
    },
    "0710/0710.3363_arXiv.txt": {
        "abstract": "{% Very high energy (VHE) $\\gamma$-ray observations have proven to be very successful in localizing Galactic acceleration sites of VHE particles. Observations of shell-type supernova remnants have confirmed that particles are accelerated to VHE energies in supernova blast waves; the interpretation of the $\\gamma$-ray data in terms of hadronic or leptonic particle components in these objects relies nevertheless strongly on input from X-ray observations. The largest identified Galactic VHE source class consists of pulsar wind nebulae, as detected in X-rays. Many of the remaining VHE sources remain however unidentified until now. With X-ray observations of these enigmatic ``dark'' objects one hopes to solve the following questions: What is the astrophysical nature of these sources? Are they predominantly electron or hadron accelerators? And what is their contribution to the overall cosmic ray energy budget? The paper aims to provide an overview over the identification status of the Galactic VHE source population.} ",
        "introduction": "Ground-based Cherenkov telescopes detect cosmic $\\gamma$-rays in the {\\em Very High Energy} (VHE, 100\\,GeV - 100\\,TeV) domain. In this frequency range, sources are visible in which charged particles are accelerated to TeV energies and beyond; those particles give rise to the detected $\\gamma$-ray emission. The energetic particles can be confined inside the objects, close to the actual acceleration site, like in young \\linebreak% shell-type supernova remnants (SNRs). In other cases, the particles diffuse out into the surrounding medium after their acceleration to high energies; in these cases, the size of the emitting region itself defines the extent of the ``source'' or object, like for example in a pulsar wind nebula (PWN).  The truly diffuse component of the high energy cosmic ray particles is predominantly being traced in the MeV-GeV domain, accessible to $\\gamma$-ray satellites. Lower energy diffuse {\\em electrons} are also traced through synchro\\-tron emission in the radio band. At VHE energies, localized sources dominate because of their intrinsically harder particle spectra compared to the diffuse component.  A successful VHE survey of the Galactic plane became possible with the high sensitivity and large field of view (FoV) of the H.E.S.S.\\ (High Energy Stereoscopic System) Cherenkov telescope system, with $F_{\\mathrm{min}}(>100\\,\\mathrm{GeV}) \\sim 4 \\times 10^{-12}\\mathrm{erg\\,cm^{-2}s^{-1}}$ for a $5\\,\\sigma$ point-source detection in 25 hrs, and a FoV of 3$^{\\circ}$ FWHM (Aharonian et al. 2006c). The H.E.S.S. telescope system became operational end of 2003, and provides through its location in Namibia an excellent view of the Galactic center region. The majority of the currently over 50 known Galactic VHE sources was discovered in the H.E.S.S. survey of the Galactic plane (Aharonian et al. 2005h, 2006g; Hoppe et al. 2007; Kosack et al. 2007), or in the H.E.S.S. FoV of observations of other targets in the Galactic plane (e.g. Aharonian et al. 2005c, 2007b).  Perhaps somewhat unexpectedly, many of the new \\linebreak% VHE-emitting sources can not readily be identified with \\linebreak% known astrophysical objects. Therefore, Galactic VHE astronomy does not only deal with the identification of particle acceleration mechanisms and efficiencies in well-known astrophysical objects, but also with the identification of \\linebreak% sources which are so far ``only'' defined through their $\\gamma$-ray emission. Follow-up observations of those sources with highly sensitive X-ray instruments such as {\\it XMM-Newton} provide the very promising possibility to identifiy those \\linebreak% enigmatic sources of high energy particles in our Galaxy. ",
        "conclusions": ""
    },
    "0710/0710.3649_arXiv.txt": {
        "abstract": "We examine the use of the CMB's TE cross correlation power spectrum as a complementary test to detect primordial gravitational waves (PGWs). The first method used is based on the determination of the lowest multipole, $\\ell_0$, where the TE power spectrum, $C_{\\ell}^{TE}$, first changes sign. The second method uses Wiener filtering on the CMB TE data to remove the density perturbations contribution to the TE power spectrum. In principle this leaves only the contribution of PGWs. We examine two toy experiments (one ideal and another more realistic) to see their ability to constrain PGWs using the TE power spectrum alone. We found that an ideal experiment, one limited only by cosmic variance, can detect PGWs with a ratio of tensor to scalar metric perturbation power spectra $r=0.3$ at $99.9\\%$ confidence level using only the TE correlation. This value is comparable with current constraints obtained by WMAP based on the $2\\sigma$ upper limits to the B-mode amplitude. We demonstrate that to measure PGWs by their contribution to the TE cross correlation power spectrum in a realistic ground based experiment when real instrumental noise is taken into account, the tensor-to-scalar ratio, $r$, should be approximately three times larger. ",
        "introduction": "Primordial gravitational waves (PGWs) polarize the cosmic microwave background (CMB) (see for example \\cite{basko, Polnarev85, Crittenden93, Frewin94, colefrewinpolnarev95, kamionkowski97, Seljak97Pol, selzal97, Kamionkowski98, baskaran06, Keating2006}). Current experiments are using the polarization of the CMB to search for this PGW background (\\cite{CloverIntro, BowdenOpt, BicepPaper}. This polarization can be used as a direct test of inflation. An alternative probe of the inflationary epoch which does not use the PGW background was studied by (\\cite{sperzal97}). This probe was used in recent analyses by the WMAP team (\\cite{peiris03}) to provide plausibility for the inflationary paradigm. This paper presents a test similar in spirit to that of \\cite{sperzal97}.  CMB polarization can be separated into two independent components: E-mode (grad) polarization and B-mode (curl) polarization. B-mode polarization can only be generated by PGWs (see for example \\cite{Seljak97Pol, selzal97, Kamionkowski98}), therefore most CMB polarization experiments which are searching for evidence of PGWs focus on measuring the BB power spectrum. However the TE cross correlation power spectrum offers another method to detect PGWs (\\cite{crittenden94}). The TE power spectrum is two orders of magnitude larger than the BB power spectrum and it was suggested that it may therefore be easier to detect gravitational waves in the TE power spectrum (\\cite{baskaran06, grishchuk07}).  In this paper we first discuss the method of detection of PGWs by measuring the TE power spectrum for low $\\ell$. This method, originally proposed in \\cite{baskaran06}, is based on a measurement of $\\ell_0$, the multipole where the TE power spectrum first changes sign. Hereafter we will call this method ``the zero multipole method''. The TE power spectrum due to density perturbations is positive on large scales, corresponding to $\\ell < \\ell_0$, changes sign at $\\ell = \\ell_0$, and then oscillates for $\\ell > \\ell_0$, while for PGWs the TE power spectrum must be negative for small $\\ell$ and then also oscillates for larger $\\ell$. The current best set of cosmological parameters, obtained in \\cite{Spergel2007WMAP}, gives, in the absence of PGWs, $\\ell_0 = 53$. Therefore, the measurement of the difference between the multipole number, $\\ell_0$, where the TE power spectrum changes sign, and $\\ell=53$ is the way to detect PGWs.  We will then consider an alternative method based on Wiener filtering, removing the contribution to the TE power spectrum due to density perturbations. Since the TE power spectrum due to PGWs is megative on large scales a test of negativity of the resulting TE power spectrum is a test of PGWs. In this paper, we present an analysis of both of these methods, based on Monte Carlo simulations.  At the present time, the main priority and the main challenge in CMB polarization observations is the detection of the PGW background via the BB power spectrum. In connection with BB experiments, the methods based on the TE cross correlation can be considered as very useful auxiliary measurements of PGWs because systematic effects in TE measurements are not degenerate with those in BB measurements. For example, T/B leakage or even E/B leakage could swamp a detection of BB, whereas T/E leakage would be small and well controlled (see \\cite{meir07}). These BB systematics could falsely indicate a detection of PGWs, but measurements of the TE power spectrum provide insurance against such a spurious detection. Additionally, galactic foreground contamination affects BB and TE in different ways, which enables us to perform powerful cross-checks and subtraction of foregrounds in BB measurements.  Another advantage of TE measurements for experiments which measure a small fraction of the sky, is related to the fact that a significant contaminant to the B modes is caused by E/B mixing. This limits the power spectrum of PGWs that can be detected (\\cite{challinor05}). The E-modes are practicly unaffected by E/B mixing so, in contrast to the BB measurements, the TE power spectrum should be nearly the same for both full and partial sky measurements.  The plan of this paper is the following. In Section \\ref{physmath}, we introduce the primordial power spectra of scalar (density) and tensor (PGW) perturbations (\\ref{primspectra}). Then following \\cite{crittenden94} and \\cite{baskaran06}, we explain why the sign of the TE power spectra for scalar and tensor perturbations is opposite for large scales (\\ref{TECross}). In Section \\ref{ell0det}, we describe in more detail the zero multipole method for the detection of PGWs. In Section \\ref{wienfilt}, we describe the method for detection of PGWs based on Wiener filtering along with the statistical tests used and a comparison of the tests. In Section \\ref{resul}, we present results of numerical Monte Carlo simulations for two toy experiments. In the first toy experiment we neglect instrumental noise and the uncertainties are limited only by cosmic variance (\\ref{toy1res}). In the second toy experiment, along with cosmic variance, we take into account instrumental noise which is comparable to real noise in current ground experiments (\\ref{toy2res}). For comparison, we also present results of simulations for the two satellite experiments, WMAP (\\ref{wmapres}) and  Planck (\\ref{planckres}). In Section \\ref{comptebb}, we compare the the signal-to-noise ratio of the TE measurements with those of BB measurements. ",
        "conclusions": "The measurement of where the TE cross correlation first changes sign can be used to detect or put constraints on PGWs. Such constraints are not as strong as the ones given by measurements of the BB power spectrum, however it is useful to have a supplementary method to detect PGWs. We have shown how well the TE mode can constrain the amount of PGWs from just a measurement of the angular scale where it first changes sign for two different toy experiments and two real satellite experiments. The absolute best limit with which we can measure $\\ell_0$ only gives us less than a $3\\sigma$ detection of the PGW component if $r=0.3$. The current confidence limits gives us $r < 0.3$ at $95\\%$ confidence level. Current and future experiments are optimized to measure the BB power spectrum if $r \\le 0.1$ even in the presence of foregrounds, which are not taken into account in this paper. Future satellite experiments should be able to detect $r<0.01$ which is $10$ times better than the sensitivity to $r$ than the result of the ideal experiment. If one neglects even cosmic variance, the discreteness of $\\ell$ limits the calculation of $\\ell_0$, and the sensitivity to $r$, to values considerably larger than $0.01$. The cosmic variance is largest at low $\\ell$ and is proportional to the total power spectrum. Since the TE cross correlation has contributions from density perturbations the errors in the measured TE power spectrum make detecting deviations of $\\ell_0$ from $53$ difficult, though they also provide insurance against a false detection or imperfect subtraction of instrumental and foreground systematic effects.  The other method described in this paper is one in which we filter out the signal due to density perturbations, leaving only the contribution to the TE power spectrum due to PGWs. We then test the resulting TE power spectrum to see if it is negative. Three different statistical tests were used to see if there was a significant detection of PGWs. The $S/N$ test can give a value for $r$ using a comparison with Monte Carlo simulations, while the Wilcoxon rank sum test can only give an allowable range for $r$. The sign test will only tell us if $r \\not= 0$.  Using the Wiener filtering method, we are unable to make as significant of a detection as using the zero multipole method. The best result was for the $S/N$ test which would give a $2.3\\sigma$ detection of $r=0.3$. To detect PGWs on the level of $3\\sigma$, the tensor-to-scalae ratio $r$ should be $r \\ge 0.4$. The sign test would give $2\\sigma$ detection for $r=0.3$ and a $3\\sigma$ detection for $r=0.45$. The Wilcoxon ranked sum test gives only a $1.2\\sigma$ detection for $r=0.3$ and a $3\\sigma$ detection for $r=0.7$. Similar results were gotten for the other three experiments tested. Thus in the sense of potential to detect PGWs, the zero multipole method is the best, next best is the $S/N$ test, then the sign test, and the worst is the Wilcoxon ranked sum test.  \\cite{baskaran06} present illustrative examples in which high $r$ is consistent with measured TT, EE, and TE correlations. The value of $r$ is so high in these examples that if PGWs with such $r$ really existed, current BB experiments would already detect PGWs. All models predict that the TE cross correlation power spectrum change sign only once for $\\ell < 100$. The fact WMAP cannot exclude several multipoles with $C_{\\ell}^{TE} > 0$ in between multipoles of $C_{\\ell}^{TE} < 0$ means that the TE cross correlation power spectrum either changes sign several times for $\\ell < 100$ or there is some instrumental noise which causes some anticorrelation measurements. Using instrumental noise consistent with WMAP, our Monte Carlo simulations give $\\Delta \\ell_0 \\approx 16$ and $\\ell_0 > 40$, which means that there is no evidence of PGWs in the TE correlation power spectrum."
    },
    "0710/0710.4296_arXiv.txt": {
        "abstract": "{ Weakly interacting massive particles (WIMPs) are one of the leading candidates for Dark Matter. So far we can use direct Dark Matter detection to estimate the mass of halo WIMPs only by fitting predicted recoil spectra to future experimental data. Here we develop a model--independent method for determining the WIMP mass by using experimental data directly. This method is independent of the as yet unknown WIMP density near the Earth as well as of the WIMP--nuclear cross section and can be used to extract information about WIMP mass with ${\\cal O}(50)$ events. \\PACS{ {95.35.+d}{Dark Matter}   \\and {29.85.Fj}{data analysis} } % } % ",
        "introduction": "\\label{intro} Today astrophysicists have strong evidence \\cite{evida}-\\cite{bullet} to believe that a large fraction (more than 80\\%) of the matter in the Universe is dark (i.e., interacts at most very weakly with electromagnetic radiation). The dominant component of this cosmological Dark Matter (DM) must be due to some yet to be discovered, non--baryonic particles. Weakly interacting massive particles (WIMPs) $\\chi$ with masses between 10 GeV and a few TeV are one of the leading candidates for DM \\cite{susydm}.  Currently, the most promising method to detect different WIMP candidates is the direct detection of the recoil energy deposited in a low--background laboratory detector by elastic scattering of ambient WIMPs off the target nuclei \\cite{detaa}, \\cite{detab}. The differential rate for elastic WIMP--nucleus scattering is given by \\cite{susydm}: \\beq \\label{eqn2101} \\dRdQ = \\calA \\FQ \\intvmin \\bfrac{f_1(v)}{v} dv\\, . \\eeq Here $R$ is the direct detection event rate, i.e., the number of events per unit time and unit mass of detector material, $Q$ is the energy deposited in the detector, $F(Q)$ is the elastic nuclear form factor, and $f_1(v)$ is the one--dimensional velocity distribution function of the WIMPs impinging on the detector. The constant coefficient $\\calA$ is defined as \\beq \\label{eqn2102} \\calA \\equiv \\frac{\\rho_0 \\sigma_0}{2 \\mchi m_{\\rm r,N}^2}\\, , \\eeq where $\\rho_0$ is the WIMP density near the Earth and $\\sigma_0$ is the total cross section ignoring the form factor suppression. The reduced mass $m_{\\rm r,N}$ is defined by \\beq \\label{eqn2103} m_{\\rm r,N} \\equiv \\frac{\\mchi \\mN}{\\mchi+\\mN}\\, , \\eeq where $\\mchi$ is the WIMP mass and $\\mN$ that of the target nucleus. Finally, $\\vmin$ is the minimal incoming velocity of incident WIMPs that can deposit the energy $Q$ in the detector: \\beq \\label{eqn2104} \\vmin = \\alpha \\sqrt{Q}\\, , \\eeq where we define \\beq \\label{eqn2105} \\alpha \\equiv \\sfrac{\\mN}{2 m_{\\rm r,N}^2}\\, . \\eeq  So far most theoretical analyses of direct WIMP detection have predicted the detection rate for a given (class of) WIMP(s), based on a specific model of the galactic halo. This can be used to estimate the mass of halo WIMPs only by fitting the predicted recoil spectra to future experimental data, e.g., \\cite{Green07}. The goal of our work is to develop methods which allow to extract information on halo WIMPs by using the experimental data directly. In our earlier work \\cite{DMDD} we used a time--averaged recoil spectrum, assumed that no directional information exists, and derived an expression for estimating moments of the normalized one--dimensional velocity distribution function of halo WIMPs: \\beqn \\label{eqn2208} \\expv{v^n} \\eqnequiv \\int_{v_{\\rm min}(\\Qthre)}^\\infty v^n f_1(v) \\~ dv \\non\\\\ \\= \\alpha^n \\bfrac{2 \\Qthre^{(n+1)/2} \\rthre+(n+1) I_n \\FQthre}{2 \\Qthre^{1/2} \\rthre+I_0 \\FQthre}\\, , \\non\\\\ \\eeqn where $Q_{\\rm thre}$ is the threshold energy of the detector, $\\rthre \\equiv (dR/dQ)_{Q = \\Qthre}$ is an estimated value of the scattering spectrum at $Q = \\Qthre$. $I_n$ can be either determined from a given expression (e.g., a fit to data) for the recoil spectrum: \\beq \\label{eqn2110} I_n = \\int_{\\Qthre}^{\\infty} \\frac{Q^{(n-1)/2}}{\\FQ} \\adRdQ \\~ dQ\\, , \\eeq or estimated directly from the measured recoil energy: \\beq \\label{eqn2210} I_n = \\sum_a \\frac{Q_a^{(n-1)/2}}{F^2(Q_a)}\\, , \\eeq where the sum runs over all events in the data set. Note that all these expressions are independent of the as yet unknown WIMP density near the Earth as well as of the WIMP--nucleus cross section. (More details about the reconstruction of the velocity distribution function of halo WIMPs and the estimate of its moments as well as all formulae needed can be found in Ref.~\\cite{DMDD}.) ",
        "conclusions": "In this paper we have presented a method which allows to extract information on the WIMP mass from the recoil energy measured in elastic WIMP--nucleus scattering experiments directly. In the long term this information can be used to constrain e.g., SUSY models in the elementary particle physics and compare with information from future collider experiments.  Our method for determining the WIMP mass by combining two (or more) experiments with different detector materials is independent of the as yet unknown WIMP density near the Earth as well as of the WIMP--nucleus cross section. The only information which one needs is the measured recoil energy.  Due to the maximal measuring energy of the detector, there will be a deviation of the reproduced WIMP mass from the true one. Nevertheless, for experiments with very few events and thus a quite large statistical error in the near future, this deviation should not affect the reproduced WIMP mass very significantly. Moreover, the numerical analysis shows also that, for WIMP masses $\\le 100~{\\rm GeV}/c^2$ some meaningful information on the WIMP mass can already be extracted from ${\\cal O}(50)$ total (each experiment ${\\cal O}(25)$) events.  The analyses of this work are based on several simplified assumptions. First, all experimental systematic uncertainties, as well as the uncertainty on the measurement of the recoil energy have been ignored. Comparing with large statistical uncertainty this should be a quite good approximation. Second, the analysis treats each recorded event as signal, i.e., background has been ignored altogether. This may in fact not be unrealistic for modern detectors. Third, each detector consists of a single isotope. This is quite realistic for the % semiconductor % detectors. For detectors containing more than one nucleus, by simultaneously measuring two signals, one might be able to tell on an event--by--event basis which kind of nucleus has been struck.  \\subsubsection*"
    },
    "0710/0710.0561_arXiv.txt": {
        "abstract": "{% We review the most important findings on AGN physics and cosmological evolution as obtained by  extragalactic X--ray surveys and associated multiwavelength observations. We briefly discuss the perspectives for future enterprises and in particular the scientific case for an extremely deep (2--3 Ms) XMM survey} ",
        "introduction": "Since their launch in 1999, both XMM--{\\it Newton} and {\\it Chandra} have performed a large ($>$ 30) number of surveys covering a wide fraction of the area vs. depth plane (see fig~1 in Brandt \\& Hasinger 2005). Thanks to vigorous programs of multiwavelength follow-up observations which, since a few years, have became customary, our understanding of AGN evolution has received a major boost. The high level of completeness in redshift determination for a large number of X--ray selected AGN (up to a few thousands) has made possible a robust determination of the luminosity function  and evolution of unobscured and mildly obscured AGN (Ueda et al. 2003; Hasinger et al. 2005; La Franca et al. 2005), which turned out to be luminosity dependent: the space density of bright QSO ($L_X > 10^{44}$ erg s$^{-1}$) peaks at z$\\sim$ 2--3, to be compared with the z$\\sim$0.7--1 peak of lower luminosity  Seyfert galaxies.  The fraction of obscured AGN is also strongly dependent from the X--ray luminosity (Ueda et al. 2003; La Franca et al. 2005). Even though the shape and the normalization of the function describing the obscured fraction vs. luminosity is still matter of debate, it is clear that absorption is much more common at low X--ray luminosities. Such a trend has been observed also in the optical (Simpson 2005) and in the near infrared (Maiolino et al. 2007) and may  be linked to the AGN radiative power which is able to ionize and expel gas and dust from the nuclear regions. More debated is the claim of an increase of the obscured fraction towards high redshifts. First suggested by La Franca et al. (2005), it has been confirmed by Treister \\& Urry (2006), while it is not required in the models discussed by Ueda et al. (2003) and Gilli et al. (2007).  A redshift dependence of the obscuring fraction would be naturally explained in the current framework of AGN formation and evolution (see, for example, Hopkins et al. 2006): the anti--hierarchical behaviour observed in AGN e\\-vo\\-lution (similar to that observed in normal galaxies), along with several other evidences, suggests that super\\-massive bl\\-ack holes (SMBH) and their host galaxies co--evolve and that their formation and evolution are most likely different aspects of the same astrophysical problem. At early times large quantities of cold gas were available to efficiently feed and obscure the growing black holes. Later on, the ionizing nuclear flux is able to \"clean\" its environment appearing as an unobscured QSO. However, the picture sket\\-ched above is likely to be much more complicated, depending on many other parameters (such as the BH mass, the Eddington ratio, the QSO duty cicle) and may not necessarily result in an increasing fraction of obscuration towards high redshifts.  It should also be remarked that sensitive X--ray observations are highly efficient to unveil weak and/or elusive accreting black holes which would be missed, or not classified as such, by surveys at longer wavelengths. Among them XBONG (Comastri et al. 2002), sources with high X--ray to optical flux ratio (Fiore et al. 2003), Extremely Red Objects (Brusa et al. 2005), Sub Millimeter Galaxies (Alexander et al. 2005). Although they are probably not representative of the sources of the X--ray background their study has allowed us to better understand the physics of accreting black holes.  Finally, the coverage of several {\\it Chandra} and XMM fie\\-lds has allowed to uncover several redshifts spikes in the distribution of X--ray sources (Gilli et al. 2005) and  demonstrated that AGN may be used as reliable tracers of the large scales structures. The underlying large scale structure may play an important role in triggering AGN activity. According to the analysis of a sample of X--ray selected AGN in the AEGIS survey, Georgakakis et al. (2006) suggest that $z \\sim$ 1 AGN are more frequent in dense environments. ",
        "conclusions": ""
    },
    "0710/0710.0757_arXiv.txt": {
        "abstract": "Analysis of the hydrodynamic and helioseismic effects in the photosphere during the solar flare of July 23, 2002, observed by Michelson Doppler Imager (MDI) on SOHO, and high-energy images from RHESSI shows that these effects are closely associated with sources of the hard X-ray emission, and that there are no such effects in the centroid region of the flare gamma-ray emission. These results demonstrate that contrary to expectations the hydrodynamic and helioseismic responses (''sunquakes\") are more likely to be caused by accelerated electrons than by high-energy protons. A series of multiple impulses of high-energy electrons forms a hydrodynamic source moving in the photosphere with a supersonic speed. The moving source plays a critical role in the formation of the anisotropic wave front of sunquakes. ",
        "introduction": "``Sunquakes\", the helioseismic response to solar flares, are caused by strong localized hydrodynamic impacts in the photosphere during the flare impulsive phase. The helioseismic waves are observed directly as expanding circular-shaped ripples in SOHO/MDI Dopplergrams, which can be detected in Dopplergram movies and as a characteristic ridge in time-distance diagrams, \\citep{Kosovichev1998, Kosovichev2006a}, or indirectly by calculating integrated acoustic emission \\citep{Donea1999, Donea2005}. Solar flares are sources of high-temperature plasma and strong hydrodynamic motions in the solar atmosphere. Perhaps, in all flares such perturbations generate acoustic waves traveling through the interior. However, only in some flares the impact is sufficiently localized and strong to produce the seismic waves with the amplitude above the convection noise level. It has been established in the initial July 9, 1996, flare observations \\citep{Kosovichev1998} that the hydrodynamic impact follows the hard X-ray flux impulse, and hence, the impact of high-energy electrons. Nevertheless, a common paradigm is that the sunquake events are caused by accelerated protons because protons carry more momentum and penetrate deeper into the solar atmosphere than electrons, which loose most of their energy in the upper chromosphere. This paradigm is not easy to test because the gamma-ray emission, which indicates the presence of high-energy protons, is rarely observed.  In a large X17 flare of October 28, 2003, the gamma-ray emission observed by RHESSI was located close to the hard X-ray sources and two of the three places of the phototospheric impacts (sunquake sources) \\citep{Kosovichev2006a}. Because of the close locations of the hard X-ray and gamma-ray sources these observations could not exclude the possibility of the proton or mixed electron-proton impacts \\citep{Zharkova2007}.  However, in one event, X4.8 flare of July 23, 2002, the hard X-ray and gamma-ray sources were significantly separated from each other. The centroid of the $\\gamma$-ray 2.233 MeV neutron-capture emission was found to be displaced by $20\"\\pm 6\"$ (with 5-sigma confidence) from that of the $0.3-0.5$ MeV X-ray emission implying a difference in acceleration and/or propagation between the accelerated electrons and ions \\citep{Hurford2003}. Therefore, this flare provides a unique opportunity to investigate the photospheric and helioseismic responses separately for high-energy electrons and protons. In this Letter, I present results of the analysis of the relationship between the hard X-ray and gamma-ray emissions and the hydrodynamic and seismic signals in the photosphere, using data from RHESSI \\citep{Lin2002} and MDI on SOHO \\citep{Scherrer1995}. RHESSI provides X-ray/gamma-ray imaging spectroscopy from 3 keV to 17 MeV with angular resolution $2.3''-3'$ ($35''$ at gamma-ray energies) over the full Sun.  MDI measures the Doppler velocity and the line-of-sight magnetic field of the photospheric plasma every minute with 2 arcsec/pixel resolution also over the full Sun. ",
        "conclusions": "The analysis of RHESSI X-ray and gamma-ray images and SOHO/MDI Dopplergrams of the July 23, 2002, X4.8 solar flare revealed that the hydrodynamic and seismic responses are closely associated the hard X-ray emission, both spatially and temporally, but showed no significant responses in the gamma-source centroid area. Because this flare was one of strongest gamma-flares, and the hard X-ray and gamma-ray sources were separated, these observations show that the accelerated protons are unlikely to be a source of the hydrodynamic response and sunquakes. Furthermore, the detailed analysis of the dynamics of sunquake sources in this Letter and in the paper by \\citet{Kosovichev2006a} reveals their close association with expanding flare ribbons and rapid HXR source motion along the ribbons, and, thus, with the magnetic reconnection process. The fast motion of these sources results in strong anisotropy of the seismic waves, clearly observed in the MDI data.  The general picture that comes from the analysis of MDI and RHESSI data is consistent with the previously developed hydrodynamic thick-target model, illustrated in Fig.~\\ref{fig2} \\citep{Kostiuk1975, Livshits1981, Fisher1985, Kosovichev1986}.  In this model, high-energy electrons heat the upper chromosphere to high temperatures generating a high-pressure region, expansion of which causes evaporation of the chromospheric plasma and a high compression shock. The shock reaches the photosphere and excites the seismic waves. However, the new results show that it is important to include effects of the multiple impact and moving source in the thick-target and sunquake models.  The photospheric and helioseismic effects observed during the impulsive phase of solar flares are closely related to the processes of acceleration and propagation of electrons and ions, and may provide new important information about these processes. \\clearpage"
    },
    "0710/0710.4518_arXiv.txt": {
        "abstract": "We investigate the birth and evolution of isolated radio pulsars using a population synthesis method, modeling the birth properties of the pulsars, their time evolution, and their detection in the Parkes and Swinburne Multibeam (MB) surveys. Together, the Parkes and Swinburne MB surveys \\citep[][]{2001MNRAS.328...17M, 2001MNRAS.326..358E} have detected nearly 2/3 of the known pulsars and provide a remarkably homogeneous sample to compare with simulations. New proper motion measurements \\citep[][]{2002ApJ...571..906B, 2003AJ....126.3090B} and an improved model of the distribution of free electrons in the interstellar medium, NE2001 \\citep[][]{NE2001}, also make revisiting these issues particularly worthwhile. We present a simple population model that reproduces the actual observations well, and consider others that fail. We conclude that: pulsars are born in the spiral arms, with the birthrate of $2.8\\pm0.5$ pulsars/century peaking at a distance $\\sim3$ kpc from the Galactic centre, and with mean initial speed of $380^{+40}_{-60}$ km s$^{-1}$; the birth spin period distribution extends to several hundred milliseconds, with no evidence of multimodality, implying that characteristic ages overestimate the true ages of the pulsars by a median factor >2 for true ages <30,000 yr; models in which the radio luminosities of the pulsars are random generically fail to reproduce the observed $P-\\dot{P}$ diagram, suggesting a relation between intrinsic radio luminosity and $(P, \\dot{P})$; radio luminosities $L \\propto \\sqrt{\\dot{E}}$ provides a good match to the observed $P-\\dot{P}$ diagram; for this favored radio luminosity model, we find no evidence for significant magnetic field decay over the lifetime of the pulsars as radio sources ($\\sim100$ Myr). ",
        "introduction": "From soon after the discovery of the pulsars \\citep[][]{1968Natur.217..709H}, their Galactic population has been the focus of numerous studies \\citep[][]{1970ApJ...160..979G, 1971IAUS...46..165L, 1977ApJ...215..885T, 1977MNRAS.179..635D, 1985MNRAS.213..613L, 1987A&A...171..152S, 1989ApJ...345..931E, 1992A&A...254..198B, 1993MNRAS.263..403L, 1997A&A...322..127H, 1997MNRAS.289..592L, 2002ApJ...568..289A, 2002ApJ...565..482G, 2004ApJ...604..775G}. Nevertheless, in spite of much progress, many outstanding questions remain. What are the birth positions, velocities, spin periods, and magnetic fields of the pulsars? How do these evolve in time and how are they related to the radio luminosities of the pulsars? One approach to answering these questions is to make use of the statistical power of the growing pulsar catalogue to study the pulsar population as a whole.\\\\ \\\\ Many recent advances in pulsar astronomy make it particularly worthwhile to revisit the above questions through population synthesis. The recently completed Parkes and Swinburne Multibeam (PMB and SMB; \\cite{2001MNRAS.328...17M, 2001MNRAS.326..358E}) surveys have detected nearly 2/3 of the known pulsars and provide a large and remarkably homogeneous observed sample to compare with simulations. New proper motion measurements \\citep[][]{2002ApJ...571..906B, 2003AJ....126.3090B} and an improved model of the interstellar medium \\citep[][]{NE2001} also provide valuable new information.\\\\ \\\\ In this work, the details of which have been reported by Faucher-Gigu\\`ere \\& Kaspi (2006) \\cite{2006ApJ...643..332F}, we investigate the birth properties of Galactic isolated radio pulsars and their time evolution. To do so, we generate an ensemble of mock galaxies populated by pulsars with prescribed birth properties (spatial locations, velocities, spin periods, radio luminosities, magnetic fields) and evolve the pulsars in time using physical models. We model the selection function of the PMB and SMB surveys using a modified version of the radiometer equation \\citep[][]{1985ApJ...294L..25D} and apply it to our mock galaxies. We compare the observed histograms of Galactic longitudes, latitudes, dispersion measures, 1.4 GHz radio fluxes, pulse periods, magnetic fields, as well as the observed $P-\\dot{P}$ diagrams, to judge how well each population model reproduces the actual observations.  \\begin{figure*} \\includegraphics[height=.25\\textheight]{ppdot_diagrams.eps} \\caption{Comparison of the $P-\\dot{P}$ diagram for the pulsars detected in the Parkes and Swinburne Multibeam surveys (left) with the corresponding diagrams in our best simulation with $L\\propto \\sqrt{\\dot{E}}$ (middle) and with a simulation in which the radio luminosities of the pulsars are uncorrelated with their other characteristics (``random''; right). While the simulation with $L\\propto \\sqrt{\\dot{E}}$ reproduces the observed diagram remarkably well, the model with random radio luminosities produces a pile-up of pulsars near the death line that is in clear disagreement with the observations.} \\label{ppdot diagrams} \\end{figure*} ",
        "conclusions": ""
    },
    "0710/0710.0494_arXiv.txt": {
        "abstract": "We study the color structure of disk galaxies in the Groth strip at redshifts $0.1<z<1.2$.  Our aim is to test formation models in which bulges form before/after the disk.   We find smooth color distributions with gentle outward blueing across the galaxy image: bulges are not distinctly redder than their disks; and bulge colors strongly correlate with global colors. The results suggest a roughly coeval evolution of bulges and disks. About 50\\% of the nuclei of galaxies with central light excesses above the outer exponential profile hold passively evolving red populations.  The remainder 50\\% are galaxies with central blue colors similar to their disks.  They may be bulges in formation, or the central parts of disks with non-exponential surface brightness profiles. ",
        "introduction": "This work addresses the simple question: \\textsl{are bulges older than their host disks?}.   The hope is to be able to falsify one or more of the hypothesis for bulge formation: bulges before disks, via mergers or primordial collapse, \\textit{vs.}\\ bulges after disks, from disk instabilities or other secular processes. The plan is admitedly na\\\"ive.  A prevalence of blue bulges would argue against an old formation age for bulges.  But bulges redder than their disks would not rule out the disk instability model, as the younger disk age might come from subsequent star formation in the disk after the bulge is formed.  However, the question is a useful guideline, especially when combined with structural and isophotal diagnostics.  For early- to intermediate-type disk galaxies in the nearby Universe, \\cite[Peletier \\& Balcells (1996)]{Peletier96} concluded that the age differences between bulges and disks was in most cases consistent with zero. But age dating of stellar populations has large uncertainties.  We have repeated the Peletier \\& Balcells study at redshifts up to $z=1$, closer to the expected formation epoch of bulges.  An extensive analysis is presented in \\cite[Dom\\'{\\i}nguez-Palmero \\& Balcells (2008)]{DominguezPalmero08}. ",
        "conclusions": ""
    },
    "0710/0710.2491_arXiv.txt": {
        "abstract": "We present optical, IR and millimeter observations of the solar-type star 13-277, also known as GM Cep, in the 4 Myr-old cluster Tr 37. GM Cep experiences rapid magnitude variations of more than 2 mag at optical wavelengths. We explore the causes of the variability, which seem to be dominated by strong increases in the accretion, being similar to EX-or episodes. The star shows high, variable accretion rates (up to $\\sim$10$^{-6}$ M$_\\odot$/yr), signs of powerful winds, and it is a very fast rotator (V$sini \\sim$43 km/s). Its strong mid-IR excesses reveal a very flared disk and/or a remnant envelope, most likely out of hydrostatic equilibrium. The 1.3 millimeter fluxes suggest a relatively  massive disk (M$_D \\sim$0.1 M$_\\odot$). Nevertheless, the millimeter mass is not enough to sustain increased accretion episodes over large timescales, unless the mass is underestimated due to significant grain growth. We finally explore the possibility of GM Cep having a binary companion, which could trigger disk instabilities producing the enhanced accretion episodes. ",
        "introduction": "}  The open cluster Tr 37 is one of the best studied intermediate-aged young star formation regions. Aged $\\sim$ 4 Myr (Sicilia-Aguilar et al. 2004, 2005; hereafter Paper I, Paper II), and located at 900 pc distance (Contreras et al. 2002), it contains a rich population of T Tauri stars (TTS; $\\sim$180 members with spectral types G to M2) of which about 48\\% still show IR excesses consistent with protoplanetary disks at different evolutionary stages (Sicilia-Aguilar et al. 2006a, b; hereafter Paper III, Paper IV). While most of the stars in Tr 37 show important evidence of disk evolution, with lower accretion rates and near-IR excesses than in younger regions, a few objects still display characteristics of much younger systems. The most remarkable one is the solar-type star 13-277, also called GM Cep (Morgenroth 1939). Because of its spatial location within the cluster, it belongs most likely to the Tr 37 main population, with average ages $\\sim$4 Myr ($\\sim$85\\% of the population is older than 2 Myr, and $\\sim$95\\% is older than 1 Myr), rather than to the young population associated to the Tr 37 globule, aged $\\sim$1 Myr (Paper II). Its continuum spectrum and the strong and broad H$\\alpha$ emission suggested an accretion rate \\.{M}$\\sim$3 10$^{-7}$ M$_\\odot$/yr, about 2 orders of magnitude over the median accretion rate of TTS in Tr 37 (Paper IV). Its bolometric luminosity  (L$\\sim$26 L$_\\odot$ in 2000) is about 1 order of magnitude higher than other stars with similar late G-early K spectral type. Finally, we notice an unusually high mid-IR flux for a solar-type star: its MIPS flux at 70 $\\mu$m is comparable to that of the Tr 37 Herbig Be star MVA-426, being one of the only four cluster members detected at this wavelength (Paper III).  All observations suggest that GM Cep is a variable star of the EX-or type, with an unstable disk and variable accretion rate, which is remarkable in an ``old'' cluster like Tr 37, where disk evolution seems ubiquitous. FU-ors and EX-or objects have been suggested to be either normal stages within the very early TTS evolution (see Hartmann \\& Kenyon 1996 for a review) or a special type of young object, probably binary (Herbig 2003 and references therein). Here we present the results of our photometry monitoring campaign during 2006-2007\\footnote{Based on observations collected at the German-Spanish Astronomical Center, Calar Alto, jointly operated by the Max-Planck-Institut f\\\"{u}r Astronomie Heidelberg and the  Instituto de Astrof\\'{\\i}sica de Andaluc\\'{\\i}a (CSIC)}, lucky imaging, millimeter continuum, and $^{12}$CO data, together with optical spectra (high- and low-resolution), and a compilation of the data available from the literature. The multiwavelength data are described in Section \\ref{observations}. In Section \\ref{analysis} we explore the characteristics of the star: optical and IR variability, spectral type, accretion, winds, disk mass, and potential companions. Section \\ref{outburst} discusses the possible causes of variability and the comparison with similar stars, and in Section \\ref{conclu} we summarize our results. ",
        "conclusions": "}  GM Cep is an extremely variable late G star in the 4 Myr-old cluster Tr 37, resembling younger EX-or objects. Its complex variability in color and magnitude is not consistent with a single typical variability mechanism in TTS (hot/cold spots, variable obscuration, or changes in the accretion rate). The amplitude of the variations, the high accretion rate, the luminous mid-IR disk, and the high stellar luminosity suggest variable accretion to be the stronger contributor, maybe mixed with variable extinction by circumstellar material (similar to UX-or variables), and some minor influence of cold spots and scattering. Increased accretion episodes are thought to produce the strong, irregular variations in young EX-or and FU-or objects, typically still surrounded  by massive infalling envelopes. Nevertheless, GM Cep has a medium-mass disk and a small or inexistent envelope, and belongs to a cluster where strong disk evolution is ubiquitous. Large changes in the accretion  rates can result in changes in the stellar and disk structure, and conversely, the disruption of the disk structure (for instance, by companions, or via gravitational instability if the disk is very massive and compact), can increase the accretion rate. The presence of companions (stellar, substellar or planetary) is a plausible mechanism to produce disk instabilities in a relatively old star like GM Cep. Simultaneous multiwavelength observations, including IRS spectra scheduled for 2007, providing coverage in the 5-35$\\mu$m range, and detailed monitoring in the coming years will help us to reveal the nature of the variability in GM Cep. Sub-millimetre observations in the future should be used to determine the SED slope and constrain the size of the grains. Old but extremely accreting stars are thus a key to understand the processes of disk accretion and evolution and the consequences for planet formation."
    },
    "0710/0710.5087_arXiv.txt": {
        "abstract": "We report on the results of spectroscopic mapping observations carried out in the nuclear region of Centaurus A (NGC5128) over the 5.2 - 15 and 20-36$\\mu$m spectral regions using the Infrared Spectrograph on the {\\it Spitzer Space Telescope}.  We have detected and mapped S(0), S(2), S(3), and S(5) pure rotational transition lines of molecular hydrogen and emissions in the fine-structure transitions of [SiII], [SIII], [FeII], [FeIII], [ArII], [SIV], [NeII], [NeV] and [OIV]. The 500 pc bipolar dust shell discovered by Quillen et al.(2006) is even more clearly seen in the 11.3$\\mu$m dust emission feature than previous broad band imaging. The pure rotational lines of molecular hydrogen other than the S(0) line are detected above the dusty disk and associated with the oval dust shell. The molecular hydrogen transitions indicate the presence of warm gas at temperatures 250--720K. The column density of the warm molecular hydrogen in the shell is $N(H_2) \\sim 10^{20}$cm$^{-2}$ and similar to that estimated from the continuum dust shell surface brightness. The ratio of the dust emission features at 7.7$\\mu$m and  11.3$\\mu$m  and the ratio of the [NeII](12.8$\\mu$m) and 11.3$\\mu$m dust emission feature are lower in the 500 pc dust shell than in the star forming disk. The clearer shell morphology at 11.3$\\mu$m, warm molecular hydrogen emission in the shell, and variation in line ratios in the shell compared to those in the disk, confirm spectroscopically that this shell is a separate coherent entity and is unlikely to be a chance superposition of dust filaments. The physical conditions in the shell are most similar to Galactic supernova remnants where blast waves encounter molecular clouds. The lines requiring the highest level of ionization, [NeV]($24.318\\mu$m) and [OIV]($25.890\\mu$m), are detected 20--25\\arcsec north-east and south-west of the nucleus and at position angles near the radio jet axis. Fine structure line ratios and limits from this region suggest that the medium is low density and illuminated by a hard radiation field at low ionization parameter. These higher S molecular hydrogen pure rotational transitions are also particularly bright in the same region as the [OIV] and [NeV] emission. This suggests that the gas associated with the dust shell has been excited near the jet axis and is part of an ionization cone. ",
        "introduction": "The nearest of all the giant radio galaxies, Centaurus~A (NGC~5128) provides a unique opportunity to observe the dynamics and morphology of an active galaxy in detail across the electromagnetic spectrum. For a recent review on this remarkable object see \\citet{israel}. In its central regions, NGC~5128 exhibits a well recognized and optically-dark band of absorption across its nucleus. Images from the {\\it Spitzer Space Telescope} (SST) with the Infrared Array Camera (IRAC) and Multiple Imaging Photometer for Spitzer (MIPS) cameras in the mid-infrared reveal a 3\\arcmin long parallelogram shape \\citep{quillen_irwarp} that has been modeled as a series of folds in a dusty warped thin disk (e.g., \\citealt{bland86,bland87,nicholson92,sparke96,quillen_irwarp}). The dusty disk is also the site of ongoing star formation at a rate of about $1M_\\odot$~yr$^{-1}$, (based on the infrared luminosity estimated by \\citealt{eckart90}). Centaurus~A hosts an active nucleus (e.g., \\citealt{whysong04,mirabel99}) that has been recently studied using infrared spectra from the SST. \\citep{weedman05,gorjian07}. Its nucleus exhibits a strong silicate absorption feature and emission from [NeV]. Here we do not study the nucleus but focus on structure exterior to it.  The IRAC and MIPS images of Centaurus~A have revealed another surprise, a bipolar shell-like structure 500\\,pc north and south of the nucleus \\citep{quillen_supershell}. This shell, seen for the first time with Spitzer imaging, could be the first extragalactic nuclear shell discovered in the infrared.  In this paper we present mapping of the central 500--600pc of Centaurus A done with the Infrared Spectrograph (IRS) on board the SST. The infrared spectral maps were obtained to test the possibility that the apparent dust shell is a separate coherent structure and not a chance superposition of dust filaments. They were also obtained to search for possible interactions between the AGN and the interstellar medium either by illumination from the AGN or caused by jets or outflows.  Based on the discussion by \\citet{israel}, we adopt a distance to Cen A of 3.4~Mpc.  At this distance, 1~kpc corresponds to 1\\arcmin on the sky. All positions reported in this manuscript are given with respect to Epoch 2000. ",
        "conclusions": "We have carried out a spectroscopic study of the central 2 arcminutes of Centaurus A using short low and long high spectral modules of the Infrared Spectrograph on board the {\\it Spitzer Space Telescope}. Most of the emission lines detected in the spectral cubes (e.g., [SIII](33.5$\\mu$m), [SiII](34.8), [FeII](26.0), [FeIII](23.9), [ArII](6.98) and H$_2$S(0)(28.2)) and dust emission features primarily trace regions of star formation in the warped disk.  Our previous study based on IRAC and MIPS imaging suggested that Centaurus A hosts an oval or bipolar dustshell at a distance approximately 500 pc  from the nucleus seen above and below the warped disk. This dust shell, if confirmed, would be the first extragalactic shell to be discovered in the infrared. Here we see the dust shell even more clearly and prominently in the 11.3$\\mu$m dust emission feature than we saw previously in the broad band IRAC images. We have found variations in the dust emission feature 7.7$\\mu$m/11.3$\\mu$m ratio and dust 11.3$\\mu$m/[NeII](12.8$\\mu$) ratio, with the oval dust shell having the lowest ratios compared to the star forming disk. The clearer shell morphology at 11.3$\\mu$m than previously seen in broad band images, the association of the molecular hydrogen emission in the shell, and the variation in line ratios in the shell compared to those in the disk, confirm spectroscopically that the shell discovered previously \\citep{quillen_supershell} is a separate coherent entity and is unlikely to be a chance superposition of dust filaments.  We find evidence for higher ionization species line emission in [NeV]($24.3\\mu$m) and [OIV]($25.9\\mu$m) near the jet axis. Emission in these two lines is seen both north-east and south-west of the nucleus along position angles $\\sim 40^\\circ$ and $\\sim -120^\\circ$.  These angles are similar to but not exactly the same as the jet axis at 55$^\\circ$ as seen at 5GHz. Outside the nucleus, the peak surface brightness in these lines is 25\\arcsec or 400~pc south-west of the nucleus. Emission line ratios and limits at the location of the [OIV] peak suggest that the emitting region is at low ionization parameter, $U \\lesssim 10^{-2}$,  and has low electron density, $n_e \\lesssim 10^2$cm$^{-3}$. We crudely estimate that the AGN can provide sufficient UV photons to account for the [OIV] luminosity 400~pc from the nucleus, as long as UV photons are not absorbed by intervening material as they travel from the AGN to the dust shell. A much more detailed photo-ionization study is required to understand the excitation of the [OIV] and [NeV] emission. Previous reports of an ionization cone in Cen A were based on near-infrared imaging of the central few arcseconds \\citep{bryant99}. Subsequent studies interpreted line emission in terms of a disk rather than ionization cone \\citep{schreier98,krajnovic06}. Unfortunately the morphology of the near-infrared images is strongly affected by extinction and the warp disk models have not been good enough to accurately predict the extinction in the central few arcsecond of the nucleus. Cen A might be the only active galaxy in which mid-infrared spectroscopy has found evidence for high ionization species such as NeV at hundred pc distances from the nucleus. As far as we know Cen A hosts the only ionization cone that has been resolved with observations from the Spitzer Space Telescope.   We see evidence for warm molecular hydrogen coincident with the peak in [OIV] in the odd pure rotational odd transitions S(3) and S(5). The S(7) and S(2) transitions are also detected but at weaker levels. The S(3) and S(5) emission also lies in the vicinity of the dust shell that is most prominent in the 11.3$\\mu$m dust (PAH) emission feature. A two temperature component model can fit the rotational line ratios and implies that there is warm molecular hydrogen with temperatures in the range 250--720K. The temperatures are warmer than seen in nucleus of non-active nearby galaxies, similar to those of LINERS and Seyferts but lower than exhibited by supernova remnants ULIRGS (as compared to studies by \\citealt{roussel07}, \\citealt{higdon06} and \\citealt{neufeld07}). Near the jet axis, the column depth of warm molecular hydrogen is $N(H_2) \\sim 10^{20}$cm$^{-2}$ similar to that estimated from the infrared continuum emission of the dust shell by \\citet{quillen_supershell}. This suggests that gas associated with the dust shell has been heated near the jet axis. There is probably more than one excitation process as [OIV] and [NeV] emission are detected primarily along the jet axis whereas the higher S pure rotational molecular hydrogen lines are detected there and in the vicinity of the dust shell.  Previous studies of the pure-rotational molecular hydrogen lines in extra-galactic objects (e.g., \\citealt{higdon06,panuzzo07,roussel07,ogle07}) have not well resolved the emission. The association of the warm molecular hydrogen gas with a shell is most similar to phenomena exhibited by Galactic supernova remnants where the blast wave encounters molecular clouds \\citep{neufeld07}. The physical conditions estimated from the molecular hydrogen observations are similar in properties (column depth, temperatures and fraction of gas in the two temperature components) to the parameters estimated by \\citet{neufeld07} in 4 Galactic supernova remnants. This suggests that theory of interstellar shock waves be applied to interpreting observations of the shell. \\citet{neufeld07} associates the warmer molecular hydrogen component responsible for the higher S pure rotational transitions with dissociative shocks. These require shock velocities $\\gtrsim 70$km/s (e.g., \\citealt{drainemckee93}). A physical scenario and model accounting for the shell's structure and energetics is currently lacking. Deep optical spectroscopic and radio studies are particularly needed to better constrain gas motions and physical conditions in this shell.    \\vskip 1.0truein We thank Dan Watson, Paul van der Werf,  Ralph Kraft, Jacqueline van Gorkom, Martin Hardcastle,  and Christine Jones-Forman for helpful discussions and correspondence. We thank Martin Hardcastle for providing us with images of Centaurus A in the radio. Support for this work was in part provided by by NASA through an award issued by JPL/Caltech, National Science Foundation grants AST-0406823 $\\&$ PHY-0552695, the National Aeronautics and Space Administration under Grant No.$\\sim$NNG04GM12G issued through the Origins of Solar Systems Program, and HST-AR-10972 to the Space Telescope Science Institute. JBH is funded by a Federation Fellowship from the Australian Research Council.  \\clearpage   \\begin{figure} \\includegraphics[angle=0,width=3.5in]{cont25.eps} \\includegraphics[angle=0,width=3.5in]{cont32.eps} \\caption{ \\label{fig:cont25} Continuum emission from the LH spectral cube from $0.4\\mu$m wide bands. Contours are evenly spaced with the lowest contour and spacing at 0.01 MJy/SR. a) Continuum centered at $25.0\\mu$m. b) Continuum centered at $32.0\\mu$m. } \\end{figure}   \\begin{figure} \\includegraphics[angle=0,width=3.5in]{PAH11.eps} \\caption{ \\label{fig:PAH11} Dust emission feature at 11.3$\\mu$m. No continuum has been subtracted as the emission feature dominates the spectrum by a factor of 3--8. The minimum contour and spacing is approximately 10 MJy/SR in the peak of the line. The black contours show the oval dust shell previously described by \\citet{quillen_supershell}. } \\end{figure}  \\clearpage  \\begin{figure} \\includegraphics[angle=0,width=3.5in]{SIIIc.eps} \\includegraphics[angle=0,width=3.5in]{SiIIc.eps} \\caption{ \\label{fig:Ssi} Continuum subtracted line emission images in the LH spectral cube, showing emission in the folded star forming disk. Contours are evenly spaced. The lowest contours and spacing for the [SIII]($33.481\\mu$m) and [SiII]($34.815$m) images are 1.0  and 1.5 $\\times 10^{-8}$erg cm$^{-2}$ s$^{-1}$ SR$^{-1}$, respectively. a) For [SIII]($33.481\\mu$m). b) For [SiII]($34.815\\mu$m). } \\end{figure}   \\begin{figure} \\includegraphics[angle=0,width=3.5in]{FeIIc.eps} \\includegraphics[angle=0,width=3.5in]{FeIIIc.eps} \\caption{ \\label{fig:Fe} Continuum subtracted line emission images in the LH spectral cube, showing emission in the folded star forming disk. Contours are evenly spaced. The lowest contours and spacings  for [FeII](25.988$\\mu$m) and [FeIII](22.925$\\mu$m) images with lowest contour are 0.05  and 0.025 $\\times 10^{-8}$erg cm$^{-2}$ s$^{-1}$ SR$^{-1}$, respectively. a) For [FeII]($25.988\\mu$m). b) for [FeIII](22.925$\\mu$m). } \\end{figure}  \\begin{figure} \\includegraphics[angle=0,width=3.5in]{H2S0c.eps} \\caption{ \\label{fig:H2S0} Continuum subtracted line emission images in the LH spectral cube, showing emission in the folded star forming disk. for the H$_2$S(0)($28.22\\mu$m)  line with lowest contour at and spacing at 0.01 $\\times 10^{-8}$erg cm$^{-2}$ s$^{-1}$ SR$^{-1}$. } \\end{figure}  \\clearpage  \\begin{figure} \\includegraphics[angle=0,width=3.5in]{OIVc.eps} \\includegraphics[angle=0,width=3.5in]{NeVc.eps} \\caption{ \\label{fig:OIVNEV} Continuum subtracted line emission images in the LH spectral cube, showing emission near the jet axis. Contours are evenly spaced.  The lowest contours and spacing for the [OIV]($25.890\\mu$m) and [NeV]($24.318\\mu$m) images are 0.5 and 0.1 $\\times 10^{-8}$erg cm$^{-2}$ s$^{-1}$ SR$^{-1}$, respectively. a) For [OIV]($25.890\\mu$m). b) For [NeV]($24.318\\mu$m). } \\end{figure}  \\begin{figure} \\includegraphics[angle=0,width=3.5in]{radioo4.eps} \\caption{ Radio emission at 5Ghz shown as grayscale with [OIV]($25.9\\mu$m) contours. The radio map is a 6\\arcsec resolution map of the inner radio lobes by \\citet{hardcastle06}. The [OIV]$25.9\\mu$ and [NeV]$24.3\\mu$m line emission are oriented approximately but not exactly along the jet axis. \\label{radioo4} } \\end{figure}  \\clearpage     \\begin{figure} \\includegraphics[angle=0,width=3.5in]{H2S3.eps} \\includegraphics[angle=0,width=3.5in]{H2S5.eps} \\caption{ \\label{fig:H2S3} Continuum subtracted line emission images from the SL spectral cube, showing the H$_2$S(3)(9.665$\\mu$m) and H$_2$S(5)(6.909$\\mu$m) lines with lowest contours at 1 $\\times 10^{-7}$erg cm$^{-2}$ s$^{-1}$ SR$^{-1}$. The higher S rotational lines from H$_2$ exhibit different morphology than the H$_2$S(0)28$\\mu$m line that was seen primarily in the folded star forming disk. Emission in the higher S lines is seen above the disk. Contours are evenly spaced. Lowest contour and spacing is  $\\times 10^{-7}$erg cm$^{-2}$ s$^{-1}$ SR$^{-1}$ for both lines. For a) H$_2$S(3)(9.665$\\mu$m). For b) H$_2$S(5)(6.909$\\mu$m). } \\end{figure}  \\begin{figure} \\includegraphics[angle=0,width=3.5in]{h2s3_o4.eps} \\includegraphics[angle=0,width=3.5in]{h2s3_pah11.eps} \\caption{ \\label{fig:h2s3_ov} a) Emission in H$_2$S(3)(9.665$\\mu$m) shown as gray scale overlayed with [OIV]($25.890\\mu$) contours. Contours are evenly spaced with lowest contours at are 0.5 $\\times 10^{-8}$erg cm$^{-2}$ s$^{-1}$ SR$^{-1}$ and a spacing 4 times this. The gray scale range for the H$_2$S(3) image is 0(white) to 7(black) $\\times 10^{-7}$erg cm$^{-2}$ s$^{-1}$ SR$^{-1}$. b) Same as a) except the H$_2$S(3) emission is overlayed with contours of the 11.3$\\mu$m  PAH dust emission feature (as shown in figure \\ref{fig:PAH11}). Contour spacing and lowest level is 20 MJy/SR. } \\end{figure}    \\begin{figure} \\includegraphics[angle=0,width=3.5in]{PAH7.eps} \\includegraphics[angle=0,width=3.5in]{PAH8.eps} \\caption{ \\label{fig:PAH} Flux at 7.7 and 8.6$\\mu$m showing dust emission features. Contours are evenly spaced.  No continuum subtraction has been done. The star forming disk is evident as the parallelogram shaped feature corresponding to the folded disk. The dust shell is seen above  and below the parallelogram feature. The minimum contour and spacing is approximately 10 MJy/SR in the peak of the line. a) The 7.7$\\mu$m dust emission feature. b) The 8.6$\\mu$m dust emission feature. } \\end{figure}   \\begin{figure} \\includegraphics[angle=0,width=3.5in]{NeII.eps} \\caption{ \\label{fig:NeII} Line emission in [NeII]($12.81\\mu$m). No continuum has been subtracted as the [NeII] line dominates the continuum by a factor greater than 2 everywhere. Contours are evenly spaced with lowest contour at 4 $\\times 10^{-7}$erg cm$^{-2}$ s$^{-1}$ SR$^{-1}$. } \\end{figure}  \\begin{figure} \\includegraphics[angle=0,width=3.5in]{PAH7_11.eps} \\includegraphics[angle=0,width=3.5in]{NeII_PAH11.eps} \\caption{ \\label{fig:ratio} a) The dust emission feature at 7.7$\\mu$m divided by that at 11.3$\\mu$m. The lowest and highest contours are shown at a ratios of 0.5 (in shell) and 1.0 (in parallelogram), with spacing of 0.1. Black is 1.0, white is 0.5. The 7.7 to 11.0$\\mu$m dust feature ratio varies with the dust shell having the lowest ratio. b) The [NeII](12.8$\\mu$m) line divided by the $11\\mu$m dust emission feature. The lowest and highest contours are shown at ratios of 1.9 (in shell) and 3.0 (in parallelogram), with spacing of 0.1. Black is 3.0, white is 1.9. The strength of the NeII](12.8$\\mu$m) compared to the 11.3$\\mu$m dust feature also varies with the dust shell having the lowest ratios. } \\end{figure}  \\clearpage  \\begin{figure} \\includegraphics[angle=0,width=3.5in]{sfrSL.eps} \\includegraphics[angle=0,width=3.5in]{jetSL.eps} \\caption{ \\label{fig:SLspec} a) Spectrum from the star forming parallelogram. D refers to a dust emission feature. b) Spectrum from the jet region. The pure rotational molecular hydrogen lines S(2)--S(7) are labeled as S2-S7. Molecular hydrogen lines are more prominent in the jet region than in the star forming parallelogram. In the star forming disk the [ArII](6.985$\\mu$m) line is brighter than the nearby pure rotational molecular hydrogen S(2) line at $6.909\\mu$m  whereas the opposite is true in the jet region. There is a change in the ratio of the dust emission features in these two spectra and in the ratio of the [NeII]/11.3 dust emission feature. } \\end{figure}  \\begin{figure} \\includegraphics[angle=0,width=3.5in]{ansi.eps} \\includegraphics[angle=0,width=3.5in]{ap1.eps} \\caption{ \\label{fig:LHspec} a) Spectra from the star forming warped disk (seen as a parallelogram in continuum) are shown at the 4 peaks in the parallelogram and at one additional point also in the parallelogram. Each spectrum is offset by $+$0.1 MJy/SR from the other. From bottom to top: The northern peak west of the nucleus, the southern peak west of the nucleus, the southern peak east of the nucleus, the northern peak east of the nucleus, a point between the southern and eastern peak and the nucleus. Positions are given in the text. In the star forming disk or parallelogram [NeV]($24.318\\mu$m) is not detected and [OIV]($25.890\\mu$) is weaker than [FeII]($25.988\\mu$m). b) Spectrum from a region near the jet axis south-west of nucleus. Near the jet axis [NeV]($24.318\\mu$m) is detected and [OIV]($25.890\\mu$) is 3-5 times brighter than [FeII]($25.988\\mu$m). The presence of [NeV] implies suggests that the radiation field is hard near the jet axis. } \\end{figure}   \\begin{figure} \\includegraphics[angle=0,width=3.5in]{h2plot.eps} \\caption{ \\label{fig:h2plot} Excitation diagram of the H$_2$ pure rotational lines. Points are based on those measured in the jet region with fluxes listed in Table \\ref{tab:hjet}. Solid squares refer to measurements and open triangles to upper limits. The solid line is for a temperature of $T_2=720K$ whereas the dotted one is $T_1 = 260K$. } \\end{figure}   \\newpage"
    },
    "0710/0710.5278_arXiv.txt": {
        "abstract": "{By re-processing the data of the second season of the OGLE survey for planetary transits and adding new mesurements on the same fields gathered in subsequent years with the OGLE telescope, we have identified 23 new transit candidates, recorded as OGLE-TR-178 to OGLE-TR-200. We studied the nature of these objects with the FLAMES/UVES multi-fiber spectrograph on the VLT. One of the candidates, OGLE-TR-182, was confirmed as a transiting gas giant planet on a 4-day orbit. We characterised it with further observations using the FORS1 camera and UVES spectrograph on the VLT. OGLE-TR-182b is a typical ``hot Jupiter'' with an orbital period of 3.98 days, a mass of 1.01 $\\pm 0.15$ M$_{Jup}$ and a radius of 1.13 $^{+0.24}_{-0.08}$ R$_{Jup}$.  Confirming this transiting planet required a large investment in telescope time with the best instruments available, and we comment on the difficulty of the confirmation process for transiting planets in the OGLE survey. We delienate the zone were confirmation is difficult or impossible, and discuss the implications for the Corot space mission in its quest for transiting telluric planets. ",
        "introduction": "Transiting extrasolar planets are essential to our understanding of planetary structure, formation and evolution outside the Solar System. The observation of transits and secondary eclipses gives access to such quantities as a planet's true mass, radius, density, surface temperature and atmospheric spectrum. The first transiting exoplanet was identified in 1999 around HD 209458. In the past three years, transiting exoplanet have been found in rapidly increasing number, both by radial velocity planet searches and by photometric surveys\\footnote{For an updated list see obswww.unige.ch/$\\sim$pont/TRANSITS.htm}. The OGLE search for transiting planets and low-mass stellar companions \\citep{ud02a} has been the first photometric transit survey to yield results. The first three seasons of photometric observations have revealed 137 transit candidates \\citep{ud02a,ud02c,ud02b,ud03}, among which 5 planets were found \\citep{ko03,bo04,po04a,kon04,bo05,ko05}, as well as two planet-sized low-mass stars \\citep{po05b,po06b}. Three further seasons of the OGLE transit survey have now been completed and await publication (Minniti et al., Udalski et al., in prep.).  The spectroscopic follow-up of most of the 137 first OGLE transit candidates, presented in \\citet{bo05} and \\citet{po05a}, has shown that the vast majority of the transit candidates were eclipsing binaries. A rate of one transiting planet for 10-20 eclipsing binaries is typical. A higher rate of planets can be found among candidates near the detection threshold. Two of the five planets from the OGLE survey, OGLE-TR-56 and OGLE-TR-132, were identified as candidates only after the application of a more sensitive transit detection algorithm \\citep{ko02}. However, lowering the detection threshold comes at the price of including some false positives of the detection procedure. The objective of the present study is to explore the regime near the detection threshold, the zone where the ratio of planets to eclipsing binaries will be more favourable than for deeper transit signals, but where the reality of the signal itself is not beyond doubt.  The exploration of this zone is relevant not only to identify new transiting planets in the OGLE survey, but also because other transit surveys will face similar issues, notably the {\\it CoRoT} and {\\it Kepler} space-based transit searches. ",
        "conclusions": "\\subsection{OGLE-TR-182b as a transiting planet}  The companion of OGLE-TR-182 has parameters typical of the planets detected by photometric transit surveys in all respects: it orbits a high-metallicity dwarf star, it has a mass comparable to that of Jupiter and a slightly larger size. Its period is close to an integer number of days, reflecting the strong selection bias due to the window function \\citep[see e.g.][about OGLE-TR-111$b$, another $P\\sim 4$ days transiting planet]{po04a}. At present, the constraint on the planetary radius is not sufficient to determine whether its radius corresponds to model expectations or whether is belongs to the set of anomalously large transiting hot Jupiters. Its position in the mass-period diagram is also similar to other known transiting planets, and reinforces the link between mass and period for close-in gas giants first pointed out by \\citet{ma05}.         \\subsection{The ``Twilight Zone'' of transit surveys}   \\begin{figure*}[ht] \\resizebox{16cm}{!}{\\includegraphics[angle=-90]{figMR.ps}} \\caption{''Observational'' mass-radius plot for known transiting exoplanets. The horizontal axis is the planet mass, scaled to $M^{2/3}P^{1/3}$, proportional to the radial-velocity semi-amplitude $K$. The vertical axis is the planet radius scaled to the host star radius, related to the depth of the photometric transit signal. \"TR-\" labels refer to OGLE candidates (closed symbols). Open symbols mark the position of other known transiting planets. The star symbol marks the best-fit location of another unsolved planet candidate from the third OGLE season. The gray bands show the near-threshold zones for the photometric detection (horizontal) and the spectroscopic confirmation (vertical) in the case of the OGLE survey. The dashed area is the ``twilight zone'' defined in the text, where confirmation is problematic.} \\label{mr} \\end{figure*}  The confirmation follow-up process for OGLE-TR-182 necessitated more than ten hours of FLAMES/VLT time for the radial velocity orbit, plus a comparable amount of FORS/VLT time for the transit lightcurve. In addition, several unsuccessful attempts were made to recover the transit timing in 2007 with the OGLE telescope, and 7 hours of UVES/VLT were devoted to measuring the spectroscopic parameters of the primary. This represents a very large amount of observational resources, and can be considered near the upper limit of what can be reasonably invested to identify a transiting planet.  Therefore, OGLE-TR-182 is a useful object to quantify the zone were neither the photometric signal nor the radial velocity signal are clear beyond doubt, the ``twilight zone'' of planetary transit candidate confirmation. In these cases, confirming the nature of the system is very difficult and time consuming. When the photometric signal is a possible false positive, a clear radial velocity orbit at the same period is an essential confirmation | as was for instance the case for OGLE-TR-132 \\citep{bo04}. On the other hand, when the radial velocity signal is marginal, a clear transit signal allows the phase and period of the orbit to be determined with confidence, reducing the radial-velocity orbit fit to a two-parameter problem ($V_0$, the systemic velocity,  and $K$, the orbital semi-amplitude) | as was the case for OGLE-TR-10 \\citep{bo05,ko05}.  However, when both the photometric and spectroscopic signals are marginal, many more observations are necessary until reasonable certainty can be achieved about the presence of a planetary companion. The uncertainties on the lightcurve make it difficult to phase the radial velocity data. The high radial velocity uncertainties hinder the identification of an orbital motion with the correct period, and the elimination of eclipsing binary blend scenarios  The OGLE survey is the first to explore this ``twilight zone'' in real conditions, since other ground-based surveys target brighter stars, for which very precise radial velocities can be obtained, so that the significance of the radial velocity signal can be established relatively easily \\citep[e.g.][]{ca07}. Based on the cases of OGLE-TR-10, OGLE-TR-132 and OGLE-TR-182,  and on the discussion in \\citet{po06}, we define the limits of the follow-up twilight zone as follows:  \\noindent -- photometric transit detection with $8< S_r < 12$, where $S_r$ is the transit significance in the presence of red noise \\citep[see definition in][]{po06}  \\noindent -- radial velocity orbital semi-amplitude 1-2 times the radial velocity uncertainties for 1-hour exposures with the facilities available: $\\sigma_{vr} < K < 2\\cdot \\sigma_{vr} $  These limits can be translated, for a circular orbit and a central transit, into limits on the radius and mass of the planet. Figure~\\ref{mr} shows the ``observational'' mass-radius diagram for the known transiting exoplanets. The horizontal axis is the planet mass divided by $M_*^{2/3}P^{1/3} $, to make it proportional to the observed radial velocity semi-amplitude (\"$K$\"). The vertical axis is the planet radius divided by the radius of the star, to make it proportional to the squareroot of the transit depth (at wavelengths where stellar limb darkening can be neglected). The units are such that a Jupiter-sized planet transiting a solar-sized star on a 4-day orbit will be placed at (1;1). Objects at similar positions in this plot will present similar challenges for confirmation. The \\mbox{ ``TR-''} labels refer to OGLE candidates, the other unmarked points are transiting planets from other surveys. The gray horizontal band is the zone where the transit detection is near the detection threshold, the vertical band is the zone where the radial velocity orbital signal is near the threshold. The intersection of the two,  delimited by the dashed lines, represents the ``twilight zone'' for the OGLE survey. We use a red noise level of $\\sigma_r = 3$ mmag  \\citep{po06}, a radial velocity uncertainty of 60-70 m/s (photon noise plus systematics), and assume that 5-10 transits are observed by the photometry. We find $ 0.85  < R_{pl}/R_* < 1.20$ and $ 0.45 < M_{pl}/{M_*}^{-2/3} P^{-1/3} < 1.05$ [$M_J {M_\\odot}^{-2/3} ({4\\ {\\rm days}})^{-1/3}$] for the zone boundaries. Candidates in this zone will be very difficult or impossible to confirm. On the left of the zone, radial velocity confirmation is out of reach, and below, the transit signal is below the photometric detection threshold.   The twilight zone for the OGLE transit survey encompasses the region in the mass-radius diagram corresponding to a normal hot Jupiter around a solar-type star. Hence the difficulty of the OGLE survey to detect transiting gas giants unless they are exceptionally heavy with a very short period like OGLE-TR-56, OGLE-TR-113 and OGLE-TR-132, or have an exceptionally high radius ratio like OGLE-TR-111.  The planetary transit system OGLE-TR-10 is also located within the zone, and indeed confirmation of its planetary nature required large investments in follow-up means both with the VLT \\citep{bo05} and Keck  \\citep{ko05} 8-10 meter telescopes.  As a further illustration of the extent of the zone, the plot shows the position of another candidate from the OGLE survey, yet unsolved despite extensive measurements with FLAMES and FORS. The radial velocity data is compatible with a planetary orbit, but many more observations would be required to confirm it securely. The best-fit planetary solution for this object is plotted on Fig.~\\ref{mr}.  In the wider context of transit searches in general, it is interesting to find where the ``twilight zone'' is located for different surveys, especially the space-based transit searches {\\it CoRoT} and {\\it Kepler}. For wide-field, small-camera surveys like HAT, WASP, XO and TrES, the twilight zone is not an important issue. Because the candidates are brighter, standard planet-search spectrographs can be used for the radial velocity follow-up, and the zone moves to the left of the mass-radius diagram, in a region were no planets are expected (low-mass, Jupiter-sized planets). In other words, if a planet is large enough to be detected by these surveys, it produces a radial velocity signal that is easily picked up by Doppler spectrographs.  Deeper surveys like SWEEPS \\citep{sa06} and planet searches in star clusters also have no twilight zone problem, for the opposite reason: their candidates are too faint to be confirmed in radial velocity for Jupiter-like planetary masses.  In the case of the {\\it CoRoT} space transit search, the zone will be located in a key position. In planet radius, it is expected to cover the 2-4 $R_\\oplus$ range \\citep[see][]{mo05} for the brightest targets. In planet mass, using the HARPS spectrograph, it will be in the 5-20 $M_\\oplus $ range, for short periods. This is a zone were planets are thought to be numerous, the domain of the ``hot Neptunes'' and ``super-Earths''. Indeed, the detection of this type of planets constitutes the main objective of the {\\it CoRoT} planetary transit search. From our experience with the OGLE follow-up, we therefore conclude that the {\\it CoRoT} mission will face similarly difficult cases in the confirmation process of transiting planets. The telescope time necessary for the follow-up of these candidates should be adequately evaluated. The OGLE follow-up process can provide some useful guidelines."
    },
    "0710/0710.1201_arXiv.txt": {
        "abstract": "We show that: (i) the long-term X-ray outburst light curve of the transient AXP XTE J1810--197 can be accounted for by a fallback disk that is evolving towards quiescence through a disk instability after having been heated by a soft gamma-ray burst, (ii) the spin-frequency evolution of this source in the same period can also be explained by the disk torque acting on the magnetosphere of the neutron star, (iii) most significantly, recently  observed pulsed-radio emission from this source coincides with the epoch of minimum X-ray luminosity. This is natural in terms of a fallback disk model, as the accretion power  becomes so low that it is not sufficient to suppress the beamed radio emission from XTE J1810--197. ",
        "introduction": "Anomalous X-ray pulsars (AXPs) and soft gamma-ray repeaters (SGRs) are systems of young neutron stars identified mainly by X-ray luminosities ($10^{34}-10^{36}$ erg s$^{-1}$) much higher than their rotational powers indicated by their observed spin period and period derivatives (Woods \\& Thompson 2005, Kaspi 2006). Brief soft gamma-ray bursts (\\la 1 s) with super-Eddington luminosities observed from SGRs and some  AXPs led to the identification of these sources as magnetars (Duncan \\& Thompson 1992). While the burst energetics are very likely to have magnetic origin, bursts themselves, being local events, do not require that the extended dipole component of the magnetic field has magnetar strength. Strong magnetic fields possibly in higher multipoles which decrease rapidly beyond the neutron star surface  is likely to be the source of these bursts.  Fallback disk models (Chatterjee et al. 2000; Alpar 2001) present a self-consistent explanation for the period clustering, soft X-ray, optical and infrared (IR) light curves of AXPs and SGRs in their persistent and enhanced phases (Ek\\c si \\& Alpar 2003; Ertan \\& Alpar 2003; Ertan \\& Cheng 2004; Ertan et al. 2006; Ertan \\& \\c{C}al{\\i}\\c{s}kan 2006, Ertan et al. 2007). These explanations entail active, viscously dissipating disks around these systems. Recent observations of AXP 4U 0142+61 in the mid-IR bands with {\\it Spitzer} Observatory (Wang et al. 2006) provided the first instance of a disk around an AXP.  Wang et al. (2006) attributed the optical and mid-IR data of this source to magnetospheric emission as proposed by magnetar models, and showed that the mid-IR data can be fitted with an irradiated passive disk. Detailed analysis of the overall data set including earlier detections in optical and near-IR bands strongly indicate the presence of an active disk around this source (Ertan et al. 2007). The difference between a passive and an active disk is that there is no viscous dissipation and therefore no mass inflow along a passive disk, while in an active disk, viscous dissipation works and leads to mass flow towards the inner disk and to the emission of X-rays through accretion onto the neutron star.  Both models expect a similar luminosity in the mid-IR bands which is emitted from the outer disk regions at which irradiation dominates dissipation. Nevertheless, estimates of these models differ in near IR and especially optical luminosities which, in the active disk model, are expected from the inner disk regions where the intrinsic dissipation becomes important. In a passive-disk model, the inner disk must be cut at a sufficiently large radius having a critical temperature below which temperatures are assumed to be not sufficient to sustain magnetorotational instability (MRI)(Balbus $\\&$ Hawley 1991), as proposed by Wang et al.(2006) for the disk around AXP 4U 0142+61.  On the other hand, the minimum temperature required for MRI which leads to viscous dissipation is uncertain. Recent work by Inutsuka $\\&$ Sano (2005) shows that MRI can be sustained in most regions of a protoplanetary disk even at temperatures as low as 300 K. Recently, Ertan et al. 2007 showed, by means of numerical fits to data, that the overall optical and IR dataset together with the X-ray luminosity of AXP 4U 0142+61 can be fitted with a single active, irradiated-disk model.  The upper limit for the dipole field strength estimated from these model fits is about $10^{13}$ G on the neutron star's surface (see Ertan et al. 2007 for details).  The X-ray enhancement of SGR 1900+14 following its giant flare can be explained by the evolution of a disk after it has been pushed back and piled up at larger radii by the burst (Ertan \\& Alpar 2003). This pushed-back disk model can also reproduce the long lasting, contemporaneous X-ray and IR enhancement light curve of AXP 1E 2259+586 for which the initial burst assumed by the model is likely to have remained below the detection limits (Ertan et al. 2006). In these models, reasonable fits to data are obtained by employing a single $\\a$ viscosity (Shakura \\& Sunyaev 1973) without invoking a disk instability. It might be  concluded that in the observed X-ray luminosity, $\\Lx$, range of these and other  persistent AXPs and SGRs ($10^{34} - 10^{36}$ erg s$^{-1})$, fallback disks do not experience global disk instabilities.   Are these persistent sources already in their quiescent states, or do they have a lower viscosity state which will finally take them into quiescence with a sharp decrease in $\\Lx$ after the mass-flow rate, $\\Mdot$, at the inner disk has decreased to below some critical level? In this paper, we suggest that the known  persistent AXPs and SGRs are not in their quiescent states, and that the transient AXP XTE J1810$-$197 is a good example for an AXP that is now, following an outburst, making a transition  back into quiescence by means of a viscous disk instability. The observations of pulsed radio emission from this source fall in the epoch following the last X-ray observation by Gotthelf and Halpern (2007) (Fig 1). Our model X-ray light curve predicts that in this epoch the $\\Lx$ remains even below the pre-outburst quiescent level. We propose that the pulsed-radio emission is a natural outcome of accretion rate decreasing below a critical value such that it is no longer sufficient to hinder pulsed-radio emission. In this picture,  all the other persistent AXPs and SGRs which are expected to be in the hot viscosity state at present would finally evolve into quiescence as their luminosity and thus the irradiation strength decrease below a critical level. The persistent and quiescent states would then correspond to two different viscosity states, and the AXPs and  SGRs could make transitions between these cold and hot states if the system undergoes a global instability triggered, for instance, by an energetic soft gamma-ray burst. Furthermore, persistent AXPs and SGRs are not expected  to show pulsed radio emission, because their accretion rates are much higher than the critical values below which pulsed-radio emission is allowed.  Three of the AXPs (of which one is a candidate AXP) and one SGR  were observed to show transient behavior in their X-ray luminosity  (Kaspi 2006, Mereghetti et al. 2006, Israel  et al. 2007, Gotthelf and Vasisht 1998, Torii et al. 1998). Of these transient sources, XTE J1810$-$197 is a confirmed AXP with measured spin period and period derivative (Ibrahim et al. 2004). This source was discovered during the early decay phase of its X-ray outburst (Ibrahim et al. 2004). The $\\Lx$  of XTE J1810$-$197 has still been decaying three years after its first detection, and  now is about two orders of magnitude less than  the observed maximum $\\Lx$. This source was not observed around the onset of the X-ray outburst. Earlier archival data show that the $L_x$ prior to the outburst is also about two orders of magnitude fainter than the observed maximum luminosity (Gotthelf \\& Halpern 2007). During the first year of the decay phase of XTE J1810$-$197, $\\Lx$ remained in the range characteristic of persistent AXPs. In this phase, the inner disk seems to be in a hot viscosity state, with temperatures above the critical level, likely to be the same as that operating in the persistent AXPs. In the present work, we show that for this first year, the  X-ray light curve of XTE J1810--197 can be simulated by a pure viscous evolution of a disk pushed back by a burst, but subsequently as $\\Lx$ decreases below a critical level around $10^{34}$ erg s$^{-1}$ the light curve deviates from the characteristic viscous relaxation model curve.  This transition can be explained by a viscous disk instability leading the system into quiescence.  The disk instability model for hydrogen disks is summarized and the possibilities for the instabilities of the fallback disks are discussed  in \\S\\ 2. In \\S\\ 3, we show that the overall decay curve of the transient AXP XTE J1810--197 can be reproduced by an irradiated active disk that undergoes a transition into quiescence below some critical X-ray luminosity that is predictable from the model fits. In \\S\\ 4 and \\S\\ 5, we discuss the X-ray light curve and its connection to pulsed radio emission respectively. In \\S\\ 6, we investigate the spin-frequency evolution of XTE J1810--197 during the post-burst fading of the X-ray flux. Our results are summarized in \\S\\ 7. ",
        "conclusions": "We have firstly shown that the X-ray outburst light curve of the transient AXP XTE J1810$-$197 lasting for about three years can be reproduced by the evolution of an irradiated disk that undergoes a transition into quiescence by means of a viscous disk instability as the X-ray luminosity decreases below a critical limit less than around $10^{34}$ erg s$^{-1}$ which is near the minimum of the observed luminosity range of AXPs and SGRs. We have assumed that the inner disk is initially pushed back by a missed soft gamma-ray burst that took a large portion of the inner disk into a hot viscosity state which is likely the state prevailing in the disks of persistent AXPs and SGRs. The mass-flow rate, $\\Mdot$,  provided by the extended disk of XTE J1810$-$197 with surface densities likely to be lower than those of persistent sources is not sufficient for the inner disk to be kept in the hot state for longer periods by fuelling the X-ray irradiation through accretion. After the X-ray luminosity has decreased below a critical level, the system evolves into quiescence, by means of a viscous disk instability. Observations indicate that the source was also in the quiescent state prior to the X-ray outburst. Unlike XTE J1810$-$197, persistent sources AXP 1E 2259+586 and SGR 1900+14 were already in the hot state with $\\Lx ~\\ga ~10^{34}$ erg s$^{-1}$ before their enhancements and their X-ray luminosities remained above the critical level during the outburst decay as well. The X-ray luminosity  of these and other persistent AXPs and SGRs are also expected to decrease slowly as the disk surface densities decrease gradually through expansion and accretion in the persistent phase. They will also finally evolve towards the quiescent regime by means of a disk instability as their accretion rate falls below critical values. Quiescent sources, on the other hand, could make an upward transition to the hot state when a viscous instability is triggered by an energetic gamma-ray burst.  Secondly, we have shown that the magnetospheric torque applied by the disk on the star is in agreement with the spin-frequency evolution of this transient AXP (see Fig. 2). Along the X-ray outburst decay, XTE J1810$-$197 is in the accretion regime with spin-down. In this phase, the system is slowly receding from the rotational equilibrium between the disk and the magnetosphere, while most of the mass flowing towards the disk is being accreted onto the neutron star. Our torque model is almost independent of the mass-inflow rate $\\Mdot$ to the inner disk in this period.  We have obtained the model curve for spin evolution directly from the evolution of the mass-inflow rate corresponding to the model X-ray light curve presented in Fig. 1 and by adjusting the field strength accordingly. The model curve given in Fig.2 is obtained with a dipolar magnetic field of $B_{\\ast }\\simeq 2\\times 10^{12}$ G. Below some critical value of $\\Mdot$, most of the inflowing matter is propelled out of the system, and the functional form of the torque is expected to change. In this regime the torque would probably be less efficient than the model torque we employ here. The dependence of the torque on $\\Mdot$ also changes as the system approaches rotational equilibrium (see \\S\\ 6), and the spin-up torque due to accretion is no longer negligible in comparison with the spin-down torque.   Finally, we have proposed that the pulsed-radio emission from XTE J1810--197 could be accounted for by the decreasing accretion rate below a critical level which switches on the beamed radio emission from the polar caps. Our model predicts a dip-like structure at the end of the decay phase such that the X-ray luminosity remains even below the pre-outburst quiescent level for about one year. This period matches the epoch of observations of pulsed-radio emission from this transient AXP (Camilo et al. 2007). This interpretation of the pulsed radio emission together with the model predictions for the X-ray light curve could be tested by future radio and X-ray  observations of XTE J1810--197 and other transient AXPs and SGRs. Transient AXPs can show up as radio pulsars when their disks make a transition to the cold state and the X-ray luminosity drops sufficiently such that their accretion power can no longer suppress the beamed radio emission.  A detailed analysis on the torque evolution and the beamed radio emission from AXPs and SGRs will be the subject of our future work."
    },
    "0710/0710.3518_arXiv.txt": {
        "abstract": "{The results from H.E.S.S. observations towards Westerlund~2 are presented. The detection of very-high-energy gamma-ray emission towards the young stellar cluster Westerlund 2 in the HII complex RCW49 by H.E.S.S. provides ample evidence that particle acceleration to extreme energies is associated with this region. A variety of possible emission scenarios is mentioned, ranging from high-energy gamma-ray production in the colliding wind zone of the massive Wolf-Rayet binary WR~20a, collective wind scenarios, diffusive shock acceleration at the boundaries of wind-blown bubbles in the stellar cluster, and outbreak phenomena from hot stellar winds into the interstellar medium. These scenarios are briefly compared to the characteristics of the associated new VHE gamma-ray source HESS~J1023--575, and conclusions on the validity of the respective emission scenarios for high-energy gamma-ray production in the Westerlund~2 system are drawn.}     \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.3803_arXiv.txt": {
        "abstract": "We present interferometric observations of $^{12}$CO and $^{13}$CO $J$=2$-$1 emission from the butterfly-shaped, young planetary nebula NGC 6302. The high angular resolution and high sensitivity achieved in our observations allow us to resolve the nebula into two distinct kinematic components: (1) a massive expanding torus seen almost edge-on and oriented in the North-South direction, roughly perpendicular to the optical nebula axis. The torus exhibits very complex and fragmentated structure; (2) high velocity molecular knots moving at high velocity, higher than 20 \\kms, and located in the optical bipolar lobes. These knots show a linear position-velocity gradient (Hubble-like flow), which is characteristic of fast molecular outflow in young planetary nebulae. From the low but variable $^{12}$CO/$^{13}$CO $J$=2$-$1 line intensity ratio we conclude that the $^{12}$CO $J$=2$-$1 emission is optically thick over much of the nebula. Using the optically thinner line $^{13}$CO $J$=2$-$1 we estimate a total molecular gas mass of $\\sim$ 0.1 M$_\\odot$, comparable to the ionized gas mass; the total gas mass of the NGC 6302 nebula, including the massive ionized gas from photon dominated region, is found to be $\\sim$ 0.5 M$_\\odot$. From radiative transfer modelling we infer that the torus is seen at inclination angle of 75$^\\circ$ with respect to the plane of the sky and expanding at velocity of 15 \\kms. Comparison with recent observations of molecular gas in NGC 6302 is also discussed. ",
        "introduction": "Low and intermediate-mass stars (1 M$_\\odot$ $\\le$ M$_*$ $\\le$ 8 M$_\\odot$) evolve through the Asymptotic Giant Branch (AGB) phase, which is characterized by copious mass loss in the form of dusty slow wind, before emerging as planetary nebulae (PNe). The mass-loss process is commonly assumed to be isotropic, resulting in the formation of a spherically symmetric envelope around the central AGB star. However, a significant fraction of the circumstellar envelope around post-AGB stars and planetary nebulae possess bipolar and even multipolar morphology. The mechanisms responsible for such departure from spherical symmetry are still unknown (e.g.\\ Balick \\& Frank 2002). The interaction between collimated fast outflows and the surrounding envelope or the influence of a binary companion are often cited as possible shaping mechanisms. Large expanding or rotating disks/tori has been frequently inferred to be present in bipolar nebulae such as the Egg Nebula (Sahai et al. 1998) and the Red Rectangle (Bujarrabal et al. 2005). These disks or tori might confine and channel the wind from the central star into bipolar directions. If the disks/tori are dense enough, the strong interaction (i.e oblique shock) between the wind and the disks/tori could focus and collimate the outflow (Frank et al. 1996). Thus, more detailed studies of the structure and kinematics of such disk/torus could provide better understanding on the formation of the nebulae.  NGC\\,6302 is a young planetary nebula and belongs to the class of the highest excitation PNe (Pottasch et al.\\ 1996).  From the presence of numerous emission lines from highly ionized species, Pottasch et al.\\ (1996) concluded that the central star is a white dwarf or approaching a white dwarf and its excitation temperature is very high, $\\sim$380,000 K. In optical images, NGC\\,6302 appears as a butterfly shaped nebula with two huge bipolar lobes separated by a dark equatorial lane (Matsuura et al. 2005), which is presumably the location of a massive disk or torus.  An expanding H\\2\\ region is detected at the center of the nebula in the radio continuum (G\\'{o}mez et al. 1989).  The presence of a massive molecular envelope is known through the detection of strong CO rotational lines in NGC\\,6302 (Huggins et al.\\ 1996, Hasegawa \\& Kwok 2003). The CO lines have very peculiar and complex shapes, with a double-peaked profile and high velocity wings, extending up to $\\sim$40 \\kms\\ from the systemic velocity. Such line shape suggests that the molecular envelope is likely non-spherical and could contain a fast molecular outflow, similar to fast bipolar outflows observed in some proto-planetary nebulae and young planetary nebulae.  Detailed observations with the VLT and HST of Matsuura et al.\\ (2005) show the presence of a large warped disk oriented perpendicular to the bipolar lobes seen prominently in the optical images. The disk is massive and has a large extinction. Strong emission from crystalline silicates together with PAH bands have also been detected in NGC\\, 6302 (Kemper et al.\\ 2002). We note that, interestingly, OH maser emission, which is usually associated with oxygen rich circumstellar envelopes, has also been detected in NGC\\,6302 (Payne et al. 1988). Thus, NGC\\,6302 is chemically peculiar, containing a mixture of carbon-rich and oxygen-rich material.  The distance to NGC\\,6302 is uncertain, with estimates ranging from 0.15 to 2.4 kpc. From the expansion proper motion of the central H\\2\\ region, Gom\\'{e}z et al. (1989) estimates a distance of 2.2 $\\pm$ 1.1 kpc. However, Matsuura et al. (2005) note that a distance much larger than 1 kpc would lead to very high luminosity and shell mass, which would exceed the highest luminosity of post-AGB stars and PNe predicted by stellar evolution models. with even the highest core mass.  More recently, Meaburn et al. (2005) detect directly the proper motion of the prominent optical lobes in NGC\\,6302 nebula and infer a distance of 1 kpc. We will adopt a distance of 1 kpc for NGC\\, 6302, similar to Matsuura et al. (2005) and Kemper et al. (2002).  In this paper we present high angular resolution observations of the $J$=2$-$1 line of $^{12}$CO and its isotope $^{13}$CO, in order to study the spatial distribution and kinematics of the molecular gas in the envelope of NGC\\,6302, especially the massive equatorial disk and the gas moving at higher velocities.  After the submission of our paper for publication, we learned of a recent paper by Peretto et al. (2007), which also presents maps of molecular gas in NGC\\,6302. Although the spatial distribution of CO emission is similar in both papers, our higher sensitivity (by a factor of 2) observations allow us to produce maps for more velocity channels and better understand the spatial kinematics of the circumstellar envelope. We are able, in particular, to image and study the fast molecular outflows, which are barely detected in lower sensivity data of Peretto et al. (2007). Where appropriate, we will compare our results with those obtained by Peretto et al. (2007) ",
        "conclusions": "We have imaged at high sensitivity and high angular resolution the envelope around the young planetary nebula NGC\\,6302 in $^{12}$CO $J$=2--1 and $^{13}$CO $J$=2--1 lines. Continuum at 1.3mm wavelegth is also imaged and seems to come from the inner H\\2\\ region.  We find that the very complex molecule-rich nebula is well resolved and can be separated into two components: a massive low-velocity torus seen nearly edge-on and high-velocity knots. The image of the dense torus is very accurately coincident with the dark lane that separates the conspicuous two lobes in the optical image of NGC\\,6302. The fast knots are located within the optical lobes and show a linear velocity gradient, which is characteristic of fast molecular gas in young planetary nebulae.  We find that the $^{12}$CO/$^{13}$CO $J$=2--1 brigthness ratio is low and varies between 2 and 5, decreasing with increasing line intensity. We conclude that the $^{12}$CO $J$=2--1 emission is optically thick over much of the nebula, but that $^{13}$CO $J$=2--1 is optically thin in most velocities and lines of sight. Using our observations of this line, we estimate masses for the different components, yielding a total molecular gas mass of $\\sim$ 0.1 M$_\\odot$. We discuss the value of the total mass of the NGC\\,6302 nebula, including the ionized gas, whose mass is comparable to that of the molecular gas, and the massive PDR, which is probably the dominant component. We find a total mass of $\\sim$ 0.5 M$_\\odot$; but we recall that the very uncertain contribution from very outer layers could be significant.  Using a radiative transfer model we infer that the torus is seen at an inclination angle of 75$^\\circ$ with respect to the plane of the sky and expanding at a velocity of 15 \\kms. The mass loss rate is found to be very high, $\\sim$ 1.5$\\times$10$^{-4}$ \\ms\\ yr$^{-1}$, resulting in a dense (average gas density $\\sim$ 2$\\times$10$^4$ cm$^{-3}$) and a massive torus.   \\bigskip We are grateful to SMA staff for carrying out the observations. We thank an anonymous referee for constructive criticisms that helped to improve our paper significantly. V.B.\\ acknowledges support from the \\emph{Spanish Ministry of Education \\& Science}, project numbers AYA2003-7584 and ESP2003-04957. Help from Sebastien Muller is gratefully acknowledged. This research has made use of NASA's Astrophysics Data System Bibliographic Services and the SIMBAD database, operated at CDS, Strasbourg, France."
    },
    "0710/0710.4081_arXiv.txt": {
        "abstract": "Accurate alignment of the radio and optical celestial reference frames requires detailed understanding of physical factors that may cause offsets between the positions of the same object measured in different spectral bands. Opacity in compact extragalactic jets (due to synchrotron self-absorption and external free-free absorption) is one of the key physical phenomena producing such an offset, and this effect is well-known in radio astronomy (``core shift''). We have measured the core shifts in a sample of 29 bright compact extragalactic radio sources observed using very long baseline interferometry (VLBI) at 2.3 and 8.6 GHz. We report the results of these measurements and estimate that the average shift between radio and optical positions of distant quasars would be of the order of 0.1--0.2 mas. This shift exceeds positional accuracy of GAIA and SIM. We suggest two possible approaches to carefully investigate and correct for this effect in order to align accurately the radio and optical positions.  Both approaches involve determining a Primary Reference Sample of objects to be used for tying the radio and optical reference frames together. ",
        "introduction": "}   Extragalactic relativistic jets are formed in the immediate vicinity of the central black holes in galaxies, at distances of the order of 100 gravitational radii, and they become visible in the radio at distances of about 1000 gravitational radii \\citep{LZ2007}. This apparent origin of the radio jets is commonly called the ``core''. In radio images of extragalactic jets, the core is located in the region with an optical depth $\\tau_s\\approx 1$. This causes the absolute position of the core, $r_\\mathrm{core}$, to vary with the observing frequency, $\\nu$, since the optical depth profile along the jet depends on $\\nu$: $r_\\mathrm{core} \\propto \\nu^{-1/k_\\mathrm{r}}$ \\citep{BlandfordKonigl79}. Variations of the optical depth along the jet can result from synchrotron self-absorption \\citep{Koenigl81}, pressure and density gradients in the jet and free-free absorption in the ambient medium most likely associated with the broad-line region (BLR) \\citep{L98}.  The core shift is expected to introduce systematic offsets between the radio and optical positions of reference sources, affecting strongly the accuracy of the radio-optical matching of the astrometric catalogues. The magnitude of the core shift can exceed the inflated errors of the radio and optical positional measurements by a large factor. This makes it necessary to perform systematic studies of the core shift in the astrometric samples in order to understand and remove the contribution of the core shift to the errors of the radio-optical position alignment.  Measurements of the core shift have been done so far only in a small number of objects \\citep[e.g.,][]{MES94,Lara_etal94,PR97,L96,L98,RL2001,Kadler_etal04,SokolovCawthorne2007}. In this paper, we present results for 29 compact extragalactic radio sources used in VLBI astrometric studies and discuss the core shift effect on the task of the radio-optical reference frame alignment. ",
        "conclusions": "}  Measurements of the frequency-dependent shift of the parsec-scale jet cores in AGN are reported for 29 bright extragalactic radio sources. It is shown that the shift can be as high as 1.4~mas between 2.3 and 8.6~GHz. We have shown that core shifts are likely to pose problems for connecting radio and optical reference frames. We have estimated from theory an average shift between the radio (4~cm) and optical (6000~\\AA) bands to be of an order of 0.1~mas for a complete sample of radio selected AGN.  The estimated radio-optical core shift exceeds the positional accuracy of GAIA and SIM. It implies that the core shift effect should be carefully investigated, and corrected for, in order to align accurately the radio and optical positions. Based on our investigation, we suggest two possible approaches, both involving determining a Primary Reference Sample of objects to be used for tying the radio and optical reference frames together. 1) In the first approach, multi-frequency VLBI measurements can be used for calculating the projected optical positions, assuming that the radio and optical emission regions are both dominated by a spatially compact component marginally resolved with VLBI and SIM and point-like for GAIA. The discrepancies between the measured optical and radio positions can then be corrected for the predicted shifts, and the subsequent alignment of the radio and optical reference frames can be done using standard procedures. 2) A more conservative approach may also be applied, by employing the VLBI observations to identify and including in the Primary Reference Sample only those quasars in which no significant core shift has been detected in multi-epoch experiments. Either of the two approaches should lead to substantial improvements of the accuracy of the radio-optical position alignment."
    },
    "0710/0710.3160_arXiv.txt": {
        "abstract": "The first stars form in dark matter halos of masses $\\sim$$10^6 \\Ms$ as suggested by an increasing number of numerical simulations. Radiation feedback from these stars expels most of the gas from their shallow potential well of their surrounding dark matter halos. We use cosmological adaptive mesh refinement simulations that include self-consistent Population III star formation and feedback to examine the properties of assembling early dwarf galaxies. Accurate radiative transport is modelled with adaptive ray tracing. We include supernova explosions and follow the metal enrichment of the intergalactic medium.  The calculations focus on the formation of several dwarf galaxies and their progenitors.  In these halos, baryon fractions in 10$^8$ \\Ms~halos decrease by a factor of 2 with stellar feedback and by a factor of 3 with supernova explosions.  We find that radiation feedback and supernova explosions increase gaseous spin parameters up to a factor of 4 and vary with time. Stellar feedback, supernova explosions, and \\hh~cooling create a complex, multi-phase interstellar medium whose densities and temperatures can span up to 6 orders of magnitude at a given radius. The pair-instability supernovae of Population III stars alone enrich the halos with virial temperatures of 10$^4$ K to approximately 10$^{-3}$ of solar metallicity.  We find that 40\\% of the heavy elements resides in the intergalactic medium (IGM) at the end of our calculations.  The highest metallicity gas exists in supernova remnants and very dilute regions of the IGM. ",
        "introduction": "The majority of galaxies in the universe are low-luminosity, have masses of $\\sim$$10^8$ solar masses, and are known as dwarf galaxies \\citep{Schechter76, Ellis97, Mateo98}.  Galaxies form hierarchically through numerous mergers of smaller counterparts \\citep{Peebles68, White78}, whose properties will inevitably influence their parent galaxy.  Dwarf galaxies are the smallest galactic building blocks, and this leads to the question on even smaller scales: how were dwarf galaxies influenced by their progenitors?  A subset of dwarf galaxies, dwarf spheroidals (dSph), have the highest mass-to-light ratios \\citep{deBlok97, Mateo98} and contain a population of metal-poor stars that are similar to Galactic halo stars \\citep{Tolstoy04, Helmi06}.  There is a metallicity floor exists of $10^{-3}$ and $10^{-4}$ of solar metallicity in dSph and halo stars, respectively \\citep{Beers05, Helmi06}.  Stellar metallicities increase with time as previous stars continually enrich the interstellar medium (ISM).  Hence the lowest metallicity stars are some of the oldest stars in the system and can shed light on the initial formation of dwarf galaxies.  This metallicity floor also suggests that metal enrichment was widespread in dark matter halos before low-mass stars could have formed \\citep[e.g.][]{Ricotti02b}.  Supernovae (SNe) from metal-free (Pop III) stars generate the first metals in the universe and may supply the necessary metallicity to form the most metal-poor stars observed \\citep{Ferrara98, Madau01a, Norman04}.  Dwarf galaxy formation can be further constrained with observations that probe reionization and semi-analytic models.  Observations of luminous quasars powered by supermassive black holes (SMBH) of mass $\\sim$$10^9 \\Ms$ \\citep{Becker01, Fan02, Fan06} and low-luminosity galaxies \\citep{Hu02, Iye06, Kashikawa06, Bouwens06, Stark07} at and above redshift 6 indicate that active star and BH formation began long before this epoch.  Semi-analytic models have argued that cosmological reionization was largely caused by low-luminosity dwarf galaxies \\citep{Haiman97, Cen03, Somerville03, Wise05, Haiman06}.  Some of the most relevant parameters in these models control star formation rates, ionizing photon escape fractions, metal enrichment, and the minimum mass of a star forming halos.  They are usually constrained using (i) the cosmic microwave background (CMB) polarization observation from Wilkinson Microwave Anisotropy Probe (WMAP) that measures the optical depth of electron scattering to the CMB \\citep{Page06}, (ii) Gunn-Peterson troughs in $z \\sim 6$ quasars, and (iii) numerical simulations that examine negative and positive feedback of radiation backgrounds \\citep{Machacek01, Machacek03, Yoshida03, Mesigner06}. Radiation hydrodynamical \\textit{ab initio} simulations of the first stars \\citep{Yoshida06a, Abel07} and galaxies can further constrain the parameters used in semi-analytic models by analyzing the impact of stellar feedback on star formation rates and the propagation of \\ion{H}{2} regions in the early universe.  Moreover, these simulations contain a wealth of information pertaining to the properties of Pop III star forming halos and early dwarf galaxies that can increase our understanding of the first stages of galaxy formation.  First we need to consider Pop III stars, which form in the progenitor halos of the first galaxies, to capture the initial properties of dwarf galaxies.  Cosmological numerical studies have shown that massive (30--300 \\Ms) Pop III stars form in dark matter halos with masses $\\sim$$10^6 \\Ms$ \\citep{Abel02, Bromm02, Yoshida06b, Gao07, OShea07}.  Recently, \\citeauthor{Yoshida06b} and \\citet{Turk08} followed the gaseous collapse of a molecular cloud that will host a Pop III star to cosmologically high number densities of $10^{16}$ and $10^{21}$ \\cubecm, respectively.  The former group thoroughly analyzed the gas dynamics, cooling, and stability of this free-fall collapse. The latter group observed a protostellar core forming with 10 Jupiter masses that is bounded by a highly asymmetric protostellar shock. Both groups found no fragmentation in the fully molecular core that collapses into a single, massive $\\sim$100\\Ms~star.  Furthermore, \\citet{Omukai03} determined that accretion may halt at the same mass scale, using protostellar models even for different mass accretion histories.  Pop III stars with stellar masses roughly between 140 and 260 \\Ms~end their life in a pair-instability SN (PISN) that releases $10^{51} - 10^{53}$ ergs of energy and tens of solar masses of heavy elements into the ambient medium \\citep{Barkat67, Bond84, Heger02}.  These explosions are an order of magnitude larger than typical Type II SNe in both quantities \\citep{Woosley86}, such explosions energies are larger than the binding energies in their low-mass hosts, e.g., $2.8 \\times 10^{50}$ ergs for a $10^6 \\Ms$ halo at redshift 20.  Thus gas structures in the host halo are totally disrupted and expelled, effectively enriching the surrounding intergalactic medium (IGM) with the SN ejecta \\citep{Bromm03, Kitayama05, Greif07}.  The combination of the shallow potential well and large explosion energy suggests that these events are good candidates for enriching the first galaxies and IGM.  Outside of the pair-instability mass range, Pop III stars die by directly collapsing into a BH \\citep{Heger03}, possibly providing the seeds of high-redshift quasars in galaxies that are associated with the rarest density fluctuations \\citep[e.g.][]{Madau01b, Volonteri05, Trenti07}.  One-dimensional calculations \\citep{Whalen04, Kitayama04, Kitayama05} and recent three-dimensional radiation hydrodynamical simulations \\citep{Yoshida06a, Abel07} have investigated how the Pop III stellar feedback affects its host halo and nearby cosmic structure.  In addition to SNe, \\ion{H}{2} regions surrounding Pop III stars, which have luminosities $\\sim$$10^6 L_\\odot$ \\citep{Schaerer02}, alone can dynamically affect gas at distances up to a few proper kpc. Ionization fronts and \\ion{H}{2} regions \\citep[see][for a review]{Yorke86} have been extensively studied in literature on star formation since \\citet{Stroemgren39}.  Stellar radiation generates an ionization front that begins as a R-type front and transforms into a D-type front when its speed slows to twice the sound speed of the ionized gas.  Then a strong shock wave forms at the front and recedes from the star at $\\sim$$30\\kms$.  The ionization front decouples from the shock wave and creates a final \\ion{H}{2} region that is 1 -- 3 proper kpc in radius for massive Pop III stars residing in low-mass halos.  The ionized gas is warm ($\\sim$$3 \\times 10^4$ K) and diffuse ($\\sim$$1 \\cubecm$).  The shock wave continues to accumulate gas and advance after the star dies.  Eventually it stalls in the IGM, but in the process, it reduces the baryon fraction of the halo below one percent \\citep{Yoshida06a, Abel07}.  Clearly the number of progenitors of a given galaxy as well as the star formation and feedback history of the progenitors will play a role in shaping all of its properties.  But how much?  If most stars of a galaxy are formed later, will the earliest episodes not be entirely negligible?  To start addressing these questions, we have carried out a suite of simulations that include accurate three dimensional radiative transfer and the SN explosions of Pop III stars and have followed the buildup of several dwarf galaxies from those Pop III star hosting progenitors. The Pop III radiative and SN feedback dramatically alter the properties of high redshift dwarf galaxies, and we discuss some of the most striking differences here.  We leave a more detailed exposition of star formation rates, star forming environments, and the beginning of cosmic reionization for a later paper.   In the following section, we detail our cosmological, radiation hydrodynamics simulations and the star formation algorithm.  Then we describe the global characteristics of dwarf galaxies that forms in our simulations in \\S\\ref{sec:reionResults}.  There we also focus on metal enrichment of star forming halos and the IGM, arising from PISNe.  In \\S\\ref{sec:reionDiscuss}, we discuss the implications of our findings on the paradigm of high-redshift galaxy formation by including \\hh~chemistry and Pop III star formation and feedback.  We summarize in the last section. ",
        "conclusions": "\\label{sec:reionDiscuss}  We find that the combination of Pop III stellar feedback and continued \\hh~cooling in \\tvir~$<$ 10$^4$ K halos alters the landscape of high-redshift galaxy formation.  The most drastic changes are as follows:   1. \\textit{Dynamic assembly of dwarf galaxies}--- A striking difference when we include Pop III radiative feedback are the outflows and gas inhomogeneities in the halos and surrounding IGM.  The outflows enrich the IGM and reduce the baryon fraction of the $10^4$ K halo as low as 0.05, much lower than the cosmic fraction \\Ob/\\Om~= 0.17 \\citep[cf.][]{Yoshida06a, Abel07}.  This substantially differs from the current theories of galaxy formation where relaxed isothermal gas halos hierarchically assemble a dwarf galaxy.  For instance in simulation A, there are remarkable filamentary structures and a clumpy ISM.  Furthermore, Pop III feedback increases the total baryonic angular momentum of the system up to a factor of 3 without SNe and up to 5 with SNe.   2. \\textit{Pop III sphere of influence}--- Pop III feedback is mainly a local phenomenon except its contribution to the UVB.  How far its \\ion{H}{2} region, outflows, and metal ejecta (if any) will predominately determine the characteristics of the next generation of stars.  Highly biased (clustered) regions are significantly affected by Pop III feedback.  The first galaxies will form in these biased regions and thus should be significantly influenced by its progenitors.   3. \\textit{Dependence on star forming progenitors}--- Although our calculations with SNe only provided an upper limit of metal enrichment, it is clear that the metallicity, therefore metal-line and dust cooling and metal-enriched (Pop II) star formation, depends on the nature of the progenitors of the dwarf galaxy.  If the galaxy was assembled by smaller halos that hosted a Pop III star that did not produce a SN, e.g., the galaxy would continue to have a top-heavy initial mass function (IMF).   4. \\textit{Complex protogalactic ISM}--- The interplay between stellar and SNe feedback, cold inflows, and molecular cooling produce a truly multi-phase ISM that is reminiscent of local galaxies.  The cool, warm, and hot phases are interspersed throughout the dwarf galaxy, whose temperatures and densities can span up to 6 orders of magnitude at a given radius.   5. \\textit{Metallicity floor}--- When the halo is massive enough to host multiple sites of star formation, the metal ejecta does not significantly increase the mean metallicity of the host halo.  There seems to be a balance between (a) galactic outflows produced from SNe, (b) inflowing metal-enriched gas, (c) inflowing pristine gas, and (d) SNe ejecta that is not blown out of the system.  In our high yield models, the metallicity interestingly fluctuates around $10^{-3} Z_\\odot$ in the most massive halo when this balance occurs at and above mass scales $\\sim$10$^7$\\Ms.   Clearly the first and smallest galaxies are complex entities, contrary to their low mass and generally assumed simplicity.  Our calculations reflect the important role of Pop III stellar feedback in early galaxy formation.  These high-redshift galaxies have a $\\sim$5--15\\% chance of being undisturbed by mergers until the present day, being ``fossils'' of reionization \\citep{Gnedin06}.  Dwarf spheroidals (dSph) galaxies are some of the darkest galaxies in the universe, having high mass-to-light ratios up to 100 \\citep{Mateo98}.  Gas loss in dSph's close to the Milky Way or M31 can be explained by gas tidal stripping during orbital encounters \\citep{Mayer07}.  However there are some galaxies (e.g. Tucana, Cetus) removed from both the Milky Way and Andromeda galaxies and cannot be explained by tidal stripping.  In addition to ultraviolet heating from reionization \\citep{Bullock00, Susa04} and intrinsic star formation \\citep{MacLow99}, perhaps stellar feedback from Pop III stars influenced the gas-poor nature of dSph's.  Even at the onset of widespread star formation in the objects studied here, the baryon fraction can be three times lower than the cosmic mean, and the dwarf galaxy may never fully recover from the early mass loss.  This initial deficit may play an important role in future star formation within these low-mass galaxies and could help explain the large mass to light ratio in isolated dSph's.  With the radiative and chemical feedback from the progenitors of the early dwarf galaxies, we have an adequate set of cosmological ``initial conditions'' to study the transition from Pop III to Pop II stars.  In this setup, the current metal tracer field can be used to include metal line cooling.  Dust cooling may induce fragmentation of solar mass fragments at metallicities as low as $\\sim$10$^{-6}$ at high densities \\citep{Schneider06}.  However, metal-line cooling might not be important at these low metallicities in diffuse gas. \\citet{Jappsen07a} showed that metal-line cooling at metallicities below 10$^{-2} Z_\\odot$ in low density gas does not significantly affect the dynamics of a collapsing halo.  In a companion paper, \\citet{Jappsen07b} found that the fragmentation of a metal-poor ($Z = 10^{-3.5} Z_\\odot$) collapsing object may depend more on the conditions, e.g. turbulence and angular momentum distributions, created during the assembly of such a halo than some critical metallicity.  Perhaps when the protogalactic gas cloud starts to host multiple sites of star formation, the associated SNe produce sufficient dust in order for a transition to Pop II.  In lower mass halos, the SN ejecta is blown out of the halo, and future star formation cannot occur until additional gas is reincorporated into the halo.  However in these halos with masses $\\gsim$10$^7$\\Ms, the SN does not totally disrupt the halo.  A fraction of the SN ejecta is contained within halo and could contribute to subsequent sites of star formation.  SN ejecta and associated dust could instigate the birth of the first Pop II stars.  The metallicities of the first generation of metal-enriched stars could depend on the metal mixing timescales in these early dwarf galaxies.  Fortunately the dispersion of heavy elements from SNe into the Galactic ISM is a rich field of study \\citep[for a review, see][]{Scalo04}.  Metallicity dispersions in local stellar clusters show fluctuations of 5--20\\% around the mean, which suggest that the ISM is well-mixed \\citep[e.g.][]{Edvardsson93, Garnett00, Reddy03}. In addition, very metal-poor ([Fe/H] $<$ --2.7) halo stars have very little scatter in elemental abundances, suggesting that the dispersal of the first metals that formed low mass stars originated from single starbursts instead of single SN explosions \\citep{Cayrel04}.  In the ISM, laminar flows and turbulence driven by SN explosions provides the impetus of metal mixing on large scales, which then cascades to smaller and smaller length scales eventually reaching a length scale associated with a Reynolds number $Re = 1$ \\citep[e.g.][]{deAvillez02}.  In structure formation, turbulence can also arise during the virialization of cosmological halos \\citep{Wise07a, Greif08}.  After the turbulent cascade creates metallicity gradients on small enough scales, molecular diffusion will homogenize the heavy elements.  \\citeauthor{deAvillez02} performed AMR simulations that include turbulent diffusivity to study the mixing timescales in the ISM.  In their simulations and the ones presented here, numerical diffusion provides the majority of the mixing at the resolution limit.  They find that the mixing time is fairly independent of the diffusion scale and could depend on some inertial scale.  The origin and magnitude of diffusion does not significantly affect the gas on large scales.  In their resolution study, mixing timescales only decrease by 20\\% when the resolution is increased by a factor of 4.  Thus we could be overestimating the mixing timescales up to a factor of a few because it is difficult to resolve the smallest turbulent scale in cosmological simulations.  Furthermore, we expect metals that are newly incorporated into the galaxy to be stirred by the turbulence in assembling halos, which is sustained in virialization and major mergers on the largest scales.   As discussed above, the metallicity of the most massive halo fluctuates around $10^{-3} Z_\\odot$.  This is intriguingly the same value as a sharp cutoff in stellar metallicities in four local dSph's: Sculptor, Sextans, Fornax, and Carina \\citep{Tolstoy04, Helmi06}. This differs with the galactic halo stars, whose metal-poor tail extends to $Z/Z_\\odot = 10^{-4}$ \\citep{Beers05}.  We must take care when comparing our results to observations since we made the simplification that every Pop III star produces a PISN \\citep[for a detailed semi-analytic model of metal enrichment, see][]{Tumlinson06}. As discussed in \\citeauthor{Helmi06}, the galactic halo may be composed of remnants of galaxies that formed from high-$\\sigma$ density fluctuations, and dwarf galaxies originate from low-$\\sigma$ peaks.  In this scenario, the objects (or its remnants) simulated here would most likely reside in galactic halos at the present day.  If we attempt to match this metallicity floor of $10^{-4} Z_\\odot$ in the galactic halo, this requires $\\sim$8\\Ms~of metals produced for every Pop III star or roughly one in ten Pop III stars ending in a PISN. More likely, the current nearby dSphs are hosted by larger dark matter halos than we have been able to simulate to date.  Hence too simple chemical evolution inspired extrapolations may be premature.    Radiative feedback from Pop III stars play an important role in shaping the first galaxies.  We studied the effects of this feedback on the global nature of high-redshift dwarf galaxies, using a set of five cosmology AMR simulations that accurately model radiative transfer with adaptive ray tracing.  Additionally, we focused on the metal enrichment of the star forming halos and their associated star formation histories.  Our key findings in this paper are   1. Dynamical feedback from Pop III stars expel nearly all of the baryons from low-mass host halos.  The baryon fractions in star forming halos never fully recover even when it reaches a virial temperature of 10$^4$ K.  The baryon fraction is reduced as low as $\\sim$0.05 with SNe feedback, three times lower than the cases without stellar feedback.  2. Baryons on average gain angular momentum as they are expelled in feedback driven outflows.  When it is reincorporated into the halo, it increased the spin parameter up to a factor of 4 with radiative and SNe feedback.  3. The accurate treatment of radiative transfer produces a complex, multi-phase ISM that has densities and temperatures that can span up to 6 orders of magnitude at a given radius.  4. Pair-instability SN preferentially enrich the IGM to a metallicity an order of magnitude higher than the overdense filaments adjacent to the sites of star formation.  5. Once a SN explosion cannot expel the gas in its host halo, the mean metallicity fluctuates around \\zstar{-2.6} as there may be a balance between SN outflows, cold inflows, and contained SNe ejecta.   We conclude that Pop III stellar feedback plays an integral part to early galaxy formation as it determines the characteristics of the first galaxies."
    },
    "0710/0710.1165_arXiv.txt": {
        "abstract": "Many physical properties of galaxies correlate with one another, and these correlations are often used to constrain galaxy formation models.  Such correlations include the color-magnitude relation, the luminosity-size relation, the Fundamental Plane, etc. However, the transformation from observable (e.g. angular size, apparent brightness) to physical quantity (physical size, luminosity), is often distance-dependent. Noise in the distance estimate will lead to biased estimates of these correlations, thus compromising the ability of photometric redshift surveys to constrain galaxy formation models. We describe two methods which can remove this bias.  One is a generalization of the $V_{\\rm max}$ method, and the other is a maximum likelihood approach.  We illustrate their effectiveness by studying the size-luminosity relation in a mock catalog, although both methods can be applied to other scaling relations as well. We show that if one simply uses photometric redshifts one obtains a biased relation; our methods correct for this bias and recover the true relation. ",
        "introduction": "The `configuration space' we use to describe galaxies is large, but galaxies do not fill it. The luminosity, color, size, surface brightness, stellar velocity dispersion, morphology, stellar mass, star formation history and spectral energy distribution of a galaxy are all correlated with one another.  These correlations encode important information about galaxy formation, and so quantifying them provides important constraints on models.  Current (e.g. SDSS, Combo-17, MUSYC, Cosmos) and planned surveys (e.g. DES, LSST, SNAP) go considerably deeper in multicolor photometry than in spectroscopy, or are entirely photometric.  For such surveys, reasonably accurate photometric redshift estimates are or will be made.  The question then arises as to which galaxy observables and correlations are affected by the noisy distance estimate associated with a photometric rather than spectroscopic redshift.  The most widely studied property is luminosity - clearly, errors in the distance result in incorrect luminosity estimates.  If not accounted for, this leads to a biased estimate of the luminosity function (e.g. Subbarao et al. 1996). Hence, there has been considerable effort devoted to the question of how to correct for this bias (e.g. Chen et al. 2003), and the problem has now been solved (Sheth 2007).  The next step is to recover an unbiased estimate of not just the luminosity function, but the joint distribution of luminosity, color, size, etc., from photometric redshift datasets. The main goal of the present work is to provide an algorithm which does this for a magnitude limited photometric redshift survey. Because the same distance error which leads to a mis-estimate of the luminosity will produce a correlated mis-estimate of the size, we have chosen to phrase the discussion in terms of the size-luminosity relation - it exhibits all the features of interest.  Section 2 illustrates the nature of the problem by showing the bias in the size-luminosity relation which results from treating photometric redshifts as though they were spectroscopic redshifts. This is done by constructing a mock galaxy sample and then perturbing the true redshifts to mimic photometric redshift errors. Section 3 places this problem in the more general context of inverse problems in statistical astronomy, and argues that a deconvolution algorithm, such as that due to Lucy (1974), is well-suited to removing the bias. It shows the result of applying this non-parametric deconvolution technique to a mock galaxy sample. Section~4 provides a maximum-likelihood formulation and solution of the problem. A final section summarizes our findings and discusses possible further studies and applications.  Where necessary, we write the Hubble constant as $H_0 = 100h~{\\rm km~s}^{-1}~{\\rm Mpc}^{-1}$, and we assume a spatially flat cosmological model with $(\\Omega_M,\\Omega_{\\Lambda}, h)=(0.3, 0.7, 0.7)$, where $\\Omega_M$ and $\\Omega_{\\Lambda}$ are the present-day densities of matter and cosmological constant scaled to the critical density. We use $D_{\\rm L}(z)$ to denote the luminosity distance; the angular diameter distance is $D_{\\rm A}(z) = D_{\\rm L}(z)/(1+z)^2$. ",
        "conclusions": "We presented two algorithms for reconstructing the intrinsic correlations between distance-dependent quantities in apparent magnitude limited photometric redshift datasets.  One was a generalization of the non-parametric $V_{\\rm max}$ method (Section~\\ref{Vmax}), and the other used a maximum-likelihood approach (Section~\\ref{ml}).  Both our reconstruction methods assume that the distribution of photo-$z$ errors is known accurately.  In practice, this means that spectroscopic redshifts are available for a subset of the data. The question then arises as to whether or not the number of spectra which must be taken to specify the error distribution reliably is sufficient to also provide a reliable (spectroscopic!) estimate of these scaling relations.  If so, what is the basis for deciding that it is worth reconstructing these relations from the photo-$z$ data? This is the subject of work in progress, although the methods presented in this paper assume that such reconstructions will indeed be necessary. E.g., if the spectra are not simply a random subset of the magnitude limited photometric sample, then it may be difficult to quantify and so correct for the selection effects associated with the spectroscopic subset.  We used the size-luminosity relation in a mock catalog which had realistic choices for the correlation to illustrate the biases which are present and must be corrected if photometric redshift datasets are to provide reliable estimates of galaxy scaling relations (Figures~\\ref{M_Me_R_Re_early_types} and~\\ref{M_R_correlation_rec}). We showed that our iterative deconvolution scheme provides a simple and reliable correction of this bias (Figures~\\ref{M_R_correlation_rec}--\\ref{p_R_given_M_rec}).  Note that although we have illustrated our methods using a 2-dimensional distribution, the extension to $n$-correlated variables is trivial.  Because our algorithm permits the accurate measurement of many scaling relations for which spectra were previously thought to be necessary (e.g. the color-magnitude relation, the size-surface brightness relation, the Photometric Fundamental Plane), we hope that our work will permit photometric redshift surveys to provide more stringent constraints on galaxy formation models at a fraction of the cost of spectroscopic surveys.  Our results may have other applications.  For example, Bernardi (2007) has highlighted a bias associated with the correlation between stellar velocity dispersion $\\sigma$ and luminosity $L$ which arises if the distance indicator used to estimate $L$ is correlated with $\\sigma$ (as may happen in the local Universe, where peculiar velocities make spectroscopic redshifts unreliable distance estimators). It may be that the methods presented here would allow an accurate reconstruction of the true relation from the biased one. This is the subject of on-going work."
    },
    "0710/0710.1023_arXiv.txt": {
        "abstract": "We report on the discovery of the $z=1.016$ cluster RzCS 052 using a modified red sequence method, followup spectroscopy and X-ray imaging.  This cluster has a velocity dispersion of $710 \\pm 150$ km s$^{-1}$, a virial mass of $4.0 \\times 10^{14}\\ M_{\\odot}$ (based on 21 spectroscopically confirmed members) and an X-ray luminosity of $(0.68 \\pm 0.47) \\times 10^{44}$ ergs s$^{-1}$ in the [1-4] keV band. This optically selected cluster appears to be of richness  class 3 and to follow the known $L_X - \\sigma_v$ relation for high redshift X-ray selected clusters. Using these data, we find that the halo occupation number for this cluster is only marginally consistent with what expected assuming a self-similar evolution of cluster scaling relations, suggesting perhaps a break of them at $z\\sim1$. We also rule out a strong galaxy merging activity between $z=1$ and today. Finally, we present a Bayesian approach to measuring cluster velocity dispersions and X-ray luminosities in the presence of a background: we critically reanalyze recent claims for X-ray underluminous clusters using these techniques and find that the clusters can be accommodated within the existing $L_X - \\sigma_v$ relation. ",
        "introduction": "Clusters of galaxies are not only a powerful tool to study galaxy evolution but can also be used to constrain cosmological parameters, resolving several parameter degeneracies (e.g.,  Allen et al. 2004;  Albrecht et al. 2006). In particular, clusters at high redshifts ($z > 1$), of which only a handful are currently known, provide the greatest leverage in determining the nature of the acceleration constant (e.g., Rapetti 2007). These determinations, however, rely on an accurate estimate of the cluster mass, whose uncertainty is arguably the dominant contributor to the error budget in deriving cosmological parameters from cluster statistics (Henry 2004; Albrecht et al. 2006).  Ideally, one wish to apply the virial theorem to get a direct measurement of cluster masses. In fact, the dark matter velocity dispersion is an extremely good tracer of the halo masses in all simulations (Evrard et al. 2007), and galaxies are nearly unbiased velocity tracer (Evrard et al. 2007 and references therein; Rines Diaferio \\& Natarajan 2007), in good agreement previous works (Biviano et al. 2006, Tormen et al. 1997). The measurement of the cluster velocity dispersion requires a large number of radial velocities, which are observationally expensive to obtain, especially for high redshift clusters. For this reason and because each mass estimator carries some key informations, more commonly the scaling between pairs of more easily observable mass-related quantities is studied, such as X-ray luminosity, temperature or the $Y_X$ (Kravtsov et al. 2006) parameter, or optical richness. These studies often look for outliers, however their search is blessed by data limitation: for example in the search of clusters X-ray dim for their optical richness, Donahue et al. (2001) and Gilbank et al. (2004), both mostly worked with putative clusters (i.e. not spectroscopically confirmed) and X-ray undetections.  \\begin{figure*} \\psfig{figure=RzCS052_150.ps,width=12truecm,clip=} \\caption[h]{True-color ($z'[3.6][4.5]$) degraded-resolution (to make galaxies not too small when printed) image of a region of a 24 Mpc$^2$ area around RzCS 052. Spectroscopically confirmed clusters and isodensity contours for red galaxies are also marked. Note the number density contrast of reddish galaxies between the cluster center and the right part of the image. The ruler is 1 arcmin long; North is up, East is to the left.} \\end{figure*}   Only few works directly address the relative quality of different mass estimators with velocity dispersion: Borgani \\& Guzzo (2001) compare the scatter of two mass estimators, X-ray luminosity and richness, and found that the former is a better mass tracer than the latter when the former is uniformly measured and the latter is taken from a 50 years old paper reporting eye-estimate of the cluster optical richness (the Abell 1958 catalog). In both CNOC and nearby clusters, mass correlates better with richness than with X-ray luminosity (Yee \\& Ellinson 2003; Popesso et al. 2005). Eke et al. (2004) found that optical luminosity is a better proxy of mass than velocity dispersion in common conditions, i.e. when velocities are available for a small sample of galaxies. A related issue, which we will examine below, is whether it exists clusters X-ray dim for their mass (velocity dispersion), (e.g. Lubin, Mulchaey \\& Postman 2004, Fang et al. 2007, Johnson et al. 2006).   The relation between richness and mass has received some recent attention in the form of the halo occupation function (Berlind \\& Weinberg 2002; Lin et al. 2004 and references therein) whose first moment is the halo occupation number (HON), the average number $N$ of galaxies per cluster of mass $M$. In order to address the evolution of the HON, velocity dispersion information is often unavailable for a large cluster sample, mass and cluster size are inferred from other mass-related quantities (for example the X-ray temperature), and assumed to evolve self similarly. The evolution of the HON with redshift is still unclear: the initial study by Lin et al. (2004) claimed that the HON increases at high redshift, but Lin et al. (2006) find evidence that it does not evolve strongly out to $z \\sim 1$, suggesting that the galaxy population in clusters was established and assembled at early epochs. Muzzin et al. (2007) confirms the above, with a sample of reduced redshift leverage and hence reduced evolution sensitivity, but available velocity dispersion information.  \\begin{figure*} \\centerline{\\psfig{figure=figpapI1.ps,width=8truecm,clip=} \\psfig{figure=figpapI2.ps,width=8truecm,clip=}} \\caption[h]{Spectra of RzCS 052 members coming the VLT run. We have vertically shifted the spectra and zoomed on a reduced wavelength range for display purpose.} \\end{figure*}  \\begin{figure*} \\psfig{figure=cmRz.ps,width=18truecm,clip=} \\caption[h]{ {\\it Left:} Colour--magnitude diagram for galaxies (close circles) within 1 arcmin from the cluster center or with a known redshift (red closed circles for members, blue crosses for interlopers) using CTIO discovery data. $R$ and $z'$ mag completeness limits are show with dashed lines. The green line is the expected CM at the cluster redshift, from Kodama \\& Arimoto (1997) {\\it Right:} Colour histograms of galaxies within 2 arcmin from the cluster center (solid histogram), and of the average control field (measured on a 0.36 deg$^2$ area, dashed histogram), normalized to the cluster area. A clear excess is seen, especially at $R-z>1.5$ mag. The hashed histogram is the colour distribution of spectroscopically confirmed member galaxies. In both panels, a few objects with spectroscopic redshift are missing because their fall on bad CCD regions or have extreme colours.} \\end{figure*}    Here, we present the photometric discovery, spectroscopic confirmation and X-ray properties of a new $z=1.016$ cluster of galaxies (RzCS 052), a cluster optically rich but undetected in the XMM-LSS survey (Pierre et al. 2007), and hence possibly X-ray dark (i.e. dim for its mass). We derive its global properties (richness, X-ray luminosity, velocity dispersion and mass) and study these in the context of cluster scaling relations ($L_X - \\sigma$, HON) at high redshift. In particular, we test the claim that the HON (the way galaxies populate cluster-scale haloes) has not changed $z \\sim 1$ (Lin et al. 2006) under far less assumptions than the original claim. We also present a Bayesian approach to the determination of cluster velocity dispersion and X-ray luminosity and use it to critically examine recent claims about the existence of underluminous X-ray clusters. A companion paper (Andreon et al. 2007) addresses the use of RzCS 052 as a laboratory for studying galaxy formation and evolution.  We adopt $\\Omega_\\Lambda=0.7$, $\\Omega_m=0.3$ and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. Magnitudes are quoted in their native photometric system (Vega for $R$, SDSS for $z'$). ",
        "conclusions": "We have identified a distant cluster from a modified red-sequence method and followed it up spectroscopically. RzCS 052 is a richness class 3 cluster at $z=1.016$ with a velocity dispersion of $710 \\pm 150$ km s$^{-1}$ and an X-ray luminosity of $0.68 \\pm 0.47 \\times 10^{44}$ ergs s$^{-1}$.  In spite of its optical detection, RzCS 052 obeys to the high redshift $L_X-\\sigma_v$ relationship as other X-ray selected clusters to the high redshift $L_X-\\sigma_v$ relationship, whereas in principle variations in the dynamical state of the clusters or in the thermal history of the intracluster medium may have moved it away from the $L_X-\\sigma_v$ relation.  Analysis of the $N-M$ scaling shows that RzCS 052 has the right number of galaxies (actually, a bit less) that it should have for its mass, ruling out intense merging (among galaxies) activities in clusters from $z=1$ to today, in agreement with Lin et al. (2006).  We present a Bayesian approach to measuring cluster velocity dispersions (most useful for sparsely sampled data and in presence of a background) and X-ray luminosities or upper limits (essential in the case of poorly determined parameters). Critical re-analysis of the data of clusters/groups claimed to be outliers of the $L_X-\\sigma_v$ relationship leads to conclude that there are no known thus far examples of clusters X-ray underluminous for their velocity dispersion. The above result is quite reassuring for ongoing X-ray surveys: there is thus far no example of cluster missed because an anomalous $L_X$ for the cluster mass.  \\begin{table*} \\caption{Luminosity and velocity dispersion of clusters at $z>0.8$} \\begin{tabular}{llcrcrr} \\hline Name & z & $\\log L_X [1-4]$ & ref & $ \\sigma_v$ & N \\hfill & ref \\\\ \\hline RXJ1716+6708   & 0.813 & $44.68\\pm0.03$ & 12 & $1522 \\pm 180 $& 37 & 1 \\\\ RXJ1821.6+6827 & 0.816 & $44.62\\pm0.01$ & 2 & $775  \\pm 122 $& 18 & 2 \\\\ MS1054-30321   & 0.830 & $44.91\\pm0.04$ & 12 & $1153 \\pm 80  $& .. & 3 \\\\ RXJ0152-1357S  & 0.830 & $44.42\\pm0.02$ & 12 & $737  \\pm 126 $& 18 & 4 \\\\ RXJ0152-1357N  & 0.835 & $44.59\\pm0.03$ & 12 & $919  \\pm 168 $& 16 & 4 \\\\ RzCS 530       & 0.839 & $44.20\\pm0.09$ & 5 & $780  \\pm 126 $& 17 & 5 \\\\ 1WGA1226+3333  & 0.890 & $45.17\\pm0.01$ & 12 & $997  \\pm 245 $& 12 & 6 \\\\ Cl1604+4304    & 0.900 & $44.00\\pm0.06$ & 14,7 & $962  \\pm 141 $& 67 & 7 \\\\ RzCS 052       & 1.016 & $43.83\\pm0.37$& this work & $710 \\pm 150 $& 21 & this work \\\\ RXJ0910+5422   & 1.106 & $44.00\\pm0.05$ & 12 & $675  \\pm 190 $& 25 & 8 \\\\ RXJ1252-2927   & 1.237 & $44.37\\pm0.07$ & 12 & $747  \\pm 79  $& 38 & 9 \\\\ LynxW          & 1.270 & $43.70\\pm0.23$ & 12 & $650  \\pm 170 $& 9  & 10 \\\\ 1WGAJ2235.3    & 1.393 & $44.57\\pm0.05$ & 11 & $762  \\pm 265 $& 12 & 11 \\\\ \\hline \\end{tabular} \\hfill \\break \\footnotesize{RzCS 530 is also known as XLSSC 003; References: 1: Gioia et al. (1999); 2: Gioia et al. (2004a); 3: Gioia et al. (2004b); 4: Demarco et al. (2005); 5: Valtchanov et al. (2004); 6: Maughan et al. (2004); 7: Gal \\& Lubin (2004); 8: Mei et al. (2006); 9: Demarco et al. (2007); 10: Stanford et al. (2001); 11: Mullis et al. (2005); 12: Ettori et al. (2004); 14: Lubin et al. (2004). \\hfill \\break } \\end{table*}"
    },
    "0710/0710.1812_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{Sec:intro} \\setcounter{equation}{0}  \\subsection{Experiments and theories} \\label{Sec:ExTh}  An ongoing and forthcoming array of experiments (e.g. WMAP, \\cite{Spergel:2006hy}, SDSS \\cite{York:2000gk, Tegmark:2003uf, Tegmark:2006az, McDonald:2004xn}, SNLS \\cite{Astier:2005qq}, ACBAR \\cite{Kuo:2006ya}, Planck \\cite{Planck}, ACT \\cite{ACT}, Spider \\cite{Spider}) is measuring the cosmic microwave background (CMB) and the large scale structure of the universe with unprecedented precision. This provides exciting opportunities to reveal the nature of the early universe and the underlying fundamental theories. The leading theoretical candidate for creating the initial conditions of our universe is inflation \\cite{Guth:1980zm,Linde:1981mu,Albrecht:1982wi}. However inflation remains a paradigm, which can be implemented by a variety of models underpinned by differing microphysical constructions; as the constraints from data tighten, there is the hope that we might identify the specific scenario that describes our universe. With the natural ingredients of such model-building being supergravity and string theory, the process of better measuring the properties of the early universe is also a process of understanding better the theory of quantum gravity.  In contrast to recent debates on the predictivity of the string theory landscape, here we use a more conventional approach to investigate the predictivity of string theory by studying the properties and exploring the dynamics of our own vacuum. We first scan the parameter space of inflationary models subject only to the requirement that they provide enough inflationary $e$-folds to solve the flatness and horizon problems. This is because the natural creation of a homogeneous and isotropic universe is the leading problem that we want to solve, and is perhaps the most attractive feature of the inflationary paradigm. After that, we study the observational consequences of all the viable parameter spaces, with the goal of looking for distinctive signatures. Some of these can be compared with observations and used to narrow down the parameter space. Despite of the vastness of all possible vacua in the string landscape, this process can be rather effective since certain observational features rely on distinctive dynamics. As we will see, such dynamics can either be field-theoretic with strong motivations from string theory, or completely stringy in nature.  There are many candidate observable signatures in inflationary models. The most generic ones are the amplitude of the primordial power spectrum and its spectral index. Since most viable models built from a fundamental theory have adjustable parameters to fit these two observables, this leaves a large number of viable models that are consistent with the data, and even leaves the nature of the inflaton field ambiguous. In principle, Nature is not obligated to provide more information within our experimental abilities, and indeed there is no evidence for further parameters required to describe the current data. But anticipating her generosity, possible distinctive observables that might be measurable in the future include the scale-dependence  (``running'') of the spectral index, departures from Gaussianity of the primordial fluctuations, a tensor contribution to the primordial power spectrum, and cosmic strings. These will be crucial to successfully carry out the program that we have outlined.  With the rapidly improving quality of cosmological data, it will become increasingly interesting to implement the above program by comparing specific models to data, starting directly from microscopic parameters of theories. Modern cosmological data analyses make use of the powerful method known as Markov Chain Monte Carlo (MCMC) to implement the comparison to data, providing an efficient way of estimating posterior distributions of the microscopic parameters. However, in practice, when directly using microscopic parameters as MCMC parameters, highly non-linear relationships between the parameters and observables may introduce severe obstacles for MCMC to efficiently search the parameter space. Therefore, a reparameterization according to the specific nature of the model often becomes necessary. So instead of a straightforward exercise, implementing MCMC becomes a rather interesting model-dependent art. It is also a purpose of this paper to use an example to illustrate this process and extract certain model-independent procedures of such reparameterization which may be of more general interest.  \\subsection{Brane inflation}  The inflationary models that we study in this paper belong to the brane inflation scenario proposed by Dvali and Tye \\cite{Dvali:1998pa,HenryTye:2006uv}.\\footnote{For recent reviews on other types of string inflation models, see Ref.~\\cite{Burgess:2007pz,Linde:2007fr,Kallosh:2007ig,Cline:2006hu}.} We are interested in these models precisely because they can give rise to a large number of distinctive observational signatures. This happens even in the simplest scenarios that provide inflation. One of the most important reasons that makes it possible is that brane inflation can be achieved via two different mechanisms, namely slow-roll and Dirac-Born-Infeld (DBI) inflation.  The original models of brane inflation \\cite{Dvali:1998pa,Burgess:2001fx,Dvali:2001fw,Kachru:2003sx} are slow-roll inflationary models \\cite{Linde:1981mu,Albrecht:1982wi}, where branes and anti-branes slowly approach each other in a flat potential. A model that uses this mechanism in the framework of the string theory flux compactification \\cite{Douglas:2006es} is studied by Kachru, Kallosh, Linde, Maldacena, McAllister and Trivedi (KKLMMT) \\cite{Kachru:2003sx}. As in the F-term inflation models in supergravity \\cite{Lyth:1998xn}, it is found that the generic shape of the potential is too steep to achieve the slow-roll inflation, in this case due to the moduli stabilization. Again, similar to those supergravity models, it is possible that several contributions to the potential manage to cancel to a certain precision so that the potential becomes sufficiently flat. There are effective parameters in the model controlling the inflaton mass that can be adjusted to fit the observed spectral index \\cite{Bean:2007hc}. The running of the spectral index, non-Gaussianities and tensor modes are all too small to be observed in the near future.\\footnote{Some observables become measurable if there are sharp features in the potential \\cite{Adams:2001vc,Peiris:2003ff,Covi:2006ci,Chen:2006xj,Hailu:2006uj}. In addition, there are other important observational possibilities of brane inflation -- cosmic strings and those related to reheating \\cite{Polchinski:2004ia,HenryTye:2006uv} -- which apply to both slow-roll and DBI inflation.}  Another inflationary mechanism that is so far uniquely found in brane inflation is the DBI inflation \\cite{Silverstein:2003hf,Alishahiha:2004eh,Chen:2004gc,Chen:2005ad}. In DBI inflation, the rolling velocity of inflaton branes is not determined by the shape of the potential but by the speed-limit of the warped internal space. Such warped spaces are naturally present in the extra dimensions due to fluxes used to stabilize the string compactification \\cite{Giddings:2001yu}.  The first model that uses such a mechanism is that of Silverstein, Tong and Alishahiha (STA) \\cite{Silverstein:2003hf,Alishahiha:2004eh}. In this model, as the branes roll into a throat from the UV side of the warped space under a quadratic potential, its velocity gets restricted by the large warping in the IR side of the warped space. However, instead of having a potential with a generic mass term, a rather steep potential, characterized by a large inflaton mass, is required to achieve this UV DBI inflation. The reason is that, when the branes enter from the UV side of the warped space in the GKP-type warped compactification \\cite{Giddings:2001yu}, the energy provided by the antibranes sitting at the IR side is not large enough to drive DBI inflation even if there is the speed-limit, since the antibrane tension has been warped down correspondingly. Therefore an extra, steep, potential has to be added to raise the inflationary energy. In addition, embedded in the same warped compactification, the model generates large non-Gaussianities that exceed the experimental bound \\cite{Bean:2007hc,Peiris:2007gz}, as well as excessive probe brane backreaction which we will address in Sec.~\\ref{Sec:UV}. This is because in this model, the levels of non-Gaussianity and probe brane backreaction sensitively depend on the inflaton value, and it is viable only if the inflaton field is of (super-)Planckian size. However, the range of the inflaton field is restricted by some geometric conditions of the compactification and is sub-Planckian \\cite{Chen:2005fe,Chen:2006hs,Baumann:2006cd,Bean:2007hc,Peiris:2007gz}.  To fully make use of the speed-limit of the warped space, it is better to make the branes roll out from the IR end, and use antibranes in other throats to provide the inflationary energy. In this way the speed-limit of the branes and the inflationary energy become relatively independent of each other, leaving a rather flexible shape of the inflaton potential which has been the main problem of model-building. This is the model proposed in Ref.~\\cite{Chen:2004gc,Chen:2005ad}. It can be generically realized in the multi-throat brane inflation scenario \\cite{Chen:2004gc}.  It happens that in this IR DBI inflation model, the large non-Gaussianities can also be small enough to satisfy the current observational bound \\cite{Chen:2005fe}. This is partly because no matter how small the warp-factor (and consequently, how big the non-Gaussianity) the branes begin with, the level of non-Gaussianity decreases as the branes roll out and approaches its minimal value at the end of the inflation. Therefore, in the segment of the warped space traversed during the last 60 $e$-folds, the level of non-Gaussianity is among the smallest in the entire DBI inflation trajectory. Moreover, the geometric conditions that put a tight constraint on the STA model are automatically satisfied in the IR DBI model and has no effect on the non-Gaussianities.  Besides providing a speed-limit to the inflaton, another important property of warped space is the reduction of the local fundamental string scale \\cite{Randall:1999ee}. This turns out to have important consequences on density perturbations in DBI inflationary models. During the epoch when the string scale is red-shifted below the Hubble parameter, the quantum fluctuations on the inflaton branes become stringy.\\footnote{Notice that such a stringy phase only happens in the inflaton sector, which is the deep IR side of a warped space with energy density of order $H^4$, so it does not backreact significantly on the Hubble expansion. We also note that such a stringy phase will backreact on the IR side of the warped geometry, but it is estimated that this still leaves a large enough portion of the geometry for DBI inflation to take place \\cite{Chen:2005ad,Chen:2006ni}. We will discuss this more in Sec.~\\ref{Sec:IR} \\& \\ref{Sec:Power}.} The density perturbations are no longer fully described by the usual field theory approximation, and acquire distinctive stringy signatures. In the IR DBI model, this stringy phase corresponds to earlier inflationary $e$-folds, and therefore larger scales in the sky. It is estimated that such a phase transition will give rise to a large transient (regional) running of the spectral index \\cite{Chen:2005ad,Chen:2005fe}. In this paper, we make this prediction more quantitative and compare it to observations.   \\subsection{Outline}  Following the strategy that we outlined in Sec.~\\ref{Sec:ExTh}, in this paper, we first summarize the overall features of brane inflation using phase diagrams that describe the parameter spaces spanned by both inflationary mechanisms, \\emph{i.e.} slow-roll vs. DBI (Sec.~\\ref{Sec:phase}), reviewing the key observational predictions in the different parts of the parameter space.  The main focus of this paper is to compare the IR DBI brane inflation model to observations (Sec.~\\ref{Sec:IRDBI}). We derive analytical and numerical model predictions for the shape of the power spectrum, non-Gaussianity, and tensor modes, giving a quantitative estimate of the effect of the Hubble-expansion-induced stringy phase transition on density perturbations (Appendix \\ref{Sec:Width}).  We then proceed to compare these results to the observational data from cosmic microwave background and large scale structure (Sec.~\\ref{Sec:MCMC}). We outline how such a comparison should be generally implemented using MCMC. The current data give a number of interesting constraints on the microscopic parameters of the model (Sec.~\\ref{conc}), including the mass of the brane moduli potential, the fundamental string scale, the charge or warp factor of throats, and the number of the mobile branes. We also quantify some distinctive observable signatures of this model, such as the level of the non-Gaussianity and the running of the spectral index. We discuss how the latter is observationally different from two other cases that may also give large running spectral index: slow-roll inflation with mild features on potential (Appendix \\ref{Sec:Mild}), and slow-roll or DBI inflation with a non-Bunch-Davies vacuum (Appendix \\ref{Sec:nonBD}). These results illustrate how string theory can make testable predictions which might be subject to observational constraints.  For convenience, all the variables used in this paper are summarized in Table~\\ref{Tb:variables}.  \\begin{table*}[!ht] \\begin{center} {\\small \\begin{tabular}{cll} \\hline Variable &  Description & Notes\\\\ \\hline \\hline $\\mpl$ & 4d reduced Planck mass & $\\mpl = (8\\pi G)^{-1/2}=2.4\\times 10^{18} {\\rm GeV} $\\\\ $m_s$ & Mass scale of fundamental strings & $m_s \\equiv \\alpha'^{-1/2}$  \\\\ $g_s$ & String coupling & $g_s<1$ \\\\ $T_3$ & D3-brane tension &  Eq.~(\\ref{RTDef}) \\\\ $R$ & Length scale of warped throat & Eq.~(\\ref{RTDef}) \\\\ $M$, $K$ & Flux numbers in warped throat & Integers \\\\ $n_A$ & Number of antibranes in A-throats &  \\\\ $n_B$ & Number of branes (inflatons) in B-throat & \\\\ $N_B$ & Effective charge of B-throat & $N_B = a_B MK$ \\\\ $a_B$ & Multiplicative factor from orbifolding & $a_B\\sim 1$ in data analysis \\\\ $\\lambda_B$ & $\\lambda_B = n_B N_B/2\\pi^2$ & Eq.~(\\ref{lambdaDef}) \\\\ $r$ & Radial coordinate of throats &  \\\\ $\\phi$ & Canonical inflaton field & $\\phi=r\\sqrt{nT_3}$ \\\\ $h_A$ & Minimum warp factor of A-throat &  \\\\ $h_B(\\phi)$ & Warp factor at location $\\phi$ in B-throat & $h\\equiv r/R=\\phi R/\\sqrt{\\lambda}$ \\\\ $\\beta$ & Characterization of shape of potential & Inflaton mass $m^2 = \\beta H^2$ \\\\ $\\gamma$ & Lorentz factor of inflaton & \\\\ $c_s$ & Sound speed in 4d & $c_s = 1/\\gamma$ for DBI inflation \\\\ $V_0$ & Inflationary energy density & \\\\ $N_e$ & Number of $e$-folds to the end of inflation\\\\ $N_e^{\\rm DBI}$ & Number of $e$-folds to the end of IR DBI inflation& \\\\ $N_{\\rm tot}^{\\rm NR}$ & Total $e$-folds of non-relativistic roll inflation & Typically fast-roll \\\\ $k_c$ & Critical scale of the stringy phase transition & Eq.~(\\ref{kcdef})\\\\ $N_c$ & Critical DBI $e$-fold at $k_c$ & Eq.~(\\ref{Ncdef})\\\\ $P(k)$ & Power spectrum &\\\\ $r_{TS}$ & Tensor to scalar ratio & \\\\ $f_{NL}^{\\rm eq}$ & Estimator of the non-Gaussianity & Equilateral shape\\\\ \\hline \\end{tabular} } \\caption{\\label{Tb:variables} \\small Description of variables. In the text, subscripts $A$ and $B$ are frequently added to some of the variables, referring to the quantities of the A- or B-throat.} \\end{center} \\end{table*} ",
        "conclusions": "\\label{conc}  We conclude by highlighting the main results and discussing their physical implications. The quoted ranges are at the $95\\%$ confidence level, and we have combined constraints from both the power spectrum and non-Gaussianity. The detailed $68\\%$ and $95\\%$ CL marginalized constraints and the maximum likelihood values are listed in Table \\ref{Tab_cosmo}.  \\subsection{Microscopic parameters}  \\begin{itemize}  \\item {\\em Shape of the inflaton brane moduli potential: $1.3<\\beta < 3.7$.} The lower bound is due to constraints from the power spectrum, while the upper bound is due to the non-Gaussianity constraint. It is encouraging that, while IR DBI inflation can happen for a range of $\\beta$ that varies over nearly 10 orders of magnitude, $0.1 \\lesssim\\beta <10^9$ (see Eq.~(\\ref{mrange})), comparison with data picks out a very small range around $\\CO(1)$ which is generically expected theoretically. This makes an explicit construction of such potentials a more interesting question.  \\item {\\em Fundamental string scale: $-9.6 < \\log_{10}(m_s/\\mpl)/g_s^{1/4} < -5.3 $.} The upper bound on the string scale is due to the large charge, and hence length scale, of the B-throat required to fit the amplitude of the density perturbations. The lower bound is due to the fact that a smaller string scale tends to increase the total number of $e$-folds of non-relativistic fast-roll inflation, and make the running of the spectral index too large (Fig.~\\ref{Fig:2d}). The model prefers an intermediate fundamental string scale, $10^8\\ {\\rm GeV} < m_s/g_s^{1/4} < 10^{13}\\ {\\rm GeV}$, and therefore an intermediate large volume compactification, $8.9\\times 10^7 < V^{1/6} \\mpl <4.8\\times 10^{13}$, where $V$ is the compactification volume.  \\item {\\em B-throat charge: $8.8< \\log_{10} N_B < 10.4 $; Number of inflaton branes: $3.9<\\log_{10} n_B< 5.1$.} In terms of the GKP-type warped compactification, this implies flux numbers $K\\sim M \\sim \\sqrt{N_B} \\sim \\CO(10^5)$. Explicit construction remains an open question as discussed in Sec.~\\ref{Sec:OpenQu}. In the multi-throat brane inflation scenario, inflaton branes are generated from flux-antibrane annihilation. The number of branes generated in this process is roughly determined by the flux number $M$. Indeed, a small number of inflaton branes is ruled out by the data.  \\item {\\em A-throat minimum warp factor: $-2.4<\\log_{10} h_A \\le 0$.} This is from combining the constraint on $n_B$ and $n_A h_A^4$, $h_A = (n_A h_A^4/n_B)^{1/4}$. A smaller $h_A$ leads to larger $N^{\\rm NR}_{\\rm tot}$ and larger running of the spectral index (Fig.~\\ref{Fig:2d}). So the A-throat tends to be short. This makes tunneling reheating possible, where many interesting phenomena can occur, such as an intermediate matter-dominated epoch.  \\end{itemize}  \\subsection{Secondary derived parameters}  \\begin{itemize}  \\item {\\em Inflationary phases.} In this model, not all $e$-folds comes from IR DBI inflation. The last $13< N^{\\rm NR}_{\\rm tot}< 24$ $e$-folds come from non-relativistic fast-rolling inflation, which is possible because inflatons are close to the top of the potential.  \\item {\\em The stringy phase transition.} The Hubble-expansion induced stringy phase transition happens at the largest scales in the sky, $-6.0<\\log_{10} k_c/{\\rm Mpc} <-2.9$. However its impact on density perturbations extends over to shorter scales, such as generating a transient large running of the spectral index.  \\item {\\em Inflation scale: $-10.0<\\log_{10}V_0^{1/4}/\\mpl<-5.1$.} This gives a very small tensor to scalar ratio $r_{TS}<10^{-13}$.  \\item {\\em Cosmic string tension: $-23<\\log_{10} G\\mu_D +\\log_{10}g_s^{1/2} < -14$.} Here the cosmic strings refer to the D-strings left over from the brane-antibrane annihilation in the A-throat, whose tension is $G\\mu_D= (m_s h_A/g_s^{1/4} \\mpl)^2/(16\\pi^2 g_s^{1/2})$. There is an unconstrained freedom coming from the additive factor $\\log_{10}g_s^{1/2}$, but it is not expected to give any significant contributions. The F-string tension differs by a factor of $g_s$, $\\mu_F = g_s \\mu_D$.  \\end{itemize}  \\subsection{Observational predictions}  \\begin{itemize}  \\item {\\em Large, but regional, running of spectral index: $-0.046< dn_s/d\\ln k(k=0.02/{\\rm Mpc}) <-0.010$.} A reconstructed full-scale power spectrum and the running of the spectral index are shown in Fig.~\\ref{Fig:Pk} \\& \\ref{Fig:nrun}.  This prediction is stringy in nature. A better understanding of the theoretical details and better measurements of both the power spectrum and non-Gaussianities on the relevant scales may reveal finer structures. In future experiments, Planck is expected to achieve $\\sigma(dn_s/d\\ln k) = 0.005$ \\cite{Planck}.  \\item {\\em Large non-Gaussianities: $-272<f_{NL}^{\\rm eq}(k=0.02/{\\rm Mpc}) <-70$.} A reconstructed full-scale prediction is in Fig.~\\ref{Fig:fnl}, which shows the running of the non-Gaussianities.  This prediction is strictly speaking field-theoretic, but with strong string theory motivations, such as warped compactification and the DBI brane action. This field theoretic regime is $k>k_c$; the theoretical analysis for non-Gaussianities at $k\\lesssim k_c$ is currently unavailable and remains an interesting open question. In future experiments, on CMB scales, Planck can achieve $\\sigma(f^{\\rm eq}_{NL}) =67$ \\cite{Smith:2006ud,Hikage:2006fe}; on large scale structure scales, some high-$z$ galaxy surveys can reach similar or better precision \\cite{Sefusatti:2007ih}.   \\end{itemize}  As seen from these results, constraints from cosmological data, and even relatively loose constraints such as the non-Gaussianity constraint, are already putting strong restrictions on models which aim to provide self-consistent microphysical descriptions of the early universe. With the bounty of precision cosmological data expected in the future, the hope of probing not just field-theoretic, but string-theoretic early universe physics burns brightly.  \\medskip"
    },
    "0710/0710.5600_arXiv.txt": {
        "abstract": "We study magnetic effects induced by rigidly rotating plates enclosing a cylindrical MHD Taylor-Couette flow at the finite aspect ratio $H/D=10$. The fluid confined between the cylinders is assumed to be liquid metal characterized by small magnetic Prandtl number, the cylinders are perfectly conducting, an axial magnetic field is imposed $\\Ha \\approx 10$, the rotation rates correspond to $\\Rey$ of order $10^2-10^3$. We show that the end-plates introduce, besides the well known Ekman circulation, similar magnetic effects which arise for infinite, rotating plates, horizontally unbounded by any walls. In particular there exists the Hartmann current which penetrates the fluid, turns into the radial direction and together with the applied magnetic field gives rise to a force.  Consequently the flow can be compared with a Taylor-Dean flow driven by an azimuthal pressure gradient. We analyze stability of such flows and show that the currents induced by the plates can give rise to instability for the considered parameters.  When designing an MHD Taylor-Couette experiment, a special care must be taken concerning the vertical magnetic boundaries so they do not significantly alter the rotational profile. ",
        "introduction": "Motion of a fluid confined between two concentric, rotating cylinders is a classical problem in hydrodynamics and, if the fluid is conducting and an external magnetic field is applied, magnetohydrodynamics (MHD). The flow of this type, usually referred to as the Taylor-Couette flow, has been first studied by Couette~\\citep{Couette1888} and later was subject of a seminal work by Taylor~\\citep{Taylor1923}, who experimentally confirmed theoretical results of a linear stability analysis.  In the field of MHD, an important work was done by Velikhov~\\citep{Velikhov59} who has shown that for the conducting fluid a weak magnetic field can play a destabilizing role and can lead to an instability which today is called magnetorotational instability (MRI~\\citep{Jietal01}).  When studying the Taylor-Couette system it is common to assume some simplifications, the small gap approximation or large aspect ratio. In the former it is assumed that the gap between the cylinders $D = \\Rout-\\Rin$ is small compared to the radii, i.e., $D/\\Rout \\ll 1$, this allows the neglect of terms of order $1/R$, $R$ being distance from the center of rotation.  When considering the large aspect ratio, one assumes that the height of the cylinders $H$ is much larger than the gap width $\\Gamma=H/D \\gg 1$, which guarantees that a secondary flow due to the plates bounding the cylinders is insignificant and does not disturb the rotational profile of the fluid.  On the other hand, there is also plenty of work done for small aspect ratio $\\Gamma \\approx 1$, where the rigidly rotating end-plates play crucial role and simply introduce a new class of problems.  When $\\Gamma$ becomes an important parameter it is possible to observe a wide family of different states (including non-axisymmetric ones or peculiar asymmetric patterns -- anomalous modes) for the same parameters, so that the observed results depend on their path through the parameters space from an initial state.  Therefore this system is an excellent subject to the bifurcation theory~\\citep{1988JFM...191....1P, LopezMarques03, Mullinetal02, Furukawaetal02, Kageyamaetal04, Youd06}.  In the present work we focus on the case of wide gap $\\Rin/\\Rout=1/2$ and $\\Gamma=10$ which is an intermediate aspect ratio, between very short and long containers, yet in purely hydrodynamical contest the influence of the vertical boundaries is small, at least for Reynolds numbers of order \\ord{10^2-10^3}.  However, if the rotation rates are large enough, so that the corresponding Reynolds number is \\ord{10^5} and larger, the plates can easily dominate the flow in the entire container. This is due to the Taylor-Proudman theorem, from which follows that in rapidly rotating systems the flow tends to align itself along the axis of rotation.  For such rotations it is necessary that $\\Gamma$ would have to be several thousand in order to obtain the rotational profile which is not profoundly altered by the end-plates~\\citep{HollerbachFournier04}.  Results of a recent MRI experiment PROMISE~\\citep{MRIexpA, MRIexpB, Stefani07}, as well as nonlinear simulations~\\citep{SzklRud06,Szklarski07} indicate that for a flow with relatively small Reynolds number $ \\approx 10^3$, and parameters resembling essentially MHD stable flow in the limit of infinitely long cylinders, there exist unexpected time-dependent fluctuation of the velocity field.  These disturbances arise as an effect of the vertical boundary conditions, moreover the simulations show that they are much stronger if the end-plates bounding the cylinders are assumed to be perfectly conducting.  The plates induce a well known hydrodynamical effect -- the Ekman circulation, which is a result of unbalanced pressure gradients in vicinity of the vertical no-slip boundary conditions.  There  the Ekman layer develops  in which the fluid velocity from the bulk of the container must match the velocity imposed by the end-plates.  It seems that for MHD Taylor-Couette flow, magnetic effects, unlike the classical hydrodynamical Ekman layer, induced by the plates have been overlooked. In this paper we argue that the rigidly rotating plates together with an imposed axial magnetic field give rise to a similar layer which develops for an infinite, rotating plate serving as a boundary for the conducting fluid.  One of the most important features of such flow is the existence of the Hartmann current (absent in the conventional Hartmann problem,~\\citep{2004JFM...504..183K}) which leaves the boundary layer and then interacts with the magnetic field. In particular, this becomes important for conducting plates which was the case for the PROMISE experiment, since one of the end-plates was made from copper.  We discuss properties of Ekman-Hartmann layers for infinite, rotating plates and relate it to the end-plates enclosing the cylinders in a Taylor-Couette setup.  It is shown that for considered radial boundary conditions, the induced current turn eventually in the radial direction and acting in concert with the imposed axial magnetic field gives rise to a body force.  We demonstrate that magnetic effects induced by the end-plates enclosing the cylinders can profoundly alter flow properties. In particular the rotational profile can become significantly different from the expected parabolic Couette solution.  Moreover, if the Hartmann current is strong enough, it is likely that the local Rayleigh criterion for stability will be violated and the flow becomes centrifugally unstable.  In an MRI experiment it is crucial to rule out such instabilities and a special care concerning the vertical boundary conditions is needed in order to obtain the desired rotational profile. ",
        "conclusions": "\\citet{GilmanBenton68} have shown with a linear theory that in vicinity of a rotating plane which serves as a border for rotating conducing fluid there develops the Ekman-Hartmann layer if $\\Omplain \\ne \\Omfluid$ and an axial magnetic field is applied.  The most important feature of Ekman-Hartmann layers is their ability to induce both mass fluxes and electric currents in the region outside the boundary layer.  If $\\Omplain < \\Omfluid$ these fluxes are directed outwards the layer (``blowing''); when $\\Omplain > \\Omfluid$ towards the layer (``suction''). For the conducting plates the fluxes are much stronger since additional currents are drawn from/into the plates.  Outside the Ekman-Hartmann layer exists the magnetic diffusion region, in which the electric current has only radial components. The current, together with the axial magnetic field, produces an electromagnetic body force acting on the fluid.  We have shown in this paper that similar effects arise for the MHD Taylor-Couette flow when the rotating cylinders are bounded by two rigidly rotating end-plates.  Near the plates the Ekman-Hartmann layer forms and, consequently, there exists the Hartmann current which penetrates bulk of the fluid. In the presence of an axial magnetic field such problem can be compared with the Taylor-Dean flow -- a flow between, possibly rotating, cylinders which is additionally driven by an azimuthal pressure gradient.  We find that under certain conditions the resulting flow becomes unstable, Taylor vortices can be observed and the rotational profile is significantly different from the standard Couette solution $\\Omegab$. The instability has essentially a centrifugal character as the Rayleigh criterion is locally vialoted.  This is an undesirable effect from the point of view of an MRI experiment.  In such experiment it is necessary to obtain a state resembling $\\Omegab$ in the major part of the container, for parameters characterizing stable MHD flows.  It is necessary to take into account the magnetic effects induced by the plates so that the MRI can be clearly identified rather then any other instability.   The fluxes induced in the Ekman-Hartmann layer are a direct consequence of a shear close to the boundaries.  Exemplary methods of reducing the shear have been proposed in~\\citep{Szklarski07}.  For rotation rates characterized by $\\Rey$ of order \\ord{10^3} all the effects can by significantly reduced by allowing the end-plates to rotate independently of the cylinders~\\citep{Abshagenetal04}.  Since for $\\Omend = \\Omin$ there is the Ekman suction, and for $\\Omend=\\Omout$ the Ekman blowing, there exists $\\Omout < \\Omend < \\Omin$ for which the generated mass and charge fluxes are minimal.  Alternatively, one can divide the plates into independently rotating rings~\\citep{JiBurin06}."
    },
    "0710/0710.0025_arXiv.txt": {
        "abstract": "{Data taken at the Pierre Auger Observatory are used to search for air showers initiated by ultra-high energy (UHE) photons. Results of searches are reported from hybrid observations where events are measured with both fluorescence and array detectors. Additionally, a more stringent test of the photon fluxes predicted with energies above $10^{19}$~eV is made using a larger data set measured using only the surface detectors of the observatory.}  \\begin{document} ",
        "introduction": "It has been suggested that the excess of cosmic-rays with energies above the predicted Greisen-Zatsepin-Kusmin (GZK) steepening, observed by AGASA, could be due to cosmic-rays being produced as the by-products of 'top-down' scenarios and not by more conventional acceleration mechanisms, such as the diffusive shock process, in astrophysical objects. If the former were the case, it would be expected that a significant proportion of the spectrum of cosmic rays at the highest energies would be UHE photons: predictions vary from model to model but are in the range 10\\% to 50\\% above 10~$\\eev$.  The extensive air showers (EAS) generated by UHE photons reach shower maximum \\xmax, the atmospheric depth at which the number of charged particles in the shower is greatest, at much greater depths than their nuclear counterparts. This is due to the much lower multiplicity in particle production in the electromagnetic-dominated photon showers than in the hadronic interactions present for nuclear primaries. The depth of maximum is further increased above 30~$\\eev$ because of the suppression of Bethe-Heitler pair production due to the LPM effect \\cite{LPMEffect_LP}\\cite{LPMEffect_M}, which is not important for other cosmic-ray primaries.  \\xmax~ can thus be used to discriminate between photon and nucleonic UHE primaries. With the Pierre Auger Observatory this can be done using hybrid events which are observed with both fluorescence and array detectors, and in which \\xmax~ is measured directly. This allows the determination of a limit to the integral fraction of photons above 10~$\\eev$. Parameters measured using only the surface array of detectors, which reflect shower maximum, can also be compared to predictions for photons from simulations to provide a stronger limit on the flux of photons at several energies \\cite{SDParametersPhotonLimit}. Two such observables are used which behave differently for nuclear primaries when compared to photons: these are the risetime of the signal in the water-Cherenkov detectors at 1000 m from the core and the radius of curvature of the shower front. ",
        "conclusions": "\\subsection{Hybrid Results} In \\cite{OriginalHybridPhotonLimit} we reported a limit to the fraction of photons in the integral cosmic-ray flux of 16\\% (95\\% c.l.) above 10~$\\eev$ based on 29 high-quality hybrid events recorded in the period Jan. 2004 - Feb. 2006. We have updated the analysis with data collected until March 2007, keeping the analysis cuts as in the original paper. In total, 58 events are now available. The measured \\xmax~distribution is shown in figure \\ref{HybridXmaxDistribution} along with the calculated distribution from 10~$\\eev$ photons made on an event-by-event basis. Even the largest observed value of \\xmax, $\\sim$~900~$\\gcm$, is well below the average value expected for photons (about $\\sim$~1000~$\\gcm$, see e.g. Table 1 in \\cite{OriginalHybridPhotonLimit}). With the updated sample, the upper limit becomes 13\\% (95\\% c.l.) above 10~$\\eev$.  \\begin{figure}[htbp] \\vspace{-10pt} \\begin{center} \\includegraphics[width=0.45\\textwidth]{icrc0602_fig03.eps} \\end{center} \\caption[font=small,labelfont=bf, aboveship=2pt]{Black: The distribution of \\xmax from 58 hybrid events with energies above 10~$\\eev$ that meet selected cuts\\cite{OriginalHybridPhotonLimit} . The dashed line shows a distribution for 10~$\\eev$ photons arriving over a range of zenith angles.  All events reach shower maximum at depths which are too shallow to be considered as a result of photon primaries.} \\label{HybridXmaxDistribution} \\end{figure}  \\begin{figure}[htbp] \\vspace{-0pt} \\begin{center} \\includegraphics[width=0.45\\textwidth, height=0.4\\textwidth]{icrc0602_fig04.eps} \\end{center}  \\caption[font=small,labelfont=bf, aboveship=2pt]{The principle component between risetime and radius of curvature as a function of energy for data (black) and simulated photons (red). Data lying above the dashed line, which is the mean of the distribution for photons, are identified as photon candidates. No events meet this requirement.} \\label{UCLAPrincipleComponentPlot} \\end{figure}   \\subsection{Surface Detector Results}  Photon candidates were searched for by combining the risetime and radius of curvature for each event into one observable using a principle component analysis, where the angle of rotation is found using 5\\% of the data. The remaining 95\\% of showers were then identified as nuclei or photon candidates using an \\emph{a priori} cut, where showers are excluded as photon-like if the principle component measurement is less than the mean of that predicted by a spectrum of photonic showers generated by Monte Carlo simulations. Events which were deemed as photonic were assigned energies using a reconstruction designed from photon simulations, reflecting the fact that these showers have low muon content and are more strongly attenuated than hadronic showers \\cite{Billior}. The spectrum of simulated photons incorporates the efficiency of photon detection and reconstructions, as well as including the possibility that photons may interact with the geomagnetic field before arriving in the atmosphere.  The results are summarised in figure \\ref{UCLAPrincipleComponentPlot} where the principle component is plotted for data (black) and the simulated photons (red) as a function of energy. There are 0 events above the dashed line, where the cut is defined, and as such there are no candidate photons for events above 10 $\\eev$. This constrains the maximum flux of photons to $3.8 \\times 10^{-3}$,~$2.5 \\times 10^{-3}$~and $2.2 \\times 10^{-3}$~\\fluxunit{} above 10, 20 and 40 $\\eev$ respectively. The upper limit on the photon fraction, based on the measured spectrum of events in \\cite{SommersSpectrum}, is calculated as 2.0\\%, 5.1\\% and 31\\% at 10, 20 and 40 $\\eev$ respectively. These limits are shown against the predictions for photons from top-down models based on the AGASA spectrum (see \\cite{review}) in figure \\ref{ModelsVsLimits}. Also shown are upper limits to the flux and fractions of photons from previous experiments. This work constrains the photon fractions and fluxes to more stringent limits than previously measured and disfavours the proposed top-down models as the sources of UHECRs and of the AGASA events.  On completion of the surface array, the Pierre Auger Observatory will reach sensitivities of $4 \\times 10^{-4}$~\\fluxunit{} for the integrated flux and 0.7\\% for the fraction of photons above 20 $\\eev$ (95\\% c.l.) after 5 years of operation.  \\begin{figure}[htbp] \\vspace{-120pt} \\subfigure[Upper limits to photon fractions] { \\centering \\includegraphics[width=0.45\\textwidth]{icrc0602_fig05.eps} } \\subfigure[Upper limits to photon flux] { \\centering \\includegraphics[width=0.45\\textwidth]{icrc0602_fig06.eps} } \\caption[font=small,labelfont=bf, aboveship=2pt]{(a) The upper limits to the fraction of photons including limits from previous experiments and predictions from top-down models based on the AGASA spectrum. The limits from this work are shown in black. The limit derived from the hybrid data is also shown and labelled FD. Labels A, HP and Y refer to the limits set by AGASA, Haverah Park and Yakutsk, for references see \\cite{review}. (b) The upper limits to the flux of photons along with predictions of top-down models.} \\label{ModelsVsLimits} \\end{figure}"
    },
    "0710/0710.2399_arXiv.txt": {
        "abstract": "Using the IRAM 30 m telescope, a mapping survey in optically thick and thin lines was performed towards 46 high mass star-forming regions. The sample includes UC~H{\\sc ii} precursors and UC~H{\\sc ii} regions. Seventeen sources are found to show \"blue profiles\", the expected signature of collapsing cores. The excess of sources with blue over red profiles ([$N_{\\rm blue}$ -- $N_{\\rm red}$]/$N_{\\rm total}$) is 29\\% in the HCO$^+$ $J$=1--0 line, with a probability of 0.6\\% that this is caused by random fluctuations. UC~H{\\sc ii} regions show a higher excess (58\\%) than UC~H{\\sc ii} precursors (17\\%), indicating that material is still accreted after the onset of the UC~H{\\sc ii} phase. Similar differences in the excess of blue profiles as a function of evolutionary state are not observed in low mass star-forming regions. Thus, if confirmed for high mass star-forming sites, this would point at a fundamental difference between low- and high-mass star formation. Possible explanations are inadequate thermalization, stronger influence of outflows in massive early cores, larger gas reserves around massive stellar objects or different trigger mechanisms between low- and high- mass star formation. ",
        "introduction": "Inflow motion is a fundamental phenomenon during stellar formation. Although the search for inflow is usually more difficult than that for outflow, studies of inflow have made great progress since the 1990s. In low-mass star forming regions, inflow motions were detected at different evolutionary stages, including Class --I, Class 0 and Class I cores \\citep{zho93, mmt97, lee99, gre00, evans03}. Recently, a number of inflow candidates were found in high mass star formation regions. Among a sample of 28 massive cores, 12 were found to show line profiles that peak at blue-shifted velocities (hereafter \"blue profiles\"; see Sect.\\,3.1), the expected signature of inflow \\citep{we03}. \\citet{fws05} (hereafter FWS05) detected such asymmetric profiles in 22 cores within a sample of 77 high-mass proto-stellar objects (HMPOs). Most recently, \\citet{wyr06} detected 9 sources with a blue profile in a sample of 12 ultracompact (UC)~H{\\sc ii} regions.  Variation of inflow motion with time is critical for high mass star formation. It has been indicated that when a protostar reaches $>$10 M$_{\\odot}$ it can generate enough radiation pressure to halt spherical infall and inhibit its mass increasing\\citep{wc87}. Observationally, however, it is not yet clear how inflow is related to the evolution of massive (proto)stars.  To study this problem, we have carried out a survey for a sample including both cores of UC~H{\\sc ii} regions and precursors of UC~H{\\sc ii} regions.  While previous surveys using single point observations provided some statistical evidence for the occurrence of infall within massive cores, blue profiles can also be caused by rotation. Therefore maps of the molecular environment are indispensable. Mapping also allows us to locate the center of the inflow and to identify cores that are simultaneously showing evidence for in- and outflow.  Therefore, we conducted a mapping survey including 46 high mass star-forming regions which were selected applying three criteria: (1) The sources must have been mapped in the submillimeter or millimeter wavelengths with continuum or spectroscopy; (2) signal-to-noise ratios should be $>$5 at 350\\,$\\mu$m (Mueller et al. 2002) and higher at other wavelengths; (3) there should be no other core within one arcmin \\citep{zhs97, hnb98, tie98, hat00, mbc00, bsp02, mse02}. With respect to their stellar content, we can divide the sample into two different groups of targets: (1) Thirty three sources lack 6 cm continuum emission and are precursors of UC~H{\\sc ii} regions or HMPOs \\citep{mbc00, bsp02}. Among these, thirty are hosting a luminous IRAS source. The remaining three are associated with IRAC (the InfraRed Array Camera on the Spitzer Space Telescope) point sources (W3-W and W3-SE) or are not hosting an IRAC source (18454--3). All 33 cores comprise `group~I'. (2) Thirteen UC~H{\\sc ii} regions are assigned to `group~II'. This letter presents  a list of the identified collapse candidates and provides the statistics of blue excesses. Detailed properties of individual cores will be analyzed in a future paper. ",
        "conclusions": "\\subsection{Blue profile identification}  For self-absorbed optically thick lines, the classical signature of inflow is a double peaked profile with the blue-shifted peak being stronger, or a line asymmetry with the peak skewed to the blue side. While optically thin lines should show a single velocity component peaking at the line center.  Among the 46 cores observed, five (05490+2658, G31.41+0.31, 18454-3, 18454-4, 19266+1745) will be ignored because they show either too complex spectral profiles, inhibiting a detailed analysis, or a lack of optically thin lines. Estimates of optical depths were obtained from line ratios between different isotopomers of CO and CS and from the relative intensities of individual hyperfine components in the case of C$^{17}$O and N$_2$H$^+$. C$^{18}$O, C$^{17}$O, C$^{34}$S and N$_2$H$^+$ tend to be optically thin, while CS is optically thick. HCO$^+$ opacities could not be estimated. However, the similarity of HCO$^+$ and CS line shapes (see Sect.\\,3.2) as well as the results of \\cite{gre00} and FWS05 clearly indicate that HCO$^+$ is also optically thick.  The 41 remaining sources were detected in at least one optically thick and one optically thin line. A blue profile caused by inflow motion with velocity $v \\propto r^{-1/2}$ in a region with higher excitation temperature ($T_{ex}$) inside requires $T_{\\rm A}$*(B)/$T_{\\rm A}$*(R) $>$ 1. Here $r$ is the radius of the collapsing core \\citep{zho93}. $T_{\\rm A}$*(B) and $T_{\\rm A}$*(R) are the blue and red peak intensities of the optically thick line. We also define a dimensionless asymmetry parameter following \\cite{mmt97}, $\\delta V$ = ($V_{\\rm thick}$-$V_{\\rm thin}$)/$\\Delta$$V_{\\rm thin}$. $V_{\\rm thick}$ is the peak velocity of the opaque line, $V_{\\rm thin}$ and $\\Delta$$V_{\\rm thin}$ denote the peak velocity and width of the optically thin line. Only for $\\delta V < -0.25$ or $> 0.25$ the line profile is rated blue or red, respectively.  Our sources (Table~2) discriminate among five main types of line shapes: (1) cores with lines showing a ``blue profile'' (in the following denoted with B); (2) cores with lines showing a \"red profile\" (R); (3) cores exhibiting blue and red profiles at different spatial positions (BRS); (4) cores where some lines show a blue profile, while others display a red profile (BRL); (5) cores without obvious asymmetric lines (S). Only cores showing at least one line of type B, but no lines of type R are identified as targets potentially undergoing inflow motion.  \\subsection{Collapse candidates and their profile ``excess''} \\label{bozomath}  With the criteria outlined in Sect.\\,3.1, seventeen inflow candidates are identified (see Table~2). Ten belong to group~I and seven are part of group~II. To provide a typical example, Fig.\\,1 shows the infall signature of the group~I core W3-SE. Fig. 1a displays the HCO$^+$\\,(1--0) spectra, showing the angular size of the core. Fig. 1b shows a number of profiles towards the central position. The HCO$^+$\\,(1--0) and (3--2) lines as well as the CS\\,(3--2) transition show the blue asymmetry. For the HCO$^+$\\,(1--0) line this is also demonstrated in the position-velocity (P-V) diagram of Fig.\\,1c. For comparison, Fig.\\,1d shows a P-V diagram of the optically thin C$^{18}$O\\,(1--0) emission.  The quantity ``excess'' as defined by \\cite{mmt97} is $E$ = ($N_{\\rm B} - N_{\\rm R}$)/$N_{\\rm T}$, where $N_{\\rm B}$ and $N_{\\rm R}$ mark the numbers of sources with blue and red profiles. $N_{\\rm T}$ is the total number of sources. For our survey the excess was calculated for the two HCO$^+$ transitions and the CS\\,(3--2) line. Fig. 2 shows the log[$T_{\\rm A}$*(B)/$T_{\\rm A}$*(R)] and $\\delta$V (see Sect.\\,3.1) distributions of the three individual lines. Statistical results are given in Table~3. The observed excess derived from the HCO$^+$\\,(1--0) and (3--2) lines is 0.29 and 0.11, respectively. Both are larger than those obtained by FWS05 for the same lines (0.15 and 0.04). For the CS transition we obtain 0.29. To evaluate the statistical significance of the determined values, we conducted the binomial test (see FWS05 and references therein). Probabilities that the excesses are a product of a random distribution are given in the last column of Table~3. These are 0.006 and 0.01 for HCO$^+$\\,(1--0) and CS (3--2) respectively. Apparently, both lines are sensitive tracers of potential inflow motion in massive cores.  To evaluate differences between the two classes of cores (I and II; see Sect.\\,1) with respect to the excess, we used the HCO$^+$\\,(1--0) line, which was mapped in the largest number of sources. The results listed in the lower part of Table~3 include 16 sources with profiles of type B. The excesses observed for group~I and II are 0.17 and 0.58, respectively.  Twenty of our 46 sources overlap with those of FWS05. Among them are 19 group~I sources (out of 33), but only one source is from group~II (out of 13). Our study includes various CO and CS lines. We also made maps. Thus we can view the common objects from a different perspective and can check, how far the choice of different molecular transitions and the presence of maps is leading to contradictions with previously published results. Differences are indeed significant. For eight of the 19 overlapping type I cores we obtain different line asymmetry classification, emphasizing the need for detailed maps. Nevertheless, the overall difference in the HCO$^+$\\,(1--0) excess is negligible (0.17 versus 0.15).  To summarize, both data sets indicate that the HCO$^+$\\,(1--0) excess is low for UC~H{\\sc ii} precursors. For  UC~H{\\sc ii} regions, our results and those of \\cite{wyr06} suggest that the excess is larger and more significant. From the binomial test for group~I and II, the probability that the blue excesses (0.17 and 0.58 respectively) are arising by chance is 0.13 and 0.008, respectively.  \\subsection{A comparison with low mass star-forming surveys}  While low mass star-forming regions show infall from the Class --I to the Class I stages of evolution, high mass star-forming regions also exhibit infall signatures from their earliest stages till a UC H{\\sc ii} region has formed \\citep{wzx05,bkl06,qzm06}. In low mass cores the profile excess was found to be 0.30, 0.31 and 0.31 for Class --I, 0 and I core samples in the HCO$^+$\\,(3--2) line (\\cite{evans03} and references therein). There seem to be no significant differences among the cores in different evolutionary phases. However, our samples show the excess of UC~H{\\sc ii} regions far surpassing that of the UC~H{\\sc ii} precursors. This may point to fundamental differences between low and high mass star-forming conditions. Possible causes to the higher blue excess in Group II sources may be: (1) The molecular gas surrounding UC H{\\sc ii} regions may be more adequately thermalized to show the blue excess, i.e. the excitation temperature of specific lines may increase more monotonically towards the center. Thus all lines may produce blue profiles indicating infall motion, while in younger cores still some lines may show red profiles. (2) The amount of dense cool gas is larger towards younger objects. Outflows of dense molecular gas may be more active around Group I objects, shaping more red profiles. (3) Low mass cores are relatively isolated and their gas supply is limited. Simulations showed that this may halt the increase of inflow \\citep{vor05}. However, high mass stars form in giant molecular clouds and their inflow motions are not easily halted by the exhaustion of molecular gas before most of it is dispelled. (4) In low mass cores, star formation may be spontaneous. In high mass cores, collapse may be trigged by extrinsic disturbances and the collapse may take more time to develop.  With respect to potential selection effects, we used the same criteria to identify the targets of the two separate groups of sources. Since this study is based on a limited number of sources, more data quantifying the blue excess as a function of evolutionary stage would be highly desirable."
    },
    "0710/0710.2350_arXiv.txt": {
        "abstract": "{Gamma-ray bursts (GRBs) have been detected up to GeV energies and are predicted by many models to emit in the very high energy (VHE, $>$ 100 GeV) regime too. Detection of such emission would allow us to constrain GRB models. Since its launch, in late 2004, the Swift satellite has been locating GRBs at a rate of approximately 100 per year. The rapid localization and follow-up in many wavelengths has revealed new and unexpected phenomena, such as delayed emission in the form of bright X-ray flares. The Milagro gamma-ray observatory is a wide field of view (2 sr) instrument employing a water Cherenkov detector to continuously ($>$ 90\\% duty cycle) observe the overhead sky in the 100 GeV to 100 TeV energy range. Over 100 GRBs are known to have been in the field of view of Milagro since January 2000, including 57 since the launch of Swift (through May 2007). We discuss the results of the searches for prompt emission from these bursts, as well as for delayed emission from the X-ray flares observed in some of the Swift bursts.}     \\begin{document} ",
        "introduction": "Almost 40 years after the detection of the first gamma-ray burst (GRB), many questions remain about these powerful explosions. Progress in the field has accelerated in the decade since the detection of the first X-ray afterglow~\\cite{costa}. The launch of Swift~\\cite{gehrels}, in late 2004, has resulted in the rapid and accurate localization of $\\sim$100 GRBs per year. This has enabled an unprecedented number of multiwavelength follow-up observations in every band of the electromagnetic spectrum, from radio to the highest energy gamma rays, paving the way for a more complete understanding of these phenomena.  The highest energy emission to be detected conclusively from GRBs was seen by the EGRET instrument. EGRET detected several bursts above 100 MeV, with no evidence for a cut-off in the spectrum~\\cite{dingus01}. One burst in particular, GRB 940217, emitted an 18 GeV photon over 90 minutes after the start of the burst~\\cite{hurley94}. Another burst, GRB 941017, was found to display a second, higher energy, spectral component which extended up to at least 200 MeV and decayed more slowly than the lower energy component~\\cite{gonzalez03}. On this evidence alone, the expectations are that the upcoming launch of GLAST, with its much higher sensitivity than EGRET, will result in a large number of detections at high energies, leading to a better understanding of the GeV properties of GRBs.  At very high energy (VHE, $>$100 GeV), there have been no conclusive detections of GRBs, although there have been several tantalizing hints of emission. Milagrito, a prototype of Milagro, searched for emission coincident with 54 BATSE bursts and reported evidence for emission above 650 GeV from GRB 970417a, at the 3$\\sigma$ level~\\cite{atkins00a,atkins03}. The HEGRA group reported evidence at the 3 sigma level for emission above 20 TeV from GRB 920925c~\\cite{padilla98}. Follow-up observations above 250 GeV by the Whipple atmospheric Cherenkov telescope~\\cite{connaughton97} failed to find any high energy afterglow for 9 bursts studied, though the delay in slewing to observe these bursts ranged from 2 minutes to almost an hour. More recent efforts by the MAGIC~\\cite{magicGRBlimits}, Whipple~\\cite{whippleGRBlimits}, and Milagro~\\cite{milagro1} telescopes have resulted only in upper limits.  There is no shortage of models predicting (or explaining) very high energy emission from GRBs (e.g.~\\cite{dermer00,pilla98,zhang01,granot03,gupta07}). Many of the proposed emission mechanisms predict a fluence at TeV energies comparable to that at keV-MeV energies. One such mechanism involves the inverse Compton upscattering of lower energy (synchrotron) photons by the energetic electrons which emitted them. The strong magnetic fields and large bulk Lorentz factors present in GRB jets can result in a high energy component peaked at TeV energies. The strength of such a component is highly dependent on the environments of the particle acceleration and the gamma ray production.  Milagro\\cite{atkins00b,atkins01} is a TeV gamma-ray detector, located at an altitude of 2630 m in northern New Mexico. Milagro uses the water Cherenkov technique to detect extensive air-showers produced by VHE gamma rays as they interact with the Earth's atmosphere. The Milagro field of view is $\\sim$2 sr and duty cycle is $>$90\\%. The effective area is a function of zenith angle and ranges from $\\sim50$ m$^2$ at 100 GeV to $\\sim10^5$ m$^2$ at 10 TeV. A sparse array of 175 4000-l water tanks, each with a PMT, was added in 2002. These ``outriggers,'' extend the physical area of Milagro to 40000 m$^2$. The combination of large field of view and high duty cycle make Milagro the best instrument currently available for conducting a search for prompt VHE emission from GRBs. Milagro is also able to operate in ``scaler'' mode, where individual tube rates can be monitored to search for GRB emission in the 1--100 GeV energy range coincident in time with known satellite bursts (see ~\\cite{taylor,cesar}). In addition to conducting a search for emission from satellite-detected GRBs, Milagro is capable of performing a blind search for emission at all locations in its field of view and many different durations (see ~\\cite{vlasios}).  \\subsection{The GRB Sample}  Twenty-five satellite-triggered GRBs occurred within the field of view of Milagro between January 2000 and December 2001. No significant emission was detected from any of these bursts~\\cite{atkins05}. In the period from January 2002 to December 2004 (post-BATSE and pre-Swift), there were only 11 well-localized GRBs within the Milagro field of view. Since the launch of Swift (through the end of May 2007), however, there have been a total of 57 bursts in the field of view of Milagro~\\footnote{A good source of information on well-localized GRBs is J. Greiner's web page {\\tt http://www.mpe.mpg.de/$\\sim$jcg/grbgen.html}}, many of them with measured redshift. Table~\\ref{grb_table} lists the 57 GRBs in this sample, along with some of their properties. Due to the absorption of high-energy gamma rays by the extragalactic background light (EBL), detections at VHE energies are only expected for redshifts less than $\\sim$0.5. The degree of gamma-ray extinction from this effect is uncertain, because the amount of EBL is not well known. In this work we use the model of \\cite{primack05}, which predicts an optical depth of roughly unity for 500 GeV (10 TeV) gamma rays from a redshift of 0.2 (0.05).  \\subsection{Bright X-ray Flares}  One of the highlights of the Swift mission has been the detection of bright X-ray flares in the afterglow phase of many GRBs~\\cite{burrows05}. These flares can sometimes be as bright as the prompt component of the GRB itself. While the exact nature of the flares is not clear, with various different proposals for what may be causing them (e.g. \\cite{guetta07,kobayashi07,krimm07}), there are reasonable expectations that these flares will result in the emission of delayed very high energy photons~\\cite{wang06}. The first survey of X-ray flares from GRBs observed by Swift in the first year of operations~\\cite{chincarini} shows that approximately one third of GRBs show significant flaring activity. Table~\\ref{flare_table} lists the 10 flares from the flare catalog of ~\\cite{chincarini,falcone} which were in the field of view of Milagro. ",
        "conclusions": ""
    },
    "0710/0710.0355_arXiv.txt": {
        "abstract": "An analytical approximation to periodic orbits in the circular restricted three-body problem is provided. The formulation given in this work is based in calculations known from classical mechanics, but with the addition of the necessary terms to give a fairly good approximation that we compare with simulations, resulting in a simple set of analytical expressions that solve periodic orbits on discs of binary systems without the need of solving the motion equations by numerical integrations. ",
        "introduction": "\\label{sec:introd}  The majority of low-mass main-sequence stars seem to be grouped in multiple systems preferentially binary (Duquennoy \\& Mayor 1991, Fischer \\& Marcy 1992). These systems have attracted more attention since the discovery that many T-Tauri and other pre-main-sequence binary stars possess both circumstellar and circumbinary discs from observations of excess radiation at infrared to millimeter waves and direct images in radio (Rodr{\\'i}guez \\etal 1998; for a review see Mathieu 1994, 2000). The improvement in the observational techniques allows, on the other hand, the testing of theories on these objects. Moreover, extrasolar planets have been found to orbit stars with a stellar companion, e.g., 16 Cygni B, $\\tau$ Bootis, and 55 $\\rho$ Cancri (Butler \\etal 1997; Cochran et al. 1997). All these facts make the study of stellar discs in binary systems, as well as the possibility of stable orbits on them that could be populated by gas or particles, a key element for better understanding stellar and planetary formation.  In particular the study of simple periodic orbits of a test particle in the restricted three-body problem results a very good approximation to the streamlines in an accretion disc in a very small pressure and viscosity regime, or for proto-planetary systems and its debris (Kuiper belt objects). Extensive theoretical work has been done in this direction (Lubow \\& Shu 1975; Paczy\\'nski 1977; Papaloizou \\& Pringle 1977; Rudak \\& Paczy\\'nski 1981; Bonnell \\& Bastien 1992, Bate 1997; Bate \\& Bonnell 1997).  Several studies are directed to find the most simple and important geometric characteristic of the discs which is the size of both circumstellar and circumbinary discs, going from analytical approximations: Eggleton (1983), who provides a simple analytical approximation to the Roche lobes; Holman \\& Wiegert (1999) radii in planetary discs; Pichardo, Sparke \\& Aguilar (2005, PSA05 hereafter) radii in eccentric binary stars.  Based in the approximation studied in classical mechanics to periodic orbits given by Moulton (1963), we provide in this work a fairly good approximation by calculating the necessary terms, for the case of circular binaries, to periodic orbits. This formulation gives not only the radii of the circumstellar and circumbinary discs, but gives a good description of any periodic orbit at a given radius for any of the discs, in the form of an analytical approximation. In this manner, the set of equations we provide can be used to find any periodic orbit or its initial conditions to run it directly in simulations of the three-body restricted problem, or as a good approximation to initial conditions in the eccentric case as well, from the point of view of particle or hydrodynamical simulations, without the need of solving the restricted three body problem numerically.  The low viscosity regime described by the periodic orbits representation, is important in astrophysics since it gives an idea of how bodies would respond only influenced by the effects of the potential exerted by the binaries, and they could permit to link the periodic orbits to physical unknown characteristics of the discs like viscosity that needs to be artificially introduced to hydrodynamical codes. It also might permit the calculation of physical characteristics like dissipation rates, etc. Studies like these result easier given in the form of approximated analytical expressions, instead of the usual numerical approximations to the periodic orbits since an analytical formulation is much faster to add to any hydrodynamical or particle code that requires the characteristics of a given family of periodic orbits, or simply their initial conditions in a more precise form than assuming Keplerian discs until the approximation fails. The periodic orbits show regions where the orbits are compressed (or decompressed), these regions could trigger density fluctuations maybe able to drive material to form important agglomerations in the discs that, depending on their positions on the disc and the density, could give origin to planets. An analytical approximation with a given density law for the discs would allow these kind of studies to be much faster and easier.  In Section \\ref{sec:method}, we show the strategy and equations used to find the approximation to the periodic orbits in circumstellar discs. In Section \\ref{sec:orbits2}, we present the equations approximating the periodic orbits in circumbinary discs, and an analysis of stability in Section \\ref{sec:stability}. A comparison of the application of the formulation and numerically calculated periodic orbits is given in Section \\ref{sec:comparison}.  Conclusions are presented in Section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have constructed an analytical set of equations based on perturbative analysis that result simple and precise to approximate the solution for the circular restricted three body problem. We choose some terms in the expanded equations of motion, which are relevant to the solution of the problem. The original equations are approximated with a set of equations that, in many cases, have an analytical solution. We show the goodness of our analytical approximation by direct comparison of geometry and rotation curves with the full numerical solution.  The set of equations we prsent can provide any periodic orbit for a binary system, either circumstellar or circumbinary, and not only the outer edge radius of circumstellar discs or the inner edge radius (gap) of the circumbinary disc, without the need to solve the problem numerically.  The relations provided are simple and straightforward in such a way that our approximation can be used for any application where initial conditions of periodic orbits or complete periodic orbits are needed, or for direct study of the three body restricted circular problem. For instance, in order to introduce on a hydrodynamical or particle code, our initial conditions would result in much more stable discs than with keplerian initial conditions, or than by constructing the discs directly from hydrodynamical simulations of accretion to binaries that will require long times to obtain stable discs. Since our approximation is completely analytical it will have the obvious additional advantage of being computationally, extremely cheap, and easier to implement than the numerical solution.  The periodic orbits respond uniquely to the binary potential and do not consider other physical factors such as gas pressure, or viscosity that work to build the fine details of discs structure. They are however, the backbone of any potential and their shape and behavior give the general discs phase space structure. In this manner, apart of all the possible hydrodynamical or particle discs applications, we can directly use them to study from the general discs geometry to rotation curves, or rarification and compression zones by orbital crowding, etc."
    },
    "0710/0710.5093_arXiv.txt": {
        "abstract": "The Unified Model of AGN predicts the sole difference between Seyfert 1 and Seyfert 2 nuclei is the viewing angle with respect to an obscuring structure around the nucleus.  High energy photons above 20 keV are not affected by this absorption if the column is Compton thin, so their 30--100~keV spectra should be the same. However, the observed spectra at high energies appear to show a systematic difference, with Seyfert 1's having $\\Gamma \\sim $2.1 whereas Seyfert 2's are harder with $\\Gamma \\sim$ 1.9. We estimate the mass and accretion rate of Seyferts detected in these high energy samples and show that they span a wide range in $L/L_{Edd}$.  Both black hole binary systems and AGN show a correlation between spectral softness and Eddington fraction, so these samples are probably heterogeneous, spanning a range of intrinsic spectral indices which are hidden in individual objects by poor signal-to-noise. However, the mean Eddington fraction for the Seyfert 1's is higher than for the Seyfert 2's, so the samples are consistent with this being the origin of the softer spectra seen in Seyfert 1's. We stress that high energy spectra alone are not necessarily a clean test of Unification schemes, but that the intrinsic nuclear properties should also change with $L/L_{Edd}$. ",
        "introduction": "The simplest version of the Unification model of Antonucci \\& Miller (1985) is that the central engine (black hole, its accretion disc and broad line region: BLR) are the same in all AGN, but that this is embedded in an obscuring torus.  The nucleus is seen directly only for inclination angles which do not intersect this material, giving the classic Seyfert 1 AGN signature of a strong and variable UV/X-ray continuum and broad emission lines. Conversely, where the obscuration is in the line of sight, these features are hidden, and the presence of an AGN can only be inferred from the high excitation lines produced in the narrow line region on much larger scales (Seyfert 2's). A key piece of evidence for this scenario is the detection of polarized broad emission lines in classic Seyfert 2 galaxies (most notably in NGC 1068), showing that the BLR is present, but can only be seen via scattered light (e.g. Antonucci \\& Miller 1985; Tran 1995, Heisler et al. 1997). This scattering medium filling the 'hole' in the torus is also detected in transmission in Seyfert 1's through the partially ionised absorption signatures seen in soft X-rays (e.g. Blustin et al 2005).  \\begin{figure*} \\begin{center} \\begin{tabular}{l} \\leavevmode \\epsfxsize=12cm \\epsfbox{panel.eps} \\end{tabular} \\end{center} \\caption{Plots showing the distribution of log mass and log accretion rate for individual instruments and the full sample (bottom panels) with Seyfert 1 sources shown by a black dashed line and Seyfert 2 sources by a solid cyan line. The log mean accretion rate for each type is given in the top right hand corner.} \\label{} \\end{figure*}  The intrinsic X-ray spectrum of a type 1 AGN can be well described by a power law (of photon spectrum $N(E)\\propto E^{-\\Gamma}$) together with its Compton reflection from the optically thick accretion disc and/or torus (e.g. Nandra \\& Pounds 1994). Again, the Unified models are supported by X-ray observations, which show that this intrinsic spectrum is substantially suppressed at low energies due to absorption in Seyfert 2's (Awaki et al 1991; Smith \\& Done 1996; Turner et al 1997; Bassani et al 1999; Cappi et al 2006). However, the strong energy dependence of photoelectric absorption means that this is unlikely to affect samples above 10 keV for columns which are not Compton thick, yet high energy experiments show that the intrinsic spectra of Seyfert 2 galaxies are systematically harder ($\\Gamma$ $\\sim$ 1.9 - 2) than than the Seyfert 1's, which have $\\Gamma > $ 2 (Zdiarski et al.  1995, Gondek et al. 1996, Perola et al. 2002, Malizia et al. 2003, Beckmann et al. 2006). This is not consistent with the idea that these nuclei are identical, unless the intrinsic emission is anisotropic, being harder in the equatorial plane.  However, all accreting black hole systems (both stellar and supermassive) generally show softer intrinsic spectra at higher accretion rates (hereafter parameterised as Eddington fraction, $L/L_{Edd}$), e.g. Laor (2000); Remillard \\& McClintock (2006). Thus inclination alone is not the sole determinant of the observed spectrum, and Seyferts at different $L/L_{Edd}$ should not be expected to have the same intrinsic emission. Here we collate estimates of black hole mass and mass accretion rate for the Seyfert 1 and 2's detected out to $>50$~keV from CGRO (OSSE), BeppoSAX (PDS) and INTEGRAL (IBIS) instruments.  We show that these high energy samples of Seyferts span a large range $L/L_{Edd}$, so should include a range of intrinisic spectral slopes. The mean $L/L_{Edd}$ for the Seyfert 1's in the sample is higher than that for Seyfert 2's, consistent with the steeper spectral slope inferred for the Seyfert 1's being due to the intrinsic spectrum softening at higher $L/L_{Edd}$. We stress that samples of Seyfert 1 and 2 AGN need to be matched on intrinsic properties such as $L/L_{Edd}$ in order to explore differences in orientation.     \\begin{table*} \\begin{center} \\begin{minipage}{125mm} \\bigskip \\begin{tabular}{ll|l|c|c} \\hline Source & Instrument & logM/M$_{\\odot}$ &  L/L$_{Edd}$   &Reference\\\\ \\hline  \\vspace{1mm}MCG--6-30-15  & {\\em OSSE }, {\\em BeppoSAX}, {\\em INTEGRAL} &  6.81 $^{+0.54}_{-0.54}$  $^{\\sigma}$  &   0.25 $^{x}$    &  6 (18) + 23\\\\ \\vspace{1mm}IC 4329A    & {\\em OSSE }&   6.70 $^{+0.56}_{-peg}$ (mean) $^{r}$   & 0.79 $^{f}$  & 20+ 20\\\\ \\vspace{1mm}&&6.85  $^{+0.55}_{-peg}$ (rms) $^{r}$ &&\\\\ \\vspace{1mm}MR 2251-178& {\\em INTEGRAL}&   6.92  $^{+0.22}_{-0.22}$ $^{m}$  &  4.68 $^{z}$          & 10 + 21 \\\\ \\vspace{1mm}NGC3783    &   {\\em OSSE }, {\\em BeppoSAX}&  6.97  $^{+0.30}_{-0.97}$ (mean) $^{r}$  & 0.23 $^{f}$   &  2 + 20 \\\\ \\vspace{1mm}&& 7.04  $^{+0.30}_{-0.96}$ (rms) $^{r}$ &&\\\\ \\vspace{1mm}NGC7469  & {\\em OSSE }, {\\em BeppoSAX}  &    6.81  $^{+0.30}_{-peg}$ (mean) $^{r}$   &  2.14 $^{f}$    & 20 + 20\\\\ \\vspace{1mm}&& 6.88  $^{+0.29}_{-peg}$ (rms) $^{r}$ &&\\\\ \\vspace{1mm}ESO 141-55  & {\\em OSSE }, {\\em BeppoSAX}  &  7.10 $^{+0.57}_{-0.57}$ $^{m}$  &   2.51 $^{ox}$     & 1\\\\ \\vspace{1mm}MCG--2-58-22 & {\\em OSSE } &    7.14 $^{+0.57}_{-0.57}$  $^{m}$  &   3.47 $^{ox}$     & 1\\\\ \\vspace{1mm}NGC6814     & {\\em OSSE }   &    7.08  $^{+0.57}_{-0.57}$ $^{r}$   &   0.05 $^{f}$      & 22 + 2\\\\ \\vspace{1mm}3C120    & {\\em BeppoSAX}   &  7.36 $^{+0.22}_{-0.28} $ (mean)$^{r}$  & 0.65 $^{f}$         & 20 + 20 \\\\ \\vspace{1mm}&&7.48 $^{+0.21}_{-0.27}$ (rms) $^{r}$  &&\\\\ \\vspace{1mm}NGC3516    & {\\em OSSE }    &    7.36   $^{+0.59}_{-0.59}$  $^{r}$  &  0.07 $^{f}$ & 2\\\\ \\vspace{1mm}Mrk279     & {\\em OSSE }   &    7.54 $^{+0.10}_{-0.13}$  $^{r}$  &  0.13 $^{z}$     & 5\\\\ \\vspace{1mm}NGC3227    & {\\em OSSE }    &    7.59 $^{+0.19}_{-peg}$ (mean) $^{r}$  &  0.01 $^{f}$      & 20 + 2\\\\ \\vspace{1mm}&&7.69 $^{+0.18}_{-peg}$ (rms) $^{r}$  &&\\\\ \\vspace{1mm}NGC5548 & {\\em OSSE }, {\\em BeppoSAX}    &   8.09   $^{+0.09}_{-0.07}$ (mean)  $^{r}$  &  0.05 $^{f}$  & 20 + 20\\\\ \\vspace{1mm}&& 7.97 $^{+0.08}_{-0.07}$ (rms) $^{r}$  &&\\\\ \\vspace{1mm}1H 1934-063 & {\\em INTEGRAL}  &  7.86  $^{+0.63}_{-0.63}$  $^{m}$ &  0.07 $^{ox}$    &  1\\\\ \\vspace{1mm}Fairall 9   & {\\em BeppoSAX}   &  7.90 $^{+0.11}_{-0.31}$ (mean) $^{r}$   & 0.16 $^{s}$   &  20 + 20\\\\ \\vspace{1mm}&&7.91 $^{+0.11}_{-0.32}$ (rms) $^{r}$  &&\\\\ \\vspace{1mm}NGC7213    & {\\em OSSE }   &    7.99  $^{+0.64}_{-0.64}$ $^{\\sigma}$   &   0.02 $^{f}$        & 2\\\\ \\vspace{1mm}MCG +8-11-11  & {\\em OSSE } &    8.06  $^{+0.64}_{-0.64}$  $^{m}$   & 0.35 $^{ox}$           & 1\\\\ \\vspace{1mm}NGC526A    & {\\em OSSE }, {\\em BeppoSAX}    &    8.11  $^{+0.65}_{-0.65}$ $^{\\sigma}$   &  0.17 $^{ox}$     &  6 (18) + 1\\\\ \\vspace{1mm}Mrk841     & {\\em OSSE }   &    8.49  $^{+0.68}_{-0.68}$ $^{r}$  &  0.17 $^{f}$           & 22 + 2\\\\ \\vspace{1mm}Mrk509    & {\\em OSSE }, {\\em BeppoSAX}, {\\em INTEGRAL}    &    8.16 $^{+0.04}_{-0.04}$  $^{r}$  &  0.45 $^{ox}$    &  5 + 1\\\\ \\vspace{1mm}III Zw 2    & {\\em OSSE }    &  8.42  $^{+0.67}_{-0.67}$ $^{r}$ &  0.13 $^{z}$   & 3\\\\ \\vspace{1mm}PG 1416-129  & {\\em INTEGRAL}  &    8.75   $^{+0.70}_{-0.70}$ $^{r}$  & 0.12 $^{z}$     & 3\\\\ \\vspace{1mm}3C111       & {\\em INTEGRAL}  &  9.56  $^{+0.76}_{-0.76}$  $^{m}$&  0.01 $^{z}$    & 9\\\\   \\hline  \\end{tabular}  \\end{minipage}  \\end{center} \\caption{Seyfert 1 sub-sample. Name superscripts indicate the hard X-ray satellite that observed the source, o: {\\it OSSE}, i: \\inte ~\\& b: \\sax. Mass superscripts refer to the measurement method, $\\sigma$: stellar velocity dispersion, r: reverberation mapping, m: other. Accretion rate superscripts denote the method by which the bolometric luminosity was determined, x: correction from X-ray, f: integration of the SED flux, ox: correction from O{\\small III}, s: SED modelling, z: other. References, 1: Wang \\etal (2007), 2: Woo \\& Urry (2002), 3: Hao \\etal (2005), 4: McHardy (1988), 5: Vestergaard \\& Peterson. (2006), 6: Garcia-Rissmann \\etal (2005), 7: Bian \\& Gu (2007), 8: Greenhill \\etal (1997), 9: Grandi \\etal (2006), 10: Morales \\& Fabian (2002), 11: Marconi \\etal (2006), 12: Whysong \\& Antonucci (2004),13: Awaki \\etal (2005), 14: Czerny \\etal (2001), 15: Tadhunter \\etal (2003), 16: van Bemmel \\etal (2003), 17: Gu \\etal (2006), 18: Tremaine \\etal (2002), 19: Risaliti \\etal (2005), 20: Kaspi \\etal (2000), 21: Monier \\etal (2001), 22: Laor \\etal (2001) 23: Dadina (2007). Where more than one reference is given, the first refers to the black hole mass and the second to the accretion rate. Where reference 18 is given in brackets, the relation from Tremaine \\etal (2002) has been used to determine the black hole mass via stellar velocity dispersion.}  \\end{table*}  \\begin{table*} \\begin{center} \\begin{minipage}{110mm} \\bigskip \\begin{tabular}{ll|l|c|c} \\hline Source & Instrument & logM/M$_{\\odot}$ &  L/L$_{Edd}$  & Reference\\\\ \\hline  \\vspace{1mm}NGC6300    & {\\em INTEGRAL}    &    5.45  $^{+0.44}_{-0.44}$ $^{m}$   &   0.91 $^{x}$    & 13\\\\ \\vspace{1mm}NGC7314    & {\\em BeppoSAX}   &    6.03  $^{+0.48}_{-0.48}$ $^{\\sigma}$   &   0.41 $^{x}$    & 17 (+18)\\\\ \\vspace{1mm}NGC4945 &  {\\em OSSE }, {\\em INTEGRAL}   &    6.15   $^{+0.49}_{-0.49}$ $^{\\sigma}$  &  0.17 $^{x}$    & 8\\\\ \\vspace{1mm}MCG--5-23-16  &  {\\em OSSE }, {\\em BeppoSAX} &    6.29 $^{+0.50}_{-0.50}$  $^{m}$  &   0.09 $^{ox}$     &  1\\\\ \\vspace{1mm}Circinus     & {\\em INTEGRAL}  &  6.42 $^{+0.51}_{-0.51}$  $^{\\sigma}$ &  0.32 $^{ox}$      & 7\\\\ \\vspace{1mm}NGC5506    &  {\\em OSSE }   &    6.65  $^{+0.53}_{-0.53}$ $^{\\sigma}$   &   2.51 $^{ox}$     & 7\\\\ \\vspace{1mm}NGC4593   & {\\em INTEGRAL}, {\\em BeppoSAX}    &    6.91  $^{+0.55}_{-0.55}$ $^{r}$ &  0.12 $^{f}$    &2\\\\ \\vspace{1mm}ESO 103-G35  & {\\em INTEGRAL}, {\\em BeppoSAX}  &  7.14 $^{+0.57}_{-0.57}$  $^{m}$ &  0.01 $^{x}$     & 14\\\\ \\vspace{1mm}NGC4388    &  {\\em OSSE }, {\\em INTEGRAL}    &    7.22   $^{+0.58}_{-0.58}$ $^{\\sigma}$  & 0.79 $^{ox}$     & 7\\\\ \\vspace{1mm}NGC1068   & {\\em INTEGRAL}   &    7.23  $^{+0.58}_{-0.58}$ $^{\\sigma}$  &  0.44 $^{f}$       & 8 + 2\\\\ \\vspace{1mm}NGC7582   &  {\\em OSSE }    &    7.25   $^{+0.58}_{-0.58}$ $^{\\sigma}$  & 0.35 $^{ox}$      & 7\\\\ \\vspace{1mm}NGC5674    & {\\em BeppoSAX}   &    7.36  $^{+0.59}_{-0.59}$ $^{\\sigma}$   &  0.17 $^{x}$    &  17 (+18)\\\\ \\vspace{1mm}ESO 323-G077  & {\\em INTEGRAL} &  7.39 $^{+0.59}_{-0.59}$ $^{m}$  &  0.28 $^{ox}$    & 1\\\\ \\vspace{1mm}NGC4507   &  {\\em OSSE }, {\\em INTEGRAL}    &    7.58   $^{+0.61}_{-0.61}$ $^{\\sigma}$  &  0.27 $^{ox}$   & 7\\\\ \\vspace{1mm}NGC4258    & {\\em BeppoSAX}   &    7.62   $^{+0.61}_{-0.61}$ $^{\\sigma}$  & 0.01 $^{f}$       & 2\\\\ \\vspace{1mm}NGC7172    &  {\\em OSSE }, {\\em BeppoSAX}   &    7.67  $^{+0.61}_{-0.61}$ $^{\\sigma}$   &   3x10$^{-3}$ $^{ox}$   & 7\\\\ \\vspace{1mm}NGC2992   & {\\em BeppoSAX}   &    7.72  $^{+0.62}_{-0.62}$ $^{\\sigma}$  &  0.01 $^{f}$   &  2\\\\ \\vspace{1mm}NGC5252    & {\\em BeppoSAX}   &    8.04  $^{+0.64}_{-0.64}$  $^{\\sigma}$  &  0.17 $^{s}$   & 2\\\\ \\vspace{1mm}Centaurus A  & {\\em INTEGRAL}  &  8.04 $^{+0.64}_{-0.64}$ $^{\\sigma}$  &  7x10$^{-4}$ $^{s}$   &  11 + 12\\\\ \\vspace{1mm}NGC1365    & {\\em BeppoSAX}  &    8.18  $^{+0.65}_{-0.65}$ $^{m}$  &  2x10$^{-3}$ $^{x}$      & 19\\\\ \\vspace{1mm}NGC2110    &  {\\em OSSE }, {\\em BeppoSAX}    &    8.30  $^{+0.66}_{-0.66}$ $^{\\sigma}$  &  5x10$^{-3}$ $^{f}$    &  2\\\\ \\vspace{1mm}NGC1275    &  {\\em OSSE }, {\\em INTEGRAL}    &    8.51  $^{+0.68}_{-0.68}$ $^{\\sigma}$  &  0.03 $^{f}$   &  2\\\\ \\vspace{1mm}Mrk3       &  {\\em OSSE }, {\\em INTEGRAL}    &    8.65 $^{+0.69}_{-0.69}$  $^{\\sigma}$  &  6x10$^{-3}$ $^{f}$    &  2\\\\ \\vspace{1mm}Cyg A      & {\\em INTEGRAL}   &  9.40 $^{+0.75}_{-0.75}$  $^{\\sigma}$ &  8x10$^{-3}$ $^{x}$    & 15 + 16\\\\  \\hline   \\end{tabular}   \\end{minipage}  \\end{center} \\caption{Seyfert 2 sub-sample.  Superscripts and references as for Table 1.}  \\end{table*} ",
        "conclusions": "The X-ray spectra of stellar mass black hole binaries can be well described by a combination of emission from the accretion disc and a Comptonised tail (see e.g. the reviews by Remillard \\& McClintock 2006; Done, Gierlinski \\& Kubota 2007). The relative importance of the disc and tail can change dramatically, as does the shape of the tail, giving rise to the well known 'spectral states' as a function of (average) $L/L_{Edd}$. In the low/hard state, generally seen at luminosities below a few percent of Eddington, the tail dominates the X-ray emission and its spectral index is fairly well correlated with (average) mass accretion rate, with $\\Gamma\\sim 1.5$ at the lowest luminosities softening to $\\Gamma\\sim 2.1$ just before the major transition to the soft, or disc dominated states. The tail in these states can be very weak, carrying less than 20 per cent of the total bolometric power, when it has $\\Gamma\\sim 2.1$ (disc dominant state). Conversely, the tail can also be much stronger and softer, with $\\Gamma\\sim 2.5$ (very high state or steep power law state).  AGN spectra should show analogous spectral states if the properties of the accretion flow scale simply with black hole mass. There is evidence for this simple scaling, from the radio-X-ray `fundamental plane' relations (Merloni, Heinz, \\& Di Matteo 2003; Falcke, K\\\"ording \\& Markoff 2004) and X-ray variability properties (McHardy et al. 2007; Gierlinski et al. 2007). In this picture (see e.g. Boroson 2002), LINERS, with their hard X-ray spectra and weak UV disc, are the analogues of the low/hard state (but see Maoz 2007), while Seyfert 1 and QSO's at higher $L/L_{Edd}$ correspond to the disc dominated states, with the characteristic strong UV disc emission and consequent broad emission lines, and the Narrow line Seyfert 1's at the highest $L/L_{Edd}$ may be the analogues of the very high state (Pounds, Done \\& Osborne 1995; Middleton, Done \\& Gierli\\'nski 2007).  Observations of type 1 AGN spectra in the 2--10~keV band are consistent with this predicted softening of the intrinsic spectrum with increasing $L/L_{Edd}$ (e.g. Laor 2000), so it seems highly likely that the high energy X--ray spectra considered here should also change with $L/L_{Edd}$. We show that current samples of Seyferts detected at high energies are clearly highly heterogeneous in this parameter (see Fig 1), so the mean spectral index will be determined by a signal--to--noise weighted average of $L/L_{Edd}$. The $L/L_{Edd}$ distributions for the total Seyfert 1 and Seyfert 2 samples presented here are not significantly different (statistically) however the individual OSSE, BeppoSAX and INTEGRAL samples considered here are all consistent with having a higher mean $L/L_{Edd}$ for the Seyfert 1's than for the Seyfert 2s. Thus, despite the limited sample size the apparently steeper spectra seen in the Seyfert 1's at high energy may be be due to an intrinsic softening with increasing $L/L_{Edd}$. Future studies with larger samples of high energy spectra of AGN should be able to test unambiguously whether $L/L_{Edd}$ is the major parameter in determining the shape of the 20--100~keV spectrum in Compton thin Seyferts.  \\label{lastpage}"
    },
    "0710/0710.1059_arXiv.txt": {
        "abstract": "In this paper we show that all supergravity billiards corresponding to $\\sigma$-models on any $\\mathrm{U/H}$ non compact-symmetric space and obtained by compactifying supergravity  to $D=3$ are fully integrable. The key point in establishing the integration algorithm is provided by an upper triangular embedding of the solvable Lie algebra associated with $\\mathrm{U/H}$ into $\\slal(\\mathrm{N},\\mathbb{R})$ which always exists. In this context we establish a remarkable relation between the arrow of time and the properties of the Weyl group. The asymptotic states of the developing Universe are in one-to-one correspondence with the elements of the Weyl group which is a property of the Tits Satake universality classes and not of their single representatives. Furthermore the Weyl group admits a natural ordering in terms of $\\ell_T$, the number of reflections with respect to the simple roots and the direction of time flows is always towards increasing $\\ell_T$, which plays the unexpected role of an entropy. ",
        "introduction": "Notwithstanding its length and its somewhat pedagogical organization, the present one is a research article and not a review. All the presented material is, up to our knowledge, new. Due to the combination of several different mathematical results and techniques necessary  to make our point, which is instead physical in spirit and relevant to basic questions in supergravity and superstring cosmology, we considered it appropriate to choose the present somewhat unconventional format for our paper. After the theoretical statement of our result, we have illustrated it with the detailed study of a few examples. These case-studies were essential for us in order to understand the main point which we have formalized in mathematical terms in part I and we think that they will be similarly essential for the physicist reader. The table of contents helps  the reader to get a comprehensive view of the article and of its structure. \\tableofcontents \\part{Theory: Stating the principles} ",
        "conclusions": "In this paper we have made a few steps forward in developing the general programme of supergravity billiards as a paradigm for superstring cosmology. Our results are both of physical and mathematical nature. \\par On the physical side, which for us means supergravity/superstring theory, the essential points are the following ones: \\begin{description} \\item[1)] We have shown that all supergravity billiards are completely integrable, irrespectively whether they are defined on a maximally split coset manifold $\\mathrm{U/H}$ as it happens in the case of maximal susy or a non maximally split $\\mathrm{U/H}$, as it happens in all lower supersymmetry cases. We have provided the explicit integration algorithm which just depends on the triangular embedding of the solvable Lie algebra $Solv(\\mathrm{U/H})$ into that of $Solv(\\mathrm{SL(N)/SO(N)})$. \\item[2)] We have discovered a new principle of time orientation of the cosmic flow which relies on the natural ordering of the Weyl group elements (or of the permutations) according to their length $\\ell_T$ in terms of transpositions. Cosmic evolution is always in the direction of increasing $\\ell_T$ which plays the role of an entropy. There is a fascinating similarity, in this context between the laws of cosmic evolution and those of black hole thermodynamics. \\item[3)] We have clarified the meaning of Tits Satake universality classes, introduced in \\cite{contoine}, at least from the vantage point of cosmic billiards.  The asymptotic states, the type of available flows and the critical surfaces in parameter space are properties of the class and do not depend on the representative manifold in the class. \\end{description} On the mathematical side the  highlights of our paper are the following ones: \\begin{description} \\item[1)] We have introduced the notion of generalized Weyl group for a non compact symmetric space $\\mathrm{U/H}$ and shown that the factor group with respect to its normal subgroup is just the Weyl group of the Tits Satake subalgebra $\\mathbb{U}_{TS} \\, \\subset \\, \\mathbb{U}.$ Moreover, we have demonstrated that not only the factor group is isomorphic to the Weyl group of the Tits Satake projection but even the generalized Weyl group is also isomorphic $\\mathcal{W}\\left(\\mathbb{U}\\right) \\, \\sim \\, \\mathcal{W}\\left(\\mathbb{U}_{\\mathrm{TS}}\\right)$. At least this is true in the considered examples and we make the conjecture that it is true in general. \\item[2)] We have established a remarkable conjecture encoded in property \\ref{kkminorprinc} of the main text: the constraints on minors of the diagonalizing orthogonal matrix for the Lax operator commute with the Toda flow. \\item[3)] We have proposed a very simple efficient method of calculating the Toda flow asymptotics at $t=\\pm \\infty$ for the Lax operator of a $\\sigma$-model  with target space any non compact-symmetric coset space ${\\mathrm{U}}/{\\mathrm{H}}$. Our algorithm requires only the knowledge of the corresponding Weyl group $\\mathrm{Weyl}(\\mathbb{U})$ as well as that of the small group $\\mathrm{H}$. \\item[4)] We have posed the question how the equations cutting out algebraic loci in compact group or coset manifolds and defined in terms of vanishing minors in the defining representation can be lifted to the abstract group level and extended to all irreducible representations. \\end{description}"
    },
    "0710/0710.1745_arXiv.txt": {
        "abstract": "Using the Green Bank Telescope (GBT) and Pulsar Spigot at 350\\,MHz, we have surveyed the Northern Galactic Plane for pulsars and radio transients.  This survey covers roughly 1000 square degrees of sky within $75^{\\circ} < l < 165^{\\circ}$ and $|b| < 5.5^{\\circ}$, a region of the Galactic Plane inaccessible to both the Parkes and Arecibo multibeam surveys.  The large gain of the GBT along with the high time and frequency resolution provided by the Spigot make this survey more sensitive by factors of about 4 to slow pulsars and more than 10 to millisecond pulsars (MSPs), compared with previous surveys of this area.  In a preliminary, reduced-resolution search of all the survey data, we have discovered 33 new pulsars, almost doubling the number of known pulsars in this part of the Galaxy.  While most of these sources were discovered by normal periodicity searches, 5 of these sources were first identified through single, dispersed bursts.  We discuss the interesting properties of some of these new sources.  Data processing using the data's full-resolution is ongoing, with the goal of uncovering MSPs missed by our first, coarse round of processing. ",
        "introduction": "Here we define the ``Northern Sky'' as the sky above the northern Arecibo declination limit of $38^{\\circ}$.  Previous pulsar surveys of the Northern Sky between $300-500$\\,MHz include an all-sky survey above $\\delta > 38^{\\circ}$ using the Jodrell Bank telescope at 428\\,MHz, which, largely due to RFI, discovered only one pulsar \\citep{nll+95}; a survey at Green Bank using the 140-ft (43-m) equatorially mounted telescope at 370\\,MHz, which found eight new pulsars \\citep{snt97}; and a survey using the Green Bank 300-ft (91-m) transit telescope at 390\\,MHz, which provided partial coverage of the Northern Sky and discovered 20 pulsars \\citep{sstd86}.  The GBT's large gain (2\\,K/Jy) and maneuverability make it the most sensitive radio telescope on Earth over a large portion of the sky. Furthermore, the primary beam width of the GBT at 350\\,MHz is $\\sim 0.6^{\\circ}$, large enough to permit very efficient large-scale, single-pixel surveys of the sky.  This motivated us to commence a 350-MHz survey, which we will refer to as the GBT350 survey, of the Northern Galactic Plane, corresponding to Galactic longitudes $75^{\\circ} < l < 165^{\\circ}$ and out to latitudes $|b| < 5.5^{\\circ}$. Here we present a short summary of the survey design and analysis techniques, and then discuss some of the new sources. ",
        "conclusions": ""
    },
    "0710/0710.3740_arXiv.txt": {
        "abstract": "We use the very large Millennium Simulation of the concordance $\\Lambda$CDM cosmogony to calibrate the bias and error distribution of Timing Argument estimators of the masses of the Local Group and of the Milky Way. From a large number of isolated spiral-spiral pairs similar to the Milky Way/Andromeda system, we find the interquartile range of the ratio of timing mass to true mass to be a factor of $1.8$, while the 5\\% and 95\\% points of the distribution of this ratio are separated by a factor of $5.7$. Here we define true mass as the sum of the ``virial'' masses $M_{200}$ of the two dominant galaxies. For current best values of the distance and approach velocity of Andromeda this leads to a median likelihood estimate of the true mass of the Local Group of $5.27\\times 10^{12}\\msun$, or $\\log M_{LG}/M_\\odot = 12.72$, with an interquartile range of $[12.58, 12.83]$ and a 5\\% to 95\\% range of $[12.26, 13.01]$. Thus a 95\\% lower confidence limit on the true mass of the Local Group is $1.81\\times 10^{12}\\msun$. A timing estimate of the Milky Way's mass based on the large recession velocity observed for the distant satellite Leo I works equally well, although with larger systematic uncertainties. It gives an estimated virial mass for the Milky Way of $2.43 \\times 10^{12}\\msun$ with a 95\\% lower confidence limit of $0.80 \\times 10^{12}\\msun$. ",
        "introduction": "\\label{intro_section}  During the 1970's it became generally accepted that most, perhaps all, galaxies are surrounded by extended distributions of dark matter, so-called dark halos \\citep{eks74,opy74}. These were soon understood to play an essential role in driving the formation and clustering of galaxies \\citep{wr78}. With the introduction of the Cold Dark Matter (CDM) paradigm, these ideas took more concrete form, allowing quantitative predictions to be made both for the population properties \\citep{blumenthal84} and for the large-scale clustering \\citep{davis85} of galaxies.  Measurements of the fluctuation spectrum of the Cosmic Microwave Background \\citep{smoot92,spergel03} and of the apparent acceleration of the cosmic expansion \\citep{riess98,perlmutter99} elevated the CDM model, in its variant with a cosmological constant ($\\Lambda$CDM), to the status of a standard paradigm. At the same time improving numerical techniques and faster computers have enabled detailed simulation of the formation and evolution of the galaxy population within this paradigm throughout a significant fraction of the observable Universe \\citep{mr_nature}. Nevertheless, direct observational evidence for halos as extended as the paradigm predicts around galaxies like our own has so far come only from statistical analyses of the dynamics of satellite galaxies \\citep[e.g.][]{zaritsky97, prada03} and of the gravitational lensing of background galaxies \\citep[e.g.][]{seljak02, mandelbaum06} based on large samples of field spirals.  The earliest observational indication that the effective mass of the Milky Way must be much larger than its stellar mass came from the Timing Argument (hereafter TA) of \\citet{kw59}. These authors noted that the Local Group is dominated by the two big spirals, and that these are currently approaching each other at about $100\\kms$. (The next most luminous galaxy is M33 which is probably about a factor of $10$ less massive than M31 and the Galaxy.) This reversal of the overall cosmic expansion must have been generated by gravitational forces, and since the distance to the nearest external bright galaxy is much greater than that between M31 and the Milky Way, these forces are presumably dominated by material associated with the two spirals themselves.  \\citeauthor{kw59} set up a simple model to analyse this situation -- two point masses on a radial orbit.  These were at pericentre (i.e. at zero separation) at the Big Bang and must have passed through apocentre at least once in order to be approaching today. Clearly this requires an apocentric separation larger than the current separation and an orbital period less than twice the current age of the Universe. Together these requirements put a lower limit on on the total mass of the pair. A more precise estimate of the minimum possible mass is obtained from the parametric form of Kepler's laws for a zero angular momentum orbit: \\begin{equation} r=a(1-\\cos\\chi) \\end{equation}  \\begin{equation} t=\\bigg(\\frac{a^{3}}{GM}\\bigg)^{1/2}(\\chi-\\sin\\chi) \\end{equation}  \\begin{equation} \\frac{dr}{dt}=\\sqrt{\\frac{GM}{a}}\\frac{\\sin\\chi}{1-\\cos\\chi} \\end{equation} where $r$ is the current separation, $dr/dt$ is the current relative velocity, $a$ is the semi-major axis, $\\chi$ is the eccentric anomaly, $t$ is the time since the Big Bang (the age of the universe) and $M$ is total mass \\citep{lb81}. Given observationally determined values for $r$, $dr/dt$ and $t$, these equations have an infinite set of discrete solutions for $\\chi$, $a$ and $M$ labelled by the number of apocentric passages since the Big Bang. The solution corresponding to a single apocentric passage gives the smallest (and only plausible) estimate for the mass, which is about $5\\times 10^{12}\\msun$ for current estimates of $r$, $dr/dt$ and $t$. Note that this is still a lower limit on the total mass, even within the simple point-mass binary model, since any non-radial motions in the system would increase its present kinetic energy and so increase the mass required to reverse the initial expansion and bring the pair to their observed separation by the present day \\citep[see][]{el82}.  As \\citeauthor{kw59} realised, this timing estimate of the total mass of the Local group exceeds by more than an order of magnitude the mass within the visible regions of the galaxies, as estimated from their internal dynamics, in particular, from their rotation curves. Thus 90\\% of the mass must lie outside the visible galaxies and be associated with little or no detectable light. Modern structure formation theories like $\\Lambda$CDM predict this mass to be in extended dark matter halos with $M(r)$ increasing very roughly as $r$ out to the point where the halos of the two galaxies meet. Such structures have no well-defined edge, so any definition of their total mass is necessarily somewhat arbitrary. In addition, their dynamical evolution from the Big Bang until the present is substantially more complex than that of a point-mass binary. Thus the mass value returned by the Timing Argument cannot be interpreted without some calibration against consistent dynamical models with extended dark halos.  A first calibration of this type was carried out by \\citet{kc91} using simulations of an Einstein-de Sitter CDM cosmogony.  Here we use the very much larger Millennium Simulation \\citep{mr_nature} to obtain a more refined calibration based on a large ensemble of galaxy pairs with observable properties similar to those of the Local Group. We find that the standard timing estimate is, in fact, an almost unbiased estimate of the sum of the conventionally defined virial masses of the two large galaxies.  \\citet{zaritsky89} attempted to measure the halo mass of the Milky Way alone by measuring radial velocities for its dwarf satellites and assuming the population to be in dynamical equilibrium in the halo potential. They noted, however, that one of the most distant dwarfs, Leo I, has a very large recession velocity and as a result provides a interesting lower limit on the Milky Way's mass by a variant of the original Timing Argument. To reach its present position and radial velocity, Leo I must have passed pericentre at least once since the Big Bang and now be receding from the Galaxy for (at least) the second time.  Applying the point-mass radial orbit Equations (1) -- (3) to this case gives a lower bound of about $1.6\\times 10^{12}\\msun$.  This seems likely to be a significant underestimate, since Leo I could not have passed through the centre of the Milky Way without being tidally destroyed so its orbit cannot be purely radial. Below we calibrate the Timing Argument for this case also, finding it to work well although with significantly more scatter than for the Local Group as a whole. This is because the $\\Lambda$CDM paradigm predicts that the dynamics on the scale of Leo I's orbit ($\\sim 200\\kpc$) is typically more complex than on the scale of the Local Group as a whole ($\\sim 700\\kpc$).  Our paper is organised as follows.  In Section~\\ref{data_section}, we briefly describe the Millennium Simulation and the selection criteria we use to define various samples of `Local Group-like' pairs and of `Milky Way-like' halos. In Section~\\ref{result_section}, we plot the ratio of true total mass to Timing Argument mass estimate for these samples, and we use its distribution to define an unbiased TA estimator of true mass with its associated confidence ranges. In Section~\\ref{app_LG_section} this is then applied to the Local Group in order to obtain an estimate its true mass with realistic uncertainties.  Section~\\ref{app_MW_section} attempts to carry out a similar calibration for the TA-based estimate of the Galaxy's halo mass from the orbit of Leo I. We conclude in Section~\\ref{conclusions} with a summary and brief discussion of our results. ",
        "conclusions": "\\label{conclusions} The statistical argument underlying the analysis of this paper is more subtle than it may at first appear, so it is worth restating it somewhat more formally in order to understand what is being assumed in deriving the mass estimates for the Local Group and for the Milky Way given above.  We believe that the mass distributions around galaxies are much more extended than the visible stellar distributions, and that these have been assembled from near-uniform ``initial'' conditions in a manner at least qualitatively resembling that in a $\\Lambda$CDM universe. Thus the assembly histories of the Local Group and of the Milky Way's halo differ in major ways from those assumed by the original Timing Arguments of \\citet{kw59} and \\citet{zaritsky89}. In addition, the meaning of the derived mass values needs clarification. We wish to use the Millennium Simulation to calibrate the TA estimates against conventional measures of halo mass, and to test the general applicability of the Timing Argument. However, we want to do this in a way which avoids any significant dependence on the details of the $\\Lambda$CDM model, for example, on the exact density profiles, abundances and substructure properties which it predicts for halos.  Our method uses the simulation to estimate the distribution of the ratio of ``true'' mass to TA mass estimate for samples of objects whose properties ``resemble'' those of the observed Local Group and Milky Way -- Leo I systems. Our restrictions on separation and radial velocity implement this similarity requirement in a straightforward way, but our constraints on $V_{max}$ have a more complex effect. Although the true $V_{max}$ values for M31 and the Milky Way are very likely within our looser range ($150\\kms \\le V_{max} < 300\\kms$) the simulation exhibits a tight correlation between $V_{max}$ and $M_{200}$ . Imposing fixed limits on $V_{max}$ is thus effectively equivalent to choosing a fixed range of $M_{200}$. As a result, we are in practice estimating the distribution of $A_{200}$ or $B_{200}$ for systems of given {\\it true} mass, subject to the assumed constraints on separation and radial velocity. However, when we apply our results to estimate true masses for the Local Group and the Milky Way, we implicitly assume that our distributions of $A_{200}$ and $B_{200}$ are appropriate for samples of given TA mass estimate, again subject to our constraints on separation and radial velocity. It is thus important to understand when these two distributions can be considered the same.  The relation can be clarified as follows. From the simulation we compile the distribution of $M_{tr}/M_{TA}$, or equivalently of $\\Delta \\equiv \\ln M_{tr} - \\ln M_{TA}$, for systems with $\\ln M_{tr}$ in a given range.  We then implicitly assume that this distribution does not depend on $M_{tr}$, at least over this range, so that the result can be taken as an estimate of the probability density of $\\Delta$ at given $M_{tr}$. Bayes Theorem then gives us the probability density function (pdf) for $\\Delta$ at fixed $M_{TA}$: \\begin{eqnarray} f[\\Delta | \\ln M_{TA}] = \\frac{f[\\Delta, \\ln M_{TA}]}{f[\\ln M_{TA}]}\\nonumber\\\\ = \\frac{f[\\Delta, \\ln M_{tr}]}{f[\\ln M_{TA}]} \\nonumber\\\\ = \\frac{f[\\Delta | \\ln M_{tr}]\\, f[\\ln M_{tr}]}{ f[\\ln M_{TA}]}\\nonumber\\\\ = f[\\Delta | \\ln M_{tr}] \\label{Bayes} \\end{eqnarray} The first line here simply writes the conditional pdf of $\\Delta$ at given $M_{TA}$ in terms of the joint pdf of the two quantities and the pdf of $M_{TA}$. The second line then rewrites the joint pdf in terms of the equivalent variables $\\Delta$ and $M_{tr}$, using the fact that the Jacobian of the transformation is unity. The third line re-expresses the joint pdf as the product of the pdf of $\\Delta$ at given $M_{tr}$ times the pdf of $M_{tr}$. The final line then follows from the normalisation condition, {\\it provided} that $f[\\ln M_{tr}]$ is constant and $f[\\Delta | \\ln M_{tr}]$ is independent of $M_{tr}$. Thus, when estimating $M_{tr}$ from $M_{TA}$, we assume a uniform prior on $\\ln M_{tr}$ and that the distribution of $\\Delta$ does not depend on true mass. Both these assumptions appear natural and appropriate.  The analysis underlying the Timing Argument (Equations (1) -- (3)) assumes that the relative orbit of the two objects is bound and has conserved energy since the Big Bang. Recently, \\citet{sales07} have shown that in $\\Lambda$CDM models this assumption is significantly violated for a substantial number of satellites within halos comparable to that of the Milky Way. In particular, they demonstrate the presence of a tail of unbound objects which are being ejected from halos as a result of $3$-body ``slingshot'' effects during their first pericentric passage. These objects are typically receding rapidly from their ``Milky Way'', as assumed in the \\citet{zaritsky89} argument, but they violate its assumption that the present orbital energy can be used to infer the period of the initial orbit (i.e. the time from the Big Bang to first pericentric passage).  Clearly such objects should also be present in the Millennium Simulation, although lack of resolution might make them under-represented in comparison to the simulations analysed by \\citet{sales07}. Thus our analysis takes the possibility of such ejected satellites into account, at least in principle. Objects of this type will show up as systems with anomalously large TA estimates for their halo mass, and Fig.~\\ref{mw_mass_scatter} shows a number of outliers which could well be explained in this way. Issues of this kind do not effect TA-based estimates of the mass of the Local Group since the two big galaxies are currently approaching for the first time.  The only kinematic information about the relative orbit of M31 and the Milky Way used in our analysis is their current approach velocity. \\citet{vdmg07} show that geometric arguments can already constrain the transverse component also, and future astrometry missions such as SIM might allow $V_{tr}$ to be measured directly. Thus it is interesting to ask if our TA mass estimate could be significantly refined by measuring the full 3-D relative motion of the two galaxies, rather than just its radial component. We address this in Fig.~\\ref{a200_vtr_scatter} which plots $A_{200}$, the ratio of true mass to TA estimate, against $V_{tr}$ for a sample of Local Group analogues with our preferred morphology, isolation and radial velocity cuts, and with $150\\kms \\le V_{max} < 300\\kms$. The median $V_{tr}$ for this sample is $86 \\kms$, comparable to the \\citet{vdmg07} estimate for the real system. There is no apparent correlation of $A_{200}$ with $V_{tr}$, and indeed, the medians of $A_{200}$ for the high and low $V_{tr}$ halves of the sample are both close to 1 and do not differ significantly.  Pairs with high $V_{tr}$ do show larger {\\it scatter} in $A_{200}$ than pairs on near-radial orbits. For given separation, radial velocity and age, the Kepler model implies a mass which increases monotonically with $V_{tr}$. The absence of a detectable trend in Fig.~\\ref{a200_vtr_scatter} shows that uncertainties in $V_{tr}$ do not dominate the scatter in our TA mass estimate, and that a measurement of $V_{tr}$ will not substantially increase the precision with which the true mass can be measured.  \\begin{figure} \\includegraphics[width=0.5\\textwidth]{a200_vtr.ps} \\caption{Scatter plot of $A_{200}$ versus transverse velocity for Local Group analogues in our sample with preferred morphology, isolation and radial velocity cuts and with $150\\kms \\le V_{max} < 300\\kms$.  The vertical dashed line indicates the median $V_{tr}$.  The distributions on either side of this line are each further split in half at the median values of $V_{tr}$ (the solid horizontal lines.) There is essentially no correlation in this plot, indicating that a measurement of the transverse velocity will not significantly improve the TA mass estimate.} \\label{a200_vtr_scatter} \\end{figure}  In conclusion, our analysis shows the Timing Argument to produce robust estimates of true mass both for the Local Group and for the Milky Way, as long as ``true mass'' is understood to mean the sum of the conventional masses of the major halos. For the Local Group as a whole, the estimate and confidence limits given in Section~\\ref{app_LG_section} and in the Abstract appear reliable given the excellent statistics provided by the Millennium Simulation, the lack of any substantial dependence on our isolation and morphology cuts, and the relatively simple dynamical situation. Although the results based on Leo I's orbit also appear statistically sound, the more complex dynamical situation offers greater scope for uncertainty, particularly when trying to place a lower limit on the mass of the Milky Way's halo.  On the other hand, our best estimate of this mass is just under half of our estimate of the sum of the halo masses of M31 and the Galaxy. Thus the picture presented by the data appears quite consistent, and gives no reason to be suspicious of the Milky Way results."
    },
    "0710/0710.2436_arXiv.txt": {
        "abstract": "The luminosities of the optical afterglows of Gamma Ray Bursts, 12 hours (rest frame time) after the trigger, show a surprising clustering, with a minority of events being at a significant smaller luminosity.  If real, this dichotomy would be a crucial clue to understand the nature of optically dark afterglows, i.e. bursts that are detected in the X--ray band, but not in the optical.  We investigate this issue by studying bursts of the pre--{\\it Swift} era, both detected and undetected in the optical.  The limiting magnitudes of the undetected ones are used to construct the probability that a generic bursts is observed down to a given magnitude limit.  Then, by simulating a large number of bursts with pre--assigned characteristics, we can compare the properties of the observed optical luminosity distribution with the simulated one.  Our results suggest that the hints of bimodality present in the observed distribution reflects a real bimodality: either the optical luminosity distributions of bursts is intrinsically bimodal, or there exists a population of bursts with a quite significant grey absorption, i.e. wavelength independent extinction.  This population of intrinsically weak or grey--absorbed events can be associated to dark bursts. ",
        "introduction": "Since the detection of the first long Gamma Ray Burst (GRB) optical afterglow (Van Paradijs et al. 1997), the non--detection of any optical source in the direction of the gamma--ray trigger of some events stimulated the interest about the possible differences between the nature of the afterglow emission of the optically bright and faint GRBs.  During the last 9 years the increasing number of optical detections and spectroscopic redshift determinations allowed us to study the intrinsic features of the optical afterglow emission of long GRBs.  Despite the improved (i.e. made more promptly) observations of optical afterglows, in almost half of the observed long GRBs no optical counterpart is still found.  These events have been called in the literature as Dark Burst, or Failed Optical Afterglows GRBs (Lazzati et al. 2002).   In Nardini et al. (2006a) \\footnote{Liang \\& Zhang (2006) independently found similar results.}  we showed the optical $R$ band luminosity light curves of a sample of 24 pre--{\\it Swift} GRBs with known spectroscopic redshift and published estimate of host galaxy dust absorption.  We found a strong clustering of that luminosities for the GRBs in our sample.  Most (i.e. 21/24) of the $R$ band luminosities at 12 hours in the source frame are clustered within a log--normal distribution centred around a mean value $\\log L_{\\nu_R}=30.65$ [erg s$^{-1}$ Hz$^{-1}$] with a dispersion $\\sigma=0.28$.  We also found 3 GRBs showing dimmer luminosities, a factor 15 (from 3.6 to 4.6 $\\sigma$) smaller than the mean of the higher luminosity distribution.  No GRB was found in the luminosity range between these two ``families''.  In Fig \\ref{istobs} we show the histogram of the $R$ band luminosities of our sample of GRBs.  In a recent update (Nardini et al. 2006b) we added 8 new GRBs detected by the {\\it Swift} satellite (Gehrels et al. 2004) for whose an estimate of the host galaxy dust extinction has been published.  This small sample of {\\it Swift} GRBs confirms both the clustering and the bimodality of the optical luminosities found by us with pre--{\\it Swift} bursts.  We also evaluated the optical luminosities for 17 other {\\it Swift} GRBs without any published $A_V^{host}$ estimate, and found that they are consistent with our previous findings.  \\begin{figure} \\vskip -0.5 true cm \\centerline{\\psfig{figure=istolumvecchi.eps,angle=0,width=10cm}} \\vskip -0.8 true cm \\caption{ Histogram of the monochromatic optical luminosities 12 hours (rest frame) after the trigger for the 24 GRBs analysed in Nardini et al. (2006a).  Data have been de-reddened both for Galactic and host extinction.  } \\label{istobs} \\end{figure}  The discovery of a family of optically dim GRBs is an important clue for the understanding of the nature of dark bursts. The few underluminous observed events could be the tip of the iceberg of a population of GRBs which are intrinsically less luminous.  Therefore, a fraction of dark GRBs could belong to this family whose distance, optical absorption, or observing conditions do not allow any optical detection. Because of its potential importance for the understanding of dark bursts, we investigate, in this paper, if the observed bimodal luminosity distribution of optically--bright bursts is due to any selection effect (related to the search/detection of GRB optical counterparts) or if it reflects the existence of two GRB populations.  To this aim we simulate through a Montecarlo method a sample of GRBs with a redshift distribution traced by the cosmic star formation rate (Porciani \\& Madau 2001), assuming different shapes of their intrinsic optical luminosity function.  We also simulate different values of dust absorption within the host galaxy to all the simulated events.  This value is calculated assuming the standard extinction curves (Pei 1992).   We infer a limiting magnitude distribution obtained by the analysis of the deepest $R$ band upper limits of all the pre--{\\it Swift} GRBs with no detection of their optical afterglow.  This is the key point of our study: the use of the upper limits on the optical flux to construct the probability that a simulated bursts would be detected or not.  It is this probability distribution that allows us to perform, meaningfully, our simulations.  We then compare the resulting luminosity distribution of the detectable simulated events with the one obtained in Nardini et al. (2006a) and shown in Fig \\ref{istobs}.  The scope of our simulation is to check if for any conceivable combination of the input assumptions (i.e. luminosity function, extinction and redshift distribution) we can reproduce a simulated sample whose $R$--band luminosity distribution (12h rest frame) is consistent with that observed with the sample of 24 pre--{\\it Swift} GRBs.   For our analysis we use only the pre--{\\it Swift} GRBs because they represent an homogeneous sample: their optical light curves were sampled from hours to days after the trigger and the estimates of the host galaxy extinction have been published. The same study can be repeated using the more recent {\\it Swift} GRBs once we will have a sufficient number of events with published estimates of the host galaxy extinction (now this number is too small, see Nardini et al.  2006b). ",
        "conclusions": "In this study we analysed the possible importance of observational selection effects on the clustering and the bimodality found in the long GRBs optical afterglow luminosity distribution.  \\begin{itemize} \\item We studied the $R$ band upper limits of all the pre--{\\it Swift} dark GRBs and we showed that they are consistent with the ``typical'' afterglow detections.  The distribution of these upper limits extrapolated at 12 h after trigger enables the definition of a Telescope Selection Function, which can define the probability, for any burst, to be pointed with a telescope and exposure time corresponding to some limiting magnitude. \\item Through Montecarlo simulations, we have studied which combinations of optical luminosity functions and absorption in the host galaxy can be consistent with the observations.  In doing so, we have found that the gap between the two observed luminosity distribution is real and it is not the result of a small number statistics. \\item If the absorption is chromatic (i.e. a ``standard\" one) no unimodal intrinsic luminosity distribution agrees with the observed one.  If a good fraction of events (but not all) is absorbed by ``grey\" (i.e. achromatic) dust, then a unimodal luminosity distribution is possible.  An achromatic absorption would not be recognised through the standard $A_V^{host}$ estimate methods.  This can lead to underestimate the intrinsic luminosity for a fraction of bursts that could thus appear as members of an underluminous family. \\item Dark bursts could then be associated either to an intrinsically optically underluminous family, or to those bursts being characterised by a relatively large achromatic absorption.  In the first case one should explain why, optically, the luminosity distribution is bimodal, while other properties are not (e.g. the distribution of the energetics of the prompt emission), while in the second case one should explain why a fraction of bursts live in a different environment, characterised by grey dust. \\end{itemize}"
    },
    "0710/0710.0096_arXiv.txt": {
        "abstract": "We have computed the size distribution of silicate grains in the outer radiative region of the envelope of a protoplanet evolving according to the scenario of \\citet{pollack96}.  Our computation includes grain growth due to Brownian motion and overtake of smaller grains by larger ones. We also include the input of new grains due to the breakup of planetesimals in the atmosphere.  We follow the procedure of \\citet{podolak03}, but have speeded it up significantly.  This allows us to test the sensitivity of the code to various parameters.  We have also made a more careful estimate of the resulting grain opacity.  We find that the grain opacity is of the order of $10^{-2}\\unit{cm^2~g^{-1}}$ throughout most of the outer radiative zone as \\citet{hubickyj} assumed for their low opacity case, but near the outer edge of the envelope, the opacity can increase to $\\about{1}\\unit{cm^2~g^{-1}}$.  We discuss the effect of this on the evolution of the models. ",
        "introduction": "The planets in our solar system were formed from material left over after the formation of the Sun, but the details of the formation process are still under investigation. It is believed that the gas giant planets of our solar system, Jupiter and Saturn, have a solid core embedded within a gaseous envelope \\citep{saumon04}. The most popular theory for the formation of these planets supposes that they were formed in two stages: A solid core was first formed by accretion of planetesimals in the protoplanetary disk, and, when the core became massive enough to gravitationally attract and capture gas from the disk, the gaseous envelope was added.  Detailed numerical simulations of these formation stages have been performed for this core accretion model \\citep{pollack96,hubickyj}.  These simulations show three distinct phases. During the first phase the planet still consists primarily of solid material, the gas accretion rate is small, and the planetesimal accretion rate increases rapidly until the planet's feeding zone is depleted. The planet then enters a second phase during which the rates of solid and gas accretion are both small and nearly constant. When the gas mass of the planet is about equal to its solid mass the third phase begins, characterized by runaway accretion of gas.  Simulations of this kind, using reasonable parameters, have been successful in producing planets with properties similar to those of the actual giant planets, however the formation time for a Jupiter-like planet at $5\\unit{AU}$ from a Sun-like star is uncomfortably close to the estimated lifetime of the protoplanetary disk in which it is to be formed. Observations of disks around young stars give estimates of 0.1--10~{Myr} for their lifetimes \\citep{haisch01}, while numerical simulations are able to produce a planet in times varying from a few to many millions of years, depending on several key parameters \\citep{hubickyj}.  One of these key parameters is the opacity of the gaseous envelope during the second of the phases described above. In particular, the contribution to the opacity from aerosols is an important factor whose value is uncertain. Grain opacity is difficult to estimate, because it depends on the size distribution as well as the composition of the aerosols.  It has been demonstrated \\citep{hubickyj} that arbitrarily lowering the opacity in the outer radiative zone by a factor of 50 can reduce the formation time by roughly a factor of 3. \\citet{podolak03} has shown that the size distribution of grains in the envelope can be quite different from the interstellar size distribution usually assumed in opacity calculations, and that this can result in a significantly lowered opacity.  The purpose of this work is to make this conclusion more robust by repeating the calculation with greater accuracy, and by exploring further the parameter space. We attempt to identify under which circumstances it is correct to assume that the grain opacity, in the context of the core accretion model, is much lower than the interstellar value, and to suggest the value that should be used instead.  We use the same microphysics model that was described in \\citep{podolak03}, but we have made the code much more efficient so that we can more readily test the sensitivity to various assumptions.  We have also done a more careful computation of the actual opacity of the computed distribution.  In section~\\ref{sec:model} we review the details of the microphysical model.  In section~\\ref{sec:results} we present the results of our calculations and some tests of the sensitivity of the model to the various input parameters, and in section~\\ref{sec:conclusion} we give our conclusions. ",
        "conclusions": "Conclusions} Before any conclusions can be drawn from the above results about the effect of grain growth on planet formation it is important to note again one important assumption of the calculations made here, the assumption of a static atmosphere. As was previously mentioned, the radius-density-temperature relation of the atmosphere remains constant throughout the calculation. This assumption was justified because the time to reach a steady state of the grain size distribution is much shorter than the time scale of the atmosphere's evolution. We must remember however that the structure of the atmosphere which is used in the present calculation was determined \\emph{without} knowledge of the grain size distribution, and with previously chosen grain opacity. The grain opacity itself will, of course, influence the evolution of the planet's atmosphere. In this sense, the calculations carried out here are not completely self consistent. It is possible that using the opacity profiles shown above in a planet evolution model would result in different atmosphere structures which, if used as background for the grain growth process, would result in different opacity profiles.  With this reservation in mind, there is at least one conclusion that can still be stated with some confidence about the effects of grain growth on the atmosphere's opacity, and that is that grain growth \\emph{does affect} the atmosphere's opacity. More precisely: The conditions during phase 2 of core accretion are such that growth and settling are efficient processes by which grains are removed from the radiative zone of the planet's atmosphere, reducing their contribution to the opacity. For most of the radiative zone the resulting grain opacity is significantly lower than it would be if grain growth were not accounted for. The implication for giant planet formation theory would be that the formation time of a Jupiter-like planet through the core accretion scenario may be shortened, alleviating one of the major points of concern about this model.  As for the actual value of grain opacity that should be used in planet evolution models, our simulations indicate that the opacity is not reduced by a constant factor throughout the atmosphere.  Instead, it is roughly equal to the interstellar grain value in the uppermost layers, where the grains reside for a relatively short time, and grain growth is not an important process.  Deeper in the atmosphere, where the grain microphysics has time to adjust the grain size distribution, the opacity can be reduced over the interstellar value by as much as three orders of magnitude.  In the uppermost parts of the atmosphere, where the optical depth is $\\tau\\leq\\sim 10$ the optical depth profile is close to that computed for the high opacity case of \\citet{hubickyj}, but as you descend into the deeper layers of the radiative zone, the optical depth remains well below the profile derived for the high opacity case.  Except for the deepest region of the radiative zone, the optical depth in our models is always higher than that derived for the low opacity case of \\citet{hubickyj}.  It seems that we must combine a microphysics calculation of the grain distribution together with an evolutionary model of the sort presented in \\citet{hubickyj} in order to determine the true effect of the grain size distribution on the contraction time of the protoplanetary envelope.  In the present work, evaporation of grains was completely neglected. If there are ice grains present, they will evaporate once the temperature exceeds $\\about 200$ K. Organic material in the grains will either evaporate or pyrolyze in the hydrogen-rich atmosphere. Ice and organics are expected to make up about $2/3$ of the high-Z material by mass, so the grain abundance throughout most of the region we have considered will be only about 30\\% of what we have considered for the baseline model.  As we showed above, such a reduction in the rate of grain accretion has only a small effect on the resulting distribution.  However, it must be remembered that in the upper layers, where the temperatures are low enough, grains can accumulate ice and organics.  These components will be lost in the deeper layers.  If the ice acts as a ``glue'' to hold the silicates together, then we would expect the larger grains to break up upon reaching the deeper, hotter layers.  At the very least the grains will develop voids and possibly a fractal structure.  This will lower their mean density and have a substantial effect on the size distribution.  Thus our conclusions about the opacity must be seen as tentative, subject to future work."
    },
    "0710/0710.4136_arXiv.txt": {
        "abstract": "The extragalactic cosmic gamma-ray background (CGB) is an interesting channel to look for signatures of dark matter annihilation. In particular, besides the imprint in the energy spectrum, peculiar anisotropy patterns are expected compared to the case of a pure astrophysical origin of the CGB. We take into account the uncertainties in the dark matter clustering properties on sub-galactic scales, deriving two possible anisotropy scenarios. A clear dark matter angular signature is achieved when the annihilation signal receives only a \\emph{moderate} contribution from sub-galactic clumps and/or cuspy haloes. Experimentally, if galactic foregrounds systematics are efficiently kept under control, the angular differences are detectable with the forthcoming GLAST observatory, provided that the annihilation signal contributes to the CGB for a fraction $\\agt$10-20\\%. If, instead,  sub-galactic structures have a more prominent role, the astrophysical and dark matter anisotropies become degenerate, correspondingly diluting the DM signature. As complementary observables we also introduce the cross-correlation between surveys of galaxies and the CGB and the cross-correlation between different energy bands of the CGB and we find that they provide a further sensitive tool to detect the dark matter angular signatures. ",
        "introduction": "Astronomical and cosmological observations provide overwhelming evidence for the presence of dark matter (DM) (see e.g.\\ \\cite{Bertone:2004pz} for a review). In particular, the combination of various cosmological data sets provides a precise measurement of the amount of DM in the universe: $\\Omega_c h^2 \\simeq 0.11$ with a $2 \\sigma$ precision of $\\sim$ 5\\% in the minimal $\\Lambda$CDM model \\cite{Tegmark:2006az,Spergel:2006hy,Seljak:2006bg} and $\\sim$ 20\\% in more extended models \\cite{Hamann:2006pf}.  However, despite the noticeable sensitivity to the cosmological abundance of matter (either dark or baryonic), such measurements only weakly constrain the properties and nature of the particle associated to DM, and very weak limits are available on the DM particle mass $m_\\chi$ and on its couplings. The simplest DM candidate is the Weakly Interacting Massive Particle (WIMP) which is characterized by having been in thermal equilibrium in the early universe (as opposed to for example the sterile neutrino or super-heavy DM), and having decoupled from equilibrium while non-relativistic. In order to get the correct DM abundance the mass of such a particle cannot be larger than $\\sim 30$ TeV \\cite{Bertone:2004pz,Griest:1989wd}. On the other hand, collider experiments provide a lower bound on the mass of $\\sim 50-100$ GeV \\cite{Bertone:2004pz}, depending on the specific particle candidate. Mass of $\\mathcal{O}$(GeV) are however possible if more exotic candidates are considered \\cite{Gunion:2005rw}. The typical thermally averaged DM annihilation cross section in the WIMP scenario is $ <\\!\\!\\! \\sigma_\\chi v \\!\\!\\! >$$\\sim 10 ^{-26}$ cm$^{3}$s$^{-1}$ \\cite{Bertone:2004pz}. However, we stress that if the DM is produced out of equilibrium in the early universe, no bounds can be given and super-massive, GUT scale, DM particles ($m_\\chi \\sim 10^{15}$ GeV) and cross sections $ <\\!\\!\\! \\sigma_\\chi v \\!\\!\\! >$ $ \\ll 10^{-26}$ cm$^{3}$s$^{-1}$ are in principle possible.  From the point of view of particle physics WIMP candidates are very appealing and emerge naturally in Supersymmetric (SUSY) extensions of the standard model or in the Universal Extra Dimensions (UED) model \\cite{Bertone:2004pz}. The sensitivity of accelerator experiments, notably the Large Hadron Collider, and of direct search experiments are approaching the levels required to test the WIMP hypothesis, and a direct discovery of DM WIMPs could happen in the not so distant future.   DM WIMP candidates have thus typically a large annihilation cross section and  pair-annihilate into standard model particles that subsequently decay and shower producing large numbers of photons and neutrinos. Such $\\gamma$-rays  from DM annihilation constitute an ideal target for astronomical searches. Thus, astrophysical and cosmological observations can provide a crucial test, complementary to a direct laboratory detection, in the search for the nature of DM particles. Various astrophysical environments have been discussed in detail as promising sites for observation of DM annihilation, among others the galactic center, satellite dwarf galaxies of the Milky Way and clumps of DM in the Milky Way halo. In the following we will focus instead on the all-sky diffuse signal expected in the extragalactic cosmic gamma-ray background (CGB) \\cite{Bergstrom:2001jj,Ullio:2002pj,Elsaesser:2004ck,Elsaesser:2004ap}.  Peculiar spectral and angular features can help in disentangling a signal produced by DM from emission by ``ordinary'' astrophysical sources. The spectrum of photons from DM annihilation is in general harder than the spectra arising from normal astrophysical processes and exhibit a pronounced cutoff at an energy near  $m_\\chi$ \\cite{Bergstrom:2001jj,Ullio:2002pj}. The resulting emission thus appear like a ``bump'' in the background astrophysical energy spectrum in the energy range in which the DM signal gives a relevant contribution. However, although this kind of signature would constitute a strong hint of DM annihilation, astrophysical processes that could mimic such behavior are possible.  Another signature, which has been widely studied, is direct annihilation into a state containing photons, resulting in a line in the background spectrum that would constitute a ``smoking gun'' signature of DM. However, by construction this process is necessarily loop suppressed and in most models the flux is quite small (see, however, \\cite{Ullio:2002pj} for a more thorough discussion).  Peculiar angular signatures thus offer a complementary signature to exclude the remaining degenerate astrophysical interpretations of a signal.  An example is the clumpiness of DM at sub-galactic scales \\cite{Berezinsky:2005py,Taylor:2002zd,Pieri:2007ir,Bergstrom:1998jj} investigated by recent zoomed high-resolution N-body simulations \\cite{Diemand:2006ik,Diemand:2005vz}: Clumpiness would result in a population of high galactic latitude extended gamma emitters with a typical annihilating DM gamma spectrum. These kinds of objects could hardly be associated to astrophysical emitters (but see \\cite{Baltz:2006sv}). In these models the size of the clumps is expected to have a characteristic distribution and thus the anisotropy of the integrated signal from all the clumps also exhibits a characteristic behavior \\cite{Pieri:2007ir}.  Likewise, the expected angular anisotropies both in the case of an astrophysical and of a DM origin of the CGB can be calculated, and have received increasing attention in the last few years \\cite{Cuoco:2006tr,Ando:2005xg,Ando:2006mt,Ando:2006cr,Miniati:2007ke}. In the following we will further pursue this issue addressing the differences expected in the two cases and their detectability in the light of the improved statistics that will be available, when the GLAST observatory is launched and start to take data in the near future. We will compare throughout the paper our findings in particular with \\cite{Ando:2005xg,Ando:2006cr} that deal specifically with anisotropies induced by DM annihilation. Already, there have been claims \\cite{Elsaesser:2004ap,deBoer:2005tm} of a DM signal in the CGB as observed by EGRET (see also \\cite{Bergstrom:2006tk,deBoer:2006ck,deBoer:2007zc}), although with the limited EGRET statistics and with the uncertainties in the galactic foregrounds, alternative astrophysical explanations cannot be ruled out. On the other hand, with the improved statistics from GLAST, a proper analysis of the anisotropy properties of the CGB should be able to prove, or disprove, the DM interpretation of features in the CGB spectrum.   Complementary to previous studies we shall employ in the following a parametric approach characterizing the expected CGB signal in terms of a few key parameters, that catch the relevant physical aspects of the problem, and varying them in order to asses the robustness and/or model dependence of the possible signatures. A further advantage of this approach is to make explicit the various assumptions employed throughout on which the final signatures depend. The relevant parameters in the following will be the degree of correlation of the CGB sources with matter and the absolute normalization of the signal, or, equivalently, the expected collected statistics. Further, we will also consider complementary anisotropy observables like the  cross-correlation between surveys of galaxies and the CGB and  the cross-correlation between different energy bands of the CGB. Together with the auto-correlations of the CGB these represent a set of independent observables that can be jointly employed improving considerably the sensitivity to the DM signal.    The paper is organized as follows: In section \\ref{CGBmain} we present a discussion of the horizons within which the CGB signal is expected to come, relevant for the determination of the intensity of the CGB anisotropies itself. In section \\ref{CGBanisotropies} we introduce the formalism to derive the CGB anisotropies in terms of the angular power spectrum. In section \\ref{GLASTforecast} we present a forecast for the expected statistics from GLAST and we discuss the possibility of disentangling the DM annhilation signal from that of astrophysical processes. In sections \\ref{crosscorr} and \\ref{crosscorr_eb} we introduce the cross-correlation between the CGB and galaxy surveys and the cross-correlation between different energy bands of the CGB and similarly we discuss the different behavior and sensitivity in the two cases of interest. In section \\ref{discussion} we discuss  how the previous conclusions apply to different possible scenarios for the CGB and DM properties. In section \\ref{final} we summarize and conclude. ",
        "conclusions": "\\label{final}  In the present work we have studied the kind of signatures that DM annihilation is expected to imprint in the anisotropies of the CGB, complementary to the signatures in the energy spectrum. We have addressed the main physical ingredients contributing to the DM signature and discussed the robustness of the signature with respect to various possible scenarios. We can summarize our findings as follows: \\begin{itemize} \\item{The DM annihilation signal traces in general the matter distribution quadratically due to its $\\rho_\\chi^2$ dependence. However, an effective linear correlation can arise if the signal is significantly enhanced by the presence of cuspy haloes or sub-galactic clumps. We have defined the two extreme ``linear'' and ``quadratic'' scenarios. The first corresponds to the case in which the cosmological DM annihilation signal is dominated by galactic or sub-galactic structures while in the second the signal is dominated by emission on scales larger than that of a galactic halo. We have chosen a phenomenological approach introducing a parameter $\\xi$ that weights the two relative contributions exploring the DM signatures for different possible choices of $\\xi$.}  \\item{The anisotropies are determined both by the intrinsic fluctuations in the source field and by the size of the emission horizon $z_{\\mathcal{H}}$. For $E_{\\rm cut}\\agt 100$ GeV the horizon $z_{\\mathcal{H}}$ is essentially fixed by photon absorption in the EBL.  The bulk of the gamma-rays is expected to originate inside $z_{\\mathcal{H}}\\approx 0.5$, independent of whether they have a DM or an astrophysical origin. For $E_{\\rm cut}\\agt 10$ GeV,  DM annihilation in the quadratic scenario has a redshift horizon $z_{\\mathcal{H}}\\approx 3-4$. The horizon is still significantly limited by PP losses at this energy, otherwise exceeding $z_{\\mathcal{H}}\\approx 10$. Blazars and DM in the linear scenario have degenerate horizons $z_{\\mathcal{H}}\\approx 1$. }  \\item{ In the quadratic scenario the DM anisotropy signal is sensibly enhanced with respect to blazars for  $E_{\\rm cut}=$10 GeV. Further, also the shapes of the angular spectra differ significantly \\cite{Ando:2005xg,Ando:2006cr}. The signature remain standing also  for $E_{\\rm cut}=$ 100 GeV despite the decreased statistics and become particularly strong, being independent of uncertainties related to the blazar-matter bias or to the evolution of blazars. This scenario can easily be detected by GLAST and would constitute a strong signature of DM annihilation. The DM linear scenario, instead, exhibit the same level of fluctuations of blazars and the two thus have almost degenerate anisotropy features.}   \\item{The above signature in the angular spectrum remains quite robust as long as the the quadratic DM signal is at least at the 10-20\\% level with respect to the linear DM or the blazar component. A further uncertainty to take into account is the normalization of the CGB (and thus the available statistics) that is likely to be reduced if part of the sources contributing to the CGB will be resolved by GLAST. If the normalization is reduced by an extreme factor of 5, (20\\% of the present EGRET value), the pure quadratic DM scenario exhibits still a relevant anisotropy transition signature. If the quadratic DM contributes for a 20\\% (thus, 4\\% of the present EGRET value) then the detection of the signature becomes quite challenging.}  \\item{The cross-correlation between the CGB and  a survey of galaxies and  the cross-correlation between different energy bands of the CGB provide further independent and sensitive observables that can be employed in combination with the CGB auto-correlation. A joint analysis of all the anisotropy observables considerably improves the sensitivity to the DM signal and, more in general, the power of the statistical analysis. In principle, the exact contribution from DM annihilation in sub-galactic clumps and cuspy haloes  can be treated as free parameters (instead of relying on a model) and inferred from the analysis.}       \\end{itemize}  The above conclusions hold exactly if a perfect cleaning of the galactic foregrounds and a lossless extraction of the CGB signal is possible. The analysis of foregrounds will be likely the main challenge in the study of the CGB. Clearly, given the above shown potential of CGB anisotropies in looking for DM signatures, it would be worth to perform further detailed studies on the issue. The launch of the GLAST satellite is expected by the middle of 2008, while the satellite AGILE \\cite{AGILE} launched in April 2007 is currently already taking data. The improvement in statistics compared to EGRET will allow for new, powerful tools to search for exotic contributions to the gamma-ray signal.  The anisotropy analysis of the CGB in particular, if foregrounds contaminations can be efficiently kept under control, promises to provide a clear signature of DM annihilation or, in the case of a negative answer, to obtain new constraints on the DM properties, complementary to a pure energy spectrum analysis.  \\vspace{1pc}"
    },
    "0710/0710.4070_arXiv.txt": {
        "abstract": "{We investigate how different models that have been proposed for solving the dark matter problem can fit the velocity dispersion observed around elliptical galaxies, on either a small scale ($\\sim$ 20kpc) with stellar tracers, such as planetary nebulae, or large scale ($\\sim$ 200kpc) with satellite galaxies as tracers. Predictions of Newtonian gravity, either containing pure baryonic matter, or embedded in massive cold dark matter (CDM) haloes, are compared with predictions of the modified gravity of MOND. The standard CDM model has problems on a small scale, and the Newtonian pure baryonic model has difficulties on a large scale, while a fit with MOND is possible on both scales. ",
        "introduction": "Measuring the velocity field in and  around galaxies is the main way to test the dark matter distribution at small and intermediate scales. The observation of what are apparently non-Newtonian rotation curves around spiral galaxies (e.g. Rubin et al. 1980) has been first solved by assuming that galaxies are embedded in dark matter haloes. Numerical simulations, however, predict a radial distribution for the CDM model much more concentrated than what is observed (Gentile et al 2004; de Blok 2005).  An alternative explanation was proposed by Milgrom (1983), as MOdified Newtonian Dynamics (MOND). When the Newtonian acceleration falls below the critical value  $a_0 \\sim 2 \\times 10^{-10}$ m s$^{-2}$, the gravity law is empirically modified and then declines in 1/$r$ instead of 1/r$^2$. Around spherical systems, the modified acceleration $g$ satisfies the relation $$ g \\mu(g/a_0)  =  g_n$$ where $g_n$ is the Newtonian acceleration. For non-spherical geometry, this is only an approximation. However, in this Letter we only consider spherical systems as representing elliptical galaxies and we adopt $\\mu(x) = x/\\sqrt{1+x^2}$.  This model is very successful on a galactic scale; in particular, it explains a large number of rotation curves of galaxies, with some exceptions (Gentile et al. 2004), and naturally the Tully-Fisher relation, (e.g. Sanders \\& McGaugh 2002). While the dark matter problem is observationally very clear around spiral galaxies, thanks to their rotation curve measured  with the cold hydrogen gas at 21cm, which is nearly in circular orbits (e.g. Bosma 1981, Verheijen \\& Sancisi 2001), the situation is much more complex around elliptical galaxies, with little or no rotation.  Recently, planetary nebulae have been used as an efficient tool for measuring the velocity field at large radii in early-type galaxies, and they complement stellar absorption kinematical studies (Romanowky et al. 2003). In typical elliptical galaxies, the velocity-dispersion profiles were found to decline with radius, up to 5 effective radii, thereby requiring no dark matter at all. Dekel et al. (2005) show that the data are still compatible with the usual dark matter models, if the planetary nebulae tracers have particularly radial orbits in the outer parts, because of a recent merger with small impact parameter. However, the recent results from Douglas et al (2007) challenge this interpretation, prolonging the decline to more than 7 effective radii.  On larger scales around early-type galaxies, from 50 to 300kpc, Klypin \\& Prada (2007), KP07, have proposed to test gravity models with satellite galaxies as a tracer.  From the Sloan Digital Sky Survey, they stack several thousand galaxies in 3 luminosity classes and determine the number density of satellites and their velocity dispersion around them. In each mass range, the radial distributions are obtained with around 1500 satellites, although about 1.5 satellite exist around each galaxy.  This large-scale galaxy neighborhood has not been widely tested yet in modified gravity. The well-known difficulty of MOND in clusters has found a possible solution with neutrinos of  2 eV mass (Sanders 2003; Angus et al 2007a), and the escape velocity around giant galaxies like the Milky Way was shown to correspond to observations, when including the external field effect (Famaey et al 2007; Wu et al 2007).  In this work, we  solve the Jeans equation for the distribution of the velocity dispersion around elliptical galaxies, and in particular we fit the NGC3379 galaxy, where the most extended data is available for the velocities. We also further explore fits at larger radii for the special case of NGC 3379, with satellite galaxies as tracers, as done by KP07 and Angus et al (2007b). This is only statistically valid around a generic early-type galaxy, with a mass comparable to NGC 3379. This is the brightest galaxy of a group, but the observed companion velocity is not statistically significant.  \\begin{figure}[ht!] \\center{ \\includegraphics[angle=-90, width=8cm]{8569-f1.ps} \\caption{ Fit of the blue-light distribution in NGC 3379. The data points  are taken from Douglas et al (2007). The curve represents the light profile derived with the mass-to-light ratio given in Table \\ref{massp} and a mass model of cored-Hernquist type (Table \\ref{tracers}). } \\label{logsigma} } \\end{figure}  \\vspace{-1 cm} ",
        "conclusions": "In contrast to the conclusion of KP07, MOND does not predict constant velocity dispersion with radius in the neighborhood of elliptical galaxies. There is a wide latitude for varying the possible anisotropy parameter according to the  scale and the tracer considered. Our best fit starts from  a tangential or isotropic configuration near the center and evolves progressively to a radial one for each of the three scales considered, which appears quite realistic.  While the Newtonian model without dark matter has problems in the outer parts, the CDM1 model encounters severe difficulties in the inner parts. The CDM2 model, with a reduced dark matter relative to the visible mass, can also fit the data quite well. Compared to the MOND model,  it requires a larger radial anisotropy on each scale, and its $\\beta$ profile  is unusual for the satellite tracer with a negative slope. For  $\\beta$= constant, the $\\sigma_{los}^{sat}$ would be too steep compared to the observations. And by introducing anisotropy increasing with radius,  the $\\sigma_{los}^{sat}$ slope is even more increased.  In the MOND regime, the external field effect (EFE, e.g. Wu et al 07) would also change the predictions. This will modulate the actual force on the particle tracers, and on the velocity. Since there was still latitude for fitting the observations with MOND, we feel that it is still possible to consider it with other anisotropy parameters. However, we note that the external field effect is not even necessary, for reproducing the observed velocity dispersion slope.   \\vspace{-0.2 cm}"
    },
    "0710/0710.1180_arXiv.txt": {
        "abstract": "\\noindent  We use a gravitational bar torque method to compare bar strengths (the maximum tangential force normalized by radial force) in $B$ and $H$-band images of 152 galaxies from the Ohio State University Bright Spiral Galaxy Survey. Our main motivation is to check how much the difference in the rest-frame wavelength could affect comparisons of bar strengths in low and high redshift observations. Between these two bands we find an average bar strength ratio \\QBH= $1.25$ which factor is nearly independent of the morphological type. We show that \\QBH $> 1$ is mostly due to reduced bulge dilution of radial forces in the $B$-band. The bar torque method needs an estimate for the vertical scale height of the galaxy, based on the radial scale length of the disk and the galaxy's morphological type. Since these two might not always be possible to determine at high redshifts in a reliable manner, we also checked that similar results are obtained with vertical scale heights estimated from the radii corresponding to the $K$-band surface brightness of 20 mag/arcsec$^2$. Also, we made a simple test of the usability of the bar torque method at high redshifts by checking the effects of image degradation (nearest neighbour sampling without any adjustment of noise levels): we found that the estimated bar strengths varied by $\\pm 10\\%$ at most as long as the total extent of the bar was at least 10 pixels. Overall, we show that the gravitational bar torque method should provide a proficient tool for quantifying bar strengths also at high redshifts. ",
        "introduction": "Bars are important drivers of galaxy evolution, which highlights the importance of studying bar fractions and the properties of bars at a large range of redshifts. Bars generally have fairly old stellar populations, thus being well visible in the near-IR, where also the obscuration by dust is minimal. However, the near-IR images we obtain from the distant galaxies are redshifted so that at redshift $z \\sim 1$ their rest-frame wavelength lies in the optical. Therefore, it is important to compare the optical and near-IR properties of bars in local galaxies.  For that purpose the $B$ and $H$-band images are very convenient, offering a wide separation in wavelengths. Namely, bars are still prominent in the $B$-band \\citep{osu,whyte,sheth06}, where they can be identified in a fairly reliable manner, whereas, due to the Balmer break, bars typically disappear in the ultraviolet \\citep{sheth}. On the other hand, the $H$-band is already fairly free of the effects of dust and thus offers a reliable tracer for the stellar mass distribution.  If bars were formed due to instabilities in the disk, one would expect a large fraction of barred galaxies at high redshifts in the framework of the hierarchical clustering models. This is because these models indicate dynamically colder disks in the distant Universe \\citep*{NFW}. Also, interactions of disk galaxies were probably more frequent in the past \\citep*[and references therein]{fer}, promoting the formation of bars \\citep*[discussed in][]{elm}. Therefore, it was surprising to find a lack of bars in galaxies at redshift of $z>0.7$ (\\citealp{abe}; \\citealp{vdb96}; \\citealp{vdb00}; \\citealp*{vdb01}). Since this might be an artifact due to noise, poorer image resolution, and redshifted wavelength, \\citet{vdb} degraded about 100 $B$-band images of the Ohio State University Bright Spiral Galaxy Survey ({\\sc osubsgs}) to study these effects. By mimicking the characteristics of the images in the Hubble Deep Field,  using the Wide Field Planetary Camera 2 on HST, they created artificial images of those galaxies as they would appear in the $I$-band at a redshift $z=0.7$. They found that two-thirds of bars initially classified as strong bars (SB), were still detectable in the degraded images, whereas weak bars (SAB) would largely disappear. They concluded that while selection effects reduce the number of observed bars, they cannot completely explain the lack of bars at high redshift.  The early studies of bar frequencies were challenged by \\citet{sheth}, who found similar bar frequencies at high and at low redshifts for very large bars. However, the number statistics were a severe concern, since only four bars were identified. These authors used galaxies in the Hubble Deep Field-North, observed at $V$, $I$ and $H$-bands, paying particular attention to the resolution limits due to which small bars cannot be detected at high redshifts. In two subsequent studies \\citep{elm,jogee}, higher resolution images were used, based on the optical Advanced Camera for Surveys (ACS) at the Hubble Space Telescope. Both studies found a constant bar fraction in the redshift range $z = 0-1$. \\citet{jogee} also found that it holds separately for the distinct intervals $z = 0-0.7$ and $z = 0.7-1.0$, using broad $V$ and $z$-bands to identify the bars. Taking into account the wavelengths used in their studies this constant bar fraction is not completely expected. Namely, if the true bar fraction is the same at all redshifts and bars are more prominent in the near-IR, one would expect a larger number of detected bars in the intermediate redshifts. There the rest-frame band shifting is not yet important, so that both strong and weak bars should be detected \\citep{sheth06}. However, in the above studies the number statistics are still too poor (less than 300 galaxies spread over the entire redshift range) for any reliable discussion of how the bar fraction evolves as a function of redshift. Also, using ACS very deep images are required for detecting weak bars at high redshift.   The above studies have been superseded recently by \\citet{sheth07}, using wide field {\\sc cosmos} observations in $I$-band at $z < 0.84$, based on deeper images than used in the previous studies. Contrary to most previous studies they found that the fraction of bars declines rapidly with redshift. This drop of bars was also found to be more dramatic for strong bars. These studies of bar frequencies clearly show, besides the difficulty of measuring bar frequencies at high redshifts, also the importance of measuring the strengths of the bars at high redshifts in a reliable manner.   In this study we use the {\\sc osubsgs} to compare bar strengths derived from the $B$ and $H$-band images of nearby galaxies, using a gravitational bar torque method. Besides bar strength, our method gives simultaneously also the length and the relative luminosity of the bar. The outline of this paper is as follows. In Sections \\ref{sample} and \\ref{method} we describe the details of our sample and the method used, and also address the factors which might affect the bar strength measurements at higher redshifts. We then show comparisons between $B$ and $H$-band bar strengths in Section \\ref{highz}; the results are discussed in Section \\ref{discussion}, and summarized in Section \\ref{conclusions}. ",
        "conclusions": "We have used a gravitational bar torque method to compare bar strengths derived from the $B$ and $H$-band images for the galaxies in the {\\sc osubsgs} sample. We presented four different ways to do the comparison. Further we subjected the method to various tests simulating a high redshift environment. We estimate $z_0$ through $r_{K_{20}}$, and degrade the $B$-band images using a nearest neighbour algorithm before calculating the bar strengths. Our main conclusions can be summarized as follows:  (i) Quite unexpectedly, we found that bars appear to have stronger torques in the optical than in the near-IR: the conversion factor between $B$ and $H$-bands is 1.25. We also showed that this is due to reduced dilution of radial forces by relatively smaller bulges in the optical.  (ii) The scaled bar lengths, $barlen$/$r_{K_{20}}$, and $r_{Q_g}/barlen$ are similar in the $B$ and $H$-bands, with the bars in the $B$-band being on average slightly shorter than bars in the $H$-band.  (iii) An approximation of the vertical scale height while calculating $Q_g$ can be estimated from $r_{K_{20}}$. Resulting bar strengths will be overestimated for early types and underestimated for late types, by less than 15\\%.  (iv) The $Q_g$ method is found to work well even on low resolution images. We degraded our $B$-band images using the nearest neighbour algorithm and find that as long as the total bar length is at least $\\sim 10$ pixels, the resulting $Q_g$ value is typically within 10\\% of the original one.  We have shown that the $Q_g$ method provides a reliable tool for measuring bar strengths at high redshifts: with the pixel resolution of $0\\farcs 05$ of the ACS, bars with $\\sim 2-3$ kpc can still be reliably measured at $z \\sim 1.0$."
    },
    "0710/0710.2847_arXiv.txt": {
        "abstract": "{M\\,87 is the first extragalactic source detected in the TeV range that is not a blazar. With the increasing performances of ground-based \\v{C}erenkov telescopes, we can now probe the variability in the $\\gamma$-ray flux at small timescales, thus putting strong constraints on the size of the emitting zone. The large scale jet of M\\,87 is misaligned with respect to the line of sight. A modification of standard emission models of TeV blazars appears necessary to account for the $\\gamma$-ray observations despite this misalignment. } {We explain TeV $\\gamma$-ray spectra and fast variability of M\\,87 by invoking an emission zone close to the central supermassive black hole, which is filled with several plasma blobs moving in the large opening angle of the jet formation zone. } {We develop a new multi-blob synchrotron self-Compton (SSC) model with emitting blobs set on a cap beyond the Alfv{\\'e}n surface in the jet, at a distance of $\\sim$100\\,$r_g$ from the central engine to interpret the high energies inferred by new TeV observations. We present a SSC model that is explicitly adapted to deal with large viewing angles and moderate values of the Lorentz factor inferred from (general relativistic) magnetohydrodynamic models of jet formation. } {This scenario can account for the recent $\\gamma$-ray observations of M\\,87 made by the High Energy Stereoscopic System (H.E.S.S.) telescope array. We find individual blob radii of about $10^{14}$\\,cm, which is compatible with the variability on timescales of days recently reported by the H.E.S.S. collaboration and is of the order of the black hole gravitational radius. Predictions of the very high energy emission for three other sources with extended optical or X-ray jet which could be misaligned blazars still with moderate beaming are presented for one Seyfert 2 radiogalaxy, namely Cen\\,A, one peculiar BL\\,Lac, PKS\\,0521$-$36, and one quasar, 3C\\,273. } {} ",
        "introduction": "M\\,87 is a well-known nearby giant elliptical galaxy \\citep[z=0.00436,][]{2000MNRAS.313..469S} close to the center of the Virgo cluster, which shows a multi-spectral jet, signature of an active galactic nucleus (AGN). Its jet is one of the best known, at all scales, thanks to its nearby location and its strong synchrotron radiation in the optical band. M\\,87 is classified as FR\\,I based on its radio morphology. \\citet{2002ApJ...568..133W} observed the jet with {\\it Chandra} on July 29 and 30, 2000 and detected it up to a distance of $\\sim$21$\\arcsec$ from the core in the X-ray band, which implies that the jet is not as strongly aligned along the line of sight \\citep[see also][]{1996MNRAS.283..873R} as in the case of blazars.  At radio wavelengths, an impressive jet, which extends up to a few tens of kiloparsecs, can be seen. The central engine is thought to be a supermassive black hole (SMBH) with a mass of $M_\\mathrm{BH} \\sim 3 \\times 10^9 M_{\\sun}$ \\citep{1997ApJ...489..579M}. The scale length is thus $r_g = G M_\\mathrm{BH} / c^2 \\sim 4.5 \\times 10^{14}$\\,cm $\\sim 1.4 \\times 10^{-4}$\\,pc. Using the {\\it Hubble Space Telescope} ({\\it HST}), \\citet{1999ApJ...520..621B} observed superluminal apparent motions of about $4c$--$6c$ beyond 400\\,pc for the internal knots, between 1994 and 1998, thus confirming that the jet is relativistic. They conclude that the jet is oriented within $19 \\degr$ of the line of sight.  Due to the presence of a SMBH in the core and the presence of the jet, M\\,87 was deemed an interesting candidate for TeV emission. \\citet{2004ApJ...610..156L} reported an upper limit observed with Whipple in 2000 and 2001, simultaneously with X-ray flares observed by {\\it RXTE}. HEGRA observed M\\,87 in 1998 and 1999 for a total exposure of 77\\,h after data quality selection \\citep{2003A&A...403L...1A,2004NewAR..48..407B}. A 4.1$\\sigma$ significance was recorded and an integrated flux ($E>730$\\,GeV) of 3.3\\% Crab was measured.  Recently, \\citet{2006Sci...314.1424A} have observed M\\,87 with the High Energy Stereoscopic System (H.E.S.S.)\\footnote{\\url{http://www.mpi-hd.mpg.de/hfm/HESS/}} between 2003 and 2006 in 89\\,h live-time with a $13\\sigma$ detection and discovered variations on timescales of about 2 days, 10 times faster than that observed in any other waveband. This shows that the emission region is very compact, with a dimension of the order of a few Schwarzschild radii. These observations, thus confirming the detection by HEGRA \\citep[][]{2005astro.ph..4395B}, are particularly interesting since M\\,87 is the first non-BL\\,Lac extragalactic object ever observed at TeV energy. Radio-loud galaxies contain AGNs with jets like blazars, but the jet emission is less strongly boosted due to larger viewing angles between the jet and the observer's line of sight. It is therefore a challenge for standard models of TeV blazars to explain the very high energy (VHE) emission of M\\,87.  In this paper, we present a modified synchrotron self-Compton (SSC) scenario to explain the VHE emission of M\\,87. Classic SSC models \\citep[e.g.][]{1979A&A....76..306G,1996ApJ...463..555I,1996ApJ...461..657B,1999MNRAS.306..551C,2001A&A...367..809K} are applied to blazars, which are beamed sources, and cannot account for the observations of radiogalaxies like M\\,87. Our goal is to further develop one of these models to reconcile beamed and unbeamed sources in the same framework of models. Such propositions for unification of AGNs have already been studied considering orientation effects \\citep[e.g.][]{1993ARA&A..31..473A,1995PASP..107..803U}, or radio/X-ray power among BL\\,Lac objects, flat-spectrum radio-loud quasars (FSRQs) and FR\\,Is \\citep[e.g.][]{1998MNRAS.299..433F,1998MNRAS.301..451G,2000MNRAS.318..493C}.  A short description of the leptonic blob-in-jet model for TeV blazars is found in Sect.~\\ref{sec:bij}, and its development and application to M\\,87 are described in Sect.~\\ref{sec:multi_blobs} and Sect.~\\ref{sec:M87}. In the framework of misaligned BL\\,Lac-like objects, we then try to predict VHE fluxes for objects with optical/X-ray extended jets in Sect.~\\ref{sec:predictions}. Implications on unification schemes of AGNs are discussed in Sect.~\\ref{sec:unification}. ",
        "conclusions": "}  We have presented a SSC model to interpret VHE emission of M\\,87 as well as other misaligned sources with extended optical/X-ray jet. This model accounts in a simple way for a differential Doppler boosting effect by modeling the emission of several blobs of plasma located in the broadened formation zone of the jet close to the SMBH, just beyond the Alfv{\\'e}n surface predicted by MHD models.  Our scenario provides a reasonable interpretation of the H.E.S.S. VHE observations of M\\,87 and provides the possibility to extend standard leptonic models of TeV blazars to other types of AGNs. However, we do not exclude other leptonic or hadronic models. For instance, \\citet{2007arXiv0704.3282N} recently interpreted the H.E.S.S. observations from 2005 of M\\,87 by invoking acceleration and radiation of electrons in the black hole magnetosphere, which is another kind of leptonic model. Hadronic models also cannot be excluded as efficient particle acceleration processes can occur in the close surroundings of the black hole.  More observations are needed to constrain the emission models and especially to distinguish between hadronic and leptonic scenarii. The upcoming {\\it GLAST} mission and the H.E.S.S.\\,II project will certainly help to understand the mechanisms at work in the AGNs by exploring spectral ranges below TeV, which is decisive to constrain the shape of the inverse Compton bump. Moreover, the lack of genuine simultaneous multiwavelength campaigns on M\\,87 needs to be filled, especially since this source is known to be variable at small timescales in VHE.  Several types of active nuclei are potential emitters of VHE photons with predicted TeV fluxes detectable by present \\v{C}erenkov arrays like H.E.S.S. and MAGIC, or by the next generation of instruments such as the CTA project. Such data will be crucial to test AGN unifying schemes."
    },
    "0710/0710.0590_arXiv.txt": {
        "abstract": "We use the magnetic butterfly diagram to determine the speed of the magnetic flux transport on the solar surface towards the poles. The manifestation of the flux transport is clearly visible as elongated structures extended from the sunspot belt to the polar regions. The slopes of these structures are measured and interpreted as meridional magnetic flux transport speed. Comparison with the time-distance helioseismology measurements of the mean speed of the meridional flows at the depth of 3.5--12~Mm shows a generally good agreement, but the speeds of the flux transport and the meridional flow are significantly different in areas occupied by the magnetic field. The local circulation flows around active regions, especially the strong equatorward flows on the equatorial side of active regions affect the mean velocity profile derived by helioseismology, but do not influence the magnetic flux transport. The results show that the mean longitudinally averaged meridional flow measurements by helioseismology may not be used directly in solar dynamo models for describing the magnetic flux transport, and that it is necessary to take into account the longitudinal structure of these flows. ",
        "introduction": "The largest-scale velocity fields on the Sun consist of differential rotation and meridional circulation. The differential rotation is defined as an integral of the zonal (East--West) component of the velocity field depending on the solar latitude, $b$, and radius. The meridional flow is calculated as an integral of the North--South component of the velocity field, generally depending again on the latitude and radius. Both the differential rotation and meridional circulation are the key ingredients of the solar dynamo. The differential rotation plays an important role in generating and strengthening of toroidal magnetic field inside the Sun, while the meridional flow transports the magnetic flux towards the solar poles resulting in cyclic polar field reversals \\citep[for a recent review, see][]{Brandenburg2005}.  The meridional flux transport seems to be an essential agent influencing the length, strength and other properties of solar magnetic cycles. Generally, the slower the meridional flows are, the longer the next magnetic cycle is expected. Dynamo models showed that the turn-around time of the meridional cell is between 17 and 21 years, and that the global dynamo may have some kind of memory lasting longer than one cycle \\citep{2006GeoRL..3305102D}.  The speed of the meridional flow and its variation with the solar cycle measured by local helioseismology in the subsurface layers of the Sun were used as an input in the recent flux-transport models \\citep{2006ApJ...649..498D}. In local helioseismology measurements (e.g. \\citeauthor{2004ApJ...603..776Z} \\citeyear{2004ApJ...603..776Z}, \\citeauthor{2006ApJ...638..576G} \\citeyear{2006ApJ...638..576G}), the meridional flow was derived from a general subsurface flow field by averaging the North-South component of the plasma velocity over longitude for a Carrington rotation period. The studies revealed that the mean meridional flow varied with the solar activity cycle. These variations may significantly affect solar-cycle predictions based on the solar dynamo models, which assume that the magnetic flux is transported with the mean meridional flow speed \\citep{2006ApJ...649..498D}.  Our goal is to verify this assumption and to investigate the relationship between the subsurface meridional flows and the flux transport. In this study, we show that the mean meridional flows derived from the time-distance helioseismology subsurface flow maps are affected by strong local flows around active regions in the activity belts. These local flows have much less significant effect on the magnetic flux transport. ",
        "conclusions": "We have compared the measurements of the meridional speed derived from two different techniques: by time-distance local helioseismology and by measuring the flux transport speed using the magnetic butterfly diagram. We have found that the results agree quite well in general, but they differ in regions occupied by local magnetic fields. The detailed flow maps from helioseismology show that this is partly due to the presence of meridional counter-cells at the equatorial side of magnetic regions, which influence the time-distance derived meridional flow profile, but does not influence the magnetic flux transport.  We have studied eleven non-consecutive Carrington rotations covering one solar cycle. The effect of the local flows around active regions and especially on their equatorial side is noticed in all the studied cases. Therefore, this behaviour seems to be a common property of the subsurface dynamics around active regions located in the activity belt. However, we have to keep in mind that both datasets are not directly comparable, since the time-distance flow maps represent the behaviour of flows during one Carrington rotation, while the butterfly diagram tracking procedure provide results averaged over few Carrington rotations. The more homogeneous data for local helioseismology are needed to study this effect in more detail.  The results show that the speed of the magnetic flux transport towards the solar poles may significantly deviate from the longitudinally averaged meridional flow speed derived from local helioseismology measurements, which are affected by local circulation flows around active regions in the activity belt. Therefore, using the longitudinally averaged meridional flow profile from helioseismology in the solar cycle models for description the flux transport is not justified. The longitudinal structure of these flows should be taken into account."
    },
    "0710/0710.5183_arXiv.txt": {
        "abstract": "We present results from {\\it Chandra} X-ray and {\\it Spitzer} mid-infrared observations of the interacting galaxy pair NGC~6872/IC~4970 in the Pavo galaxy group and show that the smaller companion galaxy IC~4970 hosts a highly obscured active galactic nucleus (AGN). The $0.5-10$\\,keV X-ray luminosity of the nucleus is variable, increasing by a factor $2.9$ to $1.7 \\times 10^{42}$\\ergs (bright state) on $\\sim 100$\\,ks timescales. The X-ray spectrum of the bright state is heavily absorbed ($N_{\\rm H} = 3 \\times 10^{23}$\\cms for power law models with $\\Gamma = 1.5-2.0$) and shows a clear $6.4$\\,keV Fe K$\\alpha$ line with equivalent width of $144-195$\\,eV. Limits on the diffuse emission in IC~4970 from {\\it Chandra} X-ray data suggest that the available power from Bondi accretion of hot interstellar gas may be an order of magnitude too small to power the AGN. {\\it Spitzer} images show that $8\\micron$ nonstellar emission is concentrated in the central $1$\\,kpc of IC~4970, consistent with high obscuration in this region. The mid-infrared colors of the nucleus are consistent with those expected for a highly obscured AGN. Taken together these data suggest that the nucleus of IC~4970 is a Seyfert $2$, triggered and fueled by cold material supplied to the central supermassive black hole as a result of the off-axis collision of IC~4970 with the cold-gas rich spiral galaxy NGC~6872. ",
        "introduction": "\\label{sec:introduction}  Supermassive black holes are known to exist at the centers of most galaxies. The strong correlation of the central supermassive black hole mass with the mass of the host galaxy's stellar bulge, as represented by the central stellar velocity dispersion (Ferrarese \\& Merritt 2000;  Gebhardt \\etal 2000) or  K-band  absolute magnitude (Marconi \\& Hunt 2003), indicates that the evolution of the central black hole and that of its host galaxy are inextricably linked. Most recent work relating the activity of the central supermassive black hole to galaxy evolution has focused on the interaction of radio-loud active galactic nuclei (AGN) with their environment, in order to resolve the cooling flow problem in galaxy clusters and bring hierarchical models of galaxy formation into agreement with observations of galaxies at low redshift (see, e.g., Dunn \\& Fabian 2006; Best \\etal 2006; Croton \\etal 2006). Hardcastle \\etal (2007) separated these AGN into two classes, low excitation and high excitation radio galaxies, based on their optical emission line characteristics, and suggested that the observed differences between the two classes are determined by the fuel source and accretion mode of the central black hole. Low excitation radio galaxies tend to be massive, dominant elliptical galaxies, residing in the deep dark matter potentials near the cores of large groups and rich clusters, and are surrounded by an ample supply of hot gas from the intracluster medium falling into their large gas halos. They show no evidence for a standard, geometrically thin, optically thick, luminous accretion disk. Radiatively inefficient Bondi accretion of hot gas accreted from their surroundings onto the central supermassive black hole is sufficient to supply the total energy, mechanical as well as radiative, observed in these systems. In contrast, high excitation radio galaxies (containing the most powerful radio galaxies known) exhibit high-excitation, narrow-line optical emission in their spectra, characteristic of emission from an accretion disk. Hardcastle \\etal (2007) argue that AGN activity in these galaxies is powered by radiatively efficient accretion of cold gas acquired in a major merger with a massive cold-gas-rich partner. The high radio power and jets observed in high excitation radio galaxies often may result from the formation of a rapidly spinning nucleus from the coalescence of the merging galaxies' two massive black holes (Wilson \\& Colbert 1995; Capetti \\& Balmaverde 2005).  Since cold-gas accretion is associated with the galaxy's interaction and merger history, and does not depend on the depth of the dark matter potential or the presence of a hot intracluster medium, galaxy collisions of lower mass galaxies with dusty, gas-rich partners in galaxy groups also would be expected to trigger AGN activity and black hole growth. Gravitational and hydrodynamical forces, induced by the galaxies' mutual interaction, funnel dust and cold gas from the gas-rich partner into the interacting, lower-mass companion galaxy's nuclear region, where it settles into an optically thick accretion disk before subsequently accreting onto the central black hole. As in the evolutionary model of Churazov \\etal (2005), accretion onto the supermassive black hole is radiatively efficient and radio emission from the AGN may be weak or absent. Such galaxy collisions were likely frequent at high redshift ($z \\sim 1.3$), when the fraction of blue, gas-rich starforming galaxies in moderately massive galaxy groups was high and galaxies were rapidly evolving (Cooper \\etal 2006; Gerke \\etal 2007). Thus, while perhaps less dramatic than radio-loud systems, cold-gas accretion in off-axis galaxy collisions and minor mergers may play an equally important role in the evolution of lower mass galaxies and moderately massive galaxy groups into what we see today.  Observations of nearby interacting galaxies in similarly moderately massive groups offer a unique window into the dynamical processes that triggered AGN activity in lower mass galaxies and influenced the coevolution of the host galaxy with its central supermassive black hole at this earlier epoch, when galaxies were rapidly transforming. The Pavo group is such a nearby ($z=0.01338$) group with a cool $0.5$\\,keV intra-group medium, $\\sim 13$ member galaxies and velocity dispersion $\\sim 425$\\kms (Machacek \\etal 2005) similar to the properties of $z \\sim 1$ groups observed in the DEEP2 survey (Gerke \\etal 2005). In the top panel of Figure \\ref{fig:pavogroup}, we present a Bj-band image of the dominant Pavo group galaxies, showing the elliptical galaxy NGC~6876 at the Pavo group center and the large, gas-rich tidally-distorted SAB(rs)c spiral galaxy NGC~6872, located $8\\farcm7$ to the northwest. The  smaller lenticular galaxy IC~4970, that is only $1/5$ to $1/10$ as massive as the spiral galaxy NGC~6872 (Mihos \\etal 1993), is located $1\\farcm12$ to the north of NGC~6872's center near a break (`knee') in NGC~6872's northern tidal arm. The small $74$\\kms line-of-sight velocity difference\\footnote{All line-of-sight velocities in this paper are taken from the CfA Redshift Survey by Martimbeau \\& Huchra. Online versions of the catalog, supporting software, and documentation on the CfA Redshift Survey are available at http://cfa-www.harvard.edu/\\~\\,huchra/zcat.} between IC~4970 and NGC~6872 suggests that the galaxy pair forms a spiral-dominated subgroup, while {\\it XMM-Newton} observations of a hot, gas trail linking, in projection, the central elliptical galaxy NGC~6876 and NGC~6872/IC~4970 indicate that the NGC~6782/IC~4970 subgroup has recently passed supersonically through the Pavo group core (Machacek \\etal 2005). Thus the Pavo group, like groups at high redshift, may be dynamically young and the interactions between galaxies representative of those important to galaxy evolution at that earlier epoch.  \\begin{figure}[t] \\begin{center} \\includegraphics[height=2.0in,width=3in]{f1acomp.ps} \\includegraphics[height=2.0in, width=3in]{f1bcomp.ps} \\caption{({\\it top}) Digitized Sky Survey  Bj band image, taken with the UK $48$inch Schmidt Telescope, of the dominant elliptical galaxy NGC~6876 and nearby NGC~6877 at the center of the Pavo group  and the large spiral galaxy NGC~6872 with its tidally interacting companion galaxy IC~4970 located $8\\farcm7$ to the northwest. ({\\it bottom}) Zoom-in from the image in the top panel to highlight the interaction of IC~4970 with the spiral galaxy NGC~6872. } \\label{fig:pavogroup} \\end{center} \\end{figure}  IC~4970 and NGC~6872 have long been known to be a tidally interacting galaxy pair (Vorontsov-Velyaminov 1959). As shown in the bottom panel of Figure \\ref{fig:pavogroup}, tidal tails extend outward more than $30$\\,kpc ($\\sim 2\\farcm$) to the east and west from the tip of NGC~6872's spiral arms. A stellar bridge connects the `knee' in NGC~6872's northern arm and tidal tail to the lenticular companion galaxy IC~4970, indicating that tidal interactions between IC~4970 and NGC~6872 are ongoing. The spiral galaxy NGC~6872 is gas rich. From CO J=1-0 line emission, Horellou \\& Booth (1997) infer $9.6 \\times 10^8\\Ms$ of molecular hydrogen in NGC~6872's central region. Although atomic hydrogen is not observed near the center of the spiral galaxy,  $1.41 \\times 10^{10}\\Ms$ of HI gas is distributed along the full extent of NGC~6872's spiral arms and associated tidal tails (Horellou \\& Booth 1997; Horellou \\& Koribalski 2007) and  $1.3 \\times 10^9\\Ms$ of HI gas is seen along the direction of IC~4970, that may be associated with the companion galaxy (Horellou \\& Koribalski 2007). H$\\alpha$ emission traces ionized gas from recent star formation along NGC~6872's outer spiral arms to the ends of the tidal tails and from the `knee' north in the stellar bridge (Mihos \\etal 1993). The distribution of H$\\alpha$ corresponds closely to the distribution of young star clusters observed in this system (Bastian \\etal 2005). However, Mihos \\etal (2003) found  no H$\\alpha$ emission from the central regions of NGC~6872 or associated with the companion galaxy IC~4970.  In this paper we focus on the effects of the IC~4970/NGC~6872 collision on the lower mass participant IC~4970, based on {\\it Chandra} X-ray and {\\it Spitzer} mid-infrared  observations that show IC~4970 hosts a highly-obscured active nucleus. Results for the two massive galaxies, NGC~6872 and NGC~6876, and the Pavo group gas from these observations will be reported in subsequent papers (Machacek \\etal 2007a,b in preparation). Our discussion is organized as follows: In  \\S\\ref{sec:obs} we briefly review the {\\it Chandra} and {\\it Spitzer} observations and our data reduction and processing procedures. In \\S\\ref{sec:analysis} we present our main analysis results. We discuss the X-ray time variability of the nuclear point source in \\S\\ref{sec:vary}, its X-ray spectral properties in \\S\\ref{sec:xrayspec}, and mid-infrared flux densities and colors in \\S\\ref{sec:iraccol}, and show that these properties are consistent with the classification of IC~4970's nucleus as a highly obscured Seyfert $2$. In \\S\\ref{sec:ULX} we briefly discuss the properties of a nearby ultra-luminous X-ray source  and in \\S\\ref{sec:diffuse} place limits on the amount of hot gas outside the central region of the galaxy. In \\S\\ref{sec:discuss} we determine the mass of the central black hole in IC~4970 and discuss possible accretion modes, concluding that nuclear activity in IC~4970 is most likely triggered and fueled by cold gas driven into the nucleus during IC~4970's ongoing  off-axis collision with the dust- and gas-rich spiral galaxy NGC~6872. We summarize our results in \\S\\ref{sec:conclude}. Unless otherwise indicated, we quote $90\\%$ confidence levels for the uncertainties in spectral parameters and $1\\sigma$ uncertainties on counts and count rates. Coordinates are J2000.0. Adopting the flat $\\Lambda$CDM cosmology from the three year WMAP results ($H_0 = 73$ \\kmsmpc, $\\Omega_m = 0.238$, Spergel \\etal 2007), the luminosity distance to the Pavo Group ($z=0.01338$) is $55.5$\\,Mpc and $1\\farcs$ corresponds to a distance scale of $0.262$\\,kpc. ",
        "conclusions": "\\label{sec:conclude}  In this paper we use {\\it Chandra} X-ray and {\\it Spitzer} IRAC mid-infrared observations to probe nuclear activity in IC~4970, the interacting companion galaxy to the large spiral galaxy NGC~6872 in the Pavo galaxy group. We find the following: \\begin{itemize} \\item{X-ray emission from the nucleus of IC~4970 is time variable with the mean $0.5-10$\\,keV count rate increasing by a factor of three on $\\sim 100$\\,ks time scales. This strongly suggests IC~4970's nucleus hosts an AGN rather than a compact nuclear starburst.} \\item{The X-ray spectrum of IC~4970's nucleus in the bright state shows a clear Fe K$\\alpha$ line and is well described by an absorbed power law plus narrow Gaussian line model with absorption column $N_{\\rm H} = 3 \\times 10^{23}$\\cms, fixed photon indices $\\Gamma = 1.5 - 2$, and an Fe K$\\alpha$ line equivalent width of $144-195$\\,eV. The $0.5-10$\\,keV X-ray luminosity of the nucleus in the bright state is $1.7^{+0.6}_{-0.3} \\times 10^{42}$\\ergs and in the faint state is  $0.6^{+0.2}_{-0.1} \\times 10^{42}$\\ergs. With these data, we find no significant difference in the power law slope or intrinsic absorption between the bright and faint nuclear states.} \\item{Nonstellar $8\\micron$ emission is concentrated in the central $\\sim 1$\\,kpc ($4\\farcs-5\\farcs$) of IC~4970, consistent with high obscuration in this region.  The mid-infrared colors of the nucleus, ${\\rm log}(S_{5.8}/S_{3.6})=0.13$, ${\\rm log}(S_{8.0}/S_{4.5})=0.6$, are consistent with the mid-infrared colors expected for highly obscured AGN candidate sources. Thus the X-ray and mid-infrared properties of the nucleus of IC~4970 suggest that it is a highly obscured Seyfert $2$. } \\item{Assuming that the $0.5-2$\\,keV X-ray emission observed in the nuclear region is all thermal, the spectrum is consistent with an absorbed $0.6^{+0.8}_{-0.4}$\\,keV APEC model with solar metallicity, a hydrogen absorbing column of $8 \\pm 4 \\times 10^{21}$\\cms, and intrinsic $0.5-10$ luminosity of $9 \\times 10^{39}$\\ergs. For uniform filling, the inferred mean electron density and hot gas mass in the nuclear region are $\\lesssim 0.2$\\cmc and $\\lesssim 3 \\times 10^{6}\\Ms$, respectively.} \\item{An ultra-luminous X-ray source with $0.5-10$\\,keV intrinsic luminosity of $\\gtrsim 3.5^{+1.8}_{-1.2} \\times 10^{39}$\\ergs is found $0.68$\\,kpc from IC~4970's nucleus.} \\item{ Little diffuse X-ray emission is observed outside the nuclear region. Assuming $0.54$\\,keV interstellar galaxy gas, the X-ray luminosity of diffuse gas outside the central $1$\\,kpc of the galaxy is $\\sim 7 \\times 10^{38}$\\ergs and gas mass is $\\sim 1-2 \\times 10^7\\Ms$. Either the gravitational potential of IC~4970 is too shallow to retain a significant hot gas halo, or that halo has been stripped by the interaction of IC~4970 with NGC~6872 and the Pavo group intra-group medium.} \\item{From the correlation between black hole mass and K-band luminosity, we expect IC~4970 to host a $5 \\times 10^7\\Ms$ black hole. Bondi accretion of hot gas in the nuclear region of IC~4970 onto the black hole can account for only $\\sim 5\\%$ of the observed X-ray luminosity of the AGN, such that the dominant power source for the active nucleus is likely the accretion of dust and cold gas driven into IC~4970's nuclear region by the galaxy's ongoing interaction with the gas-rich spiral galaxy NGC~6872.}  \\end{itemize}"
    },
    "0710/0710.0629_arXiv.txt": {
        "abstract": "We discuss three  specific modes of accretion disks around rotating magnetized neutron stars which may explain the separations of the kilo Hertz quasi periodic oscillations (QPO) seen in low mass X-ray binaries. The existence of  these modes requires that there be a maximum in the angular velocity of the accreting material, and that the fluid is in stable, nearly circular motion near this maximum rather than moving rapidly towards the star or out of the disk plane into funnel flows. It is presently not known if these conditions occur, but  we are exploring this with 3D magnetohydrodynamic simulations and will report the results elsewhere. The first mode is a corotation mode which is radially trapped in the vicinity of the maximum of the disk rotation rate and is unstable. The second mode, relevant to relatively slowly rotating stars, is a magnetically driven eccentric ($m=1$)  oscillation of the disk excited at a Lindblad radius in the vicinity of the maximum of the disk  rotation. The third mode, relevant to rapidly rotating stars, is a magnetically coupled eccentric ($m=1$) and an axisymmetric ($m=0$) radial disk perturbation which has an inner Lindblad radius also in the vicinity of the maximum of the disk rotation. We suggest that the first mode is associated with the upper QPO frequency, $\\nu_u$,  the second with the lower QPO frequency, $\\nu_\\ell =\\nu_u-\\nu_*$, and  the third with the lower QPO frequency, $\\nu_\\ell=\\nu_u-\\nu_*/2$, where $\\nu_*$ is the star's rotation rate. ",
        "introduction": "Low mass X-ray binaries often display twin kilo-Hertz quasi-periodic oscillations (QPOs) in their X-ray emissions (van der Klis 2006; Zhang et al. 2006). A wide variety of different models have been proposed to explain the origin and correlations of the different QPOs. These include the beat frequency model (Miller, Lamb, \\& Psaltis 1998; Lamb \\& Miller 2001; Lamb \\& Miller 2003), the relativistic precession model (Stella \\& Vietri 1999), the Alfv\\'en wave model (Zhang 2004), and warped disk models (Shirakawa \\& Lai 2002; Kato 2004).  A puzzling aspect of the some of the twin QPO sources considered in this work is that the  difference between the upper $\\nu_u$ and lower $\\nu_\\ell$ QPO frequencies is roughly either the spin frequency of the star $\\nu_*$ ($3$ cases where $\\nu_*=270,~330,~\\&~ 363$ Hz) or one-half this frequency, $\\nu_*/2$ ($4$ cases where $\\nu_*=401,~524,~581,~\\&~619$ Hz), for the cases where $\\nu_*$ is known, even though $\\nu_u$ and $\\nu_\\ell$ vary significantly (see, e.g., Zhang et al. 2006). A further type of behavior appears in the source Cir X-1 (Boutloukos et al. 2006), but this is not considered here. The cases where $\\nu_u-\\nu_\\ell \\approx \\nu_*$ may be explained by the beat frequency model (Miller et al. 1998), but the explanation of the cases where $\\nu_u-\\nu_\\ell \\approx \\nu_*/2$ is obscure.      Section 2.1 discusses the corotation instability, \\S 2.2  the exccentric ($m=1$) mode of the disk driven by the star's rotating magnetic field, and \\S 2.3 the coupled exccentric plus axisymmetric mode ($m=0~\\&~1$) also due to the star's rotating magnetic field. Section 4 gives the conclusions. ",
        "conclusions": "We discuss three  modes of accretion disks around rotating magnetized neutron stars which may explain the frequency separations of the twin kilo-Hertz QPOs seen in accreting X-ray binaries. The existence of  these modes requires that there be a maximum in the angular velocity of the accreting material and that the fluid is in stable, nearly circular motion near this maximum rather than moving rapidly towards the star or out of the disk plane into funnel flows. It is presently not known if these conditions occur, but  we are exploring this with 3D magnetohydrodynamic simulations and will report the results elsewhere. The first mode is a  corotation mode which is radially trapped in the vicinity of the maximum  of the disk rotation rate and is unstable. A simple dependence is assumed for the angular rotation rate of the disk $\\Omega_\\phi(r)$ which has a maximum at a radius $r_m $ outside the neutron star. The unstable mode has a frequency $\\omega_1$ somewhat less than max($\\Omega_\\phi$) and this can vary by a significant factor  depending on the state of the disk (e.g., the accretion rate). We suggest that this  mode is associated with the upper QPO frequency, $\\nu_u=\\omega_1/2\\pi$. The second mode  is a magnetically driven eccentric ($m=1$)  oscillation of the disk excited at the inner Lindblad radius which is in the vicinity of the maximum of the disk  rotation. The star's magnetic field is assumed to be the form discussed by Ruderman (2006). The Lindblad radii occur only for relatively slowly rotating stars, $\\nu_*<\\sim 380$ Hz. For these stars we suggest that the lower QPO frequency is $\\nu_\\ell =\\nu_u-\\nu_*$. The third mode, relevant to more rapidly rotating stars, is a magnetically coupled eccentric ($m=1$) and an axisymmetric ($m=0$) radial disk perturbation. It has an inner Lindblad radius also in the vicinity of the maximum of the disk rotation. For these stars the lower QPO frequency is $\\nu_\\ell=\\nu_u-\\nu_*/2$. A problem remaining for future work is the determination of the saturation amplitudes of the different modes."
    },
    "0710/0710.5526_arXiv.txt": {
        "abstract": "Absolute proper motions for six new globular clusters have recently been determined.  This motivated us to obtain the Galactic orbits of these six clusters both in an axisymmetric Galactic potential and in a barred potential, such as the one of our Galaxy. Orbits are also obtained for a Galactic potential that includes spiral arms.  The orbital characteristics are compared and discussed for these three cases. Tidal radii and destruction rates are also computed and discussed. ",
        "introduction": "Absolute proper motions have become available for six new globular clusters \\citep{CD07}, thus increasing to 54 the number of globular clusters of our Galaxy for which full space velocities and Galactic orbits can now be calculated. In a previous paper (Allen, Moreno \\& Pichardo 2006, hereafter Paper I) Galactic orbits were computed for 48 globular clusters, using both an axisymmetric and a barred Milky-Way-like potential.  We found that the effect of the bar was greatest for clusters with orbits residing mostly within the bar, and was negligible for the outermost clusters of our sample.  In the interim, \\citet{CD07} have obtained absolute proper motions for six additional clusters; they have also computed Galactic orbits in an axisymmetric potential. Since most of their cluster orbits appear to reside quite within the region of influence of the Galactic bar, it seems worthwhile to compute the orbits in a barred potential. Also, by computing orbits using different axisymmetric Galactic models some insight might be gained on the sensitivity of the results upon the potential model adopted.  In particular, \\citet{CD07} comment on the apparent ``pairing'' of the orbital parameters of NGC 2808 and NGC 4372, on one hand and NGC 4833, and NGC 5986 on the other. Since these ``pairings'' or kinematic groups may have important implications for the dynamical and merger history of our Galaxy (Kepley et al. 2007; Allen et al. 2007, and references therein), it is interesting to assess whether or not these similarities are sensitive to the Galactic model potential used.  In the present paper we compute Galactic orbits for the six new clusters, in the axisymmetric potential of \\citet{AS91}, in the barred model of \\citet{PMM04}, as well as in a potential model including spiral arms \\citep{PMME03}.  As before, we find that the orbits in the barred potential generally do not show secular changes in the total energy, $E$, or in the $z$-component of the angular momentum, $h$, both computed in the inertial Galactic frame.  Since the permanence of bar-like structures in galaxies is a matter of debate, we cannot be sure that the bar of our Milky has existed throughout the Galactic lifetime.  Therefore, we also ask ourselves how the orbits would look like if the bar has existed for only about a third of this lifetime.  To investigate the possible effects of spiral structure on the orbital characteristics of the clusters, we also include computations of the six globular clusters orbits in a Galactic potential that incorporates spiral arms. Among the six clusters selected for this computation, there is at least one cluster clearly belonging to the thick disk for which the effect of spiral perturbations is expected to be large (NGC 5927).  This paper is organized as follows. In Section \\ref{gpot} the Galactic potentials used to compute the cluster orbits are briefly described, and the initial conditions are presented.  In Section \\ref{galorb} the orbits obtained with both the axisymmetric and the barred potential are presented and compared with the ones obtained by \\citet{CD07}. In particular, we show that the effect of the bar is only negligible for one cluster NGC 3201, the most energetic one. We also discuss the effects on the orbits of a bar that has existed only during a fraction of the Galactic lifetime. In Section \\ref{spiral} the effect of spiral structure is presented for all six orbits. In Section \\ref{tirad} tidal radii are computed for orbits in the barred potential and compared with the axisymmetric case.  Destruction rates for these clusters are also computed and discussed. The final Section \\ref{conclusiones} presents a brief discussion and our conclusions. ",
        "conclusions": "We have obtained orbits for 6 additional globular clusters in both a barred and an axisymmetric Galactic potential, as well as in models including spiral arms. The total number of clusters with available orbits is now 54. Among the newly calculated cases, only the orbit of NGC 3201, an ``outer'' cluster, is unaffected by the bar. The other five clusters have orbits residing within the region of influence of the bar, and their orbits are clearly influenced by it. As in our previous study, we find that the main changes the bar causes in the orbits are larger vertical and radial excursions, and far more irregular orbits. In general, the bar causes no net global changes in the energy or the $z$-component of the angular momentum, computed in an inertial frame.  However, in four out of the six cases, jumps in these quantities do occur, even causing a temporary reversal of the sense of rotation of the orbit in the case of NGC 5986. The effect of a shorter-lived bar is found to be quite as noticeable on these orbits as that of the longer-lived one.  Orbits with spiral-arm perturbations were also computed. Contrary to expectations, even a small perturbation, accounting for only 5\\% of the local radial force ratio arm-disk, causes significant changes in the form of the orbits. Although the influence of the spiral arms on the orbits of the clusters closest to the Galactic disk are not negligible, the long-term effects are definitely dominated by the Galactic bar.  Tidal radii have been computed with the expression derived in Paper I, as well as with a numerical evaluation of the relevant quantities along the orbit. Again, we found little change due to the bar. When changes did occur, they again make the computed tidal radii somewhat larger in the presence of a bar. With the new material on hand, we confirm our earlier finding that the destruction rates due to shocks with the Galactic bulge and disk are not strongly affected when we consider a barred Galactic potential, or a barred potential with spiral arm perturbations."
    },
    "0710/0710.0904_arXiv.txt": {
        "abstract": "This paper explores the quantitative connection between globular clusters  and the ``diffuse'' stellar population of the galaxies they are associated with. Both NGC 1399 and NGC 4486 (M87)  are well suited for this kind of analysis due to their large globular cluster populations.  The main assumption of our Monte Carlo based models is that each globular cluster is formed   along with a given diffuse stellar mass that shares the same spatial distribution, chemical composition and age. The main globular clusters subpopulations, that determine the observed bimodal colour distribution, are decomposed avoiding {\\it a priori} parametric (e.g. Gaussian) fits and using a new colour (C-T$_1$)-metallicity relation. The eventual detectability of a ``blue'' tilt in the colour magnitude diagrams of the blue globulars subpopulation is also addressed.  A successful link between globular clusters and the stellar galaxy halo is established by assuming that the number of globular clusters per associated diffuse stellar mass t is a function of total abundance [Z/H] and behaves as $t=\\gamma \\exp(-\\delta[Z/H])$ (i.e. increases when abundance decreases).  The simulations allow the prediction of a surface brightness profile for each galaxy through this two free parameters approximation. The $\\gamma$, $\\delta$  parameters that provide the best fit to the observed profiles in the B band, in turn, determine several features, namely, large scale halo colour gradients, globular clusters-halo colour offset, clusters cumulative specific frequencies, and  stellar metallicity distributions, that compare well with observations.  The results suggest the coexistence of two distinct stellar populations characterised by widely different metallicities and spatial distributions. One of these populations (connected with the blue globulars) is metal poor, highly homogeneous, exhibits an extended spatial distribution and becomes more evident at large galactocentric radius contributing with some 20\\% of the total stellar mass. In turn, the stellar population associated with the red globulars is extremely heterogeneous and dominates the inner region of both galaxies.  Remarkably, and although the cluster populations of these galaxies exhibit detectable differences in colour distribution, the $\\delta$ parameter that determines the shape  of the brightness profiles of both galaxies has the same value, $\\delta \\approx$ 1.1 to 1.2 $\\pm 0.1$. ",
        "introduction": "\\label{INTRO}  The idea that globular clusters (GCs) harbour important clues in relation with the early stages of galaxy formation is a widely accepted concept. One of the most compelling arguments in favour of the existence of a connection between GCs and major star formation episodes  in the life of a galaxy is the constant cluster formation efficiency, defined in terms of total baryonic mass \\citep{b46}, in different galaxies.  However, breaking the code that leads to a detailed  quantitative link between GCs and the underlying ``diffuse'' stellar population is still an open question. Such a connection has been discussed on theoretical (e.g. \\citealt*{b3} or, more recently, \\citealt{b56b}) and observational grounds (e.g. \\citealt{b20}). In the particular case of Milky Way GCs, \\citet*{b57b} found that the chemical similarities between clusters and field stars with $[Fe/H]\\le -1$ suggests a shared chemical history in a well mixed early Galaxy.  Clarifying this issue may certainly yield some arguments in favour (or against) some predominant ideas that have been widely referenced in the literature (e.g. \\citealt{b16}; \\citealt{b72b}) and later explored within the frame of different scenarios (e.g. \\citealt{b1}, or \\citealt{b18}).  A good perspective of the complex situation in this context is given in the thorough review by \\citet{b5} and \\citet{b41}.  An initial confrontation between GCs and halo stellar populations shows more differences than similarities: a) In general, GCs exhibit more shallow spatial distributions than those characterising galaxy light (e.g. \\citealt{b59}; \\citealt{b14}); b) There is a colour offset in the sense that mean integrated globular colours appear bluer than those of the galaxy halos at the same galactocentric radius (\\citealt{b78}; \\citealt*{b22}; \\citealt{b37b}); c) GCs show frequently  bimodal colour (and hence, metallicity) distributions (see, for example, \\citealt{b53}). This feature does not seem exactly shared by the resolved stellar populations in nearby resolved galaxies (see \\citealt{b15}; \\citealt{b33}; \\citealt{b62} or \\citealt{b50}).  As discussed later in this work, those differences arise, mainly, from the fact that GCs analysis usually provide {\\bf number} weighted statistics while galaxy halos observations yield  {\\bf luminosity} weighted measurements.  A preliminary quantitative approach to the globulars-stellar halo connection was presented in \\citet*{b25} (hereafter FFG05). This last paper shows that, given the areal density distribution of the ``blue'' and ``red'' globular cluster subpopulations in NGC 1339, the galaxy   surface brightness profile, galactocentric colour gradient and  cumulative GCs specific frequency, can be matched by linearly weighting the areal density profiles. The ``weight'' corresponding to each component of the brightness profile is the inverse of the {\\bf intrinsic} GCs frequency characteristic of each cluster population.  The main argument behind that approach is that the shape of the colour (and metallicity) distribution of each globular cluster subpopulation does not change with galactocentric radius. Large angular scale studies (\\citealt{b14}; \\citealt{b2}) in fact show that the colour peaks in the GCs colour statistics of NGC 1399 keep the same position (or show very little variation) over large galactocentric ranges. A similar result is obtained for NGC 4486 by \\citet{b40} who find those peaks do not show a detectable variation in colour over 75 kpc in galactocentric radius.  It must be stressed that those subpopulations are ``phenomenologically'' defined in terms of their integrated colours but each might eventually have  a given spread in age and/or metallicity.  The presence of a ``valley'' in the globular colour statistics is usually adopted as a discriminating boundary between both subpopulations. The need to revise such a procedure was already suggested by figure 5 in FFG05. This last diagram showed that the NGC 1399 GCs bluer than the blue peak have a distinct behaviour of the areal density profile, exhibiting a flat core that disappears when all ``blue'' clusters (i.e. all GCs bluer than the colour valley) are included in the sample. That result prompted for a further analysis as discussed below.  This paper generalises the FFG05 approach trough  Monte-Carlo based models. In this frame, ``seed'' globulars are generated following a given abundance Z distribution and then associated with a ``diffuse'' stellar mass that shares its age, chemical composition and statistical spatial distribution. The luminosity associated with this mass is derived from a mass to luminosity ratio adequate for a given age and metallicity. These models aim at reproducing the features mentioned above and seek for a function that could link GCs and diffuse stellar populations keeping a minimum number of free parameters.  Both NGC 1399 and NGC 4486 appear as adequate targets in order to perform such a modelling due to their prominent globular cluster systems (GCS). Although these systems show some structural similarities in terms of their spatial distribution, they also differ markedly both in the shape of their GCs colour statistics  and specific frequencies \\citep{b24}.  This work also presents new Washington photometry, obtained and handled in a homogeneous way, that allows for a re-discussion of the GCS properties in the inner region of both galaxies and, in particular, of the behaviour of the areal density of GCs with colour. In turn, recent wide field photometric studies of the GCS associated with NGC 1399 \\citep{b2} and NGC 4486 (\\citealt{b80}; \\citealt{b81}), are well suited for extending the analysis to larger galactocentric radii. ",
        "conclusions": "\\label{Conclu}  The results presented in previous sections show that the surface brightness profiles of both NGC 1399 and NGC 4486 can be traced using a common link between GCs and the stellar halo populations in these galaxies. This link imply that the number of GCs per diffuse stellar mass, defined as {\\bf $t=\\gamma \\exp(-\\delta[Z/H])$ },increases when chemical abundance decreases. We note that \\citet{b33} had already found an increase of Sn with decreasing GCs metallicity in the case of NGC 5128.   This suggests that, on  a large scale, the dominant globular cluster subpopulations formed along major star forming episodes and following a similar pattern. However, it is not yet clear whether abundance, through the role that it plays in the t parameter, is the physical reason that governs the fraction of clustered to diffuse stellar mass or, eventually, some other ``hidden'' variable in turn correlated with abundance.  The quality of the profile fits is comparable to any other parametric approximation in a range that covers  a galactocentric radius from 10 to 100 kpc. In the inner regions, GCs fail to map the brightness distribution probably as a consequence of cluster destruction processes due to gravitational effects. However, it seems that survivor clusters, with large perigalacticon orbits, are still able to trace their associated stellar populations. It is worth mentioning that, based on far UV observations of NGC 1399 , \\citet{b42}, find arguments that support the coexistence of two widely different stellar populations, in terms of chemical abundance, in the nucleus of the galaxy.  Even though the approach only aims at reproducing the brightness profiles, other connected features as galactocentric colour gradients, GCs-halo colour offset, cumulative cluster specific frequency and inferred stellar metallicity distributions compare very well with observations.  It seems also remarkable that the common quantitative GCs-stellar halo link works in both galaxies although their cluster populations exhibit detectable differences in cluster numbers and chemical abundance. Both GCs subpopulations have larger abundance scale lengths (as defined in Section 4) in NGC 1399 than in NGC 4486 leading to total mean abundance ratios of 1.4 and 2.5 for the red and blue GCs, respectively.  As shown in Figures 10 and 14, the approach presented in this paper delivers GCs colour distributions that are not strongly different  from a two Gaussians fit requiring five free parameters, a common  procedure in the literature (e.g Ostrov, Forte \\& Geisler 1998). However, we note that our models indicate a lower ratio of the number of blue to red GCs and only require three free parameters.  Blue GCs in NGC 4486 have a  small abundance scale lenght, about four times smaller than that of the blue GCs in NGC 1399, and we speculate that this may be connected with a shorter formation time scale. In turn, that lower abundance spread may be the reason behind the presence of a blue tilt, connected with cluster mass, in the colour magnitude diagram. This feature seems absent, or probably masked, by the larger abundance scale of the blue GCs in NGC 1399, a situation that may also hold in other galaxies (e.g. NGC 4472, Strader et al. 2006).  This last result argues in favour of the idea that blue GCs ``know'' about the galaxy they are associated with (\\citealt*{b77b}). In any case, and in both galaxies, the blue GCs barely reach an abundance close to [Z/H]=-0.5. The reason for this upper metallicity cut off may be connected with some kind of sincrhonising event as the re-ionisation of the Universe (\\citealt{b11}; \\citealt{b67}; \\citealt{b63}). However, the different abundance scale lengths of the blue GCs also suggest that a local phenomenon, distinct for each galaxy, may have also played a role in modulating the the star formation rate (e.g. the onset of galaxy nuclear activity).  The overall picture seems consistent with some scenarios already discussed in the literature (\\citealt{b18}; \\citealt{b34}) that invoke two different cluster formation mechanisms and, probably, environmental conditions.  In contrast with the relative abundance homogeneity and large spatial (half density) core radii ($\\approx$ 25 kpc) of the blue GCs, the red ones exhibit a large abundance heterogeneity and much smaller core radii ($\\approx$ 5 kpc) also shared by the red diffuse stellar population. This degree of heteregeneity may be connected , for example, to mergers of different nature (e.g. \\citealt{b72}).  The total globular cluster formation efficiencies, in terms of stellar mass, indicated by the models, and adopting an average cluster mass of $2.5 \\times 10^{5 } M_\\odot $, are $2.4 \\times 10^{-3}$ for NGC 1399 and $5.6 \\times  10^{-3}$ for NGC 4486. They are comparable to the efficiency derived by \\citet{b46} although, in that case, the definition of efficiency included total baryonic mass.  Blue clusters  show a higher formation efficiency (in terms of the stellar mass they are associated with), when compared with the red ones, probably as a consequence of a lower star formation efficiency during the early phases of galaxy formation at a low metallicity regime.  Although there is no strong evidence of an age difference between the blue and red globular subpopulations (e.g. \\citealt{b37}) within the uncertainties of the measurements, a possible temporal sequence that assumes the formation of the blue population first, cannot be discarded as a result of the relatively small time scales involved at the early phases of galaxy formation (see \\citealt{b3b}). In this frame, the chemical enrichment provided by a presumably progenitor blue population might be important to boost later stellar formation efficiency through an abundance enrichment that may reach $[Z/H] \\approx $-0.60 within 100 kpc of the galaxy nucleus.  Some  scenarios suggest that blue GCs formation is associated with dark matter (\\citealt{b3}; \\citealt{b49}; \\citealt{b57}), and it is tempting to look for such a connection. For example, the results listed in Table \\ref{Model_values} show that while both galaxies have a similar total number of red GCs, NGC 4486 outnumbers NGC 1399 in a factor of about 1.8 in terms of blue GCs. Dark mass estimates within a galactocentric radius of 100 kpc are 3.4$\\times 10^{12} M_\\odot$ for NGC 1399 (extrapolating data from \\citealt{b64}) and 7.4$\\times 10^{12} M_\\odot$ for NGC 4486 \\citep{b10b}, leading to a ratio $\\approx 2.0$ comparable to that in the number of blue GCs.  The similarity of the stellar galaxy masses, and the difference in their total masses, had already been noticed by \\citet{b36} on the basis of their X ray analysis.  FFG05  found that, adopting {\\bf their} definition of blue clusters, the density profiles of the NGC 1399 GCs could be fit with a NFW profile \\citep{b51} with a scale length of 375 arcsec, coincident with that derived for the inferred dark matter halo by \\citet{b64}. \\citet{b80} also perform  a NFW profile fit to the blue clusters in NGC 4486. However, both approaches deserve a revision since, on the basis of the results presented in this work, the ``genuine'' blue GCs exhibit a rather extended inner core in their surface density profiles. As shown here, these cores have been disguised by the inclusion of the blue tail of the red subpopulation within the ``blue'' GCs sample. This overlapping should be even more severe when using colour indices less sensitive to metallicity than (C-T$_1$) and should be taken into account when doing, for example, kinematic analysis of the cluster subpopulations. \\\\"
    },
    "0710/0710.2130_arXiv.txt": {
        "abstract": "Externally Dispersed Interferometry (EDI) is the series combination of a fixed-delay field-widened Michelson interferometer with a dispersive spectrograph.  This combination boosts the spectrograph performance for both Doppler velocimetry and high resolution spectroscopy.  The interferometer creates a periodic spectral comb that multiplies against the input spectrum to create moir\\'e fringes, which are recorded in combination with the regular spectrum.  The moir\\'e pattern shifts in phase in response to a Doppler shift.  Moir\\'e patterns are broader than the underlying spectral features and more easily survive spectrograph blurring and common distortions.  Thus, the EDI technique allows lower resolution spectrographs having relaxed optical tolerances (and therefore higher throughput) to return high precision velocity measurements, which otherwise would be imprecise for the spectrograph alone. ",
        "introduction": "Externally Dispersed Interferometry (EDI) is a new technology for precision Doppler velocimetry that is both more efficient and less expensive than traditional high resolution spectrographs \\cite{ErskinePatSuperimpose,ErskinePatEDI2002,ErskineEDITheory2003,SPIEscot,HiResSPIE,UVDiegoSPIE,TediOrlandoSPIE,ResBoostApJ2003,TediDiegoSPIE,OrlandoNoiseSPIE,ErskineGe2000,G.E.R.2002,GeVirgoApJ2006,Ge2002,Ge2003,Ge51Peg,VanEykenAAS2006}.  This technique increases the responsivity of a low-dispersion spectrograph to sharp spectral features which would normally be unresolved.  The first exoplanet results are arriving from these instruments as they are starting to come online. An EDI recently discovered a new exoplanet around the star HD 102195 in Virgo\\cite{GeVirgoApJ2006}. A multi-object EDI is being tested at the 2.5 m Sloan telescope\\cite{VanEykenAAS2006}.  An NSF funded project is underway to field an EDI at the cassegrain output of the Mt.~Palomar Observatory 200 inch telescope in series with the TripleSpec near infrared (0.8--2.4 $\\mu m$) spectrograph being built by Cornell University \\cite{TripleSpecSPIE}.  This will find low mass exoplanets around low mass stars\\cite{TediDiegoSPIE,TediOrlandoSPIE,OrlandoNoiseSPIE}.  An EDI is the series combination of a fixed-delay interferometer with a dispersive spectrograph (Fig.~\\ref{fig:f1}).  The interferometer generates fringes which shift in phase in response to a Doppler velocity, and the disperser separates the fringes of different wavelengths on the detector array so that they can be detected at high visibility.  The advantages of an EDI over high-resolution spectroscopic velocimetry are three-fold:   \\begin{figure}[!tbh] \\center{\\includegraphics[height=7cm]{pre-idea_Queb4_I3.eps} \\includegraphics[width=7.5cm]{Dual_Output_overlay4c_cr.eps}} \\caption[Fig1] {% [Left] EDI concept.  A sinusoidal transmission comb created by the interferometer is multiplied against the stellar spectrum, adding moir\\'e fringes to the ordinary spectrum.  The moir\\'e and ordinary components are separated during analysis. The moir\\'e fringes change phase in proportion to the star's Doppler velocity. [Right] Example EDI apparatus, having dual outputs for high efficiency. \\label{fig:f1} \\label{fig:f2}} \\end{figure}   \\begin{enumerate} \\item {\\bf The cost, weight, and size of the spectrograph and detector are greatly reduced.}  Much lower resolution spectrographs can be used to return precision Doppler velocities.  For example, the EDI used to detect the planet around HD 102195 had a resolution\\cite{GeVirgoApJ2006} of only $\\lambda / \\Delta \\lambda = 5000$. \\item {\\bf The throughput and sensitivity of an EDI is increased.}  Lower resolution spectrographs usually have better throughput; the EDI preserves this property.  For example, the Exoplanet Tracker on the 2.1 meter Kitt Peak telescope can perform precision radial velocity (RV) measurements on sources as faint as 10th magnitude (15 minutes to achieve ~8 m/s for V=8, Ref.~\\citenum{ETnoaoWebpage2006}); this compares favorably to HIRES on the 10 meter Keck telescopes!  Wider bandwidths can be used using the interferometer signal as a wavelength fiducial linking spectral features across the band.  The increased number of spectral lines covered improves RV precision and sensitivity. Smaller telescopes can be used for a given target brightness, which often have more time available for higher cadence observations.  Alternatively, more targets are accessible to large telescopes, possibly including the faint stars targeted for deep transit searches, including the Kepler field. \\item {\\bf The EDI measurement is more robust to systematic and instrumental effects.}  The stepping manner of measuring moir\\'e pattern phases eliminates fixed pattern errors such as those due to detector or optics blemishes by acting as a built-in flatfielding---they are not synchronous with the stepping.  Common spectrograph irregularities, such as drifts in the focal spot position and diameter on the detecting array, do not strongly affect the EDI observable of fringe phase.  This is due to the differential nature of a moir\\'e pattern in comparing the stellar and interferometer sinusoidal combs.  The interferometer comb can be thought of as a set of fiducials coupled to the science signal and follow it through the instrument.  Distortions affecting the science signal are also applied to the interferometer comb.  We have estimated that the EDI is 1000 to 10,000 times more robust to translation of the focal spot, and 10 to 100 times more robust to change in focal spot diameter. \\end{enumerate} ",
        "conclusions": ""
    },
    "0710/0710.0135_arXiv.txt": {
        "abstract": "Assuming the dark matter is made entirely from neutralinos, we re-visit  the role of their  annihilation on the temperature of diffuse gas  in the high redshift universe. We consider  neutralinos  of particle mass $36 \\;  {\\rm GeV}$ and  $100\\;  {\\rm GeV}$, respectively. The former is able to produce $\\sim 7$  $\\ee$ particles per annihilation through  the fremionic channel, and the latter $\\sim 53$ particles assuming a purely bosonic channel.    High energy $\\ee$ particles up-scatter the Cosmic Microwave Background (CMB) photons into higher energies via the inverse-Compton scattering. The process produces a power-law  $\\ee$  energy spectrum  of index $-1$  in the energy range of interest, independent of the initial energy distribution. The corresponding energy spectrum of the up-scattered photons is a power-law of index $-1/2$, if absorption by the gas is not included. The scattered photons photo-heat  the gas by releasing electrons which deposit  a fraction ($~14\\%$) of their energy as heat into the ambient medium. For uniformly distributed neutralinos the heating  is insignificant. The effect is greatly enhanced by the  clumping of neutralinos into dense haloes. We use a time-dependent  clumping model which takes into account the damping of density fluctuations  on mass scales smaller than $\\sim 10^{-6}M_\\odot$. With this clumping model, the heating mechanism  boosts  the gas temperature above that of  the CMB after a redshift of $z\\sim 30$. By $z\\approx 10$ the gas temperature is  nearly 100 times its  temperature when no  heating is invoked. Similar increase is obtained for the two neutralino masses considered. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.2901_arXiv.txt": {
        "abstract": "We investigate the evolution of angular momentum in simulations of galaxy formation in a cold dark matter universe. We analyse two model galaxies generated in the N-body/hydrodynamic simulations of Okamoto et al. Starting from identical initial conditions, but using different assumptions for the baryonic physics, one of the simulations produced a bulge-dominated galaxy and the other one a disc-dominated galaxy. The main difference is the treatment of star formation and feedback, both of which were designed to be more efficient in the disc-dominated object. We find that the specific angular momentum of the disc-dominated galaxy tracks the evolution of the angular momentum of the dark matter halo very closely: the angular momentum grows as predicted by linear theory until the epoch of maximum expansion and remains constant thereafter. By contrast, the evolution of the angular momentum of the bulge-dominated galaxy resembles that of the central, most bound halo material: it also grows at first according to linear theory, but 90\\% of it is rapidly lost as pre-galactic fragments, into which gas had cooled efficiently, merge, transferring their orbital angular momentum to the outer halo by tidal effects. The disc-dominated galaxy avoids this fate because the strong feedback reheats the gas which accumulates in an extended hot reservoir and only begins to cool once the merging activity has subsided. Our analysis lends strong support to the classical theory of disc formation whereby tidally torqued gas is accreted into the centre of the halo conserving its angular momentum. ",
        "introduction": "In the standard theory of galaxy formation in the cold dark matter (CDM) cosmogony, disc galaxies form when rotating gas cools inside a dark matter halo and fragments into stars. Analytic and semi-analytic calculations of this process start by assuming that tidal torques at early times impart the same specific angular momentum to the gas and the halo and that the angular momentum of the gas is conserved during the collapse \\citep{fall,momaw98,cole00,fa00}. It is unclear whether this assumption holds in gasdynamic simulations of galaxy formation.  The first attempts to simulate the formation of a spiral galaxy from cold dark matter initial conditions generally failed, producing objects with overly centrally concentrated distributions of gas and stars. It became immediately apparent that the root cause of this problem was a net transfer of angular momentum from the baryons to the dark matter halo during the aggregation of the galaxy through mergers \\citep{NavBe,NavWhi,NavFreWhi}. This is known as the \\qm{angular momentum problem}.  It was suspected from the start that its solution was likely to involve feedback processes that would regulate the supply of gas to the galaxy.  More recent simulations within the CDM framework have produced more promising disc galaxies. \\citet{tc01} obtained a reasonably realistic disc at $z=0.5$ (when their simulation stopped) by assuming that gas cooling is strongly suppressed by feedback effects. \\citet{SN02}, using different prescriptions for star formation and feedback, found that a broad range of galaxy morphologies could be produced. In related work, \\citet{Abadi03} obtained a disc galaxy with a realistic density profile but not enough angular momentum. \\citet{sgp03} found that considerably larger energy feedback rates at early times than previously employed were required to ameliorate the angular momentum problem. They achieved this by distinguishing between {\\it early} and {\\it late} star formation modes, assuming a higher star formation efficiency and stronger feedback in the early mode. They were able to generate a variety of morphological types, including discs whose angular momentum, however, was still a factor of two smaller than observed for real disc galaxies.  \\citet{gov04} produced a disc galaxy without requiring strong feedback and ascribed the angular momentum problem to a lack of numerical resolution. However, subsequently, \\citet{gov06} were forced to invoke strong feedback in order for their simulated discs to lie closer to the Tully-Fisher relation. While resolution is clearly an important and complicated issue, \\cite{Oka1} argued that if the hot and cold gas components are decoupled from each other, then the effects of resolution in the current generation of simulations are reduced. \\citet{rob04} adopted a multiphase model for the star-forming interstellar medium (ISM) which stabilises gaseous discs against the Toomre instability, and produced a galaxy with an exponential surface brightness profile but insufficient angular momentum in the inner parts. They concluded that the fragmentation of gaseous discs caused by the Toomre instability is the main cause of the angular momentum loss in simulations that usually assume an isothermal ISM at a temperature of $\\sim 10^4$K.  \\citet{Oka2} also adopted a multiphase model for the ISM and assumed a top-heavy initial mass function (IMF) for stars formed in starbursts. The latter assumption was motivated by the semi-analytic model of \\citet{Baugh05} which requires a large relative abundance of massive stars in order to match the number counts of bright submillimeter galaxies. A top-heavy IMF is further supported by the good agreement of this model with the metal abundances measured in the intracluster medium \\citep{nag05a} and in elliptical galaxies \\citep{nag05b}.  By varying the criteria for starbursts in their simulations, \\citet{Oka2} were able to produce galaxies with a variety of morphological types, from ellipticals to spirals, starting from exactly the same initial conditions. They showed that disc galaxies can form in halos that have a relatively quiet merger history if the early collapse of baryons is inhibited by strong feedback, in this case generated by the evolution of a stellar population with a top-heavy IMF.  In this paper, we analyse two of the galaxies simulated by \\citet{Oka2}, one which ended up as a bulge-dominated galaxy and another that formed an extended stellar disc. We trace the time evolution of the angular momentum of the dark matter and the gas, and we distinguish between the evolution of the inner part of the dark halo and the halo as a whole. We show explicitly that mergers transfer angular momentum from the inner parts of the system to the outer halo. This process is one aspect of the redistribution of angular momentum during virialization investigated in N-body simulations by \\citet{don06} and \\citet{donghia07}. Thus, the spin of the {\\em central} part of the final system contains a fossil record of the merger history of the object.  This paper~is organised as follows. The simulations are briefly described in Section~2. The method we use to follow the angular momentum of the different components is explained in Section~3 where we also present our results. We conclude in section 4. ",
        "conclusions": "The inability to obtain realistic discs by the present day in simulations from cold dark matter initial conditions has been ascribed to various possible causes: (i) a lack of numerical resolution \\citep{kau07, gov04, gov06}, (ii) an artificial transfer of angular momentum from a cold disc to a hot surrounding atmosphere \\citep{Oka1}, (iii) instabilities in a heavy disc which, by causing fragments to form, end up transporting angular momentum inwards \\citep{rob04}, and (iv) the transfer of angular from merging substructures to the outer halo \\citep{NavBe,NavWhi,NavFreWhi}. The first three of these four considerations are unimportant in our simulations. \\cite{Oka1} carried out a convergence study which suggests that, so long as the cold and hot gas phases are dynamically decoupled, resolution effects are not important in calculations the size of those we analyse here. Such decoupling also suppresses process (ii). The adoption of an effective multiphase equation-of-state stabilizes a disc against gravitational instability thus alleviating (iii). Process (iv) is inevitable and is the dominant form of angular momentum transport (at least for the dark matter) in our simulations.  The series of N-body/SPH simulations of galaxy formation in a CDM universe carried out by \\citet{Oka2} led to the conclusion that galaxies with very different morphologies can form in the same dark matter halo depending on the details of the assumed star formation and feedback prescriptions. In this paper, we have investigated the evolution of the specific angular momentum as a possible explanation for why relatively small differences in the assumed baryonic processes can result in such different outcomes. We have analyzed the bulge-dominated and the disc-dominated galaxies simulated (from identical initial conditions) by \\citet{Oka2}. We trace back all the dark matter particles that end up within the virial radius of the halo ($r_v$), as well as those that end up in the central parts of the halo by virtue of being in the top 10 percentile of the binding energy distribution. We also trace back the baryonic particles that end up as a galaxy, i.e. as cold baryons (star particles and dense gas) within 10\\% of $r_v$ at $z = 0$.  The angular momentum of the dark matter halo evolves almost identically in the two simulations. The specific angular momentum of the halo as a whole grows initially as predicted by tidal torque theory in the linear regime, in proportion to $a^{3/2}$ \\citep{White, ct96}. This phase lasts until maximum expansion is reached. Thereafter, the angular momentum of the dark matter remains constant, as the system collapses to a virialized configuration. By contrast, the evolution of the most bound 10\\% of the dark matter is strikingly different. In the linear phase this material behaves just as the halo as a whole but, after maximum expansion, its specific angular momentum rapidly declines to $\\sim 10\\%$ of its maximum value. The decline occurs in episodes associated with the mergers that build up the central halo. Most of the angular momentum of these fragments is invested in their orbits and is transferred to the outer halo by tidal forces as the fragments sink by dynamical friction.  The baryonic component evolves very differently in the two simulations. In essence, in the disc-dominated galaxy, the specific angular momentum of the baryons tracks that of the halo as a whole, while in the bulge-dominated galaxy, it tracks the evolution of the central halo material. The reason for the difference can be traced back to the different strengths of feedback in the two galaxies at early times. In the bulge-dominated galaxy, feedback is weak and most of the baryons rapidly cool and condense into stars within sub-galactic fragments. As these halo fragments merge and give up their angular momentum to the outer halo, so do the baryon clumps within them. The result is a slowly-rotating central bulge. A small amount of residual gas rains in later on forming a small disc. In the disc-dominated galaxy, by contrast, feedback is strong at early times, the gas is reheated before it can make a substantial mass of stars and, instead, accumulates in an extended hot reservoir which acquires a specific angular momentum similar to that of the dark halo. By the time this gas is able to cool, much of the merger activity has subsided and so the gas is able to dissipate and collapse into an centrifugally supported disc configuration conserving its angular momentum. Not only does the angular momentum of the baryonic material track the evolution of the angular momentum of the halo, but the actual values of the specific angular momentum of both components are very similar throughout the entire history of the galaxy. By contrast, the specific angular momentum of the bulge-dominated galaxy is only about 25\\% of the halo value.  Our results confirm that the key to the formation of disc galaxies in the cold dark matter cosmology is the suppression of the early collapse of baryons into small, dense halos. By contrasting the different evolutionary histories of the angular momentum in bulge- and disc-dominated galaxies, our analysis explicitly reveals how regulating the rate at which gas is supplied determines the final morphology of the galaxy. A disc galaxy results when, as seen in Fig.~\\ref{bfrac}, a large fraction of the baryon content of the halo is ejected by winds at early times. The ejected baryons (some of which are recaptured later) are thus decoupled from the subhalos that are destined to merge and join a hot gas reservoir which is tidally torqued in the same way as the host halo and thus acquires the same net specific angular momentum. Our results provide strong support for the classical theory of disc formation by \\citet{fall} and \\citet{momaw98} in which tidal torques impart the same specific angular momentum to the dark matter and the gas and angular momentum is conserved during disc formation."
    },
    "0710/0710.3136_arXiv.txt": {
        "abstract": "This paper investigates the clustering properties of a complete sample of 1041 24$\\mu$m-selected sources brighter than $F_{24\\mu \\rm m}=400\\mu$Jy in the overlapping region between the SWIRE and UKIDSS UDS surveys. With the help of photometric redshift determinations we have concentrated on the two interval ranges $z=[0.6-1.2]$ ({\\it low-z} sample) and $z\\ge 1.6$ ({\\it high-z} sample) as it is in these regions were we expect the mid-IR population to be dominated by intense dust-enshrouded activity such as star formation and black hole accretion. Investigations of the angular correlation function produce an amplitude $A\\sim 0.010$ for the high-z sample and $A\\sim 0.0055$ for the low-z one. The corresponding correlation lengths are $r_0=15.9^{+2.9}_{-3.4}$~Mpc and $r_0=8.5^{+1.5}_{-1.8}$~Mpc, showing that the high-z population is more strongly clustered. Comparisons with physical models for the formation and evolution of large-scale structure reveal that the high-z sources are exclusively associated with very massive ($M\\simgt 10^{13} M_\\odot$) haloes, comparable to those which locally host groups-to-clusters of galaxies, and are very common within such (rare) structures. Conversely, lower-z galaxies are found to reside in smaller halos ($M_{\\rm min}\\sim 10^{12} M_\\odot$) and to be very rare in such systems.  On the other hand, mid-IR photometry shows that the low-z and high-z samples include similar objects and probe a similar mixture of AGN and star-forming galaxies. While recent studies have determined a strong evolution of the 24$\\mu$m luminosity function between $z\\sim 2$ and $z\\sim 0$, they cannot provide information on the physical nature of such an evolution. Our clustering results instead indicate that this is due to the presence of different populations of objects inhabiting different structures, as active systems at $z\\simlt 1.5$ are found to be exclusively associated with low-mass galaxies, while very massive sources appear to have concluded their active phase before this epoch. Finally, we note that the small-scale clustering data seem to require steep ($\\rho\\propto r^{-3}$) profiles for the distribution of galaxies within their halos. This is suggestive of close encounters and/or mergers which could strongly favour both AGN and star-formation activity. ",
        "introduction": "Understanding the assembly history of massive spheroidal galaxies is a key issue for galaxy formation models.  The ``naive\" expectation from the canonical hierarchical merging scenario, that proved to be remarkably successful in explaining many aspects of large-scale structure formation, is that massive galaxies generally form late and over a long period of time as the result of many mergers of smaller haloes. On the other hand, there is quite extensive evidence that massive galaxies may form at high redshifts and on short timescales (see, e.g. Cimatti et al. 2004; Fontana et al. 2004; Glazebrook et al. 2004; Treu et al. 2005; Saracco et al. 2006; Bundy et al. 2006; McLure et al. 2006), while the sites of active star formation shift to lower mass systems at later epochs, a pattern referred to as \"downsizing\" (Cowie et al. 1996; Heavens et al. 2004). In order to reconcile the observational evidence that stellar populations in large spheroidal galaxies are old and essentially coeval (Ellis et al. 1997; Holden et al. 2005) with the hierarchical merging scenario, the possibility of mergers of evolved sub-units (``dry mergers'') has been introduced (van Dokkum et al. 2005; Naab et al. 2006). This mechanism is however strongly disfavoured by studies on the evolution of the stellar mass function (Bundy et al. 2006).  Key information, complementary to optical/IR data, has come from sub-millimeter surveys (Hughes et al. 1998; Eales et al. 2000; Scott et al. 2002; Knudsen et al. 2006; Coppin et al. 2006) which have found a large population of luminous sources at substantial redshifts (Chapman et al. 2005). However, the interpretation of this class of objects is still controversial (e.g. Granato et al. 2004; Kaviani et al. 2003; Baugh et al. 2005). The heart of the problem are the masses of the objects: a large fraction of present day massive galaxies already assembled at $z\\sim 2-3$ would be extremely challenging for the standard view of a merging-driven  growth.   Measurements of clustering amplitudes are a unique tool to estimate halo masses at high $z$, but complete samples comprising at least several hundred of sources are necessary.  Recently, Magliocchetti et al. (2007) have reported evidence for strong clustering for $\\sim 800$ optically very faint ($R>25.5$), $F_{24\\mu \\rm m}\\ge 0.35$~mJy sources obtained from the Spitzer first cosmological survey (First Look Survey -- FLS; Fadda et al. 2006). Both the clustering properties and the counts of such sources are consistent with them being very massive proto-spheroidal galaxies in the process of forming most of their stars. Furthermore, by assuming a medium redshift $z\\sim 2$, their comoving number density appears to be much higher than what expected from most semi-analytic models.  The Magliocchetti et al. (2007) work however suffers from the lack of information on the redshift distribution of the optically faint Spitzer-selected sources and has to rely on models based on both template spectral energy distributions and on theoretical investigations of the issue of galaxy formation and evolution in order to go from the observed projected clustering signal to the more meaningful results in real space.  The optical-to-mid-IR depth of the UKIDSS data allows us to overcome this problem. Photometric estimates for the overwhelming majority ($\\sim 97$\\%) of 24$\\mu$m-selected sources are now available (Cirasuolo et al. 2007). Furthermore, despite the somehow poor statistics, the redshift information can also allow us for the first time to compare the clustering signal of similar sources at different epochs so to investigate possible differences and evolution in their large-scale properties. To this aim, we will concentrate on two samples, the first one which includes galaxies in the $z=[0.6-1.2]$ range, and a second one made by sources with $z\\simgt 1.6$. Diagnostics based on mid-IR photometry indicate that these two samples are likely made by a very similar mixture of active star-forming galaxies and AGN.  The layout of the paper is as follows: In \\S\\,2 we describe the parent catalogue and the sample selection.  In \\S 3.1 we derive the two point angular correlation function, while in \\S 3.2 we present the results for the spatial clustering properties of the sources in exam. \\S 4 discusses the implications of the clustering results on the source properties, and in particular for what concerns their halo masses and number density. Our main conclusions are summarized in \\S\\,5.  Throughout this work we adopt a flat cosmology with $\\Omega_m=0.3$ and $\\Omega_\\Lambda=0.7$, a present-day value of the Hubble parameter in units of $100$ km/s/Mpc $h=0.7$, and rms density fluctuations within a sphere of $8 h^{-1}$ Mpc radius $\\sigma_8=0.8$ (Spergel et al. 2003).  \\section {Sample Selection}  For this work we used the {\\it Spitzer} wide-area infrared extragalactic (SWIRE) survey (Lonsdale et al. 2003; 2004) to select sources with fluxes at 24 $\\mu$m brighter than 400 $\\mu$Jy ($5\\sigma$ completeness) in the XMM-LSS field (Surace et al. 2005). In order to obtain multi-wavelength information and accurate redshift estimates for these sources we limited our analysis to the region of the SWIRE survey which overlaps with the 0.7 square degrees covered by the UKIDSS Ultra Deep Survey (UDS - Lawrence et al. 2006) and Subaru imaging (Sekiguchi et al. 2005; Furusawa et al. in preparation).  The 5 overlapping Subaru Suprime-Cam pointings provide broad band photometry in the $BVRi'z'$ filters to typical $5 \\sigma$ depths of $B=27.5$, $V=26.7$, $R=27.0$, $i'=26.8$ and $z'=25.9$ (within a $2$-arcsec diameter aperture). The UDS is the deepest of the five surveys that constitute the UKIDSS survey (Lawrence et al. 2006) and for this work we used $J$ and $K$-band imaging from the first data release with a $5 \\sigma$ depth of $J_{AB} = 23.5$  and $K_{AB} = 23.5$ (Warren et al. 2007). Finally, the SWIRE survey also provides Spitzer-IRAC measurements at 3.6, 4.5, 5.8 and 8$\\mu$m (Surace et al. 2005).  Within the overlapping region between SWIRE and UDS we isolated 1184 sources with $F_{24\\mu \\rm m}\\ge 400\\mu$Jy, out of which 1041 have a reliable optical and/or near-infrared counterpart. We take these objects as our working sample, UDS-SWIRE hereafter. As extensively described in Cirasuolo et al. (2007), photometric redshifts for these sources have been computed by fitting the observed photometry (9 broad bands from 0.4 to 4.5$\\mu m$) with both synthetic (Bruzual \\& Charlot 2003) and empirical (Coleman, Wu \\& Weedman 1980; Kinney et al. 1996; Mignoli et al. 2005) galaxy templates, by using the public package {\\sc HYPERZ} (Bolzonella, Miralles \\& Pell\\'{o} 2000). Comparisons with spectroscopic redshifts available in the field show the good accuracy of the photometric redshifts with a $\\Delta z /(1+z) \\simeq 0.05$ over the redshift range $0<z<6$ (Cirasuolo et al. 2007).\\\\ The redshift distribution $N(z)$ for our 24$\\mu$m sources is shown by the dotted line in  Figure~{\\ref{figure:nz}}. As it is possible to appreciate from the Figure, the objects included in the UDS, $F_{24\\mu \\rm m}\\ge 400\\mu$Jy sample are found up to redshifts $z\\sim 3.5$. \\\\  \\begin{figure} \\includegraphics[width=8.0cm]{nz.ps} \\caption{Photometric redshift distribution of sources in the UDS-SWIRE, $F_{24\\mu \\rm m}\\ge 400 \\mu$Jy sample (thin dotted line). The two thick solid lines highlight the redshift intervals covered by the high-z (blue) and low-z (red) samples. For comparison, the (green) dashed line indicates the redshift distribution of all $F_{24\\mu \\rm m}\\ge 400 \\mu$Jy sources fainter than $R=24.5$. No redshift cut was performed on this last sample. \\label{figure:nz}} \\end{figure}  \\noindent In order to investigate possible evolutionary features of the large-scale structure as traced by these objects, we have concentrated on two sub-samples: a first one which includes both sources with redshifts above 1.6 and those without an optical or K-band identification (hereafter called the {\\it high-z} sample) and a second one with objects within $z=0.6$ and $z=1.2$ (hereafter called the {\\it low-z} sample). There are a number of reasons for the above choice. $z\\sim 1.6$ is in fact the redshift at which the strong spectral PAH feature centred at 7.7$\\mu$m -- indicative of a very intense star-forming activity -- enters the 24$\\mu$m band. The high-z sample is therefore expected to include a relevant fraction of galaxies undergoing an intense phase of dust-enshrouded stellar formation. Furthermore, recent studies have proved that a strong 24$\\mu$m emission combined with a very faint or no optical detection most likely originates from obscured star-forming galaxies at redshifts between about 1.6 and 2.7 (see e.g. Yan et al. 2005; 2007; Houck et al. 2005). On the other hand, results both based on IRS spectroscopy and on diagnostics of the ratio between 8$\\mu$m and 24$\\mu$m fluxes, have inferred a fraction of $z\\sim 2$ obscured AGN brighter than our chosen 24$\\mu$m flux limit of about 30\\%, figure which rapidly increases to $\\sim 100$\\% at the highest 24$\\mu$m fluxes (see e.g. Brand 2006; Magliocchetti et al. 2007). The counts presented in the right-hand panel of Figure~2 seem to confirm this expectation ($\\sim 65$\\% candidate starburst galaxies on the basis of their $F_{8\\mu \\rm m}/F_{24\\mu \\rm m}< 0.1$ ratios, with the remaining $\\sim 35$\\% being made of candidate AGN-dominated objects).  \\begin{figure*} \\includegraphics[width=8.0cm]{N_S_lowz.ps} \\includegraphics[width=8.0cm]{N_S_highz.ps} \\label{figure:N_S} \\caption{Differential number counts for candidate star-forming galaxies ($F_{8\\mu\\rm m}/F_{24\\mu\\rm m}< 0.1$; green squares) and candidate AGN-dominated sources ($F_{8\\mu\\rm m}/F_{24\\mu\\rm m}\\ge 0.1$; red circles) in the low-z (left-hand panel) and high-z (right-hand panel) samples. } \\end{figure*}   \\begin{figure*} \\includegraphics[width=8.0cm]{radec_zgt0.6.and.z.lt.1.2.ps} \\includegraphics[width=8.0cm]{radec_zgt1.6.ps} \\label{figure:radec} \\caption{Sky distribution of UDS-SWIRE, $F_{24\\mu \\rm m}\\ge 400 \\mu$Jy sources (small dots). The filled (red) circles in the left-hand panel illustrate the 350 objects included in the redshift interval $z=[0.6-1.2]$, while those in the right-hand panel represent the 210 sources with either $z>1.6$ or no optical or near-IR identification.} \\end{figure*}    The low-z sample is also expected to probe a substantial fraction of obscured star-forming galaxies and AGN, albeit with lower intrinsic luminosities since we are dealing with a flux-limited survey. The $24\\mu\\rm m$ counts presented in the left-hand panel of Figure~2 are remarkably similar to those obtained for the high-z sample ($\\sim 60$\\% and $\\sim 40$\\% respectively for candidate starburst galaxies and candidate AGN-dominated objects; we note that the lower limit of $z\\sim 0.6$ adopted for the low-z sample corresponds to the lowest redshift at which a meaningful distinction between AGN-powered and SF-dominated sources as based on the ratio between $F_{8\\mu\\rm m}$ and $F_{24\\mu\\rm m}$ fluxes can be made given the spectral properties of intense star-forming galaxies in the mid-IR). Furthermore, from recent multi-photometry studies which extend from the optical to the X-ray band, it seems that $z\\sim 1$ corresponds to the 'bulk' of obscured AGN activity (see e.g. Pozzi et al. 2007). A comparative analysis of the two samples would then allow to investigate any redshift evolution in the large-scale properties of these sources and, if present, find an evolutionary connection between objects in the low-z and high-z samples.   The {\\it high-z} sample includes 210 sources, 28 of which do not have an optical counterpart. The redshift distribution of the objects with assigned photo-z closely resembles that of $R\\ge 24.5$, 24$\\mu$m-selected galaxies (cfr. Figure~\\ref{figure:nz}), showing once again that optically-faint sources selected at 24$\\mu$m most likely reside at redshifts $z > 1.6$. The projected distribution onto the sky of the high-z sources is shown by the filled (red) circles on the right-hand side of Figure~3. Small dots indicate the positions of all $F_{24\\mu \\rm m}\\ge 400 \\mu$Jy UDS-SWIRE objects.\\\\ By making use of the redshift distribution shown in Figure~\\ref{figure:nz}, and assuming that the $N(z)$ of optically obscured sources follows that of galaxies with estimated redshift, we find that the average redshift for the above sample is $<z>=2.02$, its median is $z_{\\rm med}=1.93$, and the mean comoving number density of such sources is $\\bar{n}_G=(2.5 \\pm 0.3) \\cdot 10^{-5}$ Mpc$^{-3}$, number which decreases to $\\bar{n}_G=(2.2 \\pm 0.3)\\cdot 10^{-5}$ Mpc$^{-3}$ if one only includes objects with redshifts. The errors associated to the above quantities represent the 2$\\sigma$ confidence level as derived from Poisson statistics. Estimates of the number density at $z\\sim 2$ might be very tricky due to joint effect of the cosmological evolution of the sources under examination and to the large variance of the observed 24$\\mu$m SED in that redshift range. This is why, for a safety check, we have also repeated our calculations of $\\bar{n}_G$ by simply considering a redshift box extending from $z=1.6$ to $z=2.5$. The result is virtually identical to that quoted above both by including and excluding galaxies with no redshift information and also very similar to that obtained by extrapolating the $z\\sim 2$ Caputi et al. (2007; cfr. their Figure 8) Luminosity Function for 24$\\mu$m-selected galaxies brighter than $\\nu L_{\\nu}\\simeq 10^{11.9} L_{\\odot}$, where the limiting luminosity has been calculated for  $F_{24\\mu\\rm m}=0.4$~mJy by following an Arp220-like Spectral Energy Distribution, indicative of galaxies undergoing intense star-formation which we expect (cfr. earlier in this section and also Figure~2) to dominate our sample.  The {\\it low-z} sample instead contains 350 sources. Their sky distribution is represented by the filled (red) circles on the left-hand side of Figure~3. The average redshift of the sample is found to be $<z>=0.79$, its median $z_{\\rm med}=0.84$, and the comoving number density of these sources is $\\bar{n}_G=(9.9 \\pm 1.0) \\cdot 10^{-5}$ Mpc$^{-3}$.  Even in this case, the number density is in excellent agreement with the findings of Caputi et al. (2007) for their $z\\sim 1$ sources with $\\nu L_{\\nu}\\simgt 10^{11.3} L_{\\odot}$, where $\\nu L_{\\nu}$ was calculated as indicated before. The most relevant properties for the high-z and low-z samples are summarized in Table~2. ",
        "conclusions": "We have investigated the clustering properties of 24$\\mu$m-selected galaxies brighter than 400$\\mu$Jy, drawn from the SWIRE and UKIDSS Ultra Deep Surveys. With the help of photometric redshift determinations, we have concentrated on two datasets which include galaxies respectively in the $z=[0.6-1.2]$ ({\\it low-z} sample) and $z\\simgt 1.6$ ({\\it high-z} sample) redshift ranges. The low-z sample is made of 350 galaxies, while the high-z sample includes both 182 galaxies with estimated photo-z's and another 28 sources which do not have any optical or near-IR counterpart. Diagnostics based on the ratio between 8$\\mu$m and 24$\\mu$m fluxes indicate that these two samples are likely made by a very similar mixture of active star-forming galaxies ($\\sim 65$ per cent) and AGN (the remaining $\\sim 35$ per cent).  Results obtained by fitting with a power-law of fixed slope $\\gamma=1.8$ the observed angular two-point correlation function $w(\\theta)$ report an amplitude $A^{hz}=0.010\\pm 0.0035$ and $A^{lz}=0.0055^{+0.0015}_{-0.0020}$ ($w(\\theta)=A\\;\\theta^{1-\\gamma}$) respectively for the high-z and low-z samples. In more physical coordinates, the above results imply (comoving) correlation lengths for the spatial two-point correlation function $\\xi$ (modelled as $\\xi=(r/r_0)^{-\\gamma}$) $r_0^{hz}=15.9^{+2.9}_{-3.4}$~Mpc and $r_0^{lz}=8.5^{+1.5}_{-1.8}$~Mpc, showing that the galaxies in the high-z sample are more strongly clustered that those found at lower redshifts.   A deeper insight on the above findings is provided by the so-called Halo Occupation Scenario which connects the clustering properties of a chosen population of astrophysical objects with some of their physical properties. The key quantity for this kind of analysis is the Halo Occupation Number (HON), i.e. the average number of galaxies enclosed in a halo of given mass $M$, which can be parametrized as $\\langle N\\rangle=N_0\\left(M/M_{\\rm min}\\right)^\\alpha$, with $M_{\\rm min}$ minimum mass for a halo to host one of such galaxies and $N_0$ normalization factor which provides information on how common the galaxies under examination are.  At first we have adopted as a working hypothesis that Spitzer-selected galaxies were spatially distributed within their dark matter halos according to the dark matter (NFW) distribution. However, we find that such a model fails at reproducing the high amplitude of the observed $w(\\theta$) signal on angular scales $\\simlt 0.003$ degrees both in the low-z and high-z sample. A much better agreement with the data is obtained when we take the galaxies to be very concentrated towards their halo centres, e.g. by assuming a distribution of the kind $\\rho\\propto r^{-3}$. Such a high concentration -- not observed in local data (Magliocchetti \\& Porciani 2003) -- is suggestive of a strong interaction and close encounters between galaxies which reside in the proximity of the halo centres, and we expect such a strong interaction to eventually drive a substantial number of gas-rich mergers which could easily trigger the observed enhanced star-formation activity.  Investigations of the best-fit values for the HON report values ${\\rm Log}\\left(M_{\\rm min}^{hz}/M_\\odot\\right)=12.8^{+0.2}_{-0.4}$; ${\\rm Log} N_0^{hz}=-0.7^{+0.5}_{-0.5}$; $\\alpha^{hz}=0.7^{+0.3}_{-0.5}$ for the high-z sample and ${\\rm Log}\\left(M_{\\rm min}^{lz}/M_\\odot\\right)=11.8 ^{+0.8}_{-0.8}$; ${\\rm Log} N_0^{lz}=-2.1^{+1.1}_{-0.9}$; $\\alpha^{lz}=0.7^{+0.1}_{-0.3}$ in the case of low-z sources. The above figures indicate that galaxies belonging to the high-z sample are exclusively associated with very massive, $M\\simgt 10^{12.8}M_\\odot$, structures, identifiable with groups-to-clusters of galaxies. Furthermore, these galaxies are quite common within their halos (in the most extreme case, our results imply that more than {\\it one in two} of all the structures with masses $M\\simeq M_{\\rm min}$ host a $z\\sim 2$, $F_{24 \\mu \\rm m}\\ge 400 \\mu$Jy galaxy), and their number sensibly increases with the richness of the structure which hosts them. On the other hand, the low-z sample seems to be made of galaxies of much smaller mass ($M_{\\rm min}\\simeq 10^{11.8} M_\\odot$). These objects are also very rare; in the most unfavourable case we get that only 0.1\\% of the 'allowed' halos host a galaxy of the kind which make the low-z sample.  Such a remarkable difference in the clustering and environmental properties of active sources as seen at $z\\sim 2$ and $z\\sim 1$  can be hardly attributed to the fact that high redshift galaxies are intrinsically brighter than their lower redshift counterparts (cfr. \\S4). Indeed, despite the large uncertainties determined by the poor statistics of the considered datasets, our results indicate that the populations probed by the high-z and low-z samples are very different from each other: massive active galaxies seem to disappear when going from $z\\sim 2$ to $z\\sim 1$, and at this lower redshifts all the AGN and starforming activity appears to be segregated in much lower-mass systems (see also the results of Gilli et al. 2007 who find for their bright, $z\\sim 0.7$, 24$\\mu$m-selected galaxies in the GOODS fields a correlation length $r_0\\sim 8$~Mpc.). We stress that, while investigations of the luminosity function of 8$\\mu$m-rest frame-selected sources have already established a strong luminosity evolution of sources between redshifts $\\sim 2$ and $\\sim 1$ (e.g. Caputi et al. 2007), clustering measurements provide a unique tool to determine the physical nature of such an evolution. Our work can in fact show that the strong differential evolution of the 24$\\mu$m LF is not simply due to the fact that the less luminous sources at $z\\sim 1$ are dimmed versions of the galaxies at higher-z (i.e. pseudo-passive evolution), but indeed has to be attributed to different populations of objects inhabiting different dark matter haloes and structures.  The general picture which then emerges from the results of this paper points to a differential evolution for high-mass and low-mass systems. Massive systems form early in time at the rare peaks of the density field. Their intense activity, both in terms of star-formation and AGN accretion, is strongly favoured (if not triggered) by numerous close encounters and gas-rich mergers happening in the proximity of the centres of the massive/cluster-like structures in which these sources reside. Such an active phase is relatively short-lasted and already below $z\\sim 1.5$ these objects evolve as optically passive galaxies, in most of the cases ending up as the supermassive galaxies which locally reside in cluster centres and which show a remarkable tendency for enhanced radio activity (see e.g. Best et al. 2006; Magliocchetti \\& Br\\\"uggen 2007). Lower ($M\\simlt 10^{12}M_\\odot$) mass systems are instead characterized by an active phase which lasts down to lower redshifts. Also in this case, the AGN and star forming activities seem to be favoured by the strong concentration of these sources towards their halo centres. These objects will eventually end up as 'typical' early-type galaxies which favour relatively (but not extremely) dense environments.  At present, studies on the evolutionary properties of dust-enshrouded active systems like those we have investigated in this work are hampered by the lack of surveys which can probe their peak of emissivity to low enough fluxes and also by the extreme difficulty of getting spectroscopic redshifts for such optically faint sources. The forthcoming advent of instruments such as Herschel, SCUBA-2 and ALMA will finally fill in this gap and provide the community with 'the ultimate truth' on most of the issues connected with galaxy formation and evolution."
    },
    "0710/0710.3593_arXiv.txt": {
        "abstract": "The Hamburg/ESO quasar \\helens\\ is found to be a quadruple gravitational lens, based on observations with the twin 6.5m Magellan telescopes at the Las Campanas Observatory, and subsequently with the \\hubble. The $z_S=1.235$ quasar appears in a cross configuration, with $i'$ band magnitudes ranging from 18.0 to 18.8. With a maximum image separation of $0\\farcs{}67$, this is the smallest-separation quadruple ever identified using a ground-based optical telescope. PSF subtraction reveals a faint lensing galaxy. A simple lens model succeeds in predicting the observed positions of the components, but fails to match their observed flux ratios by up to a magnitude. We estimate the redshift of the lensing galaxy to be $z_L \\sim 0.7$. Time delay estimates are on the order of a day, suggesting that the flux ratio anomalies are not due to variability of the quasar, but may result from substructure or microlensing in the lens galaxy. ",
        "introduction": "\\label{sec:intro}  A quasar lensed by a galaxy along the line of sight presents opportunities to study both the gravitational potential that is doing the lensing and the quasar that is lensed.  In particular the inevitable microlensing of the quasar by stars in the lensing galaxies permits estimates of the stellar (as opposed to dark) matter surface density of the lensing galaxy \\citep{2004IAUS..220..103S} and of the size of the quasar accretion disk \\citep{2007ApJ...661...19P,2007arXiv0707.0003P} and emission line region \\citep{2006ApJ...639....1K}.  Quadruply imaged systems are especially useful for such purposes, giving multiple lines of sight through the lens and additional constraints on the lensing potential.  But the rate at which new systems are discovered has leveled off to something like a half dozen per year, with the most recent discoveries coming largely from the Sloan Digital Sky Survey \\citep[e.g.][]{2007arXiv0708.0871O}.  Their selection algorithm is incomplete at separations below one arcsecond and worse for quadruple systems than for doubles \\citep{2006AJ....132..999O}.  Of the new SDSS lenses, just three are quadruple.  No comparable survey has been carried out in the southern hemisphere.  We report here the discovery of a new quadruply imaged system, \\helens. It was observed as part of a program to identify lensed systems among the Hamburg/ESO quasar survey \\citep{2000A&A...358...77W}.  While included within the SDSS survey area, it has not, until the present work, been identified as a lensed system.  In \\S\\ref{sec:obs}, we report the observations made using Magellan and the \\hubble. \\S\\ref{sec:analysis} describes our analysis of the data. In \\S\\ref{sec:model} we describe a simple model of the lens, and in \\S\\ref{sec:discuss} we make a rough estimate of the lens redshift. In \\S\\ref{sec:conclusions} we discuss the conclusions we can come to regarding \\helens. Throughout this paper, we adopt a ``concordance'' cosmology with $\\Omega_{M} = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and $h = 0.7$. ",
        "conclusions": ""
    },
    "0710/0710.1243_arXiv.txt": {
        "abstract": "We present the results of a detailed study of the high mass X-ray binary 2S 0114+650 made with the pointed instruments onboard the \\emph{Rossi X-ray Timing Explorer}. The spectral and temporal behaviour of this source was examined over the pulse, orbital, and super-orbital timescales, covering $\\sim$2 cycles of the 30.7 d super-orbital modulation. Marginal evidence for variability of the power law photon index over the pulse period was identified, similar to that observed from other X-ray pulsars. If this variability is real it can be attributed to a varying viewing geometry of the accretion region with the spin of the neutron star. Variability of the neutral hydrogen column density over the orbital period was observed, which we attribute to the line of sight motion of the neutron star through the dense circumstellar environment. A reduction in the power law photon index was observed during the orbital maximum, which we speculate is due to absorption effects as the neutron star passes behind a heavily absorbing region near the base of the supergiant companion's wind. No significant variability of the column density was observed over the super-orbital period, indicating that variable obscuration by a precessing warp in an accretion disc is not the mechanism behind the super-orbital modulation. In contrast, a significant increase in the power law photon index was observed during the super-orbital minimum. We conclude that the observed super-orbital modulation is tied to variability in the mass accretion rate due to some as yet unidentified mechanism. ",
        "introduction": "The high mass X-ray binary 2S 0114+650 has been shown to be a rather unusual system. It exhibits properties consistent with both Be and super-giant X-ray binaries \\citep*{cra85}, with significant broadband temporal and spectral variability over a wide range of timescales \\citep[e.g.][]{kon83,cra85}. The derived low X-ray luminosity \\citep[L$_x$ $\\sim$1.1 $\\times$ 10$^{36}$ erg s$^{-1}$ in the 3 -- 20 keV band for a distance of 7.2 kpc;][]{hal00} implies that spherical accretion takes place via the stellar wind of the donor star \\citep{li99}. However, the HeII 4686 {\\AA} line (a common signature of the presence of an accretion disc) has been observed weakly in emission on a few occasions (\\citealt*{aab83}; \\citealt{cra85,van89}), possibly linked to the presence of a small transient accretion disc.  Coherent X-ray (and possibly faint optical) pulsations at a period of $\\sim$2.7 h have been shown to be a persistent feature of the system \\citep*{fin92,tay95} and have been variously attributed to $\\beta$ Cephei-type pulsations/tidal oscillations of the donor star and the neutron star spin period. \\citet*{soo06} analysed the evolution of the pulse period in $\\sim$9.5 yr of \\emph{Rossi X-ray Timing Explorer} (\\emph{RXTE}) All Sky Monitor (ASM) data, identifying two episodes of torque reversal. Their conclusions support the scenario that 2S 0114+650 contains a super-slow rotator -- one of the slowest spinning X-ray pulsars yet discovered.  Variability at $\\sim$11.6 d has been observed in both optical \\citep{cra85} and X-ray wavelengths \\citep*{cor99}, consistent with the orbital period of the system. \\citet{cra85} were the first to confirm the binary nature of 2S 0114+650 using radial velocity measurements (with a period of 11.588 $\\pm$ 0.003 d for an eccentricity of 0.16 $\\pm$ 0.07, and 11.591 $\\pm$ 0.003 d for an eccentricity of 0.0). Although they could not distinguish between a circular and an eccentric orbit, they concluded that the orbital motion was most likely eccentric as the system had presumably gone through a supernova stage which should lead to a perturbed orbit. \\citet{cor99} analysed $\\sim$2.5 yr of \\emph{RXTE} ASM observations of 2S 0114+650 between 1996 -- 1999, identifying X-ray modulations close to both the optically derived orbital period and the proposed pulse period. Fitting a sine wave to the ASM light curve yielded a period of 11.63 $\\pm$ 0.007 d. The inconsistency between the optical and X-ray orbital period values was explained as either due to variable contamination of absorption lines by low equivalent width emission lines (resulting from the X-ray irradiation of the donor star) or possibly variability in the X-ray orbital modulation.  The determination of the orbital period was revisited recently in both optical and X-ray wavelengths by \\citet{gru07} and \\citet{wen06} respectively. \\citet{gru07} performed a long-term spectroscopic monitoring program of the H$\\alpha$ emission in 2S 0114+650 and presented revised orbital elements with a period of 11.5983 $\\pm$ 0.0006 d and a non-zero eccentricity of 0.18 $\\pm$ 0.05. \\citet{wen06} re-analysed the ASM data and reported a new value for the orbital period of 11.599 $\\pm$ 0.005 d, consistent with the new optical period.  A super-orbital modulation at 30.7 $\\pm$ 0.1 d was reported by \\citet*{far06} from the analysis of 8.5 yr of ASM data. \\citet{wen06} subsequently reported a new improved value of 30.75 $\\pm$ 0.03 d from re-analysis of the same data. \\citet{soo06} analysed the evolution of the super-orbital period, finding it to be stable and concluding that 2S 0114+650 is only the fourth reported system to show stable super-orbital variability.  The standard model invoked to explain super-orbital variability in X-ray binaries is periodic obscuration or reflection via the precession of a warped accretion disc \\citep[e.g.][]{ogl01}. The apparent lack of a persistent disc in this system thus raises doubts about the validity of this model as an explanation for the observed 30.7 d modulation. Recently, \\citet*{kon06} presented a model by which both the pulse and super-orbital modulations in this system could be explained by tidally induced oscillations in the donor star. However, the precise mechanisms by which these oscillations may translate into a structured stellar wind are not clear. In addition, super-orbital variability is predicted for circular orbits while the eccentric orbit models predict strong variations on orbital timescales associated with periastron passage \\citep{kon06}. These results are thus inconsistent with the non-zero eccentricity recently derived by \\citet{gru07}.  In this paper we present the results of a series of X-ray observations of 2S 0114+650 performed with the Proportional Counter Array \\citep*[PCA; see][]{jah96} and High Energy X-ray Timing Experiment \\citep*[HEXTE; see][]{gru96} instruments onboard the \\emph{RXTE} satellite. The preliminary results of these studies were published in \\citet{far07}. The temporal and spectral variability of this source was studied in detail over the pulse, orbital and super-orbital timescales. In $\\S$ 2 the data reduction steps that were taken are explained. A description of the analysis methods employed and the results obtained are provided in $\\S$ 3. A discussion of our results is provided in $\\S$ 4, and $\\S$ 5 presents the conclusions that we have drawn. ",
        "conclusions": "We have performed an extensive study of the X-ray variability of 2S 0114+650 with the $\\emph{RXTE}$ satellite. The good spectral coverage provided by the PCA and HEXTE instruments have allowed us to undertake the first in-depth wide-band study of the super-orbital modulation. The overall X-ray spectrum in the 3 -- 50 keV range was found to be well described by the standard absorbed power law model with a high energy exponential cut-off, with spectral parameters consistent with those reported by other authors from other instruments and epochs.  Marginal evidence for variability in the power law photon index over the pulse period was identified, which could be attributed to the varying viewing geometry of the accretion region with the spin of the neutron star. Similar variability has been observed from other X-ray pulsars such as Her X-1, supporting the conclusion that 2S 0114+650 contains a super-slow X-ray pulsar. The near-sinusoidal pulse profile and the lack of change with energy puts 2S 0114+650 in the class of low luminosity ($\\la$ 10$^{36}$ erg s$^{-1}$) X-ray pulsars, allowing us to provide a new distance estimate of 4.5 $\\pm$ 1.5 kpc.  Variability of the neutral hydrogen column density over the orbital period was observed, indicating that the line-of-sight motion of the neutron star through the dense circumstellar environment is the mechanism behind the orbital modulation. Additional variability in the power law photon index anti-correlated with flux over the orbital period was identified, which we speculate is due to absorption effects as the neutron star passes behind a heavily absorbing region near the base of the supergiant companion's wind.  In contrast, no significant variability of the column density was observed over the super-orbital period. While super-orbital periods are commonly attributed to variable obscuration by a precessing warp in an accretion disc, the absence of super-orbital column density variability indicates that this is not the mechanism behind the super-orbital modulation. Similarly, the absence of significant iron emission or reflection components in the spectrum argues against variable reflection by a precessing warped disc. While the slope of the spectrum was relatively stable throughout the super-orbital cycle, a significant increase in the photon index was observed during the minimum phase. We conclude that the super-orbital period is tied to variability in the mass accretion rate due to some as yet unidentified mechanism."
    },
    "0710/0710.1834.txt": {
        "abstract": "{ We report on the first results of a search for molecular hydrogen emission from protoplanetary disks using CRIRES, ESO's new VLT Adaptive Optics high resolution near-infrared spectrograph. We observed the classical T Tauri star  LkH$\\alpha$ 264 and the debris disk 49 Cet, and searched for $\\upsilon= 1-0$ S(1) H$_2$ emission at 2.1218 $\\mu$m, $\\upsilon = 1-0$ S(0) H$_2$ emission at 2.2233 $\\mu$m and $\\upsilon = 2-1$ S(1) H$_2$ emission at 2.2477 $\\mu$m. The H$_2$ line at 2.1218 $\\mu$m is detected in LkH$\\alpha$ 264 confirming the previous observations by Itoh et al. (2003). In addition, our CRIRES spectra reveal the previously observed but not detected H$_2$ line at 2.2233 $\\mu$m in LkH$\\alpha$ 264. An upper limit of 5.3 $\\times 10^{-16}$ ergs s$^{-1}$ cm$^{-2}$ on the $\\upsilon = 2-1$ S(1) H$_2$ line flux in LkH$\\alpha$ 264 is derived. The detected lines coincide with the rest velocity of LkH$\\alpha$ 264. They have a FWHM of $\\sim$20 km s$^{-1}$. This is strongly suggestive of a disk origin for the lines. These observations are the first simultaneous detection of $\\upsilon = 1-0$ S(1) and $\\upsilon = 1-0$ S(0) H$_2$ emission from a protoplanetary disk. 49 Cet does not exhibit H$_2$ emission in any of the three observed lines. We derive the mass of optically thin H$_2$ at $T\\sim1500$ K  in the inner disk of LkH$\\alpha$ 264 and derive stringent limits in the case of 49 Cet at the same temperature. There are a few lunar masses of optically thin hot H$_2$ in the inner disk ($\\sim$ 0.1 AU) of LkH$\\alpha$ 264, and less than a tenth of a lunar mass of hot H$_2$ in the inner disk of 49 Cet. The measured 1-0 S(0)/1-0 S(1) and 2-1 S(1)/1-0 S(1) line ratios in LkH$\\alpha$ 264 indicate that the H$_2$ emitting gas is at a temperature lower than 1500 K and that the H$_2$ is most likely thermally excited by UV photons. The $\\upsilon = 1-0$ S(1) H$_2$ line in LkH$\\alpha$ 264 is single peaked and spatially unresolved. Modeling of the shape of the line suggests that the disk should be seen close to face-on ($i<35^{\\rm\\,o}$) and that the line is emitted within a few AU of the LkH$\\alpha$ 264 disk. A comparative analysis of the physical properties of classical T Tauri stars in which the H$_2$ $\\upsilon = 1-0$ S(1) line has been detected and non-detected indicates that the presence of H$_2$ emission is correlated with the magnitude of the UV excess and the strength of the H$\\alpha$ line. The lack of H$_2$ emission in the NIR spectra of 49 Cet and the absence of H$\\alpha$ emission suggest that the gas in the inner disk of 49 Cet has dissipated. These results combined with previous detections of $^{12}$CO emission at sub-mm wavelengths indicate that the disk surrounding  49 Cet should have an inner hole. We favor inner disk dissipation by inside-out photoevaporation, or the presence of an unseen low-mass companion as the most likely explanations for the lack of gas in the inner disk of 49 Cet. ",
        "introduction": "The discovery of extrasolar planets triggered  an increasing interest in the physical mechanisms behind the process of planet formation. Many recent efforts have been directed to the study of disks surrounding pre-main-sequence stars. Observational and theoretical evidence suggests that planets are forming in these disks. The observational characterization of the physical structure and dynamics of the gas and dust in protoplanetary disks is of paramount importance for understanding the process of planet formation.  In the inner 1 AU of protoplanetary disks, intense UV or X-ray heating can bring the gas temperatures to a few thousand Kelvin. At these high temperatures, ro-vibrational transitions of H$_2$ are excited and a rich spectrum of H$_2$ lines in the near-infrared is expected to be produced. The study of H$_2$ quiescent ro-vibrational emission\\footnote{By quiescent emission we mean emission at the rest velocity of the star. H$_2$ emission can also be produced by shocked gas associated with outflows. However, in such a case the emission is expected to be doppler shifted more than 20 km s$^{-1}$ with respect to the rest velocity of the star.} towards pre-main-sequence stars with disks offers the opportunity to address the question of the presence of hot gas in the disk, by probing the temperature and density in the innermost regions where terrestrial planets are expected to form. For example, the H$_2$ $\\upsilon =1-0$ S(1) line at 2.1218 $\\mu$m (one of the strongest H$_2$ ro-vibrational lines) is sensitive to a few lunar masses of gas. Therefore, the absence of the line would be strongly suggestive of little or no hot gas in the systems.   In this paper we present the first results of a sensitive search for near-infrared H$_2$ emission from protoplanetary disks using CRIRES, ESO's new VLT near-infrared high-resolution spectrograph. We searched for the H$_2$ $\\upsilon =1-0$ S(1) line at  2.1218 $\\mu$m, H$_2$ $\\upsilon = 1-0$ S(0) line at 2.2233 $\\mu$m and H$_2$ $\\upsilon = 2-1$ S(1) line at 2.2477 $\\mu$m, towards LkH$\\alpha$ 264, a classical T Tauri star with previously reported detections of the $\\upsilon =1-0$ S(1) line by Itoh et al. (2003), and 49 Cet, a debris disk with evidence of a large reservoir of cold gas at sub-mm wavelengths (Dent et al. 2005, Zuckerman et al. 1995). For a summary of the stellar physical properties of LkH$\\alpha$ 264 and 49 Cet see Table 1. We confirm the detection of H$_2$ emission at 2.1218 $\\mu$m in LkH$\\alpha$ 264, and announce, for the first time, the detection of H$_2$ emission at 2.2233 $\\mu$m  from a disk (LkH$\\alpha$ 264). We report the non-detections of H$_2$ emission at 2.2477 $\\mu$m in LkH$\\alpha$ 264, and at 2.1218, 2.2233 and 2.2477 $\\mu$m in 49 Cet.  The paper is organized as follows. We start with a description of the observations and the data reduction. In section 3 we present our results and calculate the mass limits for the hot ($T\\sim1500$ K) H$_2$ in the systems. In Section 4, based upon the measured 1-0 S(0)/1-0 S(1) and 2-1 S(1)/1-0 S(1) line ratios in LkH$\\alpha$ 264, we determine the excitation mechanism of the observed H$_2$ emission. By modeling of the shape of the $\\upsilon = 1-0$ S(0) H$_2$ line, we derive constraints on the H$_2$ emitting region and the inclination of the disk around LkH$\\alpha$ 264. Finally we discuss the disk properties of the stars in which H$_2$ emission has been detected and the prospects for future investigations. Our conclusions are presented in Section 5.  %%***************************** table observations *********************** \\begin{table*} \\caption{Summary of the observations.}             % title of Table \\label{table:observations}      % is used to refer this table in the text \\centering                          % used for centering table \\begin{tabular}{@{}l c l c c c c c l c c c c c c c @{}}        % centered columns (4 columns) \\hline\\hline                 % inserts double horizontal lines Star            & $\\lambda$      & Date      & UT      & %V$_{\\oplus}\\,$ & $t_{exp}$       & Airmass        & seeing    & Calibrator $^a$  & $t_{exp}$      & Airmass   & seeing  \\\\ % & [$\\mu$m]       &           & [hh:mm]  & %[km s$^{-1}$]  & [s]             &                & [arcsec]  & & [s]            &           & [arcsec] \\\\ \\hline                        % inserts single horizontal lines LkH$\\alpha$ 264 & 2.1218      &  8 Nov 2006  & 06:25   &%TBD            & 720             & 1.4            & 1.2       & HIP 13327    & 160            & 1.3       & 0.9 \\\\ & 2.2233, 2.2477       &  8 Nov 2006  & 06:52   &%              & 720             & 1.4            & 1.0       & HIP 13327    & 160            & 1.3       & 0.9 \\\\ 49 Cet          & 2.1218         &  9 Nov 2006  & 02:38   &%TBD            & 240             & 1.1            & 0.8       & HIP 8497    & 40             & 1.0       & 1.2          \\\\ & 2.2233, 2.2477       &  9 Nov 2006  & 03:07   &%              & 240             & 1.2            & 0.7       & HIP 8497     & 40             & 1.0       & 1.2          \\\\ \\hline \\addlinespace[5pt] \\multicolumn{10}{l}{$^a$ Spectrophotometric standard stars were observed immediately following the science observations.} \\end{tabular} \\end{table*} %%********************************************************** %***************************** table fluxes *********************** \\begin{table*} \\caption{H$_2$ line fluxes and upper limits measured }             % title of Table \\label{table:results}      % is used to refer this table in the text \\centering                          % used for centering table \\begin{tabular}{@{} l l c c c c c c c c @{} }        % centered columns (4 columns) \\hline\\hline \\addlinespace[5pt] & & \\multicolumn{3}{c}{LkH$\\alpha$ 264} & \\multicolumn{3}{c}{49 Cet }\\\\ \\cmidrule(r){3-5}\\cmidrule(r){6-8} & &1-0 S(1) & 1-0 S(0) & 2-1 S(1) & 1-0 S(1) & 1-0 S(0) & 2-1 S(1) \\\\ & & 2.1218 $\\mu$m & 2.2233 $\\mu$m      & 2.2477 $\\mu$m & 2.1218 $\\mu$m & 2.2233 $\\mu$m      & 2.2477 $\\mu$m \\\\ \\hline \\addlinespace[5pt] Continuum  & [ergs s$^{-1}$ cm$^{-2}~\\mu$m$^{-1}$] & $9.2\\times 10^{-11}$   &$\\,\\,1.5\\times 10^{-10}$ &$\\,\\,\\,1.6\\times 10^{-10}$ & $\\,3.5\\times 10^{-9}$  &$\\,1.7\\times 10^{-9}$    &$\\,\\,1.5\\times 10^{-9}$\\\\ 3$\\sigma$  & [ergs s$^{-1}$ cm$^{-2}~\\mu$m$^{-1}$] & $4.2\\times 10^{-12}$   &$\\,\\,8.2\\times 10^{-12}$ &$\\,\\,\\,1.0\\times 10^{-11}$ & $\\,1.0\\times 10^{-10}$ &$\\,\\,1.8\\times 10^{-10}$ &$\\,\\,\\,3.1\\times 10^{-10}$\\\\ Integrated Line Flux$^{a}$ & [ergs s$^{-1}$ cm$^{-2}$] & $3.0\\times 10^{-15}$   & $1.0\\times 10^{-15}$    & $<5.3\\times 10^{-16}$ &$<5.4\\times 10^{-15}$   & $<8.9\\times 10^{-15}$   & $<1.6\\times 10^{-14}$ \\\\ \\hline                                   %inserts single line \\addlinespace[5pt] \\multicolumn{8}{l}{$^a$ For the calculation of upper limits, we assumed that the FWHM of the line is 6.6 km s$^{-1}$. }\\\\ \\end{tabular} \\end{table*} %********************************************************** ",
        "conclusions": "We observed the classical T Tauri star LkH$\\alpha$ 264 and the debris disk 49 Cet and searched for ro-vibrational $\\upsilon =1-0$ S(1) H$_2$ emission at 2.1218 $\\mu$m, $\\upsilon =1-0$ S(0) H$_2$ emission at 2.2233 $\\mu$m, and $\\upsilon = 2-1$ S(1) H$_2$ emission at 2.2477 $\\mu$m, using CRIRES ($R \\sim 6.6$~km s$^{-1}$) at ESO-VLT. We confirmed the detection of  the $\\upsilon =1-0$ S(1) H$_2$ line in LkH$\\alpha$ 264 at the rest velocity of the star. The line has a flux of 3.0 $\\times 10^{-15}$ ergs cm$^{-2}$ s$^{-1}$, and a FWHM of 20.6 km s$^{-1}$. In addition, the enhanced sensitivity of CRIRES allowed the observation of the previously undetected $\\upsilon =1-0$ S(0) H$_2$ line in LkH$\\alpha$ 264. The line has a flux of 1.0 $\\times 10^{-15}$ ergs cm$^{-2}$ s$^{-1}$, and a FWHM 19.8 km s$^{-1}$. An upper limit of 5.3 $\\times 10^{-16}$ ergs s$^{-1}$ cm$^{-2}$ was derived for the $\\upsilon = 2-1$ S(1) H$_2$ line flux in LkH$\\alpha$ 264. The very similar FWHM of the two H$_2$ lines detected suggests that the emitting gas is located in similar regions in the disk. Both lines are spatially unresolved. The measured mean PSF's FWHM ($\\approx$ 0.36\") in the H$_2$ 1-0 S(1) spectrum indicates that the H$_2$ emitting region is located in the inner 50 AU of the disk assuming a distance of 300 pc for LkH$\\alpha$ 264. The measured 1-0 S(0)/1-0 S(1) (0.33 $\\pm$ 0.1) and the 2-1 S(1)/1-0 S(1) ($<$0.2) line ratios in LkH$\\alpha$ 264 indicate that the H$_2$ emitting gas is at a temperature lower than 1500 K and that the H$_2$ is most likely thermally excited by UV photons. The measured line ratios suggest that X-ray excitation plays only a minor role in the heating of the emitting H$_2$ in LkH$\\alpha$ 264. The flux of the $\\upsilon =1-0$ S(1) H$_2$ line in LkH$\\alpha$ 264 implies that there are a few lunar masses of hot H$_2$ gas in the inner disk of LkH$\\alpha$ 264. The $\\upsilon =1-0$ S(1) H$_2$ line in LkH$\\alpha$ 264 is single peaked. Modeling of the $\\upsilon =1-0$ S(1) line shape indicates that the disk is close to face-on ($i<35^{\\rm\\,o}$). The best model fit suggests that the disk of LkH$\\alpha$ 264 is inclined 20$^{\\rm\\,o}$ for a H$_2$ emitting region extending from 0.1 to 10 AU with a power law relation of the intensity as a function of radius with exponent $\\alpha=-2$. If the $\\upsilon = 1-0$ H$_2$ S(1) line intensity decreases with an exponent $\\alpha=-2$ as a function of radius, then 50\\% of the line flux is produced within 0.1 AU and 1 AU of the LkH$\\alpha$ 264 disk, 40\\% of the line flux is emitted within 1 and 7 AU and the rest of the flux at larger radii.  A comparative analysis of the physical properties of classical T Tauri stars in which the H$_2$ $\\upsilon =1-0$ S(1) line has been detected versus non-detected shows that there is a higher chance of observing the H$_2$ near-infrared lines in CTTS  with a high $U-V$ excess and a strong H$\\alpha$ line. This result suggests that there is a higher probability of detecting the H$_2$ $\\upsilon =1-0$ S(1) line in systems with high accretion. In contrast to weak-lined T Tauri stars, there is no apparent correlation between the X-ray luminosity and the detectability of the H$_2$ $\\upsilon =1-0$ S(1) line in classical T Tauri stars. Taken as a group, LkH$\\alpha$ 264 and the CTTS in which the H$_2$ emission has been detected exhibit typical properties of classical T Tauri stars. Therefore, we expect NIR ro-vibrational H$_2$ lines from T Tauri disks to be detected on a routine basis in the near future.  The non-detection of any of the three H$_2$ lines in 49 Cet puts stringent constraints on the amount of hot gas in the inner disk. From the upper limit for the flux of the $\\upsilon =1-0$ S(1) H$_2$ line we deduced that less than a tenth of lunar-mass of gas at $T\\sim 1500$ K is present in the inner 1 AU of the disk surrounding 49 Cet. The lack of H$_2$ ro-vibrational  emission in the  spectra of 49 Cet, combined with non detection of pure rotational lines of H$_2$ (Chen et al. 2006) and the absence of H$\\alpha$ emission suggest that the gas in the inner disk of 49 Cet has dissipated. These results together with the previous detection of $^{12}$CO emission at sub-mm wavelengths (Zuckerman et al. 1995; Dent et al. 2005) point out that the disk of 49 Cet should have a large inner hole, and it is strongly suggestive of theoretical scenarios in which the disk disappears inside-out. We favor inner disk dissipation by inside-out photoevaporation, or the presence of an unseen low-mass companion(s) as most likely explanations for the lack of warm gas in the inner disk of 49 Cet."
    },
    "0710/0710.3064_arXiv.txt": {
        "abstract": "{We present results on searches for exotic particles (relativistic magnetic monopoles and WIMPs) and for UHE neutrinos, obtained with the Baikal neutrino telescope NT200. }  \\begin{document} ",
        "introduction": "The Baikal Neutrino Telescope  is operated in Lake Baikal, Siberia,  at a depth of \\mbox{1.1 km}. The first stage telescope configuration  NT200 was put into permanent operation on April 6th, 1998. The upgraded telescope NT200+ was put into operation on April 9th, 2005. This configuration consists of the old NT200 telescope, surrounded by three new external strings. Status and description of the detector has been presented elsewhere ~\\cite{APP2}. In this paper we present selected physics results of a search for fast magnetic monopoles, neutrino signals from WIMP annihilation at the center of the Earth and a search for diffuse neutrinos with energies larger than 10 TeV, obtained with the Baikal neutrino telescope NT200. ",
        "conclusions": "The Baikal neutrino telescope NT200 is taking data since April 1998. The upper limit obtained for a diffuse ($\\nu_e+\\nu_{\\mu}+\\nu_{\\tau}$) flux with $E^{-2}$ shape is $E^2 \\Phi = 8.1 \\times 10^{-7}$cm$^{-2}$s$^{-1}$sr$^{-1}$GeV. The limits on fast magnetic monopoles and on a diffuse $\\bar{\\nu_e}$ flux at the resonant  energy 6.3$\\times$10$^6$GeV are presently the most stringent ones.   {\\it This work was supported by the Russian Ministry of Education and Science, the German Ministry of Education and Research and the Russian Fund of Basic Research (grants 05-02-17476, 05-02-16593, 07-02-10013, 07-02-00791), by the Grant of the President of Russia NSh-4580.2006.2 and by NATO-Grant NIG-9811707 (2005).}"
    },
    "0710/0710.3587_arXiv.txt": {
        "abstract": "The current \\swf\\, sample of gamma-ray bursts (GRBs) with measured redshifts allows to test the assumption that GRBs trace the star formation in the Universe. Some authors have claimed that the rate of GRBs increases with cosmic redshift faster than the star formation rate, whose cause is not known yet. In this paper, I investigate the possibility for interpreting the observed discrepancy between the GRB rate history and the star formation rate history by the cosmic metallicity evolution, motivated by the observation that the cosmic metallicity evolves with redshift and GRBs prefer to occur in low metallicity galaxies. First, I derive a star formation history up to redshift $z=7.4$ from an updated sample of star formation rate densities obtained by adding the new UV measurements of Bouwens et al. and the new UV and infrared measurements of Reddy et al. to the existing sample compiled by Hopkins \\& Beacom. Then, adopting a simple model for the relation between the GRB production and the cosmic metallicity history as proposed by Langer \\& Norman, I show that the observed redshift distribution of the \\swf\\ GRBs can be reproduced with a fairly good accuracy. Although the results are limited by the small size of the GRB sample and the poorly understood selection biases in detection and localization of GRBs and in redshift determination, they suggest that GRBs trace both the star formation and the metallicity evolution. If the star formation history can be accurately measured with other approaches, which is presumably achievable in the near future, it will be possible to determine the cosmic metallicity evolution with the study on the redshift distribution of GRBs. ",
        "introduction": "\\label{intro}  Since the discovery of the afterglows of gamma-ray bursts (GRBs) and the determination of their redshifts \\citep{met97,vanp97}, it has been firmly established that GRBs are at cosmological distances \\citep[for recent reviews see][]{pir04,zha04,mes06}. Observations on the hosts of GRBs have revealed that long-duration GRBs (hereafter GRBs) are associated with faint, blue and often irregular galaxies with high star formation rates (SFRs) \\citep[and references therein]{con05,fru06,tan07,wai07}, confirming the early speculation that GRBs occur in star-formation regions and arise from the death of massive stars (Paczy\\'nski 1998; Wijers et al. 1998; see, however, Le Floc'h et al. 2006). The discovery of the GRB-supernova connection \\citep[and references therein]{gal98,li06,woo06b} supports the collapsar model for long-duration GRBs (MacFadyen \\& Woosley 1999; MacFadyen, Woosley \\& Heger 2001).  Because of their very high luminosity, GRBs can be detected out to the edge of the visible Universe with minimal extinction by intervening gases and dust \\citep{cia00,lam00,bro02,nao07} and are hence an ideal tool for probing the formation rate and the environments of stars at high redshift, the reionization history, as well as the cosmic chemical evolution \\citep{fyn06b,pri06,sav06,tot06,bro07,cam07,gal07,pro07}. The advantage of GRBs over quasars for probing the high redshift Universe has been discussed by \\citet{bro07}. It has also been proposed that GRBs can be used as standard candles to constrain cosmological parameters \\citep[and references therein]{blo03,fri05,sch07}. However, this proposal has been seriously challenged by a recent study of \\citet{li07}.  \\begin{figure} \\includegraphics[angle=0,scale=0.51]{grb_to_sfr.eps} \\caption{The observed ratio of the GRB rate to the SFR, $R_{\\grb/\\sfr}$, as a function of $Q(z)$ and $z$, where $Q(z)$ is defined as an increasing function of redshift $z$ (equation \\ref{Q} and Fig.~\\ref{Q_z}). Normalization is chosen so that $R_{\\grb/\\sfr} = 1$ at $Q=0.5$. The GRB rate history was obtained from a sample of {\\em bright} \\swf\\ GRBs so that selection effects are minimized for $z\\la 4$ \\citep[see Section \\ref{bright} of this paper]{kis07}. The SFR is the best fit to an updated sample of measured cosmic SFRs up to $z\\sim 7.4$ (eqs.~\\ref{sfr} and \\ref{ab}, Section \\ref{sfh}). The horizontal dotted line denotes $R_{\\grb/\\sfr}=1$, an expected result of the assumption that GRBs trace the star formation unbiasedly. } \\label{rgrb_sfr} \\end{figure}  To study the relation between GRBs and the star formation, people often assume that the GRB rate is proportional to the SFR then compare the predicted distribution of the GRB redshift (or other parameters, e.g. the intensity) to the observed distribution \\citep[for a review see Coward 2007]{tot97,mao98,wij98,por01,nat05,jak06,dai07,le07}. If GRBs trace the star formation in the Universe unbiasedly, one would expect that the ratio of the GRB rate to the SFR ($R_{\\grb/\\sfr}$) does not vary with redshift. Then, an accurate measurement of the GRB rate history at high redshift would directly lead to the star formation history (SFH) in the early epoch which is otherwise hard to measure with the current technology because of the uncertainty in dust obscuration for UV photons \\citep{hop06}.  Unfortunately, recent studies show that GRBs do not seem to trace the star formation unbiasedly (Fig.~\\ref{rgrb_sfr}). Based on the current understanding on the SFH and the \\swf\\ sample of GRBs with measured redshifts, people have found that $R_{\\grb/\\sfr}$ increases with redshift significantly \\citep{dai07,le07,kis07,yuk07,cen08}. While observations consistently show that the comoving rate density of star formation is nearly constant in the interval $1\\la z\\la 4$ \\citep{hop06}, the comoving rate density of GRBs appears evolving distinctly.  Adopting a model-independent approach by selecting bright \\swf\\, GRBs with the isotropic-equivalent luminosity $L_\\iso > 10^{51}$erg s$^{-1}$, \\citet{kis07} found that there are $\\sim 4$ times as many GRBs at redshift $z\\approx 4$ than expected from star formation measurements. They claimed that some unknown mechanism is leading to an enhancement in the observed rate of high-redshift GRBs. With a more sophisticated method, \\citet{dai07} found that to reconcile the observed GRB redshift distribution with the measured SFH, the efficiency of GRB production by massive stars would be nearly six to seven times higher at $z\\sim 7$ than at $z\\sim 2$. Based on their results, Daigne et al. concluded that GRB properties or progenitors must evolve with cosmic redshift.  In this paper, I investigate the relation between the SFH and the GRB rate history, and explore the possibility that the observed enhancement in the GRB rate at high redshift is caused by the cosmic metallicity evolution. Although the possibility of leading to an enhancement in the observed GRB rate evolution by the cosmic metallicity was mentioned by \\citet{kis07}, they did not give a quantitative analysis or a detailed discussion. Instead, Kistler et al. discussed more thoroughly on other possible causes, including evolution in the fraction of binary systems which had been proposed as a channel for producing GRBs, an evolving initial mass function (IMF) of stars, and evolution in the galaxy luminosity function (LF).  There is growing evidence that metallicities play an important role in the production of GRBs. Observations on the hosts of GRBs revealed that GRBs prefer to occur in galaxies with low metallicities (Fynbo et al. 2003, 2006a; Prochaska et al. 2004; Soderberg et al. 2004; Gorosabel et al. 2005; Berger et al. 2006; Savaglio 2006; Stanek et al. 2006; Wolf \\& Podsiadlowski 2007; Modjaz et al. 2008; Savaglio, Glazebrook \\& Le Borgne 2008). Based on the observational and theoretical evidence that the mass-loss rate of Wolf-Rayet stars depends on the metallicity \\citep{cro02,vin05,cro07}, theoretical studies on the collapsar model of GRBs arising from single massive stars suggested that GRBs can only be produced by stars with metallicity $Z\\la 0.1 Z_\\odot$ since otherwise strong stellar winds will cause stars to lose too much mass and angular momentum to form a disk around a black hole of several solar masses which is essential for the production of GRBs (Hirschi, Meynet \\& Maeder 2005; Yoon \\& Langer 2005; Woosley \\& Heger 2006; Yoon, Langer \\& Norman 2006).  It is well-known that the cosmic metallicity evolves strongly with redshift, and galaxies at higher redshift tend to have lower metallicities (Pettini et al. 1999; Prochaska et al. 2003; Rao et al. 2003; Kobulnicky \\& Kewley 2004; Kewley \\& Kobulnicky 2005, 2007; Kulkarni et al. 2005, 2007; Savaglio et al. 2005, 2008; Wolfe, Gawiser \\& Prochaska 2005; Savaglio 2006; P\\'eroux et al. 2007). \\citet{nat05} considered a model where GRBs trace the average metallicity in the Universe rather than the SFR, and the GRB rate decreases with increasing metallicity. However, their model led to a GRB redshift distribution that is nearly indistinguishable from the distribution predicted by the SFR \\citep{jak06}.  Recently, \\citet{lan06} considered the effect of the cosmic metallicity evolution on the integrated production rate of GRBs in the framework of the collapsar model, assuming that the GRB rate is jointly determined by the SFR and the metallicity evolution. Adopting a simple model for the metallicity evolution and a best fitted SFR, and assuming that a GRB is produced if a progenitor star is massive enough and has a metallicity below a threshold ($Z\\la 0.1 Z_\\odot$), they showed that the observed global ratio of the GRB rate to the core-collapse supernova rate ($\\sim 0.001$) can be reproduced. \\citet{nuz07} and \\citet{cen08} investigated the host galaxies of GRBs in a cosmological hierarchical scenario with numerical simulations. They found that the observed properties of GRB hosts are reproduced if GRBs are required to be generated by low metallicity stars.  I will incorporate the model of \\citet{lan06} into a probability distribution function of the luminosity and redshift of GRBs to study the rate history of GRBs. For this purpose, I will present an updated SFH obtained by adding the new measurements of SFR densities at $z\\sim 2.3$ and $3.05$ by \\citet{red08} and at $z\\sim 3.8$, $5.0$, $5.9$ and $7.4$ by \\citet{bou07,bou08} to the data compiled by \\citet{hop06}. With the updated data I will derive an analytic formula for the SFH and show that the observed distribution of \\swf\\ GRBs can be successfully reproduced when the evolution of the cosmic metallicity is properly taken into account. Then I will argue that, after a significantly expanded and well-defined sample of GRBs with measured redshifts and luminosities is available in future and the SFH at high redshift is accurately determined with other approaches, GRBs will be a powerful tool for probing the cosmic metallicity evolution.  The paper is organized as follows. In Section \\ref{sfh}, I present the updated measurements on the SFR up to $z\\sim 7.4$ and derive an analytic formula for the cosmic SFH by fitting the updated data. In Section \\ref{meta}, I summarize the current measurements on the evolution of the metallicity in galaxies with cosmic time and argue that the data indicate a consistent picture for the cosmic metallicity evolution when the redshift-dependent relation between the metallicity and the galaxy stellar mass is considered. In Section \\ref{sample}, I describe the \\swf\\ GRB sample that is used for the current work, show the luminosity distribution of the GRBs, and discuss the selection biases involved in the detection and localization of \\swf\\ GRBs and the measurement of their redshifts. In Section \\ref{model}, the model that I adopt for calculating the GRB rate history is outlined, which includes assumptions about the probability distribution function, the SFR, the evolution of cosmic metallicity, and the form of the GRB LF. In Section \\ref{results}, I present the results calculated with the model, and fit the observed distribution of the luminosity and redshift for the whole GRB sample and a bright GRB subsample. In Section \\ref{conclusion}, I draw conclusions and discuss some implications of this work.  Throughout the paper, I assume a flat universe with $\\Omega_\\m = 0.3$, $\\Omega_\\Lambda=0.7$, and $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "\\label{conclusion}  I have presented an updated cosmic SFH up to redshift $z=7.4$. The updated sample of SFR data are obtained by adding the new UV and IR measurements on the SFR density of \\citet{bou07,bou08} and \\citet{red08} to the `good' data sample compiled by \\citet{hop06}. The two joined UV+IR measurements of \\citet{red08} at $z\\sim 2.3$ and $z\\sim 3.05$ agree well with previous measurements of SFRs in the redshift interval $1\\la z\\la 4$, and are consistent with a flat evolution SFH in this interval. The UV measurements of \\citet{bou07,bou08} at $z\\sim 3.8$, 5, 5.9, 7.4 and an estimate at $z\\sim 10$ significantly expand the SFR data sample at $z\\ga 4$, and are broadly consistent with the previous results in $3.5\\la z\\la 6$. The updated sample provides a consistent picture for the cosmic SFH up to $z\\sim 7.4$ (Fig.~\\ref{sfr_his}), although the dust correction at $z\\ga 4$ is highly uncertain which results large uncertainties in the SFH at high redshift \\citep{dro08}.  The updated SFH still under-produces the GRB rate density at high redshift when compared to the \\swf\\ GRB redshift distribution (Fig.~\\ref{rgrb_sfr}), confirming the previous claim \\citep{dai07,le07,kis07}. The discrepancy is investigated under the assumption that long-duration GRBs trace both the star formation and the metallicity evolution, as motivated by the observations that long GRBs occurred in star-forming galaxies with low metallicities and the theoretical study on the collapsar model which shows that GRBs can be produced only from low-metallicity massive stars.  Since the cosmic metallicity decreases with redshift \\citep{pet99,pro03,rao03,kob04,kew05,kew07,kul05,sav05,sav06}, it is natural to expect that the ratio of the GRB rate to the SFR must increase with redshift if the scenario that long GRBs are produced by the death of massive stars with low metallicity \\citep{mac99,mac01,hir05,yoo05,woo06a} is correct. Adopting a simple model for relating the GRB rate density to the SFR density and the cosmic metallicity evolution \\citep{lan06} and assuming a flux limit for the \\swf\\ detector, I have shown that the redshift distribution of the 64 \\swf\\ GRBs with measured redshifts and calculated luminosities can be successfully fitted by the updated SFH with a threshold in the metallicity for GRB production (Figs.~\\ref{qq_distr}--\\ref{qq_int2}).  \\citet{kis07} have considered several possibilities for the cause of the discrepancy between the the \\swf\\ GRB rate and the SFH. They have shown that the Kolmogorov-Smirnov test disfavors an interpretation as a statistical anomaly. Selection effects are also not likely to cause an increased efficiency in detecting hight-redshift GRBs. Although Kistler et al. have argued that alternative causes are possible (e.g., evolution in the fraction of binary systems, an evolving IMF of stars, etc), the results in this paper indicate that the cosmic metallicity evolution may be the simplest interpretation.  However, the results show an excess in the number of GRBs with low luminosity ($L_\\iso\\sim 3.8\\times 10^{49}$ erg s$^{-1}$; Fig.~\\ref{liso_distr}) and at low redshifts ($z\\la 0.7$; Figs.~\\ref{qq_distr2} and \\ref{qq_int2}). The existence of an excess is confirmed by the up-to-date \\swf\\ GRBs with measured redshifts, detected by 31 March 2008 (Fig.~\\ref{qq_distr2u}). Although it might simply be caused by statistical fluctuations, the observed excess could also be consistent with the speculation that there is a unique population of intrinsically faint and nearby GRBs (Section \\ref{results}).  \\begin{figure} \\includegraphics[angle=0,scale=0.51]{grb_rate2u.eps} \\caption{Distribution of $Q$ for all the 85 \\swf\\, GRBs with redshifts, detected by 31 March 2008 between and including GRBs 041220 and 080330 (the solid histogram). The dashed curve shows the distribution given by equation (\\ref{N1}) with the parameters obtained by fitting the luminosity distribution (Fig.~\\ref{liso_distr}, the dashed curve), renormalized by the total number of GRBs. The first data point has an offset about $3.0$-$\\sigma$ from the dashed curve. The overall $\\chi_\\rr^2 = 1.21$. If the first point is not included in the calculation of chi-squares, then $\\chi_\\rr^2 = 0.28$. } \\label{qq_distr2u} \\end{figure}  The GRB sample used in this paper \\citep[and in][]{kis07} is almost definitely incomplete and non-uniform, because of the complex selection biases in the redshift measurement discussed in Section \\ref{biases}. In fact, all the current works on the redshift distribution of GRBs have such a problem. However, \\citet{kis07} have argued that none of the selection biases appears to be able to increase the overall observability with redshift and account for the enhancement in the GRB rate relative to the SFR.  The number of GRBs in the sample is small, which also prohibits one from obtaining a strict constraint on the parameters in the GRB LF and the cosmic metallicity evolution.  Despite the above problems, the results of this work suggest that the rise of the observed \\swf\\ GRB rate relative to the SFR is compatible with an interpretation by the evolution of the cosmic metallicity. Once the problems are solved or significantly alleviated in future, a significantly improved and enlarged sample of GRBs with measured spectra and redshifts will become available. Then, by comparing the observed GRB rate history to the SFH determined with other approaches, it will be possible to probe the cosmic metallicity evolution with GRBs."
    },
    "0710/0710.4158_arXiv.txt": {
        "abstract": "{Measurements of the arrival directions of cosmic rays have not revealed their sources. High energy neutrino telescopes attempt to resolve the problem by detecting neutrinos whose directions are not scrambled by magnetic fields. The key issue is whether the neutrino flux produced in cosmic ray accelerators is detectable. It is believed that the answer is affirmative, both for the galactic and extragalactic sources, provided the detector has kilometer-scale dimensions. We revisit the case for kilometer-scale neutrino detectors in a model-independent way by focussing on the energetics of the sources. The real breakthrough though has not been on the theory but on the technology front: the considerable technical hurdles to build such detectors have been overcome.  Where extragalactic cosmic rays are concerned an alternative method to probe the accelerators consists in studying the arrival directions of neutrinos produced in interactions with the microwave background near the source, i.e. within a GZK radius. Their flux is calculable within large ambiguities but, in any case, low. It is therefore likely that detectors that are larger yet by several orders of magnitudes are required. These exploit novel techniques, such as detecting the secondary radiation at radio wavelengths emitted by neutrino induced showers.}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.4191_arXiv.txt": {
        "abstract": "In paper I of this series we discuss how magnification bias distorts the 3D correlation function by enhancing the observed correlation in the line-of-sight (LOS) orientation, especially on large scales. This lensing anisotropy is distinctive, making it possible to separately measure the galaxy-galaxy, galaxy-magnification {\\it and} magnification-magnification correlations. Here we extend the discussion to the power spectrum and also to redshift space. In real space, pairs oriented close to the LOS direction are not protected against nonlinearity even if the pair separation is large; this is because nonlinear fluctuations can enter through gravitational lensing at a small transverse separation (or i.e. impact parameter). The situation in Fourier space is different: by focusing on a small wavenumber $k$, as is usually done, linearity is guaranteed because both the LOS and transverse wavenumbers must be small. This is why magnification distortion of the galaxy correlation appears less severe in Fourier space. Nonetheless, the effect is non-negligible, especially for the transverse Fourier modes, and should be taken into account in interpreting precision measurements of the galaxy power spectrum, for instance those that focus on the baryon oscillations. The lensing induced anisotropy of the power spectrum has a shape that is distinct from the more well known redshift space anisotropies due to peculiar motions and the Alcock-Paczynski effect. The lensing anisotropy is highly localized in Fourier space while redshift space distortions are more spread out. This means that one could separate the magnification bias component in real observations, implying that potentially it is possible to perform a gravitational lensing measurement  without measuring galaxy shapes. ",
        "introduction": "\\label{intro}  The effect of magnification bias on the 3D galaxy/quasar correlation function was studied in paper I \\cite{3Dpaper1}. (Galaxy and quasar can be considered synonymous hereafter.) With the important exception of the classic paper by Matsubara \\cite{matsubara00}, previous work on how magnification bias modifies clustering observations has largely focused on the 2D angular correlation function \\cite{TOG84,J95,BTP95,VFC97,MJV98,MJ98,EGmag03,ScrantonSDSS05,menard,bhuv,JSS03}. The novelty of the 3D correlation function, as emphasized by \\cite{matsubara00} and \\cite{3Dpaper1}, is that magnification bias makes it anisotropic. In this paper, we extend our previous analysis by studying the anisotropy in Fourier and redshift space. At first sight, the extension to Fourier space might seem a trivial exercise. The calculations are indeed straightforward, but as we will see, the results are far from obvious: there are important qualitative differences between the results in Fourier space and real space that go beyond the usual wavenumber-position ($k$-$x$) duality.  Let us recall the situation in real space, as depicted in paper I. The anisotropy of the observed 3D correlation function can be understood intuitively as follows. The correlation function is measured by pair counts of galaxies. A pair of galaxies that are aligned along the line-of-sight (LOS) behave differently from a pair oriented transverse to the LOS. In the former case, the closer galaxy can lens the background one. The same does not happen in the transverse orientation. The net effect is an anisotropy in the observed correlation function, induced by gravitational lensing (or equivalently, magnification bias; we refer to this effect as magnification distortion). This reasoning suggests the gravitational lensing corrections are largest for a pair of galaxies oriented along the LOS. Indeed, the corrections can be quite significant: consider for instance a LOS separation of $\\sim 100$ Mpc/h; the intrinsic galaxy correlation is rather weak on such a large scale, but the lensing induced correction can be quite substantial, since for the LOS orientation, the relevant lensing impact parameter, i.e. transverse separation, is zero (keep in mind also that the lensing effect grows with the LOS separation while the intrinsic galaxy correlation generally drops with separation). In other words, for the LOS orientation, scales that otherwise would be considered {\\it linear} can in fact be secretly affected by nonlinear fluctuations via lensing. A large separation $|\\delta {\\bf x}| = \\sqrt{\\delta\\chi^2 + |\\delta {\\bf x_\\perp}|^2}$ does not guarantee linearity because nonlinear fluctuations can sneak in through lensing with a small transverse separation $|\\delta {\\bf x_\\perp}|$.  This peculiar mixing of linear intrinsic galaxy fluctuations with nonlinear lensing fluctuations does not arise in Fourier space. A small net wavenumber $|{\\bf k}|$ is sufficient to guarantee that both the LOS component $k_\\parallel$ and the transverse component $|{\\bf k_\\perp}|$ are small. Nonlinear lensing corrections cannot sneak in as long as one focuses on a small $|{\\bf k}|$, as is usually done. This immediately tells us that the anisotropy of the observed galaxy correlation must appear milder in Fourier space. Our primary goal here is to quantify this.  As discussed in paper I, the magnification bias induced anisotropy in the observed 3D galaxy correlation has two implications. First, precision measurements of the galaxy correlation must take into account such magnification distortion. These include future galaxy surveys that hope to determine the baryon oscillation scale to high accuracy \\cite{eisenstein,2dFa,2dFb,hutsi,tegmark,baotheory,baoexp}. Second, the distinctive anisotropy pattern makes it possible in principle to separately measure the galaxy-galaxy, galaxy-magnification {\\it and} magnification-magnification correlations. (The last correlation was generally ignored in previous papers that focused on angular correlation between galaxies at widely separated redshifts \\cite{MJ98,EGmag03,ScrantonSDSS05,menard,bhuv,JSS03}, where the galaxy-magnification correlation dominates.) Achieving such a separation requires that one understands other sources of anisotropy. We therefore extend our Fourier analysis here to incorporate the anisotropy due to both peculiar motions and the Alcock-Paczynski effect.  The rest of the paper is organized as follows. In \\S \\ref{distort}, we derive and numerically compute the magnification distortion of the observed galaxy clustering -- \\S \\ref{correlation} summarizes the results from paper I on the correlation function while \\S \\ref{pk} focuses on the power spectrum. Redshift space distortion due to peculiar motion is next incorporated in \\S \\ref{anisotropy}, and we conclude in \\S \\ref{discuss}. In Appendix \\ref{app:AP}, we discuss the Alcock-Paczynski anisotropy.  Before we start, it is useful to point out several related papers. Vallinoto et al. \\cite{scott} explored the impact of lensing, especially magnification bias, on the baryon oscillation signal in the real space correlation function. Their results are consistent with ours in paper I, though they focus exclusively on pair separations that are oriented transverse to the LOS, and their work is therefore more connected to our paper on the angular correlation function \\cite{wLHG}. Wagner et al. \\cite{steinmetz} examined the anisotropy of the 3D correlation that is introduced by light cone effects. A discussion of the classic paper by Matsubara \\cite{matsubara00} can be found in paper I. Both \\cite{matsubara00} and paper I focused on the real/configuration space correlation function, though peculiar motions and the Alcock-Paczynski effect are also treated in \\cite{matsubara00}. A recent paper by Zhang \\& Chen \\cite{zhangchen} explored the effects of gravitational lensing in Fourier space in the context of supernova observations. LoVerde et al. \\cite{LHGisw} examined the impact of magnification bias on integrated Sachs-Wolfe measurements. ",
        "conclusions": "\\label{discuss}  In paper I, we examined the effects of magnification distortion in real/configuration space, and here we have extended the analysis to Fourier and redshift space. The observed galaxy/quasar correlation is endowed with a distinctive lensing induced anisotropy. This is encapsulated in eq. (\\ref{scaling2}) for real space and eq. (\\ref{Pgmu}), (\\ref{Pmumu}), (\\ref{Pgg}) and (\\ref{Pobs}) for Fourier space: the linear dependence on the LOS separation $\\delta\\chi$ in the real space galaxy-magnification correlation gives rise to the multiplicative window $G$ in the galaxy-magnification power spectrum. Likewise, the independence of the real space magnification-magnification correlation on $\\delta\\chi$ accounts for the appearance of the multiplicative window $|\\tilde W_\\parallel |^2$ in the magnification-magnification power spectrum.  Qualitatively, the galaxy correlation becomes enhanced for the transverse modes in Fourier space, and for pairs oriented along the LOS in real space. Quantitatively, the degree of enhancement is rather different for the dual spaces. As explained in \\S \\ref{intro} and \\S \\ref{pk}, magnification distortion is less severe in Fourier space: as long as one focuses on modes with a small wavenumber $k$, as is usually done when obtaining cosmological constraints, both the intrinsic galaxy fluctuations and the lensing fluctuations are in the linear regime. In real space, even if the pair separation is large, one cannot help but mix up linear intrinsic galaxy fluctuations with nonlinear lensing fluctuations, as long as one considers the LOS orientation. Incidentally, this implies that a linear galaxy bias is a better approximation in Fourier space than in real space.  The above findings suggest that in precision measurements of the galaxy power spectrum, such as those that attempt to use the baryon oscillation scale to constrain dark energy \\cite{eisenstein,2dFa,2dFb,hutsi,tegmark,baotheory,baoexp}, the simplest way to immunize against magnification bias could be to go to Fourier space, and remove from consideration the small $k_\\parallel$ (transverse) modes where magnification bias has the largest effect. However, in a photometric redshift survey, these are probably the modes with the highest signal-to-noise, and the Fourier space also suffers from possible complications due to a non-Poissonian shot noise \\cite{sss}. A better strategy is perhaps to face the magnification bias corrections head on, and use the full 3D information to constrain and measure them -- after all, they contain interesting cosmological information too.  Precisely how magnification bias corrections might shift the baryon oscillation scale is a subject worthy of a separate paper. The precise shift will depend on exactly how the oscillation scale is extracted from data. A preliminary investigation in real space, where the acoustic oscillations manifest as a single local maximum whose position is easily defined, was presented in paper I: as discussed at the end of \\S \\ref{pk}, the impact on measurements of the dark energy equation of state can be significant (shifting it by up to $\\sim 15 \\%$). It should also be emphasized that the baryon oscillations are not the only large scale features of interest. The radiation-matter equality peak at around $k \\sim 0.01$ h/Mpc contains valuable cosmological information, but it can be seen from Fig. \\ref{pkall1} - \\ref{pkall3} panels (b) that the power spectrum at this scale is likely significantly affected by magnification bias.  We have incorporated the effects of peculiar motion (and redshift inaccuracy) on the power spectrum anisotropy in \\S \\ref{anisotropy} (the Alcock-Paczynski effect is further incorporated in Appendix \\ref{app:AP}). The main conclusion is that the lensing induced features remain rather robust, thanks to their localized nature to small $k_\\parallel$'s. It should be in principle possible to separately measure from data the galaxy-galaxy, galaxy-magnification and magnification-magnification power spectra, exploiting their different shapes in the $k_\\parallel - k_\\perp$ plane (see eq. [\\ref{approx}]). Exactly how to do so in an optimal fashion when faced with noisy data deserves further study. In particular, different galaxy types have a different galaxy bias and number count slope; how to best weigh their relative contributions to the observed power spectrum?  Lastly, the findings in this paper and paper I cast a new light on the well known excess correlations seen in pencil beam surveys \\cite{pencil,pencil2}. Could these be the result of enhanced correlations due to magnification bias, particularly if baryon oscillations are taken into account (see \\cite{eis})? Could the peculiar features seen in the power spectrum analysis \\cite{pencil2} be due in part to the corresponding multiplicative windows $G$ and $|\\tilde W_\\parallel |^2$? We hope to address these questions in the future."
    },
    "0710/0710.4969_arXiv.txt": {
        "abstract": "The nature of the excess near-infrared emission associated with the magnetic white dwarf commonly known as SDSS 1212 is investigated primarily through spectroscopy, and also via photometry.  The inferred low mass secondary in this system has been previously detected by the emission and variation of H$\\alpha$, and the $1-2.5$ $\\mu$m spectral data presented here are consistent with the presence of a late L or early T dwarf.  The excess flux seen beyond 1.5 $\\mu$m in the phase-averaged spectrum is adequately modeled with an L8 dwarf substellar companion and cyclotron emission in a 7 MG magnetic field.  This interesting system manifests several observational properties typical of polars, and is most likely an old interacting binary with a magnetic white dwarf and a substellar donor in an extended low state. ",
        "introduction": "Cataclysmic variables are spectrally exotic binaries, displaying a diversity and complexity of features, behavior, and variability, effectively masking the spectral signatures of their coolest secondaries.  Despite unyielding efforts to unambiguously and directly detect likely substellar companions in either non-magnetic or magnetic cataclysmic variables, these altered degenerate dwarfs remain largely elusive \\citep*{lit03}.  The mass determination of the substellar donors in two non-magnetic accreting binaries reported by \\citet*{lit06} and \\citet*{lit07} remain rare among cataclysmic variables, although the determination of their substellarity was indirect.  Cataclysmic variables containing strongly magnetic white dwarfs (called polars because of the highly polarized nature of their light) offer a slight advantage by lacking the often overwhelming luminosity from an accretion disk.  On the other hand, their magnetic fields generate copious cyclotron radiation (in addition to flux from direct accretion), which is not present in non-magnetic cataclysmic variables.  Polars are also known for exhibiting long periods of quiescence in which many secondaries have been directly observed, although none of spectral type later than M \\citep*{har05}.  The recently recognized class of low accretion rate polars may provide decent infrared hunting grounds for substellar companions, if present, as their cyclotron harmonics appear mainly in the optical due to $B\\sim50$ MG fields.  It has been suggested that no mass transfer takes place in these systems, but rather the efficient capture of secondary wind, and that their separations are consistent with being essentially detached \\citep*{sch05b}.  Yet, it is difficult to certify that these binaries are not simply in protracted low states such as seen in other polars \\citep*{sch07,har04}.  Recent {\\em Spitzer} IRAC observations of several polars indicate flat infrared spectral energy distributions which are difficult to reconcile with cool degenerate companions alone, and may indicate circumbinary dust as well as underlying cyclotron continuum emission \\citep*{bri07, how06}.  It will be interesting to see the results of similar mid-infrared observations with a larger sample of polars as well as observations at longer wavelengths where there should be little or no cyclotron emission even in $B\\approx10$ MG fields.  The first published {\\em Spitzer} IRS spectrum of a polar strongly supports the presence of circumbinary dust at EF Eridani \\citep*{hoa07b}.  This paper represents a spectroscopic effort to directly detect the almost certain substellar companion to the cool magnetic white dwarf SDSS J121209.31+013627.7 (SDSS 1212).  Its binarity was reported with optical spectra supplemented with a single $J$-band photometric observation, limiting the spectral type of the secondary to L5 or later, and was initially thought to represent a detached system \\citep*{sch05a}.  This possibility was still considered viable despite a second study which detected strong cyclotron emission -- a likely indicator of mass transfer -- in $K$-band time series photometry \\citep*{deb06}.  Both \\citet*{koe06} and \\citet*{bur06a} detected optical variability and a light curve similar to that shown by polars in low states.  \\citet*{bur06a} further showed that the amplitude of the modulations was larger at blue and ultraviolet wavelengths than in the red, consistent with a hot accretion spot responsible for the cyclotron emission.  \\citet*{bur06a} also detected x-ray emission from SDSS 1212, conclusively proving accretion is taking place.  This paper presents cross-dispersed $1-2.5$ $\\mu$m spectroscopy and spectral modeling, along with $HK$ photometry, in order to constrain the source(s) of the observed near-infrared excess emission. ",
        "conclusions": "\\subsection{Past and Future Evolution}  The gross evolution cataclysmic variables should consist of three fundamental phases:  (1) A short- to long-lived phase as a detached main sequence binary which subsequently evolves to shorter orbital periods through common envelope evolution as the primary leaves the main sequence and becomes a first ascent or asymptotic giant.  Angular momentum is transferred, via friction, into the envelope itself, ejecting it in the process \\citep*{pac76}.  If the orbit is still too large for the secondary to make contact with its Roche lobe, then three mechanisms can bring this about, in principle: gravitational radiation, magnetic braking, or nuclear evolution.  The timescales of these processes depend on the masses involved, but if one assumes a canonical 0.6 $M_{\\odot}$ white dwarf and a 0.2 $M_{\\odot}$ secondary (corresponding to the mass distribution peak among detached, unevolved, low mass companions to white dwarfs; \\citealt*{far05a}), then both gravitational radiation and magnetic braking operate within Gyr timescales for period ranges $P\\la3$ hr and $P>3$ hr respectively, whereas nuclear evolution will play no role within this timeframe \\citep*{pat84}.  (2) A phase in which the shrinking Roche lobe of the secondary becomes smaller than its radius, causing unstable mass transfer onto the white dwarf.  In principle this is a runaway inspiral until the mass ratio is close to one, but if the secondary has $M<0.5$ $M_{\\odot}$ to begin with (as it should in the vast majority of cases based on the aforementioned companion mass distribution), then conservation of angular momentum competes to partially counterbalance the losses due to any braking (magnetic or tidal) and gravitational radiation \\citep*{nel01}.  At some point the trend reverses outward as the orbital evolution becomes dominated by conservation of angular momentum and any mass lost from the system entirely (e.g. nova-like outbursts, accretion-driven winds).  (3) A final, debatable, phase consisting of a once more detached or very low accretion rate semi-detached binary comprised of two degenerate objects in an expanded orbit which persists (now subject only to gravitational radiation) up to 10 Gyr or so, depending on the final masses and separation.  This eventual state depends on the thermal and nuclear history of the secondary but requires a degenerate secondary core in order to cease expansion permanently.  Thus, a system may end its existence as a white dwarf with a close, altered substellar companion, both objects cooling gradually toward oblivion.  The ultimate fate of catacylsmic variables outlined above is uncertain for two fundamentally different reasons.  The first issue is that the outcome itself is unknown, with end predictions for secondaries including; complete stripping, tidal disruption, persistence as semi-detached substellar donors with very low mass transfer rates, and survival as detached substellar or even planetary mass objects \\citep*{lit03,kol99,pat98,how97}.  The second issue is detectability, which itself divides into two categories: searching for either erupting systems or their burnt out remnants.  \\subsection{Recently Deceased or Born Dead?}  So far, astronomers who study cataclysmic variables have concentrated primarily on those systems known to be eruptive in order to search for and understand possible end states, including substellar companions.  The means of identifying such low mass secondary systems relies heavily on semi-empirical and theoretical relations between observed periods and spectral type, mass ratio, effective temperature, and mass transfer rate \\citep*{lit06,lit07,har05,pat05,kol99, how97}. Without models, there is no means of determining whether a given substellar companion is migrating inward or outward, i.e. approaching or leaving its period minimum.  With the prior history of these interacting systems unknown, the beginning and end states may appear identical, theoretically.  Additionally, it is difficult to know with certainty what the size of the seconday star is in such systems, especially if the mass transfer rate is low -- consistent with quiescence or an underfilled Roche lobe (i.e. a non-catacylsmic variable).  If a companion does indeed have a radius smaller than its Roche lobe, then strictly speaking it is detached, and the question remains: is Roche lobe contact in its future or in its past?  The recently identified low mass transfer magnetic binaries reported in \\citet*{sch05b} are good candidates for being detached due to the combination of the inferred accretion rates, orbital periods, and extremely low or absent x-ray emission.  {\\em But if EF Eridani can turn off for nearly 10 yr, then so can other polars} \\citep*{sch07}.  SDSS 1212 sits somewhere on the border between a typical polar and the low accretion rate systems, having respectable x-ray accretion luminosity and a relatively low inferred mass transfer rate \\citep*{bur06a}. If one assumes the substellar companion to SDSS 1212 has a mass and effective temperature similar to WD 0137$-$349B (L8V, see Table \\ref{tbl3}; \\citealt*{bur06b}), then its unperturbed radius should be 0.09 $R_{\\odot}$ for ages between $1-5$ Gyr \\citep*{bar03}, but its Roche lobe should be located roughly at 0.11 $R_{\\odot}$ for a 0.60 $M_{\\odot}$ primary.  Therefore, the companion is either: 1) recently deceased, far out of thermal equilibrium and inflated as expected for an object past period minimum with the implied long history of mass-radius imbalance; or 2) born dead, intrinsically detached, and slamming the magnetic white dwarf with all of its (irradiatively driven?) wind.  This raises similar questions about the nature of all other low accretion rate magnetic binaries.  Depending on the past thermal and nuclear history of the secondary, especially if the core is degenerate, and how much mass escapes the binary completely and adiabatically, it is quite possible for the orbital period to increase significantly over the minimum into a region where no interaction (or wind accretion only) occurs \\citep*{how97,jea24}.  Without any distinct spectral features or mass-radius anomalies in such detached binaries, there may be little or no means to differentiate empirically between the pre- and post- states of these innately interacting systems.  \\subsection{The Frequency of the Doubly Dead}  Objects must first be detected as a catacylsmic variable in order to look for substellar companions; hence, selection effects play a large role in the number of systems suspected to harbor such cool degenerate companions \\citep*{lit03}.  However, this overlooks the obvious question: if the end states are detached double degenerate systems, where are they?  If one restricts a search to eruptive systems, then the potential for success is limited from the beginning for reasons stated above.  Why not search directly for systems which are detached and which may represent final states?  This seems to be the only way to reconcile the well-known discrepany between the number of predicted objects past period bounce and the number observed or suspected from observation \\citep*{pat98,how97}.  Clearly, the best place to look for these types of end states is in the infrared where cool, Jupiter-sized companions easily outshine their tiny white dwarf primaries \\citep*{far05a,far05b}. There are only two issues associated with a search around white dwarfs for extinct cataclysmics -- the white dwarf must first have been identified or previously known and one must have sufficient sensitivity to detect its cooling leftovers.  If the end state white dwarfs are very cool and far away such that $V>21$ mag (e.g. $T_{\\rm eff}<5000$ K and $d>100$ pc for $M=0.9$ $M_{\\odot}$), then only the Sloan Digital Sky Survey will detect them but will be quite insensitive to any substellar companions.  While the above conditions may not be representative, given that cataclysmic variables appear to be roughly 100 times less abundant than white dwarfs \\citep*{pat84}, a large sample of targets would be required to perform a statistically rigorous search; if all cataclysmic variables are in this state, a search would only yield 1 per 100 white dwarfs (not very encouraging).  If these close but detached systems exist in any significant numbers, they have not yet been identified from various surveys of relatively nearby white dwarfs.  Spatially unresolved L type companions, corresponding to $0.05-0.07$ $M_{\\odot}$ for unevolved ages of $1-5$ Gyr, are rare, with only 3 known among well over $N=1000$ targets \\citep*{hoa07a,bur06b,far05a,far04, far04b,wac03}.  Spatially unresolved T type companions, corresponding to $0.02-0.05$ $M_{\\odot}$ for unevolved ages of $1-5$ Gyr, are also rare, with none detected or suspected in various surveys totalling $N\\approx200$ targets (\\citealt*{mul07,han06,far05a}; Farihi, Becklin, \\& Zuckerman 2008, in preparation).  It is possible that stellar evolution has yet to produce such systems; i.e. the initial mass function for secondaries in such systems, convolved with the finite age of the Galaxy, implies that the vast majority of these systems will continue in phases (1) and (2) for up to 10 Gyr.  It could be that the lack of detected end states is an observational bias; if these systems are still marginally semi-detached then the signature of the secondary star could be masked by the accretion disk in dwarf novae and by cyclotron emission in polars. Alternatively, phase (3) never occurs and the companion is cannibalized, evaporated, stripped or tidally destroyed.  This would also explain the lack of observed end states.  For completeness, something ought to be said about the empirical frequency of zero age cataclysmic variables with a brown dwarf secondary \\citep*{pol04}.  Since their pre- states would consist of a white dwarf with a detached substellar companion, the above arguments and evidence apply ($f<0.5$\\%).  Additionally, the progenitors of such pre- states would be main sequence stars with brown dwarf companions in the few AU range, which are also extremely rare, occuring with a well-constrained frequency, $f<0.5$\\% \\citep*{but06,mar05}. Hence, there should be no doubt that cataclysmic variables form with brown dwarf secondaries quite rarely.  \\subsection{Detached Companions to Magnetic White Dwarfs}  The same argument which applies to detached substellar companions to white dwarfs also applies to detached stellar companions to magnetic white dwarfs, although in this latter case finding them should be easier in principle \\citep*{lie05}.  One might suggest a sensitive search for magnetism among known lists of close but detached, and wide common proper motion companions.  Large lists of these white dwarf plus red dwarf systems exist in the following references: for known radial velocity pairs see \\citet*{mor05}, for spatially unresolved pairs see \\citet*{hoa07a,far06,far05a,wac03}, and for wide common proper motion pairs see \\citet*{far05a, sil05}.  All of these lists contain targets which are likely to be superior to white dwarfs from the Sloan Digital Sky survey since they are brighter and, hence, more sensitive searches could be conducted. However, it should be noted that the main observational indicators of a magnetic white dwarf, Zeeman splitting and circular spectropolarimetry, are strongest in H$\\alpha$ and H$\\beta$.  Unfortunately, in unresolved white dwarf plus red dwarf binaries (i.e., the potential cataclysmic variables) these features, especially H$\\alpha$, are often completely filled-in by the red dwarf.  This may be another reason for the curious lack of detected pre-polar systems to be considered along with those already discussed by \\citet*{lie05}."
    },
    "0710/0710.1311_arXiv.txt": {
        "abstract": "Our current understanding of the star formation histories of early-type galaxies is reviewed, in the context of recent observational studies of their ultra-violet (UV) properties. Combination of UV and optical spectro-photometric data indicates that the bulk of the stellar mass in the early-type population forms at high redshift ($z>2$), typically over short timescales ($<1$ Gyr). Nevertheless, early-types of all luminosities form stars over the lifetime of the Universe, with most luminous ($-23<M(V)<-21$) systems forming 10-15\\% of their stellar mass after $z=1$ (with a scatter to higher value), while their less luminous ($M(V)>-21$) counterparts form 30-60\\% of their mass in the same redshift range. The large scatter in the (rest-frame) UV colours in the redshift range $0<z<0.7$ indicates widespread low-level star formation in the early-type population over the last 8 billion years. The mass fraction of young ($<1$ Gyr old) stars in luminous early-type galaxies varies between 1\\% and 6\\% at $z\\sim0$ and is in the range 5-13\\% at $z\\sim0.7$. The intensity of recent star formation and the bulk of the UV colour distribution is consistent with what might be expected from minor mergers (mass ratios $\\lesssim$ 1:6) in a $\\Lambda$CDM cosmology. ",
        "introduction": "The star formation histories (SFHs) of early-type galaxies have been the subject of intense and controversial debate in modern astrophysics. The classical `monolithic collapse' hypothesis for early-type evolution followed the model of Eggen et al. for the formation of the Galaxy\\citet{ELS62}. Refined and implemented by others \\citep{Larson74,Chiosi2002}, this model postulated that stellar populations in early-type galaxies form in short, highly efficient starbursts at high redshift ($z\\gg1$) and evolve purely passively thereafter. The \\emph{optical} properties of the early-type population and, in particular, their strict obedience to simple scaling relations are remarkably consistent with such a simple, largely empirical hypothesis. The small scatter in the early-type `Fundamental Plane' \\citep{Jorg1996,Saglia1997} and its apparent lack of evolution with look-back time \\citep{Forbes1998,Peebles2002,Franx1993,Franx1995,VD1996}, coupled with the homogeneity and lack of redshift evolution in their optical colours \\citep{BLE92,Bender1997,Ellis97,Stanford98,Gladders98,VD2000} is strong evidence that the bulk of the stellar population in early-type galaxies forms at high redshift. Furthermore, the over-abundance of $\\alpha$ elements in these systems \\citep{Thomas1999} indicates that the star formation timescales are shorter than the typical timescales for the onset of Type Ia supernovae (SN). For example, if Type Ia progenitors largely explode within $\\sim1$ Gyr of a starburst \\citep{Greggio1983}, then the bulk of the star formation in these galaxies probably takes place on timescales shorter than a Gyr.  While it is consistent with the optical properties of early-type galaxies, a monolithic SFH does not sit comfortably within the currently accepted $\\Lambda$CDM galaxy formation paradigm in which the stellar mass in local early-type systems is thought to accumulate over the lifetime of the Universe. Following the seminal work of Toomre\\citet{Toomre_mergers}, who showed that most galaxy collisions end in rapid merging and postulated the formation of spheroidal systems as end-products of such merger activity, the mechanics of galaxy interactions \\citep{Barnes1992a,Barnes1992b,Hernquist1993} and their link to the formation of early-type systems \\citep{Barnes1996,Bender1996} have been studied in considerable detail. Modern `semi-analytical' models of galaxy formation, within the $\\Lambda$CDM paradigm, create early-types through `major mergers', where the mass ratio of the merging progenitors is 3:1 or lower. The constituent stellar mass of early-type galaxies is predicted to form both quiescently in their progenitors and in efficient starbursts when these progenitors merge\\citep{K96,K98,Cole2000,Hatton2003,Khochfar2003}.  While theoretical arguments may be compelling, the strongest evidence for the role of interactions in shaping early-type galaxy evolution and inducing coincident star formation comes from observation, both in the local Universe and at high redshift. 40\\% of local ellipticals contain dust lanes \\citep{Sadler1985}, while $\\sim75$\\% contain nuclear dust and by implication gas \\citep{Tomita2000,Tran2001}. The gas is often kinematically decoupled from the stars, indicating, at least in part, an external origin, e.g. through the accretion of a gas rich satellite \\citep{sauron5}. Up to two-thirds of nearby early-types contain shells, ripples and morphological disturbances \\citep{Malin83,VD2005} and a significant fraction exhibit kinematically decoupled cores \\citep{sauron2}, both of which are evidence for interactions in the recent past. While the detection of such spatially resolved fine structure is possible only for galaxies in our local neighbourhood, clear signatures of recent star formation have been found in individual early-type systems out to modest redshifts \\citep{Trager2000a,Trager2000b,Fukugita2004}.  Deep ground and space-based imaging are increasingly providing access to galaxy populations over the last ten billion years of look-back time. A significant fraction ($\\sim$30\\%) of luminous early-type systems at high redshift ($0.4<z<0.8$) exhibit blue cores, indicative of merger-driven starbursts triggered by the accretion of low-mass gas-rich companions \\citep{Ferreras2005,Menanteau2001a}. Furthermore, most blue cores are predominantly contained in early-type galaxies \\citep{Ferreras2005} and, not unexpectedly, they are typically accompanied by spectral lines characteristic of recent star formation \\citep{Ellis2001}. A fundamental consequence of early-type evolution in the $\\Lambda$CDM model is the gradual loss of late-type progenitors \\citep{VD2001a,Kaviraj2005a} and a corresponding increase in the fraction of early-type galaxies. Numerous observational studies have detected such an evolution in the morphological mix of the Universe. While $\\sim$80 percent of galaxies in the cores of local clusters have early-type morphology \\citep{Dressler80}, a higher fraction of spiral (blue) galaxies have been reported in clusters at high redshift \\citep{BO84,Dressler97,Couch98,VD2000,Margoniner2001,Andreon2004,Borch2006,Bundy2006}, combined with increased rates of merger and interaction events \\citep{VD99}. Similar results have been found in large-scale survey data which suggest that the mean mass density on the red sequence (which is dominated by early-type systems at all redshifts) has at least doubled since $z=1$ (e.g. Bell et al. 2004 \\citep{Bell2004}; Faber et al. 2007 \\citep{Faber2007}). A significant body of observational evidence thus indicates that the SFH of \\emph{at least} some early-type galaxies, and perhaps the early-type population as a whole, deviates significantly from the expectations of the classical monolithic collapse hypothesis, both in terms of their structural evolution and the star formation experienced by them over the lifetime of the Universe.  Although the majority of early-type studies in the past have focussed on the optical spectrum, a significant drawback of optical photometry is its lack of sensitivity to moderate amounts of \\emph{recent star formation} (RSF). While red optical colours do imply a high-redshift formation epoch for the bulk of the stellar mass in early-type systems, the optical spectrum remains largely unaffected by the \\emph{minority} of stellar mass that is expected to form in these objects at low and intermediate redshifts\\citep{Kaviraj2005a}. As a result, it is difficult to resolve early-type SFHs over the last 8 billion years (where the predictions of the two competing models diverge the most) using optical colours alone.  An ideal route to quantifying early-type SFHs at low and intermediate redshift is to exploit a sensitive indicator of RSF. If the early-type population could be studied, over a large range redshift, using such an indicator, then the steady accumulation of stellar mass could be robustly quantified and strong constraints applied to the predictions of galaxy formation models over the last 8 billion years of look-back time. Spectroscopic indicators of RSF already exist, such as the commonly used $H_{\\beta}$ index, higher order Balmer lines such as $H_{\\gamma}$ and $H_{\\delta}$ and the D4000 break. With the advent of large spectroscopic surveys at low redshift, such as the Sloan Digital Sky Survey (SDSS)\\citep{SDSSDR4}, it is possible to employ spectroscopic line indices to study the local galaxy population. However, equivalent data at intermediate redshift - certainly on the same scale - remains scant, although this is gradually changing with efforts such as the DEEP2 survey\\citep{Davis2003}.  Rest-frame UV photometry provides an attractive alternative. While its impact on the optical spectrum is relatively weak (and virtually undetectable, given typical observational and theoretical uncertainties), a small mass fraction ($<3$\\%) of young ($<1$ Gyr old) stars strongly affects the rest-frame UV shortward of $3000\\AA$. Furthermore, the UV remains largely unaffected by the age-metallicity degeneracy \\citep{Worthey1994} that typically plagues optical analyses \\citep{Kaviraj2006c}, making it an ideal \\emph{photometric} indicator of RSF. The sensitivity of the UV to young stars is demonstrated in Figure \\ref{fig:uv_sense}. We assume two instantaneous bursts of star formation, where the first burst is fixed at $z=3$ and the second burst is allowed to vary in age and mass fraction. The near-UV (NUV; $2300\\AA$) colour of the composite stellar population is plotted as a function of the age (symbol type) and mass fraction (x-axis) of the second burst. It is apparent that even a very small mass fraction ($\\sim$1\\%) of young stars ($\\sim0.1$ Gyrs old) causes a dramatic change in the $NUV-r$ colour compared to what might be expected from a purely old stellar population ($NUV-r\\sim 6.8$). Given that typical observational uncertainties in the $NUV-r$ colours from modern instrumentation are $\\sim0.2$ mag, the usefulness of the UV in detecting residual amounts of RSF becomes quite apparent.  \\begin{figure} \\begin{center} \\includegraphics[width=0.8\\textwidth]{uv_sense} \\caption{The sensitivity of the UV to young stars. We assume two instantaneous bursts of star formation, where the first burst is fixed at $z=3$ and the second burst is allowed to vary in age and mass fraction. The near-UV (NUV) colour of the composite stellar population is plotted as a function of the age (symbol type) and mass fraction (x-axis) of the second burst. It is apparent that even a small mass fraction ($\\sim$1\\%) of young stars ($\\sim0.1$ Gyrs old) causes a dramatic change in the $NUV-r$ colour compared to what might be expected from a purely old stellar population ($NUV-r\\sim 6.8$).} \\label{fig:uv_sense} \\end{center} \\end{figure}  The advent of the GALEX UV space telescope \\citep{Martin2005} and deep optical surveys (which can be used to trace the rest-frame UV spectrum at high redshift) has provided us, over the last few years, with an unprecedented opportunity to quantify the SFHs of early-type galaxies over the last 8 billion years by exploiting their (rest-frame) UV properties. This review describes the results of recent efforts that incorporate UV photometry to study the evolution of the early-type galaxy population across the redshift range $0<z<1$. While the emphasis is on studies that have quantified the buildup of stellar mass in these systems - without necessarily exploring the processes that drive this star formation - possible sources of RSF are explored, both in the context of monolithic collapse and in the standard $\\Lambda$CDM cosmology. ",
        "conclusions": "The advent of the GALEX space telescope and deep optical surveys have revolutionised our understanding of the properties of the early-type galaxy population in the rest-frame UV. The sensitivity of the UV to recent star formation provides an unprecedented opportunity to quantify the SFHs of early-type galaxies over the last 8 billion years and put strong constraints on models for their formation.  The first generation of studies that have harnessed the UV to directly constrain the RSF in the early-type population indicate that the early-type population, in general, forms stars over the lifetime of the Universe. The SFH of luminous ($M(V)<-21$) early-types is \\emph{quasi-monolithic}, with more than 80\\% of their stellar mass forming at $z>1$.  Fainter early-types experience more prolonged star formation, potentially forming a substantial fraction (30-60\\%) of their stellar mass at recent epochs ($z<1$).The derived values of RSF are consistent with the observed values of alpha-enhancements of luminous early-type galaxies in the local Universe.  The rest-frame UV CMR shows a typical scatter of several magnitudes over the entire redshift range $0<z<1$, indicating that low-level recent star formation is widespread in the early-type population. The intensity of recent star formation and the \\emph{bulk} of the UV colour distribution is consistent with what might be expected from simulations of minor mergers (mass ratios $\\lesssim$ 1:6) in a $\\Lambda$CDM cosmology.  Forthcoming studies with COSMOS\\citep{COSMOS} and DEEP2 \\citep{Davis2003} data will significantly consolidate our understanding of the UV properties of the high-redshift early-type population, alleviate the effect of cosmic variance and allow us to the split the observed sample of galaxies not only by luminosity (mass) but also by environment (using spectroscopic redshifts), which is a key driver of galaxy evolution\\citep{Kaviraj2006a}."
    },
    "0710/0710.3122_arXiv.txt": {
        "abstract": "Oscillations of magnetic structures in the solar corona have often been interpreted in terms of magnetohydrodynamic waves. We study the adiabatic magnetoacoustic modes of a prominence plasma slab with a uniform longitudinal magnetic field, surrounded by a prominence-corona transition region (PCTR) and a coronal medium. Considering linear small-amplitude oscillations, the dispersion relation for the magnetoacoustic slow and fast modes is deduced assuming evanescent-like perturbations in the coronal medium. In the system without PCTR, a classification of the oscillatory modes according to the polarisation of their eigenfunctions is made in order to distinguish modes with fast-like or slow-like properties. Internal and external slow modes are governed by the prominence and coronal properties respectively, and fast modes are mostly dominated by prominence conditions for the observed wavelengths. In addition, the inclusion of an isothermal PCTR does not substantially influence the mode frequencies, but new solutions (PCTR slow modes) are present. ",
        "introduction": "\\label{Introduction}  Despite recent observational and theoretical advances (see the reviews by \\opencite{martin}; \\opencite{patso}), the intimate physical nature of solar prominences is still enigmatic. It is clear that prominences can only exist at locations where the coronal magnetic field is ``peculiar'' enough to allow for the equilibrium of the dense prominence plasma against the force of gravity. Nevertheless, there remain many unknowns in the birth, evolution, and death of these objects.  Prominence seismology attempts to contribute to the research of prominence properties. Observations of waves and oscillations in prominences have become increasingly complex in the last decade and a rich field has emerged. Oscillatory periods ranging from less than a minute \\cite{balthasar} to 12 hours \\cite{foullon} have been detected. It has also been found that perturbations are typically attenuated with a damping time of a few periods \\cite{terradas02} and in some cases the phase speed and wavelength of propagating perturbations have also been characterised ({\\it e.g.} \\opencite{terradas02}). Finally, although there is a considerable gap between observations and theory, some attempts of doing prominence seismology have been undertaken and reasonable values of the prominence physical parameters have been obtained \\cite{pouget}. The observational background has been reviewed at length in various publications \\cite{oliverballester02,wiehr,engvold,banerjee} to which the reader is referred for more information.  The theoretical magnetohydrodynamic models developed so far start from very crude representations of the prominence and usually focus on a few ``ingredients'' whose relevance wants to be tested. For example, \\inlinecite{joarderroberts93} studied the adiabatic modes of oscillation of a plasma slab threaded by a skewed magnetic field. In \\inlinecite{oliverballester96} the magnetic geometry was simplified by taking a transverse magnetic field to the prominence slab, but a prominence-corona transition region (PCTR) was included; this is the only theoretical study of prominence oscillations in which the influence of the PCTR has been taken into account. The temporal and spatial  damping of perturbations has been investigated in some other works ({\\it e.g.} \\opencite{terradas05}). \\inlinecite{oliverballester02} and \\inlinecite{banerjee} kave reviewed this subject.  The purpose of this work is to address the importance of the PCTR on the oscillations of a prominence slab threaded by a straight, uniform magnetic field parallel to the slab axis (and not perpendicular to it as in \\opencite{oliverballester96}). The PCTR is a narrow layer through which the physical parameters (temperature, density, ionisation degree, {\\it etc.}) change abruptly from prominence to coronal values \\cite{cirigliano}. The availability of telescopes for the observation of spectral lines sensitive to PCTR temperatures has allowed the detection of waves in this region \\cite{blanco, bocchialini, pouget}. The first of these works reports the propagation of a disturbance travelling at $\\approx 170$~km~s$^{-1}$ and with a period in the band 3.2\\,--\\,6.4~minutes. \\inlinecite{pouget} reported periodicities detected in three solar filaments in the He\\,{\\sc i} line at 5843.3~\\AA~(corresponding to 20\\,000~K) and identified some of the periods with the six fundamental oscillatory modes according to the model proposed by \\inlinecite{joarderroberts93}. Incorporating the PCTR in our model will allow us to test whether it has an observable effect on the oscillations of a prominence.  When the PCTR is removed, our equilibrium configuration is identical to that of \\inlinecite{ER82}, hereafter ER82. ER82 studied a wide range of values for the physical conditions of the slab and the outer medium, plotting dispersion diagrams for each situation, and also discussed the dispersion relation in the long- and short-wavelength limits. \\inlinecite{JR92}, hereafter JR92, extended the work of ER82 by considering the possibility of motions and propagation in the $y$-direction and made a more refined classification of the oscillation modes according to their phase speed. Our purpose is to revise the work done by JR92 in the case of no motions and no propagation in the $y$-direction (i.e. the case analysed by ER82), taking into account realistic physical conditions similar to those in a prominence and its surrounding PCTR and corona, and to study the polarisation of motions and the other perturbations. Additionally, we explore the parameter space in an attempt to understand how the variation of the temperature, density and magnetic field affects the oscillatory period.  This paper is organised as follows: Section~2 contains a description of the equilibrium model and the linear, adiabatic wave equations, besides the derivation of the dispersion relation of magnetoacoustic modes. The oscillatory modes are studied first for a system without PCTR (as a revision of ER82) in Section~3. Next, an isothermal PCTR is included in the equilibrium and its effects are investigated in Section~4. Finally, our conclusions are given in Section~5. ",
        "conclusions": "The investigation performed in the present paper is a simple study about the influence of the presence of a prominence-corona transition region on the magnetoacoustic oscillatory modes supported by the prominence body. The reader must be aware that the wave modes investigated here are not all the possible solutions supported by our equilibrium due to the restriction $k_y = 0$. For this reason, the case with $k_y \\neq 0$ should be investigated in a future work.  This study is divided into two parts. In the first one (Section~\\ref{res:nopctr}), the JR92 prominence model for $k_y = 0$ (ER82 case) has been revised and the oscillatory modes have been reclassified according to their magnetoacoustic properties. This has allowed us to find a wave mode with slow-like properties (the external slow mode), which passed unnoticed in ER82 and JR92. In the second part of the paper (Section~\\ref{res:pctr}), an isothermal PCTR has been considered in the system, and its effect on the previously existing wave modes has been assessed. In addition, a new solution has been obtained (the PCTR slow mode), which possesses slow-like properties and produces large longitudinal motions inside the transition region.  The main conclusions of this paper are summarised next. \\begin{itemize}  \\item Internal slow modes are dominated by prominence conditions whereas the behaviour of the external slow mode is controlled by coronal physical conditions. The latter only exists as an evanescent-like solution for small wavenumber.  \\item Fast modes are dominated by the prominence but show a slight dependence on coronal properties. The influence of the corona on fast modes grows for small $k_x$ (large wavelength) since perturbations are poorly confined.  \\item The presence of an isothermal PCTR does not greatly affect the frequency of the oscillatory modes of the equilibrium without transition region if the PCTR sound speed is less than the internal Alfv\\'en speed. In this situation, a new type of slow-like oscillatory modes (PCTR slow modes) which produce large displacements of the plasma inside the PCTR appear in an isolated window in the phase speed diagram. These new solutions do not interact with the previously existing solutions of the system without PCTR.  \\item When an isothermal transition region with a tube speed greater than the internal Alfv\\'en speed is considered, the PCTR slow modes interact with the previously existing fast modes by means of couplings.  \\end{itemize}  The large longitudinal displacements of the plasma inside the transition region produced by the PCTR slow modes could be detected as oscillations in the intensity or Doppler shift of spectral lines associated with typical PCTR temperatures ({\\it e.g.} \\opencite{pouget}). As mentioned in the Section~\\ref{Introduction}, there are very few observations of oscillations linked to prominences in transition region spectral lines, so one must be cautious about the observational values and their theoretical interpretation. Despite these remarks, the addition of the PCTR in theoretical investigations of prominence oscillations is important because it gives a way to understand observations obtained with wavelengths corresponding to different plasma temperatures. The consideration of an isothermal, homogeneous PCTR is a very rough approximation to the real situation but this simple theoretical study has allowed us to obtain some interesting results. From our point of view, the most important of all these results is that global fast modes of the system can interact with PCTR slow modes, which are essentially confined in the narrow PCTR. This fact could become of especial relevance if a more realistic, smooth PCTR is assumed. Then, a slow continuum would be present instead of individual PCTR slow modes. The interaction of the fast mode with this slow continuum could have important repercussions on the wave behaviour since a resonance phenomenon could occur. The consideration of a smooth PCTR is therefore the next step to be undertaken in future investigations, the present isothermal case being a simple first approach to the problem."
    },
    "0710/0710.5532_arXiv.txt": {
        "abstract": "The aim of this paper is to investigate ways to optimize the accuracy of photometric redshifts for a SNAP like mission. We focus on how the accuracy of the photometric redshifts depends on the magnitude limit and signal-to-noise ratio, wave-length coverage, number of filters and their shapes and observed galaxy type. We use simulated galaxy catalogs constructed to reproduce observed galaxy luminosity functions from GOODS, and derive photometric redshifts using a template fitting method. By using a catalog that resembles real data, we can estimate the expected number density of galaxies for which photometric redshifts can be derived. We find that the accuracy of photometric redshifts is strongly dependent on the signal-to-noise (S/N) (i.e., S/N$>$10 is needed for accurate photometric redshifts). The accuracy of the photometric redshifts is also dependent on galaxy type, with smaller scatter for earlier type galaxies. Comparing results using different filter sets, we find that including the U-band is important for decreasing the fraction of outliers, i.e., ``catastrophic failures''. Using broad overlapping filters with resolution $\\sim4$~gives better photometric redshifts compared to narrower filters (resolution $\\gsim5$) with the same integration time. We find that filters with square response curves result in a slightly higher scatter, mainly due to a higher fraction of outliers at faint magnitudes. We also compare a 9-filter set to a 17-filter set, where we assume that the available exposure time per filter in the latter set is half that of the first set. We find that the 9-filter set gives more accurate redshifts for a larger number of objects and reaches higher redshift, while the 17-filter set is gives better results at bright magnitudes. ",
        "introduction": "In recent years, there have been unprecedented progress in observational astronomy due, in large part, to the advent of large format and highly sensitive optical/infrared detectors. Installation of these cameras on 8m ground-based telescopes and space-borne facilities have enabled planning large and deep galaxy surveys, increasing the discovery space by over an order of magnitude. In particular, wide-area multi-waveband imaging from space, complemented by follow up ground-based observations, have provided extremely valuable datasets for studying diverse topics in observational astronomy and cosmology. For example, installation of the Advance Camera for Survey (ACS) on the Hubble Space Telescope (HST) has resulted in multi-waveband surveys of galaxies, including the Great Observatories Origins Deep Survey (GOODS - Giavalisco et al. 2004), COSMOS (Scoville et al. 2007) and the Hubble Ultra-Deep Field (HUDF- Beckwith et al. 2006). These surveys provide deep multi-waveband data covering large areas, used to study a number of issues concerning formation and evolution of galaxies, including: study of rest-frame properties of different populations of galaxies (e.g., Bundy et al. 2005; Grogin et al. 2005; Dahlen et al. 2007), search for the highest redshift (Kneib et al. 2004; Bouwens et al. 2005) and new population of galaxies (Wiklind et al. 2007), mapping of the dark matter distribution in strong lensing clusters (Smith et al. 2005; Broadhurst et al. 2005; Limousin et al. 2007; Richard et al. 2007), cosmological constraints from weak lensing (Massey et al. 2007a) and clustering of galaxies (McCracken et al. 2007) and the 3D large scale structure dark matter distribution (Massey et al. 2007b).  Among the most important outcomes from these studies was the first ever space-borne search for Supernovae Type Ia in the GOODS fields (Riess et al. 2004). These observations have darker sky background, leading to deeper images, and significantly narrower PSFs, leading to better spatial resolution and hence, identification of more distant supernovae in galaxies, compared to ground-based images. This allowed discovery of 21 high redshift SNe Ia at $z > 1$, which includes almost all (but one) of the highest redshift SNe Ia known at the time (Strolger et al. 2004). Combining these high-$z$~and nearby SNe Ia, the Hubble diagram was established, allowing significant constraints on properties of Dark Energy and its equation of state (Riess et al. 2004, 2007). An essential component of this study was measurement of photometric redshifts of the hosts of SN candidates to identify objects of the highest interest for follow up with subsequent spectroscopic observations. This was possible due to availability of multi-waveband data from space- and ground-based observations.  The SuperNova Acceleration Probe (SNAP\\footnote{http://snap.lbl.gov/}) mission is aimed at finding thousands of SNe of various types to redshift $z\\sim$1.7, allowing detailed study of reliability of SNe Ia as standard candles (i.e. dependence of their observed properties on the type of their host galaxy and its redshift, effect of dust, their frequency of appearance in early-type galaxies). This proposed mission will improve the use of of SNe Ia as standard candles and provide significant constraints on properties of Dark Energy and its nature. Moreover, it will provide a large sample of SNe Type II, which are used as diagnostics for star formation activity in galaxies, allowing a statistically large sample of these objects to examine evolution of star formation and metallicity with redshift (Dahlen et al. 2004). As part of the SNAP's primary mission, a wide (at least 1000 square degree per year), multi-band survey will be conducted.  This survey will be optimized for the detection of weak gravitational lensing, a powerful probe of dark energy. Weak lensing provides a direct way for measuring the distribution of dark matter in the Universe. The evolution of dark matter structures over cosmic time is governed by the nature of the dark energy. Thus, accurate photometric redshifts, which are required to measure the 3-dimensional distribution of dark matter, are necessary to exploit weak lensing as a probe of dark energy. Furthermore, these multi-waveband deep data will be extremely useful in studying formation and evolution of galaxies, groups and clusters as a function of their redshift, morphology, environment, color and star formation properties. It will also be a unique mission to probe the first quasars and the first luminous galaxies in the Universe, thus probing the epoch of cosmic reionization.  Future Dark Energy probes based on SN Ia and weak lensing ideally require accurate redshifts for individual galaxies. However, given the size of the planned surveys and their depth, it is not practical (nor feasible) to measure spectroscopic redshifts for all the observed galaxies. Therefore, a critical evaluation of the photometric redshift capabilities of any of the JDEM experiments is key to optimize the design of the respective mission. In the case of SNAP, this requires an optimization of the number of filters used, their spectral resolution and shapes and thus throughputs as well as the overall wavelength coverage, to allow most accurate measurement of photometric redshifts. To investigate these problems, we use observational data from the GOODS fields to create a mock galaxy catalog containing a large set of objects with known properties, i.e., redshift, spectral type, luminosity and amount of internal extinction. We then run photometric redshift codes on the multi-waveband data for galaxies in the mock catalog. The aim of this investigation is to study how the accuracy of galaxy photometric redshifts depends on a range of factors including redshift, signal-to-noise, magnitude limits, galaxy spectral type and, in particular, the number of filters, filter shapes and band-widths. For an investigation on how to optimize filters for Type Ia cosmology investigations, see Davis et al. (2006). In a following paper (Jouvel et al. 2007, in preparation), we will investigate in further detail the impact of calibration, SED evolution and size of the spectroscopic surveys on the photomtric redshift determination.  Throughout this paper we use $\\Omega_{\\Lambda}=0.7$, $\\Omega_M=0.3$~and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. Magnitudes are in the AB system. ",
        "conclusions": "We have used simulations to investigate how the behavior of photometric redshifts for a SNAP-like mission depends on e.g., magnitude limits, filters choices, wavelength coverage and galaxy types. We have also discussed methods for decreasing the expected fraction of outliers, i.e., galaxies with significant disagreement between the photometric and spectroscopic redshifts. We note that our investigation is not primarily focussed on getting exact numbers for e.g., the photometric redshift accuracy and outlier fractions, but to investigate how these diagnostics are affected when changing e.g., S/N and filter sets. Our main conclusions are: \\begin{itemize} \\item{We find that including the U-band significantly decreases the fraction of outliers and results in an increase photometric redshift accuracy.} \\item{A 17-filter set results in larger scatter compared to a 9-filter set except at bright magnitudes. However, after excluding outliers, the 17-filter set gives more accurate redshifts at all magnitudes.} \\item{A 9-filter set with narrower filter functions (resolution $\\sim$5 instead of $\\sim$4), results in an increase in the scatter of the photometric redshifts. This is caused by the larger photometric errors when fewer photons are detected.} \\item{Using a filters with a square transmission curves  decreases the scatter in the photometric redshifts for the majority of the objects. At the same time, however, the fraction of outliers is doubled. Therefore, if possible outliers could be efficiently flagged, the square filter set would be preferred.} \\item{The accuracy of the photometric redshifts depends on both magnitude (or efficiently the S/N) and galaxy spectral type, with better results at high S/N and for earlier type galaxies.} \\item{Using the D95 method can significantly decrease the number of outliers, while only decreasing the total number of objects moderately. Using this method is therefore preferred compared to using only the S/N as a cut to decrease scatter.} \\end{itemize}"
    },
    "0710/0710.0317_arXiv.txt": {
        "abstract": "{ Validation is often defined as the process of determining the degree to which a model is an accurate representation of the real world from the perspective of its intended uses. Validation is crucial as industries and governments depend increasingly on predictions by computer models to justify their decisions. In this article, we survey the model validation literature and propose to formulate validation as an iterative construction process that mimics the process occurring implicitly in the minds of scientists. We thus offer a formal representation of the progressive build-up of trust in the model, and thereby replace incapacitating claims on the impossibility of validating a given model by an adaptive process of constructive approximation. This approach is better adapted to the fuzzy, coarse-grained nature of validation. Our procedure factors in the degree of redundancy versus novelty of the experiments used for validation as well as the degree to which the model predicts the observations. We illustrate the new methodology first with the maturation of Quantum Mechanics as the arguably best established physics theory and then with several concrete examples drawn from some of our primary scientific interests: a cellular automaton model for earthquakes,  an anomalous diffusion model for solar radiation transport in the cloudy atmosphere, and a computational fluid dynamics code for the Richtmyer--Meshkov instability.}   \\newpage ",
        "introduction": "\\subsection{Introductory Remarks and Outline}  At the heart of the scientific endeavor, model building involves a slow and arduous selection process, which can be roughly represented as proceeding according to the following steps: \\begin{enumerate} \\item start from observations and/or experiments; \\item classify them according to regularities that they may exhibit: the presence of patterns, of some order, also sometimes referred to as structures or symmetries, is begging for ``explanations'' and is thus the nucleation point of modeling; \\item use inductive reasoning, intuition, analogies, and so on, to build hypotheses from which a model \\footnote{ By model, we understand an abstract conceptual construction based on axioms and logical relations developed to extract logical propositions and predictions. } is constructed; \\item test the model obtained in step 3 with available observations, and then extract predictions that are tested against new observations or by developing dedicated experiments. \\end{enumerate} The model is then rejected or refined by an iterative process, a loop going from step~1 to step~4. A given model is progressively validated by the accumulated confirmations of its predictions by repeated experimental and/or observational tests.  Building and using a model requires a language, i.e., a vocabulary and syntax, to express it. The language can be English or French for instance to obtain predicates specifying the properties of and/or relation with the subject(s). It can be mathematics, which is arguably the best language to formalize the relation between quantities, structures, space and change. It can be a computer language to implement a set of relations and instructions logically linked in a computer code to obtain quantitative outputs in the form of strings of numbers. In this later version, our primary interest here, validation must be distinguished from verification. Whereas {\\it verification}\\/ deals with whether the simulation code correctly solves the model equations, {\\it validation}\\/ carries an additional degree of trust in the value of the model vis-\\`a-vis experiment and, therefore, may convince one to use its predictions to explore beyond known territories \\cite{Roa98a}.  The validation of models is becoming a major issue as humans are increasingly faced with decisions involving complex tradeoffs in problems with large uncertainties, as for instance in attempts to control the growing anthropogenic burden on the planet within a risk-cost framework \\cite{Cos97,Pim97} based on predictions of models. For policy decisions, national, regional, and local governments increasingly depend on computer models that are scrutinized by scientific agencies to attest to their legitimacy and reliability. Cognizance of this trend and its scientific implications is not lost on the engineering \\cite{Bab04} and physics \\cite{Pos05} communities.  Our purpose here is to clarify {\\it from a physics-based perspective}\\/ what validation is and to propose a roadmap for the development of systematic approach to physics-based validation with broad applications.  We will focus primarily on the needs of computational fluid dynamics and particle/radiation transport codes.  In the remainder of this section, we first review different definitions and approaches found in the literature, positioning ourselves with respect to selected topics or practices pertaining to validation; we then show how the validation problem is related to the mathematical statistics of hypothesis testing and discuss some problems associated with emergent behaviors in complex systems. In section~2, we list and describe qualitatively the elements required in our vision of model validation as an iterative process where one strives to build trust in the model going from one experiment to the next; however, one must also be prepared to uncover in the model a flaw, which may or may not be fatal. We offer in sections~3--4 our quantitative physics-based approach to model validation, where the relevance of the experiment to the validation process is represented explicitly. (An appendix explores the model validation problem more formally and in a broader context.) Section~5 demonstrates the general strategy for model validation using the historical development of quantum physics---a remarkably clear ideal case. Section~6 uses some research interests of the present authors to further illustrate the validation procedure using less-than-perfect models in geophysics,  computational fluid dynamics (CFD), and radiative transfer.  We summarize in section~7.   \\subsection{Standardized Definitions}  The following definitions are given by the American Institute of Aeronautics and Astronautics \\cite{AIA98}: \\begin{itemize} \\item {\\it Model}: A representation of a physical system or process intended to enhance our ability to predict, control and eventually to understand its behavior. \\item {\\it Calibration}: The process of adjusting numerical or physical modeling parameters in the computational model for the purpose of improving agreement with experimental data. \\item {\\it Verification}: The process of determining that a model implementation accurately represents the developer's conceptual description of the model and the solution of the model. \\item {\\it Validation}: The process of determining the degree to which a model is an accurate representation of the real world from the perspective of the intended uses of the model. \\end{itemize} Figure \\ref{f:VV}, sometimes called a Sargent diagram, shows where validation and several other of the above constructs and stages enter into a complete modeling project.   \\begin{figure} \\begin{center} \\includegraphics[width=6cm]{Fig_1_VV.eps} \\caption{Schematic representation of the conventional position of validation in model construction according to Schlesinger \\cite{Sch79} and Sargent  \\cite{Sar98,Sar01}.} \\label{f:VV} \\end{center} \\end{figure}   In the concise phasing of Roache \\cite{Roa98a}, {\\it ``Verification consists in solving the equations right while validation is solving the right equations.''} In the context of the validation of astrophysical simulation codes, Calder et al. \\cite{Cal02} add: {\\it ``Verification and validation are fundamental steps in developing any new technology. For simulation technology, the goal of these testing steps is assessing the credibility of modeling and simulation.''}  {\\it Verifications}\\/ of complex CFD codes usually comprise a suite of standard test problems in the field of fluid dynamics \\cite{Cal02}. These include Sod's test \\cite{Sod78}, the strong shock tube problem \\cite{Rid99}, the Sedov explosion problem \\cite{Sed59}, the interacting blast wave problem \\cite{Woo84}, a shock forced through a jump in mesh refinement, and so on.  {\\it Validations}\\/ of complex CFD codes is usually done by comparison with experiments testing a variety of physical phenomena, including instabilities, turbulent mixing, shocks, etc. Validation requires that the numerical simulations recover the salient qualitative features of the experiments, such as the instabilities, their nonlinear development, the determination of the most unstable modes, and so on. See, for instance, Gnoffo et al. \\cite{Gno98}.  Considerable work on verification and validation of simulations has been done in the field of CFD, and in this literature the terms verification and validation have precise, technical meanings \\cite{AIA98,Roa98a,Roa98b,Sar98,Sar01}. Verification is taken to mean demonstrating that a code or simulation accurately represents the conceptual model. Roache~\\cite{Roa02} stresses the importance of distinguishing between (i) verification of codes and (ii) verification of calculations. The former is concerned with the correctness of the code. The later deals with the correctness of the physical equations used in the code. The programming and methods of solution can be correct (verification (i) successful) but they can solve erroneous equations (verification (ii) failure). Validation of a simulation means demonstrating that the simulation appropriately describes Nature. The scope of validation is therefore much larger than that of verification and includes comparison of numerical results with experimental or observational data. In astrophysics, where it is difficult to obtain observations suitable for comparison to numerical simulations, this process can present unique challenges. Roache [op. cit.] goes on to offer the optimistic prognosis that {\\it ``the problems of Verification of Codes and Verification of Calculations are essentially solved for the case of structured grids, and for structured refinement of unstructured grids. It would appear that one higher level of algorithm/code development is required in order to claim a complete methodology for Verification of Codes and Calculations. I expect this to happen. Within 10 years, and likely much less, Verification of Codes and Calculations ought to be settled questions. I expect that Validation questions will always be with us.''} We fully endorse this last sentence, as we will argue further on that validation is akin to the development of ``trust'' in theories of real phenomena, a never-ending quest.   \\subsection{Impossibility Statements}  For these reasons, the possibility of validating numerical models of natural phenomena, often endorsed either implicitly or identified as reachable goals by natural scientists in their daily work, has been challenged; quoting from Oreskes et al. \\cite{Ore94a}: {\\it ``Verification and validation of numerical models of natural systems is impossible. This is because natural systems are never closed and because model results are always non-unique.''} According to this view, the impossibility of ``verifying'' or ``validating'' models is not limited to computer models and codes but to all theories that rely necessarily on imperfectly measured data and auxiliary hypotheses.  As Sterman \\cite{Ste94} puts it: {\\it ``Any theory is underdetermined and thus unverifiable, whether it is embodied in a large-scale computer model or consists of the simplest equations.''} Accordingly, many uncertainties undermine the predictive reliability of any model of a complex natural system in advance of its actual use. \\footnote{ For further debate and commentary by Oreskes and her co-authors, see refs. \\cite{Ryk94,Ore94b,Ore98}; also noteworthy is the earlier paper by Konikov and Bredehoeft \\cite{Kon92} for a statement about validation impossibility in the context of groundwater models. }  Such ``impossibility'' statements are reminiscent of other ``impossibility theorems.'' Consider the mathematics of algorithmic complexity \\cite{Cha87}, which provides one approach to the study of complex systems. Following reasoning related to that underpinning G\\\"odel's incompleteness theorem, most complex systems have been proved to be computationally irreducible, i.e., the only way to predict their evolution is to actually let them evolve in time. Accordingly, the future time evolution of most complex systems appears inherently unpredictable. Such sweeping statements turn out to have basically no practical value. This is because, in physics and other related sciences, one aims at predicting coarse-grained properties. Only by ignoring most of molecular detail, for example, did researchers ever develop the laws of thermodynamics, fluid dynamics and chemistry. Physics works and is not hampered by computational irreducibility because we only ask for approximate answers at some coarse-grained level \\cite{Buc05}. By developing exact but {\\it coarse-grained} procedures on computationally irreducible cellular automata, Israeli and Goldenfeld \\cite{Isr04} have demonstrated that prediction may simply depend on finding the right level for describing the system. More generally, we argue that only coarse-grained scales are of interest in practice but their description requires ``effective'' laws which are in general based on finer scales. In other words, real understanding must be rooted in the ability to predict coarser scales from finer scales, i.e., a real understanding solves the universal micro-macro challenge. Similarly, we propose that validation is possible, to some degree, as explained further on.   \\subsection{Validation and the Mathematical Statistics of Hypothesis Testing}  Calder et al. \\cite{Cal02} also write: {\\it ``We note that verification and validation are necessary but not sufficient tests for determining whether a code is working properly or a modeling effort is successful. These tests can only determine for certain that a code is not working properly.''} This last statement is important because it points to a bridge between the problem of validation and some of the most central questions of mathematical statistics \\cite{Bor98}, namely, hypothesis testing and statistical significance tests. This connection has been made previously by several others authors \\cite{Col97,Hil99,Hil00,Eas01}. In showing the usefulness of the concepts and framework of hypothesis testing, we depart from Oberkampf and Trucano \\cite{Obe02a} who mistakenly state that hypothesis testing is a true or false issue, only. Every test of significance begins with a ``null'' hypothesis H$_0$, which represents a theory that has been put forward, either because it is believed to be true or because it is to be used, but has not been proved. \\footnote{ We refer the reader to V.J. Easton and J.H. McColl, {\\it Statistics Glossary}, http://www.cas.lancs.ac.uk/glossary\\_v1.1/main.html, from which we have borrowed liberally for this brief summary. }  For example, in a clinical trial of a new drug, the null hypothesis might be: ``the new drug is no better, on average, than the current drug.'' We would write H$_0$: ``there is no difference between the two drugs on average.'' The alternative hypothesis H$_1$ is a statement of what a statistical hypothesis test is set up to establish. In the example of a clinical trial of a new drug, the alternative hypothesis might be that the new drug has a different effect, on average, to be compared to that of the current drug. We would write H$_1$: the two drugs have different effects, on average. The alternative hypothesis might also be that the new drug is better, on average, than the current drug. Once the test has been carried out, the final conclusion is always given in terms of the null hypothesis. We either ``reject H$_0$ in favor of H$_1$'' or ``do not reject H$_0$.'' We never conclude ``reject H$_1$,'' or even ``accept H$_1$.'' If we conclude ``do not reject H$_0$,'' this does not necessarily mean that the null hypothesis is true, it only suggests that there is not sufficient evidence against H$_0$ in favor of H$_1$; rejecting the null hypothesis then suggests that the alternative hypothesis may be true, or is at least better supported by the data. Thus, one can never prove that an hypothesis is true, only that it is wrong by comparing it with another hypothesis. One can also conclude that ``hypothesis H$_1$ is not necessary and another, more parsimonious, one $H_0$ should be favored.'' The alternative hypothesis H$_1$ is not rejected, strictly speaking, but is found unnecessary or redundant with respect to H$_0$. This is the situation when there are two (or several) alternative hypotheses H$_0$ and H$_1$, which can be composite, nested, or non-nested. \\footnote{ The technical difficulties of hypothesis testing depend on these nested structures of the competing hypotheses; see, for instance, Gourieroux and Monfort \\cite{Gor94}. }  Within this framework, the above-mentioned statement by Oreskes et al. \\cite{Ore94a} that verification and validation of numerical models of natural systems is impossible is hardly news: the theory of statistical hypothesis testing has taught mathematical and applied statisticians for decades that one can never prove an hypothesis or a model to be true. One can only develop an increasing trust in it by subjecting it to more and more tests that ``do not reject it.'' We attempt to formalize below how such trust can be increased to lead to an asymptotic validation.   \\subsection{Code Comparison}  The above definitions are useful in recasting the role of code comparison in verification and validation (Code Comparison Principle or CCP). Trucano et al. \\cite{Tru03} are unequivocal on this practice: {\\it ``the use of code comparisons for validation is improper and dangerous.''} We propose to interpret the meaning of CCP for code verification activities (which has been proposed in this literature) as parallel to the problem of hypothesis testing: Can one reject Code \\#1 in favor of Code \\#2? In this spirit, the CCP is nothing but a reformulation in the present context of the fundamental principle of hypothesis testing. Viewed in this way, it is clear why CCP is not sufficient for validation since validation requires comparison with experiments and several other steps described below. The analogy with hypothesis testing illuminates what CCP actually is: CCP allows the selection of one code among several codes (at least two) but does not help one to draw conclusions about the validity of a given code or model when considered as a unique entity independent of other codes or models. ~\\footnote{ We should stress that the Sandia Report \\cite{Tru03} by Trucano et al. presents an even more negative view of code comparisons because it addresses the common practice in the simulation community that turns to code comparisons rather than bone fide verification or validation, without any independent referents. } Thus, the fundamental problem of validation is more closely associated with the other class of problems addressed by the theory of hypothesis testing, which consists in the so-called ``tests of significance'' where one considers only a single hypothesis H$_0$, and the alternative is ``all the rest,'' i.e., all hypotheses that differ from H$_0$. In that case, the conclusion of a test can be the following: ``this data sample does not contradict the hypothesis H$_0$,'' which is not the same as ``the hypothesis H$_0$ is true.'' In other words, an hypothesis cannot be excluded because it is found sufficient at some confidence level for explaining the available data. This is not to say that the hypothesis is true. It is just that the available data is unable to reject said hypothesis. Restating the same thing in a positive way, the result of a test of significance is that the hypothesis H$_0$ is ``compatible with the available data.''  It is implicit in the above discussion that, to compare codes quantitatively in a meaningful way, they must solve the same set of equations using different algorithms, and not just model the same physical system.  Indeed, there is nothing wrong with ``validating'' a numerical implementation of a knowingly approximate approach to a given physical problem.  For instance, a (duly verified) diffusion/P$_1$ transport code can be validated against a detailed Monte Carlo or S$_n$ code. The more detailed model must in principle be validated against real-world data. In turn, it provides validation ``data'' to the coarser model. Naturally, the coarser (say, P$_1$ transport) model still needs to establish its relevance to the real world problem of interest, preferably by comparison with real observations, or at least be invoked only in regimes where it is known a priori to be sufficiently accurate based on comparison with a finer (say, Monte Carlo transport) model.  Two noteworthy initiatives in transport model comparison for non-nuclear applications are the Intercomparison of 3D Radation Codes (I3RC) \\cite{Cah05} (i3rc.gsfc.nasa.gov) and the RAdiation Model Intercomparison (RAMI) \\cite{Pin01,Pin04} (rami-benchmark.jrc.it).  The former is focused on the challenge of 3D radiative transfer in the cloudy atmosphere while the later is about 3D radiative transfer inside plant canopies; both efforts are motivated by issues in remote sensing (especially from space) and radiative energy budget estimation (either in the framework of climate modeling or using observational diagnostics, which typically means more remote sensing). \\footnote{ In remote sensing science, transport theory (for photons) plays a central role and ``validation'' has a special meaning, namely, the estimation of uncertainty for remote sensing products based on ``ground-truth,'' i.e., field measurements of the very same geophysical variables (e.g., surface temperature or reflectivity, vegetation productivity, soil moisture) that the satellite instrument is designed to quantify. These data are collected at the same location as the imagery, if possible, at the precision of a single pixel.  This type of validation exercise will test both the ``forward'' radiation transport theory and its ``inversion.'' Atmospheric remote sensing, particularly of clouds, poses a special challenge because, strictly-speaking, there is no counterpart of ground-truthing. One must therefore often make do with comparisons of ground-based and space-based remote-sensing (say, of the column-integrated aerosol burden) to quantify uncertainty in both operations.  In-situ measurements (temperature, humidity, cloud liquid water, etc.) from airborne platforms---balloon or aircraft---are always welcome but collocation is rarely close enough for point-to-point comparisons; statistical agreement is then all that is to be expected, and residuals provide the required uncertainty. } Much has been learned by the modelers participating in these code comparison studies, and the models have been improved on average \\cite{Wid07}. Although not connected so far to the engineering community that is at the forefront of V\\&V standardization and methodology, the I3RC and RAMI communities talk much about ``testing,'' and sometimes ``certification,'' and not so much about ``verification'' (which would be appropriate) or ``validation'' (which would not).  What about multi-physics codes such as those used routinely in astrophysics, nuclear engineering, or climate modeling?  CCP, along with the stern warnings of Trucano et al. \\cite{Tru03}, applies here, too.  Even assuming that all the model components are properly verified or even individually validated, the aggregated model is likely to be too complex to talk about clean verification through output comparison. Finding some level of agreement between two or more complex multi-physics models will naturally build confidence in the whole (community-wide) modeling enterprise. However, this is not to be interpreted as validation of any or all of the individual models.  There are many reasons for wanting to have not just one model on hand but a suite of more or less elaborate ones.  A typical collection can range from the mathematically and physically exact but numerically intractable to the analytically solvable, possibly even on the proverbial back-of-an-envelope. We elaborate on and illustrate this kind of hierarchical modeling effort in section \\ref{s:4pbs} of the Appendix, offering it as an approach where model development is basically simultaneous with its validation.   \\subsection{Relations Between Validation, Calibration and Data Assimilation} \\label{s:cal+DA}  As previously stated, validation can be characterized as the act of quantifying the credibility of a model to represent phenomena of interest. Virtually all such models contain numerical parameters, the precise values of which are not known a priori and, therefore, must be assigned.  Calibration is the process of adjusting those parameters to optimize (in some sense) the agreement between the model results and a specific set of experimental data.  Such data necessarily have uncertainties associated with them, e.g., due to natural variability in physical phenomena as well as to unavoidable imprecision of diagnostics. Likewise, there are intrinsic errors associated with the numerical methods used to evaluate many models, e.g., in the approximate solutions obtained from discretization schemes applied to partial differential equations.  The approach of defensibly prescribing parameters for complex physical phenomena while incorporating the inescapable variability in these values is called ``calibration under uncertainty,'' \\cite{Tru06} a field that poses non-trivial challenges in its own right.  However calibration is approached, it must be undertaken using a set of data---ideally from specifically chosen calibration experiments/observations \\cite{Obe02b}---that differs from the physical configurations of ultimate interest (i.e., against which the model will be validated).  In order to ensure that validation remains independent of calibration, it is imperative that these data sets be disjoint.  In the case of large, complex, and costly experiments encountered in many real-world applications, it can be difficult to maintain a scientific ``demilitarized zone'' between calibration and validation. To not do so, however, risks undermining the scientific integrity of the associated modeling enterprise, the potential predictive power of which may rapidly wither as the validation study devolves into a thinly disguised exercise in calibration.  For complex systems, there are many choices to be made regarding experimental and numerical studies in both validation and calibration. The high-level approach of the Phenomena Identification and Ranking Table (PIRT) \\cite{Tru02} can be used to heuristically characterize the nature of one's interest in complicated systems. This approach uses expert knowledge to identify the phenomenological components in a system of interest, to rank their (relative) perceived importance in the overall system, and to gauge the (relative) degree to which these component phenomena are perceived to be understood.  This rough-and-ready approach can be used to target the choice of validation experiments for the greatest scientific payoff on fixed experimental and simulation budgets.  To help guide calibration activities, one can apply the quantitative techniques of sensitivity analysis to rank the relative impact of input parameters on model outcome.  Such considerations are particularly important for complex models containing many adjustable parameters, for which it may prove impossible to faithfully calibrate all input parameters.  Saltelli et al. \\cite{Sal00, Sal04} have championed ``sensitivity analysis'' methods, which come in two basic flavors and many variations. One class of methods uses exact or numerical evaluation of partial derivatives of model output deemed important with respect to input parameters to seek regions of parameter space that might need closer examination from the standpoints of calibration and/or validation.  If the model has time dependence, one can follow the evolution of how parameter choices influence the outcome.  The alternate methodology uses adjoint dynamical equations to determine the relative importance of various parameters. The publications of Saltelli et al. provide numerous examples illustrating the value and practical impact of sensitivity analysis, as well as references to the wide scientific literature on this subject.  The results of numerical studies guided by sensitivity analysis can be used both to focus experimental resources on high-impact experimental studies and to steer future model development efforts.  In dynamical modeling, initial conditions can be viewed as parameters and, as such, they need to be determined optimally from data. If the dynamical system in question is evolving continuously over time and data become available along the trajectory of the dynamical system, the problem of finding {\\it a single} initial condition over the entire trajectory becomes increasingly and exceedingly difficult as the time window of the trajectory extends. In fact, it is practically impossible for the systems like the atmosphere or ocean whose dynamics is highly nonlinear, high-dimensional model is undoubtedly imperfect, and inhomogeneous and sporadic data are subject to (poorly understood) errors.  Data assimilation is an approach that attends to this problem by breaking up the trajectory  over (fixed-length) time windows and solving the initialization problem sequentially over one time window at a time as data become available. A novelty of data assimilation is that, rather than solving the initialization problem from scratch, it uses the model forecast as the first guess (the prior) of the initialization (optimization) problem. Once the optimization is completed, the optimal solution (the posterior) becomes the initial condition for the next model forecast.  This iterative Bayesian approach to data assimilation is most effective when the uncertainties in both the prior and the data are accurately quantified, as the system evolves over time and the data assimilation iterates one cycle after another. This is a non-trivial problem, because it requires the estimate of not only the model state but also the uncertainties associated with it, as well as the proper description of the uncertainties in data.  Numerical weather prediction (NWP) is one of the most familiar application areas of data assimilation---one with major societal impact. The considerable progress in skill of the NWP in recent decades has been due to improvements in all aspects of data assimilation  \\cite{Kal03}, i.e., modeling of the atmosphere, quality and quantity of data, and data assimilation methods. At the time of writing, most operational NWP centers use the so-called the ``three-dimensional variational method'' (3D-Var) \\cite{Par92}, which is an economical and accurate statistical interpolation scheme that does not include the effect of uncertainty in the forecast. Some centers have switched to the ``four-dimensional variational method'' (4D-Var) \\cite{Rab00}, which incorporates the evolution of uncertainty in linear sense by the used of the adjoint model of the highly nonlinear model. These variational methods always call for the minimization of a cost function (cf. Appendix) that measures the difference between model results and observations throughout some relevant region of space and time. Currently active research areas in data assimilation include the effective and efficient quantification of the time-dependent uncertainties of both the prior and posterior in the analysis. To this end, the ensemble Kalman filter methods have recently received considerable attention motivated by future integration into operational environments \\cite{Hou05,Szu07,Whi07}. As the importance of the uncertainties in data assimilation have become clear, many NWP centers perform ensemble prediction along with the single analysis obtained by the variational methods \\cite{Tot93,Hou95,Mol96}.  Clearly, considerable similarities exist between the data assimilation problem and the model validation problem. Can successful data assimilation be construed as validation of the model?  In our opinion, that would be unjustified because the objectives are clearly different for these problems. As stated above, data assimilation admits the imperfection of the model. It explicitly makes use of the knowledge from the previous data assimilation cycle. As the initialization problem is solved iteratively over relatively short time windows, deviation of the model trajectory from the true evolution of the dynamical system in question tend to be small and data could be assimilated into the model without much discrepancy. Moreover, the operational centers perform careful quality-control of data to eliminate any  isolated ``outliers'' with respect to the model trajectory. Thus, the data assimilation problem differs from the validation problem by design. Nevertheless, it is important to recognize that the resources offered by data assimilation can ensure that models perform well enough for their intended use.   \\subsection{Extension of the Meaning of Validation}  A qualitatively new class of problems arise in fields such as the geosciences that deal with the construction of knowledge of a unique object, planet Earth, whose full scope and range of processes can be replicated or controlled neither in the laboratory nor in a supercomputer. This has led recently to championing the relevance of ``systemic'' (meaning ``system approach'') also called ``complex system'' approaches to the geosciences. In this framework, positive and negative feedbacks (and even more complicated nonlinear multiplicative noise processes) entangle many different mechanisms, whose impact on the overall organization can be neither assessed nor understood in isolation. How does one validate a model using the systemic approach? This very interesting and difficult question is at the core of the problem of validation.  How does one validate a model when it is making predictions on objects that are not fully replicated in the laboratory, either in the range of variables, of parameters, or of scales? For instance, this question is crucial \\begin{itemize} \\item in the scaling the physics of material and rock rupture tested in the laboratory to the scale of earthquakes; \\item in the scaling the knowledge of hydrodynamical processes quantified in the laboratory to the length and time scales relevant to the atmospheric/oceanic weather and climate, not to mention astrophysical systems; \\item in the science-based stewardship of the nuclear arsenal, where the challenge is to go from many component models tested at small scales in the laboratory to the full-scale explosion of an aging nuclear weapon. \\end{itemize}  \\begin{figure} \\begin{center} \\includegraphics[width=12cm]{Hefner_text_low_resolution.eps} \\end{center} \\end{figure}  The same issue arises in the evaluation of electronic circuits. In 2003, Allen R. Hefner, Founder and Chairman of the NIST/IEEE Working Group on Model Validation, writes in its description: {\\it ``The problem is that there is no systematic way to determine the range of applicability of the models provided within circuit simulator component libraries.''} See full-page boxed text for the complete version of this interesting text, as well as Ref.~\\cite{Ber98}. This example of validation of electronic circuits is particularly interesting because it stresses the origin of the difficulties inherent in validation: the fact that the dynamics are nonlinear and complex with threshold effects and does not allow for a simple-minded analytic approach consisting in testing a circuit component by component. Extrapolating, this same difficulty is found in validating general circulation models of the Earth's climate or computer codes of nuclear explosions. The problem is thus fundamentally a ``system'' problem. The theory of systems, sometimes referred to as the theory of complex systems, is still in its infancy but has shown the existence of surprises. The biggest surprise may be the phenomenon of ``emergence'' in which qualitatively new processes or phenomena appear in the collective behavior of the system, while they cannot be derived or guessed from the behavior of each element. The phenomenon of ``emergence'' is similar to the philosophical law on the ``transfer of the quantity into the quality.'' How does one validate a model of such a system? Validation therefore requires an understanding of this emergence phenomenon.  From another angle, the problem is that of extrapolating a body of knowledge, which is firmly established only in some limited ranges of variables, parameters and scales, beyond this clear domain into a more fuzzy zone of unknowns. This problem has appeared and appears again and again in different guises in practically all scientific fields. A particularly notable domain of application is risk assessment; see, for instance, Kaplan and Garrick's classic paper on risks \\cite{Kap81}, and the instructive history of quantitative risk analysis in US regulatory practice \\cite{Rec00}, especially in the US nuclear power industry \\cite{Kee91,Hel94,Hel99,Hel00}. An acute  question in risk assessment deals with the question of quantifying the potential for a catastrophic event (earthquake, tornado, hurricane, flood, huge solar mass ejection, large meteorite, industrial plant explosion, ecological disaster, financial crash, economic collapse, etc.) of amplitude never yet sampled from the knowledge of past history and present understanding.  To tackle this enduring question, each discipline has developed its own strategies, often being unaware of the approaches of others. Here, we attempt a formulation of the problem, and outline some general directions of attack, that hopefully will transcend the specificities of each discipline. Our goal is to formulate the validation problem in a way that may encourage productive crossings of disciplinary lines between different fields by recognizing the commonalities of the blocking points, and suggest useful guidelines. ",
        "conclusions": "The validation of numerical simulations continues to become more important as computational power grows, as the complexity of modeled systems increases, and as increasingly important decisions are influenced by computational models. We have proposed an iterative, constructive approach to validation using \\emph{quantitative measures} and \\emph{expert knowledge} to assess the relative state of validation of a model instantiated in a computer code.  In this approach, the increase/decrease in validation is mediated through a function that incorporates the results of the model vis-\\`a-vis the experiment together with a measure of the impact of that experiment on the validation process. While this function is not uniquely specified, it is not arbitrary: certain asymptotic trends, consistent with heuristically plausible behavior, must be observed. In four fundamentally different examples, we have illustrated how this approach might apply to a validation process for physics or engineering models. We believe that the multiplicative decomposition of trust gains or losses (given in Eq. \\ref{mglerdffd}), using a suitable functional prescription (such as Eq. \\ref{mgler222}), provides a reasoned and principled description of the key elements---and fundamental limitations---of validation. It should be equally applicable to biological and social sciences, especially since it is built upon the decision-making processes of the latter. We believe that our procedure transforms the paralyzing criticisms in Popper's style that ``we cannot validate, we can only invalidate'' \\cite{Ore94a} into a practical constructive algorithm. This strategy addresses specifically both problems of distinguishing between competing models and transforming the vicious circle of lack of suitable data into a virtuous spiral path: each cycle is marked by a quantified increment of the evolving trust we put in a model based on the novelty and relevance of new data and the quality of fits.  We have also surveyed and commented extensively on the V\\&V literature. We hope this digest will help the reader as much as its collation helped us deepen our understanding of the challenge of model validation, including a new perspective on some of our own work.  We close with these far-reaching thoughts by Patrick J. Roache \\cite{Roa04}:  \\begin{quote}  {\\it In an age of spreading pseudoscience and anti-rationalism, it behooves those of us who believe in the good of science and engineering to be above reproach whenever possible. Public confidence is further eroded with every error we make. Although many of society's problems can be solved with a simple change of values, major issues such as radioactive waste disposal and environmental modeling require technological solutions that necessarily involve computational physics. As Robert Laughlin~\\cite{Lau02} noted in this magazine,} ``there is a serious danger of this power [of simulations] being misused, either by accident or through deliberate deception.'' {\\it Our intellectual and moral traditions will be served well by conscientious attention to verification of codes, verification of calculations, and validation, including the attention given to building new codes or modifying existing codes with specific features that enable these activities.}  \\end{quote}"
    },
    "0710/0710.5704_arXiv.txt": {
        "abstract": "We present an analysis of archival \\textit{Chandra} and VLA observations of the E0 galaxy NGC\\,1399 and the E2 galaxy NGC\\,4649 in which we investigate cavities in the surrounding X-ray emitting medium caused by the central AGN. We calculate the jet power required for the AGN to evacuate these cavities and find values of $\\sim 8\\times 10^{41}$ erg s$^{-1}$ and $\\sim 14\\times 10^{41}$ erg s$^{-1}$ for the lobes of NGC\\,1399 and $\\sim 7\\times 10^{41}$ erg s$^{-1}$ and $\\sim 6\\times 10^{41}$ erg s$^{-1}$ for those of NGC\\,4649. We also calculate the $k/f$ values for each cavity, where $k$ is the ratio of the total particle energy to that of electrons radiating in the range of 10 MHz to 10 GHz, and $f$ is the volume filling factor of the plasma in the cavity. We find that the values of $k/f$ for the lobes of NGC\\,1399 are $\\sim93$ and $\\sim190$, and those of the lobes of NGC\\,4649 are $\\sim15000$ and $\\sim12000$. We conclude that the assumed spectrum describes the electron distribution in the lobes of NGC\\,1399 reasonably well, and that there are few entrained particles. For NGC\\,4649, either there are many entrained particles or the model spectrum does not accurately describe the population of electrons. ",
        "introduction": "\\label{sec:intro} Within nearby galaxies, groups, and clusters, embedded active galactic nuclei (AGN) have been found to interact with the surrounding hot, X-ray emitting thermal gas, causing such disturbances as shocks, ripples, and cavities (\\eg \\citealp{Fabian00, Fabian03b, Fabian06, Forman05}). The cavities, which appear as X-ray surface brightness depressions, consist of low density relativistic plasma and are a consequence of the thermal gas being displaced by the jets of the AGN. They have been found to be associated with the central AGN of many clusters, such as Perseus (\\citealp{Bohringer93,Fabian00, Fabian03b, Fabian06}), Hydra A (\\citealp{McNamara00}), M\\,87 (\\citealp{Churazov01, Forman05}), and Centaurus (\\citealp{Taylor02,Taylor06a}). AGN-induced disturbances are also present within the hot interstellar medium (ISM) in the haloes of normal elliptical galaxies \\citep{Diehl06a}. Many early-type galaxies, such as M\\,84 (\\citealp{Finoguenov01}), have been found to harbor cavities in their X-ray haloes. Our interest in these AGN-induced cavities stems from the fact that they can be used as calorimeters. The kinetic energy of the jets emanating from the central black hole can be estimated by calculating the amount of energy required to inflate the observed cavities (\\citealp{Birzan04, Rafferty06, Dunn06f}). Using this estimate, the jet power can be calculated once an approximate age of the cavity has been deduced.  For nearby large, elliptical galaxies, accurate measurements of the density and temperature profiles of the thermal gas can be obtained from high resolution \\textit{Chandra} observations. For those systems that are sufficiently close, such that radii within an order of magnitude of the accretion radius are resolved, these measurements can be used to calculate the Bondi accretion rate, given estimates of the black hole mass. \\citet{Allen06}, in a study of 9 X-ray luminous sources that harbor cavities, found that a tight physical correlation exists between the power available from Bondi accretion of the hot gas and the observed jet power. These results are significant, especially for studies of accretion and jet formation, as well as the formation and evolution of galaxies. In particular, building upon earlier studies by \\citet{DiMatteo03} and \\citet{Taylor06a}, they confirmed that the bulk of the energy produced by the central AGN in these systems, which are all \\citet{Fanaroff74} type-I sources, is released in the relativistic jets. The energy injected back into the environment is in principle sufficient to balance radiative cooling of the hot X-ray gas. This possibility has also been investigated by \\citet{Best05, Best06b}, who showed that the cooling of hot gas from the atmosphere of the host galaxy can feasibly provide the fuel for low power, radio-loud AGN, and that the heating by AGN feedback can balance this cooling, except perhaps in the most massive clusters. However, the manner in which the energy provided by the AGN is coupled to the thermal energy of the hot gas in order to offset cooling, as well as the efficiency of this process, are as yet unknown.  Because of the significance of AGN activity, it is desirable to investigate further the population of galaxies which exhibit interactions between the central AGN and the surrounding medium. In particular, we are interested in those sources which harbor cavities, as they offer insight into the physical properties of the AGN itself.  We analyzed \\textit{Chandra} and VLA\\footnote{The VLA (Very Large Array) is operated by the National Radio Astronomy Observatory (NRAO).} data for NGC\\,1399 and NGC\\,4649, both of which are large, elliptical galaxies that harbor cavities and are located in nearby clusters. Because of their proximity, we are able to determine the density of the X-ray emitting gas down to radii within a factor of 10 of the accretion radius, making these two galaxies ideal for future analysis concerning Bondi accretion. At present, we use the cavities in these systems to estimate the kinetic energy of the jets and jet power. We also investigate the particle content of the cavities by determining $k/f$, where $k$ is the ratio of the total particle energy contained in the cavity to the energy accounted for by electrons emitting synchrotron radiation in the range of 10 MHz to 10 GHz, and $f$ is the volume filling factor of the relativistic plasma in the cavity. The cavities are filled with relativistic charged particles and magnetic fields, which together determine the synchrotron emission from the lobes. The relative energy densities of these two components, and hence $k$ and $f$, cannot be disentangled from the emission alone. We must therefore consider the ratio $k/f$, which, under the assumption that the gas pressure is equal to the total pressure from the radio plasma, can be related to the magnetic field strength.  The maximum field strength is obtained for the minimum value of $k/f=1$, and would indicate a pure electron-positron plasma uniformly filling the cavity.  Our discussion will proceed as follows. We begin in \\S \\ref{sec:radio} by describing the radio data selection and processing. In \\S \\ref{sec:xray}, we describe the X-ray data and preparation, as well as the identification of cavities. In \\S \\ref{sec:interaction} we present \\textit{Chandra}/VLA comparisons for NGC\\,1399 and NGC\\,4649 and determine the power of the jets in these systems. We present the $k/f$ ratios in \\S \\ref{sec:kf}, followed by conclusions in \\S \\ref{sec:concl}.  Throughout this discussion, we assume a flat $\\Lambda$CDM cosmology with $\\Omega_M=0.3$ and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "\\label{sec:concl}  In analyzing \\textit{Chandra} and VLA data for NGC\\,4649, we have found cavities in the surrounding X-ray emitting gas that have not been previously reported. We used these cavities and those in NGC\\,1399 to estimate the power of the jets emanating from the central black hole of these systems. The $P_{\\rm jet}$ values for NGC\\,1399 using the sound speed expansion timescale are $\\sim8\\times10^{41}$ erg s$^{-1}$ for the northern lobe and $\\sim14\\times10^{41}$ erg s$^{-1}$ for the southern lobe. The values using the buoyancy rise time are on average a factor of 2 lower for both lobes, but are roughly in agreement with the aforementioned values within the errors. As the cavities in NGC\\,1399 are detached from the core but presently powered by the AGN jets, it is difficult to determine which timescale is the relevant one in this case. In NGC\\,4649, the $P_{\\rm jet}$ values using the sound speed expansion timescale are $\\sim7\\times10^{41}$ erg s$^{-1}$ for the northern lobe and $\\sim6\\times10^{41}$ erg s$^{-1}$ for the southern lobe. Calculations using the buoyancy rise time yield similar values, within the errors.  Comparing these jet powers to those calculated in \\citet{Allen06}, we find that our sources have jet powers that are lower than most of those in their sample. Examples are M\\,87 and NGC\\,4696, both of which are located within dense cluster environments. M\\,87 is near the center of the Virgo cluster and has jet powers which are roughly an order of magnitude larger than those of our sources. The jet powers for the lobes of NGC\\,4696, which is in the center of Centaurus, are roughly 3 to 4 times those of our sources. We find that our sources have similar jet powers to those of NGC\\,4552 in Virgo, which has values of $\\sim7\\times10^{41}$ erg s$^{-1}$ and $\\sim8\\times10^{41}$ erg s$^{-1}$, as well as those of NGC\\,5846, which dominates a group of intermediate mass and has jet powers of $\\sim4\\times10^{41}$ erg s$^{-1}$ for each lobe.  We also investigated the particle content of the cavities by calculating $k/f$. We find that the $k/f$ values for NGC\\,4649 are much larger than those of NGC\\,1399 by about 2 orders of magnitude. \\citet{Dunn05} found that there is a large range of values for $k/f$ in cavities found in cluster galaxies. In general, $k/f$ is in the range of $\\sim 1$ to $\\sim 1000$ for cavities that are active. The values for NGC\\,1399 fall into this range; however, in the case of NGC\\,4649, the values for $k/f$ using either timescale are more than a factor of 10 larger than the upper limit of what has been previously observed.  The minimum value of $k/f=1$ corresponds to a lobe which is filled with a uniform electron-positron plasma with energies corresponding to emission in the range of 10~MHz to 10~GHz. Large values of $k/f$, such as those found in NGC\\,4649, can be achieved by large values for $k$ or small values for $f$. Large values for $k$ imply that, in order for there to be pressure balance, there must be particles present within the lobe which are not accounted for by the assumed spectrum. These particles can be electrons whose emission is outside the range of 10~MHz to 10~GHz or thermal protons that either exist in the jets during their formation or are entrained as the cavities travel through the surrounding medium. The filling factor $f$ describes the fraction of the volume of the cavity that is occupied by radio emitting plasma. The remaining volume fraction can either be occupied by thermal gas or plasma whose relativistic electrons have aged. \\citet{Dunn05}, in studying ghost cavities as well as active cavities in Centaurus and Perseus, found that the older cavities have larger $k/f$ values than the active ones and have attributed this to the aging of the relativistic electrons.  The $k/f$ values for NGC\\,4649 are similar to those of ghost cavities in which the relativistic electrons have aged, though the values can also suggest that the the lobes contain entrained particles. If the electrons in the lobes of NGC\\,4649 have aged, the model spectrum is simply not an accurate description of the electron distribution. This scenario seems likely, as the radio flux is so small that NGC\\,4649 is expected to soon become a ghost cavity system should clear cavities remain. For NGC\\,1399, the model spectrum appears to be an appropriate description of the lobe contents, and there are probably not many entrained particles. As was proposed by \\citet{Dunn04}, higher power radio sources are less likely to pick up extra particles from the surrounding medium. Our results agree with this description in that NGC\\,1399 has lower $k/f$ values than NGC\\,4649 and also higher radio power. However, there are many physical parameters to consider in investigating these values, and no clear correlation has been found between $k/f$ and any one parameter.  In calculating $k/f$ we have assumed a simple power-law spectrum. Using high frequency observations to obtain radio flux measurements of a lobe whose relativistic electrons have aged would tend to cause this power-law approximation to underestimate the energy of the relativistic electrons in the lobe and therefore overestimate $k/f$. Our results, then, could be said to suggest that the lobes in NGC\\,4649 are likely to be older than those in NGC\\,1399. However, these results could also suggest that the power-law approximation may be poor in both cases. In order to determine whether the power-law assumption is reasonable, it would be desirable to obtain sensitive lower frequency VLA observations of these galaxies, such as at 330~MHz, that would allow for accurate measurements of the radio flux in the regions corresponding to the cavities."
    },
    "0710/0710.2062_arXiv.txt": {
        "abstract": "{For the SUSY 2007 conference, I was asked to review the topic of indirect searches for dark matter. As part of that talk, I summarized several observations which have been interpreted as the product of dark matter annihilations. In my contribution to the proceedings, I have decided to focus on this aspect of my talk. In particular, I will discuss the cosmic positron spectrum measured by HEAT, the 511 keV emission from the Galactic Bulge measured by INTEGRAL, the diffuse galactic and extragalactic gamma ray spectra measured by EGRET, and the microwave excess from the Galactic Center observed by WMAP. \\PACS{ {95.35.+d}{Dark matter} } % } % ",
        "introduction": "\\label{intro}  The quest to uncover dark matter's particle identity is a multifaceted one. In addition to collider searches for stable, weakly interacting particles, a wide range of astrophysical searches for dark matter are underway. These astrophysical searches can be classified as direct detection experiments, which hope to observe the elastic scattering of dark matter particles with nuclei, and indirect experiments, which search for the annihilation products of WIMPs, including gamma rays, neutrinos, positrons, antiprotons, antideuterons and synchrotron radiation.  Over the past several years, a number of observations have been interpreted as possible products of dark matter annihilations. Here, I take the opportunity to summarize and discuss five of these observations. After briefly reviewing the basic features of these signals, I will discuss the particle physics and astrophysics that is required to generate them through the process of dark matter annihilations. ",
        "conclusions": "\\label{summary}  I have summarized here five different astrophysical signals which have been interpreted as possible products of dark matter annihilations. Each of these has their own strengths and weaknesses. Attempting to remain as objective as possible, I would like offer my own assessment of these various observations.  \\begin{table*} \\caption{A comparison of the particle physics and astrophysics required in each of the five scenarios I have discussed.} \\label{tab:1}       % \\begin{tabular}{lll} \\hline\\noalign{\\smallskip} Signal & Required Particle Physics & Required Astrophysics  \\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} HEAT positron spectrum &Non-thermal production mechanism or... & Boost factor of 50 or more  \\\\ & Electroweak-scale WIMP & \\\\ \\hline\\noalign{\\smallskip} INTEGRAL 511 keV line & MeV mass dark matter particle & Mildly cusped halo profile \\\\ & p-wave annihilation cross section & \\\\ \\hline\\noalign{\\smallskip} Diffuse galactic $\\gamma$-rays & $\\sim$50-100 GeV WIMP & Two massive dark matter rings \\\\ &&Anisotropic diffusion (convention, etc.) \\\\ \\hline\\noalign{\\smallskip} Diffuse extragalactic $\\gamma$-rays & $\\sim$500 GeV WIMP & Dense cusps common  \\\\ && No cusp in the Milky Way  \\\\ \\hline\\noalign{\\smallskip} WMAP Haze & Electroweak-scale WIMP & Cusped halo profile  \\\\ &Cross section predicted for a thermal relic & No boost factors/exotic astrophysics  \\\\ \\noalign{\\smallskip}\\hline \\end{tabular} \\vspace*{1cm}  % \\end{table*}  In Table~\\ref{tab:1}, I have summarized the nature of the particle physics and astrophysics which must be introduced to generate the observed signals in each of the five scenarios I have discussed. Based on these requirements, I would like to make the following points:  \\begin{itemize}  \\item{In the case of the positron excess and each of the two EGRET excesses, there is a significant degree of difficulty involved in generating an annihilation rate sufficiently large to produce the observed signal. As a result, some combination of large annihilation cross sections or cusped profiles (or boost factors) are required. That being said, such rates might be possible, even though they are higher than our naive expectations suggest.}  \\item{The two EGRET signals each require somewhat more complicated astrophysical setups if they are to be interpreted as products of dark matter annihilation. In the case of the Galactic excess, the diffusion model has to be modified to avoid the overproduction of antiprotons. Also, a rather exotic (although not impossible) dark matter distribution also has to be adopted. Regarding the extragalactic signal, the annihilation rate has to be very large in most galaxies, but not in the Milky Way. While not impossible, it requires us to find ourselves living somewhere marginally exceptional.}  \\item{The INTEGRAL signal stands out among these five in that it requires a very different dark matter particle than is generally considered. The conventional wisdom is that it is not easy to construct a viable particle physics model containing an MeV dark matter candidate. I leave it to you to judge the merits of this argument.}  \\item{The WMAP Haze stands alone among these five signals in that it requires neither unexpected particle physics or astrophysics in order to be generated through dark matter annihilations. The observed angular distribution, intensity and spectrum are consistent with a vanilla electroweak-scale particle which was produced thermally in the early universe and is distributed with a cusped halo profile in our Galaxy.}  \\item{It is worth noting that {\\it none} of the signals I have discussed here are particularly distinctive -- there is not yet a ``smoking gun''. Further observations will be required to draw any conclusive statements. Fortunately, new data from PAMELA, GLAST, Planck and even the LHC are not too far off. In a few years time, most (if not all) of these signals will have either gone away or become much more interesting.}  \\end{itemize}"
    },
    "0710/0710.0298_arXiv.txt": {
        "abstract": "The INTEGRAL/IBIS survey was performed collecting all the GPS and GCDE data together with all the available public data . The second catalogue, published in 2006 by \\citet{Bird1}, is dominated by detection of 113 X-ray binaries, with 38  being high--mass and 67 low--mass. In most systems the compact object is a neutron star, but the sample also contains 4 confirmed Black Holes and 6 LMXB black hole candidates (BHC). There are also, in additional, 6 tentative associations as BHCs based simply on spectral and timing properties. In the sample of 12 sources (BHC and tentatively associated BHC) there are 7 transient sources that went into outbursts during the INTEGRAL survey observations. We present here the monitoring of the time and spectral evolution of these 7 outbursts. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.0251_arXiv.txt": {
        "abstract": "We investigate the effects of neutrino-nucleus interactions (the $\\nu$-process) on the production of iron-peak elements in Population III core-collapse supernovae. The $\\nu$-process and the following proton and neutron capture reactions produce odd-$Z$ iron-peak elements in complete and incomplete Si burning region. This reaction sequence enhances the abundances of Sc, Mn, and Co in the supernova ejecta. The supernova explosion models of 15 $M_\\odot$ and 25 $M_\\odot$ stars with the $\\nu$-process well reproduce the averaged Mn/Fe ratio observed in extremely metal-poor halo stars. In order to reproduce the observed Mn/Fe ratio, the total neutrino energy in the supernovae should be $3 - 9 \\times 10^{53}$ ergs. Stronger neutrino irradiation and other production sites are necessary to reproduce the observed Sc/Fe and Co/Fe ratios, although these ratios increase by the $\\nu$-process. ",
        "introduction": "The abundance-metallicity relation of low mass extremely metal-poor (EMP) stars ($-4 \\la [{\\rm Fe/H}] \\la -3$; $[{\\rm X/Y}] \\equiv \\log_{10} (N_{\\rm X}/N_{\\rm Y}) - \\log_{10} (N_{\\rm X}/N_{\\rm Y})_\\odot$ where $N_{\\rm X}$ and $N_{\\rm Y}$ are the abundances of elements X and Y, respectively) and very metal-poor (VMP) stars ($-3 \\la [{\\rm Fe/H}] \\la -2$) has been clarified by observations of very high quality spectra \\citep[e.g.,][]{nr01,ca04}. This relation is expected to provide the information of supernovae (SNe) of Population (Pop) III (the first generation) massive stars \\citep[e.g.,][]{nt06} and the first stage of Galactic chemical evolution (GCE) \\citep[e.g.,][]{kn06}. The EMP stars are considered to have suffered the injection of heavy elements from one or a few SNe \\citep[e.g.,][]{st98}. A halo interstellar medium becomes gradually homogeneous with metallicity \\citep{as00}. Although recent studies of the star formation in the first stage of GCE have been in progress, the characteristics of first generation stars, such as initial mass function, have not been clarified. In order to clarify such characteristics, we should investigate nucleosynthesis in the candidates of the first generation stars as well as the observed variations of the abundance distributions of the EMP and VMP stars.  Recent observations have indicated that the abundance ratios [X/Fe] averaged in the metal-poor stars with $-4 \\la [{\\rm Fe/H}] \\la -2$ are between $-0.5$ and 0.5 for most of observed elements \\citep{ca04}. Their scatter around the average values deduced from the new observation is much smaller than found in earlier studies. The elemental abundance distribution of the metal-poor stars is not so different from that of the solar-system composition. Therefore, the small scatter is considered to suggest primordial burst of high-mass stars or the very rapid mixing of matter from different bursts of early star formation.   From theoretical viewpoint, \\citet{tu07} compared the Pop III hypernova (HN) yields with the abundance pattern of EMP stars with the metallicity range $-4.2 < [{\\rm Fe/H}] < -3.5$ in \\citet{ca04}. They adopted the mixing-fallback model for the HNe. The HN yields give good agreement with the observed abundances for C, Na, Mg, Al, Si, Ca, Ni, and Zn. They also compared the Pop III SN and HN yields integrated over the Salpeter initial mass function (IMF) with the abundances of VMP stars in the metallicity range $-2.7 < [{\\rm Fe/H}] < -2.0$ in \\citet{ca04}. They reproduced the observed trends of abundance ratios for C, O, Na, Mg, Al, Si, Ca, Cr, Ni, and Zn. The SN and HN yields of N, K, Sc, Ti, Mn, and Co are short for explaining the observed yield ratios, although Sc and Ti may be improved in the high-entropy explosion models. The yields of Mn and Co may be reproduced by the modification of electron fraction $Y_e$, of which value is still uncertain in the innermost region of SNe and HNe. Additional nucleosynthesis processes also remain open for these elements.     During a SN explosion, the collapsed core becomes extremely high temperature and high density, and neutrinos are produced through pair creation. The neutrinos are almost thermalized by neutrino-nucleus scattering at the center. They pass away through the neutrino sphere with carrying the gravitational binding energy of the core. The neutrino irradiation lasts for about $\\sim 10$ s. The total number of neutrinos emitted in a SN is very huge, i.e., about $N_\\nu \\sim 10^{58}$. The neutrinos interact with nuclei in the surrounding stellar material to produce new species of nuclei. This is called the $\\nu$-process.  The $\\nu$-process is expected to be an important production process for some elements in Pop III SN explosions. The $\\nu$-process is important for the synthesis of light elements, Li and B \\citep{de78,wh90,ww95,yt04,yk05,yk06a,yk06b}, F \\citep[e.g.,][]{wh90,ww95}, and some heavy neutron-deficient nuclei such as $^{138}$La and $^{180}$Ta \\citep{ga01,hk05}. The $\\nu$-process products are synthesized through the spallations by neutrinos from abundant seed nuclei. It also produces protons and neutrons and their capture reactions also enhance the abundances of some elements. For Pop III SNe, the seed nuclei are $\\alpha$-nuclei abundantly produced in complete and incomplete Si burning during the explosions. Thus, the $\\nu$-process in Pop III SNe will produce additional some odd-$Z$ elements.   We note that most of the nucleosynthesis calculations including the $\\nu$-process have been carried out using spherical explosion models \\citep{ww95,yu05}. Although the total neutrino energy should be close to the gravitational binding energy released from the collapsing core, properties of the neutrinos are still uncertain. The properties strongly depend on explosion mechanism but details of the mechanism have not been clarified. On the other hand, aspherical explosions have been investigated recently \\citep[e.g.][]{ks04,sp04}. In such a case, aspherical structure of neutrino sphere can be formed and the properties of the neutrinos are also still uncertain. Therefore, the dependence of the abundances of Pop III SNe and HNe on neutrino properties should be investigated.   In this study, we focus on the effect of the $\\nu$-process on the production of odd-$Z$ iron-peak elements, Sc, Mn, and Co, during Pop III SN and HN explosions. We calculate the SN nucleosynthesis including the $\\nu$-process. Then, we investigate the dependence of the abundances of Sc, Mn, and Co in the SN ejecta on the energy of neutrinos, the stellar mass, and the explosion energy. Sc, Mn, and Co are expected to be at least partly synthesized through the $\\nu$-process. We also indicate whether the abundance ratios of Sc, Mn, and Co to Fe reproduce those observed in the EMP stars when we take into account the $\\nu$-process.   In \\S 2 we explain stellar evolution models and the explosion models of SNe and a HN. We also describe the nucleosynthesis models including the $\\nu$-process and the neutrino irradiation models. In \\S 3 we show the abundance distributions of Sc, Mn, and Co in the ejecta of the SN and HN models. We explain additional production of these elements through the $\\nu$-process. In \\S 4 we explain the effect of the $\\nu$-process on the abundance ratios of Sc, Mn, and Co to Fe and the dependence on the neutrino irradiation strength. We also compare our results with the abundance ratios observed in the EMP stars. In \\S 5 we discuss other production processes of Sc, Mn, and Co proposed in recent studies and the effect of the $\\nu$-process. Finally, we summarize this study in \\S 6. ",
        "conclusions": "\\subsection{Abundance Pattern of VMP Stars}  The abundance pattern of VMP stars are considered to reflect partly mixed interstellar materials in the early Universe. Recently, the averaged abundance pattern of VMP stars with $-2.7 < [{\\rm Fe/H}] < -2.0$ in \\citet{ca04} has been reproduced by the Pop III SN and HN yields integrated over the Salpeter IMF function \\citep{tu07}. However, the evaluated [Sc/Fe] and [Mn/Fe] are still short of the observed ratio. We discuss the enhancement of these ratios by the $\\nu$-process.  \\citet{tu07} evaluated the Sc/Fe ratio as [Sc/Fe] = $-1.5$ whereas the observed value is [Sc/Fe] = $0.12 ^{+0.24}_{-0.14}$. In this study, we showed that the $\\nu$-process raises the Sc yield by 1.1 dex at the maximum. However, this enhancement is still short for reproducing the observed value. Production sites other than the $\\nu$-process is needed to reproduce it (see \\S 5.2).  The [Mn/Fe] ratio observed in VMP stars is $-0.48 ^{+0.10}_{-0.24}$. The evaluated value of [Mn/Fe] is $-1.14$. When we take into account the $\\nu$-process, this value increases by $0.24 - 0.4$ dex with $E_\\nu = 3 \\times 10^{53}$ ergs and $0.49 - 0.73$ dex with $E_\\nu = 9 \\times 10^{53}$ ergs. The increase by the $\\nu$-process with $E_\\nu = (3-9) \\times 10^{53}$ ergs will provide better fitting of the Pop III SN and HN yields to [Mn/Fe] in VMP stars.  The Pop III SN and HN models in \\citet{tu07} indicate [Co/Fe] close to the lowest ratio in VMP stars; [Co/Fe] = $0.29 ^{+0.17}_{-0.27}$. Taking into account the $\\nu$-process with $E_\\nu = 3 \\times 10^{53}$ ergs, [Co/Fe] increases by $0.37 - 0.53$ dex in SNe and it does not in HNe. This enhancement is enough to reproduce the average [Co/Fe] ratio in VMP stars. The $\\nu$-process of the Pop III SNe and HNe is one of the preferable sites of the Mn and Co in VMP stars.   \\subsection{Other Production Sites of Sc, Mn, and Co}  We have shown that Sc, Mn, and Co are produced through the $\\nu$-process in the Pop III SNe and HNe. For other than Mn, the $\\nu$-process in the SNe and HNe does not produce large enough abundances to reproduce the abundances observed in low mass EMP stars. So, we will discuss other synthesis processes proposed in recent studies.   \\subsubsection{Sc}  \\begin{center} {\\it Aspherical Explosions} \\end{center}  As shown in previous studies \\citep[e.g.,][]{un02}, spherically symmetrical SN and HN models do not provide enough amount of Sc to reproduce the observed abundance of low mass EMP stars. On the other hand, aspherical explosion and jet-like explosion produce larger amount of Sc \\citep{mn03}. In the case of spherical explosions, low-density structure enables to supply larger amount of Sc \\citep{un05}. Thus, higher temperature and low density, which are realized by aspherical explosions, are one of the favorable environments to produce Sc. The amount of Sc produced in SNe and HNe is also sensitive to $Y_e$ in complete Si burning region. Large $Y_e$ value ($Y_e > 0.5$) brings about a large amount of Sc \\citep{iu06}.   \\begin{center} {\\it Nucleosynthesis in Innermost Region of SNe} \\end{center}  Detailed thermal evolution of the innermost region of SN ejecta, i.e., the neighbor of the ^^ ^^ mass cut'' has not been solved strictly. However, recent progress of hydrodynamical calculations has gradually revealed thermal evolution in such a deep region. At the same time, detailed nucleosynthesis in such a deep region has also been investigated. \\citet{pw05,ph06} investigated the nucleosynthesis in convective bubbles in the innermost region of SN ejecta and in early stage of neutrino-driven winds. They used results of two-dimensional hydrodynamical calculations by \\citet{jb03}. \\citet{fr06a} examined the simulations of one-dimensional self-consistent treatment of core-collapse SNe with modified neutrino cross sections to explode. They also investigated the nucleosynthesis in the innermost region just above the mass cut produced self-consistently.   Both of the studies showed that the electron fraction $Y_e$ exceeds 0.5 in such regions, which enhances the production of  Sc. In the latter study, they took into account all charged-current weak-interactions that changes $Y_e$ in the ejecta. They obtained that the amount of Sc with $10^{-6} M_\\odot$ is produced in the innermost region with $Y_e > 0.5$. The Sc/Fe ratio evaluated in their study is slightly smaller than the averaged abundances observed in low mass EMP stars \\citep{ca04}. Thus, the Sc/Fe ratio is larger than our result. They also showed that a significant amount of nuclei with mass of $A > 64$ in this region, especially light $p$-nuclei such as $^{92,94}$Mo and $^{96,98}$Ru \\citep[the $\\nu p$-process;][]{fr06b}.   It is noted that the region of $Y_e > 0.5$ is convectively unstable so that the region will mix with outer smaller $Y_e$ region in dynamical time scale. However, they expected that the $Y_e$ remains high in an average sense. Indeed, the former study took into account the convection in two-dimensional simulations and obtained $Y_e > 0.5$. We also note that neutral-current $\\nu$-process reactions were not included in \\citet{fr06a,fr06b}. They also pointed out that neutrino-induced spallation reactions can change the final abundances of some nuclei. We expect that these reactions enhance the abundance of Sc and that the Sc/Fe ratio becomes close to the observed ratio.  It is also noted that materials blown off in neutrino-driven winds in early time after a SN explosion can achieve large $Y_e$ value. Accordingly, the $rp$-process can take place in the materials. This synthetic process is strongly preferable for the production of the light $p$-nuclei \\citep{wa06}. However, the contribution of the $rp$-process in neutrino-driven winds to Sc production is small.   \\subsubsection{Mn}   The amount of Mn produced in SNe and HNe is strongly sensitive to electron fraction $Y_e$ in incomplete Si burning region \\citep{un02,un05}. Their HN models modifying $Y_e$ value to $Y_e \\simeq 0.4995 - 0.4997$ produce enough amount of Mn to reproduce the observed one in low mass EMP stars. However, HN models, whose $Y_e$ in incomplete Si burning region is not modified, indicated that the amount of Mn is less than the observed value. Aspherical explosions suppress Mn production \\citep{mn03}. We have shown that the Mn production during SN explosions is mainly due to the $\\nu$-process rather than to the incomplete Si burning. The $\\nu$-process with normal neutrino irradiation ($E_\\nu = 3 \\times 10^{53}$ ergs) in SNe reproduces the Mn/Fe ratio averaged in low mass EMP stars even without $Y_e$ modifications. In the case of $E_\\nu = 9 \\times 10^{53}$ ergs, the Mn amount marginally reproduces the upper limit of the observed Mn abundance. In the case of $E_\\nu = 3 \\times 10^{54}$ ergs, however, the SN models overproduce Mn. Thus, the $\\nu$-process in the Pop III SNe is one of the main production processes of Mn observed in low mass EMP stars. The observed abundance constrains the total neutrino energy $E_\\nu$ in SNe. The observational constraint is $E_\\nu \\la 3 - 9 \\times 10^{53}$ ergs.   \\subsubsection{Co}  A large amount of Co can be produced through complete Si burning in aspherical explosion \\citep{mn03}. For spherical explosions, \\citet{un05} pointed out that the Co abundance is significantly enhanced for $Y_e > 0.5$ in the case of HNe. However, the Co/Fe ratio in \\citet{fr06a} is still smaller than the observed ratio in \\citet{ca04}. Their model corresponds to a normal SN explosion. Therefore, HNe ($E_{51} > 10$) with $Y_e \\ga 0.5$ produce Co of which amount reproduces observed Co/Fe ratio in low mass EMP stars. Neutral-current $\\nu$-process reactions may help additional production in the deep region just above the mass cut. However, in the case of HNe, large amount of protons are also produced through $\\alpha$-rich freeze out. This effect may hinder the proton production by the $\\nu$-process.   \\subsection{Comparison with Other SN Nucleosynthesis Models}  The nucleosynthesis in Pop III SNe have been conducted in several groups. Here, we briefly compare with the abundances of Sc, Mn, and Co of 15 $M_\\odot$ and 25 $M_\\odot$ Pop III SN models in \\citet{ww95}, hereafter abbreviated to WW95, (Z15A and Z25B models) and in \\citet{cl04}, abbreviated to CL04. The SN nucleosynthesis models in WW95 include the $\\nu$-process. The temperature of $\\nu_{\\mu,\\tau}$ in WW95 is assumed to be 8 MeV, which is larger than the one adopted in this study. The SN nucleosynthesis with updated physics inputs of WW95 has been shown in \\citet{rh02} but the Pop III SN nucleosynthesis has not been calculated. The nucleosynthesis models in CL04 do not include the $\\nu$-process. The explosion models in CL04 were set to be the ejected mass of $^{56}$Ni equal to 0.1 $M_\\odot$.  \\citet{ww95} indicate [Sc/Fe]=$-0.23$ for $15 M_\\odot$ model; the Sc/Fe is much larger than our result. On the other hand, CL04 indicates Sc/Fe slightly larger than our results of SN models without the $\\nu$-process ([Sc/Fe]=$-1.58$ for 25 $M_\\odot$ model). Although the Sc/Fe in WW95 is much larger than the other models, it is still smaller than the observed ratio. Therefore, other synthesis processes such as discussed above are necessary for Sc production.  The Mn/Fe shown in WW95 is smaller than the observed ratio ([Mn/Fe]$\\le -0.50$). The 15 $M_\\odot$ model in CL04 also indicates the Mn/Fe smaller than the observed one ([Mn/Fe]$=-0.56$). The 25 $M_\\odot$ model in CL04 indicates larger Mn/Fe ratio, which is comparable to the observed one ([Mn/Fe]$=-0.33$). We showed that the $\\nu$-process increases the Mn abundance, but difference of stellar evolution models may affect the abundance.  The 15 $M_\\odot$ models of WW95 and CL04 show [Co/Fe] of $-0.22$ and $-0.11$, which are larger than the ratio of our SN models. Their 25 $M_\\odot$ models indicate the Co/Fe similar to our results without the $\\nu$-process. The evaluated Co/Fe ratios are still short of the observed ratio, so that HNe will have an important role for Co production in the early Universe.     \\subsection{The $\\nu$-Process for Other Elements}  \\subsubsection{Li and B}  Among Li, Be, and B, $^7$Li and $^{11}$B are mainly produced through the $\\nu$-process. When the $\\nu$-process is not considered, the Li/Fe and B/Fe ratios are much smaller than the corresponding solar ratio. SN explosions with the neutrinos of $E_{\\nu} = 3 \\times 10^{53}$ ergs bring about Li yield ratio of [Li/Fe] $\\ga -1$. They also produce B of which yield ratio to Fe is larger than the solar ratio. Therefore, SN explosions may contribute to the B production in the early universe. The $^7$Li is mainly produced as $^7$Be in the He-rich region. The $^{11}$B is produced as $^{11}$C and $^{11}$B in C-enriched oxygen layer and C-enriched He layer. The Li and B production in HNe is less effective than that in SNe. Higher maximum temperature in the He layer decomposes $^7$Be and $^{11}$C produced through the $\\nu$-process and the following $\\alpha$-captures. In addition, strong explosion makes the region of C-rich oxygen layer small.   \\subsubsection{F}  The $\\nu$-process is a main F production process in SNe \\citep{wh90}. The F/Fe ratio including the $\\nu$-process is much larger than the one without the $\\nu$-process. The seed nucleus of F in the $\\nu$-process is $^{20}$Ne. Most of F is produced in O and Ne-enriched region. Stellar mass dependence of the F/Fe ratio is small in the SN models. On the other hand, the HN model shows smaller F/Fe ratio. The HN explosion brings about stronger explosive Ne burning, so that the O and Ne-enriched region becomes smaller.  \\subsubsection{Na and Al}  Figures $4-6$ show that the $\\nu$-process scarcely affects the yields of Na and Al. Furthermore, the yield ratios are much smaller than the observed ratios even in the case of the largest neutrino irradiation. Although the $\\nu$-process enhances the abundances of Na and Al in incomplete Si burning region, Na and Al are mainly produced in carbon and neon shells. The yields produced in the outer two shells strongly depend on the nucleosynthesis in stellar evolution. \\citet{iu05} successfully reproduced relatively large [Na/Fe] and [Al/Fe] observed in two hyper-metal-poor stars. They considered that the effect of overshooting during preSN evolution and extensive matter mixing and fallback during SN explosions with small explosion energies.   \\subsubsection{V}  Vanadium is mainly produced as $^{51}$Mn in incomplete Si burning region in SNe. In the case of the HN model, it is mainly produced in complete Si burning region. Additional production through the $\\nu$-process from $^{52}$Fe in incomplete Si burning region contributes to the enhancement of V in SNe. Low metallicity stars with $-3 \\la {\\rm [Fe/H]} \\la -2$ indicate the V abundance ratios of $-0.2 \\la {\\rm [V/Fe]} \\la 0.6$ \\citep[e.g.,][]{ha04,kn06}. The average abundance ratio is reproduced by the 15 and 25 $M_\\odot$ SN models with the $\\nu$-process of $E_\\nu \\sim 9 \\times 10^{53}$ ergs. In order to reproduce the V/Fe ratios, strong neutrino irradiation in the SNe is favored. In the case of the HN model, V is also produced through the $\\nu$-process. However, the produced abundance is much smaller than that produced in complete Si burning region."
    },
    "0710/0710.1194_arXiv.txt": {
        "abstract": "{} {This paper reports on the detection of optical novae in our neighbour galaxy M\\,31 based on digitized historical Tautenburg Schmidt plates. The accurate positions of the detected novae lead to a much larger database when searching for recurrent novae in M\\,31.} {We conducted a systematic search for novae on 306 digitized Tautenburg Schmidt plates covering a time span of 36 years from 1960 to 1996. From the database of both $\\sim 3 \\times 10^5$ light curves and $\\gtrsim 10^6$ detections on only one plate per colour band, nova candidates were efficiently selected by automated algorithms and subsequently individually inspected by eye.} {We report the detection of 84 nova candidates. We found 55 nova candidates from the automated analysis of the light curves. Among these, 22 were previously unknown, 12 were known but not identified on Tautenburg Schmidt plates before, and 21 novae had been discovered previously on Tautenburg plates. An additional 29 known novae could be confirmed by the detailed investigation of single detections. One of our newly discovered nova candidates shows a high position coincidence with a nova detected about 30 years earlier. Therefore, this object is likely to be a recurrent nova. Furthermore, we re-investigated all 41 nova candidates previously found on Tautenburg plates and confirm all but two. Positions are given for all nova candidates with a typical accuracy of $\\sim 0\\,\\farcs4$. We present light curves and finding charts as online material. } {The analysis of the plates has shown the wealth of information still buried in old plate archives. Extrapolating from this survey, digitization of other historical M\\,31 plate archives (e.g. from the Mount Wilson or Asiago observatories) for a systematic nova search looks very promising.} ",
        "introduction": "Classical novae (CN) are thermonuclear explosions on the surface of white dwarfs (WD) in cataclysmic binary systems. These explosions are caused by the transfer of matter from the companion to the WD. The transferred hydrogen-rich matter is accumulated on the surface of the WD until hydrogen ignition starts as a thermonuclear runaway process in the degenerated matter of the envelope. The resulting expansion of the hot envelope causes the luminosity of the star to rise by more than 9 magnitudes within a few days (see Hernanz \\cite{Hernanz05}; Warner \\cite{Warner95}; and references therein).  A comprehensive compilation of novae in nearby galaxies is desirable for various issues such as nova physics (e.g., Pietsch et al. \\cite{Pietsch05}) or the distributions of stellar populations (e.g., Ciardullo et al. \\cite{Ciardullo87}; Hatano et al. \\cite{Hatano97}; Yungelson et al. \\cite{Yungelson97}). In our galaxy, the investigation of the nova population is hampered by the large area (namely the whole sky) to be scrutinised and by our unfavourable position close to the Galactic Plane. For nearby extragalactic systems, the situation is comparatively more promising. In particular, our huge neighbour galaxy M\\,31 is very suitable for different kinds of statistical studies of novae because of its proximity and the high number of optical nova outbursts per year.  Pietsch et al. (\\cite{Pietsch05}) find that optical novae represent the major class of supersoft X-ray sources (SSS) in M\\,31. More than 30\\% of the optical novae in the M\\,31 centre area show an SSS phase within a year (Pietsch et al. \\cite{Pietsch07}; hereafter PHS07). The X-ray monitoring of optical novae allows the mass of the ejecta, the burned mass, and the mass of the WD in the system to be constrained.  There have been various systematical surveys of novae in M\\,31 for more than 80 years starting with the pioneering work of Hubble (\\cite{Hubble29}) and Arp (\\cite{Arp56}). Among the many other contributors, we refer here only to the work of Rosino et al. (\\cite{Rosino64}, \\cite{Rosino73}, \\cite{Rosino89}), Ciardullo and co-workers (e.g., \\cite{Ciardullo87}), Sharov \\& Alksnis (\\cite{Sharov91}, \\cite{Sharov92}, \\cite{Sharov97}), and Shafter and co-workers (e.g., Shafter \\& Irby \\cite{Shafter01}). More recent efforts include, in particular, the AGAPE project (Ansari et al. \\cite{Ansari04}), the POINT-AGAPE project (Darnley et al. \\cite{Darnley04}), and the WeCAPP project (Fliri et al. \\cite{Fliri06}). The result of this huge effort is a relatively large number of known novae in M\\,31. Recently, PHS07 collected 719 optical novae and nova candidates in M\\,31 with outbursts before the end of 2005 based on their own research and a search in the literature. This catalogue is regularly updated\\footnote{available via the Internet at \\hfill \\\\ http://www.mpe.mpg.de/~m31novae/opt/m31}. PHS07 point out that position information (for some also time of outburst) is poor for many of the early nova detections. The positions could be significantly improved by re-analysing of the original plates. Accurate positions - also for historical novae - are important for identifying novae in different wavelength regimes, and especially for secure identification of recurrent novae.  A successful search for optical novae over the whole body of M\\,31 requires a large set of observations that have to be deep enough and have to cover a sufficiently large ($\\sim\\,10$ square degrees) field. For such an aim, the plate archives of large Schmidt telescopes provide very useful observational material. The present study is concerned with the M\\,31~plates in the archive of the Tautenburg Schmidt telescope (Sect.\\,2). Some plates, in particular the oldest ones, have been used already for nova searches by eye (Moffat \\cite{Moffat67}; B\\\"orngen \\cite{Boerngen68}; Meinunger \\cite{Meinunger75}). Here we describe the first automised, systematic nova search in the M\\,31 field on a large and complete selection of digitized Tautenburg Schmidt plates of the M\\,31 field.  Preliminary results of the present study have already been briefly summarised (Henze et al. \\cite{Henze06}). The present paper includes a more detailed description of the observational data (Sect.\\,2), the data reduction (Sect.\\,3), the calibration (Sect.\\,4), and the procedures for selecting nova candidates (Sect.\\,5). The results are presented in Sect.\\,6 and discussed in Sect.\\,7. Finally, a summary is given in Sect.\\,8.  \\\\ ",
        "conclusions": "\\label{sec:discuss} Before we discuss individual nova candidates, we want to emphasise that we searched explicitly for objects that were visible only on a short time span. Consequently this means that the seldom novae in globular clusters (like in Quimby et al. \\cite{Quimby07}) are not recognised by this procedure if those clusters are bright enough to be detected. A search for such novae could be made in the context of an analysis of long-term variable objects, which we plan to conduct in the future.  In addition, we want to point out that the results presented in the present paper contain a few modifications compared to the preliminary results published by Henze et al. (\\cite{Henze06}) where 19 new historical nova candidates in M\\,31 have been announced. Meanwhile, we identified one of these objects (nova 32, discussed below) with an already known nova and found four additional new nova candidates (novae 35, 36, 42, and 43).  \\subsection{Rejected novae} We have re-investigated all nova candidates previously found on Tautenburg plates in considerable detail. Most of them were confirmed as nova candidates and are listed in Tables\\,\\ref{tab:tauno} or \\ref{tab:confno}. However, for two objects we found that they have very likely been misclassified as novae, whatever they are instead. These objects will be discussed briefly below.  \\subsubsection*{Nova M31N 1966-10a = Nova B\\\"orngen 21} This nova candidate is described by B\\\"orngen (1968) as ``nova like, weakly confirmed\" and was detected on only three plates. We find that there is a very faint object visible on all of these three plates; however, we noticed position differences of up to $4\\,\\farcs3$ between the three objects and therefore do not consider this object as a nova. (There is also no pattern recognisable in this position's differences that can be interpreted as caused by a high-proper motion star.)  \\subsubsection*{Nova M31N 1967-08c = Nova Meinunger S10753} This object was classified as a nova by Meinunger (\\cite{Meinunger75}) because of a significant ($\\geq 1$ mag) outburst that is confirmed in our data. However, there is a very faint object visible at that position on many other plates of different epochs. Moreover, there is clearly an object visible on POSS II plates. We checked the possibility of this object being a nova in a globular cluster (similar to the one reported by Quimby et al. \\cite{Quimby07}), but there is no entry in the catalogue of M\\,31 globular clusters (Galetti et al. \\cite{Galleti06}) at this position. Consequently we consider this object not to be a nova in M\\,31.    \\subsection{Remarks on individual nova candidates} \\label{sec:dis_rem} Here we briefly present remarks on individual objects. We discuss a candidate recurrent nova and a rejected candidate recurrent nova, give details on two novae with large position corrections, and discuss three objects that we detected much earlier than the previously known outbursts. Furthermore, we note a nova that might belong to M\\,32, discuss three novae with lower position accuracy, and give remarks on six very fast or very bright novae and on one very slow nova.  \\subsubsection*{Nova 4 = M31N 1963-08a = Nova Moffat 2} Fast nova, $t_2 \\sim 20$ days.  \\subsubsection*{Nova 5 = M31N 1963-09b = Nova Moffat 1} Moffat (\\cite{Moffat67}) gave wrong coordinates for this object. In fact, the nova Moffat 1 coincides with nova Rosino 47 (Rosino \\cite{Rosino73}). This is clearly seen on the plates where the nova candidate originally was found by Moffat. The date of outburst is 2 days earlier than given by Rosino and is very well constrained (1 day). This nova is very bright (15.4 mag) and very fast ($t_2 \\sim 3$ days).  \\subsubsection*{Nova 7 = M31N 1963-10b = Nova B\\\"orngen 8} Fast nova, $t_2 \\sim 13$ days.  \\subsubsection*{Nova 17 = M31N 1967-10b = Nova Meinunger S10731} Very fast nova, $t_2 \\sim 8$ days.  \\subsubsection*{Nova 18 = M31N 1967-10c = Nova B\\\"orngen 26} Brightest nova found in this work (15.2 mag).  \\subsubsection*{Nova 24 = M31N 1973-10b = Nova ShAl 15} Very fast nova, $t_2 \\sim 8$ days.  \\subsubsection*{Nova 28 = M31N 1985-09g = Nova Rosino 139} We detected this nova 21 days before the outburst date reported in Rosino et al. (\\cite{Rosino89}). But Rosino also noted a \\textit{possible} detection of this nova 2 days before our detection date. Because of the good agreement with our data, let us suppose that this possible detection indeed was real and that Rosino detected the outburst of the nova at nearly the same date as we did.  \\subsubsection*{Nova 31 = M31N 1992-12b = Nova Shafter \\& Irby 1992-01} This nova was found on Tautenburg plates even 84 days before the outburst reported by Shafter \\& Irby (\\cite{Shafter01}). First we note that Shafter \\& Irby found this nova at the beginning of an observation period with previous observations one year ago. Furthermore, Shafter \\& Irby conducted their survey in the H$\\alpha$ band, in which the brightness of novae declines more slowly (Shafter \\& Irby \\cite{Shafter01}). Therefore, we assume that this nova is a slow one and still visible in H$\\alpha$ 100 days after outburst. Unfortunately there were no M\\,31 plates taken in Tautenburg at this time (see Table \\ref{tab:lightcurves9}).  \\subsubsection*{Nova 32 = M31N 1996-08d = Nova ShAl 48} Despite a position difference of $46\\,\\farcs1$, this nova candidate is identified with nova ShAl 48, since the dates of outburst coincide very well and also the finding chart given by Sharov \\& Alksnis (1997) evidently shows the same object. In agreement with Alksnis (2006), we assume that the large position difference is due to a typing error in Sharov \\& Alksnis (\\cite{Sharov97}).  \\subsubsection*{Nova 40 = M31N 1972-01b} Very fast nova, $t_2 \\sim 7$ days.  \\subsubsection*{Nova 43 = M31N 1975-09a} This nova candidate was found in our search for recurrent novae using correlations of single-detection objects with the PHS07 catalogue. Originally we identified this object with nova M31N 1999-01a found by Rector et al. (\\cite{Rector99}), because the distance of the two objects is only $3\\,\\farcs4$. However, a finding chart provided by Rector (\\cite{Rector07}) clearly shows that both nova candidates have different positions. However, we confirmed this object as a new nova candidate using additional plates.  \\subsubsection*{Nova 44 = M31N 1975-11a} The position of this newly found candidate correlates (distance $\\sim$ $1\\,\\farcs0$) with nova M31N 1945-09c detected by Baade \\& Arp (\\cite{Baade64}). Hence we classify this nova as recurrent, although the position accuracy of $0\\,\\farcs3$ given by Baade \\& Arp may be too optimistic, and there are no finding charts available for this object to confirm this assumption. Unfortunately Baade \\& Arp gave no date of outburst for their nova 32/Table 4. For naming purposes, PHS07 set the date of outburst to 1 September 1945 but the actual date can be between 1945 and 1949. Therefore, the recurrence time is estimated to 26 - 30 years. As for nova 41, we classify this object as a new nova because the particular outburst in 1975 was not known before.  \\subsubsection*{Nova 57 = M31N 1962-08a = Nova Meinunger S10730} This candidate is near to the centre of M\\,32 (distance $\\sim$ $1\\,\\farcm8$). Therefore, we note that this nova might belong to the galaxy M\\,32.  \\subsubsection*{Nova 58 = M31N 1962-07a = Nova B\\\"orngen 2} After the detected outburst in September 1962, this object became visible for a year. In agreement with Rosino (\\cite{Rosino64}), we note that this is a very slow nova ($t_2 \\sim 350$ days).  \\subsubsection*{Nova 62 = M31N 1963-10a = Nova B\\\"orngen 7} This object was not detected in the automatic search routine due to the proximity to a bright star. For the same reason, no photometry is given here. We confirmed this nova candidate by eye and computed the coordinates separately. Therefore, the position accuracy is only $\\sim 1\\arcsec$.  \\subsubsection*{Nova 64 = M31N 1964-08b = Nova B\\\"orngen 12} This object was not detected in the automatic search routine because the plates it is detected on are of lower photometric quality and have therefore not been included in the 306 selected plates of the present study. We confirmed this nova candidate by eye and computed the coordinates separately. Therefore, the position accuracy is only $\\sim 1\\,\\farcs5$.  \\subsubsection*{Nova 68 = M31N 1966-09c = Nova B\\\"orngen 18} This object was detected only on broken plates. Therefore, the position accuracy is only $\\sim 0\\,\\farcs8$  \\subsubsection*{Nova 82 = M31N 1985-09b = Nova Ciardullo 20} We detected this nova 19 days before the outburst reported by Ciardullo et al. (\\cite{Ciardullo87}). Similar to Nova 31, the date of outburst reported by Ciardullo et al. is at the beginning of an observation period of their H$\\alpha$ survey. Therefore, we interpret this object as a moderately fast nova that is visible much longer in the H$\\alpha$ band."
    },
    "0710/0710.3969_arXiv.txt": {
        "abstract": "We present a model to explain the mass segregation and shallow mass functions observed in the central parts of dense and young starburst stellar clusters. The model assumes that the initial pre-stellar cores mass function resulting from the turbulent fragmentation of the proto-cluster cloud is significantly altered by the cores coalescence before they collapse to form stars. With appropriate, yet realistic parameters, this model based on the competition between cores coalescence and collapse reproduces the mass spectra of the well studied Arches cluster. Namely, the slopes at the intermediate  and high mass ends are reproduced, as well as the peculiar bump observed at $6$ M$_{\\odot}$. This coalescence-collapse process occurs on short timescale of the order of one fourth the free fall time of the proto-cluster cloud (i.e., a few $10^{4}$ years), suggesting that mass segregation in Arches and similar clusters is primordial. The best fitting model implies the total mass of the Arches cluster is $1.45 \\times 10^{5}$ M$_{\\odot}$, which is slightly higher than the often quoted, but completeness affected, observational value of a few $10^{4}$ M$_{\\odot}$. The derived star formation efficiency is $\\sim 30$ percent which implies that the Arches cluster is likely to be gravitationally bound. ",
        "introduction": "Mass segregation is observed in many young stellar clusters, both compact and sparse, with the most massive stars being preferentially located in their central parts (e.g., Pandey et al. 1992; Malumuth \\& Heap 1994;  Brandl et al. 1996; Hillenbrand \\& Hartmann 1998; Fisher et al. 1998; Figer et al. 1999b, Sagar et al. 2001; Stolte 2002; Le Duigou \\& Kn\\\"{o}delseder 2002; Sirianni et al. 2002; Gouliermis et al. 2004; Sharma et al. 2007). Among young clusters, the class of very dense clusters, harboring large numbers of massive stars and known as starburst clusters such as the Arches cluster, NGC 3603, the Quintuplet cluster, Westerlund 1 and 2 and R 136 offers a very challenging testbed for the theories of massive star formation not only because of the extreme mass segregation observed in their central parts, but also due to the fact that their radially integrated stellar mass function is observed to be top heavy (Moffat et al. 1994; Figer et al. 1999; Stolte et al. 2002; Stolte et al. 2005; Stolte et al. 2006; Kim et al. 2006; Harayama et al. 2007; Brandner et al. this volume). Similar trends are observed for the super star cluster NGC 1705-1 (Sternberg 1998).  In particular, the structure of the young (age $\\sim 1-2$ Myrs), and dense ($\\sim 10^{5}$ stars/pc$^{3}$) Arches cluster which is located at a projected distance of only 25 pc from the Galactic center has been studied extensively over the past few years. Its mass function has been determined using a variety of space borne and ground based instruments in the near infrared such as the {\\it Hubble Space Telescope (HST)} NICMOS camera (Figer et al. 1999a), the Hokupa AO system on Gemini (Stolte et al. 2002), the NAOS/CONICA camera on the VLT (Stolte et al. 2005), and the NIRC2 instrument on Keck (Kim et al. 2006). Star counting in the inner annulus of the cluster (i.e., between the center and 1 time the core radius which is $\\sim 0.2$ pc) is affected by incompleteness effects due to severe crowding, and the outer parts might still be contaminated by field stars. Thus, an interesting and well resolved area to test theoretical models is the second annulus. The converging results of these observations yield a very flat mass function with $\\alpha \\sim 1$ ($dN/d M=M^{-\\alpha}$; The Salpeter value is $\\alpha=2.35$, Salpeter 1955) in the central annulus of the cluster, a value of $\\alpha \\sim 1.7-1.9$ in the second annulus (i.e., between 1 and 2 times the core radius), and a nearly Salpeter-like or slightly steeper exponent in the outer areas.  The mass segregation and shallow IMFs observed in starburst clusters such as Arches have been interpreted as the result of dynamical processes (Kim et al. 2006, Portegies-Zwart et al. 2007). Although dynamical processes will inevitably play a role in enhancing the fraction of massive stars in the central regions of a cluster, two points remain unclear: The first one is that the primordial IMF might be shallow in nature which leaves little room for the subsequent effects induced by dynamical friction. The second issue is that for young cluster as the Arches, it is not clear whether dynamical processes have time to play a significant role (see Freitag et al. this volume). Moreover, models based on mass segregation by dynamical processes do not reproduce, to date, some features of the Arches cluster mass function such as the peculiar bump at $\\sim 6$ M$_{\\odot}$.  In this work, we discuss a model to explain the observed mass function of the Arches and similar starburst clusters based on the coalescence of pre-stellar cores (PSCs) in the proto-cluster cloud and their subsequent collapse into stars. ",
        "conclusions": "The origin of the shallow mass functions and mass segregation observed in young starburst stellar clusters is explained by means of a model based on the efficient coalescence of pre-stellar cores and their subsequent collapse to form stars. The coalescence of cores causes a strong flattening of the turbulence generated, primordial core mass function. Once the reservoir of intermediate mass cores is depleted, and the radii of cores are significantly reduced by gravitational contraction, coalescence proceeds at a much slower pace until all cores in the whole mass range collapse to form stars. The entire process is fast, lasting for a fraction ($\\sim 0.25$) of the proto-cluster cloud free-fall time, i.e., a few $10^{4}$ years. When applied to the Arches cluster, the model is able to reproduce the observed mass slopes of $\\alpha \\sim 2$ and $\\sim 1.7$ at the intermediate and high mass ranges, respectively, as well as the peculiar bump at $\\sim 6$ M$_{\\odot}$. Using this model, we estimate a total mass of Arches of $1.45 \\times 10^{5}$ M$_{\\odot}$, which is slightly higher than the often quoted, but completeness affected, observational estimates of a few $10^{4}$ M$_{\\odot}$. The star formation efficiency in this model is $\\sim 30$ percent which implies that the Arches cluster is gravitationally bound according to the results of Geyer \\& Burkert (2001)."
    },
    "0710/0710.3228_arXiv.txt": {
        "abstract": "We report on the discovery of a surprising observed correlation between the slope of the low-mass stellar global mass function (GMF) of globular clusters (GCs) and their central concentration parameter $c=\\log(r_t/r_c)$, i.e. the logarithmic ratio of tidal and core radii. This result is based on the analysis of a sample of twenty Galactic GCs, with solid GMF measurements from deep HST or VLT data, representative of the entire population of Milky Way GCs. While all high-concentration clusters in the sample have a steep GMF, low-concentration clusters tend to have a flatter GMF implying that they have lost many stars via evaporation or tidal stripping. No GCs are found with a flat GMF and high central concentration. This finding appears counter-intuitive, since the same two-body relaxation mechanism that causes stars to evaporate and the cluster to eventually dissolve should also lead to higher central density and possibly core-collapse. Therefore, severely depleted GCs should be in a post core-collapse state, contrary to what is suggested by their low concentration. Several hypotheses can be put forth to explain the observed trend, none of which however seems completely satisfactory. It is likely that GCs with a flat GMF have a much denser and smaller core than suggested by their surface brightness profile and may well be undergoing collapse at present. It is, therefore, likely that the number of post core-collapse clusters in the Galaxy is much larger than thought so far. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.3534_arXiv.txt": {
        "abstract": "{} {We investigate the influence of high--energy stellar radiation at close--in orbits on atmospheric mass loss during stellar evolution of a G--type star.} {High--energy stellar luminosity varies over a wide range for G field stars. The temporal evolution of the distribution of stellar X--ray luminosity and its influence on the evolution of close--in exoplanets is investigated. X--ray luminosity distributions from the Pleiades, the Hyades and the field are used to derive a scaling law for the evolution of the stellar X--ray luminosity distribution. A modified energy--limited escape approach is used to calculate atmospheric mass loss for a broad range of planetary parameters.} {We show that the evolution of close--in exoplanets strongly depends on the detailed X--ray luminosity history of their host stars, which varies over several orders--of--magnitude for G stars. Stars located at the high--energy tail of the luminosity distribution can evaporate most of its planets within 0.5 AU, while for a moderate luminosity a significant fraction of planets can survive. We show the change on an initial planetary mass distribution caused by atmospheric escape.} {} ",
        "introduction": "Since the discovery of the first close--in Hot Jupiter in 1995 (Mayor and Queloz, 1995) the investigation of the structure and the temporal evolution of their highly irradiated atmospheres was a main topic for both, modelers and observers. In 2002, Charbonneau et al. (2002) reported the first detection of an atmosphere around HD209458b by observing absorption in the NaI D lines during transits. Soon after that Vidal-Madjar et al. (2003, 2004) reported the detection of several other species and pointed out that the planet is losing mass at a rate of more than 10$^{10}$ g/s, as indicated by an expanded atmosphere likely due to heating by incoming stellar radiation. Recently, a layer of excited hydrogen atoms with a temperature of about 5000 K slightly above the visual radius of the planet was reported (Ballester et al., 2007), further strengthening the hypothesis of a strongly heated atmosphere. Shortly after the discovery of Hot Jupiters, theoretical models were established to verify the evaporation conditions of these close--in planets (Guillot et al., 1996). Lammer et al. (2003) used an energy--limited escape approach for investigating the atmospheric escape from close--in exoplanets. After that, several hydrodynamic models were established in order to analyze the atmospheric conditions in more detail (Yelle, 2004, 2006; Tian et al., 2005; Garcia Munoz, 2007). Cecchi--Pestellini et al. (2006) developed an accurate heat transfer model to determine the heating of planetary atmospheres due to X--rays. Mass loss calculations were also included in different models for giant planet evolution (Lecavelier des Etangs et al., 2004; Baraffe et al., 2004, 2005, Hubbard et al., 2006, 2007). Recently, Lecavelier des Etangs (2007) presented a diagram representation where the influence of atmospheric evaporation on the evolution of the exoplanets can be estimated for a wide range of planetary parameters and for different spectral types of host stars.\\\\ However, in all these works the present--day solar luminosity was scaled to the considered orbital distances, or, in the evolutionary models, the scaling law derived by Guinan and Ribas (2002) from the \\textit{Sun in Time} program was used. This program was based on a small sample of solar--type stars, allowing to determine an average scaling law for the temporal evolution of the stellar radiation for a range of wavelength from 1--1200 \\AA. However, in reality, G stars show a broad distribution of the luminosity which varies over a few orders--of--magnitude. This distribution can be observed only for the very extreme UV and X--ray (1--200 \\AA) (e.g., Preibisch and Feigelson, 2005) because of interstellar absorption and lack of sensitive instruments. In this paper, we will focus on the influence of radiation in this wavelength range, which is justified since Cecchi--Pestellini et al. (2006) showed that X--ray and very extreme UV radiation has significant influence on planetary atmospheres. Also the X--ray luminosity of clusters with a given age like the Pleiades (Micela, 2001) and the Hyades (Stern et al., 1995) is known, therefore the temporal evolution of the population of G stars can be derived. Using this X--ray luminosity distribution as an input, we apply a modified energy--limited approach formula (Erkaev et al., 2007) to study the influence of atmospheric loss for close--in exoplanets around solar--like G--stars and discuss the expected deviations from an initially known planetary mass distribution due to atmospheric loss. ",
        "conclusions": "We studied the influence of the stellar luminosity distribution on atmospheric escape of exoplanets at orbits less than 0.1 AU. We showed that a significant amount of planets can be evaporated over time scales of 4 Gyr and that the final mass distribution is different from the initial one. For a Jupiter--mass planet with a density of 0.4 g/cm$^{3}$, about 5 \\% of the planets are lost after 4 Gyr at an orbit of 0.02 AU. For Neptune--mass planets this number increases to about 60 \\% for a density of 0.8 g/cm$^{3}$. We also show that for close--in orbits a large number of planets can be eroded to Super--Earth and/or may lose dense hydrogen envelopes. Further, we present the resulting mass distribution of an initially flat mass distribution of planets between 0.2 M$_{nep}$ and 10 M$_{jup}$. For close--in orbits about 32 \\% of the initial planets are lost and the distribution is shifted to smaller masses.\\\\ Our work is a first step to understand the evolution of planetary mass distribution due to the stellar activity evolution. Our approach is subject to several assumptions that will be verified with future observations and/or modelling. First of all we assume a constant density in time and a heating function of 100 \\%. Better knowledge of the densities of exoplanets will be achieved by further modelling supported by observational evidence, while the latter assumption needs to be verified with detailed radiative transfer calculations. The scaling law of the L$_{X}$ evolution in Eq. \\ref{scaling} is derived from a few points. This potential source of uncertainty is however mitigated by the fact that most of the effects occur in the first Gyr (see Fig. \\ref{age}) where the X--ray evolution is better known. The UV flux and evolution should be added in the future. We plan to extend our work to lower mass stars where the X--ray effects may influence the habitability zone."
    },
    "0710/0710.1531_arXiv.txt": {
        "abstract": "We present far-ultraviolet observations of the Antlia supernova remnant obtained with Far-ultraviolet IMaging Spectrograph (FIMS, also called SPEAR). The strongest lines observed are \\Civ{} \\dl1548,1551 and \\Ciii{} $\\lambda$977. The \\Civ{} emission of this mixed-morphology supernova remnant shows a clumpy distribution, and the line intensity is nearly constant with radius. The \\Ciii{} $\\lambda$977 line, though too weak to be mapped over the whole remnant, is shown to vary radially. The line intensity peaks at about half the radius, and drops at the edge of the remnant. Both the clumpy distribution of \\Civ{} and the rise in the \\Civ{} to \\Ciii{} ratio towards the edge suggest that central emission is from evaporating cloudlets rather than thermal conduction in a more uniform, dense medium. ",
        "introduction": "\\label{intro} The Antlia supernova remnant (SNR) is a nearby ($d_A\\simeq60-340$ pc), large ($\\sim24\\arcdeg$ in diameter), old ($t_A\\sim1$ Myr) SNR recently discovered by \\cite{McCullough(2002)ApJ_576_L41}, centered on  ($l=276\\fdg52$, $b=+19\\fdg05$). It shows a center-filled X-ray morphology, which is partially enclosed by prominent \\Ha{} and weak radio features (cf.~Fig.~\\ref{antlia-mm}). The remnant abuts the Gum Nebula to its southwest, and the boundary lies along the radio ridge seen in Fig.~\\ref{antlia-mm}. Based on the X-ray and radio morphologies, we classify the Antlia SNR as a mixed-morphology (MM) SNR \\citep{Rho(1998)ApJ_503_L167}. The X-ray is clearly center-filled, and though the radio shell is weak, there is no central radio emission. Furthermore there is no known pulsar near the center of the Antlia SNR so it is unlikely to be pulsar-driven nebula.  To explain the physical characteristics of MM SNRs, in particular their center-filled X-ray morphologies, two models have been invoked. Both focus on how the interior density, and thereby the emissivity may be increased. In the first model \\citep[hereafter ``cloudlet'' model]{White(1991)ApJ_373_543}, thermal evaporation of small clouds overrun by the supernova shock and engulfed by the hot gas bubble increases the density of the interior gas. In the second model (hereafter ``dense medium'' model), applied to the SNR W44 by \\cite{Cox(1999)ApJ_524_179}, thermal conduction in a dense medium alone increases the interior density of the remnant as heat is transported outward from the center. These two models assume different distributions of the denser component that would show radiative-cooling features as the remnant evolves.  Here, we report the detection of such cooling features in the far-ultraviolet (FUV) wavelengths from the Antlia SNR. We observed the Antlia SNR with the Far-ultraviolet IMaging Spectrograph (FIMS, also called SPEAR), and detected several emission lines of which \\Civ{} \\dl1548,1551 and \\Ciii{} $\\lambda$977 are the strongest. In the following sections we describe the observations, present the analysis and results, and discuss what we have learnt about MM SNRs. ",
        "conclusions": "We have presented and analyzed FUV observations of the Antlia SNR obtained with FIMS/SPEAR. The Antlia SNR was identified as an MM SNR based on its X-ray and radio morphology, and its boundary was assumed as a 24$\\arcdeg$-diameter circle at the center. The clumpy \\Civ{} map without an edge-concentration and the radial variation of I(\\Civ)/I(\\Ciii) within the remnant suggest that the ``cloudlet'' model, rather than the ``dense medium'' model, is plausible for the formation of the Antlia SNR's center-filled X-ray morphology.  The Antlia SNR is one of about 25 MM SNRs in the Galaxy, and perhaps the only one that can be studied in ultraviolet---all the others are suffering high extinction. Hence, the results presented here are unique, showing that the ambient cloudlets may be an important component for energy exchange between MM SNRs and their surroundings. However, it is clearly impossible to generalize from this one example to all objects in the class. The ``dense medium'' model could be suitable in other situations where the ambient interstellar medium is different from that around the Antlia SNR.  We note that, at an angular resolution of $2\\fdg5$, much of the sub-structure in the \\Civ{} map is almost certainly hidden. Thus, an FUV study of the Antlia SNR at much higher angular resolution (if and when that becomes possible) would allow the detailed structure of the radiatively cooling gas to be mapped and its relationship with the emission at other wavelengths to be characterized more precisely."
    },
    "0710/0710.5312_arXiv.txt": {
        "abstract": "We discuss the cross section formula both for massless and massive neutrinos on stable and radioactive nuclei. The latter could be of interest for the detection of cosmological neutrinos whose observation is one of the main challenges of modern cosmology. We analyze the signal to background ratio as a function of the ratio $m_{\\nu}/\\Delta$, i.e. the neutrino mass over the detector resolution and show that an energy resolution $\\Delta \\le 0.5$~eV would be required for sub-eV neutrino masses, independently of the gravitational neutrino clustering. Finally we mention the non-resonant character of neutrino capture on radioactive nuclei. ",
        "introduction": "Modern big-bang cosmology firmly predicts the existence of a relic neutrino background, and relates its temperature to the temperature of the background microwave radiation \\begin{equation}\\label{eq:1} T_{\\nu}/T_{\\gamma} = (4/11)^{1/3} ~, \\end{equation} see, e.g. the basic texts \\cite{Peebles,Wein_c}\\footnote{ Corrections to the above ratio caused by the incomplete neutrino decoupling are only at a few percent level \\cite{mangano}.}.  Verifying the existence of the relic neutrino sea represents one of the main challenges of modern cosmology.  Clearly, in contrast to the study of the  background microwave radiation that has a long history and has reached an unprecendented accuracy (see, e.g. the latest results in \\cite{Spergel:2006hy}), detection of  relic neutrinos remains an unfulfilled dream. Various strategies have been proposed so far, based on laboratories searches \\cite{orpher,Lewis:1979mu,Cabibbo:1982bb,Langacker:1982ih,Stodolsky:1974aq,Hagmann:1998nz,Smith:1985fz} and astrophysical observations \\cite{Weiler:1982qy,Weiler:1983xx,Eberle:2004ua,Weiler:1997sh,Fargion:1997ft}, such as absorption dips in the flux of Ultra High Energy neutrinos (for a review see e.g. \\cite{Gelmini:2004hg}). As far as their detection in laboratory experiments is concerned, one needs to overcome two main obstacles: the low cross section characteristic of weak interactions and the low energy of relic neutrinos. The second obstacle  can be overcome if the corresponding detection reactions have vanishing thresholds. Therefore we discuss here the possibility of detecting the relic neutrinos by the charged current reactions using radioactive unstable nuclei as targets.  The paper is organized as follows: In the next section we derive expressions for the charged current neutrino induced reaction cross sections involving nonrelativistic neutrinos. We show that such cross sections, when the corresponding reaction has a vanishing threshold, scale with $c/v_{\\nu}$ so that the number of events converges to a constant for $v_{\\nu} \\rightarrow 0$. In the following section we discuss the possibility to use a tritium target to detect the cosmological $\\nu_e$. We show that the main challenge is the separation of the produced electrons with energies just above the endpoint of the $\\beta$ spectrum from the overwhelming flux of the electrons from the tritium $\\beta$ decay that extends just below the 18.6 keV endpoint.  We also discuss the possibility of a resonance enhancement of reactions involving cosmological neutrinos. We show that  the charged current reactions, included the radiative ones, do not have a resonance character. In the conclusion we summarize our findings and stress the need for an extreme energy resolution if sensitivity to detect sub-eV mass relic neutrinos should be reached. ",
        "conclusions": "We have shown that the charged current cross section of nonrelativistic neutrinos scales like $1/v_{\\nu}$ in agreement with the generic behavior of the cross sections for slow particles.  That means, in particular, that the charged current reactions with vanishing threshold  of nonrelativistic neutrinos have a rate that converges to a finite  value as $v_{\\nu} \\rightarrow 0$. With that in mind we evaluate the cross section and reaction rate of the relic $\\bar{\\nu}_e$ on tritium nuclei. If a modest gravitational clustering enhancement of the relic neutrino number density is present, the number of events is large enough that it might be potentially observable. The more important issue is the elimination of the overwhelming background of the electrons from tritium $\\beta$ decay. We show that  signal/background ratio of the order of unity can be achieved only if the neutrino mass $m_{\\nu}$ exceeds the characteristic experimental resolution width $\\Delta$ by a factor of two or more. Thus  an energy resolution $\\Delta \\le 0.5$~eV would be required in order to be possible, even in an ideal experiment, to detect relic neutrinos with sub-eV neutrino masses, independently of the gravitational neutrino clustering.  \\bigskip"
    },
    "0710/0710.0701_arXiv.txt": {
        "abstract": "s{ We present the first constraints on pure-gravity sector Standard-Model Extension (SME) parameters using Lunar Laser Ranging (LLR).  LLR measures the round trip travel time of light between the Earth and the Moon.  With 34+ years of LLR data, we have constrained six independent linear combinations of SME parameters at the level of $10^{-6}$ to $10^{-11}$.  There is no evidence for Lorentz violation in the LLR dataset. } ",
        "introduction": "Two of us (C.W.S. and J.B.R.B.) are members of the Apache Point Observatory Lunar Laser-ranging Operation (APOLLO), a next-generation LLR facility, capable of millimeter-precision lunar range measurements (see the article by T.W. Murphy in these proceedings).  The APOLLO project was motivated by the realization that an order-of-magnitude improvement in fundamental physics constraints (e.g. equivalence principle, gravitomagnetism, gravitational $1/r^2$ law and $\\dot{G}$, to name a few) could be achieved with straightforward improvements to the standard LLR apparatus.  With the recent description of the pure-gravity sector of the SME,\\cite{BK2006} we learned that LLR can also provide incisive constraints on Lorentz Violation.  The predicted LLR observable under Lorentz Violation is a periodic perturbation to the Earth-Moon range with the leading order effects occuring at four distinct frequencies: $2\\omega$, $\\omega$, $2\\omega-\\omega_0$ and $\\Omega_{\\earth}$.  Here $\\omega$ is the lunar orbital (sidereal) frequency, $\\omega_0$ is the anomalistic lunar orbital frequency and $\\Omega_{\\earth}$ is the mean Earth orbital (sidereal) frequency.  Although the APOLLO program is still in the data collection phase, there are more than three decades of freely-available archival LLR data on a public archive.\\cite{ILRS} In this article, I present the SME parameter constraints that result from our analysis of archival LLR data.  These are the first LLR-based constraints on pure-gravity SME parameters. ",
        "conclusions": "We have analyzed 34+ years of LLR data and have derived constraints on six SME parameters combinations (see Table \\ref{tab:smeconstraints}). We find no deviation from Lorentz Symmetry at the $10^{-6} - 10^{-11}$ level.  This work provides the first LLR-based constraints of SME parameters.  There are several ways in which these constraints could be improved. First of all, by incorporating auxiliary solar system data (e.g. planetary radar ranging) model parameter correlations can be reduced, and $F$ decreased.  This would allow for tighter constraints on the SME parameters using the same LLR dataset.  In addition, APOLLO data, which is about 10 times more precise than the archival data, will soon be ready for analysis.  With this improved dataset, we will further tighten the SME parameter constraints.  \\begin{table}[t] \\tbl{SME parameter estimates and their realistic (scaled) uncertainties ($F\\sigma$) with $F=20$.  } {\\footnotesize \\begin{tabular}{cr} \\hline {} &{}\\\\[-1.5ex] Parameter & Estimate \\\\[1ex] \\hline {} &{}\\\\[-1.5ex] $\\bar{s}^{11}-\\bar{s}^{22}$  & $\\left( 1.3 \\pm 0.9 \\right) \\times 10^{-10}$  \\\\[1ex] $\\bar{s}^{12}$               & $\\left( 6.9 \\pm 4.5 \\right) \\times 10^{-11}$  \\\\[1ex] $\\bar{s}^{02}$               & $\\left(-5.2 \\pm 4.8 \\right) \\times 10^{-07}$  \\\\[1ex] $\\bar{s}^{01}$               & $\\left(-0.8 \\pm 1.1 \\right) \\times 10^{-06}$  \\\\[1ex] $\\bar{s}_{\\Omega_\\earth c}$  & $\\left( 0.2 \\pm 3.9 \\right) \\times 10^{-07}$  \\\\[1ex] $\\bar{s}_{\\Omega_\\earth s}$  & $\\left(-1.3 \\pm 4.1 \\right) \\times 10^{-07}$  \\\\[1ex] \\hline \\end{tabular}\\label{tab:smeconstraints} } \\vspace*{-13pt} \\end{table}"
    },
    "0710/0710.3560.txt": {
        "abstract": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% {} {We examine the ionization, abundances, and differential dust depletion of metals, the kinematic structure, and the physical conditions in the molecular hydrogen-bearing sub-damped Ly~$\\alpha$ system toward HE~0515-4414.} {We used the STIS and VLT UVES spectrographs to obtain high-resolution recordings of the damped Ly~$\\alpha$ profile and numerous associated metal lines. Observed element abundances are corrected with respect to dust depletion effects.} {The sub-damped Ly~$\\alpha$ absorber at redshift $z=1.15$ is unusual in several aspects. The velocity interval of associated metal lines extends for $700$~km\\,s$^{-1}$. In addition, saturated \\ion{H}{i} absorption is detected in the blue damping wing of the $N_\\mathrm{\\ion{H}{i}}=8\\times10^{19}$~cm$^{-2}$ main component. The column density ratios of associated \\ion{Al}{ii}, \\ion{Al}{iii}, and \\ion{Fe}{ii} lines indicate that the absorbing material is ionized. 19 of in total 31 detected metal line components are formed within peripheral \\ion{H}{ii} regions, while only 12 components are associated with the predominantly neutral main absorber. The bimodal velocity distribution of metal line components suggests two interacting absorbers. For the main absorber the observed abundance ratios of refractory elements to Zn range from Galactic warm disk $[\\mathrm{Si}/\\mathrm{Zn}]_\\mathrm{g}=-0.40\\pm0.06$, $[\\mathrm{Fe}/\\mathrm{Zn}]_\\mathrm{g}=-1.10\\pm0.05$ to halo-like and essentially undepleted patterns. The dust-corrected metal abundances indicate a nucleosynthetic odd-even effect and might imply an anomalous depletion of Si relative to Fe for two components, but otherwise do correspond to solar ratios. The intrinsic average metallicity is almost solar $[\\mathrm{Fe}/\\mathrm{H}]_\\mathrm{m}=-0.08\\pm0.19$, whereas the uncorrected average is $[\\mathrm{Zn}/\\mathrm{H}]_\\mathrm{g}=-0.38\\pm0.04$. The ion abundances in the periphery conform with solar element composition.} {The detection of \\ion{H}{ii} as well as the large variation in dust depletion for this sight line raises the question whether in future studies of damped Ly~$\\alpha$ systems ionization and depletion effects have to be considered in further detail. Ionization effects, for instance, may pretend an enrichment of $\\alpha$ elements. An empirical recipe for detecting \\ion{H}{ii} regions is provided.} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  The study of QSO absorption lines provides vital information on the nucleosynthetic history of the universe by complementing the compositional analysis of stars and interstellar space in local galaxies with element abundances at higher redshift. In particular, interests are focused on extragalactic structures termed damped Ly $\\alpha$ (DLA) systems, essentially comprised of neutral hydrogen with column densities $N_\\ion{H}{i}\\ge2\\times10^{20}$ atoms cm$^{-2}$ \\citep[for a review see][]{WolfeGP_2005}. Absorbers in the sub-DLA range with column densities $N_\\ion{H}{i}\\ge10^{19}$ atoms cm$^{-2}$ might be mainly neutral when the ionizing background is reduced \\citep{PerouxDKMD_2002,PerouxDDKM_2003}. The aim of these examinations is to establish accurate element abundances for the aggregations of neutral gas that are examples of interstellar environments in the high-redshift universe. Since the measurement of metal column densities is straightforward, the only problem is their correct interpretation.  The true nature of DLA systems is unknown and the underlying population, being constituted of hierarchical structures with different morphologies, chemical enrichment histories, and physical environments, is multifarious. The diversity is attested by the disparate values obtained for metal abundances at any given redshift.  The metallicity of DLA systems is not correlated with their column density, however, there is an upper bound for distribution of column densities versus metallicity \\citep{BoisseLBD_1998}. Though there are high column density DLA absorbers with high metallicity in the foreground of the star forming hosts of gamma-ray bursts \\citep{Watson_2005}, similar absorbers are not detected toward QSOs. The cosmic mean metallicity of DLA absorbers increases with cosmic time \\citep{ProchaskaGWCD_2003,KulkarniFL_2005,RaoPHW_2005}, but is an order of magnitude lower than predicted by cosmic star formation history \\citep[see the discussion of the missing metals problem by][]{WolfeGP_2005}. The solution to this problem is a matter of debate. Conclusive evidence of enriched material ejected from DLA absorbers into the intergalactic medium or of active star formation restricted to compact regions is missing. The latter possibility is closely linked to the physical properties of the interstellar medium and its molecular content \\citep{WolfePG_2003,WolfeGP_2003}. Molecular gas is uncommon in DLA absorbers. If found, the fraction of molecular hydrogen, usually between $10^{-6}$ and $10^{-2}$, is not correlated with the column density of atomic hydrogen \\citep{LedouxPS_2003}. %\\cbstart However, \\citet{PetitjeanLNS_2006} have demonstrated that the presence of molecular hydrogen at high redshift is strongly correlated with the metallicity. %\\cbend  Since the spectroscopic analysis is restricted to the gaseous phase of the absorbing medium, observed element abundances are potentially distorted by dust removing atoms in varying amounts, depending on their affinity to the solid state. In particular high-metallicity and molecule-bearing absorbers are affected by dust \\citep{PetitjeanSL_2002,LedouxPS_2003}. Depletions are largely lower than in the Galactic halo, but increase with metallicity \\citep{VladiloG_2004}. In practice, the observed element abundances are corrected ad hoc, using Galactic interstellar depletion patterns as reference \\citep{VladiloG_2002a,VladiloG_2002b}. Another aspect of dust is the possibility that DLA absorbers may elude detection because the background QSOs are obscured \\citep{FallS_PeiY_1993}. The selection effects are complicated since obscurement is counteracted, but not compensated, by amplification due to gravitational lensing \\citep{SmetteCS_1997}. The effect of dust is subject of several studies \\citep{MurphyM_LiskeJ_2004,QuastRS_2004,AkermanEPS_2005,SmetteWL_2005,VladiloG_PerouxC_2005,WildHP_2005}. A further difficulty are ionization effects. Examples of DLA-associated metal line components formed within mainly ionized material are given by \\citet{ProchaskaHO_2002} and \\citet{DessaugesPD_2006}.  The column density distribution and kinematic structure of absorbers provide important constraints on hierarchical structuring \\citep[e.g.][]{CenOPW_2003,NagamineSH_2004} and immediate insight into the processes of galaxy formation \\citep{WolfeA_ProchaskaJ_2000a}. The extended multicomponent velocity structure and characteristic asymmetry of DLA-associated metal lines is consistent with galaxy formation models in hierarchic cold dark matter cosmologies, and reproducible by the hydrodynamical simulation of rotation, random motion, infall, and merging of irregular protogalactic clumps hosted by collapsed dark matter halos \\citep{HaehneltSR_1998}. The velocity structure of sub-DLA absorbers compares to that of the higher column density systems \\citep{PerouxDDKM_2003}, which is unexpected since semianalytic galaxy formation models \\citep{MallerPSP_2001,MallerPSP_2003} predict markedly different kinematic properties. The absorption velocity intervals of both sub-DLA and DLA absorbers typically extend for $100$~km\\,s$^{-1}$. More extended systems tend to higher metallicities and lower hydrogen column densities \\citep{WolfeA_ProchaskaJ_1998}. In particular the latter property is unexpected and difficult to interpret in terms of rotating disks models. The most extended systems, however, are probably due to interacting or merging galaxies \\citep{PetitjeanSL_2002,RichterLPB_2005}. %\\cbstart Strong observational evidence for a correlation between DLA metallicity and absorption profile velocity spread, which probably is the consequence of a mass-metalicity relation, has recently been provided by \\citet{LedouxPFMS_2006}. %\\cbend  In this study we revisit the $z=1.15$ sub-DLA system toward HE~0515--4414 \\citep{ReimersHRW_1998,VargaRTBB_2000}. The main components of associated metal lines exhibit excited neutral carbon and molecular hydrogen \\citep{QuastBR_2002,ReimersBQL_2003}. Most outstanding, the absorption velocity interval extends for $700$~km\\,s$^{-1}$. Based on refined spectroscopy, we examine the ionization, abundances, and differential dust depletion of metals as well as the kinematic structure, and physical conditions of this unusual absorption line system.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  Based on high-resolution spectra obtained with STIS and the VLT UVES we have presented a reanalysis of the chemical composition, kinematic structure, and physical conditions of the H$_2$-bearing sub-DLA system toward HE~0515--4414. The sub-damped system is unusual in several aspects: \\begin{enumerate} \\item The velocity interval of associated metal lines extends for $700$~km\\,s$^{-1}$. The velocity distribution of metal line components is bimodal, indicating the presence of two interacting absorbers. \\item Most of the associated metal line components are formed within \\ion{H}{ii} regions, only one third of the components associated with the predominantly neutral main absorber. \\item For the main components 23-28 the observed abundance ratios of refractory elements Si, Cr, Mn Fe, Ni to Zn show a distinct gradient along the sight line. The differential depletion of refractory elements ranges from Galactic warm disk to halo-like and essentially undepleted patterns. The variation in the dust-to-metal ratio indicates the multiphase structure of the absorbing medium. The dust-corrected metal abundances show the nucleosynthetic odd-even effect and might imply an anomalous depletion of Si relative to Fe, but otherwise do correspond to solar abundance ratios. The intrinsic average metallicity is almost solar, $[\\mathrm{Fe}/\\mathrm{H}]_\\mathrm{m}=-0.08\\pm0.19$, whereas the uncorrected average is $[\\mathrm{Zn}/\\mathrm{H}]_\\mathrm{g}=-0.38\\pm0.04$. For the H$_2$-bearing components 23 and 24 the dust-to-metal and dust-to-gas ratios (relative to Galactic warm disk ratios) are $\\rho=1.02\\pm0.02$ and $\\kappa=0.64\\pm0.31$, respectively. The ion abundances in the periphery conform with solar element composition. \\item The diagnostics of fine-structure lines is not conclusive. Adequate recordings of the H$_2$ lines are needed to provide reliable results. \\end{enumerate}  The presence of \\ion{H}{ii} regions might have consequences for the DLA abundance diagnostics in general. If any metal line components are connected with \\ion{H}{ii} rather than \\ion{H}{i} regions, the usual averaging of element abundances is incorrect. In particular, ionization effects can pretend an enrichment of $\\alpha$ elements. We have obtained a diagnostic diagram (Fig.~\\ref{fg:twocolor}) which allows to detect \\ion{H}{ii} region-like ionization conditions from empirical \\ion{Al}{ii}, \\ion{Al}{iii}, and \\ion{Fe}{ii} column densities. If both $N_\\ion{Al}{iii}/N_\\ion{Al}{ii}>0.25$ and $N_\\mathrm{Al}/N_\\ion{Fe}{ii}>0.40$, the absorbing material is largely ionized. In this context, it is interesting to note that high column densities can be attained by the interception of relatively compact regions. For the present case \\citet{ReimersBQL_2003} have already pointed out that the absorption path length contributed by the H$_2$-bearing components 23 and 24 is less than 1~lyr when the number density of hydrogen atoms is about $100$~cm$^{-3}$.  Our analysis shows that sub-DLA systems can exhibit solar metallicities. If the highest-metallicity sub-DLA absorbers prove to be regular DLA absorbers having consumed large amounts of neutral hydrogen due to massive star formation, their detection is important. Modern surveys of DLA systems setting the cut-off below the traditional column density limit may provide interesting insights."
    },
    "0710/0710.5548_arXiv.txt": {
        "abstract": "We present $\\sim3^{''}$ resolution maps of CO, its isotopologues, and HCN from in the center of Maffei~2. The J=1--0 rotational lines of $^{12}$CO, $^{13}$CO, C$^{18}$O and HCN, and the J=2--1 lines of $^{13}$CO and C$^{18}$O were observed with the OVRO and BIMA arrays. The lower opacity CO isotopologues give more reliable constraints on H$_{2}$ column densities and physical conditions than optically thick $^{12}$CO.  The J=2--1/1--0 line ratios of the isotopologues constrain the bulk of the molecular gas to originate in low excitation, subthermal gas.  From LVG modeling, we infer that the central GMCs have $n_{H_{2}}\\sim 10^{2.75}~cm^{-3}$ and T$_{k}~\\sim 30$ K. Continuum emission at 3.4 mm, 2.7 mm and 1.4 mm was mapped to determine the distribution and amount of \\ion{H}{2} regions and dust. Column densities derived from C$^{18}$O and 1.4 mm dust continuum fluxes indicate the standard Galactic conversion factor overestimates the amount of molecular gas in the center of Maffei~2 by factors of $\\sim$2-4.  Gas morphology and the clear ``parallelogram'' in the Position-Velocity diagram shows that molecular gas orbits within the potential of a nuclear ($\\sim$220 pc) bar.  The nuclear bar is distinct from the bar that governs the large scale morphology of Maffei~2.  Giant molecular clouds in the nucleus are nonspherical and have large linewidths, due to tidal effects.  Dense gas and star formation are concentrated at the sites of the $x_{1}-x_{2}$ orbit intersections of the nuclear bar, suggesting that the starburst is dynamically triggered. ",
        "introduction": "Concentrations of molecular gas are common in the centers of large spiral galaxies, often in the form of nuclear bars.  Bars represent likely mechanisms by which rapid angular momentum loss and gas inflow concentrate gas in the centers of galaxies \\citep[eg.][]{SOIS99,SVRTT05,K05}.  Secondary nuclear bars can be important in controlling the dynamics of the innermost few hundred parsecs of galaxies \\citep[eg.,][]{SFB89,FM93,HSE01,MTSS02,SH02,ES04}. In addition to being an avenue by which nuclear gas and stellar masses are built up, these bars can also significantly influence the physical and chemical properties of molecular clouds within them \\citep[eg.][]{MT01,PW03,MT04,MT05}.  We have observed the two lowest rotational lines of $^{13}$CO and C$^{18}$O, plus new high resolution images of $^{12}$CO(1-0) and HCN(1-0) in the nearby galaxy, Maffei 2, with the Owens Valley Millimeter (OVRO) Array and the Berkeley-Illinois-Maryland Association Array (BIMA).  Because they are optically thin, or nearly so, CO isotopologues more directly trace the entire molecular column density than does optically thick $^{12}$C$^{16}$O (hereafter ``CO'').  Also ratios amongst isotopologues are more sensitive to changes in density and temperature throughout the clouds.  HCN, which has a higher critical density than CO, constrains the dense molecular cloud component.  Here we aim to characterize the properties of molecular clouds in the center of this barred galaxy, and to connect those properties with the dynamics of the nucleus.  Maffei 2 is one of the closest large spirals (D$\\simeq$3.3 Mpc, $\\S$2; Table \\ref{Tmaf}), but lies hidden behind more than 5 magnitudes of Galactic visual extinction \\citep[][]{M68}.  A disturbed, strongly barred spiral galaxy \\citep[][]{HMGT93a,BM99}, Maffei 2 has an asymmetric HI disk and tidal arms that suggest a recent interaction with a small satellite \\citep[][]{HTH96}.  The interaction may be responsible for the bright nuclear CO emission \\citep[eg.,][]{RTP77,I89,MW04} and active nuclear star formation of L$_{OB}~\\simeq ~1.7\\times 10^{9}~$L$_{\\odot}$ \\citep[][corrected for distance]{TH94}. ",
        "conclusions": "New aperture synthesis maps are presented for emission in the J=2--1 and 1--0 transitions of $^{13}$CO and C$^{18}$O, as well as the J=1--0 lines of HCN and CO in the central arcminute ($\\sim$ 1 kpc) of Maffei 2.  The H$_{2}$ column density as traced by optically thin CO isotopologues is similar in morphology to what is implied from $^{12}$CO, except that the emission from the isotopologues is more closely confined to the two extended molecular arm ridges and more uniformly distributed across the central ring.  The dense gas traced by HCN(1-0) is more confined to the center of the galaxy than the CO emitting gas.  The central molecular bar contains five main peaks that resolve into at least 17 distinct GMCs, with radii of $\\sim$40--110 pc and linewidths $\\gsim$ 40 km s$^{-1}$.  In the two innermost molecular cloud complexes, at galactocentric radii of $\\sim 5$\\arcsec\\ (80 pc from the dynamical center), the GMCs are distinctly nonspherical, elongated along the nuclear bar, with linewidths as large as 100 km s$^{-1}$.  These GMCs are probably being tidally stretched due to the nuclear potential.  The H$_{2}$ column density for the central GMCs is $N_{H_{2}}~\\simeq ~4.4 - 10\\times 10^{22}~ cm^{-2}$ ($\\overline{\\Sigma} ~\\sim 950 - 2200$ M$_{\\odot}$ pc$^{-2}$), corresponding to mean optical extinctions of $A_v \\sim$40--100.  The molecular mass within the central 20\\arcsec\\ galactocentric radius ($\\sim$300 pc) is 2.1$\\times10^{7}~M_{\\odot}$, while the dynamical mass in the same region is M$_{dyn}(20^{''})=7.3\\times 10^{8}~M_{\\odot}$.  The molecular mass is only a few percent of the dynamical mass. Excitation temperatures, assuming $^{13}\\tau (^{18}\\tau) ~\\sim$ 1 ~($\\ll$1), are T$_{ex}~\\sim ~ 3~-~6$ K over much of the central 500 pc for both $^{13}$CO and C$^{18}$O.  These T$_{ex}$ values are low compared with the brightness temperature observed in CO ($\\gsim 30$ K) indicating subthermal excitation, and that the average densities of the GMCs are probably only moderate.  Single component LVG analysis of the GMCs in CO, $^{13}$CO, and C$^{18}$O yield best-fit solutions of $n_{H_{2}}~\\simeq ~ 10^{2.75}~ cm^{-3}$ and T$_{kin} ~\\simeq ~ 20 - 30$ K.  Average densities estimated from the total C$^{18}$O column densities are consistent with these values.  The $^{13}$CO and C$^{18}$O lines are weaker than expected from CO(1-0), which appears to be overluminous per unit gas mass across the starburst region.  Column densities derived from both C$^{18}$O and 1.4 mm dust continuum emission imply that $X_{CO}$(Maf~2) is about 2--4 times lower than the Galactic value, similar to $X_{CO}$ values found for the centers of other large spirals, including our own.  The weakness of the isotopologues at large galactocentric radii and in the ``off-arm'' spray regions of Maffei 2, suggest that in these regions either the isotopologues cease to effectively trace molecular gas or that the Galactic conversion factor overestimates the molecular column.  The lack of applicability of the Galactic $X_{CO}$ to the clouds in the center of Maffei 2 is probably due to the effect of bar motions and strong tides on the structure and dynamics of these clouds.  Millimeter continuum emission reveals three prominent locations of star formation with the most intense occurring where the molecular bar intersects the nuclear ring.  Lyman continuum rates of N$_{Lyc}\\sim 3$--$5 \\times 10^{51}$ s$^{-1}$ are implied for individual regions. The total rate for the entire nucleus is $_{Lyc}\\sim 2.6 \\times 10^{52}$ s$^{-1}$, or SFR $\\sim$0.26 M$_{\\odot}$ yr$^{-1}$.  A P-V diagram of the nucleus of Maffei 2 shows a distinct ``parallelogram'' indicating molecular gas response to a barred potential.  The morphological and kinematic data confirms Maffei 2 as true double barred galaxy.  We suggest a bar model where the nuclear gas distribution and velocity is governed by a small nuclear bar of $r = \\sim$110 pc.  An upper limit to the mass inflow rate along the nuclear bar is $dM/dt~ \\lsim ~0.7$ M$_{\\odot}$ yr$^{-1}$, enough to drive the current star formation rate seen at the end of the bar arms and populate the nuclear ring with gas.  The locations of star formation and the dense gas in the central region appear to coincide with the location of the $x_{1}-x_{2}$ orbit crossings of the nuclear bar, consistent with dynamical triggering of the the star formation."
    },
    "0710/0710.0647_arXiv.txt": {
        "abstract": "{We present multi-wavelength observations of bright quasar \\HS: the optical data taken with the MMT and Keck telescopes, with the 40-50~\\kms\\ resolution, and X-ray data taken by the \\emph{Chandra X-ray Observatory} satellite.} {The optical spectra contain a very large number of absorption lines from numerous heavy elements. Our goal is to analyse those features to obtain constraints on the properties of associated absorbers, to be used in modeling of the quasar intrinsic flux and properties of the clouds.} {We have determined the properties -- column densities and redshifts -- of the individual components. We derived the X-ray properties of \\HS\\ and the optical-to-X-ray slope index \\aox.} {We found \\aox\\ of 1.70, which is on the high end of the typical range for a radio quiet QSO.  We found 49 individual heavy element absorption clouds, which can be grouped into eleven distinct systems. Absorbers from the associated system which likely is the one spatially closest to the QSO show large \\CIV\\ to \\HI\\ column density ratio, reaching $\\sim$20.} {Intrinsic X-ray properties of the quasar are typical. Determination of column densities of ions (including hydrogen) gives a strong foundation for modeling of the quasar ionising flux.} ",
        "introduction": "\\label{sec:intro}  \\object{HS1603+3820}\\ is a $\\zem=2.54$ quasar discovered during the Hamburg/CfA Bright Quasar Survey \\citep{dobrzycki1996}. With $B=15.9$, it is among the top few brightest quasars known in the $\\zem=2-3$ redshift interval \\citep{veron2006}. The most striking feature of \\HS\\ is, however, the extreme richness of its heavy element absorption spectrum, and in particular the $\\zabs\\approx\\zem$ absorption.  The absorption spectrum of \\HS\\ was first described in \\citet[][hereafter D99]{dobrzycki1999}, \\defcitealias{dobrzycki1999}{D99} based on medium-resolution (90-100~\\kms) spectrum obtained in 1997 April with the MMT and the Blue Channel spectrograph. Although in this spectrum many of the absorption systems were blended, it was clear that the quasar held great promise for the analysis of the properties of associated absorbers.  The absorption systems which are spatially close to the quasar can be used as probes for the intrinsic emission of the quasar itself, since the conditions in the systems are undoubtedly heavily influenced by the QSO flux \\citep[e.g.][and references therein]{crenshaw2003,gabel2005,scott2005}. Systems which contain lines originating from various species and various ionisation levels are particularly interesting, since they can provide independent way of determining physical properties of the absorbing medium, which can cast light at the intrinsic quasar continuum in far UV band as well as the absorbing cloud formation.  With that as the principal goal, we initiated a project of multi-wavelength observations of \\HS, aiming at obtaining spectral information at various wavelengths. In this paper we present the results of our X-ray and optical observations of the quasar. In the forthcoming paper (R\\'o\\.za\\'nska et al.\\ 2007, in preparation; Paper~II) we use those observations to constrain models of the quasar broad band continuum and the cloud confinement.  In the meantime, \\citet[][hereafter referred to as M03 and M05, respectively]{misawa2003,misawa2005} \\defcitealias{misawa2003}{M03} \\defcitealias{misawa2005}{M05} observed \\HS\\ in 2002 March and 2003 July in very high resolution (6.7~\\kms), using the Subaru telescope with the High Dispersion Spectrograph. The primary goal of those observations was to establish the nature of the absorbers (intervening or associated), using the analysis of line time variability, partial coverage and profile shapes. High resolution of their observations enabled resolving many of the systems into a number of narrow components. \\citetalias{misawa2003,misawa2005} grouped the absorbers into several systems, designated with the letters from A to H. In our optical spectroscopy data we see all those systems and three additional intervening systems.  For consistency, and to enable easy comparisons, wherever possible we will use those designations throughout this paper.  In a very recent paper, \\citet[][hereafter M07]{misawa2007} \\defcitealias{misawa2007}{M07} expanded the analysis of absorber variability, using new spectra from the Subaru and Hobby-Eberly telescopes, and confirmed the finding from \\citetalias{misawa2005} that \\CIV\\ in system A is highly variable. This paper also described one system (labeled I) which they did not have in their older data. This system was seen in our data, and for consistency with \\citetalias{misawa2007} we rewrote parts of this paper to also use this designation.  Our optical spectra are of lower resolution (FWHM of 40-50~\\kms), but they complement the Subaru data in two important areas: they have higher signal-to-noise ratio (100-140 per pixel in the \\CIV\\ and \\Lya\\ regions) and they cover much larger wavelength range ($\\sim$3700 to $\\sim10000$~\\AA). The former reduces uncertainty in blended regions containing thick lines and helps in finding weak lines, while the latter allows analysing lines from species inaccessible in the Subaru spectra.  The most important addition is detection and analysis of hydrogen lines, not seen in the Misawa et al.'s spectra, as well as detection of heavy-element lines other than \\CIV\\ in system A.  This paper is organised as follows. In Section~\\ref{sec:obsxray} we discuss the \\emph{Chandra} X-ray observation of \\HS. In Section~\\ref{sec:obsopt} we present our optical observations of \\HS, and we discuss the individual systems and compare our results with those of Misawa et al., discussing the possible sources of discrepancies. We comment on the absorber variability in Section~\\ref{sec:variability}. We summarise the results in Section~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  The following summarises our findings:  \\begin{itemize}  \\item The intrinsic X-ray properties of \\HS\\ are in the typical range. The available X-ray data were insufficient to provide meaningful constraints on the properties of the intrinsic absorber. The optical-to-X-ray slope,  $\\aox=1.70$, which is at the high end of the typical range at this redshift.  \\item In addition to systems seen in \\citetalias{dobrzycki1999} and Misawa et al.'s papers, we identify two previously unknown intervening systems at $\\zabs=1.9266$, and 2.2622. With eleven complexes of heavy element absorbers, which together contain $\\sim$50 (or even more) individual clouds -- and that not counting \\Lya\\ forest systems -- \\HS\\ is truly a spectacular object.  \\item Some of the clouds in system A have a high ratio of \\CIV\\ to \\HI\\ column densities, reaching $\\sim$20.  Clouds from other systems show $N_{\\CIV}/N_{\\HI}<1$.  \\item The data presented in this paper are not well suited for proper analysis of line temporal variability. We note, however, that the data are consistent with absence of variability for all systems.  \\item The optical data presented in this paper contain identifications and measurements of several lines not seen in the Subaru spectra, in particular the hydrogen and magnesium lines.  \\item Our data contain measurements of column densities of several ions and therefore provide excellent constraints for modeling of the cloud internal properties and QSO intrinsic flux. We will present such modeling in the forthcoming paper.  \\end{itemize}"
    },
    "0710/0710.2368_arXiv.txt": {
        "abstract": "We discuss the intriguing possibility that dark energy may change its equation of state in situations where large dark energy fluctuations are present. We show indications of this dynamical mutation in some generic models of dark energy. ",
        "introduction": "There is ample evidence that supports the existence of dark matter (DM) and dark energy (DE) in the universe \\cite{Seljak}. Due to its charge neutrality and dust-like equation of state (i.e., negligible pressure), dark matter starts to cluster gravitationally very early in the history of the Universe, and is crucial for the formation of large scale structure. Dark energy, on the other hand, becomes relevant only more recently, and is presumed to be a smooth component with a negative equation of state in order to fuel the accelerated expansion of the Universe \\cite{Reviews}.  At the background level, dark energy (or any other fluid) is completely determined by its equation of state $w=p_e/\\rho_e$, where $p_e$ is the pressure and $\\rho_e$ is the energy density of dark energy. Already at this level dark energy can affect large scale structures \\cite{Linder05} -- see also \\cite{LiberatoRosenfeld} for the effect in different parameterizations of the equation of state of dark energy.  However, if dark energy is in fact a manifestation of a dynamical mechanism such as a scalar field, then it will also develop inhomogeneities due to its gravitational interactions with itself and with dark matter \\cite{Cobleetal}. In linear perturbation theory, besides the energy density perturbation $\\delta\\rho$, we need two extra degrees of freedom to characterize cosmological perturbations: the pressure perturbation $\\delta p$ and the scalar anisotropic stress $\\pi$ \\cite{Bardeen,KodamaSasaki}. Alternatively, one can also use the velocity potential $\\theta=\\vec\\nabla \\cdot \\vec{v}$ and the anisotropic stress \\cite{Bertschinger}.  The inhomogeneities of dark energy are often quite small, particularly in the case of ordinary (canonical) scalar field models with almost $\\Lambda$-like behaviour -- that is, when $w \\simeq -1$ \\cite{Cobleetal,DuttaMaor,MotaShawSilk}. In fact, as $w \\rightarrow -1$ the perturbations in all dark energy models are suppressed in relation to those of dark matter. However, if that is not the case then the dark energy density contrast $\\delta_e \\equiv \\delta\\rho_e/\\rho_e$ can be either small or large, depending on the pressure perturbations of dark energy \\cite{US}.  Here we present further evidence of an intriguing possibility which was first pointed out in supergravity-motivated scalar field models of dark energy \\cite{MotaBruck,NunesMota}: that dark energy can mutate into a fluid with clustering properties similar to those of dark matter. We will show that this effect is a generic feature of dark energy, and that it has a simple origin: when pressure perturbations are large, the effective equation of state inside a collapsed region can be completely different from the equation of state of its homogeneous component.   \\section {Gravitational collapse with dark energy}  In the following, in order to specify the properties of the dark energy perturbations we will neglect the anisotropic stress. Moreover, we will characterize the pressure perturbation in a simplified manner, using the so-called effective (or non-adiabatic) sound velocity \\cite{Hu}, defined as $c_{eff}^2 \\equiv \\delta p_e/ \\delta \\rho_e$, which we will assume to be a function of time only. Notice that this assumption lacks a formal basis in perturbation theory, since $\\delta p_e$ is a perturbed variable whose time and spatial dependences can be, and often are, independent of the variations of $\\delta\\rho_e$. Nevertheless, in a particular gauge (the so-called ``rest frame'' of the fluid, where $T^i_0 = 0$), the effective sound speed coincides with the phase velocity of linear relativistic perturbations, $c_X^2$ \\cite{Hu,Mukhanov}.  Describing the pressure perturbation as $\\delta p_e = c_{eff}^2 \\, \\delta\\rho_e$ allows us to treat a wide variety of dark energy models and, crucially, it also allows us to compute non-linear structure formation using the Spherical Collapse model (SC) \\cite{GunnGott}. Furthermore, in this case the SC equations (derived in a simplified relativistic framework) are identical to the equations of pseudo-Newtonian cosmology \\cite{US2}. This means that the physics of gravitational collapse of structures such as galaxy clusters is well described within this framework.  Consider then, in the spirit of the SC model, a spherically symmetric region of constant dark energy overdensity (the so-called ``top-hat'' density profile.) Let us call $\\rho^c_e = \\rho_e + \\delta \\rho_e$ and $p^c_e = p_e + \\delta p_e$ the energy density and the pressure of this region, which are modified with respect to the corresponding background quantitites, $\\rho_e$ and $p_e$ by the perturbations $\\delta \\rho_e$ and $\\delta p_e$. The equation of state $w$ is defined as the ratio of the total pressure to the total energy density, and hence it will be different for the background and the interior of the collapsed region. A simple calculation shows that the equation of state inside the collapsed region, $w^c$, is given by: \\begin{equation} w^c = \\frac{p_e + \\delta p_e}{\\rho_e + \\delta\\rho_e} = w + (c_{eff}^2 - w) \\frac{\\delta_e}{1+ \\delta_e} \\; . \\label{deltaw} \\end{equation} For small density contrasts $|\\delta_e| \\ll 1$, the equation of state inside the overdense region does not change appreciably. However, if $c_{eff}^2 \\not= w$, then in the nonlinear regime, where $\\delta \\gtrsim 1$ (halos), there could be a substantial modification in $w^c$ with respect to the background equation of state. Even in underdense regions (voids), where $\\delta \\approx -1$, there could be large modifications of the equation of state. Hence, in principle dark energy could even effectively mutate into dark matter inside halos and voids.   The above argument is completely general. What remains to be shown is whether there are models in which this dramatic situation is actually realized. This requires a non-linear analysis of the evolution of pertubations for two gravitationally coupled fluids, dark energy and dark matter (we will neglect radiation and baryons in what follows). Unfortunately, at present there are no totally rigorous methods for performing this analysis -- except in the case of canonical scalar fields, but even then only approximately \\cite{DuttaMaor,MotaShawSilk}.  Here we employ a generalization of the SC model for the case of a relativistic fluid with pressure. In the next section we present the relevant equations and mention under which conditions they are equivalent to a pseudo-Newtonian approach. We will then analyse the evolution of the coupled system of perturbations for a wide variety of dark energy models for which the pressure perturbations are characterized by some homogeneous effective sound speed. ",
        "conclusions": "We have shown that it is possible to change radically the clustering properties of dark energy in collapsed regions (halos and voids.) We exemplified this behaviour with a few models for the dark energy perturbations, and showed that it happens not only in scalar field models, but also in generic models of dark energy -- in particular the Generalized Chaplygin Gas and Silent Quartessence models.   Since the physics of most observed collapsed structures, such as galaxy clusters, is well approximated by quasi-Newtonian physics, this dynamical mutation should be a general phenomenom. Clearly, this is a crucial issue for all attempts to compute the influence of dark energy on the formation of large scale structures. More detailed studies, including a relativistic approach and using different realistic parameterizations of the dark energy equation of state are currently under way \\cite{US2}."
    },
    "0710/0710.1040_arXiv.txt": {
        "abstract": "{On 4 July 2005 at 05:52\\,UT, the impactor of NASA's Deep Impact (DI) mission crashed into comet 9P/Tempel 1 with a velocity of about 10\\,\\kms.  The material ejected by the impact expanded into the normal coma, produced by ordinary cometary activity.} {Based on visible and near-IR observations, the characteristics and the evolution with time of the cloud of solid particles released by the impact is studied in order to gain insight into the composition of the nucleus of the comet. An analysis of solid particles in the coma not related to the impact was also performed.} {The characteristics of the non-impact coma and cloud produced by the impact were studied by observations in the visible wavelengths and in the near-IR. The scattering characteristics of the \"normal\" coma of solid particles were studied by comparing images in various spectral regions, from the UV to the near-IR. For each filter, an image of the \"normal\" coma was then subtracted from images obtained in the period after the impact, revealing the contribution of the particles released by the impact.} {For the non-impact coma the Af$\\rho$, a proxy of the dust production, has been measured in various spectral regions. The presence of sublimating grains has been detected. Their lifetime was found to be $\\sim 11$ hours.  Regarding the cloud produced by the impact, the total geometric cross section multiplied by the albedo, SA, was measured as a function of the color and time. The projected velocity appeared to obey a Gaussian distribution with the average velocity of the order of 115\\,\\ms. By comparing the observations taken about 3 hours after the impact, we have found a strong decrease in the cross section in $J$ filter, while that in Ks remained almost constant. This is interpreted as the result of sublimation of grains dominated by particles of sizes of the order of some microns.} {} ",
        "introduction": "The Deep Impact mission (hereafter DI) to the Jupiter family comet 9P/Tempel 1 (hereafter 9P) was aimed at studying the cratering physics in minor bodies in the solar system and the primordial material preserved inside cometary nuclei. On July 4, 2005 the impactor of the DI experiment produced a high-speed (about 10\\,\\kms) impact in the nucleus of 9P excavating a considerable amount of cometary material that was observed and measured both in-situ by the DI fly-by spacecraft and remotely by Earth-based instrumentation. First results of the mission are described in \\citet{AHearn2005} and \\citet{Sunshine2006}. Early earth-based and other space-based measurements of the event have been published by \\citet{Meech2005}, \\citet{Sugita2005}, \\citet{Harker2005}, \\citet{Lisse2006}, and \\citet{Schleicher2005}.  At the European Southern Observatory (ESO) DI received considerable observing time allocated to observe the event at their Chilean observatory, at Cerro La Silla and at Cerro Paranal sites \\citep{Kaeufl2005a}. Here, we summarize results from the visible and near-IR measurements of the dust in the cometary coma obtained both shortly before and after the DI event. We focus on the dust ejecta properties such as scattering properties, projected velocity, and spatial distribution and their evolution with time.  Complementary data from the ESO DI campaign on polarimetric and mid-IR observations as well as on the cometary gas emission and the large-scale coma activity of the comet are described elsewhere (see e. g., Boehnhardt et al., 2007). Pre-impact monitoring of the cometary activity is described by \\citet{Kaeufl2005b} and \\citet{Lara2006}. ",
        "conclusions": "From observations of gas emission-free regions of the comet 9P/Tempel\\,1 made before and after the Deep Impact event, the scattering characteristics and the velocity of the ejecta cloud produced by the impact have been measured.  Seventeen and a half hours after the impact, the total area covered by the grains of the ejecta cloud, multiplied by their albedo, was $27-35$\\,km$^2$ in $JH\\Ks$ . Three hours later, it dropped to $\\approx 18$\\,km$^2$ in $J$, but remained almost constant in \\Ks. During this interval of time, the $J-\\Ks$ color gradient with the nucleocentric distance $\\rho$ also changed significantly, increasing by a factor of three in the direction of the Sun.  The projected average velocity of the ejected cloud measured in the near-IR was $115 \\pm 16$\\,\\ms, and was found to be independent of the filter used for the observations and the position angle. Its distribution was very similar to a Gaussian.  It has been shown that all these results cannot be explained assuming that grain size has a power distribution nor it is possible to assume that grains are ejected by gas drag.  While the mechanism of grain ejection is difficult to explain, the behavior afterward can be justified only with the presence of $\\approx 1$\\,$\\mu$m size organic grains that are sublimated by the solar radiation.  From the pre-impact and late post-impact observations, the presence of a sublimating component has been detected and interpreted in terms of organic grains that sublimate because of the solar radiation."
    },
    "0710/0710.3874_arXiv.txt": {
        "abstract": "The three-flavor crystalline color-superconducting (CCS) phase of quantum chromodynamics (QCD) is a candidate phase for the ground state of cold matter at moderate densities above the density of the deconfinement phase transition. Apart from being a superfluid, the CCS phase has properties of a solid, such as a lattice structure and a shear modulus, and hence the ability to sustain multipolar deformations in gravitational equilibrium. We construct equilibrium configurations of hybrid stars composed of nuclear matter at low, and CCS quark matter at high, densities. Phase equilibrium between these phases is possible only for rather stiff equations of state of nuclear matter and large couplings in the effective Nambu--Jona-Lasinio Lagrangian describing the CCS state. We identify a new  branch of stable CCS hybrid stars within a broad range of central densities which, depending on the details of the equations of state,  either bifurcate from the nuclear sequence of stars when the central density exceeds that of the deconfinement phase transition or form a new family of configurations separated from the purely nuclear sequence by an instability region. The maximum masses of our non-rotating hybrid configurations are consistent with the presently available astronomical bounds. The sequences of hybrid configurations that rotate near the mass-shedding limit are found to be more compact and thus support substantially larger spins than their same mass nuclear counterparts. ",
        "introduction": "Matter in the interiors of neutron stars is compressed by gravity to densities which are by factors 5 to 10 larger than the density of an ordinary nucleus. At such densities baryons are likely to lose their identity and dissolve into their constituents - deconfined quarks. It has been suggested long ago that quark matter may exist in the interiors of compact objects~\\cite{DECONFINMENT}. Within the deconfined phase unbound quarks form  Fermi seas and attractive interactions bind them into Cooper pairs thus giving rise to quark superconductivity~\\cite{CSC,Casalbuoni:2003wh}. Along with the deconfinement phase transition the dense matter undergoes a second transition related to the  restoration of  chiral symmetry (the dynamical symmetry of the strong interaction). This symmetry may remain broken in some color-superconducting phases, such as the color-flavor-locked phase.   If compact (hybrid) stars featuring quark cores surrounded by a nuclear mantle exist in nature, they could provide a unique window on the properties of QCD at high baryon densities. Even if the relevant densities and sufficiently low temperatures could be reached in a laboratory in the future, the conditions prevailing in compact stars are different from those produced in accelerators: the matter is long-lived, charge neutral and in $\\beta$-equilibrium with respect to weak interactions.  A central question of the theory of compact stars, which we will address in this work, is whether the equation of state of matter at high densities admits stable configurations of self-gravitating objects in General Relativity featuring deconfined quark matter, and if so, are the gross parameters of these objects, like their mass and radius, compatible with the known astronomical bounds? The deconfinement phase transition from baryonic to quark matter leads to a softening of the equation of state which could lead to an instability towards a collapse into a black hole. The details of a stability analysis may depend on the theoretical model of deconfined, color-superconducting quarks. The Nambu--Jona-Lasinio (NJL) model is a  non-perturbative low-energy approximation to  QCD, which is anchored in the low-energy phenomenology of the hadronic spectrum. While  dynamical symmetry breaking, by which quarks acquire mass, is incorporated in this model, it lacks confinement~\\cite{Klevansky:1992qe}. Previous studies of hybrid color-superconducting stars within the three-flavor NJL model displayed a general instability of hybrid stars towards  collapse into a black hole~\\cite{Buballa:2003et}. Stable stars featuring two-flavor superconducting (2SC) phases were obtained with typical maximum masses $M\\sim 1.7 M_{\\odot}$ under a favorable choice of constituent quark masses (which are related to the value of the baryo-chemical potential at zero pressure)~\\cite{Shovkovy:2003ps} or by replacing the hard NJL cut-offs by soft form factors with parameters fitted to the same set of data~\\cite{Grigorian:2003vi,Blaschke:2007ri}. Heavier objects are obtained if a repulsive vector interaction is introduced in the NJL Lagrangian~\\cite{Klahn:2006iw}.   More phenomenological studies of hybrid stars based on the MIT bag model where carried out using a generic parameterization of the quark matter equation of state~\\cite{Alford:2004pf}. Non-perturbative QCD corrections to the equation of state of the Fermi gas were accounted through the parameter $a_4=1-c$, where $c\\simeq 0.3$ is a reasonable estimate~\\cite{Fraga:2001id}. For this value of the correction stable hybrid stars containing color-flavor-locked superconducting quark matter exist with maximum mass 1.9~$M_{\\odot}$.   Because of $\\beta$ equilibrium in the light quark sector the chemical potentials of  $u$ and $d$ quarks obey the constraint $\\mu_d \\simeq \\mu_u+\\mu_e$~\\cite{COMMENT1}, from which it is evident that the Fermi surfaces of the light flavor quark flavor are shifted by an amount corresponding to the chemical potential of electrons.  Furthermore, the Fermi surface of strange quarks is mismatched with the Fermi surfaces of light $u$ and  $d$ quarks because of the large strange quark mass. A generic feature of fermionic systems with mismatched Fermi surfaces is the existence of phases with broken spatial symmetries. A particular realization is the Larkin-Ovchinnikov-Fulde-Ferrell (LOFF) phase~\\cite{LO,FF}, which is theoretically predicted  in different fermionic systems ranging from dilute atoms~\\cite{LOFF_ATOMS}, nuclear matter~\\cite{LOFF_NUCL} to dense quark matter~\\cite{Alford:2000ze,Casalbuoni:2003wh,Rajagopal:2006ig, Casalbuoni:2005zp,Ippolito:2007uz,Kiriyama:2006ui,Mannarelli:2006fy,Fukushima:2007bj,he:2007}.    The rotational and translation symmetries are broken by the superconducting LOFF phase, since the condensate carries a non-zero momentum~\\cite{LO,FF}. The study of crystalline color superconductivity (CCS) in the context of QCD reveals that this phase could be the true ground state of deconfined quark matter at intermediate densities~\\cite{Alford:2000ze,Casalbuoni:2003wh,Rajagopal:2006ig, Casalbuoni:2005zp,Ippolito:2007uz,Kiriyama:2006ui,Mannarelli:2006fy}. Later on we shall describe a particular realization of the three-flavor LOFF phase, which is used in our search for stable sequences of hybrid stars~\\cite{Casalbuoni:2005zp,Ippolito:2007uz}.  The nuclear equation of state, as is well known, can be constructed starting from a number of different principles \\cite{WEBER_BOOK,Sedrakian:2006mq}. We tested a large number of equations of state to construct hybrid star configurations. These belong to the classes of (i) non-relativistic variational and Bruckner-Hartree-Fock theories which use as an input a non-relativistic potential fitted to the elastic nucleon-nucleon scattering data; (ii) relativistic mean-field models which are fitted to the bulk properties of nuclear matter, and (iii) relativistic models which include correlations at the level of the covariant scattering amplitude (Dirac-Bruckner-Hartree-Fock theories). For present purposes, it is sufficient to characterize these equations of states by their stiffness; as we shall see, only the stiffest equations of state are admissible for phase equilibrium between nuclear and quark matter.   The paper is organized as follows. In Sec.~\\ref{sec:phases} we discuss the input equations of state with the emphasis on the CCS phase of quark matter. Section~\\ref{sec:results} contains our results for the sequences of non-rotating and rapidly rotating hybrid stars. Our conclusions are collected in Sec.~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We constructed a set of equations of states of hybrid stars which are composed of color-superconducting quark matter at high, and purely nuclear matter at low densities. High-density quark matter is described in terms of the semi-microscopic NJL model which includes pair correlations that lead to the three-flavor LOFF phase as the ground state of QCD at moderate densities. The low-density nuclear phase is described in terms of hard relativistic equations of state based on the Dirac-Bruckner-Hartree-Fock theory. The sequences of stellar configurations constructed on the basis of these equations of state reveal a new branch of {\\it stable hybrid configurations} featuring CCS matter. These are analogous to the twin configurations obtained for non-superconducting quark matter in Refs.~\\cite{Glendenning:1998ag} and are in contrast to the previously reported NJL-model based sequences of non-rotating three-flavor color-superconducting hybrid stars, which were either entirely unstable or had a narrow range of stable configurations adjacent to the point where they bifurcate from the nuclear ones.   A further search in the parameter space is required in order to judge how generic our constructions are; a quite general conclusion of our analysis is that the nuclear equations of states need be {\\it hard} to enable the matching to the quark equations of state. Furthermore, matching is only achieved if the ratio of the couplings in the diquark and quark-anti-quark channels is large ($\\eta =1$) compared to the value obtained through the Fierz transformation ($\\eta = 0.75$). The lack of confinement in the NJL model leaves the low-density behavior of the quark matter equation of state somewhat uncertain. Small variations in the bag constant help in achieving a phase transition, but do not affect largely the global properties of hybrid stars (cf. our models A and A1). The maximum masses of our stable hybrid configurations {\\it are consistent with the presently available astronomical bounds on masses and the mass-radius relationship of compact stars}.  We provided a first discussion of {\\it rapidly rotating} configurations of CCS hybrid stars and mapped out their masses and orbital Keplerian frequencies at which they are destabilized by mass-shedding from the equator. Rapidly rotating configurations support larger masses (as expected). The mass-central density relation for this class of objects is ``flatter'' than for their non-rotating counterparts, i.~e., smaller variations of mass can drive the system unstable. The Keplerian frequencies of massive members of the sequences, in particular hybrid configurations, are by a factor of two larger than the rotation frequency 716 Hz of the fastest millisecond pulsar known to date. We find that the hybrid configurations, being more compact, can rotate at substantially larger rates than their nuclear counterparts of the same mass. The moment of inertia of a hybrid star is smaller than that of a nuclear star of the same mass. At the point of bifurcation it drops abruptly giving rise to a sequence of stable hybrid stars with almost the same masses but different moments of inertia; this implies that simultaneous mass and moment-of-inertia measurements are useful in distinguishing hybrid configurations from their nuclear counterparts."
    },
    "0710/0710.3429_arXiv.txt": {
        "abstract": " ",
        "introduction": "The space of inflationary models is vast and varied. While it may be impossible to reconstruct the precise inflationary potential from data, narrowing the classes of models to those that give rise to our corner of the universe is a vital step in describing our cosmology. A promising avenue for doing so in the near future will be the measurement of the level of tensor fluctuations from inflation \\cite{Efstathiou:2006ak}, which is typically characterized by the ratio $r=T/S$, where $T$ and $S$ are the amplitudes of tensor and scalar perturbations respectively  Experimental constraints on $r$ come from measurements of the microwave background polarization.  The current bound from WMAP and SDSS is $r < 0.3$ \\cite{Tegmark:2006az}. We can expect this limit on $r$ to be driven down by the next generation of CMB experiments. In addition to the soon to be launched Planck satellite, there are quite a few experiments on the horizon that have been specifically designed to measure B-mode polarization and thus inflationary gravity waves (IGW).\\footnote{In this paper  we consider only B-modes from tensor fluctuations during inflation. We will not discuss B-modes from cosmic strings,  which might be produced during the exit from inflation. The observational signatures for these two types of B-modes are quite different, see for example \\cite{Pogosian:2006hg}.} These include Spider, SPUD, EBEX, polarBear, QUIET, and Clover. BICEP and QUAD are already taking data now. All of these hope to detect $r \\geq 0.1$ (or even $r \\geq 0.02$ for BICEP2) or place the corresponding upper bound, and should have results by 2011 \\cite{Netterfield}. Further into the future the NASA CMBpol satellite project \\cite{Bock:2006yf} is expected to be launched by 2018 and hopes to achieve a detection or a bound at $r < ~0.01$. The ESA Bpol project has an analogous goal.  There are known inflationary models predicting detectable levels of IGW with $ 10^{-3} <r < 0.3$. These include chaotic inflationary models \\cite{Linde:1983gd,Destri:2007pv,Kallosh:2007wm,Linde:2007fr}, with quadratic, cubic and quartic potentials where one expects $r\\geq 2\\cdot 10^{-2}$ \\cite{Destri:2007pv}, and natural inflation models \\cite{Freese:1990rb} where one expects $r\\geq 10^{-3}$, \\cite{Savage:2006tr}.  The potential for natural inflation is $\\Lambda^4(1-\\cos (\\phi/f))$ where $\\phi$ is the canonically normalized field and $f$ is the axion decay constant.  Quite generally, as shown by Lyth (and exemplified by chaotic inflation), single-field models predicting measurable IGW require a super-Planckian variation of the inflaton over the 60 e-foldings corresponding to the observable part of our universe; $\\Delta \\phi >  M_{Pl}$ \\cite{Lyth:1996im}. For axion models this implies $f > M_{Pl}$ \\footnote{We use the reduced Planck mass, $M_{Pl}=(\\sqrt {8\\pi G_N})^{-1}$}.  The variation of individual fields can be somewhat smaller in models of assisted inflation \\cite{Liddle:1998jc}, with many fields increasing the frictional force during inflation. In such models the value of the VEVs of individual fields may be reduced; for example by a factor of ${1\\over \\sqrt N}$ in assisted $m^2\\phi^2$ inflation, where $N$ is the number of fields.  Our concern in this paper is the study of models derivable from string theory and effective supergravity which have appreciable levels of IGW. Such models are rare and, as discussed recently in \\cite{Kallosh:2007wm,Kachru:unpub,Kallosh:2007ig}, they are not amongst those scenarios that have been well studied in string theory. For example, brane inflation models were shown to be IGW free in \\cite{Baumann:2006cd,Bean:2007hc} and the same is true of modular inflation proposals (see \\cite{Kallosh:2007ig} for a review of these models).  Part of the problem with embedding IGW in string theory is that the value of $r$ correlates directly with the energy scale of inflation. For a given value of $r$, the Hubble constant during inflation is roughly $H\\sim r^{1/2}\\times 10^{14} ~{\\rm GeV}$. This dependence means that, for measurable values of $r$, inflation would have to have occurred at close to GUT-scale energy densities -- $V \\sim (10^{16} {\\rm GeV})^4$. Now, since the moduli of a string theory compactification need to be stabilized above the energy scale of inflation to obtain a robust model (as discussed in many papers such as \\cite{Kachru:2003sx}), it is clear that generating a measurable value $r$ translates into the need for moduli stabilization at very high energy scales.  Unfortunately, achieving this is a challenging problem. In the simplest methods of moduli stabilization, such as the one advocated in the KKLT paper \\cite{Kachru:2003aw}, there is a constraint relating the mass of gravitino to the Hubble parameter, $m_{3/2}\\geq H$ \\cite{Kallosh:2004yh}. This would predict $r\\sim 10^{-24}$ for $m_{3/2} \\sim  \\textrm{TeV}$ and require an extremely heavy gravitino for detection at $r\\sim 10^{-2}$, as shown in \\cite{Kallosh:2007wm}. It is possible to avoid this problem by modifying the original KKLT construction as shown in \\cite{Kallosh:2004yh}. However, achieving $r \\sim 10^{-2}$ in  such models for $m_{3/2} \\sim  \\textrm{TeV}$ would require significant fine-tuning \\cite{Kallosh:2007wm}.    An additional difficulty with finding IGW in string theory arises from considering the aforementioned Lyth bound (the general requirement of trans-Planckian field motion to get measurable gravity waves) in string theory. Not much is known about the behavior of a variation of a scalar modulus over a range $\\Delta \\phi > M_{\\rm string}$ in a string theory moduli space and what we suspect is not encouraging. In general, maintaining the flatness of the potential over a distance $\\Delta \\phi \\gg M_{\\rm string}$  appears to be difficult. However, by using axions, whose potential is to a significant extent determined by symmetries, one would evade some of the issues that appear to prevent flatness \\cite{Binetruy:1986ss,Gaillard:1995az}. Unfortunately, though, it was pointed out in \\cite{Banks:2003sx} that the super-Planckian axion decay constants $f > M_{Pl}$ required for desired regime seem to be unavailable in controlled limits of string theory. While this constraint is met by many of the models we consider in this paper, we note that the argument (reviewed in section \\ref{Banks}) has been shown in \\cite{Kim:2004rp} to be avoidable for a particular choice of two axions. Furthermore, we will demonstrate in section \\ref{sugramulti} that it is possible to avoid the problem even for one axion field, with a particular racetrack-type superpotential. In addition, we will discuss some  numerical factors of $2^m$ and $\\pi^n$ which need clarification to make theoretical string cosmology compatible with cosmological precision data.  Even without $f > M_{Pl}$ axions can still provide IGW-generating potential through the assistance effect discussed above. Assisted inflation has been implemented with stringy axions in the recent N-flation scenario \\cite{Dimopoulos:2005ac}. As mentioned previously, in assisted models one tries to benefit from the Hubble friction generated by N different fields rolling towards their minima. If one can enforce, via symmetries, that the cross-couplings between the fields are suppressed, then one finds that the inflationary slow roll parameters scale with inverse powers of $N$. It was argued in \\cite{Dimopoulos:2005ac} that string axions provide excellent candidates for assisted inflation.  In many models, they are numerous (with the number being determined by the topology of the compact manifold) and they have independent shift symmetries which are broken by distinct instantons.  This provides a rationale for the lack of cross-couplings.  For sufficiently large N, and with a proper choice of scales, one can then hope to make a stringy model of axion assisted inflation that generates IGW. However, although N-flation circumvents several obstacles to inflationary model-building in string theory, it does so at the cost of making certain assumptions. One of our goals in this paper is to see if these assumptions can be realized in controlled string compactifications.  One of the main assumptions in the analysis of \\cite{Dimopoulos:2005ac} was that the scalar superpartners of the axions (which are K\\\"ahler moduli in many classes of models), get effective masses $>H$ during inflation and decouple from the dynamics. Accordingly, we will concern ourselves with the necessity of finding models of moduli stabilization where the volume moduli have much steeper potentials than their axionic partners. If such a model can be found, then quite plausibly the physics will be that of many massive axions acting in concert to give some combination of assisted chaotic inflation and natural inflation. If not, while there is no general proof that inflation does not occur, the story is certainly more complicated than the one described in \\cite{Dimopoulos:2005ac}.  In \\cite{Kallosh:2007ig} a class of ${\\cal N}=1$ supergravity potentials with this mass hierarchy (for the specific case of a single complex modulus) were described, and given the name ``axion valley'' potentials. A natural inflation model \\cite{Freese:1990rb} with a pseudo-Nambu-Goldstone boson (pNGb) is a slice of the bottom of the 2d axion valley potential of the supergravity model and, very roughly speaking, an N-flation model involves decoupled dynamics along the bottom of $N$ such valleys. Thus if axion valleys exist in the string landscape we can obtain measurable IGW by realizing either pNGb inflation or N-flation -- the latter has the obvious advantage of having only sub-Planckian VEVs. In \\cite{Kallosh:2007ig} the preliminary analysis suggested that such valleys are hard to find for axions that are superpartners of K\\\"ahler moduli in Calabi-Yau threefolds (or orientifolds), but they may exist for axions or brane position modes which enjoy quadratic shift symmetric K\\\"{a}hler potentials (in some approximation scheme). Here we will make a thorough analysis of this issue, describing some classes of known models which can give rise to the promising K\\\"ahler potentials and test carefully several possibilities for deriving appropriate axion valley potentials from string theory.  Our explorations of axions and inflation in string theory will follow a structure based on the observations made above. We begin in section  \\ref{nflation} by considering the status of axions in string theory and the apparent restrictions on their decay constants. We also discuss the details of N-flation in more depth. We then loosely divide the models we consider into the following categories. First, in section \\ref{onemodulus}, we consider models with K\\\"{a}hler moduli that have a shift symmetric logarithmic K\\\"{a}hler potential (associated with the cubic prepotential) and a non-perturbative superpotential, such as one would expect in KKLT-type scenarios in IIB theories. We demonstrate, in a number of examples, that both the  axion and its partner get masses of the same order -- i.e. we fail to meet the criteria for N-flation;\\, $N$ moduli would not give $N$ decoupled axions.  In section \\ref{kahlerstab} we consider similar models, but use a more general form for the K\\\"{a}hler potential, with $\\alpha'$ corrections to the tree level form. Naively this K\\\"{a}hler potential is such that one might expect it stabilize the volume part of the moduli, leaving a relatively flat axion direction for each modulus, even with a constant superpotential. However, by examining the scenario in detail we prove that stabilization cannot take place except far in the interior of moduli space, where our formulae are suspect and we have little control. We follow this analysis in section \\ref{othermodels}, by looking at related examples in the literature, but with a non-trivial superpotential. Here, though the situation is somewhat different, the examples we study do not provide us with the framework necessary for assisted axion inflation.  We continue our studies in section \\ref{sugramodels} where we first describe the supergravity axion valley model proposed in \\cite{Kallosh:2007ig},  which leads to a natural axion inflation \\cite{Freese:1990rb} due to a quick stabilization of a radial partner of the axion-inflaton. We present an example of an ``racetrack fix'' mechanism  for this model where by using two exponents in the superpotential one can avoid an argument of  \\cite{Banks:2003sx} which puts a constraint on the axion decay constant. We further look at some models in string theory which may lead to a regime of the effective supergravity of the type required in \\cite{Kallosh:2007ig} and conclude that more detailed studies will be required.  In section \\ref{discussion} we stress that investigation  of IGW in string cosmology is very complicated. We have just began the exploration of string theory models where one is able to perform an analysis. Although the best studied models fail to lead to IGW, we do have some indications of classes of models where this is not the case. These models are less understood and require further study, but may well provide predictions of a measurable scalar-tensor ratio. ",
        "conclusions": "\\label{discussion} We have explored the possibility of deriving axion inflation (in particular N-flation) in string theory. We do so, noting that such models offer some hope of implementing stringy inflation giving substantial tensor fluctuations and thus measurable levels of CMB polarization. We began in section \\ref{nflation} by highlighting some of the problems of embedding inflation (with measurable tensor fluctuations) in string theoretic settings. Following this we examined N-flation in more detail, noting that in most reasonable implementations it is assumed that the volume moduli have significantly steeper potentials than the axions that drive inflation.  In section \\ref{onemodulus} we considered the simplest KKLT-type models. We found that, despite the shift symmetry of the K\\\"{a}hler potential, the curvature of the potential in the volume direction was of the same order as in the axion direction -- i.e. both axion and volume moduli had the same order mass.  In section \\ref{kahlerstab} we considered an $\\alpha'$ corrected K\\\"{a}hler potential for many K\\\"{a}hler moduli. Our motivation here was to have the shift-symmetric $K$ itself stabilize the moduli. By construction, such a scenario could only give a mass to the volume-type moduli, potentially allowing non-perturbative corrections to give small masses to the axions and creating the hierarchy required for N-flation. However, we found several theoretical obstacles; with a constant superpotential, neither SUSY nor (with some caveats) non-SUSY stable vacua were found. Furthermore it appears to be only possible to stabilize one volume modulus using the K\\\"{a}hler potential alone.  Therefore in this setting it is difficult to justify the   multiple axions rolling in concert to inflate the universe. We also examined some specific examples with a non-constant superpotential, \\cite{Conlon:2005ki,Westphal:2006tn}. Although in the large volume example of \\cite{Conlon:2005ki} a large mass hierarchy is obtained, it is only for one  of moduli, and even in that case the mass of the axion is too light to provide a realistic model of inflation.  Therefore, at present we find it  difficult within  the standard corner of the stringy landscape\\footnote{Under standard corner of the stringy landscape we mean the well studied Calabi-Yau compactifications with unbroken supersymmetry and known structure of the K\\\"{a}hler potentials. Supersymmetric vacua do not have to come from Calabi-Yau manifolds; some more general examples can be found by considering ``generalized complex'' manifolds.  However, to find  the effective theory around such vacua is difficult \\cite{Grana:2005ny}.} to remove the volume moduli from the dynamics and to keep only nearly flat axion directions. More precisely, it seems impossible to  give the axion a substantially lower mass than its partner. We have pointed out some caveats in our discussions of these issues; these may be studied in future and could, perhaps, change our current conclusions.  Having had no luck with constructible stringy models, we examined supergravity scenarios in section \\ref{sugramodels}. Here we use a quadratic  shift-symmetric (axion independent) K\\\"{a}hler potential and find that  axion pNGb-type inflation with IGW  is possible, as proposed in \\cite{Kallosh:2007ig}. Moreover, we have found here a possible solution to the problem of the multi-instanton corrections \\cite{Banks:2003sx}: one needs a racetrack-type superpotentials with 2 or more exponents, so that each exponent is not small but the small difference between them may provide the correct parameters for inflationary models with IGW.  Also in section 6 we argued that it may be useful to look for type IIB orientifold models which may have features different from generic Calabi-Yau compactifications: such models provide, in certain approximations, the required quadratic shift-symmetric (axion independent) K\\\"{a}hler potentials. Often in such orientifold models the axions originate from hypermultiplets.  Whether such models will be valid in the regime  predicting inflationary gravity waves has to be studied and remains a project we postpone for the future.    There are also string models of inflation with IGW which we have not been discussed in this paper\\footnote{ For example, a model in \\cite{Krause:2007jr} based on heterotic M-theory predicts observable $r$. However, the problem of  stabilization of the orbifold-length and Calabi-Yau volume moduli has not been solved so far. In \\cite{Becker:2007ui}, \\cite{Kobayashi:2007hm} the DBI-type model \\cite{Alishahiha:2004eh}  of a D5 brane wrapped on a 2-cycle in Klebanov-Strassler throat geometry predicts IGW, according to \\cite{Becker:2007ui}. However, the problems of back reaction and unusually large orbifolding in  \\cite{Becker:2007ui} still have  to be resolved.}. In all cases, more study will be required.  Given both the results discussed in this paper and other work in this area, we must conclude that finding measurable IGW in stringy models remains problematic. Accordingly, should the next generation of CMB experiments measure appreciable B-mode polarization and thus infer a substantial scalar-tensor ratio, this would pose a significant challenge for string theory."
    },
    "0710/0710.4515_arXiv.txt": {
        "abstract": "KPD0005+5106 is the hottest known helium-rich white dwarf. We have identified \\ion{Ne}{viii} lines in UV and optical spectra and conclude that it is significantly hotter than previously thought, namely $T_{\\rm eff}=200\\,000$~K instead of 120\\,000~K. This is a possible explanation for the observed hard X-ray emission as being of photospheric origin. Concerning its evolutionary state, we suggest that KPD0005+5106 is not a descendant of a PG1159 star but more probably related to the O(He) stars and RCrB stars. ",
        "introduction": "lines and new temperature determination}  The hottest helium-rich white dwarfs exhibit lines of ionized helium in their optical spectra and they are classified as DO white dwarfs. Currently forty DOs are known. The coolest one has $T_{\\rm eff}=40\\,000$~K and the hottest one is KPD0005+5106.  From an analysis of optical and \\emph{Hubble Space Telescope} FOS spectra $T_{\\rm eff}=120\\,000$~K and $\\log g=7$ has been derived (Werner et al.\\@ 1994) for KPD0005+5106. The Sloan Digital Sky Survey has recently revealed another DO with $T_{\\rm eff}=120\\,000$~K (H\\\"ugelmeyer et al.\\@ 2006).  In the \\emph{Far Ultraviolet Spectroscopic Explorer} (FUSE) spectra of the hottest PG1159 stars ($T_{\\rm eff}>140\\,000$~K) we have recently discovered absorption lines of \\ion{Ne}{viii} (Werner et al.\\@ 2007). The same features were also discovered in KPD0005+5106 (Fig.\\,1). Emission lines in the optical spectrum are also from \\ion{Ne}{viii} and not, as previously thought, from superionized (i.e., non-thermally excited) \\ion{O}{viii}. Non-LTE line profile fits to the \\ion{Ne}{viii} lines of KPD0005+5106 yield $T_{\\rm eff}=200\\,000$~K and $\\log g=6.5$. A reassessment of the \\ion{He}{ii} line spectrum confirms these parameters (Werner et al.\\@ 2007). Hence, KPD0005+5106 is by far the hottest He-rich white dwarf. In fact, it is the hotter than any other white dwarf or central star of planetary nebula. It is only rivaled by the exotic PG1159 star H1504+65 (Werner et al.\\@ 2004, and these proceedings). ",
        "conclusions": ""
    },
    "0710/0710.3335_arXiv.txt": {
        "abstract": "We present results of non-linear, 2D, numerical simulations of magneto-acoustic wave propagation in the photosphere and chromosphere of small-scale flux tubes with internal structure. Waves with realistic periods of three to five minutes are studied, after applying horizontal and vertical oscillatory perturbations to the equilibrium model. Spurious reflections of shock waves from the upper boundary are minimized thanks to a special boundary condition. This has allowed us to increase the duration of the simulations and to make it long enough to perform a statistical analysis of oscillations. The simulations show that deep horizontal motions of the flux tube generate a slow (magnetic) mode and a surface mode. These modes are efficiently transformed into a slow (acoustic) mode in the $v_A < c_S$ atmosphere. The slow (acoustic) mode propagates vertically along the field lines, forms shocks and remains always within the flux tube. It might deposit effectively the energy of the driver into the chromosphere. When the driver oscillates with a high frequency, above the cut-off, non-linear wave propagation occurs with the same dominant driver period at all heights. At low frequencies, below the cut-off, the dominant period of oscillations changes with height from that of the driver in the photosphere to its first harmonic (half period) in the chromosphere. Depending on the period and on the type of the driver, different shock patterns are observed. ",
        "introduction": "\\label{sect:Introduction}   Solar magnetic structures show a continuous distribution of fluxes and sizes. In active regions, the field strength concentrated in large-scale structures (sunspots and pores) can be as large as $2-4$ kG. In plage and network areas, the average flux decreases to several hundreds Gauss. However, individual flux tubes in these areas can have intrinsic magnetic field strength similar to that of solar pores, \\ie\\ $1-2$ kG. The decrease of the average flux is caused mainly by the decrease of the size of the magnetic features, \\ie\\ their filling factor.  The observed photospheric brightness of the magnetic structures is in close relationship with their flux and size (see for example \\opencite{Berger+etal2007}). Large-scale features like sunspots and pores appear dark in photospheric observations. Plage and facular flux tubes appear as bright features (see for example Figure 7 in \\opencite{Berger+etal2007}).  Due to their larger size, sunspots show horizontal gradients of the magnetic field strength and inclination. Instead, small-scale flux tubes possess a more homogeneous field with a rather sharp transition between the magnetized and non-magnetized surroundings. From the point of view of wave propagation, it is important to realize that the temperature at a given geometrical height is different in the different structures. Temperature and magnetic field define the characteristic wave propagation speeds, \\ie\\ sound speed, $c_S$, and Alfv\\'en speed, $v_A$, as well as the acoustic cut-off frequency, $\\omega_C$, and the height of the transformation layer where $c_S$ is equal to $v_A$ and different wave modes can interact. These parameters can be rather different in large-scale dark structures and small-scale bright structures. Thus, it is not surprising  that the observed properties of waves in sunspots, network, and plage areas are rather different.  \\subsection{Sunspots and Pores}  Several decades of studies of sunspot oscillations can be summarized as follows. Waves observed at different layers of the atmosphere of sunspots umbrae and penumbrae seem to be the manifestation of the same phenomenon (see \\eg\\/ \\opencite{Maltby+etal1999}, \\citeyear{Maltby+etal2001}, \\opencite{Brynildsen+etal2000}, \\citeyear{Brynildsen+etal2002}, \\opencite{Christopoulou+etal2000}, \\citeyear{Christopoulou+etal2001}, \\opencite{Rouppe+etal2003}, \\opencite{Tziotziou+etal2006}). In the photosphere, five-minute oscillations are observed in sunspot umbrae and penumbrae and in solar pores \\cite{Bogdan+Judge2006}. Wave propagation is linear with amplitudes of several hundreds of \\ms\\ and is parallel to the magnetic field lines (see \\eg\\/ \\opencite{Gurman+Leibacher1984}, \\opencite{Lites1984}, \\opencite{Collados2001}, \\opencite{Centeno+etal2006a}, \\opencite{Bloomfield+etal2007b}). Magnetic field oscillations of a few Gauss are detected in the deeper layers and attempts have been made to interpret these oscillations in terms of fast and slow MHD waves (\\eg\\/ \\opencite{Ruedi+Solanki+Stenflo+Tarbell+Scherrer1998}, \\opencite{Lites+Thomas+Bogdan+Cally1998}, \\opencite{BellotRubio+Collados+RuisCobo+RodriguezHidalgo2000}, \\opencite{Khomenko+Collados+BellotRubio2003}).  In the chromosphere, the dominant period of oscillations decreases and shock waves are observed with a three-minute periodicity in sunspot umbrae and in pores (\\opencite{Lites1984}, \\citeyear{Lites1986}, \\citeyear{Lites1988}, \\opencite{Socas-Navarro+etal2000}, \\opencite{Centeno+etal2006a} \\opencite{Tziotziou+etal2007}). The amplitudes of the observed shocks are in the range of $5-15$ \\kms, depending on the size of the structure. In pores, shocks are weaker \\cite{Centeno+etal2006b}. Umbral flashes also show a three-minute periodicity. In contrast, in the penumbra, the shock behavior of waves with five-minute periodicity is detected for running penumbral waves even at chromospheric heights \\cite{Tziotziou+etal2006}. Waves observed in the transition region and the corona conserve these properties (\\opencite{DeMoortel+etal2002}, \\opencite{Marsh+Walsh2005}, \\citeyear{Marsh+Walsh2006}, \\opencite{Maltby+etal1999}, \\citeyear{Maltby+etal2001}, \\opencite{Brynildsen+etal2000}, \\citeyear{Brynildsen+etal2002}, \\opencite{Christopoulou+etal2000}, \\citeyear{Christopoulou+etal2001}). Wave propagation from the photospheric pulse into the chromosphere and higher layers is along the magnetic field lines for both three and five minutes perturbations.   \\subsection{Plages, Facular regions, and Network}   Bright magnetic field structures like plages, facular regions, and network manifest five-minute linear oscillations in the photosphere and the oscillations maintain the five minute periodicity in the chromosphere and larger heights (\\opencite{Lites+Rutten+Kalkofen1993}, \\opencite{Krijer+etal2001}, \\opencite{DePontieu+etal2003}, \\opencite{Centeno+etal2006b}, \\opencite{Bloomfield+etal2006}, \\opencite{Veccio+etal2007}). Shock occurrence is much lower in these quiet Sun regions with enhanced flux (\\opencite{Krijer+etal2001}, \\opencite{Centeno+etal2006b}). The velocity oscillation amplitudes in the chromosphere do not reach values as large as in sunspots and stay within $2-4$ \\kms. Waves at photospheric and chromospheric layers seem to be correlated and vertical propagation along the vertical magnetic fields seem to dominate (\\opencite{Krijer+etal2001}, \\opencite{Centeno+etal2006b}, \\opencite{Veccio+etal2007}). Curiously, there are areas with enhanced three-minute power surrounding the strong magnetic field concentrations, possibly related to flux tube canopies (aureoles, see \\opencite{Braun+Linsday1999}, \\opencite{Thomas+Stanchfield2000}, \\opencite{Krijer+etal2001}, \\opencite{Judge+etal2001}). In the higher coronal layers, loops connecting sunspots oscillate with three minutes and the loops connecting network and plage regions oscillate with five minute periods \\cite{DeMoortel+etal2002}.  Thus, the wave properties are distinctly different in large-scale and small-scale magnetic features. The difference in oscillation properties is present through the whole solar atmosphere and should be a consequence of the thermal and magnetic properties of these structures.   \\subsection{Theoretical Models}  A number of theoretical mechanisms have been proposed to explain the oscillation spectra at different heights in magnetic structures and to interpret the observed oscillations in terms of MHD waves. \\inlinecite{Zhugzhda+Locans1981}, \\inlinecite{Gurman+Leibacher1984}, and \\inlinecite{Zhugzhda2007} argue that the observed spectrum of sunspot umbral oscillations in the chromosphere is due to the temperature gradients of the atmosphere acting as an interference filter for linear three-minute period acoustic waves. On the other hand, it has been demonstrated by \\inlinecite{Fleck+Schmitz1991} that the change of the period with height from five to the three minutes is a basic phenomena occurring even in an isothermal atmosphere for linear waves due to the resonant excitation at the atmospheric cut-off frequency. In the case of the solar atmosphere, the temperature minimum gives rise to the three-minute cut-off period. Later, the same authors found that, due to non-linear shock interaction and overtaking, the frequency of acoustic oscillations also changes with height (\\opencite{Fleck+Schmitz1993}). The response of the solar atmosphere to an adiabatic shock wave leads to a natural appearance of the three-minute peak in the oscillation power spectra, in the case the underlying photosphere has a five-minute periodicity. A third mechanism was shown to explain the three-minute periodicity of calcium bright grains in the quiet Sun: with a spectrum of periods in the photosphere with a peak at five minutes, this distribution changes to one with a maximum at the acoustic cut-off because at longer periods the energy falls off exponentially with height (\\opencite{Carlsson+Stein1997}). These mechanisms may explain the shift with height of the period of oscillations from five to three minutes observed in sunspots.  The question remains: why is there not such a shift in small-scale structures of plages and network regions? \\inlinecite{DePontieu+etal2004} argue that the inclination of the magnetic field may play an important role. If (acoustic or slow MHD) waves have a preferred direction of propagation defined by the magnetic field, the effective cut-off frequency is lowered by the cosine of the inclination angle with respect to the vertical. This allows evanescent waves to propagate. It should be recalled, however, that mainly vertical propagation is observed in the photosphere and chromosphere in plage regions (see \\eg\\/ \\opencite{Krijer+etal2001}, \\opencite{Centeno+etal2006b}), thus it is unclear whether the mechanism suggested by \\inlinecite{DePontieu+etal2004} is at work. Alternatively, the decrease of the effective acoustic cut-off frequency can be produced by taking into account the radiative losses of acoustic oscillations \\cite{Roberts1983, Centeno+etal2006b}. If the radiative relaxation time is effectively small (as expected in small-scale magnetic structures, transparent to radiation), the cut-off frequency can be reduced and evanescent waves can propagate.   All works cited above apply the theory of acoustic waves in idealized atmospheres in order to explain the observed oscillation properties in magnetic structures. In magnetized atmospheres, different wave types can exist and different modes can be observed depending on the magnetic field configuration and the height of the transformation layer ($c_S=v_A$) relative to the height of the formation of the spectral lines used in observations. \\inlinecite{Bogdan+etal2003}, \\inlinecite{Rosenthal+etal2002}, \\inlinecite{Hasan+Ulmschneider2004}, Hasan {\\it et al.} (2003, 2005), \\nocite{Hasan+etal2003} \\nocite{Hasan+etal2005} addressed this question in their simulations of magneto-acoustic waves in non-trivial magnetic configurations. They pointed out the importance of defining the mode type and the mode conversion height when interpreting the observations. Due to mode mixing at the transformation layer, no correlation may be observed between perturbations below and above the $c_S=v_A$ height.  All these arguments point of towards the need to have realistic, self-consistent, and systematic models of the wave spectrum in magnetic structures. It is necessary to include non-trivial magnetic field configurations in order to allow the existence of different mode types and mode conversion. Non-linearities should be taken into account for waves propagating into the chromosphere. The thermal structure should be realistic in order to make possible the reflection of waves with frequencies below the cut-off. Radiative damping also plays an important role, in particular to estimate the amount of energy that may deposited in the higher atmosphere due to waves. In the present work we address several of the above questions. We perform non-linear numerical modeling of magneto-acoustic waves in small-scale magnetic flux tubes with internal structure. We consider waves with realistic periods of three and five minutes, which has never been done before for non-trivial magnetic configurations. We study the mode transformation and oscillation spectra at different heights from 0 to 2000 km as a function of the type of the driver and its period. At present, radiative damping of oscillations is not taken into account. This question will be addressed in a separate paper.   \\begin{figure} \\centerline{\\includegraphics[width=1.0\\textwidth]{fig1.eps}} \\caption{Second-order flux tube magnetostatic model, calculated after Pneuman, Solanki, and Stenflo (1986). The tube radius is equal to 100 km in the photosphere. Left panel: magnetic field lines and contours of plasma $\\beta$ (defined as $c_S^2/v_A^2$), with values indicated on each contour. Right panel: distribution of the magnetic field strength.} \\label{fig:tube} \\end{figure} ",
        "conclusions": "We have performed simulations of waves in flux tubes with non-trivial magnetic field configuration. The model flux tube has horizontal and vertical variations of magnetic field strength and gas pressure and of the ratio between the characteristic wave speeds $v_A$ and $c_S$. These properties have allowed us to study wave propagation and mode transformation inside the flux tube. Waves are excited by a photospheric driver with a period of three and five minutes, for the first time with a magnetic field configuration of this kind. Different wave patterns are observed depending on the period and on the type of the driver. The following items summarize our findings: \\begin{itemize} \\item Horizontal motions of the flux tube at the bottom photospheric boundary generate a slow magneto-acoustic mode inside the tube and a surface mode at the magnetic/field free interface. \\item After the slow magneto-acoustic mode and the surface mode reach the height where $v_A \\approx c_S$, their energy is effectively transformed into a slow acoustic mode in the high atmosphere where $v_A > c_S$. Only a small part of the driver energy is returned back to the photosphere by the fast magneto-acoustic mode. \\item The slow acoustic mode propagates vertically along the magnetic field lines in the atmosphere where $v_A > c_S$, forms shock waves above 1000 km with amplitudes of $5-15$ \\kms\\ and remains always within the same flux tube. Thus, it can deposit effectively most part of the energy of the driver into the chromosphere. \\item If the frequency of the horizontal driver is above the acoustic cut-off, non-linear wave propagation in the tube occurs with the dominant period of the driver at all heights. Non-linear effects produce higher harmonics, but their amplitude is small. \\item If the frequency of the horizontal driver is below the acoustic cut-off, the dominant period of the longitudinal velocity changes with height from five to three minutes due to non-linear generation of the higher harmonics. \\item The five-minute vertical perturbation of the flux tube at the bottom photospheric boundary generates a fast acoustic mode that propagates upward through the transformation $v_A = c_S$ layer, without changing its nature and forming only eventual shocks in the chromosphere. \\item After reaching the stationary state in the simulations with a vertical driver at 300 seconds, both three and five minute oscillations co-exist inside the flux tube at chromospheric heights. The dominant period at the axis is five minutes and the off-axis dominant period is three minutes. \\end{itemize}  Our simulations suggest once again that the properties of waves observed in magnetic structures are the direct consequence of their magnetic configuration, temperature, magnetic field strength, and the height where the mode transformation occurs. When analyzing observations, it is important to know whether they correspond to the level below or above the layer $v_A = c_S$. If the motions that excite oscillations inside flux tubes are purely horizontal, no correlation may be observed between velocity variations measured in the photosphere and the chromosphere. This is because only an acoustic mode can be detected in observations (at least at disc center) since it produces significant vertical velocities, unlike the slow magnetic mode existing in the photosphere. The acoustic mode is only generated above $700-1000$ km in our simulations with a horizontal driver. In contrast, if the driver that excites oscillations has a vertical component, the acoustic mode is generated already in the photosphere. In this case, the oscillations in vertical velocity are coherent at photospheric and chromospheric heights.  In general, both simulations with horizontal and vertical drivers with 300-second periodicity show the change of the dominant wave period with height. This property is different from observations of plage and network regions where the five minute periodicity is preserved also in the chromosphere (\\opencite{Lites+Rutten+Kalkofen1993}, \\opencite{Krijer+etal2001}, \\opencite{DePontieu+etal2003}, \\opencite{Centeno+etal2006b}, \\opencite{Bloomfield+etal2006}, \\opencite{Veccio+etal2007}). While our study may help to identify the wave modes observed in small-scale magnetic structures, further work is needed in order to explain the behavior of the oscillatory spectra with height."
    },
    "0710/0710.2333_arXiv.txt": {
        "abstract": "We model the massive dark object at the center of the Galaxy as a Schwarzschild black hole as well as Janis-Newman-Winicour naked singularities, characterized by the mass and  scalar charge  parameters, and study  gravitational lensing (particularly time delay, magnification centroid, and total  magnification)  by them. We find that the lensing features are qualitatively similar (though  quantitatively different) for the Schwarzschild black holes, weakly naked, and marginally strongly naked singularities. However, the lensing  characteristics of strongly naked singularities are qualitatively very different from those due the Schwarzschild black holes. The images produced by  Schwarzschild black hole lenses and weakly naked and marginally strongly naked singularity lenses always have positive time  delays. On the other hand, the strongly naked singularity lenses can give rise to images with positive, zero, or negative time delays. In particular, for a large angular source position  the direct image (the outermost image on the same side as the source) due to strongly naked singularity lensing  always  has negative time delay. We also found  that the  scalar field  decreases the time delay and increases the magnitude of magnifications of images; this result could have important implications for cosmology. As the Janis-Newman-Winicour metric also describes the exterior gravitational field of a scalar star, naked singularities as well as scalar star lenses, if these exist in nature, will serve as more efficient cosmic telescopes than  regular gravitational lenses. ",
        "introduction": "Introduction}  A naked (visible) singularity is defined as a spacetime singularity which can be seen by some observer and also lies to the future of some point of the spacetime \\cite{Pen98}. The well-known weak cosmic censorship hypothesis (WCCH) of Penrose essentially states  that, generically, spacetime singularities of physically realistic gravitational collapse are hidden within black holes \\cite{Pen98,Wal97}. The concept of visible singularities is objectionable to many scientists, as their existence is thought to have alarming  astrophysical implications. On the other hand, a failure of the WCCH will give us the great opportunity  to probe the extremely strong gravitational fields that will help in the discovery of the physical laws of quantum gravity. Despite many industrious efforts,  we are still far from having a general proof (or disproof) of this hypothesis,  and Penrose \\cite{Pen98} expected that  radically new mathematical techniques might be required for this purpose. As a proof or  disproof  of this hypothesis appears to be  inordinately difficult, it may be easier to find a persuading  counterexample to demonstrate that the hypothesis is not correct.  Numerous diligent efforts have been put in this direction in the last four decades; however, we still do not have a single convincing counterexample  to the  WCCH (see references in \\cite{Pen98,Wal97,Vir99}). In a seminal review,  Penrose \\cite{Pen98} concluded that  the question of the cosmic censorship  is still very much open and considered  this to be possibly the  most important unsolved problem in classical general relativity.   Given that we  have neither a proof (or disproof) nor a convincing counterexample  of the WCCH, it is important to explore whether or not this hypothesis could be tested observationally. To this end, Virbhadra {\\it et al.} \\cite{VNC98} introduced  a theoretical research project using the gravitational lensing phenomena and encouraging results came out of that. Further, Virbhadra and Ellis \\cite{VE00} obtained a new gravitational lens equation that allows  large deflection of light and therefore it can be used to study strong gravitational field lensing. They used this lens equation to study the gravitational lensing due to light deflection close to the photon sphere of the supermassive ``black hole'' at the center of the Galaxy. They found that the presence of a photon sphere gives rise to   a theoretically infinite sequence of  highly demagnified images on both sides of the optical axis (the line joining the lens and the observer) and they termed these {\\em relativistic images}. Virbhadra and Ellis \\cite{VE02} further extended the previous studies of Virbhadra {\\it et al.} \\cite{VNC98} in detail. They also organized the investigations by classifying   naked singularities in two groups: {\\em weakly naked} and {\\em strongly naked} singularities. They modeled  massive dark objects at the  centers of a few galaxies as  Schwarzschild black holes and Janis-Newman-Winicour weakly as well as strongly naked singularities, and studied  gravitational lensing by  them. The Schwarzschild black holes as well as weakly naked singularities have qualitatively similar lensing characteristics : Both have one Einstein ring and no radial critical curves. On the other hand, the strongly naked singularities have  qualitatively different lensing features, i.e., they give rise to two  or none  Einstein rings and one radial critical curve.  After publication of these results \\cite{VNC98, VE00,CVE01,VE02}, there has been  a growing curiosity in black hole lensing and many interesting papers have appeared in  last few years (see \\cite{Per,Zaketal,Amoreetal, Bozzaetal, Eioraetal, KNPF, MajuMukh,KP05, KPIW, ExoticLens, BHLens, Scalar} and references therein).    In this paper, we study the time delay, magnification centroid, and total magnification of images due to  gravitational lensing by Schwarzschild black holes and  Janis-Newman-Winicour naked singularities. One of the most striking results in this paper is that the strongly naked singularities can give rise to images  with negative time delays.  We use geometrized units (i.e., $G=1, c=1$) throughout this paper; however, we finally compute  time delays in  terms of minutes. We use MATHEMATICA \\cite{Math} for computations. ",
        "conclusions": "Discussion and Conclusion}  The naked singularities are classified in three categories: {\\em WNS}, {\\em MSNS}, and {SNS}. We modeled the Galactic MDO as the SBH, and Janis-Newman-Winicour WNS, MSNS and SNS lenses, and studied point source gravitational lensing by them. We found that the  gravitational lensing effects due to the SBH, WNS, and MSNS are qualitatively similar (but these differ slightly quantitatively) to  each other; however, they differ qualitatively from SNS lensing. Therefore, it will be easier to observationally differentiate a SNS (compared to  a WNS or a MSNS) from a SBH.   {\\em SBH, WNS and MSNS lensing}. These do not give rise to any radial caustics; however, they do produce one Einstein ring when the lens components are aligned (i.e. $\\beta=0$).  When $\\beta$ increases, the Einstein ring splits into two images, one on each side of the optical axis. The time delays for both images are positive for all values of $\\beta$. For a given value of $\\beta$, a decrease in $\\nu$, i.e., an increase in $(q/M)^2$, decreases the absolute angular image positions, time delays, magnification centroid, and magnification centroid shift; however, it increases the magnitude of  magnifications of images and therefore also the total  magnification. The differences are however very small. The deflection angle $\\hat{\\alpha}$ becomes unboundedly large as the impact parameter $J \\rightarrow J_{ps}$ for the SBH and WNS, and $ J \\rightarrow 0$ for the MSNS \\cite{VE02}.    {\\em SNS lensing}. There are  two types of lensing in this category. In the first, for example for the case of  $\\nu = 0.02$,  there are double concentric Einstein rings (when $\\beta = 0$) and one  radial critical curve (when $\\beta \\approx 210.2934$   arcseconds). As the angular position of the source  increases from the alignment position of the lens components (i.e., $\\beta = 0$), the two Einstein rings ``break'' into four images,  giving two images on each side of the optical axis.  The separation between images on the same and opposite side from the source, respectively, increases and decreases as $\\beta$ increases, and eventually the two images on the opposite side from the source coalesce to form a single highly magnified image. For any  further increase in $\\beta$, there are only two images on the same side as the source. The second category of SNS lensing, for example for  $\\nu=0.001$, there is one RC; however, there is no Einstein ring when $\\beta=0$. Moreover, there is no image for small values of $\\beta$. As $\\beta$ increases,  a highly magnified image (RCC) first appears  on the same side as the source. A further increase in $\\beta$ splits this into two images and the separation between them keeps increasing  (both images remaining on  the same side as the source). The time delay of  images of SNS lensing may be positive, zero, or negative depending on the values of $\\nu$ and $\\beta$. However, the time delay of a direct image is negative for any SNS lensing if $\\beta$ is large. As shown in \\cite{VE02}, the deflection angle $\\hat{\\alpha}$ approaches $-\\pi$ as the impact parameter approaches its minimum value zero. Therefore, if a light ray is sent radially toward a SNS, it will ``bounce''; however, it remains to be computed whether or not the ``reflected'' light has enough magnification to be observed by present instruments or by those likely to be available in near future. This may serve as a crucial test for the existence of SNS.   The naked singularity lensing gives rise to images of smaller time delay and stronger magnification than the black hole lensing. Therefore, if naked singularities indeed exist in nature, then these will serve as {\\em better cosmic telescopes} and  will help probe the universe more efficiently. The results obtained in this paper also helps understand the effects of scalar field on gravitational lensing, which could have valuable implications for research in  cosmology.   The Janis-Newman-Winicour metric also describes the exterior gravitational field of a scalar star. Therefore, results obtained in this paper for naked singularities are also applicable to scalar stars. The scalar star however must be compact enough for the images not to be obstructed.   The  metric we considered in this paper may or may not be physically realistic. However,  gravitational lensing studies with this metric serves as a  stepping stone to understand the distinctive lensing features of black holes and naked singularities. It would be indeed of great astrophysical significance to obtain distinguishing  {\\em qualitative} lensing characteristics of Kerr black holes and naked singularities, so that the  weak cosmic censorship hypothesis could be tested observationally  without any ambiguity."
    },
    "0710/0710.2105_arXiv.txt": {
        "abstract": "Utilizing {\\it Chandra} X-ray observations in the All-wavelength Extended Groth Strip International Survey (AEGIS) we identify 241 X-ray selected Active Galactic Nuclei (AGNs, $L_{2-10{\\rm keV}} > 10^{42}$ ergs s$^{-1}$) and study the properties of their host galaxies in the range $0.4 < z < 1.4$.  By making use of infrared photometry from Palomar Observatory and $BRI$ imaging from the Canada--France--Hawaii Telescope, we estimate AGN host galaxy stellar masses and show that both stellar mass and photometric redshift estimates (where necessary) are robust to the possible contamination from AGNs in our X-ray selected sample.  Accounting for the photometric and X-ray sensitivity limits of the survey, we construct the stellar mass function of X-ray selected AGN host galaxies and find that their abundance decreases by a factor of $\\sim$2 since $z \\sim 1$, but remains roughly flat as a function of stellar mass.  We compare the abundance of AGN hosts to the rate of star formation quenching observed in the total galaxy population.  If the timescale for X-ray detectable AGN activity is roughly 0.5--1 Gyr---as suggested by black hole demographics and recent simulations---then we deduce that the inferred AGN ``trigger'' rate matches the star formation quenching rate, suggesting a link between these phenomena.  However, given the large range of nuclear accretion rates we infer for the most massive and red hosts, X-ray selected AGNs may not be {\\em directly} responsible for quenching star formation. ",
        "introduction": "Recent observations of the galaxy population and its evolution since $z \\approx 2$ reveal a pattern in which the most massive galaxies appear to shut down star formation activity at early times with increasingly less massive galaxies following later \\citep[e.g.,][]{juneau05, treu05, bundy06, borch06, cimatti06}.  This pattern of ``quenching'' in the star formation history of galaxies---thought to be largely responsible for the growing abundance of galaxies on the red sequence \\citep[e.g.,][]{faber07}---is commonly referred to as ``downsizing'' \\citep{cowie96}.   Given that the dark matter halos of galaxies are expected to assemble hierarchically in the $\\Lambda$CDM paradigm, understanding the physical mechanisms responsible for downsizing remains an important challenge. We seek a process capable of quenching star formation, driving galaxies onto the red sequence, and preventing further star formation.  Work with the DEEP2 Galaxy Redshift Survey \\citep{davis03} has shown that such a process must operate over a range of environmental densities \\citep{bundy06, cooper07, gerke07}, suggesting an internal component to quenching that acts in addition to the suppression of star formation expected in high-density environments.  Meanwhile, theoretical work has focused on one such internal process, namely the potential role played by AGN feedback \\citep{silk98}, both as a way of explosively initiating the quenching event \\citep{granato04, scannapieco05, hopkins05b} and preventing hot gas in already passive galaxies from cooling to form stars \\citep{croton06, bower06, cattaneo06, de-lucia07}.  These scenarios remain largely untested because observational evidence has been difficult to obtain.  One of the most promising ways forward is to examine the properties of AGN host galaxies and search for signatures of this feedback.  Such observations are challenging however because no selection method finds all galaxies that host active nuclei \\citep{mushotzky04} and brighter AGNs (quasars) can easily outshine their hosts.  With such difficulties in mind, previous studies suggest that most lower luminosity AGNs tend to be found in massive, mostly spheroidal galaxies \\citep{dunlop03, kauffmann03b, grogin05, pierce07} although some ``transition'' sources exhibit disturbed morphologies \\citep[e.g.,][]{canalizo01, hutchings06, conselice07}.  AGNs detected through emission line diagnostics and X-rays appear to prefer host galaxies on the red sequence and ``green valley'' \\citep{nandra07, martin07, salim07}.   Further progress on testing the link between AGN activity and the downsizing of star formation requires that we understand the stellar mass distribution and redshift evolution of AGN hosts as compared to evolutionary patterns in the general galaxy population.  In this paper we use deep {\\em Chandra} observations to identify AGN hosts in the AEGIS field \\citep{davis07} where spectroscopic and photometric redshifts (photo-$z$'s) as well as infrared photometry provide reliable stellar mass estimates out to $z = 1.4$.  We then compute the AGN host stellar mass function and use it to estimate the rate at which AGN activity is triggered in the galaxy population.  We will show that this mass-dependent rate is consistent with the rate of star formation quenching of all galaxies at the same epochs at which the AGN activity is observed.    The structure of the paper is as follows.  We describe the multi-wavelength observations and properties of the sample in $\\S$\\ref{data} and $\\S$\\ref{samples}.  The way in which stellar masses and restframe colors are determined is given in $\\S$\\ref{methods} while our methods for constructing mass functions and the resulting AGN host mass function are presented in $\\S$\\ref{mfn}.  In $\\S$\\ref{discussion} we present evidence for a link between AGN activity and quenching in the context of estimates of the X-ray AGN timescale and explore whether feedback from X-ray selected AGNs causes quenching.  We summarize in $\\S$\\ref{summary}.  Where necessary, we assume a standard cosmological model with $\\Omega_{\\rm M}=0.3$, $\\Omega_\\Lambda=0.7$, $H_0=70 h_{70}$ km~s$^{-1}$~Mpc$^{-1}$. ",
        "conclusions": ""
    },
    "0710/0710.5513_arXiv.txt": {
        "abstract": "We present the results of Chandra observations of 13 optically-selected clusters with $0.6 < z < 1.1$, discovered via the Red-sequence Cluster Survey (RCS). All but one are detected at S/N$>3$; though three were not observed long enough to support detailed analysis. Surface brightness profiles are fit to $\\beta$-models. Integrated spectra are extracted within $\\rm{R}_{2500}$, and $T_X$ and $L_X$ information is obtained. We derive gas masses and total masses within $\\rm{R}_{2500}$ and $\\rm{R}_{500}$. Cosmologically corrected scaling relations are investigated, and we find the RCS clusters to be consistent with self-similar scaling expectations. However, discrepancies exist between the RCS sample and lower-$z$ X-ray selected samples for relationships involving $L_X$, with the higher-$z$ RCS clusters having lower $L_X$ for a given $T_X$. In addition, we find that gas mass fractions within $\\rm{R}_{2500}$ for the high-$z$ RCS sample are lower than expected by a factor of $\\sim2$. This suggests that the central entropy of these high-$z$ objects has been elevated by processes such as pre-heating, mergers, and/or AGN outbursts, that their gas is still infalling, or that they contain comparatively more baryonic matter in the form of stars. Finally, relationships between red-sequence optical richness (\\bgc) and X-ray properties are fit to the data. For systems with measured $T_X$, we find that optical richness correlates with both $T_X$ and mass, having a scatter of $\\sim30\\%$ with mass for both X-ray and optically-selected clusters. However, we also find that X-ray luminosity is not well correlated with richness, and that several of our sample appear to be significantly X-ray faint. ",
        "introduction": "}   The Extended Medium Sensitivity Survey~\\citep[EMSS]{gioia90} sparked renewed interest in the ongoing search for clusters of galaxies at high redshift. Since then, numerous high redshift surveys have been carried out, both optically~\\citep[and others]{gilbank04,donahue02,postman,bower}, and in the X-ray~\\citep[and others]{valtchanov,bauer,wilkes}. The motivations for such searches are multifaceted, but the most compelling of these are cosmological in nature.  Clusters of galaxies are an important source of information about the underlying cosmology of the universe. They are considered to be essentially ``closed boxes'', meaning that the primordial matter that they were initially assembled from has remained trapped in their deep potential wells since they were formed. This makes them ideal objects with which to study galaxy formation and evolution. In addition, clusters are the largest virialized objects in the universe.  By virtue of this fact we are able to, through high redshift samples, investigate the growth of large scale structure. A firm knowledge of the evolution of the cluster mass function would provide an enormous contribution to constraining cosmological parameters such as $\\sigma_8$ (the normalization of the density perturbation spectrum) and $w$~\\citep[the dark energy equation of state; e.g.,][]{voit05}.  Two ingredients are required to achieve this goal. First, a statistically significant sample of clusters in multiple redshift bins is needed. Second, reliable mass estimates of the clusters in that sample must be obtained. Difficulties in reaching the first requirement include the vast amount of telescope time required to carry out such a search in the X-ray, and the propensity for false detections due to projection effects in optical surveys. The primary challenge in reaching the second part of this goal is again the high cost of observing time to achieve either X-ray or dynamical mass estimates.  The Red-sequence Cluster Survey~\\citep[RCS;][]{gladders00,gladders05,yee07} has attempted to evade such difficulties -- RCS is an optical survey which uses the color-magnitude relationship of cluster ellipticals to find galactic overdensities in small slices of redshift space. This technique has been estimated to bring false detection rates down to $\\sim 5-10\\%$~\\citep{gladdersphd,blindert07,cohn07}. The chosen filters ($R_c$ and $z'$), optimize this finding algorithm for the redshift range $0.2<z<1.2$ and provide photometric redshift information with accuracies of $\\sim 10 \\%$. In addition, optical richness information is immediately available from the survey data and, if sufficiently calibrated, this information could provide a highly efficient way to estimate the masses of cluster candidates.  The first phase of the Red-Sequence Cluster Survey~\\citep[RCS-1;][]{gladders05}, from which our cluster sample was drawn, covers 90 square degrees and was performed at CFHT and CTIO. RCS-1 has identified 6483 cluster candidates in the redshift range $0.2 <z< 1.2$, over 1000 of which are at least as optically rich as Abell class 0 clusters~\\citep{gladders05}.  The motivations for this work are to characterize high-redshift optically-selected cluster samples, probe cluster evolution, and move forward in attempts to calibrate a robust relationship between optical richness and cluster mass. This paper presents a detailed analysis of the {\\it{Chandra}} data of thirteen RCS-1 clusters with redshifts in the range $0.6<z<1.1$ Our analysis investigates the temperatures and gas distributions of ten of these clusters, and provides mass estimates for use in the calibration of relationships between optical richness and cluster mass.  We also use our results to investigate the X-ray scaling laws of our sample, and thereby probe redshift evolution in these relationships. To facilitate comparisons between the RCS clusters and lower-redshift X-ray selected samples, we make use of our previous {\\it{Chandra}} analysis of the Canadian Network for Observational Cosmology (CNOC) subsample of the Extended Medium Sensitivity Survey (EMSS)~\\citep{hickscnoc,yee96,gioia90}. This sample, with redshifts in the range $0.1<z<0.6$, was chosen based on X-ray luminosity \\citep[$L_X\\geq2\\times10^{44}$ erg s$^{-1}$;][]{gioia90}.   This paper is organized as follows:  In Sections~\\ref{s:obs} and~\\ref{s:sig}, we introduce our sample and describe the basic properties of our data.  In Sections~\\ref{s:sb2} and~\\ref{s:spec}, we investigate gas distributions and obtain cluster temperatures.  We derive masses for our sample in Section~\\ref{s:mass}.  High-$z$ X-ray scaling relationships are examined in Section~\\ref{s:laws2}, while correlations between optical richness and cluster X-ray properties are explored in Section~\\ref{s:or2}.  In Section~\\ref{s:selection} we investigate possible sources of bias in cluster sample selection.  A summary and discussion of our results is presented in Section~\\ref{s:sum}. Unless otherwise noted, this paper assumes a cosmology of $\\rm{H}_0=70~\\rm{km}~\\rm{s}^{-1}~\\rm{Mpc}^{-1}$, $\\Omega_{\\rm{M}}=0.3$, and $\\Omega_{\\Lambda}=0.7$.  All errors are quoted at 68\\% confidence levels. ",
        "conclusions": "}  We have performed an in-depth X-ray investigation of 13 high redshift ($0.6<z<1.1$) optically selected clusters of galaxies from the Red-sequence Cluster Survey (RCS; Table~\\ref{table1}). All but one of these clusters was detected by {\\it{Chandra}} at a signal-to-noise ratio of greater than 3 (Table~\\ref{table2}), though two additional clusters in our sample (RCS1417+5305 and RCS2112-6326), did not posess enough signal to support further analysis. Initial imaging of the objects reveals that the RCS sample spans a wide range in cluster morphology (Figure~\\ref{fig1}), from very regular objects (e.g., RCS1419+5326) to more disturbed systems (e.g., RCS2112-6326).  Surface brightness profiles were extracted for ten clusters in $1-2\\arcsec$ annular bins, and were reasonably well fit by single $\\beta$ models (Figure~\\ref{fig2} and Table~\\ref{table3}). Cluster emission was modeled with XSPEC, beginning with a spectral extraction region of 300 kpc radius. The results of single temperature spectral fits, combined with best fit $\\beta$ models, were used to determine $\\rm{R}_{2500}$. Spectra were re-extracted from regions of that radius for further temperature fitting and $\\rm{R}_{2500}$ luminosity estimates, until extraction regions and R$_{2500}$ estimates were in agreement. Results of this process are given in Table~\\ref{table4}. We have also used the $\\sigma-T_X$ relationship to compare the X-ray temperatures of three of our objects to currently available velocity dispersions~\\citep{gilbank07,gilbankprep}. We find consistency in all cases (Table~\\ref{table6}), suggesting that these three objects are at least relatively undisturbed.  Using the results of both spectral fitting and surface brightness modeling, X-ray masses were calculated for ten clusters in our sample out to $\\rm{R}_{2500}$ (Table~\\ref{table7}), with extrapolation to $\\rm{R}_{500}$ (Table~\\ref{table8}). Canonical X-ray scaling laws were investigated for nine clusters and compared to those of the moderate redshift ($0.1<z<6$) CNOC sample (Tables~\\ref{table9} and~\\ref{table10}). For the $L_X-T_X$ relationship (Figure~\\ref{fig3}), both RCS and CNOC fits have slopes consistent with self-similar predictions; however, their normalizations disagree. Interestingly, the slope of our combined RCS-CNOC sample agrees with that of X-ray selected clusters at similar redshift~\\citep{ettori04}, suggesting that evolution in the normalization of the $L_X-T_X$ relationship may lie behind the observed steeper slopes. Results from $L_X-$M$_{\\rm{tot}}$ fits are qualitatively very similar to those of the $L_X$-$T_X$ relationship (Figure~\\ref{fig4}), RCS and CNOC slopes are consistent with self-similarity at R$_{2500}$, but disagree in normalization due to differing ICM densities.  The most notable outcome of our mass estimations is that the $\\Delta=2500$ gas mass fractions of RCS clusters are lower than expected by a factor of $\\sim2$~\\citep{vikhlinin06}. Low gas mass fractions are also reported in the findings of other high redshift cluster studies, both observation and theory~\\citep{lubin02,ettori04b,kravtsov05,sadat,ettori06,nagai07,afshordi07}. Physical explanations for low gas fractions include suggestions that much of the baryonic mass has been converted into stars, that the gaeous component of these clusters is still infalling, or that some mechanism (i.e., galaxy formation, AGN, mergers, radio jets) is responsible for raising the entropy of the gas.  Though we do see some evidence for higher entropy in the RCS sample (Figure~\\ref{fig7}), it may not be enough to explain the entire f$_{\\rm{gas}}$ discrepancy. Further study will have to be undertaken to determine whether the overall lower gas fractions that permeate this sample are ubiquitous at high redshift, or an outcome of sample selection.   Explanations aside, the growing evidence that massive ($T_X \\sim6$ keV) clusters may have an evolving or broad range of central gas mass fractions may have important consequences for the interpretation of future cluster surveys in the microwave and X-ray bands, which select in part on the basis of central gas density.  Scatter in this parameter will tend to reduce completeness and, if not properly accounted for, inject bias in the cluster samples such surveys produce.  In addition, since SZ mass determinations tend to rely on the assumption of a constant gas mass fraction~\\citep{reid06}, complimentary data may be required to avoid systematic errors in SZ total mass estimates.   Using red-sequence optical richness measurements of both samples the relationships between \\bgc~and global cluster properties ($T_X$, $L_X$, $\\rm{M}_{2500}$, and M$_{200}$) were investigated (Table~\\ref{table11}). We find that \\bgc~is poorly correlated to X-ray luminosity, with average scatter in the relationships reaching $\\sigma_{\\rm{log Y}}\\sim0.36$. Temperature, however, is nicely predicted by optical richness. The scatter in the $T_X-$\\bgc~relationship is lower even than the scatter we find in $L_X-T_X$ $(\\sigma_{{\\rm{log Y}}}\\sim0.16$ compared to 0.22), though this measurement may exclude several of the strongest outliers in our sample. This can be explained by the lack of hot baryons in RCS clusters, which would affect cluster $L_X$, but not temperature or \\bgc. Richness-mass relationships are generally consistent with one another and with previous studies~\\citep{yee03,blindert07}, and we find an average scatter of $\\sigma_{{\\rm{log Y}}}\\sim0.30$ for these relationships (Figure~\\ref{fig11}). Though this scatter is on average somewhat higher than the other mass proxies investigated here (by $0.07-0.19$ in $\\sigma_{\\rm{logY}}$), the comparative speed and ease with which it can be obtained still recommend it as a promising tool for mass estimations of large high redshift cluster samples."
    },
    "0710/0710.0100_arXiv.txt": {
        "abstract": "{ We estimate the limiting angular resolution and detection area for an array of 3 large-aperture Imaging Atmospheric Cherenkov Telescopes. We consider an idealized IACT system in order to understand the limitations imposed by the intrinsic nature of the atmospheric showers and geometry of the detector configuration. The idealization includes the assumptions of a perfect optical system and the absence of the night sky background with the goal of finding the optimum camera geometry and array configuration independent of detailed assumptions about the telescope design.  The showers are simulated using the ALTAI code for the altitude of 2700 m corresponding to one of possible future sites for a new northern-hemisphere array. The optimal design depends on the target energy range; for each energy we vary both the cell length (telescope spacing) and the image processing parameters in order to maximize the signal-to-noise ratio. We then present the resulting values of the detection area and the angular resolution for this energy dependent optimization. We discuss the dependence of these quantities on the field of view of the telescopes and pixel size of the camera.}  \\begin{document} ",
        "introduction": "Future experiments based on Imaging Atmospheric Cherenkov Technique are likely to benefit from advances in technology, but the technology development strategy depends critically on performance and cost tradeoffs for parameters such as pixel size, field of view and cell length. A likely realization of the next generation IACT observatory is a km$^2$ array of a large number of telescopes equipped with $>100$ m$^2$ mirrors. The basic properties of such an array can be derived partially from the studies of a single cell of 3 telescopes. In this paper we determine the best performance of a system of 3 telescopes with the aperture of 250 m$^2$. The idealizations used here include: \\begin{itemize} \\item a perfect optical system, \\item ability to measure the exact coordinates of individual photoelectrons in the focal plane of the camera, \\item zero level of night sky background (NSB), \\item field of view of a telescope that can be as large as $180^\\circ$. \\end{itemize}  These idealizations allow us to study the limitations of the IACT technique imposed by the nature of the atmospheric shower itself and geometry of the telescopes rather than by the specific details of the telescope hardware.  We determine the optimal spacing between the telescopes for each target energy. For the optimization we chose the figure of merit to be the highest signal-to-noise ratio achieved for a point-like $\\gamma$-source, or \\begin{equation} S(E,l)=\\frac{{dN_{\\gamma}(E, l)}/{dE}}{\\sqrt{R_{\\mathrm{CR}}(l)}} \\label{snr} \\end{equation} where $R_{\\mathrm{CR}}$ stands for the integral detection rate for the cosmic-ray background and $dN_{\\gamma}/dE$ is the differential detection rate of photons from a Crab-like spectrum. The function is closely related to the definition of the angular resolution we use in the paper, which is the radius around the source position in camera coordinates that provides the highest value for the signal-to-noise ratio. ",
        "conclusions": "Telescopes with an $8^\\circ$ field of view, achieve a better sensitivity than telescopes with a field of view of $4^\\circ$, however our simulations show a smaller effect than previous studies \\cite{perez-2006-26}. Thus, our study shows that the main motivation for building telescopes with a large $FoV$ is to observe extended sources and to survey the TeV $\\gamma$-ray sky more efficiently.  No significant differences in the signal-to-noise ratios are observed for pixel size up to $0.15^\\circ$ compared with the perfect detector. This conclusion follows from the properties of $\\gamma$-shower itself and intrinsic rejection of hadronic showers. Optimization of angular resolution at high energies, as well as the desire to reduce the NSB background in individual pixels may favor smaller pixels. The values obtained here are optimal if night-sky noise is negligible. For a future array this may indeed be the case, if the signal integration windows are further shortened, and sites with lower night sky background noise are used. The results presented apply for an idealized 3-telescope cell. Work is in progress to extend our results to a km$^2$ array with the constraint on the total number of telescopes."
    },
    "0710/0710.4273_arXiv.txt": {
        "abstract": "} \\nc{\\eab}{ \\noindent  We extend previous work on the minimal D-term inflation model modified by Right-Handed (RH) sneutrino fields to include additional inflaton-sector SUGRA corrections and two-field inflation effects. We show that SUGRA corrections simultaneously allow $n_{s}$ to be within 3-year WMAP limits and the cosmic string contribution to the CMB power spectrum to be less than 5$\\%$. For gauge coupling $g \\leq 1$, the CMB contribution from cosmic strings is predicted to be at least 1$\\%$ while the spectral index is predicted to be less than 0.968 for a CMB string contribution less than 5$\\%$. Treating the inflaton-RH sneutrino system as a two-field inflation model, we show that the time-dependence of the RH sneutrino field strongly modifies the single-field results for values of RH sneutrino mass $m_{\\Phi} > 0.1 H$. The running spectral index is $\\alpha \\approx -0.0002$ when $m_{\\Phi} < 0.1 H$ but increases to positive values as $m_{\\Phi}/H$ increases, with $\\alpha > 0.008$ for $m_{\\Phi} > 1.0 H$. ",
        "introduction": "One clue to the nature of supersymmetric (SUSY) inflation is its compatibility with supergravity (SUGRA). It has long been known that in the case of inflation models driven by an F-term, SUGRA corrections generate a Hubble scale inflaton mass which spoils the flatness of the inflaton potential \\cite{eta}. Elimination of such corrections is possible in F-term hybrid inflation \\cite{fti}, but only if there is a significant suppression of the higher-order terms in the K\\\"ahler potential \\cite{etahi}. In contrast, D-term hybrid inflation models \\cite{dti} automatically evade Hubble corrections and are therefore naturally compatible with SUGRA.   However, there are problems with D-term inflation. Minimal D-term inflation ends with a phase transition which produces local cosmic strings \\cite{dtics}. WMAP observations \\cite{wmap} bound the contribution of local cosmic strings to the cosmic microwave background (CMB) power spectrum to be less than about 10$\\%$ \\cite{wyman}, which results in an upper bound on the cosmic string tension which excludes minimal D-term inflation unless the gauge and superpotential couplings are suppressed \\cite{endo,csb,csb2}. In addition, the latest 3-year WMAP value of the spectral index, $n_{s} = 0.958 \\pm 0.016$ (1-$\\sigma$) \\cite{wmap}, is significantly smaller that the minimal D-term inflation prediction, $n_{s} = 0.983$.  D-term inflation has two distinct regimes corresponding to large and small couplings. If we consider unsuppressed gauge and superpotential couplings then the CMB power spectrum is completely dominated by the cosmic string contribution. However, if we consider small superpotential and gauge couplings ($ \\lae 10^{-5}$ and $10^{-2}$ respectively) then it is possible to suppress the cosmic string contribution to O$(10)\\%$ \\cite{csb,csb2}. (Similar but somewhat less restrictive upper bounds were obtained in \\cite{battye}.) The adiabatic density perturbation from inflation in this case has $n_{s} = 1$. However, the addition of the cosmic string contribution can lower the apparent spectral index of the CMB sufficiently for it to be consistent with the observed value \\cite{battye,aur,bevis}. This may allow minimal D-term inflation to be consistent with the observed CMB, but at the cost of small couplings \\footnote{A alternative approach has recently been suggested in \\cite{bj}.}.  In the following we will consider the case of unsuppressed gauge and superpotential couplings. In this case, if minimal D-term inflation is correct, it should be modified to reduce both the string tension and the spectral index.    How can this be achieved within a well-motivated framework? A possible solution was proposed in \\cite{dlin1}. In general, the minimal D-term inflation model should be considered as a sector of a complete theory which also includes the Minimal SUSY Standard Model (MSSM), extended to include neutrino masses ($\\nu$MSSM). In this case, modification of the minimal D-term inflation model could come from corrections due to $\\nu$MSSM sector fields. In the context of see-saw neutrino masses, it was shown in \\cite{dlin1} that a SUGRA correction due to the F-term energy density of a RH sneutrino field can lower both the string tension and spectral index. In this model, RH Sneutrino-Modified (RHSM) D-term inflation, it was found that if the RH sneutrino field is in the range (0.1-1)$M$ ($M = M_{Pl}/\\sqrt{8 \\pi}$) and has mass of the order of (0.1-1)$H_{I}$, where $H_{I}$ is the Hubble parameter during inflation, then it is possible to satisfy the CMB upper bound on the string tension for values of the spectral index that are within the lower bound from WMAP. This provides the only example of a minimal D-term hybrid inflation model with unsuppressed couplings that is consistent with the observed CMB. However, the analysis of the RHSM D-term inflation model given in \\cite{dlin1} was restricted to the case where inflaton-sector SUGRA corrections are negligible and the slow-rolling RH sneutrino field can be treated as a constant. The purpose of this paper is to generalise this analysis.  SUGRA corrections to minimal D-term inflation without RH sneutrino modification were considered in \\cite{osyok}. Although these allowed a reduction of both the string tension and the spectral index, it was shown that such corrections cannot lower the string tension sufficiently to be consistent with the observed CMB \\cite{csb2}. Here we will show that once RH sneutrino modification is included, such corrections can substantially increase the range of spectral index for which RHSM D-term inflation is consistent with the observed CMB.    The RHSM D-term inflation model also provides a well-motivated example of a two-field inflation model. In order to study the effects of two-field inflation we will apply the $\\delta N$ method \\cite{deltan,deltan2}. We will show how such effects modify the string tension and spectral index and determine the conditions under which the RH sneutrino field may be treated as constant.  The paper is organised as follows. In Section 2 we review the original RHSM D-term inflation model and update WMAP and cosmic string bounds to take account of recent developments. In Section 3 we consider the effect of combining SUGRA corrections from the inflaton sector fields with the RH sneutrino correction. In Section 4 we apply the $\\delta N$ method to study two-field inflation effects of the inflaton-RH sneutrino system and establish the conditions under which the constant RH sneutrino field approximation is valid. We also compute the running spectral index due to the evolution of the RH sneutrino field.  In Section 5 we present our conclusions. ",
        "conclusions": "We have shown that including SUGRA corrections allows the RH sneutrino-modified variant of minimal D-term inflation to be consistent with CMB bounds on the spectral index even when the CMB contribution from cosmic strings is small. The value of $n_{s}$ is predicted to be low, less than 0.968 when $g \\leq 1.0$ for a cosmic string contribution less than 5$\\%$, while cosmic strings are expected to contribute at least 1$\\%$ to the CMB power spectrum for $g \\leq 1$. The combination of low spectral index plus greater than $1 \\%$ percent CMB cosmic string contribution may be regarded as a prediction of the RH sneutrino-modified D-term inflation model.  We have anaylsed the model as a two-field inflation model using the $\\delta N$ method. The main two-field inflation effect is from the time evolution of the classical RH sneutrino field, with quantum fluctuations of the RH sneutrino field having a negligible effect on the curvature perturbation.  Two-field inflation effects strongly diminish the RH sneutrino suppression of the cosmic string tension once $m_{\\Phi}/H \\gae 0.2$, in which case the model will violate CMB bounds on cosmic strings. However, the effect on RH sneutrino suppression of the spectral index is smaller, which may be significant for D-term inflation models which evade cosmic string formation but still have a large spectral index. For $m_{\\Phi}/H \\lae 0.1$ it is a good approximation to treat the model as a single-field inflation model with constant $\\phi$.  The running spectral index in the limit $m_{\\Phi}/H < 0.1$ is negative and of the same magnitude as in conventional hybrid inflation models, $\\alpha \\approx -0.0002$. For larger $m_{\\Phi}/H$, $\\alpha$ increases to positive values due to the rolling of the $\\Phi$ field. Values of $\\alpha$ greater than 0.008 are found for $m_{\\Phi}/H > 1.0$. For these values of $m_{\\Phi}/H$ there is no significant suppression of the CMB contribution from cosmic strings. However, the spectral index is still reduced relative to conventional D-term inflation, so a large positive $\\alpha$ could occur in non-minimal D-term inflation models without cosmic strings. Observation of a positive running spectral index together with a reduced spectral index may therefore provide a signature of non-minimal D-term inflation with RH sneutrino modification.  In conclusion, we have shown that RH sneutrino-modified D-term inflation provides a version of the original minimal D-term inflation model which is completely consistent with CMB observations. The model requires no suppression of couplings and no additional fields beyond those that already exist in the minimal D-term inflation model and the MSSM with neutrino masses. It therefore provides a predictive candidate for SUGRA inflation with all the advantages of conventional D-term inflation with respect to naturalness. RH sneutrino modification can also play a role in bringing non-minimal D-term inflation models without cosmic strings into agreement with the observed spectral index."
    },
    "0710/0710.1510_arXiv.txt": {
        "abstract": "We consider the linear growth of matter perturbations in various dark energy (DE) models. We show the existence of a constraint valid at $z=0$ between the background and dark energy parameters and the matter perturbations growth parameters. For $\\Lambda$CDM $\\gamma'_0\\equiv \\frac{d\\gamma}{dz}|_0$ lies in a very narrow interval $-0.0195 \\le \\gamma'_0 \\le -0.0157$ for $0.2 \\le \\Omega_{m,0}\\le 0.35$. Models with a constant equation of state inside General Relativity (GR) are characterized by a quasi-constant $\\gamma'_0$, for $\\Omega_{m,0}=0.3$ for example we have $\\gamma'_0\\approx -0.02$ while $\\gamma_0$ can have a nonnegligible variation. A smoothly varying equation of state inside GR does not produce either $|\\gamma'_0|>0.02$. A measurement of $\\gamma(z)$ on small redshifts could help discriminate between various DE models even if their $\\gamma_0$ is close, a possibility interesting for DE models outside GR for which a significant $\\gamma'_0$ can be obtained. ",
        "introduction": "There is growing observational evidence for the late-time accelerated expansion of our universe \\cite{SS00}. This radical departure from conventional decelerated expansion is certainly a major challenge to cosmology. This non standard expansion could be due to an exotic non clustered component yet to be determined with a sufficiently negative pressure called Dark Energy (DE). By analogy all models trying to explain the accelerated expansion are called DE models but many models go now well beyond this simple picture. While the usual Friedmann equations in the presence of a cosmological constant term $\\Lambda$ seem to be in good agreement with the data, it is clear that other models with a variable equation of state are allowed as well \\cite{SS00}. While a cosmological constant universe is appealing because of its simplicity it nonetheless poses the problem of the magnitude of the cosmological constant $\\Lambda$. This is the basic incentive to look for other models where DE has a variable equation of state. An additional incentive comes from the possibility to have phantom dark energy at low redshifts as this excludes the quintessence models, models with a minimally coupled scalar field inside GR \\cite{RPW88}. It might also be that one should change the theory of gravity, as for example in scalar-tensor models \\cite{ST,BEPS00}, and a lot of research has focused recently on other modified gravity models and higher dimensional models. Sometimes, the background expansion going back to high redshifts is enough to rule out some models \\cite{APT07}, but typically this is not the case: models of a very different kind will be able to have a viable background expansion where the low redfshift expansion is in accordance with SNIa data.  Depending on the gravity theory one is considering, the growth of the perturbations, even at the linear level, will be affected. Indeed, while distance luminosity measurements probe the cosmic expansion, matter perturbations probe in a independent way (see e.g. \\cite{B06}) the gravity theory responsible for their growth (and of course also for the cosmic expansion). The growth rate of matter perturbations could be probed with three dimensional weak lensing surveys (see e.g. \\cite{HKV07}). Hence two DE models based on different gravitation theories can give the same late-time accelerated expansion and still differ in the matter perturbations they produce \\cite{S98}. This fact could provide an additional important way to discriminate between various models (see e.g. \\cite{HL06}) and it is therefore important to characterize as accurately as possible the growth of matter perturbations which is the aim of the present work. ",
        "conclusions": "Considering the linear growth of matter perturbations in various models, we give a constraint at $z=0$, eq.(\\ref{dgamma0b}) valid for all models, including modified gravity DE models that satisfy $\\frac{G_{{\\rm eff},0}}{G_{N,0}}=1$. This constraint implies that the quantity $\\gamma'_0$ is completely fixed by the remaining parameters $\\gamma_0,~w_{DE,0}$ and $\\Omega_{m,0}$. For the models considered here inside GR, $|\\gamma'_0|\\lesssim 0.02$. Interestingly for models inside GR with constant $w_{DE}$, $\\gamma'_0$ is quasi-constant with $\\gamma'_0\\approx -0.02$ as the variation of $w_{DE,0}$ is compensated by a simultaneous variation of $\\gamma_0$ (for given $\\Omega_{m,0}$).  We have generically $\\gamma'_0\\ne 0$ and we emphasize that a significant $\\gamma'_0$ could help discriminate between models, even if their $\\gamma_0$ values are close. We have illustrated this schematically on Figure 2b. This potential resolution improves as $\\Omega_{m,0}$ goes up and/or $w_{DE,0}$ goes down and could be important when dealing with DE models outside General Relativity. We will give elsewhere specific models where this is the case \\cite{GP07}. Generally, this approach could be very fruitful whenever $\\gamma(z)$ is close to linear on small redshifts $0\\le z\\le 0.5$ so that the slope is essentially given by $\\gamma'_0$. So we feel it would be useful to try to measure $\\gamma(z)$ on small redshifts, and not just $\\gamma_0$. Finally it is important to realize that neglecting a small but nonvanishing $\\gamma'_0$ can induce a large error on the parameters $\\Omega_{m,0},~w_{DE,0}$ that one could extract from the growth of matter perturbations."
    },
    "0710/0710.0385_arXiv.txt": {
        "abstract": "In the past decade, surveys of the stellar component of the Galaxy have revealed a number of streams from tidally disrupted dwarf galaxies and globular clusters.  Simulations of hierarchical structure formation in $\\Lambda$CDM cosmologies predict that the dark matter halo of a galaxy like the Milky Way contains hundreds of subhalos with masses of $\\sim10^{8}$ M$_{\\odot}$ and greater, and it has been suggested that the existence of coherent tidal streams is incompatible with the expected abundance of substructure.  We investigate the effects of dark matter substructure on tidal streams by simulating the disruption of a self-gravitating satellite on a wide range of orbits in different host models both with and without substructure.  We find that the halo shape and the specific orbital path more strongly determine the overall degree of disruption of the satellite than does the presence or absence of substructure, i.e., the changes in the large-scale properties of the tidal debris due to substructure are small compared to variations in the debris from different orbits in a smooth potential.  Substructure typically leads to an increase in the degree of clumpiness of the tidal debris in sky projection, and in some cases a more compact distribution in line-of-sight velocity.  Substructure also leads to differences in the location of sections of debris compared to the results of the smooth halo model, which may have important implications for the interpretation of observed tidal streams.  A unique signature of the presence of substructure in the halo which may be detectable by upcoming surveys is identified.  We conclude, however, that predicted levels of substructure are consistent with a detection of a coherent tidal stream from a dwarf galaxy. ",
        "introduction": "\\label{sec:intro}  In recent years, surveys of the stellar component of the Milky Way have revealed rich structure and a number of new tidal streams \\citep{belokurov_etal_06field,grillmair_06,grillmair_dionatos_06cold,belokurov_etal_07orphan}. In addition to the well-studied Sagittarius (Sgr) stream \\citep{ibata_irwin_lewis_stolte_01, majewski_etal_03, newberg_etal_03, martinez-delgado_etal_04} and the Monoceros stream \\citep{newberg_etal_02, yanny_etal_03, penarrubia_etal_05}, streams from the tidal disruption of globular clusters such as Palomar 5 \\citep{odenkirchen_etal_01, rockosi_etal_02, odenkirchen_etal_03, grillmair_dionatos_06pal} and NGC 5466 \\citep{grillmair_johnson_06ngc5466, belokurov_etal_06ngc} have also been identified. Tidal streams probe the gravitational potential on large scales, and consequently may prove to be a valuable tool for constraining the mass distribution of the Galaxy.  Simulations of structure formation in $\\Lambda$CDM cosmologies  predict as many as several hundred dense clumps with masses $\\gtrsim 10^8$ M$_{\\odot}$ in the halo, roughly an order of magnitude more than the observed abundance of dwarf galaxies surrounding the Milky Way \\citep{klypin_etal_99, moore_etal_99}. One class of solutions to this `missing satellites' problem involves modifications of the properties of the dark matter particle, e.g., warm dark matter \\citep{hogan_dalcanton_00} or self-interacting dark matter \\citep{spergel_steinhardt_00}, in order to reduce the amount of substructure.  Suppression in the small-scale end of the primordial power spectrum has also been invoked to reduce the predicted abundance of subhalos \\citep{kamionkowski_liddle_00, zentner_bullock_03}.  Alternatively, mechanisms by which only a fraction of the dark matter subhalos have visible stellar components at the present epoch have been proposed to explain this discrepancy \\citep{bullock_kravtsov_weinberg_01,kravtsov_gnedin_klypin_04}.  In these scenarios, a large number of dark substructures are predicted to populate the halo.  In the past few years the number of known satellites of the Milky Way has increased dramatically with the discovery of 11 new, faint objects which are likely to be dwarf galaxies or globular clusters \\citep{willman_etal_05glob, willman_etal_05ursa, zucker_etal_06ursa, zucker_etal_06canes,belokurov_etal_06bootes, belokurov_etal_07cats, irwin_etal_07}, hinting that a significant number of very faint or dark satellites could populate the halo. If such a population does exist, these structures may be detectable in several ways.  Gamma rays from the annihilation of dark matter particles in substructures may reveal their presence \\citep[e.g.,][]{bergstrom_etal_99, calcaneo-roldan_moore_00, tasitsiomi_olinto_02, stoehr_etal_03, koushiappas_etal_04, diemand_kuhlen_madau_07}. Dark substructures may also be detectable via flux anomalies in multi-image gravitational lens systems \\citep{mao_schneider_98, metcalf_madau_01, dalal_kochanek_02, chiba_02} or via spectroscopic lensing \\citep{moustakas_metcalf_03}.  However, recent work favors microlensing by stars in the lens galaxy over dark matter substructure as an explanation for flux anomalies  \\citep{pooley_blackburne_rappaport_etal_07}.  In this study we take a theoretical approach to determine whether substructure as predicted by $\\Lambda$CDM models may be detectable by its influence on the formation and properties of tidal streams from dwarf galaxies.  We select a variety of orbits in several host galaxy models and simulate the tidal disruption of a progenitor satellite on these orbits both with and without dark matter substructure to evaluate the influences of the orbital path, host potential, and presence of substructure on the tidal debris. In \\S\\ref{sec:orbits} we discuss the phase space properties of orbits in the context of tidal disruption, describe our host galaxy models, and outline our method for selecting a diverse range of orbits.  We give the details of our numerical simulations in \\S\\ref{sec:sims}, and present our results in \\S\\ref{sec:results}.  We discuss our study in relation to previous work on this subject in \\S\\ref{sec:disc}, and finally summarize our conclusions in \\S\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have investigated the effect of halo substructure on the tidal debris from a dwarf galaxy.  Our conclusions can be summarized as follows: \\begin{enumerate}  \\item{Our results suggest that the shape of the host potential and the path of the orbit in question have a stronger influence on the overall spatial distribution of tidal streams than does the presence of substructure in the halo.  In light of the enormous amount of variation in the properties of the tidal debris produced by simulations in a smooth halo, it is clear that dispersal of the debris is not an effect that can be attributed exclusively, or even primarily, to substructure.  We find that the detection of a coherent tidal stream from a dwarf galaxy does not rule out the abundance of substructure considered in this work.}  \\item{Substructure can noticeably shift the location of sections of debris relative to the debris in the corresponding smooth halo simulation.  This is likely to be an important complication for constructing models to fit observational data.}  \\item{Interactions with a population of subhalos can lead to significant changes in the small-scale properties of the tidal debris.  Satellites on some orbits develop rich structure in the $r_{\\rm sun}$-$v_{\\rm los, rest}$ distribution in a smooth halo which is largely smeared out when simulated with substructure.  The presence of substructure also tends to lead to more clustered debris in a sky projection, and in some cases smaller velocity dispersions.}  \\item{Stars at large distances with a small range of positive line-of-sight velocities with respect to the rest frame of the Galaxy which increase with increasing distance may be a unique signature of an interaction with a relatively massive halo substructure.  We find that particles in this configuration typically have experienced a single strong encounter with a subhalo of mass $\\gtrsim 10^{10}$ M$_{\\odot}$.  The formation of this feature thus depends on the orbits of the most massive subhalos in relation to the spatial and velocity distribution of stellar material.  The large number of tidal streams expected to be present in a halo like that of the Milky Way, coupled with the fact that stars from the stellar halo of the Galaxy would similarly be susceptible to such interactions, implies a high likelihood for producing these features.  Despite the fact that these particles are observed at large radial velocities and distances, they are not unbound and will thus eventually return.  Each individual feature is therefore transient.}  \\item{While satellites on resonantly trapped orbits appear to produce more coherent tidal debris than non-resonant orbits in some individual cases, we do not consistently see this trend in our simulated debris.  Similarly, we do not find that streams from satellites on resonantly trapped orbits are less susceptible to the effects of substructure.}  \\end{enumerate}   With the wealth of data expected from upcoming missions, we are optimistic about the prospects of using tidal streams to gain a better understanding of the mass distribution of the Milky Way.  If the halo of the Galaxy contains a population of substructure as predicted by $\\Lambda$CDM, clear signatures of its presence may be detectable in the near future."
    },
    "0710/0710.2455_arXiv.txt": {
        "abstract": "\\noindent Experimental evidence is reviewed for the existence of superfluid turbulence in a differentially rotating, spherical shell at high Reynolds numbers ($\\Rey\\gsim 10^3$), such as the outer core of a neutron star. It is shown that torque variability increases with $\\Rey$, suggesting that glitch activity in radio pulsars may be a function of $\\Rey$ as well. The $\\Rey$ distribution of the $67$ glitching radio pulsars with characteristic ages $\\tau_c \\leq 10^6$ {\\rm yr} is constructed from radio timing data and cooling curves and compared with the $\\Rey$ distribution of all $348$ known pulsars with $\\tau_c \\leq 10^6$ {\\rm yr}. The two distributions are different, with a Kolmogorov-Smirnov probability $\\geq 1 - 3.9 \\times 10^{-3}$. The conclusion holds for (modified) Urca and nonstandard cooling, and for Newtonian and superfluid viscosities. ",
        "introduction": "The known number of glitching radio pulsars has quadrupled in the last four years, as has the total number of glitches, in the wake of the Parkes Multibeam Survey and improved multifrequency timing solutions \\citep{peralta06}. Out of a total population of $1730$ pulsars, $101$ have experienced a total of $287$ glitches \\citep{wan00,Kra03,catalog04,js06}. The enlarged data set invites an updated statistical study. Previous analyses identified certain empirical trends, e.g., glitch activity and healing parameter versus characteristic age \\citep{sl96,uo99,wan00}, and glitch amplitude versus recurrence time \\citep{mmwgz06}, while refuting certain theoretically predicted correlations, e.g., glitch amplitude versus spin period $P$, period derivative $\\dot{P}$, and total spin down since the last glitch (``reservoir model\") \\citep{sl96,wan00}.  On the theoretical front, the global pattern of superfluid circulation in the outer core of a neutron star was recently computed numerically for the first time \\citep{pmgo05a,pmgo06a,pmgo06b}. In general, a fluid contained in a differentially rotating, spherical shell flows unsteadily under a wide range of conditions, especially in thick shells and at high Reynolds numbers $\\Rey \\gsim 10^3$ \\citep{je00}. Such unsteady flow is observed in numerical and laboratory experiments with viscous fluids, e.g., transitions between Taylor vortex states \\citep{mt87b}, Taylor vortex oscillations \\citep{h98}, relaminarization \\citep{nst05}, and nonaxisymmetric modes on the route to high-$\\Rey$ turbulence, like Taylor-G\\\"ortler vortices \\citep{li04}. Time-dependent flow is also observed in superfluids like helium II, e.g., torque ``steps\" and quasiperiodic torque oscillations in Couette and spin-up experiments \\citep{tsatsa80}.  In this Letter, we combine theory and observation to test whether superfluid turbulence exists in neutron stars, as originally proposed by \\citet{g70}. In \\S\\ref{sec:hydro}, we present numerical simulations of a $^1$S$_0$-paired neutron superfluid in the outer core of a neutron star as the crust spins down electromagnetically. We show that the torque is steady for $\\Rey \\lsim 10^3$ and unsteady (e.g., oscillatory) otherwise. In \\S\\ref{sec:reynolds}, we combine standard cooling curves with measurements of the characteristic ages of radio pulsars to compute the temperature, viscosity, and hence $\\Rey$ of the star. We show that the $\\Rey$ distributions of those pulsars that glitch, and those that do not, are different --- a rare instance of regularity in the glitch phenomenon. The implications for the existence of superfluid turbulence in neutron stars are weighed critically in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion} A superfluid (or, indeed, a Newtonian fluid) contained in a differentially rotating, spherical shell circulates meridionally and exerts a time-dependent viscous torque on the outer shell. The vigour of the circulation, and the variability of the torque, increase with $\\Rey$;  the flow is steady for $\\Rey \\lsim 10^2$ and chaotic for $\\Rey \\gsim 10^5$. Spherical Couette flow is an idealised model of the superfluid outer core of a neutron star, in which the outer sphere is the stellar crust, which spins down electromagnetically. It is therefore interesting to test whether the incidence of pulsar rotational irregularities like glitches depends on $\\Rey$. We show in \\S\\ref{sec:reynolds} that this is indeed the case: the distribution of glitching and nonglitching pulsars with $\\tau_c \\leq 10^6$ {\\rm yr} is different, with K-S probability $\\geq 1 - 3.9 \\times 10^{-3}$. This new fact joins the $A_g$-versus-$\\tau_c$ correlation as one the few observed regularities in glitch phenomenology.  An immediate worry is that the distributions in Figure \\ref{fig:re13} are contaminated by some observational selection effect. Our estimates of $\\Rey$ are uncertain in two ways. First, $\\tau_c$ is the characteristic spin-down age of the pulsar, not its true age. For example, five young pulsars with $\\tau_c \\leq 1.7 \\times 10^3$ {\\rm yr} have braking index $n < 3$, whereas the formula $\\tau_c = P/(2\\dot{P})$ assumes $n=3$ \\citep{m97,lkgk06}. Second, the theoretical cooling curves depend sensitively on the superfluid energy gap in the outer core, which is poorly known (although admittedly calibrated by $11$ pulsars with observed thermal emission). However, the above uncertainties do {\\it not} affect the conclusion that the $\\Rey$ distributions of the glitching and nonglitching populations in Figure \\ref{fig:re13} are different, because $T$ and hence $\\Rey$ are calculated in the same way for both samples. Indeed, it is hard to imagine that our ability to detect glitches in radio timing data is a function of $\\Rey$.  Physically, the conclusions from Figure \\ref{fig:re13} are at once natural yet surprising. It is well known that torque variability increases with $\\Rey$ in a differentially rotating superfluid in laboratory and numerical experiments (\\S\\ref{sec:hydro}). However, it is strange that pulsars behave differently at $\\Rey \\lsim 10^{10}$ (few glitches) and $\\Rey \\gsim 10^{11}$ (many glitches). Naively, Kolmogorov turbulence should be scale-free, fully developed, and {\\it indistinguishable} at these Reynolds numbers. There are at least two ways to explain this. First, processes involving microscopic superfluid turbulence, e.g., pinning \\citep{apas84} and vortex tangle formation \\citep{pmgo06b}, may depend sensitively on $\\Rey$ in the range $10^{10} \\leq \\Rey \\leq 10^{11}$. Second, it is possible that theoretical estimates of $\\Rey$ \\citep{acg05,mm05} are $\\sim 10^6$ times too high. If so, this reduces the peak of the glitching pulsar distribution to $\\Rey \\lsim 10^5$, exactly where spherical Couette flow breaks up into nonaxisymmetric modes before becoming chaotic (\\S\\ref{sec:hydro}). We have no reason to doubt the $\\nu_n$ carefully derived in the literature. However, if the flow is turbulent, Reynolds stresses may completely overwhelm viscous stresses, such that $\\nu_n$ is replaced by $\\nu_n+A$, where $A$ satisfies $A \\partial v_i/\\partial x_j \\approx \\langle \\delta v_i \\delta v_j \\rangle$, $v_i$ is the mean velocity, and $\\delta v_i$ is the fluctuating velocity (zero average) \\citep{pedlosky_book}. Typically, we have $A/\\nu_n \\sim (\\delta v/v)^2 \\Rey \\gg 1$. Turbulent Reynolds stresses drastically reduce $\\Rey$ in a range of geophysical and astrophysical applications \\citep{pedlosky_book}.  Plainly, there exist several high-$\\Rey$ pulsars that do not glitch. Observational selection effects aside, there exist several theoretical mechanisms that may be responsible. First, for spherical Couette flow in the range $10^3 < \\Rey < 10^7$, there are sizable intervals of $\\Rey$ where the flow is almost laminar, interposed between intervals where it is turbulent \\citep{nt05}. Second, for  $\\Rey > 10^7$, where the turbulence is fully developed (Kolmogorov) and ``isotropic\" when averaged over a turn-over time, the torque on the star is steadier than for organised modes (e.g., Taylor-G\\\"ortler vortices) near the onset of turbulence ($\\Rey \\sim 10^5$), which are not isotropic on average. Third, if atmospheres and planetary interiors are any guide, the turbulent Reynolds stresses depend sensitively on what nonlinear modes are excited in the turbulence, and the effective $\\Rey$ may vary greatly (and unpredictably) from object to object, in ways not captured by (\\ref{eq:rey2}). Incidentally, none of the pulsars in our sample have low $\\Rey$, even after adjusting for turbulent Reynolds stresses; the minimum is $\\Rey \\sim 10^3$. Hence there are no examples of low-$\\Rey$ objects that glitch unexpectedly.  In closing, we emphasize that the results of this paper do {\\it not} prove that superfluid turbulence exists in a neutron star, let alone that it controls glitch behaviour. However, considerable effort has been expended to identify patterns in glitch behaviour, with limited success. The undeniable empirical fact that the $\\Rey$ distributions in Figure \\ref{fig:re13} are different is therefore intriguing, especially because $\\Rey$ is a {\\it dimensionless} quantity with {\\it deep hydrodynamical significance}, and because the K-S comparison is robust with respect to how $T$, $\\nu_n$, and hence $\\Rey$ are estimated observationally. One hopes it will offer an enlightening clue to glitch physics, whether or not superfluid turbulence turns out to play a central role. We also hope that the results presented here will stimulate the ongoing observational campaign to measure pulsar temperatures."
    },
    "0710/0710.0499_arXiv.txt": {
        "abstract": "{} {We aim at getting high spatial resolution information on the dusty core of bipolar planetary nebulae to directly constrain the shaping process.} {We present observations of the dusty core of the extreme bipolar planetary nebula Menzel\\,3 (Mz\\,3, Hen 2-154, the Ant) taken with the mid-infrared interferometer MIDI/VLTI and the adaptive optics NACO/VLT.} {The core of Mz\\,3 is clearly resolved with MIDI in the interferometric mode, whereas it is unresolved from the Ks to the N bands with single dish 8.2m observations on a scale ranging from 60 to 250mas. A striking dependence of the dust core size with the PA angle of the baselines is observed, that is highly suggestive of an edge-on disk whose major axis is perpendicular to the axis of the bipolar lobes. The MIDI spectrum and the visibilities of Mz\\,3 exhibit a clear signature of amorphous silicate, in contrast to the signatures of crystalline silicates detected in binary post-AGB systems, suggesting that the disk might be relatively young. We used radiative-transfer Monte Carlo simulations of a passive disk to constrain its geometrical and physical parameters. Its inclination (74\\deg$\\pm$3\\deg) and position angle (5\\deg$\\pm$5\\deg) are in accordance with the values derived from the study of the lobes. The inner radius is 9$\\pm$1AU and the disk is relatively flat. The dust mass stored in the disk, estimated as $1 \\times 10^{-5}M_{\\odot}$, represents only a small fraction of the dust mass found in the lobes and might be a kind of relic of an essentially polar ejection process. } {} ",
        "introduction": "It is commonly accepted that disks can be an essential ingredient to the shaping of planetary nebulae, but the spatial resolution of astronomical instruments is usually not sufficient for detecting and studying these disks. The geometry of the disk and the mass stored are key parameters for constraining the models of nebula formation, and for tracing back the evolution of the central star (\\cite{2002ARA&A..40..439B}). Its inner edge, as seen from the star, could be thick and dense enough to collimate stellar winds into lobes, and knowing its geometry enables us to better understand the distribution of illuminated and shadowed regions in the extended nebula. The VLT Interferometer is a powerful `disk hunter' with its two instruments: AMBER operating in the near-IR (\\cite{2007A&A...464....1P}) and MIDI in the mid-IR (\\cite{2003Ap&SS.286...73L}) providing a typical spatial resolution of 2 and 10mas, respectively. It must be stressed that the number of visibility measurements usually recorded do not allow any image reconstruction from the interferometric observations because this would be too time-consuming. Therefore, the interpretation uses models that depend critically on the information already gathered on the sources by other techniques. As an example, a disk was detected by HST and MIDI in the center of the complex multipolar nebula CPD-56$^\\circ$8032 (\\cite{2002ApJ...574L..83D}, \\cite{2006A&A...455.1009C}). The disk exhibits a bright and extended (on an interferometric scale, $\\sim$70mas) inner rim whose PA and inclination do not seem directly related to extended structures in the HST images. Definitely, the disk detection based on three visibility measurements provided a lot of  information, but also demonstrated the need for an interferometric imaging campaign involving wide $uv$ coverage for this complex object.  In contrast, the Butterfly (M2-9) and the Ant (Mz\\,3, Hen 2-154) nebulae were chosen as perfectly suited targets to a snapshot study with the VLTI, as they are among the most impressive bipolar PNs, highly collimated on large scales with tightly pinched waists as seen in Hubble Space Telescope (HST) images (\\cite{2002ARA&A..40..439B} and references therein; \\cite{2003MNRAS.342..383S}; \\cite{2005AJ....129..969S}; \\cite{2005AJ....130..853S}). For these targets, assuming that the putative disk hidden in their center shares the symmetry of the extended nebula constitutes a good starting point for preparing and interpreting the VLTI observations. The nebula Mz\\,3 shows complex structures revealing different stages of mass-loss events that share the same axis of symmetry (\\cite{2004A&A...426..185S}, \\cite{2004AJ....128.1694G}). Of importance is also the fact that M2-9 and Mz\\,3 are spectroscopic twins at visual and near-IR wavelengths. Each nucleus shows a rich optical spectrum with the various iron and helium lines, that are usually associated with a dense circumnuclear cocoon which reprocesses most of the stellar light emitted in our line of sight that hides any close companion from direct view. Each object requires an ionization source with a temperature on the order of 30,000\\, K, and the dust mass detected in their nebulae suggests intermediate-mass progenitors (\\cite{2005AJ....129..969S}). Published distances to each object vary widely, hovering around 1-2 kpc. A detailed imaging study of the dust in the core of these nebulae has been presented recently by Smith (2003) and Smith \\& Gehrz (2005). The spatial resolution of the mid-IR images was 1$\\arcsec$, allowing them to isolate the central engines from the surrounding nebulae. The observed mid-IR SED of the unresolved central sources is flatter than a single gray-body, implying that a range of dust temperatures and an additional component of hot dust are both required to fit the near-IR continuum spectrum. This suggests the existence of a circumstellar disk.  The study of Smith \\& Gehrz (2005) motivated us to use the unique MIDI/VLTI capabilities for detecting the disks in the center of Mz\\,3 and M2-9. The present Letter is entirely devoted to the disk detected in the center of Mz\\,3, while M2-9 observations are described in another paper. The observations are presented in Section 2. In Section\\,3 we derive the physical parameters of the dusty environment by radiative-transfer models and discuss the results in Section\\,4. ",
        "conclusions": "The MIDI observations provide evidence of a flat, nearly edge-on disk primarily composed of amorphous silicate whose parameters are described in Table 2. The inner rim is at sublimation temperature and the disk is optically thick ($\\tau_N\\sim3.5$). The dust chemistry agrees with \\cite{2005A&A...431..523B} and \\cite{2005MNRAS.362.1199C} who found  evidence in the ISO spectra of weak signatures of amorphous silicate and no obvious signature of carbon dust in the form of PAHs, in accordance with the $C/O$ ratio of 0.83 determined by \\cite{2002MNRAS.337..499Z} in the gas phase. The match of the SED and the visibility curves is impressively good showing that such a simple model of a passive disk can account for these observations. That the silicate is in the form of amorphous grains is important, since the long-lived disks found in binary post-AGBs exhibit strong signatures of {\\it crystalline} silicate (\\cite{2007A&A...467.1093D, 2007BaltA..16..145D}). This suggests that the disk may be relatively young, in line with the youth of the PN. The dust mass stored in the disk, estimated to be $1\\times 10^{-5} M_{\\odot}$ is two orders of magnitude below the mass inferred in the lobes ($2.6\\times 10^{-3} M_{\\odot}$ from \\cite{2005AJ....129..969S}). The low mass of the disk seems insufficient for pinching the waist of the bipolar nebula hydrodynamically, pointing instead to an essentially polar ejection process. The small disk opening angle and the limited angle sustained by the lobes also favor such a scenario. % The detection of a bright X-ray core and a possible jet, reported by \\cite{2003ApJ...591L..37K}, is probably a good indication that a compact accretion disk is hidden from direct view by the more extended dusty disk presented here. Using near-IR color-color diagram and the similarity between the locus of symbiotic stars and Mz\\,3, \\cite{2001A&A...377L..18S} suggested that a Mira should be hidden in the center of the system. Our preferred interpretation is that this kind of diagram is useful for probing {\\it dusty disks}, but is not convincing evidence of a Mira star. Only the temporal variability of the near-IR flux would represent such evidence but it is currently lacking  for Mz\\,3. The constraint provided by our observations is that the orbit of the system must be well within the inner radius of the disk, at a distance of at most 1-2 AU from the primary, which is typical of S-type symbiotics (\\cite{2003ASPC..303....9M}) but much too small for D-type symbiotic Miras. Therefore a Mira seems improbable, but a less luminous cool giant companion might remain hidden by the disk. The constraints from the SED are tight, limiting the cool companion luminosity to less than 100-150L$_\\odot$ (\\cite{2003MNRAS.342..383S}). This object could be a young link to the binary post-AGBs (\\cite{2006A&A...448..641D}) that probably experienced non conservative mass loss. The current estimate of the mass stored in the Mz\\,3 disk is in line with the mass of long-lived (crystalline) disks found in many of the binary post-AGBs."
    },
    "0710/0710.5057_arXiv.txt": {
        "abstract": "We examine published observations of dwarf nova oscillations (DNOs) on the rise and decline of outbursts and show that their rates of change are in reasonable agreement with those predicted from the magnetic accretion model. We find evidence for propellering in the late stages of outburst of several dwarf novae, as shown by reductions in \\emph{EUVE} fluxes and from rapid increases of the DNO periods. Reanalysis of DNOs observed in TY PsA, which had particularly large amplitudes, shows that the apparent loss of coherence during late decline is better described as a regular switching between two nearby periods. It is partly this and the rapid deceleration in some systems that make the DNOs harder to detect.  We suggest that the 28.95~s periodicity in WZ Sge, which has long been a puzzle, is caused by heated regions in the disc, just beyond the corotation radius, which are a consequence of magnetic coupling between the primary and gas in the accretion disc. This leads to a possible new interpretation of the `longer period DNOs' (lpDNOs) commonly observed in dwarf novae and nova-like variables. ",
        "introduction": "Dwarf nova oscillations (DNOs) were first recognized in outbursts of dwarf novae and in nova-like variables; that is, during stages of high mass transfer rates ($\\dot{M}$) from companion to the white dwarf primary of cataclysmic variables (Warner \\& Robinson 1972). Their short time scales ($\\sim 5$--50~s) indicate originations close to the surface of the primary, but their large variations in period distinguish them from the short period highly stable magnetic accretors such as DQ Her and AE Aqr (for a general review see Warner 1995a). Paczynski (1978) suggested that these variations result from the torque acting on an accreted equatorial belt on the primary during phases of changing $\\dot{M}$. This model has been extended in Warner (1995b) and the earlier papers in this series (Woudt \\& Warner 2002, Warner \\& Woudt 2002, 2006; Warner, Woudt \\& Pretorius 2003; Pretorius, Warner \\& Woudt 2006); a review is given in Warner (2004).  In this contribution we reinvestigate some early observations, in the light of later developments, looking specifically for additional evidence relevant to the Low Inertia Magnetic Accretor (LIMA) model (Warner \\& Woudt 2002; hereafter WW2002). In Section 2 we compare predicted and observed rates of change of spin on the rise and fall of dwarf nova outbursts; in Section 3 we look at possible incidences of a propellering phase late in dwarf nova outbursts; in Section 4 we consider the question of `coherence' of DNOs; in Section 5 we suggest a model for the multiple DNOs in WZ Sge, their behaviour during outburst, and their possible relevance to the distinct group of longer period DNOs (lpDNOs). Section 6 contains a short general conclusion. ",
        "conclusions": "A principal result of this study is further evidence that the variations in period or phase of DNOs are best described by switching between two periods -- the sidereal period, determined by the rotation period of an accreted equatorial belt, and a synodic period, probably generated by reprocessing of the radiation emitted from accretion zones on the equatorial belt. This effect is only seen where the DNO amplitudes are large enough to enable detailed analysis, but may be suspected to hold in general.  We point out that propellering of accreting gas may commonly occur towards the end of dwarf nova outbursts, depending on the existence of a sufficiently strong magnetic field on the primary. We suggest that the longer period DNOs might be located at regions just beyond the corotation radius, rather than on the surface of the primary, and that previously observed but not understood periodicities in WZ Sge and AE Aqr arise from this source."
    },
    "0710/0710.5261_arXiv.txt": {
        "abstract": "I outline a mechanism, akin to Weibel instabilities of interpenetrating beams, in which the neighboring current sheets in a striped wind from an oblique rotator interact through a two stream-like mechanism (a Weibel instability in flatland), to create an anomalous resistivity that heats the sheets and causes the magnetic field to diffusively annihilate in the wind upstream of the termination shock. The heating has consequences for observable unpulsed emission from pulsars. ",
        "introduction": "}  Observations over the last 15 years \\cite{gaens06, slane05} and dynamical models \\cite{bogo05, delzanna04, delzanna06, komiss03,  spit04} have confirmed the early suggestions \\cite{emmer87, ken84, rees74} that PWNe behave as if the MHD wind from the underlying pulsars are weakly magnetized as the plasma emerges from the wind's termination shock (TS). If some form of shock acceleration underlies the conversion of the upstream flow's energy flux to the observed nonthermal spectra of synchrotron emitting electrons and positrons, with particle energies from 10s of MeV to PeV, the MHD shock jump conditions imply that $\\sigma_1 = B_1^2 /4\\pi \\rho_1 \\gamma_1 c^2$, the magnetization just upstream of the shock, must be substantially less than unity.  Studies of the flow dynamics in the nebulae, especially just outside of the termination shocks, arrive at the same conclusion.  The most advanced comparisons of dynamical theory (both MHD and beyond MHD) to date suggest that the latitude averaged upstream $\\sigma$ in the Crab Nebula is $\\sim 0.02$ \\cite{delzanna06, spit04}.  Low $\\sigma$, and, from energy conservation, high wind Lorentz factor $\\Gamma_w$, are puzzles, since elementary ideal MHD applied to relativistic wind outflows suggest small wind acceleration and $\\sigma \\gg 1$ in the asymptotic wind. As a result, there are several interpretations of the low $\\sigma$ conclusion, depending on precisely what one means by the shock transition to which the jump conditions are applied.  Begelman \\cite{begel99} suggested that the shock really is a jump in a high $\\sigma$ flow, but the toroidal magnetic field is MHD unstable with respect to 3D kink motions, whose result is to introduce downstream magnetic dissipation that rapidly coverts magetic energy into heat, hopefully in the form of the observed nonthermal particle spectra.  If this layer is thin, then the shock jump conditions apply to the overall layer, including both the true shock and the downstream dissipative layer.  The viability of this idea requires 3D MHD modeling, which has yet to appear; from an observational perspective, kink instability driven motions which can lead to actual dissipation also tangle the field, which creates the risk that the polarization of the synchrotron emission in the final model is less than that observed \\cite{veron93}.  Lyubarsky \\cite{lyubarsky03}, drawing upon a specific model of  the striped upstream wind of the oblique rotator \\cite{lyubarsky01} which suggested slow dissipation of the magnetic field and its current sheets, has suggested that the wind arrives at the shock with $\\sigma_1 \\gg 1$ but appears as if it is low sigma with respect to the downstream nebula because magnetic dissipation (driven by reconnection) occurs within the shock front itself \\cite{lyubarsky03, petri07} - the wavelength of the stripes is much less than the Larmor radii of the $e^\\pm$ in the shock's magnetic field, which facilitates rapid dissipation.  Coroniti \\cite{coroniti90} originally suggested that the current sheets of the striped wind disipate rapidly in the upstream wind, so that the MHD shock really does form in a $\\sigma_1 \\ll 1$ flow.  As pointed out in \\cite{kirk03}, for causality reasons this can happen only if the flow has four velocity $c \\beta_{w1} \\Gamma_{w1}$ just upstream of the TS that is relatively low compared to the early estimate $\\Gamma_{w1} \\sim 10^6 $ obtained in \\cite{ken84, rees74} - these early models considered a wind whose plasma content is inadequate to model the radio emission from PWNe, where the inefficiency of synchrotron emission from lower energy particles requires a wind with much larger mass loading than is required to explain the emission of optical and higher energy photons from the nebulae.  I reconsider Coroniti's model, by introducing a new model for a two-stream like filamentation instability of the curent sheets in the striped wind that draws its free energy from the interaction between neighboring sheets.  The instability has close kinship with the Weibel instability that occurs between interpenetrating plasma streams in an unmagnetized, homogeneous plasma. I use a simple magnetic trapping model of the instability's saturation to estimate the anomalous resistivity that appears within the sheets, and I find the dissipation rate to be closely related to Bohm diffusion.  The resulting dissipation is rapid, with the sheet's broadening speed exceeding $0.1c$, which leads to annihilation of the striped magnetic field within a broad equatorial sector of the outflow in the upstream wind,  thus offering a viable model for the origin of low sigma in the plasma arriving at the TS. I also suggest the radiation from the heated upstream plasma is the origin of the unpulsed optical emission observed in the Crab pulsar \\cite{kanbach05}. ",
        "conclusions": "I have argued that due to unstable interaction between the beam currents in the {\\it neighboring} sheets that separate the magnetic stripes in the oblique rotator's wind, anomalous resistivity develops in the sheets which causes them to heat, expand and consume the magnetic field of the stripes well upstream of the termination shock of the Crab pulsar's wind, and quite likely in the winds of other pulsars.   The dissipation mechanism put forward here falls in Kirk \\& Skj{\\ae}raasen's \\cite{kirk03} ``fast'' category, with the sheet broadening velocity approaching c/3 toward the end of the process. Contrary to the conclusion of Lyubarsky and Kirk \\cite{lyubarsky01}, who constructed a model in Kirk \\& Skj{\\ae}raasen's ``slow'' category, I find that dissipation of the magnetic field in the equatorial, striped zone can and probably does occur in the flow upstream of the termination shock, as was suggested by Coroniti \\cite{coroniti90}.  More complex dissipative flows, such as relativistic tearing and drift-kink instabilities \\cite{zenitani07} would only enhance this conclusion - estimates indicate that simple thickening of the sheets dominates over these more complex flows once the sheets have substantially broadened, although they may play a role in the early non-linear dynamics of the sheets.  My conclusion rests on the winds being heavily mass loaded, as is indicated by radio observations of PWNe, therefore having asymptotic wind 4-velocities much less \\cite{gallant02} than the estimates in \\cite{ken84, rees74}, who neglected the implications of the radio emission in their pioneering modeling efforts. Theoretically, such large mass loading is not understood. Theoretical models of pair creation in pulsar magnetospheres with any pretense of self-consistency all underpredict pulsars' mass loss rates by one, two or more orders of magnitude. Combined with the basic conflict between the poloidal currents found in force-free models of the underlying magnetospheres with the currents implied by the extant pair creation models (both ``polar cap'' and ``outer gap''), these problems suggest a substantial rethinking of pair creation physics and wind formation is in order \\cite{arons08}.  Most of the dissipation happens far from the star, where $\\Gamma_w$ is close to to its maximum. Because of relativistic beaming, radiation from the inefficiently emitting, resistively heated plasma would appear as a steady point source superposed on the pulsar.  Optical, polarimetric observations of the Crab pulsar have amply demonstrated the existence of unpulsed emission with flux $\\sim$ 1-2\\% of the pulse peak intensity \\cite{kanbach05}.  Preliminary estimates indicate that synchrotron emission from the resistively heated beams and the pair plasma  from the stripes engorged by the expanding sheets may well be the origin of this emission.  If so, the radiatively dark winds upstream of their termination in the surrounding nebulae may be subject to observational investigation.  \\begin{theacknowledgments}  During the slow and continuing unfolding of this project, I have particularly benefitted from discussions with S. Cowley, F. Coroniti and B. Schmekel.  My research efforts on this and other topics have been supported by NSF grant AST-0507813,  NASA grant NNG06G108G and DOE grant DE-FC02-06ER41453, all to the University of California, Berkeley; by the Department of Energy contract to the Stanford Linear Accelerator Center no. DE-AC3-76SF00515; and by the taxpayers of California.  \\end{theacknowledgments}"
    },
    "0710/0710.5782_arXiv.txt": {
        "abstract": "We study the evolution of the intracluster medium (ICM) with a uniformly analyzed sample of 70 galaxy clusters spanning $0.18 < z < 1.24$ and observed with \\Chandra. We find that X-ray luminosity and ICM mass at a fixed temperature evolve with redshift in a manner inconsistent with either the standard self-similar model of cluster formation or a model that assumes no evolution of cluster structure. Both luminosity and ICM mass evolve more slowly than the self-similar prediction, i.e., clusters have lower luminosity and ICM mass at fixed emission-weighted temperature than expected at higher redshifts. We find that evolution in these two observables can be modeled by a simple evolution in the cluster gas mass fraction, evolving as $(1+z)^{-0.39\\pm0.13}$ when measured using core-subtracted observables. Excluding cluster cores from measurements results in evolution more consistent with the self-similar model than when the entire cluster is used, indicating that the fraction of clusters with cool cores increases with time, or that cool cores become more developed over time in those clusters that have them; this is supported by direct study of the redshift dependence of central surface brightness, which increases in scatter and magnitude at low redshift. We find that isophotal size--temperature relations evolve differently according to which isophote is used, indicating that the central and outer regions of cluster ICM evolve differently. We show that  constraints on the evolution of the gas fraction and isophotal size--temperature relations constraints can be combined to measure cluster distances, and thus to constrain cosmological parameters in a way complementary to other techniques. Scatter in scaling relations is considerably reduced by using either core-subtracted quantities or three-parameter relations including the central surface brightness; in addition, there are indications that scatter decreases at higher redshift, suggesting that merging is not the dominant source of cluster structural variation. Our results provide constraints for simulations attempting to model cluster physics, indicate some difficulties for cosmological studies that assume constant cluster gas fractions, and point toward other potentially more robust  uses of clusters for cosmological applications. ",
        "introduction": "\\label{sec:intro}  Scaling relations among bulk properties of galaxy clusters provide a powerful means to test models of the large-scale structure and evolution of the universe. These correlations among properties such as X-ray luminosity, intracluster medium (ICM) mass, mean ICM temperature, and cluster virial mass reflect gravitational and non-gravitational processes involved in the formation of structure in an expanding universe. Scaling relations also provide the means to readily estimate masses of clusters from much more easily measured properties such as luminosity, an essential component of cosmological studies that use X-ray observations to determine the redshift evolution of the cluster mass function.  Simple models of cluster formation via gravitational collapse predict particular forms for the redshift evolution of cluster scaling relations \\citep[][]{kaiser86}. Adding additional cluster physics such as radiative cooling of the ICM, and energy injection by active galactic nuclei (AGN), supernovae, and star formation,  modifies these predictions \\citep[e.g.,][]{cavaliere98,ettori04b,muanwong06,kay07}. Observational studies of scaling relation evolution are required to properly constrain models of cluster evolution and to understand the effects of non-gravitational processes on the scaling relations that will be used to study cosmology. X-ray studies of the ICM are complementary to studies of the evolution of the cluster galaxy population \\citep[e.g.,][]{depropris99,lin06}, helping to constrain the overall evolution of cluster baryons and their distribution in various forms within clusters.  Several studies of X-ray scaling relation evolution  have been carried out in recent years \\citep[e.g.,][]{vikhlinin02,ettori04a,kotov05,maughan06,morandi07,branchesi07b}, but no clear consensus has emerged. In this paper we will address scaling relation evolution using a systematic analysis of a\t \\Chandra\\ sample of 70 clusters covering $0.18 < z < 1.24$, the largest sample yet used for this purpose.  Our study addresses two difficulties which may affect scaling evolution measurements. The first arises from the fact that radiative cooling of the ICM leads to the development of cool, dense (and hence very luminous) cores in many clusters; these relatively small cores bias cluster measurements such as X-ray temperature and luminosity to an extent that they are not representative of the overall cluster structure. This introduces significant scatter into scaling relations; indeed,  there is evidence that cool core clusters, which are traditionally regarded as ``relaxed,\" actually exhibit greater structural variation than non-cool core clusters, which are often thought to have recently undergone major mergers \\citep{ohara06}. Studies of scaling relations commonly attempt to ``correct\" for the impact of cool cores on cluster properties by one of several methods, such as simply leaving clusters with evidence for strong cool cores out of the sample \\citep[e.g.,][]{arnaud99}, or excising central regions within a fixed metric radius \\citep[e.g.,][]{morandi07} or a fraction of the virial radius \\citep[e.g.,][]{maughan07}, and perhaps ``correcting\" measured luminosity by some factor determined from a model of the cluster surface brightness distribution \\citep[e.g.,][]{vikhlinin02}. In this paper we measure temperatures with and without cores defined as fractions of the virial radius, and we also measure luminosities with and without the same core. By using relations both with and without core subtracted quantities, we can examine the effects that core development has on cluster scaling relation slopes and evolution.  The other issue usually faced by scaling relation studies is the use of scaling relation slopes and normalizations from low-redshift studies carried out with different instruments. The relatively small fields of view of \\Chandra\\ and XMM-{\\it Newton} make measurements of local samples quite challenging with those instruments; hence, studies using older X-ray instruments are used as references for $z=0$ relations. Unfortunately, differences in spectral and imaging results among X-ray instruments are well established, making such approaches subject to instrument-related systematics; indeed, even the same instrument has produced results differing by the author, as calibrations change and varying reduction and analysis methods are adopted. By using a large sample (70 clusters), we can avoid the use of outside references for scaling relation parameters or the direct inclusion of data from other samples, in favor of a single, homogeneously analyzed sample. While this approach is not entirely new---for example, \\citet{branchesi07b} studied evolution using their own 17 cluster sample both with and without the inclusion of data from other studies; and \\citet{morandi07} studied a homogeneously reduced 24 cluster sample---the size of our sample leads to significantly smaller uncertainties on scaling relation parameters than have otherwise been obtained.   In \\S~\\ref{sec:theory} we provide a brief overview of scaling relations and their predicted evolution, and in \\S~\\ref{sec:data} we explain our data reduction and measurement procedures. We test for scaling relation evolution with respect to expectations from the self-similar theory and from a scenario of no evolution in cluster parameters in \\S\\S~\\ref{sec:scaling_SS} and \\ref{sec:scaling_noev}, respectively, and provide an explanation for observed evolution in scaling relations via a simple evolution in the gas mass fraction \\S~\\ref{sec:ficm}. In \\S~\\ref{sec:size} we examine the evolution of isophotal size, and discuss the implications for studying cosmology using size measurements, and in \\S~\\ref{sec:scatter} we discuss the effectiveness of two different methods of reducing the scatter in measured scaling relations. In \\S~\\ref{sec:discuss} we compare our results to previous observations and simulation results, and discuss some implications of our findings. Finally, we list our conclusions in \\S~\\ref{sec:concl}.   We adopt the {\\it WMAP} + LRG $\\Lambda$CDM  cosmology from \\citet{spergel07}, which combines the third year {\\it WMAP} data with results from the SDSS luminous red galaxy survey \\citep{eisenstein05} to give $H_0 = 70.9$ ~km~s$^{-1}$~Mpc$^{-1}$,  $\\Omega_M=0.266$, and $\\Omega_\\Lambda = 0.734$. All uncertainties are 68\\% confidence, unless specified otherwise. ",
        "conclusions": "\\label{sec:concl}  We study the evolution of the ICM using X-ray scaling relations measured from a large, homogeneously analyzed sample of clusters spanning $0.2 \\lesssim z \\lesssim 1.2$. We use luminosity-- and ICM mass-temperature relations, including both relations with and without core subtracted quantities, to test scenarios of standard ``self-similar evolution\" and of ``no evolution\". We also study the evolution of isophotal size--temperature relations, for which (under certain assumptions) these two scenarios are identical. Finally, we compare the scatter in scaling relations after attempting to reduce cool core-induced scatter in two different ways. Our principal results appear below:  \\begin{enumerate}  \\item{Luminosity-- and ICM mass-temperature relations evolve less rapidly than expected in the self-similar evolution scenario; that is, clusters at higher redshifts have systematically lower luminosity and ICM mass at a given temperature than would be expected if clusters evolved self-similarly. The core subtracted relations have a combined consistency with the self-similar prediction of $<$0.1\\%; non-core subtracted relations are even more inconsistent with the self-similar prediction.}  \\item{The data are also inconsistent with the no evolution scenario, though not at as strongly as in the self-similar scenario. The core subtracted relations  evolve more rapidly than expected at higher redshift in this scenario, with combined probability of consistency with no evolution of 1\\%.}  \\item{The evolution in the \\LXCS--\\Txcs\\ and \\Micm--\\Txcs\\ relations is consistent with a simple evolution in gas fraction, with  evolution in \\ficm\\  at $>99$\\% confidence ($\\gamma_{f_g} = -0.39 \\pm 0.13$) in  the self-similar evolution   scenario when using core subtracted observables measured within \\rtfive.}  \\item{Isophotal size evolves with redshift at a rate that depends on the isophote used, reflecting evolution in the ICM spatial distribution in clusters. Evolution of isophotal size at a low isophote (i.e., well away from the core) is consistent with that expected given the measured \\ficm\\ evolution.}  \\item{Relations with core subtracted quantities in general have more positive evolution than relations with the cores included, suggesting that either the cool core fraction decreases with increasing redshift, or that the cool core fraction remains constant but the cores that do exist are weaker at high redshift. This is supported by direct observations of the redshift dependence of central surface brightness, a good indicator of cool core development; the scatter and magnitude of \\Io\\  increase at low redshift.}  \\item{Core subtracted relations generally have temperature dependences that are shallower than non-core subtracted relations, and thus are more consistent with the slopes predicted by the self-similar model for each scaling relation.}  \\item{The use of core subtracted quantities for scaling relations and the use of non-core subtracted quantities with the addition of a third parameter, the central surface brightness, both significantly reduce scaling relation scatter by compensating to some extent for cool core-related effects. }  \\item{Scatter in observables at fixed temperature appears to decrease with redshift. This could indicate an increase in the cool core fraction, an increase in the strength of cool cores in those clusters that have them, or both.}  \\end{enumerate}  Cluster simulations are still improving with regard to their ability to accurately model non-gravitational processes and thus to directly test specific models by comparison to observational data. However, our results of negative evolution with respect to self-similar expectations in \\LX\\ and \\Micm, and consequently in \\ficm, provide important constraints for future computational studies. Our findings provide new warnings with regard to the assumptions made when using \\ficm\\ measurements to study cosmology. It has long   been established that \\ficm\\ varies with radius inside clusters and   varies with cluster mass when measured within \\rfive\\ \\citep[e.g.,][]{david93,mohr99}.  Our results strongly suggest that \\ficm\\ varies with redshift as well.  Given the differences in behavior of collisionless   dark matter and the ICM (particularly the ICM's sensitivity to radiative   cooling and feedback from AGN and supernovae), perhaps it should not be   surprising that these components vary differentially with radius,   cluster mass, and even redshift.  At the same time, the combination of isophotal size measurements with measurements of the evolution of \\ficm\\ from \\LX\\ and \\Micm\\ relations provides a promising tool for measuring angular diameter distances. Our proposed cosmic consistency test would   allow one to use cluster structure and its evolution to constrain   cosmology in a manner complementary to more established techniques.   Finally, our results underscore the need to directly calibrate (or   self-calibrate) mass--observable scaling relations in large cluster   survey cosmology experiments.  Cluster structural evolution is   subject to a wide range of interesting physics, and determining that   mix reliably enough for even the most sophisticated simulations to   precisely predict cluster mass--observable scaling relations and   their evolution will remain enormously challenging for the   foreseeable future."
    },
    "0710/0710.0819_arXiv.txt": {
        "abstract": "In this paper we address important issues surrounding the choice of variables when performing a dynamical systems analysis of alternative theories of gravity. We discuss the advantages and disadvantages of compactifying the state space, and illustrate this using two examples. We first show how to define a compact state space for the class of LRS Bianchi type I models in $R^n$-gravity and compare to a non--compact expansion--normalised approach. In the second example we consider the flat Friedmann matter subspace of the previous example, and compare the compact analysis to studies where non-compact non--expansion--normalised variables were used. In both examples we comment on the existence of bouncing or recollapsing orbits as well as the existence of static models. ",
        "introduction": "Over the past decade a number of models have been proposed to account for the shortcomings of the standard model of cosmology based on General Relativity (GR). In most of these models, some modification is made to GR to explain recent observations such as the cosmic acceleration. These modifications include the adding of extra dimensions like in the brane world models \\cite{Brax03}, the adding of a minimally or non-minimally coupled scalar field \\cite{STGrav}, or modifications of the underlying field equations by either adding higher order corrections to the curvature \\cite{DEfR} or changing the equation of state \\cite{NLEoS}. In general, these modified theories of gravity have more complicated effective evolution equations, and it can be more difficult to find exact analytical solutions.  The implementation of the theory of dynamical systems (see for example \\cite{Dynamical,DSCosmo,Wainwright04}) has proven to be useful to gain a qualitative understanding of a given class of cosmological models. This dynamical systems approach does not require the knowledge of any exact solutions. However, the equilibrium points of the dynamical system correspond to the interesting cosmological solutions. This approach helps identifying exact solutions with special symmetries, which is particularly useful when studying complicated field equations.  In recent times, this approach has been used to investigate alternative theories of gravity such as Brans-Dicke theory \\cite{Burd88,Kolitch95,Kolitch96,Santos97,Holden98}, scalar-tensor theories \\cite{Foster98,Amendola90,Billyard99,Holden00,Gunzig00,Gunzig01,Saa01,Faraoni05,Carloni07,Agarwal07}, and higher order gravity \\cite{Carloni05,Leach06,Abdel07,Clifton05,Barrow06c,Goheer07a,Amendola07a,Amendola07b,Cognola07,Carloni07a}. It has also proven useful in theories with non-linear equations of state \\cite{Ananda06a,Ananda06b} and brane world models \\cite{Campos01a,Campos01b,Coley02a,Coley02b,Dunsby04,Goheer04}.  In this paper, we consider the various frameworks in which dynamical systems theory can be applied to cosmology.  In section 2 we discuss the characteristics of non-compact and compact state spaces in general. We point out the advantages of compactifying the state space, emphasising the aspect of static and bounce type solutions. In section 3, we proceed to give the specific example of LRS Bianchi I models in $R^n$--gravity and compare the results of \\cite{Goheer07a} and \\cite{Leach06}, where compact and non--compact expansion--normalised variables were used respectively. In section 4 we consider the flat Friedmann models in $R^n$--gravity and compare the compact formalism of \\cite{Goheer07a} to that of Clifton {\\it et al.} \\cite{Clifton05}, where non-compact non--expansion--normalised variables were used. We summarise our results in section 5, pointing out the discrepancies between the three approaches \\cite{Goheer07a}, \\cite{Leach06} and \\cite{Clifton05}. ",
        "conclusions": "In this work, we compared the use of compact and non-compact variables for a dynamical systems analysis of alternative theories of gravity. We first considered state spaces where expansion--normalised variables were used. These expansion--normalised variables were first introduced in the context of flat FLRW models in GR \\cite{Collins71}, where they define a compact state space. This compactness is desirable for determining the global behaviour of cosmological models and was one of the original reasons for introducing expansion--normalised variables in dynamical systems theory applied to cosmology. In \\cite{Goliath99} a method to compactify the state space of more general cosmologies was introduced, which  was successfully applied to a modified theory of gravity in \\cite{Goheer07a}.  We here showed that when non-compact expansion normalised--variables are used, we are restricted to expanding or contracting cosmological models only. Static solutions and bouncing or recollapsing type solutions lie at or approach infinity in this framework. This was illustrated in section 3, where we compared the expansion--normalised compact \\cite{Goheer07a} and non-compact \\cite{Leach06} state spaces of LRS Bianchi I models in $R^n$-gravity. While the works agree for the finite points in the non--compact analysis with $y\\neq 0$, discrepancies were found for the points with $y=0$ and the points at infinity. For $y=0$, which corresponds to the limit of vanishing Ricci curvature, differences with respect to the existence of solutions at the given points were observed. At infinity in the non--compact analysis \\cite{Leach06}, both the occurrence of \\equi points and the exact solutions at the \\equi points differs in places from the results obtained in \\cite{Goheer07a}. For example, we found that the asymptotic points in \\cite{Leach06} only have analogs in the compact analysis for specific values of the parameter $n$. The asymptotic \\equi point $\\mathcal{D}_\\infty$ does not have a counterpart in \\cite{Goheer07a} at all - it only appears to be an \\equi point in \\cite{Leach06} because the collapsing part of the state space is not included in this case.  We resolved these problems and found the approach used in the compact analysis \\cite{Goheer07a} more straightforward to analyse: Since there are no infinities in this framework, the opportunity to miss relevant information is reduced.  In section 4 we extended our comparison of formalisms by considering the non-compact non--expansion--normalised variables used in \\cite{Clifton05}. We compared the results of \\cite{Clifton05} with the compact analysis of \\cite{Goheer07a} for the flat Friedmann models. As in the first example, the points at infinity only have corresponding \\equi points in the compact analysis for specific values of $n$ and/or $w$. Unlike in \\cite{Leach06}, bounce and recollapse behaviours could be investigated in the framework of \\cite{Clifton05}, and we recover these results in the compact formalism.  We note that in both non-compact formalisms \\cite{Leach06} and \\cite{Clifton05}, the \\equi points at infinity are associated with a divergence in the respective dynamical systems time variable $\\tau$. When expansion--normalised variables are used, $\\tau$ diverges when $\\Th \\rightarrow 0$, while in framework of \\cite{Clifton05} $\\tau$ diverges when the matter density becomes negligible. This divergence however does not seem to effect the stability of the \\equi points at infinity, since the same results were found in the compact analysis.  Finally, we observe that in both non-compact analyses considered here, there are more \\equi points found than in the corresponding compact analysis. In the non-compact analysis of \\cite{Leach06} for example, five individual \\equi points and a line of points were found for the expanding LRS Bianchi I vacuum models, while in the corresponding compact analysis \\cite{Goheer07a} only one \\equi point and a line of \\equi points were found in the expanding subset. Similarly, for the flat Friedmann model, five pairs of (expanding and collapsing) \\equi points were found in \\cite{Clifton05}, while only three pairs of points were found in the compact analysis \\cite{Goheer07a}. The detailed comparison in sections 3 and 4 of this paper shows that there are two main reasons for the additional \\equi points in the non-compact analyses \\cite{Leach06} and \\cite{Clifton05}. The first is the duplication of \\equi points at infinity: we find that in the non--compact analysis \\cite{Clifton05} two copies of same point in the finite analysis (here points $5,6$) are created, and while static points exist only for special parameter values in \\cite{Goheer07a}, they have analogs at infinity in \\cite{Leach06} for all values of $n$ (see Tables \\ref{points:vac}, \\ref{points:mat} and \\ref{points:mat2}). Secondly, some points which are classified as \\equi points in the non-compact analysis, are not \\equi points in the compact analysis. In the examples studied here, this applies to points $\\mathcal{D}_\\infty$ in \\cite{Leach06} and $5,6$ in \\cite{Clifton05}, which correspond to $\\tilde{\\mathcal{M}}$ and $\\tilde{\\mathcal{N}}$ respectively.  Its worth noting that the analysis considered here is also applicable to more general theories of gravity such as $f(R)=R+\\alpha R^n$. In principle, compact variables can be constructed for such theories, but at the expense of a large number of sectors to be analysed (as can be see from equation (18) in \\cite{Barrow06c}).  In conclusion, we have shown that it is advantageous to compactify the state space whenever possible. The use of appropriately constructed compact variables allows for a clear and complete analysis including static, bouncing and recollapsing solutions and avoids the complications caused by \\equi points at infinity.  \\ack  This research was supported by the National Research Foundation (South Africa) and the Italian {\\it Ministero Degli Affari Esteri-DG per la Promozione e Cooperazione Culturale} under the joint Italy/ South Africa Science and Technology agreement. NG is funded by the Claude Leon Foundation. We thank the referees for their useful comments."
    },
    "0710/0710.3501_arXiv.txt": {
        "abstract": "The dynamics of cosmological gravitating system is governed by the Euler and the Poisson equations. Tiny fluctuations near the big bang singularity are amplified by gravitational instability into the observed structure today. Given the current distribution of galaxies and assuming initial homogeneity, dynamical reconstruction methods have been developed to derive  the cosmic density and velocity fields back in time. The reconstruction method described here is based a least action principle formulation of the dynamics of collisionless particle (representing galaxies). Two observational data sets will be considered. The first is the distribution of galaxies which is assumed to be an honest tracer of the mass density field of the dark matter. The second set is measurements of the peculiar velocities (deviations from pure Hubble flow) of galaxies. Given the first data set, the reconstruction method recovers the associated velocity field which can then be compared with the second data set. This comparison  constrains the nature of the dark matter and the relation between mass and light in the Universe. ",
        "introduction": "Introduction}  Cosmology is concerned with observing and modeling the universe on large scales: from our own Milky Way, other galaxies, galaxy clusters, super clusters up to the largest scales as probed by measurements of the cosmic microwave background radiation (CMB). These observations span a huge  range of scales and all strongly suggest that: 1) the dominant form of matter is dark (a factor of 6 in mass over the normal baryonic matter), 2) the clustering amplitude decreases with scale, and 3) structure forms by gravitational amplification of ting initial fluctuations. These are some of the main component of the standard paradigm in cosmology. Violation of any of them or all of them is consistent  with only a very limited set of observations, if any. Cosmology has had a great impact on other fields of physics and science in general. The shear existence of the  gravitationally dominant dark matter has stimulated scientists' (and others')  vivid imagination for a few decades now. Abundance and masses of non-standard particles have been constrained from the observed clustering pattern alone. In addition to gravity, hydrodynamical processes can greatly influence the formation and evolution of galaxies, groups and clusters of galaxies. Hydrodynamical effects, however, play a minor role in shaping the observed distribution of galaxies on scales a few times larger than the size of galaxy clusters. Therefore, gravitational instability theory directly relates the present-day large scale structure to the initial density field and provides the frame-work within which the observations are analyzed and interpreted.  Gravitational instability is a non-linear process.  Analytic solutions exist only for configurations with special symmetry, and approximate tools are limited to moderate density contrasts.  So, numerical methods are necessary for a full understanding of the observed large scale structure of the universe.  There are two complementary numerical approaches.  The first approach relies on $N$-body techniques designed to solve an initial value problem in which the evolution of a self-gravitating system of massive particles is determined by numerical integration of the Newtonian differential equations.  Combined with semi-analytic models of galaxy formation,  $N$-body simulations have become an essential tool for comparing the predictions of cosmological models with the observed properties of galaxies.  Because the exact initial conditions are unknown, comparisons between simulations and observations are mainly concerned with general statistical properties. The second approach aims at finding the past orbits of mass tracers (galaxies) from their observed present-day distribution. The orbits must be such that the initial spatial distribution is homogeneous.  This approach is very useful for direct comparisons between different types of observations of the large scale structure.  Most common are the velocity-velocity (hereafter {\\it v-v}) comparisons between the observed peculiar velocities of galaxies and and the velocity field inferred from the galaxy distribution in redshift surveys.  This types of analysis yield the cosmological mass density parameter $\\Omega_m$. Any systematic mismatch between the fields serves as an indication to the nature of galaxy formation and/or the origin of galaxy intrinsic scaling relations used to measure the distances, provided that errors in the calibration have been properly corrected for. This second approach also allows to perform back-in-time reconstructions of the density field on scales $\\sim 5 \\hmpc$  \\cite{Frisch2002}.    Finding the orbits that satisfy initial homogeneity and match the present-day distribution of mass tracers is a boundary value problem. This problem naturally lends itself to an application of Hamilton's variational principle where the orbits of the objects are found by searching for stationary variations of the action subject to the boundary conditions.  The use of the Principle of Least Action in a cosmological frame-work has been pioneered by Peebles (1989) \\cite{Peebles1989} and has long been restricted to small systems such as the Local Group \\cite{Peebles1994} and the Local Supercluster \\cite{Shaya}. Early applications to large galaxy redshift surveys have been hampered by the computational cost of handling the relatively large number of objects.  Subsequent numerical applications speeded up the method and allowed the reconstruction of the orbits of $\\sim 10^3$ particles (Shaya, Peebles \\& Tully 1995).  However, it was only recently that the improvement of the minimization techniques and the use of efficient gravity solvers made it possible to deal with more than $10^4$ objects\\cite{NusserBranchini}, comparable to the number of objects contained in the largest all-sky galaxy catalogs. ",
        "conclusions": "The rapid rotation of galactic disks revealed the existence of dark matter halos which engulf the luminous component. The measured virial motions of galaxies in clusters of galaxies also require the a gravitationally dominant dark component. Away from bound systems of galaxies and galaxy clusters, field galaxies show coherent flow pattern which deviates from a pure Hubble expansion. This coherent  velocity field is a direct probe of the large scale dark matter distribution in as much as rotational speeds and virial motions are a measure of the dark matter in galaxies and clusters. Indeed, the cosmic gravitational field responsible for the motions of galaxies, mainly  depends on the gravitationally dominant mass density field of the dark matter. The actual distribution of galaxies may well be quite different from the dark matter distribution. Recent analysis of the galaxy surveys, however, reveal a good match between the statistical properties of the galaxy distribution and the corresponding properties for the dark matter as inferred from numerical simulations of dark matter evolution in the universe. This is encouraging, but there may still be significant deviations between the distribution of the dark and luminous components, which are not reflected in statistical comparisons. The only way to detect such deviations is via direct detailed comparisons between the measured velocities of galaxies and velocities estimated from the  galaxy distribution. These comparisons have been done in the linear regime. The overall  agreement between the fields is impressive, but minor persisting mismatch is detected in some  regions in the local volume. It is possible that non-linear analysis a al the least action principle could mitigate some of the disagreement. This remains to be seen. The least action principle could also be used to recover the initial conditions \\cite{Mohayeehere},  allowing  us  to answer one of the fundamental question of whether or not  initial fluctuations were gaussian \\cite{ndy}.     The program is not without flaws. Many physical effects need to be addressed in detail. Most pressing is incorporating the assembly (or merging) history of galaxies. Galaxies reside in dark matter halos which form in a hierarchical manner from small to large. Thus our own Milky Way galaxy for example, is likely to have had  a major merging activity some 8 Gyr ago. All reconstruction methods assume that galaxies are point tracers of the mass density field and do not account for merging effects."
    },
    "0710/0710.1504_arXiv.txt": {
        "abstract": "We present very deep HST/ACS images of five QSO host galaxies, classified as undisturbed ellipticals in earlier studies. For four of the five objects, our images reveal strong signs of interaction such as tidal tails, shells, and other fine structure, suggesting that a large fraction of QSO host galaxies may have experienced a relatively recent merger event. Our preliminary results for a control sample of inactive elliptical galaxies do not reveal comparable fine structure. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0710/0710.3995_arXiv.txt": {
        "abstract": "We study how the determination of the Hubble constant from cosmological distance measures is affected by models of dark energy and vice versa. For this purpose, constraints on the Hubble constant and dark energy are investigated using the cosmological observations of cosmic microwave background, baryon acoustic oscillations and type Ia suprenovae. When one investigates dark energy, the Hubble constant is often a nuisance parameter, thus it is usually marginalized over.  On the other hand, when one focuses on the Hubble constant, simple dark energy models such as a cosmological constant and a constant equation of state are usually assumed.  Since we do not know the nature of dark energy yet, it is interesting to investigate the Hubble constant assuming some types of dark energy and see to what extent the constraint on the Hubble constant is affected by the assumption concerning dark energy.  We show that the constraint on the Hubble constant is not affected much by the assumption for dark energy.  We furthermore show that this holds true even if we remove the assumption that the universe is flat.  We also discuss how the prior on the Hubble constant affects the constraints on dark energy and/or the curvature of the universe. ",
        "introduction": "Recent precise cosmological observations provide us with a large amount of data to probe the evolution and the present state of the universe.  We now customarily extract the information from them by constraining cosmological parameters.  The Hubble constant, the expansion rate of the universe at present, is one of the most important cosmological parameters and measured in various ways.  In most studies which consider the constraint on the Hubble constant from cosmological observations, a cosmological constant is often assumed as dark energy. However, the nature of dark energy is not understood yet, which is one of the challenging problem in cosmology today, thus the constraints on the Hubble constant should also be considered assuming some possible dark energy models other than a cosmological constant. Since the energy density of dark energy can change some distance measures which have been used to constrain the Hubble constant, the assumption concerning dark energy is expected to have an influence on the constraint on it. In some works, the constraint on the Hubble constant is obtained without assuming a cosmological constant (e.g. Refs.~\\cite{Spergel:2006hy,Tegmark:2006az}) but a constant equation of state is still assumed. However, a lot of models of dark energy proposed so far have a time-varying equation of state and most of recent works of dark energy accommodate its time dependence in some way.  Thus, in light of these considerations and the importance of the Hubble constant, it is interesting to study constraints on the Hubble constant including time-evolving dark energy equations of state. This is one of the issues which we are going to investigate in this paper.  On the other hand, notice that the Hubble constant is often treated as a nuisance parameter over which we marginalize when we investigate dark energy since the parameters for dark energy themselves are those of interest in such a case. However, it should be mentioned that, when one studies the nature of dark energy, observations of distance measures such as the angular diameter distance which is relevant to the position of acoustic peaks in cosmic microwave background (CMB) power spectrum and the scale of baryon acoustic oscillations (BAO) are often used.  Since the Hubble constant affects such distance measures, it is expected that the prior on the Hubble constant can affect constraints on the nature of dark energy. Furthermore, the Hubble constant can also be determined with the cosmic distance ladder measurements independently from the cosmological distance measurements mentioned above. Hence investigating dark energy with some priors on the Hubble constant is also an interesting subject, which is discussed in this paper too.  To study the issues mentioned above, we study the constraints on the Hubble constant and dark energy using the observations of CMB, BAO and type Ia supernovae (SN).  For dark energy, we assume some types of time-varying equation of state as well as the case with a constant equation of state.  By assuming several priors on the Hubble constant and dark energy equation of state, we can see how the determination of one of them can affect that of the other.  In addition, we study them allowing a non-flat universe to make our analysis more general.  The structure of this paper is as follows. In the next section, we summarize our method for obtaining constraints from cosmological observations.  Some parametrizations of dark energy adopted in this paper are also briefly explained.  The data used for the analysis are mentioned there too.  In Sec.~\\ref{sec:result}, we present our results and discuss some implications of our results for the study of the Hubble constant, dark energy and the curvature of the universe.  In the final section, we summarize our results and give the conclusion.  In the appendix, we give some quantitative explanations how the combinations of CMB, BAO and SN can break degeneracies among the Hubble constant, matter density, dark energy parameters and the curvature of the universe. ",
        "conclusions": "We studied the constraints on the Hubble constant from CMB, BAO and SN assuming several types of dark energy parametrization.  Although the Hubble constant and dark energy are both important issues in cosmology today, these two subjects have not been investigated much simultaneously.  First we investigated the constraints in the $\\Omega_m$--$h$ plane assuming several dark energy parametrizations. The constraints on the Hubble constant from the combination of CMB, BAO and SN observations obtained under the different dark energy models and priors are summarized in Table~\\ref{tab:margDE} when a flat universe is assumed and in Table~\\ref{tab:margDE_Ok} when the flatness assumption is dropped respectively.  It is noticed that the constraints are not affected drastically by the dark energy model assumed and/or the assumption of the flatness of the universe.  It is rather more affected by the choice of the the SN data set: the Gold06 set gives slightly lower value of $h$ than the Davis07 set.  We can conservatively conclude that $H_0 < 59$ and $H_0 > 76$ are highly unlikely from these cosmological observations: these parameter regions are not allowed at 2$\\sigma$ level for any dark energy parametrization even if we do not restrict ourselves to a flat universe. Since the distance ladder estimations of $H_0$ have somewhat large systematic errors at present, we are not at the stage of arguing any possible discrepancies among these measurements now.  Nevertheless, it is worth mentioning in passing that the Sandage's central value $H_0 = 62$ \\cite{Sandage:2006cv} and the Macri's value $H_0 = 74$ \\cite{Macri:2006wm} are fairly close to the limit we have obtained here.  We have also investigated the constraints on some dark energy parameters assuming several priors on the Hubble constant.  The constraints are derived with and without the assumption of the flatness of the universe.  Using the present cosmological observations assuming no Hubble prior, we have found that a cosmological constant and a flat universe can fit all the data satisfactorily.  When we impose a prior on the Hubble constant, we have adopted Gaussian priors $h=0.72 \\pm 0.02$ and $h= 0.62 \\pm 0.02$.  The central values are those of Freedman's \\cite{Freedman:2000cf} and Sandage's \\cite{Sandage:2006cv} but the error  are taken to be a hypothetical value to give a meaningful effect on the dark energy parameter estimation.  We have found that, even with some limited options discussed in this paper for the assumptions with respect to the prior on the Hubble constant, a parametrization of the dark energy equation of state, the curvature of the universe, they can have a great influence on determining the nature of dark energy.  It should also be mentioned that the choice of the SN data set affects the allowed region. We demonstrate these points by taking a cosmological constant as a reference dark energy model because it is the simplest and most conventional model which can fit the observational data sets. For example, when we adopt the parametrization Eq.~\\eqref{eq:eos} in a flat universe, the prior $h= 0.62 \\pm 0.02$ reject a cosmological constant at 2$\\sigma$ level for the Gold06 data set, but the allowed region broadens to allow a cosmological constant if we do not adopt the flatness assumption or if we instead adopt the prior $h= 0.72 \\pm 0.02$. Moreover, if we adopt the parametrization Eq.~\\eqref{eq:eos2} with $z_\\ast = 0.5$ and the prior $h= 0.72 \\pm 0.02$, the Gold06 data set is in good agreement with a cosmological constant in a flat universe, but if we drop the flatness condition and marginalize over the curvature, a cosmological constant is disfavored at 2$\\sigma$ level. These examples imply that our understanding of the nature of dark energy can be varied by the assumptions on the Hubble constant, a parametrization of dark energy equation of state and the curvature of the universe.  Finally, we have investigated the constraints on the curvature of the universe assuming several types of dark energy and the priors on the Hubble constant.  In contrast to the constraints on the Hubble constant, we see the result depends on the dark energy parametrization we adopt. For the cases with a cosmological constant and a constant equation of state, the allowed region occupies larger area in a closed universe, whereas for the case with a time-varying equation of state parametrized as Eq.~\\eqref{eq:eos2} with $z_\\ast = 0.5$, the allowed region extends into a region of an open universe.  Since there is an obvious positive correlation between $h$ and $\\Omega_k$, the prior $h= 0.62 \\pm 0.02$ exaggerates the preference for a closed universe and the prior $h= 0.72 \\pm 0.02$ for an open universe.  This is what we see in Fig.~\\ref{fig:Ok_h_margOm_w0}.  We also reconfirmed that the preference of an open universe is caused by the equation of state for dark energy which is close to $0$ at earlier time.  This is because, we have found that $\\Omega_k$ as large as 0.2 is allowed in our previous papers Refs.~\\cite{Ichikawa:2006qb,Ichikawa:2006qt} under the weak prior $w_X \\le 0$, whereas $\\Omega_k$ is found to be well below 0.1 in this paper under the stronger prior $w_X \\le -1/3$.  Since dark energy is one of the most important problems in science today, a large number of works are focusing on dark energy itself. However, when one tries to probe the nature of dark energy, other cosmological parameters such as the Hubble constant, which was discussed in this paper, should necessarily be involved in various manners.  In light of precise measurements of cosmology that we are having now, the works from this kind of viewpoint should be done to check our understanding of cosmology and may also give insights to probe the present state and the evolution of the universe.    \\bigskip \\bigskip  \\noindent {\\bf Acknowledgments:} This work was supported in part by the Grant-in-Aid for Scientific Research from the Ministry of Education, Science, Sports, and Culture of Japan, No.\\,18840010 (K.I.) and No.\\,19740145 (T.T.).  \\bigskip \\bigskip    \\pagebreak  \\appendix  \\noindent {\\bf \\Large Appendix}"
    },
    "0710/0710.5488_arXiv.txt": {
        "abstract": "We present {\\it Spitzer} images of the relatively sparse, low luminosity young cluster L988e, as well as complementary near-infrared (NIR) and submillimeter images of the region.  The cluster is asymmetric, with the western region of the cluster embedded within the molecular cloud, and the slightly less dense eastern region to the east of, and on the edge of, the molecular cloud. With these data, as well as with extant H$\\alpha$ data of stars primarily found in the eastern region of the cluster, and a molecular $^{13}$CO gas emission map of the entire region, we investigate the distribution of forming young stars with respect to the cloud material, concentrating particularly on the differences and similarities between the exposed and embedded regions of the cluster. We also compare star formation in this region to that in denser, more luminous and more massive clusters already investigated in our comprehensive multi-wavelength study of young clusters within 1 kpc of the Sun. ",
        "introduction": "Lynds 988e (L988e =IRAS 21023+5002) was identified as one of six IRAS point sources (L988a-f) observed toward the L988 dark cloud, which is located on the edge of the Cygnus OB7 molecular cloud association \\citep{dobashi94} and which contains several bright, pre-main sequence objects. Bipolar molecular outflows were discovered \\citep{clark86} to be associated with the vicinity of L988e, as well as the vicinity of L988a and L988f, and there is also high velocity CO emission (blue-shifted lobe only) to the west of L988c. There are two LkH$\\alpha$ objects and related nebulosity seen both at visible and infrared wavelengths, LkH$\\alpha$ 324-SE and LkH$\\alpha$ 324, the brightest cluster members of a region of $\\sim$50 H$\\alpha$ emission-line stars \\citep{herbig06} on the edge of the heavy extinction associated with  L988e's parent molecular cloud.  Most of these emission-line stars lie to the east of the molecular cloud or on its eastern edge.   LkH$\\alpha$ 324 is a Herbig Ae/Be (HAeBe) star, close (in projection) to the IRAS coordinate, and LkH$\\alpha$ 324-SE (misidentified as LkH$\\alpha$ 324 in the the Herbig and Bell catalog \\citep{herbig88}) is probably also an HAeBe star \\citep{herbig06}. Both Classical T Tauri stars (CTTS) and Weak-line T Tauri stars (WTTS) are included in the emission-line stellar population.  Employing an average extinction of A$_V \\sim$2.5 mag for most of the population, and measured extinctions for a few objects, \\citet{herbig06} estimate median ages of 0.6 and 1.7 Myr for the emission-line stellar population from $V$ and $I$ photometry: for these estimates they assume evolutionary isochrones from \\citet{DM97}, and \\citet{siess00} respectively.  In this paper, we examine properties of the entire cluster, including both the emission-line stellar population and an obscured population to the west within the molecular cloud which we have identified here, and refer to as the embedded population.  At a distance of $\\sim$700~pc and a far infrared luminosity of $\\sim200~L_{\\odot}$, the cluster of young pre-main sequence stars associated with L988e is relatively sparse, and in the bottom quartile in far-infrared luminosity of the northern young cluster sample surveyed in $^{13}$CO and C$^{18}$O molecular gas \\citep{ridge03} (see Table~\\ref{littable}). Despite present interest in young cluster evolution as discussed, for example, in a recent review \\citep{ll03}, with the exception of papers by  \\citet{hodapp94} and  \\citet{herbig06}, the young cluster associated with L988e has been relatively unstudied at infrared wavelengths to date. Hereafter we refer to the young cluster as the L988e cluster, or just L988e.  In this paper, we utilize mid-infrared (MIR) images of L988e obtained with the Infrared Array Camera (IRAC) and 24 $\\mu$m images obtained with the Multi-Band Imaging Photometer (MIPS) on the {\\it Spitzer} Space Telescope\\footnote{This work is based [in part] on observations made with the Spitzer Space Telescope, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under a contract with NASA.} (see a composite three-color {\\it Spitzer} image of the L988e region, also including L988f and c in Fig.~\\ref{13mips_fig}; also see Fig.~\\ref{cover_fig} where the various fields of view are shown on a reproduction of the POSS V-band image of the region), as well as complementary near-infrared (NIR) ground-based images, to identify cluster members via their infrared excess emission from protostellar envelopes, and from circumstellar disks.  The {\\it spatial distribution} of identified protostars and disk objects in L988e is compared with the 850 $\\mu$m continuum maps of cold dust emission obtained with Submillimeter Common-User Bolometer Array (SCUBA) on the James Clerk Maxwell Telescope (JCMT), with dust extinction maps generated using NIR data, and with a $^{13}$CO map of the associated molecular clouds \\citep{ridge03}. These data were gathered as part of a coordinated multi-wavelength Young Stellar Cluster Survey, with the goal of providing the spatial distribution of a complete census of YSOs -- young stellar objects (all objects with an infrared excess, including stars with disks and protostars) in $\\sim$30 nearby embedded clusters chosen from a list of 73 young groups and clusters within 1 kpc of the Sun \\citep{porras03}. Program clusters are representative of quite diverse star-forming regions which vary in number of cluster members, member number density, luminosity, and environment.  This coordinated study, which includes NIR, mid-infrared (MIR), submillimeter and millimeter CO imaging of the regions, is designed to advance our understanding of clustered star formation.  A comprehensive paper examining YSO spatial distributions in this survey is in preparation: in a few instances, where unique data and/or characteristics distinguish the object, a separate paper is written.  Here, the availability of deep NIR images and a published catalog of H$\\alpha$ emission-line stars \\citep{herbig06} distinguishes this object from the others in the survey.  Several authors have described efforts to define IRAC-only \\citep{allen04}, or combined IRAC/MIPS \\citep{muzerolle04} protocols which differentiate between stars with dusty circumstellar accretion disks (class II), and protostars or objects dominated by dust emission from envelopes (class I).  All class I and II objects exhibit infrared excess over photospheric emission typical of more evolved young cluster members (diskless or optically thin disk class III) as well as of foreground/background normal stars. Since substantial and/or variable interstellar extinction can cause ambiguity in identification of these objects using IRAC-only methods \\citep{gutermuth05}, and the MIPS 24 $\\mu$m images may not detect the faintest YSOs, a more robust classification method is required.  \\citet{gutermuth05} used seven-band photometry from  $J$-band to IRAC 8.0 $\\mu$m to systematically determine de-reddened SED slopes.  Using this classification as a standard, methods requiring five-band detection at $J$ (1.25 $\\mu$m), $H$ (1.65 $\\mu$m), $K_{s}$ (2.15 $\\mu$m), [3.6] (3.6 $\\mu$m), [4.5] (4.5 $\\mu$m), or  four-band detection at $H$,$K_{s}$,[3.6], [4.5] also have been demonstrated in the same publication.  Here we primarily employ a combination of the five-band and four-band methods, but we also use slopes deduced from IRAC photometry in combination  with MIPS 24 $\\mu$m photometry, where available, to classify protostars and transition disk objects (objects with a pronounced inner disk gap) in L988e. Alternate methods are exploited for the young cluster IC 348 \\citep{lada06}: they utilize a simple power-law fit to the observed spectral energy distribution (SED) from the 4 IRAC bands to determine a slope $\\alpha_{IRAC}$, and as well study the SED from visual to far-infrared wavelengths for all cluster members as a function of spectral class. ",
        "conclusions": "1.  {\\it Spitzer} photometry of the asymmetric L988e cluster in combination with new NIR photometry reveals the presence of an embedded population associated with the molecular cloud (the cloud is defined by CO emission, cool dust emission, and A$_{Ks}$ maps of extinction). This is in addition to the previously studied exposed region of the cluster, associated with the optically studied H$\\alpha$ emission-line stars in a cloud cavity on the eastern edge, sculpted presumably by the HAeBe stars LkH$\\alpha$ 324 and LkH$\\alpha$ 324 SE.  2.  The disk fraction (class II/(class II + class III) of ONLY the emission-line objects, primarily located in the exposed region, is $\\sim$ 67\\%.  This is consistent with either of the two age estimates deduced from $V$ and $I$ photometry \\citep{herbig06}.  3.  The embedded population exhibits larger surface density than the exposed population of the cluster, and is fully enshrouded within L988e's natal molecular cloud.  This, as well as the higher $\\underline{mean}$ disk fraction calculated in the conventional way for all cluster members in the field (74\\% vs. 55\\%) points to the probability of a somewhat younger age for the embedded population as compared with the exposed population, although the uncertainties of these disk fractions overlap.  The exposed population of the cluster lies primarily in a region where the extinction A$_{Ks} < $ 0.5 mag.  On the other hand, the embedded population suffers a peak extinction of A$_{Ks}$ = 3, and extinctions A$_{Ks} >$ 0.5 throughout the molecular cloud.  4.  In addition to the asymmetric cluster, there exists a diffuse halo population of YSOs surrounding both the exposed and embedded regions of the cluster.  Although the halo is dominated by class IIs, class Is are also found in the diffuse halo.  Whether these halo objects are formed in situ is yet to be determined.  5.  Comparing these populations to six regions in the Young Cluster Program studied in a similar manner \\citep{gutermuth05}, we confirm the conclusion that the spatial distribution within the cluster is not strongly coupled to either the FIR luminosity or the numbers of cluster members.  The younger clusters more faithfully follow the cloud density structure in which the cluster is born, and exhibit higher surface density.  Therefore, L988e is older than the youngest clusters in the survey because the Eastern part of the cloud has had time to be excavated.  It is becoming increasingly obvious that class Is are found in the haloes around clusters in regions where the molecular column density is relatively low.  6.  {\\it Spitzer}/IRAC 3.6$\\mu$m, 5.8$\\mu$m and {\\it Spitzer}/MIPS 24$\\mu$m photometry reveal another embedded cluster associated with the object L988c."
    },
    "0710/0710.3117_arXiv.txt": {
        "abstract": "{We examine the applicability of effective temperature scales of several broad band colours to T Tauri stars (TTS). We take into account different colour systems as well as stellar parameters like metallicity and surface gravity which influence the conversion from colour indices or spectral type to effective temperature. \\\\ For a large sample of TTS, we derive temperatures from broad band colour indices and check if they are consistent in a statistical sense with temperatures inferred from spectral types. There are some scales (for $V-H$, $V-K$, $I-J$, $J-H$, and $J-K$) which indeed predict the same temperatures as the spectral types and therefore can be at least used to confirm effective temperatures.\\\\ Furthermore, we examine whether TTS with dynamically derived masses can be used for a test of evolutionary models and effective temperature calibrations. We compare the observed parameters of the eclipsing T Tauri binary V1642 Ori A to the predictions of evolutionary models in both the H-R and the Kiel diagram using temperatures derived with several colour index scales. We check whether the evolutionary models and the colour index scales are consistent with coevality and the dynamical masses of the binary components. It turns out that the Kiel diagram offers a stricter test than the H-R diagram. Only the evolutionary models of \\cite{BCAH98} with mixing length parameter $\\alpha=1.9$ and of \\cite{DM94,DM97} show consistent results in the Kiel diagram in combination with some conversion scales of \\cite{HBS00} and of \\cite{KH95}.% ",
        "introduction": "Only in a few cases the effective temperature of T Tauri stars (TTS) can be measured directly. In most other cases, conversions from spectral type to effective temperature are used \\citep[see for example][]{CK79,HSS94}. However, the spectral type of TTS is not always well determined due to their variability. For example, the classical TTS \\object{DL~Tau} is classified as GVe\\,--\\,K7Ve \\citep{CGCVS}. Thus, we recommend to verify the spectral type temperature with at least one colour index conversion scale. In this paper we want to carefully examine which colour indices are most suited on variable young stars with possible UV and IR excesses. \\par To illustrate the effects of the variability of TTS, two examples shall be given. The classical TTS (cTTS) most variable in $V$, \\object{DR~Tau}, has a variability in $(B-V)_\\mathrm{J}$ in the range $-0.76\\leq(B-V)_\\mathrm{J}\\leq1.11$ \\citep[for details, see][]{variability}. This difference corresponds roughly to the difference between an O-type and a K-type star. As expected, the variability of weak-line TTS (wTTS) is much smaller. The most variable wTTS, \\object{V410~Tau}, has a variability of $0.87\\leq(B-V)_\\mathrm{J}\\leq1.32$ \\citep[][]{variability} which corresponds roughly to the difference between an early K-type and a late K-type star. In contrast, the typical measurement error of a colour index is only about 0.02 magnitudes.\\par The applied colour index scales are summarised in Sect. \\ref{scales}; the way we calculated the colour index temperatures is explained in Sect. \\ref{Tfi}.\\par In Sect. \\ref{test_single}, a statistical test for a sample of apparently single TTS is done by comparing their colour index temperatures with their spectral type temperatures. We have not found any similar test in the literature.\\par If, in addition to the temperature, another stellar parameter like $\\log{g}$ or $L_\\mathrm{bol}$ is known, one can estimate the mass and the age of a star by means of  evolutionary models. If furthermore a mass determination independent from evolutionary models is known, one can check whether predicted and measured mass are consistent. For binaries, one can furthermore check whether both components are coeval. By using various conversion scales, we test combinations of evolutionary model and conversion scale (for details, see Sect. \\ref{test_evolve}). ",
        "conclusions": "We compiled several conversion scales which allow to derive effective temperatures from broad band colour indices, in order to examine their applicability to TTS.\\par These scales were first tested with a large sample of apparently single TTS. For this purpose we used four statistical criteria. As a result, we found ten scales for $V-H$ and $V-K$ as well as for the infrared colours $I_\\mathrm{C}-J$, $J-H$, and $J-K$ which are consistent with the temperatures derived from spectral type.\\par Furthermore we compared the colour index temperatures and the dynamically derived masses of the components of the eclipsing T Tauri binary V1642~Ori~A with predictions of pre-main sequence evolutionary models \\citep{BCAH98,DM94,DM97,PS99,SDF00,Yi03}, both in the HR diagram and the Kiel diagram.\\par In the HR diagram all evolutionary models give consistent results in combination with at least one conversion scale. As well, every conversion scale gives consistent results in combination with at least one evolutionary model. In the more decisive Kiel diagram, the evolutionary models of \\cite{BCAH98} with $\\alpha=1.9$, \\cite{DM94}, and \\cite{DM97} yield consistent results in combination with at least some conversion scales. In this diagram the scales HBS00d $(V-K)_\\mathrm{CIT}$, HBS00d $(V-K)_\\mathrm{JG}$, and ``combined $V-K$'' appear to be most suitable. But it is important to keep in mind that only combinations of evolutionary model and conversion scale are tested -- neither evolutionary models nor conversion scales alone. \\par As the Kiel diagram offers stricter constraints on evolutionary models than the H-R diagram, we recommend to use the Kiel diagram whenever possible."
    },
    "0710/0710.1781_arXiv.txt": {
        "abstract": "{} {We present a 900 sec, wide-field $U$ image of the inner region of the Andromeda galaxy obtained during the commissioning of the blue channel of the Large Binocular Camera mounted on the prime focus of the Large Binocular Telescope.} {Relative photometry and absolute astrometry of individual sources in the image was obtained along with morphological parameters aimed at discriminating between stars and extended sources, e.g. globular clusters. } {The image unveils the near-ultraviolet view of the inner ring of star formation recently discovered in the infrared by the Spitzer Space Telescope and shows in great detail the fine structure of the dust lanes associated with the galaxy inner spiral arms. The capabilities of the blue channel of the Large Binocular Camera at the Large Binocular Telescope (LBC-Blue) are probed by direct comparison with ultraviolet GALEX observations of the same region in M31. We discovered 6 new candidate stellar clusters in this high-background region of M31.  We also recovered 62 \\textit{bona-fide}  globulars and 62 previously known candidates from the Revised Bologna Catalogue of the M31 globular clusters, and firmly established the extended nature of 19 of them.} {} ",
        "introduction": "As the nearest giant spiral galaxy, Andromeda (M31) provides a unique opportunity to study the structure and evolution of a massive twin galaxy of the Milky Way (MW) from the ``outside'', and, by comparison with the MW, to address the question of variety in the evolutionary histories of massive spirals. In particular, the M31 Globular Cluster (GC) system provides the crucial bridge to connect the integrated-light-based methods that are used to study GCs in distant galaxies to the resolved-star-based methods that can be applied to the MW clusters (see e.g. Galleti et al. 2004, 2006a, hereafter G04, G06a).  Hierarchical structure formation models predict that the halo of a large galaxy is assembled through mergers of smaller subsystems. It has been suggested that M31 originated as an early merger of two or more relatively massive metal-rich progenitors (van den Bergh 2000, 2006). The prominent tidal stream in M31 (Ibata et al. 2001) extending several degrees from the centre of the galaxy (McConnachie et al. 2003), and the newly discovered arc-like over-density connecting M31 to its dwarf elliptical companion NGC 205 (McConnachie et al. 2004), are indeed the most spectacular interaction signatures presently known in the Local Group (LG) among the few nearby examples of tidal debris trails from galaxies currently merging.  The M31 outer disk is warped, and the halo contains numerous loops and ripples. Moreover, infrared photometry of the disk by IRAS (Habing et al. 1984) and by the Spitzer Space Telescope has revealed two rings of star formation off-centred from the galaxy nucleus Gordon et al. (2006). The two rings appear to be density waves propagating in the disk. Numerical simulations show that both rings may result from a companion galaxy plunging head-on through the centre of the M31 disk about 210 Myr ago (Block et al. 2006). The Galaxy Evolution Explorer (GALEX, Thilker et al. 2005) obtained far (1350-1750 $\\AA$, FUV) and near (1750-2750 $\\AA$, NUV) ultraviolet mosaic imaging of the M31 nucleus at a spatial resolution of 5$^{\\prime\\prime}$ in NUV. The M31 inner ring region is not well detected by GALEX. The Large Binocular Telescope (LBT), a forefront observational facility built and operated by an Italian-German-American collaboration, is at present the only telescope capable to unveil the near UV counterpart of the Spitzer inner ring in M31, thanks to the optimal combination of telescope size, spatial resolution (a factor 5 higher than in GALEX), field of view and sensitivity in the blue bands. The LBT is located on Mount Graham, Arizona, and in its final configuration will have two 8.4 m primary mirrors on a common mount, feeding various instruments (Salinari 1999,  Hill et al. 2006). Science commissioning at prime focus was carried out during October - December 2006 with a single mirror feeding the blue channel of the Large Binocular Camera (LBC-Blue; Ragazzoni et al. 2006, Giallongo et al. 2007) , a UV-optimized wide-field mosaic camera consisting of four CCD chips of 2048 $\\times$ 4608 pixels, with a mean pixel scale of $0.225\\arcsec/$pixel. The mosaic covers a field of view (FOV) of  24$^{\\prime} \\times 25^{\\prime}$; once the inter-chip gaps are taken into account, the effective sampled area is $\\sim 23^{\\prime} \\times 23^{\\prime}$.  In this paper we present results from a 900 sec wide-field $U$ band image of the M31 inner regions, reaching the limiting magnitude $U \\sim$ 25 mag. This image was obtained with the LBT during the commissioning phase, with the purpose of testing the $U$ capabilities of the blue channel of the LBC. The image is compared to the GALEX and Spitzer views of the same area. Results are also presented on new M31 candidate stellar clusters that were identified on the image. ",
        "conclusions": ""
    }
}